{"Identifier":"2021ApJ...914..136C__Choe_et_al._2017_Instance_1","Paragraph":"PYU+ formed from HC3N and H2NCO+, having at least 88 kJ mol\u22121 of internal energy, can redissociate or be stabilized by radiative relaxation. Collisional stabilization in the ISM is not effective due to an extremely low gas density. To compare rates of redissociation and radiative relaxation, the microcanonical rate constant for PYU+ \u2192 HC3N + H2NCO+ was calculated using the statistical Rice\u2013Ramsperger\u2013Kassel\u2013Marcus (RRKM) theory (Marcus & Rice 1951). The method has been described elsewhere (Baer & Hase 1996; Yim & Choe 2011; Choe et al. 2017; Choe & Kim 2019). At 10\u201350 K, energies in the range of 0.2\u20131 kJ mol\u22121 (5\/2RT) of thermal translational and rotational energies of HC3N and H2NCO+ are transferred to an vibrational internal energy of PYU+, which are added to 88 (87.8, more precisely) kJ mol\u22121 of the energy of PYU+ for redissociation. RRKM rate constants calculated at this energy range were (3\u20137) \u00d7 103 s\u22121 (Table 1). On the other hand, rate constants for typical radiative relaxations of vibrationally excited ions have been found to be lie within the range 101\u2013103 s\u22121 (Herbst 1985; Dunbar et al. 1996; Smith et al. 2001; Wakelam et al. 2010). This indicates that a considerable number of the excited PYU+ ions would be stabilized by comparing rate constants. At lower temperatures, stabilization would be more effective because rate constants for dissociation and radiative relaxation are expected to show positive and negative dependences on temperature, respectively (Wakelam et al. 2010). The stabilized PYU+ ions can undergo a further association reaction with NH3, H2O, or CH3OH. We will call this case A. These association reactions with NH3 and H2O, which finally lead to CyH+ and UrH+, respectively, can occur without activation barriers. The latter reaction is barrierless when assisted by a H2O molecule. On the other hand, the reaction of stabilized PYU+ ion with CH3OH to form ThH+ can hardly occur thermally at low temperatures due to its overall activation energy of 13 kJ mol\u22121.","Citation Text":["Choe et al. 2017"],"Citation Start End":[[530,546]]}
{"Identifier":"2022AandA...657A..96A__K\u00fcpper_et_al._2002_Instance_1","Paragraph":"The availability of two rotational transitions with different upper level energies allows us to constrain the rotational temperature of the propargyl radical in TMC-1. We built a rotation diagram using the velocity-integrated intensities given in Table 1, and we derive a rotational temperature of 9.9 \u00b1 1.5 K (see Fig. 2). We therefore confirm the assumption made by Ag\u00fandez et al. (2021a) that the rotational levels of CH2CCH are thermalized at the gas kinetic temperature of TMC-1, ~10 K (Feh\u00e9r et al. 2016). This fact is expected based on the low dipole moment of CH2CCH (0.150 D; K\u00fcpper et al. 2002), which implies low critical densities, probably a few 102 cm\u22123 (i.e., well below the volume density of H2 in TMC-1, afew 104 cm\u22123; Pratap et al. 1997; Cordiner et al. 2013). The column density derived from the rotation diagram for ortho CH2CCH is (8.2 \u00b1 1.7) \u00d7 1013 cm\u22122. A more precise determination of the column density can be obtained by fitting the observed spectra with synthetic spectra calculated under local thermodynamic equilibrium. For this calculation we adopted a rotational temperature of 9.9 K, as derived from the rotation diagram, a full width at half maximum of 0.72 km s\u22121 for the 20,2\u201310,1 lines and 0.57 km s\u22121 for the 50,5\u201340,4 lines, which are the arithmetic mean of the values derived for the hyperfine components of each transition (see Table 1), and assumed that the emission is distributed in the sky as a circle with a radius of 40\u2033, as observed for various hydrocarbons in TMC-1 (Foss\u00e9 et al. 2001). The observed spectra at 37.5 GHz and 93.6 GHz are well reproduced when adopting a column density of 7.5 \u00d7 1013 cm\u22122 (see Fig. 1). Assuming an ortho-to-para ratio of three, the column density of CH2CCH (including ortho and para) in TMC-1 is (1.0 \u00b1 0.2) \u00d7 1014 cm\u22122, which is slightly higher than the value derived previously by Ag\u00fandez et al. (2021a). The column density of the closed-shell counterpart CH3CCH in TMC-1 is (1.1\u20131.3) \u00d7 1014 cm\u22122 (Gratier et al. 2016; Cabezas et al. 2021). Therefore, in this study we confirm that the propargyl radical is thermalized to the gas kinetic temperature of TMC-1 and revise the abundance ratio CH2CCH\/CH3CCH to nearly one.","Citation Text":["K\u00fcpper et al. 2002"],"Citation Start End":[[585,603]]}
{"Identifier":"2016AandA...592A..19C__Johansson_et_al._2012a_Instance_1","Paragraph":"Figure 24 shows that the median Re increases for increasing mass (from ~5 Kpc to ~20 Kpc) and, given a stellar mass, our passive ETGs have median sizes smaller than those of the parent sample even by ~15% (at the highest masses). Furthermore, the entire Re distribution is extended to smaller radii in the case of passive ETGs, especially for log\u2009(M\/M\u2299) \u2273 11.5. Remembering that small differences in galaxy size imply large differences in stellar mass density, the derived trends suggest that the ETGs analyzed in this work should have formed from higher density progenitors which, in the hypothesis of no coeval mergers, do not increase their mass during the evolution. On the other hand, we cannot exclude the possibility that galaxies in the parent sample have experienced more dry mergers than our massive and passive galaxies, which have increased their size across cosmic time (Naab et al. 2009; Johansson et al. 2012a). The observed trend is also in agreement with some literature studies at low redshift which also suggest that more compact galaxies contain older stellar populations than larger ones (e.g. Saracco et al. 2009; Shankar & Bernardi 2009; Williams et al. 2010; Poggianti et al. 2013; McDermid et al. 2015). In general, this is qualitatively consistent with the high zF inferred by our analysis (z> 5) since, according to the cosmological evolution of the baryonic matter density, gas was denser at these cosmic epochs. These findings also agree with recent observations of compact systems (with various level of SF) at z ~ 2\u22123, which could be identified as local massive ETGs progenitors (e.g. Daddi et al. 2004; Cattaneo et al. 2013; Finkelstein et al. 2013; Marchesini et al. 2014; Nelson et al. 2014; Williams et al. 2015). Moreover, recent models of elliptical galaxy formation also predict the formation of compact and dense progenitors at high redshifts (e.g. Johansson et al. 2012a; Naab et al. 2014 and references therein). ","Citation Text":["Johansson et al. 2012a"],"Citation Start End":[[902,924]]}
{"Identifier":"2016MNRAS.462.1415C__Gonzalez-Perez_et_al._2014_Instance_1","Paragraph":"Over the last 15 yr, our understanding of how galaxies form and evolve has improved substantially. The combination of technological and theoretical progress has brought this field into a new era: advances in observational techniques (e.g. multi-object spectroscopy, efficient near-infrared CCDs) have enabled multiwavelength observations of large samples of galaxies out to the highest redshifts, while the steady rise of computational power and refinement of numerical techniques have fostered new approaches (e.g. semi-analytic models, hydro-dynamic simulations) to model the formation and evolution of galaxies. This progress has led to a general consensus about the main physical ingredients required to describe the evolution of the galaxy population (e.g. Gonzalez-Perez et al. 2014; Lu et al. 2014; Vogelsberger et al. 2014; Henriques et al. 2015; Schaye et al. 2015): collapse and hierarchical growth of dark matter haloes; accretion of baryons on to these haloes; conversion of baryons into stars; feedback of massive stars and active galactic nuclei (AGN) on star formation; supernova- and AGN-driven outflows of metal-enriched gas; infall of both pristine and metal-enriched gas on to galaxies. The large-scale environment can also affect galaxy properties, in particular, by providing quenching mechanisms (e.g. tidal or ram-pressure stripping, strangulation; e.g. Lagos et al. 2014; Rafieferantsoa et al. 2015), and through its influence on the merger rate (e.g. Lackner et al. 2012; Rafieferantsoa et al. 2015) and galactic spins (e.g. Hahn, Teyssier & Carollo 2010; Codis et al. 2012). Although these different ingredients are present in many galaxy formation models, we still lack a detailed quantification of their respective roles in shaping the properties of galaxies. This is because of the complexity inherent in galaxy physics, which combines gravity, radiation hydro-dynamics, magnetic fields and high-energy physics, acting on scales from less than a pc (e.g. for the formation of proto-stellar cores) to over a Mpc (e.g. for environmental effects). For this reason, \u2018first-principles\u2019 simulations of galaxy formation remain far beyond the reach of current computational capabilities. Instead, small-scale baryonic physics is generally subsumed into sub-grid prescriptions, which vary from model to model (e.g. Scannapieco et al. 2012; Haas et al. 2013a,b; Vogelsberger et al. 2013; Torrey et al. 2014; Crain et al. 2015). The appropriateness of such prescriptions, and hence, our ability to understand galaxy formation, must be assessed by comparing simulated and observed galaxy properties.","Citation Text":["Gonzalez-Perez et al. 2014"],"Citation Start End":[[762,788]]}
{"Identifier":"2018ApJ...854...26L___2015a_Instance_2","Paragraph":"The hot emission line of Fe xxi 1354.09 \u212b and the cool emission line of Si iv 1402.77 \u212b have been used in many spectroscopic studies to investigate chromospheric evaporation (e.g., Tian et al. 2014, 2015; Li et al. 2015b, 2017a, 2017b; Brosius et al. 2016; Zhang et al. 2016a, 2016b). It is widely accepted that the forbidden line of Fe xxi 1354.09 \u212b is a hot (log T \u223c 7.05) and broad emission line during solar flares (Doschek et al. 1975; Cheng et al. 1979; Mason et al. 1986; Innes et al. 2003a, 2003b). Meanwhile, IRIS spectroscopic observations show that Fe xxi 1354.09 \u212b is always blended with a number of cool and narrow emission lines, which are from neutral or singly ionized species. Those blended emission lines can be easily detected at the position of the flare ribbon, including known and unknown emission lines, such as the C i line at 1354.29 \u212b, the Fe ii lines at 1353.02 \u212b, 1354.01 \u212b, and 1354.75 \u212b, the Si ii lines at 1352.64 \u212b and 1353.72 \u212b, and the unidentified lines at 1353.32 \u212b and 1353.39 \u212b (e.g., Li et al. 2015a, 2016a; Polito et al. 2015, 2016; Tian et al. 2015, 2016; Young et al. 2015; Tian 2017). In order to extract the hot line of Fe xxi 1354.09 \u212b and the cool line of C i 1354.29 \u212b (log T \u223c 4.0; Huang et al. 2014), we apply a multi-Gaussian function superimposed on a linear background to fit the IRIS spectrum at the \u201cO i\u201d window (e.g., Li et al. 2015a, 2016a), which has been pre-processed (i.e., IRIS spectral image deformation, bad pixel despiking and wavelength calibration) with the standard routines in Solar Soft Ware (SSW; Freeland et al. 2000). In short, the line positions and widths of these blended emission lines are fixed or constrained, and their peak intensities are tied to isolated emission lines from similar species. More details can be found in our previous papers (Li et al. 2015a, 2016a). On the other hand, the cool line of Si iv 1402.77 \u212b (log T \u223c 4.8) at the \u201cSi iv\u201d window is relatively isolated, and it can be well fitted with a single-Gaussian function superimposed on a linear background (Li et al. 2014, 2017a). Using the relatively strong neutral lines (i.e., \u201cO i\u201d 1355.60 \u212b and \u201cS i\u201d 1401.51 \u212b), we also perform an absolute wavelength calibration for the spectra at the \u201cO i\u201d and \u201cSi iv\u201d windows, respectively (Tian et al. 2015; Tian 2017). Finally, the Doppler velocities of Fe xxi 1354.09 \u212b, C i 1354.29 \u212b, and Si iv 1402.77 \u212b are determined by fitting line centers removed from their rest wavelengths (Cheng & Ding 2016b; Guo et al. 2017; Li et al. 2017a). As the hot Fe xxi line is absent in the non-flaring spectrum, the rest wavelength for the Fe xxi line (i.e., 1354.09 \u212b) is determined by averaging the line centers of the Fe XXI profiles which were used in the previous IRIS observations (Brosius & Daw 2015; Polito et al. 2015, 2016; Sadykov et al. 2015; Tian et al. 2015; Young et al. 2015; Brosius et al. 2016; Lee et al. 2017), while the rest wavelengths for the C i and Si iv lines, i.e., 1354.29 \u212b and 1402.77 \u212b, respectively, are determined from their quiet-Sun spectra (Li et al. 2014, 2015a).","Citation Text":["Li et al. 2015a"],"Citation Start End":[[1373,1388]]}
{"Identifier":"2018AandA...611A..78G__Cruddace_et_al._2002_Instance_1","Paragraph":"In the 2D density distribution (Fig. 5, upper panel) we can distinguish only the primary structure. Taking into account the measured mass according to the fitted shear profile for this component, we compute the corresponding projected density profile, \u03a3(r), assuming a NFW distribution according to the equations given by Wright & Brainerd (2000). All of the obtained values are below the 3\u03c3 detection level (250 h70 M\u2299 pc\u22122) and only the inner region, corresponding to a radius of 60 kpc would present a density larger than the 2\u03c3 detection level (180 h70 M\u2299 pc\u22122). Thus, we consider that the mass of this structure is near the detection threshold of the 2D density contrast distribution. Also, we do not detect significant X-ray emission in the secondary structure region above the threshold adopted to build the brightness contours. In order to establish if this is due to the observing detection limit, we consider the lowest detected flux of RASS-based catalogs according to Piffaretti et al. (2011), which contains the lowest X-ray emission clusters identified using this data. This corresponds to a flux Flim = 1.5 \u00d7 10\u221212 erg s\u22121 cm\u22122, in the 0.1\u20132.4 keV band of the SGP catalog (Cruddace et al. 2002). Considering this flux and the redshift of the secondary structure, we obtain a limiting luminosity of 1.5 \u00d7 1042 \n\n\n$h^{-2}_{70}$\n\n\n\nh\n\n70\n\n\u22122\n\n\n\n erg s\u22121. Taking into account the obtained lensing mass, we compute the expected X-ray luminosity for the system according to Leauthaud et al. (2007) M_ LX relation, obtaining (3.7 \u00b1 2.4) \u00d7 1043 \n\n\n$h^{-2}_{70}$\n\n\n\nh\n\n70\n\n\u22122\n\n\n\n erg s\u22121, only 1.5\u03c3 above the detection limit. For this calculation we use the derived SIS mass since it has the lowest \u03c72 value and errors correspond to the propagation of the M_LX parameter errors. Therefore, according to the lensing estimated mass the secondary structure would be roughly at the detection limit of the X-ray observation. We considered using more sensitive X-ray data, therefore we checked in the available databases of Chandra and XMM; however this component is not present in the field-of-view of these surveys. Considering the X-ray luminosity emission upper limit (6.1 \u00d7 1043 \n\n\n$h^{-2}_{70}$\n\n\n\nh\n\n70\n\n\u22122\n\n\n\n erg s\u22121) the emission is 2.5\u03c3 above the detection limit given by Flim (Piffaretti et al. 2011). Therefore, it is important to highlight that a low X-ray emissivity could be explained by a low density of the intracluster gas, which might be produced by a past interaction between the structures.","Citation Text":["Cruddace et al. 2002"],"Citation Start End":[[1188,1208]]}
{"Identifier":"2021ApJ...922..149D__Paczynski_1992_Instance_1","Paragraph":"Since the ground-breaking observation of radio pulsars by Hewish et al. (1968), NSs remain one of the most studied astrophysical objects; they are assumed to be possible sources of high-energy emission. The typical values of the surface magnetic field as inferred from simple magnetic dipole models and spin-down rates are in the range 108\u20131013 G (Taylor et al. 1993; Alpar et al. 1982). Note that among radio pulsars, PSR J1847\u22120130 exhibits a strong magnetic field of B = 9.4 \u00d7 1413 G (McLaughlin et al. 2003). However, besides the X-ray luminosities observed from the anomalous X-ray pulsars (AXPs), the inferred periods of AXPs and soft-\u03b3 repeaters suggest that such NSs have even larger surface magnetic fields of 1014\u22121015 G (Duncan & Thompson 1992; Paczynski 1992; Thompson & Duncan 1996; Melatos 1999). NSs with such high surface magnetic fields are popularly known as magnetars. Several interesting studies (Usov 1992; Klu\u017aniak & Ruderman 1998; Wheeler et al. 2000; Starling et al. 2009; Cenko et al. 2010) have predicted that magnetars are the probable source of \u03b3-ray bursts, and they require a higher magnetic field such as 1016\u20131017 G to initiate Poynting-flux-dominated jets. Although until now only around 30 magnetars have been detected, researchers have speculated that these astrophysical objects may account for 10% of the NS population (Kouveliotou et al. 1998). For NSs, the effects of a strong magnetic field on the ultra-dense electron gases in their interiors have been studied in several papers (Canuto & Ventura 1977; Fushiki et al. 1989; Abrahams & Shapiro 1991; Fushiki et al. 1992; Roegnvaldsson et al. 1993). Studies of dense and strongly magnetized nuclear matter have also been carried out by Chakrabarty et al. (1997), Bandyopadhyay et al. (1998), Broderick et al. (2000), Suh & Mathews (2001), Harding & Lai (2006), Chen et al. (2007), Rabhi et al. (2008), and references therein. Finally, we mention that studies of NSs with different magnetic-field configurations, viz., toroidal, poloidal, or mixed were carried out by Bocquet et al. (1995), Cardall et al. (2001), and Pili et al. (2014).","Citation Text":["Paczynski 1992"],"Citation Start End":[[756,770]]}
{"Identifier":"2018MNRAS.478.2074B__S\u00e1nchez-Cruces_et_al._2018_Instance_1","Paragraph":"The interaction of the pulsar wind with the surrounding environment gives rise to a pulsar wind nebula (PWN) (Gaensler & Slane 2006; Bucciantini 2008; Olmi et al. 2016). A bubble of relativistic pairs and magnetic field, that shines with a non-thermal broad-band spectrum, extending from radio to X-rays and \u03b3-rays, via synchrotron and inverse Compton emission. The Crab nebula is the prototype of PWNe inside the parent supernova remnant (SNR) (Hester 2008). Given that the typical pulsar speed at birth, the so called kick velocity, is usually of the order of 100\u2013400\u2009km\u2009s\u22121 (Cordes & Chernoff 1998; Arzoumanian, Chernoff & Cordes 2002; Sartore et al. 2010; Verbunt, Igoshev & Cator 2017), much smaller than the expansion speed of young SNRs [typically a few thousands km s\u22121 (Truelove & McKee 1999; Hughes 1999, 2000; DeLaney & Rudnick 2003; Borkowski et al. 2013; Tsuji & Uchiyama 2016)], for the first few thousands of years, the PWN is going to remain confined within the parent SNR (van der Swaluw et al. 2003; van der Swaluw, Downes & Keegan 2004; Temim et al. 2015). However the SNR expansion speed will drop in time as the expansion proceeds (Cioffi, McKee & Bertschinger 1988; Leahy, Green & Tian 2014; S\u00e1nchez-Cruces et al. 2018), sweeping up more and more interstellar medium (ISM) material, such that on a typical time-scale of a few tens of thousands of years, the pulsar can escape from the SNR, and begin to interact directly with the ISM. Given the typical sound speed of the ISM, the pulsar motion is highly supersonic, and the ram pressure balance between the pulsar wind and the ISM flow gives rise to a cometary-like nebula (Wilkin 1996) known as bow-shock pulsar wind nebula (BSPWN), of shocked pulsar and ISM material (Bucciantini & Bandiera 2001; Bucciantini 2002b). If the ISM is cold and its neutral fraction is high, these nebulae can be observed in H \u03b1 emission (Kulkarni & Hester 1988; Cordes, Romani & Lundgren 1993; Bell et al. 1995; van Kerkwijk & Kulkarni 2001; Jones, Stappers & Gaensler 2002; Brownsberger & Romani 2014; Romani, Slane & Green 2017), due to charge exchange and collisional excitation, taking place in the shocked ISM downstream of the forward bow shock (Chevalier, Kirshner & Raymond 1980; Hester, Raymond & Blair 1994; Bucciantini & Bandiera 2001; Ghavamian et al. 2001), or alternatively in the UV (Rangelov et al. 2016), and IR (Wang et al. 2013). On the other hand the shocked pulsar material is expected to emit non-thermal synchrotron radiation, and indeed many such systems have been identified in recent years either in radio or in X-rays (Arzoumanian et al. 2004; Gaensler et al. 2004; Chatterjee et al. 2005; Gaensler 2005; Li, Lu & Li 2005; Yusef-Zadeh & Gaensler 2005; Kargaltsev et al. 2008, 2017; Misanovic, Pavlov & Garmire 2008; Hales et al. 2009; Ng et al. 2009; Ng et al. 2010; De Luca et al. 2011; Ng et al. 2012; Marelli et al. 2013; Jakobsen et al. 2014; Klingler et al. 2016; Posselt et al. 2017). There is also an interesting dichotomy between H \u03b1 emitting systems and synchrotron emitting ones, suggesting that energetic pulsars, more likely to drive bright synchrotron nebulae, might pre-ionize the incoming ISM suppressing line emission from neutral hydrogen.","Citation Text":["S\u00e1nchez-Cruces et al. 2018"],"Citation Start End":[[1214,1240]]}
{"Identifier":"2022MNRAS.515L..39Z__Trinquier_et_al._2009_Instance_1","Paragraph":"As we discussed in the introduction, different choices of age anchors for the 26Al\u201326Mg decay system, CAIs versus the D\u2019Orbigny angrite, result in different ages with a difference of 1\u20131.5\u2009Ma. The differences can be interpreted as a reflection of different initial 26Al\/27Al in the CAI accretion region, closest to the Sun (MacPherson, Wark & Armstrong 1988; Sossi et al. 2017) compared with the initial 26Al\/27Al in the angrite formation region. In fact, CAIs also show large mass-independent isotopic anomalies for multiple elements, e.g. O, Cr, and Ti (Trinquier et al. 2009; Krot et al. 2020), which suggest that they formed in different nebula environments compared to achondrites like angrites and EC 002. In Fig. 3 and Table S6, we show U\u2013Pb, 53Mn\u201353Cr, and 26Al\u201326Mg age comparisons for several achondrites. The data indicates that for all NC-like achondrites [\u03b554Cr values  0.3; Zhu et al. (2021a)], their U\u2013Pb ages, and 53Mn\u201353Cr and 26Al\u201326Mg ages anchored to D\u2019Orbigny are consistent. This consistency may indicate that the abundance of initial 26Al\/27Al value of NC achondrites at 4567\u2009Ma may have been lower in the non-carbonaceous region relative to the canonical initial Solar system value of 26Al\/27Al derived from CAIs (Schiller et al. 2015). However, for NWA 6704, which has a CR chondrite-like \u03b554Cr, its 53Mn\u201353Cr ages are more consistent if its 26Al\u201326Mg age is anchored to CAIs, which suggests that the 26Al abundance of the CC bodies in the outer Solar system may follow that of CAIs. This conclusion is also consistent with the olivine grains (with Al\/Mg ratio close to 0) in carbonaceous chondrites having the canonical initial \u03bc26Mg (\u223c\u221235\u2009ppm), indicating that the initial 26Al abundance in CCs at 4567\u2009Ma was similar to that of CAIs (Gregory et al. 2020). Thus, the age difference of 26Al\u201326Mg ages relative to 53Mn\u201353Cr ages can be caused by a different initial abundance of 26Al in the inner relative to the outer Solar system. We note that the U\u2013Pb age of NWA 2976 overlaps the Al\u2013Mg ages anchored to both CAIs and D\u2019Orbigny. This issue warrants more detailed assessment and discussion, e.g. the U\u2013Pb dating (Krestianinov, Datta & Amelin 2021); however, mineral fractions used for dating should avoid xenocrystic material that may have formed at different times and from different sources.","Citation Text":["Trinquier et al. 2009"],"Citation Start End":[[556,577]]}
{"Identifier":"2015MNRAS.449.2910C__Maller,_Dekel_&_Somerville_2002_Instance_1","Paragraph":"The second problem is that of the angular momentum distribution. Even if angular momentum is assumed to be conserved, one cannot explain the exponential form of disc galaxies. Using cold dark matter (CDM) numerical simulations (Bullock et al. 2001), found that if discs are formed from gas with an angular momentum distribution similar to that of dark matter, this results in excess mass near the centre, in comparison with an exponential disc. Specifically, the models predict too much material with low angular momentum, and this makes it particularly difficult to explain the origin of bulgeless dwarf galaxies (van den Bosch, Burkert & Swaters 2001). Responding to this difficulty Maller & Dekel (2002) proposed a model in which the angular momentum is built up by a sequence of mergers. While the acquisition of angular momentum in dissipationless N-body simulations has been well studied (Bullock et al. 2001; Maller, Dekel & Somerville 2002; Vitvitska et al. 2002; Avila-Reese et al. 2005; Bett et al. 2007, 2010; D'Onghia & Navarro 2007), our understanding of the extent to which gas behaves similarly has only recently started to be explored via computations (Chen, Jing & Yoshikaw 2003; Sharma & Steinmetz 2005; Brook et al. 2011; Kimm et al. 2011; Pichon et al. 2011; Sharma & Nath 2012). In the last decade galaxy redshift surveys (see e.g. Colless et al. 2001) have enabled us to map the universe on increasingly large scales, and to study systematically the properties of large-scale structures. These surveys show that the galaxies are not distributed evenly, but develop within a pattern of filaments and sheets which are separated by immense voids. The latter are defined observationally as large regions with very low galaxy number density, which occupy some 40\u2009per\u2009cent of the overall cosmic volume (Hoyle & Vogeley 2004). It is generally believed that the voids originate from the local minima of the primordial density fluctuations, and as their mean density is lower than that of their surroundings, they expand more quickly. If tidal shear plays a significant role in their development, it should not only cause them to depart from sphericity, but should entail their acquisition of spin angular momentum (Lee & Park 2006). The specific angular momentum of cold dark matter haloes in a \u039bCDM universe depends strongly on their merging histories (D\u2018Onghia & Burkert 2004). Haloes with a quiet merging history, dominated until the present epoch by minor mergers and accretion, acquire by tidal torques an average of only 3\u2009per\u2009cent of the angular momentum required for their rotational support. This is in direct conflict with the observational data for a sample of late-type bulgeless galaxies, which indicate that these galaxies reside in dark haloes with exceptionally high values of specific angular momentum. Although minor mergers and accretion preserve or slowly increase the specific angular momentum of dark haloes as a function of time, this mechanism is not efficient enough to explain the observed spin values, especially for late-type dwarf galaxies. Energetic feedback processes have been invoked to solve the problem implied by the loss of a large fraction of its specific angular momentum by the gas during infall. This is predicated on the assumption that dark haloes which host bulgeless galaxies acquire their mass via quiescent accretion. D\u2018Onghia & Burkert (2004) find an even more serious version of this problem: the specific angular momentum gained during the formation of these objects is not enough to provide rotational support even if no angular momentum is lost during gas infall.","Citation Text":["Maller, Dekel & Somerville 2002"],"Citation Start End":[[916,947]]}
{"Identifier":"2016MNRAS.462.1508G__Gaur_et_al._2012c_Instance_2","Paragraph":"The new sample of HSPs gave us an opportunity to see the optical IDVs of HSPs and compare its properties with optical IDVs of LSPs. We started a dedicated project to search for optical IDV in HSPs and after doing 62 nights of IDV observations of HSPs which gave us 144 LCs (41 in B band, 62 in R band, and 41 in B\u2212R colour) of five HSPs (Mrk 421, 1ES 1426+428, 1ES 1553+113, 1ES 1959+650, and 1ES 2344+514). Interestingly, we found that, four HSPs did not show any IDV (Gaur, Gupta & Wiita 2012a; Gaur et al. 2012b,c), but only one HSP 1ES 1426+428 for which we have the least observations have shown IDV in six LCs out of eight LCs (Gaur et al. 2012c). Our this pilot project gave us 6 IDV LCs out of 144 LCs searched for IDV i.e \u223c4 per cent LCs have shown IDV. We explained it by density inhomogeneities and bends in the bases of the jets by Kelvin\u2013Helmholtz instabilities (Romero, Cellone & Combi 1999). We gave an alternative explanation i.e. since in HSPs, the optical band lies below the SED peak, hence, we should see changes in the efficiency of acceleration of, and\/or in the rates at which energy is radiated by, the highest energy electrons available for synchrotron emission would have a more retarded effect on optical variability in HSPs (Gaur et al. 2012b). In LSPs, the optical band is dominated by highest energy electrons emitting synchrotron radiation and probably the X-ray emission is dominated by the comparatively lower energy electrons emitting the inverse Compton radiation, hence their X-ray variability is less pronounced than optical variability. If SED peak is really responsible for IDV properties, then we suspected that X-ray IDV LCs in LSPs should not show any IDV at all or show on rare occasions. With this motivation, here we present the X-ray IDV study of almost complete sample of 10 LSPs and 2 ISPs observed by XMM\u2013Newton since its launch and we found that the LSPs show very less IDV 2 out of 50 LC i.e. 4 per cent in X-ray bands. We have reported above the similar finding for HSPs in optical bands.","Citation Text":["Gaur et al. 2012c"],"Citation Start End":[[634,651]]}
{"Identifier":"2021ApJ...911...50U__Kotake_et_al._2004_Instance_1","Paragraph":"Many classical theories are presented exploring the dynamics of various rotating plasma instabilities, such as those of Vedenov et al. (1961) and Rosenbluth & Simon (1965). Initially studied by Velikov (1959), work in this area starts with the nondissipative Couette flow of an ideal conducting fluid between two differentially rotating cylinders. It is observed that the fluid is destabilized by applying an axial magnetic field. This work was confirmed by Chandrasekhar (1960). By now, research on magnetorotational instability (MRI) is well understood and reported in both experimentally (liquid metals; R\u00fcdiger et al. 2003; Ji et al. 2004) and theoretically based studies of astrophysical plasmas (Blaes & Balbus 1994; Hawley et al. 1995; Quataert et al. 2002). In astrophysical environments, MRI is now believed to be the main factor deriving plasma turbulence and also to account for the accretion process that feeds stars and black holes (Kotake et al. 2004). After its discovery, MRI remained outside the general theory of plasma instabilities for almost four decades, and was revisited with its application to astrophysical objects by considering the anomalous viscosity problem in the accretion disks surrounding astrophysical compact objects; see Balbus & Hawley (1991). MRI is supposed to be a destructive factor for the hydrodynamical Rayleigh\u2019s stability criterion, making it magnetohydrodynamically unstable and leading directly to orbital angular momentum transport and disk turbulence (Hawley & Balbus 1991, 1992; Balbus & Hawley 1998). A two-fluid model for MRI in differentially rotating plasma was presented by Ren et al. (2011) to describe the dynamical behavior of ions and electrons in both low- and high-frequency regimes. They have demonstrated the instability criteria that differ significantly from the one derived in the single-fluid model (Balbus & Hawley 1991). Mikhailovskii et al. (2008a) studied MRI in dusty plasmas using a single-fluid MHD model. They derived a generalized relation by including the effect of dust particulates, and due to the induced electric field the rotation of equilibrium plasma was modified and gave rise to an instability called a dust-induced rotational instability. In their preceding work (Mikhailovskii et al. 2008b) with an immobile dust grain assumption, a comprehensive analysis was carried out for the collective phenomenon of instability which provides a more detailed theory about MRI. The MRI in differential rotating dusty plasmas with dissipative effects was investigated by using local linear analysis with effect of immobile dust grains (Ren et al. 2009).","Citation Text":["Kotake et al. 2004"],"Citation Start End":[[946,964]]}
{"Identifier":"2020AandA...641A.123H__Mikal-Evans_et_al._(2019)_Instance_2","Paragraph":"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\u00d7 to 30\u00d7 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\u20131.1 \u03bcm), 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 \u03bcm feature is caused by VO emission. The secondary eclipse was observed at 2 \u03bcm (Kov\u00e1cs & Kov\u00e1cs 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 \u03bcm, prompting the authors to entirely discard their previous interpretation of emission caused by VO.","Citation Text":["Mikal-Evans et al. (2019)"],"Citation Start End":[[831,856]]}
{"Identifier":"2022MNRAS.517.2801W__Gallo_et_al._2014_Instance_1","Paragraph":"During the X-ray \u2018hard\u2019 state, the radio and X-ray behaviour of BHXBs is correlated and has been studied in depth for many sources using quasi-simultaneous observations (e.g. Corbel et al. 2003; Gallo, Fender & Pooley 2003; Coriat et al. 2011; Corbel et al. 2013), and is known as the radio:X-ray plane. It was thought that all XBs followed a relation in the form of $L_{\\rm Radio} \\propto L_{\\rm X-ray}^{0.6}$, based upon early observations of GX 339\u22124 (Hannikainen et al. 1998; Gallo et al. 2003; Corbel et al. 2003, 2013). This relation extends down to very low luminosities, i.e. into quiescence (Corbel et al. 2003, 2013; Plotkin et al. 2017; Tremou et al. 2020), and has been observed in other sources such as V404 Cygni (e.g. Corbel, Koerding & Kaaret 2008) and XTE J1118+480 (Gallo et al. 2014). However, further observations revealed the presence of another population of BHXBs which are less radio luminous than this relation, a so-called \u2018radio-quiet\u2019 branch, which followed $L_{\\rm R} \\propto L_{\\rm X}^{1.4}$, such as H 1743\u2212322 (e.g. Coriat et al. 2011; Williams et al. 2020). In some of these \u2018radio-quiet\u2019 objects, they are then found to re-join the \u2018radio-loud\u2019 branch when they go back into quiescence (Coriat et al. 2011; Carotenuto et al. 2021). The underlying cause of the split tracks (see Gallo, Miller & Fender 2012; Gallo et al. 2014; Gallo, Degenaar & van den Eijnden 2018, for a clustering analysis into the statistical robustness of this split) for BHXBs is not known, but it may be due to differences in the radiative efficiency of the accretion flow (Coriat et al. 2011; Koljonen & Russell 2019), an inclination effect of the source (Motta, Casella & Fender 2018), differences in the accretion disc contribution (Meyer-Hofmeister & Meyer 2014) or changes in the magnetic field (Casella & Pe\u2019er 2009). For the purposes of this paper, we will refer to the original $L_{\\rm R} \\propto L_{\\rm X}^{0.6}$ correlation sources as \u2018radio-loud\u2019 objects, and those that diverge on to the $L_{\\rm R} \\propto L_{\\rm X}^{1.4}$ track as \u2018radio-quiet\u2019 sources.","Citation Text":["Gallo et al. 2014"],"Citation Start End":[[784,801]]}
{"Identifier":"2016ApJ...818..180C__Carollo_et_al._1997_Instance_1","Paragraph":"The difference in the early-type fractions of star-forming and quenched satellites may be taken to indicate that quenching is associated with structural changes in the galaxies, that is, changes affecting their underlying mass distribution, such as substantial growth of stellar mass in the bulge (or inner disk) components. It is well established that quenched galaxies have radial surface brightness profiles that are steeper than those of star-forming galaxies and are well described by high S\u00e9rsic indices (e.g., Kormendy et al. 2009 and references therein), reflecting high central mass densities (e.g., Fang et al. 2013 and references therein). Also, since the pioneering work of Mihos & Hernquist (1994), inward flows of gaseous material have long been recognized to take place during galaxy mergers and through disk instabilities. Such inward gas flows are likely to contribute to the growth of stellar mass in galactic bulges (see, e.g., Courteau et al. 1996; MacArthur et al. 2003; Immeli et al. 2004; our own work, i.e., Carollo et al. 1997, 1998, 2001, 2007 and Carollo 1999, 2004; the comprehensive review by Kormendy & Kennicutt 2004 and references therein; Bournaud et al. 2011, and many others). This idea has been further developed in more recent theoretical studies (e.g., Dekel & Burkert 2014) and is suggested to be observed at high redshifts (e.g., Genzel et al. 2006; Cameron et al. 2011) and, at least for low-mass bulges, certainly observed in the local universe, where small bulges show signs of \u201crejuvenation\u201d of their stellar populations, with up to 10%\u201330% of their total mass consistent with having formed in the past few gigayears (Thomas & Davies 2006; Carollo et al. 2007). If mergers and disk instabilities were directly or indirectly connected with galactic quenching, it would not be implausible to draw a connection between quenching and the growth of bulges. Establishing the direction of causality in any such change would remain, however, very difficult. Quenching could be directly linked to the growth of the bulge (e.g., Martig et al. 2013; Genzel et al. 2014), or the quenching process itself could be associated with bulge growth, or the prominence of the bulge could be linked to a third property (e.g., galactic mass or the presence of an AGN) that itself controls the quenching of star formation.","Citation Text":["Carollo et al. 1997"],"Citation Start End":[[1032,1051]]}
{"Identifier":"2021ApJ...914L...1C__Rodr\u00edguez-Kamenetzky_et_al._2017_Instance_1","Paragraph":"Large-scale molecular outflows are also commonly associated with the earliest stages in the formation of massive (B- and O-type) stars. However, bipolar outflows observed in high-mass star-forming regions tend to be less collimated than those observed in low-mass regions (e.g., Arce et al. 2007). Recent radio observations have established that very young massive protostars are also commonly associated with compact free\u2013free emission, being the signpost of mass ejected near the protostar (e.g., Moscadelli et al. 2016; Purser et al. 2016, 2021; Rosero et al. 2016, 2019; Sanna et al. 2018; Obonyo et al. 2019). However, it is not yet clear whether this emission traces the bases of jets as well collimated as those observed in low-mass star-forming regions. To date, we know of only a handful of highly collimated jets from massive protostars extending over tens of thousands of au (e.g., Cep A HW 2: Curiel et al. 2006; HH 80\u201381: Rodr\u00edguez-Kamenetzky et al. 2017; IRAS 16547\u22124247: Rodr\u00edguez et al. 2005; G035.02+0.35: Sanna et al. 2019b; IRAS 16562\u22123959: Guzm\u00e1n et al. 2010; G016.59\u22120.05: Moscadelli et al. 2019), and, although evidence exists for magnetic fields playing an important role in their collimation (e.g., Carrasco-Gonz\u00e1lez et al. 2010; Surcis et al. 2014; Sanna et al. 2015; Maud et al. 2019), it is not yet clear whether collimation takes place at small distances from massive protostars. One possibility is that different mechanisms could be at work to produce jets around massive protostars, as they are expected to accrete at higher rates from much larger, massive and turbulent disks than those around low-mass protostars (e.g., Carrasco-Gonz\u00e1lez et al. 2012; Sanna et al. 2019a; A\u00f1ez-L\u00f3pez et al. 2020). Under these conditions, it could be very difficult to form strong and well-structured helical magnetic fields capable of highly collimate jets at au scales. Indeed, wide-angle or even spherical winds are also commonly found in high-mass star-forming regions (e.g., Torrelles et al. 1997, 2001, 2011; Moscadelli et al. 2007). Interestingly, it has been proposed that initially poorly collimated winds could be externally collimated at large (\u227310\u2013100 au) distances from the protostar by strong ambient medium pressure (e.g., Carrasco-Gonz\u00e1lez et al. 2015) or by a large-scale ordered magnetic field (e.g., Albertazzi et al. 2014). Therefore, there is the possibility that mass ejection from massive protostars could be mostly isotropic at au scales, and, under favorable conditions, it might become collimated at larger distances from the driving source.","Citation Text":["Rodr\u00edguez-Kamenetzky et al. 2017"],"Citation Start End":[[935,967]]}
{"Identifier":"2021AandA...653A..32D__Pozzetti_&_Mannucci_2000_Instance_1","Paragraph":"Since the expected number density of high-z massive quiescent galaxies is low, large fields with deep photometric coverage are required to identify these systems and reliably assess their SEDs. For this reason, we exploited the 2 deg2 COSMOS field. Sources with Ktot\u2004\u200422.5 were extracted from the catalogue of McCracken et al. (2010), limiting the selection to those satisfying the observed-frame BzK colour criterion for passive systems (Daddi et al. 2004). Targets formally classified as star-forming BzK (sBzK) with a signal-to-noise ratio S\/N\u2004\u20045 in the B and z bands were retained since these photometric candidates are degenerate with quiescent galaxies becoming fainter in these bands with increasing redshift and decreasing mass. Photometric redshifts specifically calibrated for high-z QGs were derived with EAZY (Brammer et al. 2008) as in Onodera et al. (2012) and Strazzullo et al. (2015). This calibration was based on the sample of 34 spectroscopically confirmed passive galaxies at 1.3\u2004\u2004zspec\u2004\u20042.1 observed with VLT\/VIMOS that later appeared in Gobat et al. (2017) and the sample of 18 passive galaxies at 1.4\u2004\u2004zspec\u2004\u20041.9 of Onodera et al. (2012) observed with Subaru\/MOIRCS. The calibrated zphot were used to select galaxies within 2.5\u2004\u2264\u2004zphot\u2004\u2264\u20043.5 and to remove objects with UVJ rest-frame colours inconsistent with passive evolution (Pozzetti & Mannucci 2000; Labb\u00e9 et al. 2005; Williams et al. 2009). SED fitting was performed using FAST (Kriek et al. 2009), allowing for constant and delayed exponentially declining SFHs. Optical dust attenuation was left free to vary up to AV\u2004=\u20045 mag assuming a Calzetti et al. (2000) attenuation law. Fits were repeated by adopting purely quiescent templates only. All passive UVJ candidates whose SED fits to optical-NIR photometry could not reject dusty star-forming solutions at high confidence were further discarded. Contamination from dusty star-forming galaxies was further minimised by removing objects with Spitzer\/MIPS 24 \u03bcm S\/N\u2004\u2265\u20044 detections in the Le Floc\u2019h et al. (2009) catalogue, except for galaxies with high-confidence passive SEDs, which are indicative of mid-infrared (MIR) emission caused by a central dusty AGN torus. This selection provided a total of 174 passive candidates with UVJ-quiescent colours (47 of which were pBzK in the original selection, plus 67 and 60 uncertain sBzK with and without significant MIPS detections, respectively). The maximum required magnitude to obtain HST\/G141 spectra with sufficient S\/N in order to secure spectroscopic confirmation within one to two orbits was assessed by simulating their grism spectra based on their best-fitting SED templates. This yielded a magnitude cut of HAB\u2004\u200422, which narrowed the sample down to 23 objects. A total of 10 galaxies were targeted for HST WFC3\/IR G141 near-IR observations: 9 randomly drawn galaxies with HAB\u2004\u200422 (M\u22c6\u2004>\u20041.1\u2005\u00d7\u20051011\u2006M\u2299) plus 1 robust candidate with HAB\u2004=\u200422.9\u2006(M\u22c6\u2004=\u20048\u2005\u00d7\u20051010\u2006M\u2299) selected to be the highest-z candidate based on its high-confidence passive SED. More details about the selection are available in Lustig et al. (2021).","Citation Text":["Pozzetti & Mannucci 2000"],"Citation Start End":[[1352,1376]]}
{"Identifier":"2019MNRAS.488.1012M___2011_Instance_1","Paragraph":"As a prediction of the \u039b cold dark matter (\u039bCDM) cosmological model (Springel, Frenk & White 2006), hierarchical clustering (White 1997; Steinmetz 2001) and major mergers (Toomre 1977) build elliptical galaxies and classical (elliptical-like) bulges of disc galaxies on a short time-scale through the complete and violent collapse of the protogalactic clouds. In this schema, mergers get less common through the expansion of the Universe and consequently the evolution of galaxies, which are getting now more and more isolated, gradually changing to a more secularly driven one under the influence of internal rather than external actors (Toomre 1977; Le F\u00e8vre et al. 2000; Conselice, Chapman & Windhorst 2003; Kormendy et al. 2010; Athanassoula, Machado & Rodionov 2013; Knapen 2013). This slow evolution of galaxies has been introduced as \u2018secular evolution\u2019 (Kormendy & Illingworth 1982; Kormendy & Kennicutt 2004; Kormendy 2008). However, mergers and accretion must also occur along the life of galaxies, even in low-density environments, and such external processes are known to push disc galaxies along the Hubble sequence towards Sa\/S0 (Aguerri, Balcells & Peletier 2001; Eliche-Moral et al. 2006, 2011). Much observational and numerical work has been devoted to discriminating between internal and external evolutionary mechanisms (for a review see Kormendy 2013). Bars in disc galaxies are considered as key drivers of the secular evolution by redistribution of the angular momentum, triggering of star formation, and the morphological transformation of galaxies. During the secular evolution phase, bars can significantly alter the structure and kinematics of disc galaxies and develop a complex rotation pattern within the central parts of galaxies, called \u2018cylindrical rotation\u2019 (Kormendy & Illingworth 1982). This feature is defined as almost constant stellar rotation speed within the bulge of a galaxy in the direction perpendicular to the disc plane, such that \u03b4$v$\/\u03b4|$z$| \u223c 0, where $v$ is the line-of-sight velocity and $z$ is the distance from the disc plane in edge-on projection. These feature has been confirmed in several observational studies of nearby barred galaxies observed in the edge-on view (e.g. Pence 1981; Kormendy 1983; Bettoni & Galletta 1997; Emsellem et al. 2001; M\u00e1rquez et al. 2003; P\u00e9rez, S\u00e1nchez-Bl\u00e1zquez & Zurita 2009; Williams et al. 2011; Molaeinezhad et al. 2016, 2017), including our Milky Way (see Barbuy, Chiappini & Gerhard 2018, and references therein). N-body simulations of barred galaxies confirm the link between the presence of a bar and a tendency to rotate on cylinders for stars within the bar, as a natural consequence of the bar buckling processes (see Combes & Sanders 1981; Athanassoula 2003, 2005; Martinez-Valpuesta & Shlosman 2004; Bureau & Athanassoula 2005). In this scenario, after the bar develops, it experiences a period of buckling instability that thickens it in the axial direction (e.g. Combes & Sanders 1981), giving birth to a boxy or peanut-shape bulgy structure in the central parts of the bar, termed boxy\/peanut bulge (B\/P, hereafter). From the galactic-dynamics point of view, bar formation makes the orbits in the bar more eccentric and more boxy and aligns their principal axes without greatly affecting the motion perpendicular to the plane; consequently, the cylindrical rotation follows naturally (Bureau & Freeman 1997). As demonstrated in numerous N-body simulation studies of barred galaxies, the resulting velocity field due to this orbital reconfiguration is compatible with the observed cylindrical rotation in B\/P bulges (see Molaeinezhad et al. 2016, and references therein).","Citation Text":["Eliche-Moral et al.","2011"],"Citation Start End":[[1179,1198],[1205,1209]]}
{"Identifier":"2019ApJ...883L..10A__Desai_et_al._2006_Instance_1","Paragraph":"Also as noted in Mason et al. (2008), He\/O, and to a lesser degree Ne\/O, ratios show a downward slope with respect to Fe\/C. This cannot by explained by variations of solar wind composition between fast and slow streams. The relative abundances of C and O are similar between slow and fast wind, and since the Fe, He, and Ne are more abundant in the slow than the fast solar wind, this would lead to a positive correlation (von Steiger et al. 2000). However, it is possible that the variations along the power-law slopes are, at least partially, due to a solar cycle effect on the solar wind composition, with the He\/O ratios and Ne\/O ratios varying throughout the solar cycle (see Desai et al. 2006; Mason et al. 2008). In order to investigate solar cycle phase dependencies in the abundance ratios, the right column of Figure 3 illustrates the ratios by solar cycle phase. To define the phase, the smoothed sunspot number was normalized to the maximum for each solar cycle and broken up into times when the normalized sunspot number was less than 1\/3 (solar minimum, blue triangles), greater than 2\/3 (solar maximum, red diamonds), or in between (ascending\/descending phase, green rectangles). As would be expected from the Fe\/O ratio in Figure 1, all of the times the Fe\/C ratios were greater than approximately 0.3 occurred during solar maximum. However, the ratios during solar minimum and the declining\/ascending phase are very similar. While the solar maximum Fe\/C ratios extend to much higher values than for CIRs during either solar minimum or the declining\/ascending phases, some of the solar maximum CIRs are also observed to have Fe\/C ratios well within the range of values observed for the other phases. This seems to suggest that the process(es) that is\/are energizing Fe preferential to C can be more efficient during solar maximum, but the process(es) does\/do not require all CIR-associated Fe\/C ratios during solar maximum to be necessarily elevated above that during the other phases. An example of this is the Fe\/O enhancement during the declining phase (Figure 1).","Citation Text":["Desai et al. 2006"],"Citation Start End":[[681,698]]}
{"Identifier":"2018MNRAS.478.3890B__Steidel_et_al._2010_Instance_1","Paragraph":"The circumgalactic medium (CGM) is the interface between cold flows from the intergalactic medium on to a galaxy, and hosts hot halo gas and material ejected from galaxies (for reviews, see Putman, Peek & Joung 2012; Tumlinson, Peeples & Werk 2017). With various processes in galaxy evolution consuming (e.g. star formation) and removing (e.g. winds) gas, the CGM is shaped by the processes internal to the galaxy. Early progress in the study of the CGM came from connecting the absorption lines in quasar (QSO) spectra with galaxies imaged in the foreground, tracing the extent and properties of the CGM gas as a function of the host galaxy\u2019s properties (e.g. Bergeron 1986; Bowen, Blades & Pettini 1995; Lanzetta et al. 1995; Adelberger et al. 2005; Chen et al. 2010; Steidel et al. 2010; Bordoloi et al. 2011; Prochaska et al. 2011; Turner et al. 2014). Building on these foundations, our understanding of the CGM has been significantly improved through the surveys with the Hubble Space Telescope (HST) Cosmic Origins Spectrograph (COS; Green et al. 2012). The first of several surveys of the CGM surrounding low-redshift galaxies was the COS-Halos survey (Tumlinson et al. 2013) which targetted the CGM around 44 \u223cL\u22c6 galaxies, demonstrating that the properties of the CGM differ depending on whether the central galaxy is passive or star forming [defined using a specific star formation rate cut of specific star formation rate (sSFR) $\\rm {=10^{-11} yr^{-1}}$; Tumlinson et al. 2011; Werk et al. 2013; Borthakur et al. 2016]. The COS-Halos team found a distinct lack of O\u2009vi around passive galaxies, while H\u2009i was found at the same strength around all galaxies (Tumlinson et al. 2011; Thom et al. 2012). Additionally, connections have been made between the CGM and properties of the host galaxy, including increased H\u2009i content of the CGM with larger interstellar medium (ISM) gas masses (COS-GASS; Borthakur et al. 2015), the presence of extended gas reservoirs around galaxies of all stellar mass (COS-Dwarfs; Bordoloi et al. 2014b), and enhanced metal content around starbursting hosts (COS-Burst; Borthakur et al. 2013; Heckman et al. 2017).","Citation Text":["Steidel et al. 2010"],"Citation Start End":[[770,789]]}
{"Identifier":"2020ApJ...895...81R__Chapman_et_al._2005_Instance_1","Paragraph":"DSFGs in the z > 5 tail are thought to be rare, but their level of rarity is subject to debate (e.g., B\u00e9thermin et al. 2015, 2017; Asboth et al. 2016; Ivison et al. 2016; see also Simpson et al. 2014, 2020; Dudzevi\u010di\u016bt\u0117 et al. 2020). A significant challenge in determining the space density of such sources is the difficulty in finding them in the first place. Given their distance, classical techniques combining optical and radio identifications have been largely unsuccessful due to the faintess or lack of detection at these wavelengths, commonly leading to misidentifications given the significant positional uncertainties of the classical sub\/millimeter single-dish surveys in which they are the most easily seen (e.g., Chapman et al. 2005; Cowie et al. 2009, and references therein). Also, due to the strong negative K correction at sub\/millimeter wavelengths (e.g., Blain et al. 2002), it remained challenging to pick out the most distant DSFGs among the much more numerous specimens at z  3.5. Over the past decade, many of these challenges were overcome through new observational capabilities and selection techniques, such as direct identifications based on interferometric observations of the dust continuum emission (e.g., Younger et al. 2007; Smol\u010di\u0107 et al. 2012; Simpson et al. 2014, 2015; Brisbin et al. 2017; Stach et al. 2018; see also earlier works by, e.g., Downes et al. 1999; Dannerbauer et al. 2002), redshift identifications through targeted molecular line scans (e.g., Wei\u00df et al. 2009; Riechers 2011), and target selection based on sub\/millimeter colors or flux limits (e.g., Cox et al. 2011; Riechers et al. 2013, 2017; Dowell et al. 2014; Vieira et al. 2010; Wei\u00df et al. 2013 ). Nevertheless, all of the current studies only provide incomplete censuses of the z > 5 DSFG population due to biases in the selection, limited sensitivity in the parent sub\/millimeter surveys, and incomplete redshift confirmations of existing samples.","Citation Text":["Chapman et al. 2005"],"Citation Start End":[[726,745]]}
{"Identifier":"2020MNRAS.496.1051A__Garc\u00eda-Rojas_2018_Instance_1","Paragraph":"PNe provide crucial information about the nucleosynthesis and chemical evolution of different elements. It is thought that PNe progenitors do not synthesize O and other \u03b1-elements as Ne, S, or Ar and, in principle, they can be considered as probes of the chemical composition of the ISM at the moment of the birth of the progenitors. Therefore, they can be used as tracers of the temporal evolution of the abundance gradients of the Galaxy. However, the recent results by Delgado-Inglada et al. (2015) indicate the production of O in C-rich PNe at near-solar metallicities that can increase the O\/H ratio up to 0.3\u2009dex. This result calls into question the suitability of O in PNe as tracer of the ISM composition. Other critical problems in the determination of the abundance gradients from PNe data is the large uncertainties in the distances of these objects (Frew, Parker & Boji\u010di\u0107 2016; Stanghellini & Haywood 2018; Garc\u00eda-Rojas 2018; Pe\u00f1a & Flores-Dur\u00e1n 2019), the different range of ages of their progenitors and the radial migration along the lifetime of the progenitors. Most of the determinations of the Galactic O\/H gradient based on data of young PNe \u2013 with ages between 1 and 4\u2009Gyr \u2013 provide slopes which tend to be flatter than our value (Liu et al. 2004; Maciel, Costa & Cavichia 2015; Stanghellini & Haywood 2018; Pe\u00f1a & Flores-Dur\u00e1n 2019), with typical values in the range between \u22120.019 and \u22120.048\u2009dex\u2009kpc\u22121 for RG between 0 and 16\u2009kpc. This behaviour has led some authors to propose the steeping of the O gradient with time (e.g. Stanghellini & Haywood 2018). However, other authors as Henry et al. (2004) or Maciel & Costa (2013) claim that the time variation of the O\/H gradient may be masked by the aforementioned uncertainties related to the abundance and gradient determination with PNe. As \u03b1-elements, Ne, S, and Ar should have abundance gradients similar to that of O\/H, but as it has been discussed in Section 4.2, their value may be affected by the ICF scheme used. For Ne and Ar, Pe\u00f1a & Flores-Dur\u00e1n (2019) reported a slope of \u22120.026 \u00b1 0.014\u2009dex\u2009kpc\u22121 and \u22120.010 \u00b1 0.006\u2009dex\u2009kpc\u22121, respectively, for a sample of young PNe, which are flatter than our slopes for these elements in H\u2009ii regions. Henry et al. (2004) studied a large sample of PNe providing the radial gradients for the different elements studied here but for a sample of objects covering progenitors of a wide range of ages. Those authors reported values of the slope for the gradients of Ne, S, Cl, and Ar abundances fairly similar to ours, from \u22120.030 to \u22120.048\u2009dex\u2009kpc\u22121. On the other hand, the results obtained by Stanghellini & Haywood (2018) for their compilation of PNe data from the literature show slopes of Ne and Ar which are somewhat flatter than ours, \u22120.018 to \u22120.030\u2009dex\u2009kpc\u22121, respectively. As we can see, the comparison with available data from PNe does not give conclusive evidence of a time variation of the abundance gradients in the Milky Way.","Citation Text":["Garc\u00eda-Rojas 2018"],"Citation Start End":[[920,937]]}
{"Identifier":"2022AandA...660A.108Z__P\u00e1l_2012_Instance_1","Paragraph":"The Chamaeleon I star-forming region, including CR Cha, was covered in Sectors 11 and 12 of the Transiting Exoplanet Survey Satellite (TESS, Ricker et al. 2015) in 2019. The observations started on 2019 April 22 and were completed on June 19, providing almost uninterrupted broadband optical photometric observations with a 30-min cadence. Details of the TESS data reduction and photometry can be found in Plachy et al. (2021) and P\u00e1l et al. (2020). We only summarize the main steps here. Photometry of the source was performed via differential image analysis using the ficonv and fiphot tools of the FITSH package (P\u00e1l 2012). This requires a reference frame, which we constructed as a median of 11 individual 64\u2005\u00d7\u200564 subframes obtained close to the middle of the observing sequence. The convolution-based approach used here exploits all of the information in the images by minimizing the difference and simultaneously correct for the various temporal aberrations (e.g., differential velocity aberration, variations in the point spread function, PSF, and pointing jitter corrections). A more detailed description of these methods can be found in Plachy et al. (2021) or P\u00e1l et al. (2020). This method is therefore equivalent to an ensemble analysis and requires basically no further post-processing of the light curves such as cotrending or similar types of decorrelartion methods. The photometry process also requires a reference flux to correct for various instrumental and intrinsic differences between the target and the reference frames. For this, we used the median of our SMARTS I-band photometry (see below) taken over the same time period as the TESS data. We obtained aperture photometry in the TESS images using an aperture radius of 2 pixels and a sky annulus between 5 and 10 pixels (the pixel scale is about 20\u2033). We note that we also inspected the most recent TESS observations, obtained between 2021 April 2 and June 24, and those data were reduced with the same procedure.","Citation Text":["P\u00e1l 2012"],"Citation Start End":[[616,624]]}
{"Identifier":"2022ApJ...924...56S__Lapi_et_al._2020_Instance_1","Paragraph":"The second ingredient is the probability distribution of stellar mass at given SFR and redshift:\n5\n\n\n\ndpdlogM\u22c6(M\u22c6\u2223\u03c8,z)\u221dM\u22c6M\u22c6M\u22c6,MS(\u03c8,z)M\u22c6,MSexp\u2212logM\u22c6\u2212logM\u22c6,MS\u03c8,z22\u03c3logM\u22c62M\u22c6\u2265M\u22c6,MS(\u03c8,z),\n\nwhere M\n\u22c6,MS(\u03c8, z) is the observed redshift-dependent galaxy main sequence with log-normal scatter \n\n\n\n\u03c3logM\u22c6\u22480.2\n\n dex (we adopt the determination by Speagle et al. 2014 for an anlytic fit, see their Equation (28)). The main sequence is a relationship between SFR and stellar mass followed by the majority of star-forming galaxies, apart from some outliers located above the average SFR at given stellar mass (see Daddi et al. 2007; Rodighiero et al. 2011, 2015; Sargent et al. 2012; Speagle et al. 2014; Whitaker et al. 2014; Schreiber et al. 2015; Caputi et al. 2017; Bisigello et al. 2018; Boogaard et al. 2018). The expression in Equation (5) holds for an approximately constant SFR history, which is indicated both by in situ galaxy formation scenarios (see Mancuso et al. 2016b; Pantoni et al. 2019; Lapi et al. 2020) and by observations of ETG progenitors (that have on overage slowly rising star formation history with typical duration of \u22721 Gyr; see Papovich et al. 2011; Smit et al. 2012; Moustakas et al. 2013; Steinhardt et al. 2014; Cassar\u00e1 et al. 2016; Citro et al. 2016) and late-type galaxies (that have on the average slowly declining star formation history over a long timescale of several gigayears; e.g., see Chiappini et al. 1997; Courteau et al. 2014; Pezzulli & Fraternali 2016; Grisoni et al. 2017). In this vein, off-main-sequence objects can be simply viewed as galaxies caught in an early evolutionary stage that are still accumulating their stellar mass (which grows almost linearly with time for a constant SFR), and are thus found to be preferentially located above the main sequence or, better, to the left of it. As time goes by and the stellar mass increases, the galaxy moves toward the average main-sequence relationship, around which it will spend most of its lifetime before being quenched due to gas exhaustion or feedback processes.","Citation Text":["Lapi et al. 2020"],"Citation Start End":[[991,1007]]}
{"Identifier":"2021ApJ...912..140S__Clement_et_al._2001_Instance_1","Paragraph":"In this paper, we will use a SAD map of the Large Magellanic Cloud (LMC) to derive the DTD of RR Lyrae stars\u2014pulsating horizontal branch stars with periods between 0.2 and 1 day (Smith 2004). We chose to test the DTD method on RR Lyrae for several reasons. First, the sample size of RR Lyrae in the LMC is quite large (see Section 2.1), allowing us to measure a DTD with high significance. Second, there is strong evidence that RR Lyrae are mostly ancient stars, older than 10 Gyr, given their pulsational properties (Smith 2004; Marconi et al. 2015) and abundance in old globular clusters (Clement et al. 2001; Soszy\u0144ski et al. 2014, 2016). Measuring an RR Lyrae DTD therefore provides a rigorous test of the DTD method for the recovery of progenitor age distributions as well as measurements of the star formation history of old resolved stellar populations. Lastly, a DTD analysis provides an opportunity to test stellar evolution models of RR Lyrae in a completely new way. While there is consensus on the interpretation of RR Lyrae as ancient stars, the lower limit on their ages has been somewhat unconstrained. For example, the absence of RR Lyrae in the SMC cluster Lindsey 1 (\n\n\n\n\nt\n\u2248\n9\n\n\n Gyr), compared to their presence in NGC 121 (\n\n\n\n\nt\n\n\u2248\n\n\n 10\u201311 Gyr), is generally cited as evidence of a lower limit of 10 Gyr for the progenitor age of RR Lyrae (Olszewski et al. 1996; Glatt et al. 2008). However, growing evidence of thin-disk RR Lyrae in the Milky Way raises questions about whether this limit might be lower, and whether an intermediate-age progenitor channel can exist (Layden 1995; Zinn et al. 2019; Prudil et al. 2020). Additionally uncertainties may exist in the evolutionary models of RR Lyrae stars themselves: since RR Lyrae are horizontal branch stars, their positions on the color\u2013magnitude diagram depend on an unknown combination of factors (commonly known as the second-parameter phenomenon) such as metallicity, age, mass loss on the red giant branch, stellar rotation, core structure, and chemical abundance (see Fusi Pecci & Bellazzini 1997; Catelan 2009; Dotter 2013, for reviews).","Citation Text":["Clement et al. 2001"],"Citation Start End":[[591,610]]}
{"Identifier":"2022ApJ...934..100S__Smak_1969_Instance_1","Paragraph":"\nAccretion disks around compact objects. Compact low-luminous massive objects are expected to lurk in the CGM of galaxies. From lone low-mass stars (e.g., Helmi 2020), to black holes (BHs) arising from stellar evolution (e.g., Fender et al. 2013), including remnants of Population III stars (e.g., Madau & Rees 2001; Filho & S\u00e1nchez Almeida 2018), and primordial BHs (e.g., Carr & Hawking 1974; Clesse & Garc\u00eda-Bellido 2015). The number density of these objects is assumed to be high; for example, there are claims that primordial BHs account for all the dark matter in the universe (Clesse & Garc\u00eda-Bellido 2015). Even if this is not the case (e.g., Carr & K\u00fchnel 2020), the expectations clearly overwhelm the density required to account for all the emission signals that we observe, of the order of one clump per central galaxy (Section 4.1). When these compact objects are surrounded by accretion disks, they should emit in H\u03b1, with the rotation of the disk giving rise to two-horn H\u03b1 profiles when observed with the appropriate viewing angle (Smak 1969). This kind of double-peak H\u03b1 emission is observed in X-ray binaries (Grundstrom et al. 2007; Zamanov et al. 2013; Casares & Torres 2018; Monageng et al. 2020) and cataclysmic variables (e.g., Zolotukhin & Chilingarian 2011), with peak separations of up to hundreds of km s\u22121 (Zamanov et al. 2013). At a completely different mass scale, double-peak H\u03b1 emission is sometimes observed in the broad-line region of AGN, where it is also supposed to trace a rotating disk (e.g., Eracleous & Halpern 1994; Ho 2008). The double-peak H\u03b1 emission observed around stellar-mass compact objects comes together with continuum emission (e.g., McSwain et al. 2010; Zamanov et al. 2019). There are hints of continuum emission in our stacks of observed spectra (Figure 8), but the question arises as to whether this very faint continuum is consistent or not with an accretion disk around a compact object. A number of arguments show that the observed lack of continuum emission does not discard the accretion disk scenario. First, large H\u03b1 equivalent widths (EWs) are sometimes observed in X-ray binaries. H\u03b1 EWs larger than 100 \u00c5 are not rare during quiescence (Fender et al. 2009; Casares & Torres 2018), with maxima reaching 2000 \u00c5 (Mu\u00f1oz-Darias et al. 2016), which implies line to continuum flux ratios enough to account for our observations (Figure 8). Second, H\u03b1 photons are to be produced by recombination of H atoms photoionized by the accretion disk (e.g., Matthews et al. 2015). Under particular circumstances, the nebular emission produced by a photoionizing source can emit an H\u03b1 line with very little underlying continuum. For example, young starbursts produce H\u03b1 with an EW in excess of 103 \u00c5 (Leitherer et al. 1999), and galaxy-integrated spectra with EW of hundreds of angstrom are not uncommon (e.g., Morales-Luis et al. 2011). The main physical ingredient for the photoionization source to produce little continuum is presenting a hard spectrum with an ionization flux greatly exceeding the flux in the optical. Accretion disks can easily match this requirement since they can be extremely hot with spectra peaking in the far-UV and X-ray.","Citation Text":["Smak 1969"],"Citation Start End":[[1047,1056]]}
{"Identifier":"2022ApJ...926...21B__Brown_et_al._2011_Instance_1","Paragraph":"Some studies have used the 2.5D mean-field dynamo approach to do so, extending solar mean-field dynamo models to other stellar spectral types (Chabrier & K\u00fcker 2006; Jouve et al. 2010; K\u00fcker et al. 2011; Kitchatinov et al. 2018, and references therein). While these studies are very helpful, most of them lack the full nonlinearity and genuine parametric dependence of 3D magnetohydrodynamic (MHD) simulations. Recent developments by Pipin (2021) are starting to overcome these limits and have extended the work of Rempel (2006) on the Sun to solar-type stars with various rotation rates. Nevertheless, with the arrival of more powerful supercomputers, other authors have used instead global 3D MHD simulations to model DR and stellar magnetism in the convection zone of solar-like stars (Glatzmaier & Gilman 1982; Miesch et al. 2000, 2006; Brun et al. 2004, 2011; Brown et al. 2008, 2010; Ghizaru et al. 2010; K\u00e4pyl\u00e4 et al. 2011, 2014; Gastine et al. 2014; Augustson et al. 2015; Karak et al. 2015). These studies pointed out the large magnetic temporal variability and the critical effect of stellar rotation and mass on magnetic field generation through dynamo mechanism, leading in some parameter regimes to configurations with cyclic activity (Gilman & Miller 1981; Gilman 1983; Glatzmaier 1985a; Brown et al. 2011; Racine et al. 2011; Augustson et al. 2013, 2015; K\u00e4pyl\u00e4 et al. 2013; Nelson et al. 2013; Beaudoin et al. 2016; Guerrero et al. 2016, 2019; Strugarek et al. 2017, 2018; Viviani et al. 2018, 2019; Warnecke 2018; Matilsky & Toomre 2020). Several studies pointed out the positive effect of a stable region underneath the convection zone (Parker 1993) on the efficient storage of intense toroidal field and the lengthening of the stellar dynamo cycle period (Glatzmaier 1985b; Browning et al. 2006; Lawson et al. 2015; Beaudoin et al. 2016; Guerrero et al. 2016, 2019; K\u00e4pyl\u00e4 et al. 2019; Bice & Toomre 2020). Over the last decade, significant progress has been made in successfully simulating large-scale mean flows and stellar activity cycle using different numerical codes and methods (Jones et al. 2011). This is quite reassuring that a global consensus is growing on the nature of solar-like star dynamos. It is common knowledge that there are still key transitions in Rossby number (at low and high values of this parameter) that need to be understood further, as well as what is the exact type of convective dynamos realized in solar-like stars as their global parameters are varied. This study continues this effort by doing an even broader systematic parametric study of solar-like star dynamos coupled to a stably stratified layer below than what have been published so far. It extends the work published in Varela et al. (2016) and Brun et al. (2017) with the MHD anelastic spherical harmonic code (ASH) (Brun et al. 2004). In particular, we wish to better characterize energy transfers and how much of a star\u2019s energy (luminosity) is converted into magnetic energy by nonlinear global convective dynamos over a wide range of Rossby numbers, generalizing to solar-like stars the work by Starr & Gilman (1966) and Rempel (2006).","Citation Text":["Brown et al. 2011"],"Citation Start End":[[1302,1319]]}
{"Identifier":"2019AandA...627A.172R__Rozitis_&_Green_(2013)_Instance_1","Paragraph":"For comparisons with the light curve YORP constraints, the YORP effect acting on Cuyo could be predicted by computing the total recoil forces and torques from reflected and thermally emitted photons from the asteroid surface using the ATPM. These calculations were made for both a smooth and rough surface, and were averaged over both the asteroid rotation and elliptical orbit (see Rozitis & Green 2012, 2013, for methodology). As demonstrated in Rozitis & Green (2012), the inclusion of rough-surface thermal-infrared beaming effects in the YORP predictions tends to dampen the YORP rotational acceleration on average but can add uncertainties of up to several tens of per cent if the roughness was varied across the surface. Since the light curve inversion produced convex shape models only, then shadowing and self-heating effects inside global-scale concavities (see Rozitis & Green 2013) were not possible to model. However, a study of non-convex shape models for fast two to four hour rotators in Rozitis & Green (2013) indicated that such asteroids have rather minimal levels of global-scale concavities, and the ~ 2.7 h rotation period of Cuyo implies that its shape could be similar. Furthermore, the Tangential-YORP effect, that is, a predicted rotational acceleration caused by temperature asymmetries within exposed rocks and boulders on the surface of an asteroid (Golubov & Krugly 2012), was also not included in the ATPM predictions. However, the very low thermal inertia value measured for Cuyo implies the absence of rocks and boulders on its surface of the quantity and size that are necessary to induce a significant Tangential-YORP component. As Cuyo is likely to be an S-type rubble-pile asteroid, a bulk density equivalent to that measured for the S-type rubble-pile asteroid (25143) Itokawa (Abe et al. 2006) of 2 g cm\u22123 was assumed for the YORP computations. Using the thermo-physical properties derived earlier, the ATPM predicts YORP rotational acceleration of (\u22126.39 \u00b1 0.96) \u00d7 10\u221210 rad day\u22122 for the nominal shape model. The uncertainty given here corresponds to the standard deviation of results when the degree of surface roughness israndomly varied across the surface of Cuyo (see Lowry et al. 2014, for details of the Monte Carlo methodology used). These values lie well within the light curve rotational acceleration constraints determined previously.","Citation Text":["Rozitis & Green","2013"],"Citation Start End":[[383,398],[405,409]]}
{"Identifier":"2016AandA...588A..44Y__Malinen_et_al._(2012)_Instance_1","Paragraph":"As described in Sect. 1, cloudshine is observed in the near-IR, usually in the three photometric bands J, H, and K. For instance, using the Wide Field CAMera of the United Kingdom InfraRed Telescope (WFCAM of UKIRT), Malinen et al. (2013) observed a 1\u00b0 \u00d7 1\u00b0 field in the Taurus Molecular Complex (d ~ 140 pc). This field is centred on the dense TMC-1N cloud (RA (J2000) 4h39m36s and Dec (J2000) + 26\u00b039\u203222\u2032\u2032) north of TMC-1. After smoothing the data to 40\u2033 resolution on a 8\u2033 pixel grid and removing background emission, Malinen et al. (2013) were able to extract J, H, and K coreshine median radial profiles across the TMC-1N filament (see their Fig. 12). This filament was previously observed with PACS and SPIRE instruments on board Herschel: using the colour temperature of the dust submm emission and assuming that the dust opacity varies as 0.1 cm2\/g (\u03bd\/ 1000 GHz)\u03b2 with \u03b2 = 2, Malinen et al. (2012) found that the density distribution of the cloud is accurately described by a Plummer-like function with \u03c1C = 4 \u00d7 104 cm-3, Rflat = 0.03 pc, and p = 3 (see Eq. (1)). To model this cloud, we start from this density distribution and the dust population mixtures presented in Sect. 4.1. The cloud is then illuminated by the ISRF+CM radiation field, extinguished or not by a layer of CM grains with \\hbox{$A_V^{\\rm ext} = 0.5$}AVext=0.5 to 3 (see Fig. 5). The resulting dust scattered emission maps4 are finally smoothed to the same resolution as the data presented by Malinen et al. (2013). The best fit is achieved for the CMM+AMMI populations illuminated by the ISRF+CM radiation field with \\hbox{$A_{{V}}^{\\rm ext} = 1.5^{{\\rm mag}}$}AVext=1.5mag. The results are presented in Fig. 11. As Malinen et al. (2013) subtracted the cloud background, it is not surprising to find \\hbox{$A_{{V}}^{\\rm ext} \\neq 0$}AVext\u22600; it probably reflects the fact that the incident radiation field is extinguished by the removed envelope before reaching the dense cloud. The peak intensity and the profile shapes are reproduced well for the three bands, even if the modelled profiles for the J and K-bands are slightly broader than the observed values. ","Citation Text":["Malinen et al. (2012)"],"Citation Start End":[[884,905]]}
{"Identifier":"2017ApJ...849..109P__Lee_et_al._2014_Instance_2","Paragraph":"Lastly, we model the evolution of the ejecta discussed in Section 2.2 into circumstellar profiles discussed in Section 2.3. We use our cosmic-ray hydrodynamics code, hereafter called ChN to model the evolution of the ejecta to an age of \n\n\n\n\n\n\nt\n\n\nSNR\n\n\n=\n400\n\n\n yr. ChN is a Lagrangian hydrodynamics code that includes a prescription for diffusive shock acceleration (DSA; Ellison et al. 2007; Lee et al. 2012). We have modified the code to include the effects of DSA on non-equilibrium ionization (Patnaude et al. 2009, 2010) and have coupled the code to SN ejecta models (Lee et al. 2014; Patnaude et al. 2015). We have also included radiative losses via forbidden line cooling (Lee et al. 2015). This effect will be important in the evolution of the SN shock with a nearby CSM shell, or if we choose to model the radiative shock that could form in the ejecta during early SN evolution (Nymark et al. 2006). However, we begin our simulations at an age of 5 yr, and over the lifetime of the simulation the shocks remain adiabatic, so we do not consider the radiative shock model presented in our previous work here. Since ChN couples nonlinear particle acceleration to the SNR shock dynamics, we are able to reproduce the broadband thermal and nonthermal emission (Ellison et al. 2010, 2012; Castro et al. 2012; Slane et al. 2014; Lee et al. 2013). The diffusive shock acceleration process is an integral part of ChN, and some injection of thermal particles into the acceleration process is always assumed. Here we set the injection parameter to the test particle limit, though we note that the interaction of a strong shock with a massive CSM shell or cloud will lead to enhanced particle acceleration (e.g., Ellison et al. 2012; Lee et al. 2014), and the differing CSM configurations, combined with the differing ejecta profiles and compositions, may result in differences in the broadband nonthermal emission. The study of nonthermal emission in evolving SNe is sufficiently broad that we defer its study to future papers.","Citation Text":["Lee et al. 2014"],"Citation Start End":[[1733,1748]]}
{"Identifier":"2019MNRAS.486.3741H__Machida_&_Hosokawa_2013_Instance_1","Paragraph":"As the initial state of star-forming clouds, a critical Bonnor\u2013Ebert (B.E.) density profile (Ebert 1955; Bonnor 1956) is adopted for each model. Note that the B.E. density profile or B.E. sphere is usually used as the initial condition of star-forming clouds (e.g. Matsumoto & Tomisaka 2004; Banerjee & Pudritz 2006; Machida et al. 2006a; Machida, Inutsuka & Matsumoto 2006b). The B.E. density profile is determined by the central density nc, 0 and isothermal temperature Tcl. The initial central density is set to $n_{\\rm c,0}=10^4\\, {\\rm cm}^{-3}$ for all models. The temperature Tcl of each cloud is determined as the result of a one-zone calculation (for details, see Susa et al. 2015, and Paper I), and the results Tcl are given in Table 1. The cloud radius rcl, which depends on the initial cloud temperature, is also given in Table 1. To promote cloud contraction, the density is set to 1.8 times to the critical B.E. density profile (Machida & Hosokawa 2013). The initial cloud mass for each model is also listed in Table 1. Although the initial clouds have different radii and masses with different metallicities, the ratio \u03b10 of thermal to gravitational energy, which significantly affects the cloud collapse (e.g. Miyama, Hayashi & Narita 1984; Tsuribe & Inutsuka 1999a,b), is the same for all models (\u03b10 = 0.47). In addition, the ratio of rotational to gravitational energy in the initial cloud is set to \u03b20 = 1.84 \u00d7 10\u22122 for all models (Goodman et al. 1993; Caselli et al. 2002). The initial magnetic field strength in each cloud is defined to satisfy \u03bc0 = 3 (Troland & Crutcher 2008; Crutcher et al. 2010; Ching et al. 2017). The parameter \u03bc0 is the mass-to-flux ratio of the initial cloud normalized by the critical value and is defined as \n(1)\r\n\\begin{eqnarray*}\r\n\\mu _0 &=& \\frac{\\left(M\/\\Phi \\right)}{\\left(M\/\\Phi \\right)_{\\rm cri}}, \r\n\\end{eqnarray*}\r\nwhere M and \u03a6 are the mass and magnetic flux of the initial cloud, respectively, and (M\/\u03a6)cri is the ratio of the critical values of these parameters, which is (M\/\u03a6)cri \u2261 (2\u03c0G1\/2)\u22121 (Nakano & Nakamura 1978). The direction of the magnetic field vector is parallel to the rotation vector (z-axis) in the initial cloud, in which a uniform magnetic field and rigid rotation are imposed.","Citation Text":["Machida & Hosokawa 2013"],"Citation Start End":[[942,965]]}
{"Identifier":"2020ApJ...903L..22T__Vuitton_et_al._2007_Instance_4","Paragraph":"While the Loison et al. (2015) CH3C3N model corroborates the upper atmospheric abundance of C4H3N inferred by Vuitton et al. (2007) from the T5 INMS measurements (a factor of 2 higher than those derived from T40 in Vuitton et al. 2019), a large disparity between the photochemical models (and within the ensemble of models produced by Loison et al. 2015) arises in the lower atmosphere due to the poorly constrained C4H3N branching ratios and reaction rate coefficients at temperatures appropriate for Titan. Aside from electron dissociative recombination of C4H3NH+ (Vuitton et al. 2007), neutral production of CH3C3N can occur in a few ways, as found through crossed beam experiments and theoretical and photochemical modeling studies (Huang et al. 1999; Balucani et al. 2000; Zhu et al. 2003; Wang et al. 2006; Loison et al. 2015). First, through the reactions of larger hydrocarbons with CN radicals,\n1\n\n\n\n\n\n\n\n2\n\n\n\n\n\nSimilarly, with CCN radicals following their formation through H + HCCN (Takayanagi et al. 1998; Osamura & Petrie 2004) and subsequent reactions with ethylene,\n3\n\n\n\n\n\nor through the chain beginning with acetylene,\n4\n\n\n\n\n\nWhile both reactions (3) and (4) are found to be equally likely by Loison et al. (2015), the production of CCN via H + HCCN is not well constrained, and the synthesis of CH3C3N through CN radicals (Equations (1) and (2)) are not included in their photochemical model. Additionally, cyanoallene may be produced through reactions (1)\u2013(4) instead of (or in addition to) methylcyanoacetylene. CH3C3N itself may form the protonated species, C4H3NH+, through reactions with the HCNH+ and C2H5+ ions producing HCN and C2H4, respectively (Vuitton et al. 2007). The other mechanism for forming C4H3NH+ is through the combination of HCN and l-C3H3+, though the reaction rate coefficient for this reaction and the abundance of l-C3H3+ are unknown (Vuitton et al. 2007). As such, the production and loss pathways for both C4H3NH+ and CH3C3N require further investigation.","Citation Text":["Vuitton et al. 2007"],"Citation Start End":[[1879,1898]]}
{"Identifier":"2017AandA...606A..50D__Coupeaud_et_al._(2011)_Instance_1","Paragraph":"Although in the MIR, the dust models are more absorbant than the measured samples, the opposite is true in the FIR. Table 2 gives the value of the MAC of the measured samples and of the modeled MAC at selected wavelengths in the 100 \u03bcm\u22121 mm range. In this range, the experimental MAC at 300 K is more than five times greater than the modeled MAC for all samples. At 10 K, the experimental MAC is lower than at 300 K and the factor of enhancement compared to the models is smaller, depending on the sample and on the wavelength, however it is always higher than two and usually of the order of four to five. The MAC value of the ferromagnesium silicate analogues in this study is very close the MAC of the pure Mg-rich samples from Demyk et al. (2017) and from Coupeaud et al. (2011). This shows that the enhancement of the measured MAC compared to the modeled MAC is not related to the iron content of the grains, that is, to differences in composition between the studied samples and the analogues used in the cosmic dust models. As discussed in detail by Demyk et al. (2017), an enhancement factor greater than two cannot be explained by the effect of grain size, grain shape, or by grain coagulation within the pellets. Grain size and grain shape effects are illustrated in Fig. 10 and the increase of the MAC due to large (micronic) spherical grains is negligible in the FIR. Coagulation might happen during the process of fabrication of the pellets and it is taken into account during the analysis of the experimental data following the method explained in Mennella et al. (1998) and based on the Bruggeman theory. More detailed treatment of dust coagulation by methods such as DDA have shown that it may increase the MAC by a factor of two at most (K\u00f6hler et al. 2011; Mackowski 2006; Min et al. 2016), that is, not enough to fully account for the discrepancy of the cosmic dust models and the experimental data. The enhancement of the emissivity of iron-rich analogues compared to the MAC of cosmic dust models commonly used for interpreting FIR\/submm dust emission observations is related to the disordered nature of the samples, to the number, distribution, and nature of defects in their structure at microscopic scale and to the existence of absorption processes added to the classical Debye model (we refer to Demyk et al. 2017, for more details). These additional absorption processes are more or less important depending on the structural state of the material at the microscopic scale. ","Citation Text":["Coupeaud et al. (2011)"],"Citation Start End":[[760,782]]}
{"Identifier":"2015ApJ...813..103M__Koss_et_al._2012_Instance_3","Paragraph":"The stochastic accretion of gas and galaxy merger-driven gas inflows are both known triggers of supermassive black hole (SMBH) growth and nuclear activity, but the relative contributions of each is still unclear. Simulations of galaxy mergers show that they drive gas to the centers of merger-remnant galaxies (e.g., Springel et al. 2005; Hopkins & Hernquist 2009), predicting that merger-driven SMBH mass growth occurs when the black hole nears the center of the merger remnant. Observations have shown that the AGN fraction does increase with decreasing distance between two merging galaxies (Ellison et al. 2011; Koss et al. 2012; Ellison et al. 2013), but this has not been well tested at the very centers of merger-remnant galaxies because of the observational difficulty of detecting and resolving two AGNs with separations 10 kpc. This is known as the \u201cdual AGN\u201d phase.4\n\n4\nThe separation scale expected for dual AGNs is between 0.1 and 10 kpc. The SMBHs in a merger stay at these separations for a few hundred megayears before evolving into a gravitationally bound, parsec-scale separation binary AGN system (Begelman et al. 1980).\n Hundreds of AGN pairs with >10 kpc separations have been discovered (Myers et al. 2008; Hennawi et al. 2010; Liu et al. 2011). However, there are only a few confirmed kiloparsec-scale dual AGNs (Junkkarinen et al. 2001; Komossa et al. 2003; Hudson et al. 2006; Rodriguez et al. 2006; Bianchi et al. 2008; Fu et al. 2011b; Koss et al. 2011, 2012; Mazzarella et al. 2012; Liu et al. 2013; Comerford et al. 2015). Dual AGNs are an intermediate evolutionary stage between first encounter and final coalescence of two merging gas-rich galaxies (e.g., Comerford et al. 2009; Liu et al. 2012), in which strong tidal interactions are more likely to influence the nuclear accretion and star formation in both galaxies (Barnes & Hernquist 1996). Indeed, galaxy merger simulations and observations clearly show that the dual AGN phase is the critical stage when SMBH growth and star formation activity are the most vigorous (e.g., Koss et al. 2012; Van Wassenhove et al. 2012; Blecha et al. 2013).","Citation Text":["Koss et al. 2012"],"Citation Start End":[[2061,2077]]}
{"Identifier":"2016MNRAS.455.2918H__iya,_Cenko_et_al._2012a_Instance_1","Paragraph":"A number of candidate TDEs were discovered by UV and X-ray surveys (e.g. NGC5905; Komossa & Greiner 1999, RX J1624+75; Grupe et al. 1999, A1795; Donato et al. 2014, Swift J164449.3+573451; Bloom et al. 2011; Burrows et al. 2011; Levan et al. 2011; Zauderer et al. 2011, Swift J0258.4+0516; Cenko et al. 2012b, and GALEX candidates D1\u20139, D3\u201313, and D23H-1; Gezari et al. 2008, 2009). Many of these cases did not show correspondingly strong optical emission, and in some cases they were discovered in sparse archival data, limiting their usefulness as probes of TDE physics. In recent years, however, a number of candidates have been found by high-cadence optical surveys, including the ASASSN-14ae (Holoien et al. 2014a), the Palomar Transient Factory (PTF10iya, Cenko et al. 2012a; PTF09ge, PTF09axc, and PTF09djl, Arcavi et al. 2014), Pan-STARRS (PS1-10jh, Gezari et al. 2012; PS1-11af, Chornock et al. 2014), and the Sloan Digital Sky Survey (SDSS; TDE 1 and TDE 2; van Velzen et al. 2011). These transients typically showed strong UV and optical emission but no associated X-ray emission, and were often better-studied than previous candidates, as they were discovered and followed up by survey projects searching their data in real-time. Optically discovered candidates have provided new opportunities to study these rare transients in greater detail than previously possible, and recent estimates based on optical TDE discoveries put the TDE rate at $\\dot{N}_{\\text{TDE}}=$ (1.5\u20132.0)$^{+2.7}_{-1.3}\\times 10^{-5}$ yr\u22121 per galaxy (van Velzen & Farrar 2014). However, there is still tension between this estimate and the rates predicted by modelling two-body scattering of stars in galactic nuclei (typically \u223c10\u22124 yr\u22121 per galaxy), which may in part be due to the fact that only a small fraction of TDEs are optically luminous (Stone & Metzger 2014; Metzger & Stone 2015). Selection effects from optical surveys may also play a role in this discrepancy, and a careful determination of TDE detection efficiency and completeness is needed to determine the degree to which the actual TDE rate may be higher than the rates inferred from optical surveys.","Citation Text":["Cenko et al. 2012a"],"Citation Start End":[[762,780]]}
{"Identifier":"2020ApJ...897...73M__Molkov_et_al._2019_Instance_1","Paragraph":"AstroSat data enabled us to derive the X-ray spectrum in the 0.8\u201370 keV energy band as shown in Figure 7. The spectrum was fitted reasonably well using two models. The first model was defined as an absorbed Fermi\u2013Dirac cutoff model along with a blackbody, an Fe emission line, and three Gaussian absorption lines introduced to model cyclotron scattering features and its two higher harmonics as observed in the spectrum. The CompTT model was used as the second model in combination with an Fe emission line and three Gaussian absorptions lines as defined in the first model. The CompTT model is generally used for neutron-star-based low-mass X-ray binaries such as Z-type and atoll sources with a relatively lower magnetic field (\n\n\n\n\n\n G) of the neutron star (Ferinelli et al. 2008). However, the model could also successfully define the spectrum of some of the pulsars in Be binaries, for example, Cep X-4 (Jaiswal & Naik 2015), 4U 1907+09 (Varun et al. 2019), and GRO J2058+42 (Molkov et al. 2019). The parameters derived from the two models are tabulated in Table 2. The first model estimated a blackbody temperature of 0.83 \u00b1 0.04 keV and detected the presence of a cyclotron resonance scattering feature and its harmonics. The Comptonization model, on the other hand, enabled us to determine the input photon Wien temperature of 0.52 \u00b1 0.02 keV, the plasma temperature of 8.22 \u00b1 0.10 keV, and the plasma optical depth of 5.21 \u00b1 0.12 for the phase-averaged spectra. The Wien temperature, per the CompTT model, originates far from the neutron star surface and closer to the inner accretion disk i.e., at the outer transition layer; hence, the Wien temperature is always found to be relatively lower than the neutron star blackbody temperature as it originates closer to the inner transition layer, i.e., near the surface of the neutron star (Ferinelli et al. 2008). For the CompTT model, bulk Comptonization occurs in the innermost part of the transition layer region, while thermal Comptonization is dominant in the outer transition layer and presumably within some extended region located above the accretion disk. Similar deviations were also observed in the case of Cep X-4 fitted with the CompTT model and a blackbody combined with the FDCUT model (Table 1 of Jaiswal & Naik 2015). However, the centroid energy of the cyclotron absorption features and its detected harmonics are found to be consistent within errors for the two models (Table 2).","Citation Text":["Molkov et al. 2019"],"Citation Start End":[[981,999]]}
{"Identifier":"2020AandA...644A..64D__Bron_et_al._2018_Instance_1","Paragraph":"We aim to measure the mass associated with each molecular structures described in the previous section. We defined spatial boundaries enclosing these structures and independently studied the spectral data corresponding to each subregion of the field. The spatial boxes defined for the cloudlet (A), ring-like structure (B), shocked clump (C), and shocked knot (D) are shown in Fig. 6, and the average spectra obtained in these boxes are presented in Figs. 8 and 7 for every line from CO and its isotopologs, which are available in our IRAM 30 m and APEX data cubes. The choice of the boundaries is based on our morphological classification, but we carefully checked that the brightest spectral features are coherent across the different boxes that we defined (coordinates of these boxes are given in Table A.1). We performed that selection manually, as the size of our sample is not large enough to apply statistical methods (e.g., clustering, see Bron et al. 2018). Based on the analysis of the emission of 12CO, 13CO and C18O lines, our description of these spectral features is the following:\n\n1.Cloudlet: toward box A (Fig. 7, left-panel), the line profile of 12CO and 13CO lines are similarly double peaked and best modeled by the sum of two Gaussian functions centered on the systemic velocities vLSR = \u22125.7 \u00b1 0.3 km s\u22121 (associated with the cloudlet) and vLSR = \u22123.3 \u00b1 0.1 km s\u22121 (associated with the ambient cloud).\n2.Ring-like structure: toward box B (Fig. 7, right-panel), the 12CO and 13CO lines are double peaked as well. The use of two Gaussian functions to model the line profile yields the systemic velocities vLSR = \u22125.6 \u00b1 0.2 km s\u22121 (associated with the ring-like structure) and vLSR = \u22123.3 \u00b1 0.1 km s\u22121 (associated with the ambient cloud). The Gaussian decomposition is very similar to that of the cloudlet, suggesting that the apparent ring-like structure might be incidental despite its remarkable features in the first moment map (Fig. 5).\n3.Shocked clump: considering the geometry of the SNR and the locally perpendicular direction of propagation of the SNR shockwave (van Dishoeck et al. 1993), the high-velocity emission arises from at least two shock waves, if not a collection of transverse shocks propagating along the molecular shell. In other words, the projection along the line of sight of several distinct shocked knots with distinct systematic velocities could contribute to the broadening of the 12CO lines. We measure vs \u2243 27 km s\u22121 and vs \u2243 21 km s\u22121 for the blueshifted and redshifted transverse shocks, respectively. Except for the J = 1\u20130 spectrum for which the emission of the ambient gas contributes to the average spectra, all spectra of 12CO lines exhibit a significant absorption feature around the vLSR of IC443G, suggesting that there is strong self-absorption of the emission lines. Evidence of line absorption is found in the velocity range \u22126 km s\u22121 vLSR  \u22122 km s\u22121, which is where we detect the spatially extended features associated with the NW-SE complex of molecular gas described by Lee et al. (2012). Hence, it is possible that the foreground cold molecular cloud is at the origin of the absorption of the 12CO J = 1\u20130 and J = 2\u20131 lines. A faint and thin emission line is detected around v = 6.5 km s\u22121 both in the 12CO and 13CO spectra. This signal is associated with the NE\u2013SW complex of molecular gas described in Sect. 3.2.\n4.Shocked knot: the shock signature of this line is distinct from the shocked clump. As hinted by the moments map (Fig. 5), its fainter high-velocity wings are displaced toward negative velocities. A self-absorption feature is also observed in this structure. Between v = \u22125.5 km s\u22121 and v = \u22122 km s\u22121 a bright andthin feature traces the ambient gas shown on the channel maps in Fig. 2 (second row from bottom; first, second and third panels from left).\n","Citation Text":["Bron et al. 2018"],"Citation Start End":[[948,964]]}
{"Identifier":"2020ApJ...888...72K__Alessio_et_al._2006_Instance_1","Paragraph":"In our calculation, we introduce two dust components: small dust grains and large dust grains. The minimum and maximum sizes are 0.0025 and 0.2 \u03bcm for small dust grains and 0.01 \u03bcm and 1 mm for large dust grains, respectively. The size distribution of both dust components is n(a) \u221d a\u22123.5, where n(a) is the number density per unit radius and a is the radius of dust grains. The opacity models for both small and large dust grains are described in Hashimoto et al. (2015) (see also Kim et al. 1994; Wood et al. 2002; Dong et al. 2012). We adopt Equation (16) as the surface density distribution of large dust grains and assume that the surface density of the small dust grains is 10 times smaller than that of the large dust grains (\u03a3small = 0.1\u03a3large; e.g., D\u2019Alessio et al. 2006). We note that the small-to-large dust mass ratio for the size distribution of n(a) \u221d a\u22123.5 is 1%. However, many practical calculations (e.g., Andrews et al. 2011; Dong et al. 2015b) use the size distribution and change the small-to-large dust mass ratio independently because the results are not very different even if the mass ratio is assumed as 10%. To obtain the density distributions of small and large dust grains, we set the scale height of them,\n17\n\n\n\n\n\n\n\n18\n\n\n\n\n\nwhere Hl and Hs are the scale heights of large and small dust grains, respectively. We assume that the scale height of small dust grains is similar to the gas scale height. The scale height of the large dust grains is assumed to be about three times smaller than that of the small dust grains.15\n\n15\nIf we assume that the gas surface density is 10 g cm\u22122 and turbulence strength \u03b1 = 10\u22122 at 70 au, we obtain that the scale height of 1 mm dust, which is the maximum size of the large dust, is about three times smaller than the gas scale height. In reality, this ratio will depend on the radius, but here we assume the constant ratio for simplicity.\n Using these scale heights, we evaluate the density of small and large dust grains, whose density distributions in the z-direction are proportional to \n\n\n\n\n\n. We note that the dependence of model parameters on the resulting disk structures is beyond the scope of this work and remains to be explored.","Citation Text":["D\u2019Alessio et al. 2006"],"Citation Start End":[[759,780]]}
{"Identifier":"2016MNRAS.462.4067P__Ma_&_Yan_2015_Instance_1","Paragraph":"To infer starburst and AGN luminosities, we use the multicomponent SED fitting technique described in Hatziminaoglou et al. (2008) and Hatziminaoglou, Fritz & Jarrett (2009). In the most general case, the method fits fluxes and their errors with three separate models: stellar, AGN, and starburst. Our sample however consists entirely of luminous type 1 quasars with extremely high SFRs. This allows us to make two key simplifications to the general method. First, it is almost certain that the stellar emission will be outshone by the AGN and\/or the starburst at all wavelengths, so we do not include a stellar template in the fit. Second, since our sources are type 1 quasars, it is reasonable to conclude that the optical to mid-infrared (MIR) coverage will be dominated by the AGN (Hatziminaoglou et al. 2005), and that the FIR coverage will be dominated by the starburst (e.g. Hatziminaoglou et al. 2010; Ma & Yan 2015; Harris et al. 2016). The infrared emission reprocessed by the AGN torus strongly depends on the dust geometry. In particular, the ratio between the outer and inner radius of the torus plays a major role in determining the AGN SED at longer wavelengths (in the models used in this paper the maximum value of this ratio is 150). As a result, the peak of the infrared emission of the AGN torus models lies at 10\u201330 \u03bcm (fig. 5 of Feltre et al. 2012) and the maximum width of the infrared bump (as defined in Feltre et al. 2012) for these models is \u223c50 \u03bcm. We therefore do not expect any significant contribution from the AGN emission at wavelengths larger than \u223c50 \u03bcm. Complementary, observationally motivated discussion on this point can be found in section 5.3 of Harris et al. (2016). We note the caveat though that, with our data, which do not cover the range between 22 and 250 \u03bcm, we cannot completely rule out significant FIR emission from the AGN. We thus use the shorter wavelength data to constrain the properties of the AGN, and, having done so, then use the SPIRE data to constrain the properties of the starburst. Two example SED fits are given in Fig. 3; one for a classical quasar and one for a BAL quasar.","Citation Text":["Ma & Yan 2015"],"Citation Start End":[[910,923]]}
{"Identifier":"2021MNRAS.508..637S__Pereira,_Bryan_&_Gill_2008_Instance_1","Paragraph":"The NLA analysis on these new satellite\/central disaggregated data vectors are shown in Fig. 15. As expected, the galaxy bias is consistent between the two, and constrained primarily by wgg, which is the same in the two data vectors. By construction the large scale fits to these data are each sensitive to a particular combination of two halo IA power spectra. Specifically $\\mathbf {D}^{c}$ is sensitive to ($P^{2h, s}_{\\rm GI}$, $P^{2h, ss}_{\\rm II}$), and $\\mathbf {D}^{s}$ probes ($P^{2h, c}_{\\rm GI}$, $P^{2h, cc}_{\\rm II}$). Again, we assume that on two halo scales, the satellite\/central composition of the density tracer is not relevant. Notably, the amplitude of large-scale central alignments in Fig. 15 is stronger than that of satellites by a factor of \u223c2, at the level of a few \u03c3. The subject of satellite alignments has been discussed quite extensively in the literature, and the overall picture fits with our results here. A number of theoretical studies point to satellite IAs being dominated by tidal torque induced radial alignments within their haloes (Faltenbacher et al. 2007; Knebe et al. 2008; Pereira, Bryan & Gill 2008), which scale rapidly with separation, and tend to wash out on very large scales. There is also now evidence from various observations on both cluster and galaxy scales supporting the same picture (Sif\u00f3n et al. 2015; Singh et al. 2015; Huang et al. 2018). This, again, is consistent with Johnston et al. (2019) and Fortuna et al. (2021), who suggest satellite shapes are effectively random on sufficiently large scales. Centrals, on the other hand, tend to align with the host halo, and so trace the large scale correlations in the background large-scale structure (Catelan et al. 2001; Kiessling et al. 2015). Although not shown in Fig. 15, it is also worth noting that the central galaxies show a clear monotonic increase in IA amplitude with redshift, a trend which is not replicated in satellite galaxies. There is also some evidence to suggest not only that satellites are distributed in a significantly anisotropic way within their haloes (Zentner et al. 2005; Piscionere et al. 2015; Butsky et al. 2016), but also that this anisotropy can have a non-trivial impact on the IA signal, even on large scales (Huang et al. 2016; Samuroff, Mandelbaum & Di Matteo 2020).","Citation Text":["Pereira, Bryan & Gill 2008"],"Citation Start End":[[1118,1144]]}
{"Identifier":"2016ApJ...816...41Y__Martin_1998_Instance_1","Paragraph":"The most striking characteristic of this event is displayed by the AIA 304 \u212b movie in which a twisted coherent structure is gradually shown to have underwent a rolling motion during the fadeaway of the brightening (panels (d) and (e)). Interestingly, it is clearly evident from the H\u03b1 observations that the twisted structure is a filament, which consists of three sections with an approximate length of about 45 Mm (panel (k)). Taking the magnetic field polarity around the filament into account, it is obvious that the filament roughly separated oppositely polarized magnetic fields, with the western ends rooted in a negative-polarity region and the eastern ends rooted in a positive-polarity region (panel (l)). As a consequence, the axial field component of the filament can be determined as pointing to the right when viewed from its positive-polarity side, and the filament was identified as dextral, which is consistent with the preferential filament pattern in the northern hemisphere (Martin 1998). Careful inspection of the 304 \u212b observations also shows that the filament probably consists of two sets of mutually intertwined dark threadlike structures (panel (e)), which may naturally explain why it displays three sections in the H\u03b1 observations. One set of the dark threadlike structures may include its western and eastern parts (see the twin arrows). They first appeared above the brightening and then separated from each other and lifted up. Finally, they formed a coherent reverse s-shaped structure, with its dipped part blocked by the other sets of dark threadlike structures (panels (b) and (e)), which appeared to be co-spatial with the spread brightening and showed up immediately after the fade away of the brightening (panels (d) and (e)). The dynamic evolution of these dark threadlike structures is also displayed on the H\u03b1 images (panels (i)\u2013(k)). The above observations suggested that the newly formed filament is a flux rope with twisted magnetic structures, and its formation is again attributed to reconnections between T1 and T2. Its apparent rolling motion might be thus understood in terms of a change in the orientation of the connecting T1 and T2. All these observations, including the two-sided spread brightenings, the newly formed EUV loop, and the dynamic formation of the flux rope seem to support the previous idea that the formation of the flux rope is the result of a tether-cutting type reconnection (Moore et al. 2001; Liu et al. 2010; Chen et al. 2014) between two sets of sheared arcade cool loops. However, our observation further reveals the formation of a filament by magnetic reconnection among T1 and T2.","Citation Text":["Martin 1998"],"Citation Start End":[[994,1005]]}
{"Identifier":"2021ApJ...920..139M__Koyama_et_al._1991_Instance_1","Paragraph":"Studies of the pulse characteristics of X-ray pulsars offer vital information about their nature and the geometry of the binary system. The shape of the pulse profile offers insight into the mode of accretion inflows, the pulsar luminosity, the geometry of the accretion column and the magnetic field configuration (Parmar et al. 1989a). Therefore, such studies offer an understanding of pulsar systems and the process of mass accretion in the presence of the strong magnetic field of a neutron star. Detailed studies of pulse characteristics were conducted during outburst activities of some of pulsars such as EXO 2030+375 (Parmar et al. 1989a), Cepheus X-4 (Koyama et al. 1991; Mihara et al. 1991; Mukerjee & Agrawal et al. 2000), XTE J1946+274 (Paul et al. 2001), etc., which offered valuable information and understanding of the physical processes responsible for the observed properties of these pulsars. Spectroscopic studies, however, reveal information about the environment surrounding a pulsar and the underlying mechanism responsible for energy generation in these systems. Detailed studies of cyclotron absorption features, if present in the pulsar spectrum, not only enable us to determine its surface magnetic field, but they also offer an insight into the line-producing region, its structure, and the geometry of the accretion column (Staubert et al. 2019). Therefore, detailed studies of cyclotron absorption features have been pursued by researchers since their discovery (Truemper et al. 1978), as they are an important diagnostic probe for pulsars. Cyclotron absorption features were detected in the spectrum of many Be binaries covering a wide range of energies, starting from a lower energy of \u223c10 keV (Jun et al. 2012; DeCesar et al. 2013) to a higher energy at \u223c100 keV (La Barbara et al. 2001). Detailed studies of cyclotron line sources and their properties were reported by Staubert et al. (2019) and Maitra (2017). The Be binaries, particularly because of their transient nature, offer us opportunities for detailed study of some of these interesting properties as a function of changes in source luminosity and in time. The studies of cyclotron line energy with respect to its pulse phase, source luminosity, and time, showed wide variations for some sources, such as Vela X-1, Cen X-3, and Her X-1 (Staubert et al. 2019). These interesting properties help us to systematically investigate and understand their underlying physical properties.","Citation Text":["Koyama et al. 1991"],"Citation Start End":[[661,679]]}
{"Identifier":"2022ApJ...940L..13A__Chhiber_et_al._2021_Instance_1","Paragraph":"Since 2018 the Parker Solar Probe (PSP) mission is collecting solar wind plasma and magnetic field data through the inner heliosphere, reaching the closest distance to the Sun ever reached by any previous mission (Fox et al. 2016; Kasper et al. 2021). Thanks to the PSP journey around the Sun (it has completed 11 orbits) a different picture has been drawn for the near-Sun solar wind with respect to the near-Earth one (Bale et al. 2019; Kasper et al. 2019; Chhiber et al. 2020; Malaspina et al. 2020; Bandyopadhyay et al. 2022; Zank et al. 2022). Different near-Sun phenomena have been frequently encountered, with the emergence of magnetic field flips, i.e., the so-called switchbacks (Dudok de Wit et al. 2020; Zank et al. 2020), kinetic-scale current sheets (Lotekar et al. 2022), and a scale-invariant population of current sheets between ion and electron inertial scales (Chhiber et al. 2021). Going away from the Sun (from 0.17 to 0.8 au), evidence of radial evolution of different properties of solar wind turbulence (Chen et al. 2020) as the spectral slope of the inertial range (from \u22123\/2 close to the Sun to \u22125\/3, at distances larger than 0.4 au), an increase of the outer scale of turbulence, a decrease of the Alfv\u00e9nic flux, and a decrease of the imbalance between outward (z\n+) and inward (z\n\u2212) propagating components (Chen et al. 2020) has been provided. Although the near-Sun solar wind shares different properties with the near-Earth one (Allen et al. 2020; Cuesta et al. 2022), significant differences have been also found in the variance of magnetic fluctuations (about 2 orders of magnitude) and in the compressive component of inertial range turbulence. In a similar way, Alberti et al. (2020) first reported a breakdown of the scaling properties of the energy transfer rate, likely related to the breaking of the phase-coherence of inertial range fluctuations. These findings, also highlighted by Telloni et al. (2021) and Alberti et al. (2022) analyzing a radial alignment between PSP and Solar Orbiter, and PSP and BepiColombo, respectively, have been interpreted as an increase in the efficiency of the nonlinear energy cascade mechanism when moving away from the Sun. More recently, by investigating the helical content of turbulence Alberti et al. (2022) highlighted a damping of magnetic helicity over the inertial range between 0.17 and 0.6 au suggesting that the solar wind develops into turbulence by a concurrent effect of large-scale convection of helicity and creation\/annihilation of helical wave structures. All these features shed new light onto the radial evolution of solar wind turbulence that urges to be considered in expanding models of the solar wind (Verdini et al. 2019; Grappin et al. 2021), and also to reproduce and investigate the role of proton heating and anisotropy of magnetic field fluctuations (Hellinger et al. 2015).","Citation Text":["Chhiber et al. 2021"],"Citation Start End":[[879,898]]}
{"Identifier":"2020MNRAS.495.4508E__Heinke_et_al._2014_Instance_2","Paragraph":"Several qLMXBs have been identified in GCs and in the Galactic field (for some examples, see table 4 in Guillot et al. 2009 and references therein). While LMXBs in the field were detected following the onset of a bright accretion outburst, most qLMXBs in GCs, including all those with the highest flux at Earth, have not shown accretion activity.3 Most of these sources have only been spectrally identified based on their similarities to field LMXBs, observed during quiescence (e.g. Cen X-4 or Aql X-1). Previous works have confirmed that H-atmosphere models accurately describe the spectra of qLMXBs, with radii in the range 10\u201315 km, as expected for NSs, either from single sources (e.g. Heinke et al. 2006a; Webb & Barret 2007; Guillot, Rutledge & Brown 2011; Heinke et al. 2014; Bogdanov et al. 2016), or from statistical analyses of multiple qLMXBs (e.g. Guillot et al. 2013; Guillot & Rutledge 2014; Lattimer & Steiner 2014; Guillot 2016; Steiner et al. 2018). However, in some cases the accreted material may not be hydrogen, but helium (e.g. Servillat et al. 2012; Catuneanu et al. 2013; Heinke et al. 2014). One way to circumvent this is to identify the nature of the donor star, i.e. to determine the nature of the material transferred on to the NS (e.g. with the detection of an H\u2009\u03b1 emission line, presumably originating in a faint accretion disc, Haggard et al. 2004). The possibility of helium (or heavier element) atmospheres is well-founded on the existence of ultracompact X-ray binaries (UCXB), with white dwarfs or helium-dominated donors4 (e.g. Zurek et al. 2009; Altamirano et al. 2010; Sanna et al. 2017; Cadelano et al. 2019). In fact, around 1\/3 of the LMXBs in GCs with constraints on the companion nature, possess a white dwarf donor (Bahramian et al. 2014). Since NS He-atmosphere models have harder spectra than H-atmosphere models, using the incorrect composition for the observed thermal emission can result in biases of the inferred radii (Servillat et al. 2012; Heinke et al. 2014).","Citation Text":["Heinke et al. 2014"],"Citation Start End":[[1097,1115]]}
{"Identifier":"2022ApJ...929...36K__Carlsten_et_al._2019b_Instance_1","Paragraph":"We overplot the observed number of member galaxies and the host stellar mass in Figure 11(a) with colored symbols. Colors indicate the five groups of \u0394m\n12. The circles and diamonds indicate the LV galaxy groups and SAGA survey sample, respectively. The LV sample, most of which is from the compilation by Carlsten et al. (2021), consists of 13 galaxy groups of which satellite membership is determined by measuring their distances using either the TRGB or SBF method: the MW (McConnachie 2012, for compiled data), M31 (McConnachie 2012; Martin et al. 2016; McConnachie et al. 2018, for compiled data), the Local Group (including the MW and M31 subgroups and quasi-isolated outlying members classified by McConnachie 2012), NGC 2403 (Carlin et al. 2021), NGC 4258 (Kim et al. 2011; Spencer et al. 2014; Carlsten et al. 2021), NGC 4631 (Tanaka et al. 2017; Carlsten et al. 2021), M51 (Carlsten et al. 2021), M101 (Bennet et al. 2017; Danieli et al. 2017; Carlsten et al. 2019b; Bennet et al. 2019), M94 (Smercina et al. 2018), NGC 1023 (Trentham & Tully 2009; Carlsten et al. 2021), M104 (Javanmardi et al. 2016; Carlsten et al. 2021), M81 (Chiboucas et al. 2009, 2013), and NGC 5128 (Crnojevi\u0107 et al. 2014; M\u00fcller et al. 2017; Crnojevi\u0107 et al. 2019; M\u00fcller et al. 2019). The other galaxy groups that are in the Carlsten et al. (2021) compilation but not this study are excluded because either the spatial coverage is too low (less than 20% of the projected virial area is covered) or the SBF S\/N is too low to confirm satellites as faint as M\n\nr\n > \u221212.3 mag. The SAGA survey sample (Mao et al. 2021) consists of galaxy groups with 36 MW-like host galaxies. The MW-like host galaxies are selected by their K-band magnitudes (\u221224.6  M\n\nK\n  \u221223 mag). The SAGA survey covered more than 80% of the projected R\n\nh\n = 300 kpc area and confirmed satellite memberships from their redshifts. It is considered complete to satellites brighter than M\n\nr\n = \u2212 12.3 mag. Note that while the spatial coverage of the LV sample significantly differs from galaxy to galaxy, and most of them cover less than the virial area, the SAGA survey sample has a relatively consistent spatial coverage.","Citation Text":["Carlsten et al. 2019b"],"Citation Start End":[[954,975]]}
{"Identifier":"2018MNRAS.476..814H__Scoccimarro_2000_Instance_2","Paragraph":"The most general third-order statistics is the three-point correlation function (hereafter referred to as 3PCF), which is defined in configurations space. Alternatively, one can study its Fourier space counterpart, the bispectrum. These two statistics contain, in principle, the same information. However, their analyses implicate different limitations and challenges, which can affect the physical interpretation of the results. A main advantage of the bispectrum is that an analysis in Fourier space allows for a clear exclusion of high-frequency modes in the density fluctuations, which are difficult to interpret theoretically due to their highly non-linear evolution. In configuration space, these high-frequency modes contribute to the 3PCF, in principle, at all scales. In practice, one therefore needs to restrict the analysis to large scales, where their contribution is negligible, lavishing a lot of valuable data. Another advantage of the bispectrum is that its covariance is diagonal for Gaussian density fluctuations. This approximation works well, even for evolved density fields, while deviations from Gaussianity can also be taken into account (Scoccimarro 2000; Sefusatti et al. 2006; Chan & Blot 2017). The covariance of the 3PCF, on the other hand, is not diagonal, even for Gaussian fluctuations, which makes the modelling more difficult (Srednicki 1993; Slepian & Eisenstein 2015; Byun et al. 2017; Gualdi et al. 2017). An additional difference in the analysis of the bispectrum and the 3PCF lies in the fact that the computation of the latter is more expensive. However, this aspect can be tackled by employing advanced algorithms and appropriate computational resources, as done in this work (see also, Barriga & Gazta\u00f1aga 2002; McBride et al. 2011a; Jarvis 2015; Slepian & Eisenstein 2015, and references therein). Besides its disadvantages, there are some arguments that speak for the 3PCF. One of them is the fact that the amplitude of the 3PCF (but not its errors) is not affected by shot-noise, whereas the latter affects the bispectrum amplitude at all scales and hence needs to be modelled for correcting the measurements. In addition, an analysis in configuration space has the advantage that complicated survey masks can be easily taken into account in the analysis of observational data, while in Fourier space such masks impose complicated effects on the measured bispectrum, which are difficult to model (e.g. Scoccimarro 2000). A more general consideration is that it is easier to interpret effects such as redshift space distortions or baryon acoustic oscillations (BAOs) on the statistics in configuration space, since that is where the physical processes that cause these effects happen. Studies of third-order correlations in the literature usually focus on either Fourier or configuration space (e.g. McBride et al. 2011b; Mar\u00edn et al. 2013; Gil-Mar\u00edn et al. 2015). However, it is worthwhile studying both statistics and cross-check the results, since their different advantages and disadvantages are quite complementary.","Citation Text":["Scoccimarro 2000"],"Citation Start End":[[2446,2462]]}
{"Identifier":"2020ApJ...896...59A__Sitnova_et_al._2018_Instance_1","Paragraph":"We used the code detail (Butler & Giddings 1985) based on the method of accelerated \u039b iteration (Rybicki & Hummer 1991). The detail opacity package was updated by Przybilla et al. (2011) by including bound\u2013free opacities of neutral and ionized species. The calculated departure coefficients, bi = nNLTE\/\n\n\n\n\n\n, were used by the code synthV_NLTE (Ryabchikova et al. 2016) to calculate the synthetic NLTE line profiles. Here, nNLTE and nLTE are the statistical equilibrium and thermal (Saha Boltzmann) number densities, respectively. The binmag code11\n\n11\n\nhttp:\/\/www.astro.uu.se\/oleg\/binmag.html\n\n (Kochukhov 2018) was used for automatic spectral line fitting and comparison with the observed spectrum. The typical uncertainties in the fitting procedure with the observed profile are less than 0.02 dex for weak lines and 0.03 dex for strong lines. For consistency with our NLTE studies of C i\u2013C ii (Alexeeva et al. 2016), Ti i\u2013Ti ii (Sitnova et al. 2016), and Ca i\u2013Ca ii (Sitnova et al. 2018), here we use exactly the same model atmospheres for 21 Peg, HD 22136, \u03c0 Cet, and \u03b9 Her, calculated with the code LLmodels (Shulyak et al. 2004), as in the earlier papers. For the remaining stars, the calculations were performed using plane-parallel (1D), chemically homogeneous model atmospheres from the Kurucz\u2019s grid12\n\n12\n\nhttp:\/\/www.oact.inaf.it\/castelli\/castelli\/grids.html\n\n (Castelli & Kurucz 2003). Table 1 lists the atomic data for the Ne i and Ne ii lines used in the present line-formation analysis. For Ne i lines belonging to the 2p53p\u20132p53s transition array, we adopted data of atomic oscillator strengths from the experimental measurements deduced using a neon-filled hollow cathode lamp in conjunction with two spectrographs (Piracha et al. 2015). For 2p53d\u20132p53p transitions at 8377 and 8495 \u212b, we also used experimental transition probabilities measured with the shock tube (Doherty 1962) and renormalized by the NIST group (Wiese et al. 1966). These values of transition probabilities are supported by the accurate theoretical calculations of Seaton (1998a). For the remaining Ne i lines and for the Ne ii lines, the oscillator strengths were taken from the NIST Handbook of Chemistry and Physics (Fuhr & Wiese 1998\u2014Ne i) and from Kurucz\u2019s website. Kurucz\u2019s theoretical calculations agree very well with the theoretical calculations by Breit\u2013Pauli R-matrix method used in the earlier papers on neon abundance analyses (see Section 3.4). Stark collisional data were adopted from the Kurucz\u2019s website.","Citation Text":["Sitnova et al. 2018"],"Citation Start End":[[972,991]]}
{"Identifier":"2021ApJ...909..114N__Margutti_et_al._2017_Instance_1","Paragraph":"Almost all the efforts to reveal the geometry of GW170817 have been based on model fitting to the EM afterglow light curve3\n\n3\nDhawan et al. (2020) suggested a method to estimate the viewing angle using the macronova\/kilonova signal. However, while the afterglow light curve is based on rather clear physics the macronova\/kilonova light curve involves uncertainties in all aspects, ranging from the composition of the ejecta to its three-dimensional distribution, unknown opacities of the relevant heavy elements and complicated radiation transfer calculations (see Nakar 2020 for a discussion of the various uncertainties). Thus, a concern arises that due to these numerous uncertainties it might be extremely difficult to control the systematics that may arise using this method.\n (Alexander et al. 2017, 2018; Haggard et al. 2017; Margutti et al. 2017; Troja et al. 2017, 2019; D\u2019Avanzo et al. 2018; Dobie et al. 2018; Gill & Granot 2018; Granot et al. 2018b; Lazzati et al. 2018; Lyman et al. 2018; Margutti et al. 2018; Mooley et al. 2018a; Troja et al. 2018; Fong et al. 2019; Hajela et al. 2019; Lamb et al. 2019; Wu & MacFadyen 2019; Ryan et al. 2020; Takahashi & Ioka 2020). All recent studies agree that GW170817 have launched a relativistic jet that broke out of the ejecta successfully and then interacted with the circum-merger medium to produce the afterglow. They all find that the jet had an angular structure with a narrow highly energetic core, with a jet opening \u03b8j, which is observed from a viewing angle \u03b8obs \u226b \u03b8j. Each of these studies obtained a different estimate of \u03b8obs, and \u03b8j (see Section 5 and Figure 3). The viewing angles found by different studies are in the range of 14\u00b0\u201338\u00b0 and the jet opening angle is in the range of 25\u20138\u00b0. The typical errors quoted in these studies are a few degrees on the viewing angle and a fraction of a degree on the jet opening angle. Thus, the various estimates are often highly inconsistent with each other. An exception is the work of Mooley et al. (2018b) who used the motion of the radio image as measured in VLBI observations (see also Ghirlanda et al. 2019; Hotokezaka et al. 2019) in addition to the light curve to constrain the system geometry.","Citation Text":["Margutti et al. 2017"],"Citation Start End":[[834,854]]}
{"Identifier":"2017AandA...599A...4K__hand,_Hotta_et_al._(2016)_Instance_2","Paragraph":"Our results appear to stand apart from similar studies in full spherical shells (e.g., Nelson et al. 2013; Hotta et al. 2016) in that the differential rotation is strongly quenched as a function of the magnetic Reynolds number. However, in Nelson et al. (2013) the values of Rm\u2032 (=2\u03c0ReM) correspond to a range of 8...32 in ReM in our units where the radial and latitudinal differential rotation decrease by about 30 per cent. This is roughly consistent with our results. On the other hand, Hotta et al. (2016) reached higher values of ReM than in the present study, but no strong quenching was reported. The reason might be that their models are rotating substantially slower than ours, leading to weaker magnetic fields and a weaker back-reaction to the flow. Furthermore, in these models, the differential rotation is strongly influenced by their SGS heat flux, which transports one third of the luminosity. Another obvious candidate for explaining the difference is the wedge geometry used in the current simulations. However, we note that earlier simulations with a similar setup did not show a marked trend in the energy of the differential rotation as the azimuthal extent of the domain was varied (see Table 1 of K\u00e4pyl\u00e4 et al. 2013). However, results of Boussinesq simulations of convective dynamos have shown a similar change as a function of the magnetic Prandtl number (Schrinner et al. 2012). The drop in the amplitude of the differential rotation was associated with a change in the dynamo mode from an oscillatory multipolar solution to a quasi-steady dipolar configuration (cf. Fig. 15 of Schrinner et al. 2012) that prevents strong differential rotation from developing. We do not find a strong dipole component in our simulations (see Sect. 4.2.3). However, the strong suppression of the differential rotation often coincides with the appearance of a small-scale dynamo (see Table 1 and the discussion in the Sect. 4.1.2) or a change in the large-scale dynamo mode as discussed above. ","Citation Text":["Hotta et al. (2016)"],"Citation Start End":[[490,509]]}
{"Identifier":"2019AandA...626A..64H__Gies_et_al._2003_Instance_1","Paragraph":"Already soon after the identification of the optical counterpart of Cyg X-1, X-ray light curves were found to show a strong orbital modulation of the X-ray absorption column, NH, due to absorption of X-rays from the black hole in the stellar wind (Li & Clark 1974; Remillard & Canizares 1984; Ba\u0142uci\u0144ska-Church et al. 2000; Poutanen et al. 2008; Mi\u0161kovi\u010dov\u00e1 et al. 2016; Grinberg et al. 2015). Together with observations of the orbital modulation of optical lines from HDE 226868 (Gies & Bolton 1986a,b; Gies et al. 2003), these phenomena led to the picture of the stellar wind of HDE 226868 as a line driven wind or CAK wind after Castor et al. (1975, see also Friend & Castor 1982 and Morton 1967) with an asymptotic velocity of v\u221e ~ 2000 km s\u22121 (Muijres et al. 2012). This wind is disturbed by the gravitational potential of the black hole, which leads to a focusing of the wind toward the black hole; see Friend & Castor (1982) and Gies & Bolton (1986b) for early models and El Mellah et al. (2019) for a modern treatment ofthis process. In addition, the wind is also affected by the X-rays from the compact object: strong orbital modulation of NH is seen during the canonical hard state of the black hole, where the X-ray spectrum is dominated by a Comptonized power law (Parker et al. 2015; Nowak et al. 2011; Wilms et al. 2006, and references therein). The typical bolometric luminosity of Cyg X-1 in this state is around 2 \u00d7 1037 erg s\u22121 (e.g., Wilms et al. 2006; Nowak et al. 1999), albeit with a large uncertainty due to our lack of knowledge of the UV spectral shape. Only very little NH modulation is seen during the thermally dominated X-ray soft state (Wen et al. 1999; Boroson & Vrtilek 2010), in which the typical bolometric luminosity is at most a factor two higher than in the hard state (e.g., Tomsick et al. 2014; Zhang et al. 1997). Optical spectra show the stellar wind to be strongly photoionized during the latter state (Gies et al. 2003, 2008).","Citation Text":["Gies et al. 2003","Gies et al. 2003"],"Citation Start End":[[504,520],[1945,1961]]}
{"Identifier":"2015MNRAS.450.4364N__Wu_et_al._2004_Instance_2","Paragraph":"Low- and intermediate-mass stars are formed by the gravitational collapse of the parental giant molecular cloud (GMC), followed by the accretion process (Palla 1996). During the accretion phase, material is ejected as well via collimated bipolar jets. However, when a YSO reaches 8 M\u2299, the radiative flux becomes so intense (using \u03d5 = L\/4\u03c0d2, the ratio between the radiative fluxes of an O5 and a B3 star \u2013 masses of \u223c40 and \u223c8 M\u2299, respectively \u2013 is \u2248250) that it may interrupt the accretion flow. A process that constrains the outcoming radiation field to narrower angles may leave some room for the accretion process to continue in some directions. This seems to be the case for the outflows driven by young stars from a very broad mass range, as previous reported by several authors (Bachiller 1996; Bontemps et al. 1996; Shepherd & Churchwell 1996; Beuther et al. 2002; Wu et al. 2004). Outflows associated with high-mass objects are expected to be more energetic than the outflows observed in lower mass YSOs (Beuther et al. 2005; Zhang et al. 2005; L\u00f3pez-Sepulcre et al. 2009), with velocities greater than \u223c100\u2009km\u2009s\u22121 (Mart\u00ed, Rodr\u00edguez & Reipurth 1998). Some authors have found evidences that outflows associated with massive stars are scaled up versions of their low-mass counterparts (Vaidya et al. 2011; Codella et al. 2013) while other works have reported that no well-collimated outflows have been found towards MYSOs (Shepherd, Testi & Stark 2003; Sollins et al. 2004). Massive YSO outflows mapped in high-velocity CO lines have collimation factors R = length\/width \u223c2.05 \u00b1 0.96 as compared to R \u223c 2.81 \u00b1 2.16 for low-mass stars (Wu et al. 2004), indicating a weak tendency that outflows associated with massive stars are less collimated than those from low-mass stars as previously thought (Richer et al. 2000). Besides the degree of collimation, these massive outflows would work removing mass from the plane of the accretion disc, lowering the density on the plane and, therefore, facilitating the accretion flow to reach the stellar core as shown in the recent 3D simulations presented by Krumholz et al. (2009). Although these authors have not included the outflow activity on their simulations, they argue that the presence of outflows would decrease the star formation efficiency from 70\u2009per\u2009cent (considering purely radiation effects) to 50\u2009per\u2009cent.","Citation Text":["Wu et al. 2004"],"Citation Start End":[[1643,1657]]}
{"Identifier":"2021MNRAS.505.1046X__McNish_&_Lincoln_1949_Instance_1","Paragraph":"As we all know, sunspots that can represent solar activity have a long history of research. There are many researches on sunspot variation and solar cycle prediction. The prediction of the sunspot cycle can be classified as following three types: The first is the time series method. In this method, the sunspot number is regarded as a function of time for statistical prediction, in which the observed value or smoothed value of the sunspot number is regarded as a non-stationary random time series (Solanki 2003). The traditional approach is to transform the non-stationary into stationary time series, such as the famous McNish\u2013Lincoln method (McNish & Lincoln 1949) and Box\u2013Jenkins method (Box & Jenkins 1990). These kinds of the method have not yet found an ideal way to accomplish the transformation. It can only get the general trend of solar activity with the help of the periodicity of the series itself but cannot get satisfying prediction accuracy. The neural network method developed in the last decade is also used in forecasting. For example, Chattopadhyay et al. (2011) found that the effect of the auto-regressive neural network (ARNN) model in forecasting the solar cycle is better than the traditional time series method. With the sunspot number from 1749 to 2019, Pala & Atici (2019) used a long-term memory algorithm and neural network auto-regressive (NNAR) algorithm to predict the solar activity in the 25th cycle, showing that sunspots will reach the peak between 2022 and 2023. Although the neural network method sometimes performs better than many classical time series methods, however usually it is just like a black box without physical meaning and the performance also needs further validated. Similar to the above method is the time curve fitting method, Wilson (1988,1992) proposed the concept of long and short sunspot cycles, by using three sinusoidal curves to fit the sunspot number curve and extrapolate the predicted value. Some people also used other curves or multiple line regression methods to predict the sunspot number (Nordemann 1992; Hathaway, Wilson & Reichmann 1994). In particular, Hathaway et al. (1994) proposed a simple function (HWR function) to approximate the solar cycle profile with a certain predictive function. This kind of method aims at the profile prediction of a solar cycle, which is suitable for the prediction of the solar activity in the declining phase of the cycle, but the prediction for the peak value of the cycle is poor and not robust.","Citation Text":["McNish & Lincoln 1949"],"Citation Start End":[[647,668]]}
{"Identifier":"2021AandA...647A..35B__Dartois_et_al._2018_Instance_1","Paragraph":"It is expected that photons and\/or cosmic rays coming from various sources could trigger the ejection of methanol from the icy dust grains into the gas phase and participate in the overall gas-to-ice balance of these cold regions. For the specific case of photons, this mechanism is known as photodesorption. Several experimental studies have been conducted to quantitatively constrain these processes in order to explain astrophysical observations, especially for methanol-containing ices. In one of these experiments, heavy ion 136Xe23+ irradiations were performed on methanol in pure ice and embedded in a water-ice matrix (Dartois et al. 2019). A sputtering yield of methanol close to that of the main water-ice matrix (Dartois et al. 2018), which is ~ 104 sputtered molecule per incident ion, was estimated for each studied ice. When it was embedded in a CO2 ice, Dartois et al. (2020) found that this sputtering yield is about six times higher. Experimental studies of UV photodesorption in the 7\u201310.5 eV range were first conducted for pure methanol ice at 20 K by \u00d6berg et al. (2009). More recent experiments by Bertin et al. (2016) in the 7\u201314 eV range suggested an efficiency for methanol UV photodesorption from pure methanol ice of ~ 10\u22125 molecule\/photon. This desorption was found to be below the detection threshold (10\u22126 molecule\/photon) when methanol was mixed with CO ice for a wide range of dilution factors (from 1 in 4 to 1 in 50; Bertin et al. 2016). Accordingly, Cruz-Diaz et al. (2016) derived an upper limit of 3 \u00d7 10\u22125 molecule\/photon for the UV photodesorption yield of methanol from pure methanol ice from 8 to 130 K (in the 6.88\u201310.9 eV range). In addition to these previous mechanisms, chemical desorption, which is the desorption induced by exothermic reactions, is a possible route for explaining the gas-phase enrichment in the ISM (Cazaux et al. 2016; Minissale et al. 2016b; Ligterink et al. 2018). However, chemical desorption of methanol by H addition onto CO, H2 CO, and CH3 OH ices was not detected (upper limit 5%; Minissale et al. 2016b), and more generally, chemical desorption of several other molecules appears to be very low when experiments are made on water substrates (Minissale et al. 2016a).","Citation Text":["Dartois et al. 2018"],"Citation Start End":[[724,743]]}
{"Identifier":"2017ApJ...844L..19P__Shapiro_2000_Instance_1","Paragraph":"We model the NS structures using the EOSs summarized in Table 1. These have been chosen to span a representative range of maximum NS masses, but all with maximum gravitational masses \n\n\n\n\n\n given current observational limits (Demorest et al. 2010; Antoniadis et al. 2013). For each of these EOSs, we solve for the hydrostatic structure including general relativity and compute a grid of NS models spanning a wide range of masses. We calculate the structures for both a non-rotating NS and a maximally rotating NS for use in treating the merger remnant (further discussed in Section 3). For the latter case, we assume solid-body rotation because magnetic braking eliminates differential rotation on an Alfv\u00e9n timescale of \u223c10\u2013100 ms (Baumgarte et al. 2000; Shapiro 2000). All models have been computed using the LORENE library and in particular its publicly available codes rotseq and nrotstar. The former can easily compute sequences of non-rotating NSs, while the latter has been used to compute equilibrium configurations at the mass-shedding limit (Stergioulas 2003). The different EOSs were imported in LORENE in a table format either produced by using a piecewise polytropic approximation (APR4, Endrizzi et al. 2016; H4, Kawamura et al. 2016; MS1, Ciolfi et al. 2017), provided by the EOS authors (SHT; Kastaun et al. 2016), or given by the Compose project (GM16\n\n6\n\nhttp:\/\/compose.obspm.fr\/\n\n; Glendenning & Moszkowski 1991; Douchin & Haensel 2001). This large grid of models gives us, for each mass and EOS, a relation between the gravitational mass\n1\n\n\n\n\n\nwhere e is the energy density and R is the NS radius, and the baryonic mass (Shapiro & Teukolsky 1983)\n2\n\n\n\n\n\nwhere m(r) is the enclosed mass at a given radius and \u03c1 is the rest-mass density (both equations refer to non-rotating models, but similar ones can be derived for rotating models). The important point is that in general \n\n\n\n\n\n because of the gravitational redshift factor. In the fast rotating limit, we find that Mg is greater by \u22481.18\u20131.20, consistent with the calculations by Breu & Rezzolla (2016).","Citation Text":["Shapiro 2000"],"Citation Start End":[[756,768]]}
{"Identifier":"2017ApJ...838...67E__Engle_et_al._2014_Instance_1","Paragraph":"In the present era of \u201chigh precision\u201d cosmology (see Riess et al. 2016), it is important to exploit the full potential of Cepheids as precise extragalactic distance indicators for determining the expansion rate of the universe and setting constraints on cosmology models. To achieve these goals, a deeper understanding and characterization of Cepheids is needed. Recent discoveries such as circumstellar environments (Nardetto et al. 2016), infrared excesses (M\u00e9rand et al. 2015), and ultraviolet emission line variability and possible more recent X-ray emissions (Engle et al. 2014) show that some important aspects of Cepheids may not be well understood. Cepheids have also been found to show additional complications that include cycle-to-cycle variations in their light and radial velocity curves (see Evans et al. 2015b; Anderson 2016; Anderson et al. 2016; Smolec & \u015aniegowska 2016; Derekas et al. 2017). These newly discovered properties and time-dependent phenomena of Cepheids, unless better understood and accounted for, could place impediments on achieving the challenging goal of determining the local Hubble constant (H0) with a precision of \u223c1%, as suggested by Suyu et al. (2012). Great efforts are being undertaken to achieve this level of precision, and hopefully resolve the developing \u201cHubble Discrepancy\u201d (see Riess et al. 2016), where theoretical values of the Hubble constant (H0) derived via the Lambda-cold dark matter (\u039b-CDM) cosmology model (\u039b = the cosmological constant), including cosmic microwave background data (e.g., Planck, WMAP (Wilkinson Microwave Anisotropy Probe)), show a small (albeit statistically significant) disagreement with the value of H0  will be derived via standard candles (e.g., Cepheids, SNe Ia (type Ia Supernovae)). As discussed by Suyu et al. (2012) and more recently by Riess et al. (2016), improved measurements of H0 provide critical independent constraints on dark energy and the validity of the present \u039b-CDM model.","Citation Text":["Engle et al. 2014"],"Citation Start End":[[566,583]]}
{"Identifier":"2022AandA...663A..11L__Haiman_et_al._2000_Instance_1","Paragraph":"Several scenarios have been suggested to explain the presence of magnesium-enriched gas in between the galaxies of a group. Magnesium atoms are \u03b1 elements released in the ISM and CGM by core collapse supernovae. One can therefore naturally imagine that the magnesium gas residing outside galaxies has been ejected through strong supernovae-driven outflows. The P Cygni shape of the Mg\u202fII doublet in the integrated spectrum of galaxy A (see Fig. 2 top) is in good agreement with this scenario. Indeed, according to radiative transfer (RT) models, this line profile is a signature of resonant scattering in an optically thick outflowing medium where the redshifted emission corresponds to back-scattered photons reflected by the receding medium (Haiman et al. 2000; Dijkstra et al. 2006; Verhamme et al. 2006; Kollmeier et al. 2010). The Mg\u202fII lines of galaxies D and E seem to be redshifted compared to their systemic redshifts (Fig. 2), however, the continuum is too faint to detect significant absorption lines and, thus, too faint to detect the P Cygni profiles, namely, the presence of outflows for those galaxies. We detected > 2\u03c3 Mg\u202fII emission aligned with the minor axis of galaxy E and near galaxy D (Fig. 3). These regions could have been enriched via outflows after supernovae feedback. However, the S\/N of our data is too poor in these areas to carry out a resolved kinematics study. The resolved Mg\u202fII map (Fig. 6) shows that the two regions along the minor axis of galaxy A have different velocities with opposite signs (\u2248\u2005+\u2005120 and \u221280 km s\u22121 with respect to the systemic redshift of galaxy A), providing additional evidence of the presence of an outflow. Even if we do not eventually find clear evidence of an outflow in the four other galaxies, we cannot exclude the deposition of metals in the CGM by past feedback processes explaining the observed diffuse Mg\u202fII emission detected around each of the galaxy members. In that scenario, we could imagine that the low-surface-brightness gaseous bridge in-between the galaxy subgroups [A,B,C] and [D,E] could be due to galactic wind transfer between the group members as observed in the FIRE simulation (Angl\u00e9s-Alc\u00e1zar et al. 2017).","Citation Text":["Haiman et al. 2000"],"Citation Start End":[[744,762]]}
{"Identifier":"2021ApJ...923L..22A__Sesana_et_al._2016_Instance_1","Paragraph":"Pulsar timing experiments (Sazhin 1978; Detweiler 1979) allow us to explore the low-frequency (\u223c1\u2013100 nHz) part of the gravitational-wave (GW) spectrum. By measuring deviations from the expected arrival times of radio pulses from an array of millisecond pulsars, we can search for a variety of GW signals and their sources. The most promising sources in the nanohertz part of the GW spectrum are supermassive binary black holes (SMBHBs) that form via the mergers of massive galaxies. Orbiting SMBHBs produce a stochastic GW background (GWB; Lommen & Backer 2001; Jaffe & Backer 2003; Volonteri et al. 2003; Wyithe & Loeb 2003; Enoki et al. 2004; Sesana et al. 2008; McWilliams et al. 2012; Sesana 2013; Ravi et al. 2015; Rosado et al. 2015; Kelley et al. 2016; Sesana et al. 2016; Dvorkin & Barausse 2017; Kelley et al. 2017; Bonetti et al. 2018; Ryu et al. 2018), individual periodic signals or continuous waves (CWs; Sesana et al. 2009; Sesana & Vecchio 2010; Mingarelli et al. 2012; Roedig & Sesana 2012; Ravi et al. 2012, 2015; Rosado et al. 2015; Schutz & Ma 2016; Mingarelli et al. 2017; Kelley et al. 2018), and transient GW bursts (van Haasteren & Levin 2010; Cordes & Jenet 2012; Ravi et al. 2015; Madison et al. 2017; Islo et al. 2019; B\u00e9csy & Cornish 2021). We expect to detect the GWB first, followed by detection of individual SMBHBs (Siemens et al. 2013; Rosado et al. 2015; Taylor et al. 2016; Mingarelli et al. 2017) that stand out above the GWB. Detection of GWs from SMBHBs will yield insights into galaxy mergers and evolution not possible through any other means. Other potential sources in the nanohertz band include cosmic strings (Damour & Vilenkin 2000, 2001; Berezinsky et al. 2004; Damour & Vilenkin 2005; Siemens et al. 2006, 2007; \u00d6lmez et al. 2010; Sanidas et al. 2013; Blanco-Pillado et al. 2018; Chang & Cui 2021; Ghayour et al. 2021; Gorghetto et al. 2021; Wu et al. 2021a; Blanco-Pillado et al. 2021; Lin 2021; Chiang & Lu 2021; Lazarides et al. 2021; Chakrabortty et al. 2021; Ellis & Lewicki 2021), phase transitions in the early universe (Witten 1984; Caprini et al. 2010; Addazi et al. 2021; Arzoumanian et al. 2021; Di Bari et al.2021; Borah et al. 2021; Nakai et al. 2021; Brandenburg et al.2021; Neronov et al. 2021), and relic GWs from inflation (Starobinski\u01d0 1979; Allen 1988; Lazarides et al. 2021; Ashoorioon et al. 2021; Yi & Zhu 2021; Li et al. 2021; Poletti 2021; Vagnozzi 2021; Sharma 2021), all of which would provide unique insights into high-energy and early-universe physics.","Citation Text":["Sesana et al. 2016"],"Citation Start End":[[761,779]]}
{"Identifier":"2021AandA...647A..81K__Verwichte_&_Kohutova_2017_Instance_1","Paragraph":"Small-amplitude oscillations are observed without any clearly associated driver and can persist for several hours (Wang et al. 2012; Nistic\u00f2 et al. 2013; Anfinogentov et al. 2015). Their amplitudes are of the order of 100 km and they show no observable damping. Even though it is clear they must be driven by some small-scale abundant process, the exact excitation mechanism of this oscillation regime is unclear. Several excitation mechanisms have been proposed to explain the persistent nature of these oscillations. These include mechanisms acting at coronal heights such as the onset of thermal instability in the corona (Kohutova & Verwichte 2017; Verwichte & Kohutova 2017), the presence of siphon flows in the coronal loop (Kohutova & Verwichte 2018), self-oscillation of coronal loops due to interaction with quasi-steady flows (Nakariakov et al. 2016; Karampelas & Van Doorsselaere 2020), and footpoint-concentrated drivers. A commonly assumed excitation mechanism is associated with turbulent vortex flows which are ubiquitous in the lower solar atmosphere (Carlsson et al. 2010; Shelyag et al. 2011; Liu et al. 2019) and result in random footpoint buffeting. This is typically modelled as a stochastic driver at the footpoints (e.g., Pagano & De Moortel 2019). Such a driver however does not reproduce the observational characteristics of this type of oscillation, in particular the stability of the oscillation phase and lack of any observable damping (Nakariakov et al. 2016). It has been proposed that a footpoint driver of the form of a broadband noise with a power-law fall-off can in principle lead to excitation of eigenmodes along coronal loops (Afanasyev et al. 2020), but this has not been tested through MHD simulations. Finally, excitation of transverse oscillations through solar p-mode coupling has also been proposed in several studies (Tomczyk & McIntosh 2009; Morton et al. 2016, 2019; Riedl et al. 2019). As p-modes correspond to acoustic waves generated by turbulent flows in the solar convection zone, they are subject to acoustic cut-off at the \u03b2\u2004=\u20041 equipartition layer. High magnetic field inclinations can however lead to p-mode leakage into chromosphere even for frequencies below the cut-off frequency (De Pontieu et al. 2004), and to subsequent coupling to different magnetoacoustic wave modes (Santamaria et al. 2015). However, three-dimensional (3D) MHD simulations of coronal fluxtubes with a gravity-acoustic wave driver at the footpoints did not reveal any evidence of the acoustic driving leading to significant transverse displacement of the fluxtubes at coronal heights (Riedl et al. 2019).","Citation Text":["Verwichte & Kohutova 2017"],"Citation Start End":[[653,678]]}
{"Identifier":"2021AandA...647A..67B___2007_Instance_1","Paragraph":"The expected pressure profile of the interstellar medium inside the Bondi sphere is p\u2004\u221d\u2004z\u22122 and, according to analytical and numerical models (Tchekhovskoy et al. 2008; Komissarov et al. 2009; Beskin et al. 2017), a jet propagating in a medium with such a pressure gradient develops the characteristic parabolic shape. Based on Eq. (3), estimating the Bondi radius requires the knowledge of the black hole mass and of the temperature of the accretion flow in the nuclear region. A detailed analysis of the X-ray emission in NGC 315 was carried out by Worrall et al. (2003, 2007) using sensitive Chandra data. The jet, which presents several X-ray-bright knots of synchrotron emission, is embedded in a hot gaseous atmosphere, as inferred from the presence of an X-ray thermal emission component in the spectrum. Such a feature is also detected within a circle of 1 arcsec radius centered around the nucleus, in addition to a dominant, mildly absorbed power-law component possibly associated to the jet. The nuclear hot atmosphere has a temperature \n\n\n\nk\nT\n=\n0\n.\n\n44\n\n\u2212\n0.04\n\n\n+\n0.08\n\n\n\n\n$ kT=0.44^{+0.08}_{-0.04} $\n\n\n keV. Assuming a black hole mass of 1.3\u2005\u00d7\u2005109 M\u2299, we then estimate rB\u2004=\u200492 pc (dark blue vertical line in Fig. 3). Even if the black hole mass was smaller by a factor of several, this radius would still be much larger (two orders of magnitude) than the distance at which we observe the jet shape transition from parabolic to conical, zt\u2004=\u20040.58\u2005\u00b1\u20050.28 pc. This result would also not be affected by varying the gas temperature within the given small uncertainty. By assuming slightly different parameters, a similar value for the Bondi radius in NGC 315 was recently obtained by Inayoshi et al. (2020), who estimated an uncertainty of 50%3. Then, if the transition from parabolic to conical shape in NGC 315 is induced by a change in the external pressure gradient, this must occur not in the proximity of the black hole sphere of influence, but in the vicinity of the black hole, on sub-parsec scales. Based on the analysis of the VLA data in Fig. 3, right panel, we also note that no discontinuity is observed in the expansion profile after the jet crosses the Bondi radius: the jet shape is close to conical both inside and outside the Bondi sphere.","Citation Text":["Worrall et al.","2007"],"Citation Start End":[[551,565],[573,577]]}
{"Identifier":"2021MNRAS.505..515Z__Naul_et_al._2018_Instance_1","Paragraph":"For variable star classification, both convolutional neural networks (CNNs; LeCun, Bengio & Hinton 2015) and recurrent neural networks (RNNs; Hochreiter & Schmidhuber 1997; Cho et al. 2014) have been shown to be competitive to the traditional RF-based methods. Naul et al. (2018) used an RNN autoencoder network to learn low-dimensional representations of period-folded light curves in an unsupervised fashion. This representation was then, in a supervised context, used as feature inputs to an RF classifier. They showed that the learned features are at least as good as, and often better than, two sets of state-of-the-art hand-crafted features (Richards et al. 2011; Kim & Bailer-Jones 2016), in terms of downstream classification accuracy. Becker et al. (2020) used an RNN for which instead of period-folding, each input light curve is grouped with a moving time-sample window of size 50 and stride 25. Although period-folding improves performance (Naul et al. 2018), Becker et al. (2020)\u2019s time-space RNN does not require the period to be calculated, and is thus less computationally expensive in terms of preprocessing. Again, they found similar performance to an RF classifier with the Nun et al. (2015) features over three data sets, although lower accuracy was seen for many sub-classes with the OGLE data set (Table 2; see Section 3.3 for data description). More recently, Jamal & Bloom (2020) systematically benchmarked the performance of different configurations of RNN and CNN network architectures on variable star classification. Aside from other work (e.g. Aguirre, Pichara & Becker 2018; Tsang & Schultz 2019) evaluating neural network (NN) performance retrospectively on previously labelled data sets, D\u00e9k\u00e1ny & Grebel (2020) used an RNN classifier to identify a new sample of fundamental-mode RR Lyrae (RRab) stars. Similarly, D\u00e9k\u00e1ny et al. (2019) found Classical and Type II Cepheids with a CNN classifier, also using the VISTA Variables in the Via Lactea (VVV) survey (Minniti et al. 2010) and using period-folded light curves.","Citation Text":["Naul et al. (2018)"],"Citation Start End":[[261,279]]}
{"Identifier":"2022MNRAS.509.3339K__Valenti_&_Piskunov_1996_Instance_1","Paragraph":"We determine the spectroscopic parameters by using a relatively high S\/N (\u223c 65 per pixel) spectrum taken without the iodine cell at TCES-TLS. As a first step, we applied the empirical software SpecMatch-Emp code (Hirano et al. 2018) that compares well-characterized FGKM stars observed with Keck\/HIRES to the data, in our case, the TCES-TLS spectra with R = 67\u2009000. We obtain a stellar effective temperature, Teff = 5804 \u00b1 110 K, a stellar radius, R*  = 2.078 \u00b1 0.180R\u2299, and the abundance of the key species iron relative to hydrogen, [Fe\/H] = 0.29 \u00b1 0.09 (dex) as mentioned in Table 4. To perform a more detailed model, we used the spectral analysis package SME (Spectroscopy Made Easy; Valenti & Piskunov 1996; Piskunov & Valenti 2017) version 5.22. This code is based on grids of recalculated stellar atmospheric models that calculates a synthetic spectrum. The best-fitted stellar parameters are derived through a \u03c72 \u2212minimization iteration that compares the synthetic and observed spectra for a given set of parameters. The spectra was synthesised based on the atomic and molecular line data from VALD (Ryabchikova et al. 2015) and the Atlas12 (Kurucz 2013) atmosphere grids. We refer to Fridlund et al. (2017) and Persson et al. (2018) for a more detailed description of the modelling. Briefly, we chose to model spectral features sensitive to different photospheric parameters: Teff from the H\u03b1 line wings, and the surface gravity, log\u2009g*, from the Ca\u2009i \u03bb\u03bb6102, 6122, and 6162 triplet, and the \u03bb6439 line. The abundances of iron and calcium, and the projected stellar rotational velocity, v sin\u2009i were fitted from narrow, unblended lines between 6100 and 6500 \u00c5. We derive the vsin\u2009i as 7.0 \u00b1 0.5 km s\u22121. We fixed the micro- and macro-turbulent velocities, Vmic and Vmac, to 1 and 3 km s\u22121 (Bruntt et al. 2010; Doyle et al. 2014), respectively. The results from SME are in good agreement within the uncertainties with SpecMatch-Emp (Table 4). The SME derived stellar parameters are supplied as priors in the global modelling. The finally adopted stellar parameters obtained through global modelling of the RV and transit data are indicated in the last row of Table 4 in bold font. For more details on the global modelling, see Section 3.3.2. One thing to note here is that we find the TOI-1789 is a late F-type star with our spectral analysis. Nevertheless, SIMBAD reported it as \u2018K0\u2019 spectral type.","Citation Text":["Valenti & Piskunov 1996"],"Citation Start End":[[688,711]]}
{"Identifier":"2022AandA...665A..22B__Stecklum_et_al._2021_Instance_1","Paragraph":"The gravitational collapse of the G345.88-1.10 hub provides the means for rapid mass provision to the centre, along with non-steady accretion flows which might open the cavity on short timescales through an accretion burst. This can be the result of e.g. jet precession (Rosen & Krumholz 2020) and opening of the jet itself (Cesaroni et al. 2018). The modelling by Cesaroni et al. (2018) indicates that the jet opening angle can increase to values as high as 40\u201350\u00b0, but it should also be noted that this is model dependent. The simulations by Rosen & Krumholz (2020) indicate that a jet precession can be very rapid at the early stages of high-mass star formation with rapid angle changes that on average reach 40\u00b0 kyr\u22121 with apparent peaks as high as 250\u00b0 kyr\u22121. However, it should be noted that the opening of the cavities by an accretion burst also faces the issue that this is generally associated with a strong increase of the core brightness in the far-infrared (e.g. Stecklum et al. 2021; Hunter et al. 2021), which is not seen in G345.88-1.10. It could be that, after an accretion burst has occurred, there is a time delay between the dimming of the core and the dimming of the cavity. Such a time delay would be of the order of the jet\u2019s crossing time of the cavity, i.e. \u22721000 yr. We could therefore be catching G345.88-1.10 right at that moment where the cavity is still bright but the core has dimmed. Although this would be a noteworthy coincidence, it can fit with the cooling timescale we estimate for the massive H1 fragment which should be lower than 20 yr and can be as low as a couple of days or months, see Appendix H. These timescales also seem to be in agreement with surveys of continuum variability, at least in low-mass star forming cores that typically find very rapid changes on the timescales of weeks to several years (e.g. Johnstone et al. 2013; Mairs et al. 2017; Lee et al. 2021). Because of these very short cooling timescales, it is even possible that multiple accretion bursts spread over few 10-100 yr are at the origin of the heated cavities. In Appendix H, we also estimated the cooling timescale in the cavities for which we find values between 1.7 \u00d7 103 and 2.7 \u00d7 104 yr which is significantly longer than for the core. Furthermore, the lower end of this estimated timescale is of the same order of magnitude as the crossing timescale for the jet in the cavities. This might explain why the full cavity is bright instead of only several subregions in the cavity.","Citation Text":["Stecklum et al. 2021"],"Citation Start End":[[975,995]]}
{"Identifier":"2017ApJ...839...56T___2013_Instance_3","Paragraph":"Mid-infrared and millimeter polarimetric observation has so far been considered as the best method to probe the magnetic field. This is because if aspherical grains in disks become aligned with the magnetic field as is the case in the interstellar medium (ISM), the polarization vector arising from thermal emission of the aligned grains becomes perpendicular to the local magnetic field line (Cho & Lazarian 2007, henceforth CL07, Matsakos et al. 2016; Yang et al. 2016b; Bertrang et al. 2017). At mid-infrared wavelengths, Li et al. (2016) performed a polarimetric imaging observation of AB Aur using CanariCam. As a result, they detected a centrosymmetric polarization pattern, and the degree of polarization was as high as 1.5% at large radii. At millimeter wavelengths, polarimetric observations of disks have been performed (e.g., Hughes et al. 2009, 2013; Rao et al. 2014; Stephens et al. 2014; Cox et al. 2015; Kataoka et al. 2016b). Polarized emission from a circumstellar disk has been detected in the Class 0 phase (Rao et al. 2014; Cox et al. 2015). More evolved disks do not show a degree of linear polarization larger than 0.5% (Hughes et al. 2009, 2013). It should be mentioned that Stephens et al. (2014) detected polarized flux from HL Tau, which is classified as a Class I-II, with an average degree of linear polarization of 0.9%. More recently, Kataoka et al. (2016b) reported the first submillimeter polarization observation of a disk obtained with ALMA, and they clearly detected polarized flux from HD 142527. The polarization fraction at a peak position of polarized intensity was 3.26%, and the maximum polarization fraction was as high as 13.9%. The disk reveals radial polarization vectors; however, they flip by 90\u00b0 in its northeast and northwest regions. In addition, the detected polarization fraction is much larger than the stringent limit set by Hughes et al. (2009, 2013), and further polarimetric observations by ALMA will reconcile this discrepancy.","Citation Text":["Hughes et al.","2013"],"Citation Start End":[[1879,1892],[1900,1904]]}
{"Identifier":"2020MNRAS.492.1295P__Bonning_et_al._2012_Instance_1","Paragraph":"The evolution of colour or the spectral index, \u03b1, (F(\u03bd) \u221d \u03bd\u2212\u03b1 where \u03bd is the radiation frequency and F(\u03bd) is the flux density provides an insight into the particle distribution giving rise to the observed flux density and its variability. In particular, at the synchrotron frequencies, within the simplest scenario of single-zone emission models with homogeneous magnetic field distributions, clear patterns between the spectral index and the total intensity are predicted, i.e. a \u2018spectral hysteresis\u2019, depending on the relative lengths of the radiative cooling time-scale and the escape time-scale of the accelerated particles from the emission zone (e.g. Kirk, Rieger & Mastichiadis 1998). A significant fraction of long-term multiband flux monitoring studies have revealed bluer-when-brighter trends for BL Lac objects but frequently redder-when-brighter trends for FSRQs (Gu et al. 2006; Osterman Meyer et al. 2009; Hao et al. 2010; Rani et al. 2010; Ikejiri et al. 2011; Bonning et al. 2012; Sandrinelli, Covino & Treves 2014; Li et al. 2018; Meng et al. 2018; Gupta et al. 2019) while \u2018achromatic\u2019 flux variability (no colour evolution, Stalin et al. 2006; Bonning et al. 2012; Gaur et al. 2019), and erratic patterns (Wierzcholska et al. 2015) have also been reported. It has been argued that particles accelerated to higher energies are injected at the emission zone before being cooled radiatively in BL Lac sources leading to their overall SEDs being bluer-when-brighter; however, the \u2018redder\u2019 and more variable jet-component can overwhelm the \u2018bluer\u2019 contribution from the accretion disc, leading to redder-when-brighter trends for FSRQ type sources (Gu et al. 2006). Achromatic variability is often ascribed to changes in the Doppler boosting factor (\u03b4) as each frequency notes the same special relativistic multiplication of flux (Gaur et al. 2012). However, erratic colour trends together with the opposite behaviours, i.e. redder-when-brighter changes for BL Lacs (Gu & Ai 2011) and bluer-when-brighter trends for FSRQs (Wu et al. 2011), indicate that more complex scenarios, presumably involving the dominance of the relative contributions of the Doppler boosted jet emission component and the accretion disc component, respectively, are particularly relevant for blazars with peak synchrotron frequencies in the range of 1013\u201315\u2009Hz (low-frequency peaked blazars; Isler et al. 2017; Gopal-Krishna, Britzen & Wiita 2019).","Citation Text":["Bonning et al. 2012"],"Citation Start End":[[977,996]]}
{"Identifier":"2022MNRAS.510.3039K__Sakai_et_al._2018_Instance_1","Paragraph":"Finally, our models assume the absence of planetary magnetic fields. The early paradigm considering the evolution of terrestrial planets has implied that the planetary magnetic field is necessary to protect planetary atmospheres and reduce the atmospheric mass loss (see e.g. Dehant et al. 2007, and references therein). The later studies, however, show that this point of view is ambiguous. Thus, the effect of the magnetic field on the atmospheric escape can be considered as a result of the two concurring processes: reducing the escape by capturing the ionised atmospheric species within the closed magnetic field lines, and enhancing the escape of the atmospheric ions through the regions of the open magnetic lines (polar cusps, in the case of a dipole field) and the reconnection on the night-side (see, e.g. Khodachenko et al. 2015; Sakai et al. 2018; Carolan et al. 2021). Thus, for planets in the Solar System, it was shown both in the observations (Gunell et al. 2018; Ramstad & Barabash 2021) and by modelling (Sakai et al. 2018; Egan et al. 2019) that the presence of a weak magnetic field can intensify atmospheric escape. These results, however, should be taken with caution for young planets, and in particular those in the sub-Neptune range, because of the different atmospheric structures and the non-thermal mechanisms dominating the atmospheric mass loss in the Solar System, which are contrary to the planets considered in this study (see, e.g. Scherf & Lammer 2021, for the discussion). For hot Jupiters, Khodachenko et al. (2015) predict a significant suppression of escape for intrinsic magnetic fields larger than 0.3 G. The model with the closest setup to this study by Carolan et al. (2021) predicts, however, for the 0.7Mjup planet experiencing XUV (thermally) driven atmospheric escape, a small increase in the atmospheric mass-loss rate with increasing dipole field strength (about twice between 0 and 5\u2009G). We therefore expect that the possible effect from the planetary intrinsic magnetic field depends largely on the strength and configuration of the planetary and stellar magnetic fields, but, according to the numbers reported in the literature, might not affect our results dramatically. The lack of studies for close-in sub-Neptune-like planets, however, holds us from making final conclusions.","Citation Text":["Sakai et al. 2018"],"Citation Start End":[[841,858]]}
{"Identifier":"2016MNRAS.463L..79T__Perri,_Carbone_&_Veltri_2010_Instance_1","Paragraph":"The high-frequency break often found in interplanetary magnetic field power density spectra separates the fluid from the kinetic regime of fluctuations. This break is located around scales which are typical of protons kinetics like the proton inertial length \u03bbi = c\/\u03c9p and the proton Larmor radius \u03bbL = vth\/\u03a9p, where \u03c9p is the local plasma frequency while \u03a9p is the local gyro-frequency, with vth and c the thermal speed and the speed of light, respectively. Several authors tried to reproduce the location of the frequency break according to various models (Leamon et al. 1998; Perri, Carbone & Veltri 2010; Bourouaine et al. 2012) but Markovskii, Vasquez & Smith (2008) showed that none of the available model could predict a value for the frequency break in good agreement with the observations. Only recently, Bruno & Trenchi (2014), studying the radial behaviour of the frequency break location within the expanding high-speed wind, were able to conclude that the resonant condition for outward parallel propagating Alfv$\\acute{\\textrm{e}}$n\/ion-cyclotron waves (ICWs) was the mechanism able to provide the best agreement with the observations. This result, although not expected on the basis of anisotropy predictions by any turbulent cascade (Chen et al. 2014), gave new relevance to the ion-cyclotron resonance mechanism in the frame of turbulence collisionless heating without cancelling the role of other possible mechanisms that might be at work as well. Leamon et al. (1999) proposed Landau damping of obliquely propagating kinetic Alfv$\\acute{\\textrm{e}}$n waves (KAWs) resulting in a frequency break corresponding to the scale of the Larmor radius \u03bbL; for Dmitruk, Matthaeus & Seenu (2004) 2D turbulence dissipation through turbulence reconnection process and generation of current sheets of the order of the ion inertial length \u03bbi enhances the role of this scale, which is the most relevant one also in the framework of incompressible Hall magnetohydrodynamics used by Galtier (2006) to explain the break.","Citation Text":["Perri, Carbone & Veltri 2010"],"Citation Start End":[[579,607]]}
{"Identifier":"2016ApJ...829..120A__Perron_et_al._1988_Instance_2","Paragraph":"Annealing is sometimes used to avoid effects of temperature fluctuations during the etching and\/or to remove a background consisting of tracks of light ions. We did not anneal samples before etching. The reason is that nanometric structure transformations of olivine along the heavy projectile trajectory provide enhanced etching of this region. Figure 3 and simulations made in Gorbunov et al. (2015) demonstrate that the diameter of an emerging amorphized track core is up to about 10 nm in the trajectory sector where the Bragg peak of the electronic stopping of heavy ions is realized. The chemical activity of this track core may be reduced due to recrystallization during annealing. To stimulate such recrystallization within etching, the etching temperature must reach thresholds activating (a) fast diffusion of atoms\/structure defects supplying structure modifications at times much shorter than the etching time, or (b) melting of the track core followed by its rapid solidification. Because of olivines\u2019 high melting temperatures (1800\u00b0C\u20131850\u00b0C) the second scenario cannot be realized at the etching temperatures used (110\u00b0C) or during the hand-polishing of samples before etching. Such a temperature increase arising during treatments of samples cannot stimulate a fast diffusion of atoms, either, due to their high migration barriers (e.g., migration barriers of vacancies and interstitials in oxides with covalent binding exceeding 1\u20132 eV). This is well-illustrated in some experiments (Perelygin et al. 1985) when the procedure of track annealing is applied to study ancient tracks from GCR in olivine crystals from meteorites. Dissipation of \u201cbackground\u201d tracks of light iron group nuclei from GCR (initial density 1010\u20131011 cm\u22123) was detected in Perelygin & Stetsenko (1989) after annealing these crystals at higher temperatures (430 \u00b1 1)\u00b0C for 32 hr before etching, and in a 6\u20138-fold decrease of track lengths for nuclei with Z \u2265 54. This correlates with the analysis (Perron et al. 1988), which demonstrated that the preliminary track annealing led to unpredictable changes in track lengths, resulting in a lower accuracy of nuclear charge determination. For example, path length variations of accelerated Kr and Xe nuclei (with energies of 12.5 and 10.0 MeV per nucleon, respectively), decelerated in olivine crystals from Marjalahti pallasite, depend on annealing time (Lal et al. 1969). The etched lengths of tracks of these nuclei are reduced by 2\u20133 times for the first 10\u201320 hr of annealing (382\u00b0C). A further increase of annealing time (up to 240 hr) is not followed by any significant decrease in track length, but these final lengths of tracks of Kr ions vary from 18 \u00b1 3 \u03bcm to 11 \u00b1 3 \u03bcm (40% difference), i.e., the dispersion of the measured lengths is too high. Annealing of tracks of U, Au, and Xe decelerated in olivine crystals from Marjalahti pallasite at temperatures of 430\u00b0C, 435\u00b0C, and 450\u00b0C resulted in a similar distribution of etched track lengths (Perron et al. 1988). Dispersion of L values measured in individual olivine crystals from Marjalahti pallasite sometimes reaches a 3\u20134-fold value. This effect has been observed, in particular, for tracks from U ions, annealed for 5 hr at a temperature of 450\u00b0C, when the L value measured in the same crystals varied within the range of Lmin = (217 \u00b1 52) \u03bcm up to Lmax = (762 \u00b1 77) \u03bcm. The annealing of tracks from U ions held for 5 hr at T = 435\u00b0 gave the L values within the range of Lmin = (440 \u00b1 100) \u03bcm to Lmax = (869 \u00b1 53) \u03bcm (Perelygin & Stetsenko 1989). Similarly, almost twofold intervals of L variation were obtained for Xe and Au ion tracks. Taking into account these causes, the technique without preliminary annealing at a higher temperature is used in the presented work, i.e., we did not apply annealing of samples before their etching at a temperature of 110\u00b0C. The search for the heavy component in GCR within the framework of the OLIMPIYA project is based on the registration and measurement of the dynamic and geometric parameters of chemically etched tracks generated by nuclei with Z > 40 in combination with calibration experiments at heavy ion accelerator facilities. The detection method is an annealing-free technique based on layer-by-layer grinding and chemical etching. This technique provides for the geometrical parameters of tracks and the lengthwise track etching rate along the ion trace, as an additional parameter for identification of charges Z of the particle producing tracks.","Citation Text":["Perron et al. 1988"],"Citation Start End":[[2989,3007]]}
{"Identifier":"2017ApJ...851..127W__Kosteleck\u00fd_&_Russell_2011_Instance_1","Paragraph":"Because of the high-energy extent of their emission, their large cosmological distances, and their fast variabilities, gamma-ray bursts (GRBs) have been deemed as the most promising sources for searching for LIV-induced vacuum dispersion (Amelino-Camelia et al. 1998; Jacob & Piran 2007; Amelino-Camelia & Smolin 2009; Amelino-Camelia 2013; Wei et al. 2016). To date, various limits on LIV have been obtained by studying the dispersion of light in observations of individual GRBs or a large sample of GRBs (Amelino-Camelia et al. 1998; Coleman & Glashow 1999; Schaefer 1999; Boggs et al. 2004; Pavlopoulos 2005; Kahniashvili et al. 2006; Jacob & Piran 2008; Abdo et al. 2009a, 2009b; Xiao & Ma 2009; Shao et al. 2010; Chang et al. 2012; Nemiroff et al. 2012; Kosteleck\u00fd & Mewes 2013; Vasileiou et al. 2013; Zhang & Ma 2015; Xu & Ma 2016, see also Kosteleck\u00fd & Russell 2011; Liberati 2013 and summary constraints for LIV therein). Although these limits on LIV have reached high precision, most were obtained by relying on the rough time lag of a single highest-energy photon. Performing a search for LIV using the true time lags of high-quality and high-energy light curves in different energy multi-photon bands is therefore crucial. Furthermore, the method of the arrival time difference used for testing LIV is tempered by our ignorance concerning the intrinsic time delay that depends on the unknown emission mechanism of GRBs. Most previous studies concentrated on the time delay induced by LIV, while neglecting the intrinsic time delay, which would impact the reliability of the resulting constraints on LIV. Most recently, however, Wei et al. (2017a, 2017b) provided some solutions to disentangle the intrinsic time delay problem. They first proposed that only GRB 160625B, which so far is the only burst with a well-defined transition from positive to negative spectral lags,3\n\n3\nThe spectral lag is defined as the arrival time difference between high- and low-energy photons and is considered to be positive when high-energy photons precede low-energy photons.\n provides a good opportunity to distinguish the possible LIV effect from any source-intrinsic time delay in the emission of photons of different energy bands. By fitting the true multi-photon spectral lag data of GRB 160625B, they obtained both a reasonable formulation of the intrinsic energy-dependent time delay and robust limits on the QG energy scale and the Lorentz-violating coefficients of the Standard-Model Extension.","Citation Text":["Kosteleck\u00fd & Russell 2011"],"Citation Start End":[[847,872]]}
{"Identifier":"2021MNRAS.503.5658M__Gupta_&_Schlichting_2019_Instance_1","Paragraph":"The opacity at the radiative\u2013convective boundary controls the atmosphere\u2019s cooling rate, L\u221d1\/\u03ba, and as such is key to understanding the evolution of super-Earth atmospheres. As discussed in Section 2, we use a constant value of the atmosphere\u2019s opacity, \u03ba = 0.1 cm2g\u22121, throughout this work. This choice is based on the typical order of magnitude value expected for H\/He-dominated atmospheres (e.g. Freedman et al. 2008). In reality, the atmospheric opacity depends on many factors which are likely to vary over the course of super-Earth atmospheric evolution. More complex opacity models can include a power law increase in opacity with density (e.g. Gupta & Schlichting 2019), and power law increases with temperature and atmospheric metallicity in addition to density (e.g. Lee & Chiang 2015). We found that implementing an opacity dependence on density did not have a major effect on our final atmospheric masses. The strongest effect, on the critical radiative convective boundary to which a planet cools during spontaneous mass-loss, is only logarithmic (equation 18). However, if the overall composition of the atmosphere changes significantly as the super-Earths lose atmospheric mass, an increase in opacity with metallicity over time could affect our results. To test whether we expect the hydrodynamic outflow to significantly fractionate the atmosphere, thereby increasing the metallicity of the remnant atmosphere, we compute the cross-over mass, $\\mu _\\mathrm{c} = \\mu + k_\\mathrm{B} T_\\mathrm{eq} \\dot{M}\/(4 \\pi b \\mu G M_\\mathrm{c})$, where b is the binary diffusion coefficient between the two species. This is the molecular weight above which a species cannot be liberated from the planet by the hydrodynamic wind, as the drag forces cannot overcome gravity (Hunten, Pepin & Walker 1987). Using b \u223c 5 \u00d7 1017T0.75 cm\u22121s\u22121, a value typical of the binary diffusion coefficients of common secondary species in hydrogen gas (Zahnle & Kasting 1986), we find \u2013 given our mass-loss rates \u2013 the smallest value for the cross-over mass for the planets considered in this paper is about \u03bcc \u223c 103 amu. This is much larger than any potential heavy species in these atmospheres. We therefore conclude that secondary species are efficiently carried away by the outflow and that the atmosphere does not become significantly enhanced in heavy elements due to the wind itself for the planets considered in this paper.","Citation Text":["Gupta & Schlichting 2019"],"Citation Start End":[[652,676]]}
{"Identifier":"2019ApJ...886...34F__Murty_et_al._2007_Instance_1","Paragraph":"If the variation in 10Be\/9Be ratios of CAIs reflects those episodic accretion events, 10Be\/9Be ratios of CH\u2013CB CAIs observed in this study would give important constraints on the evolution of the solar protoplanetary disk. Astronomical observations suggest that FUori-type outbursts are confined to the first few hundreds of thousands of years, which correspond to the class I stage of the protoplanetary disk evolution (e.g., Schulz 2012). We propose that the high and variable 10Be\/9Be ratios recorded in CH\u2013CB CAIs reflect episodic cosmic-ray fluxes caused by FUori-type outbursts. On the other hand, relatively low and less variable 10Be\/9Be ratios recorded in CV CAIs may reflect less intensive episodic accretion events, possibly the EXori-type outbursts, which are confined to the evolutional stage of a few million years after the formation of the protoplanetary disk (=class II). Note that CH\u2013CB CAIs studied here show no (or very low) signs of 26Al-derived 26Mg excesses, while most CV CAIs show clear evidence for the past presence of 26Al. If 26Al was introduced into the solar system at the earliest stage of the disk evolution (e.g., Sahijpal & Goswami 1998), differences in Be\u2013B and Al\u2013Mg systematics between CH\u2013CB and CV CAIs imply that the injection of 26Al have occurred between the evolutionary stages class I and class II of the solar protoplanetary disk. This scenario is in agreement with arguments by other authors that the 26Al-free CAIs formed prior to injection and homogenization of 26Al in the early solar system (Sahijpal & Goswami 1998; Sahijpal et al. 2000; Krot et al. 2008a see more discussion in Krot et al. 2012a). Importantly, as mentioned in the introduction, CH\u2013CB chondrites may have accreted a significant amount of outer solar system materials (Murty et al. 2007; Ivanova et al. 2008; Briani et al. 2009; Bonal et al. 2010; Olsen et al. 2016; Van Kooten et al. 2016), suggesting that CH\u2013CB chondrites formed at outer parts of the solar protoplanetary disk relative to CV chondrites. In this case, our new Be\u2013B and Al\u2013Mg data set implies that the earliest formed CAIs tend to be transported into the outer part of the solar protoplanetary disk, where the parent bodies of CH\u2013CB chondrites likely accreted. Yang & Ciesla (2012) modeled the evolution of the protoplanetary disk and material transport in the protoplanetary disk. Interestingly, Yang & Ciesla (2012) showed that outward radial transport in class I would have been greater than that of later stages of YSO evolution, suggesting that the earliest formed CAIs could be preserved in primitive bodies that accreted in the outer part of the disk. This model is consistent with our interpretation for the Be\u2013B and Al\u2013Mg systematics on CH\u2013CB CAIs. It should be noted, however, that it is possible that 26Al were heterogeneously distributed in the CAI-forming regions at the earliest stage of the solar system evolution (e.g., Krot et al. 2008a; Holst et al. 2013; Park et al. 2017 and reference therein). Because no Pb\u2013Pb ages of CH\u2013CB CAIs are available at present, we cannot discard that possibility. Very recently, K\u00f6\u00f6p et al. (2018) found helium and neon excesses in the 26Al-free hibonite-rich CAIs, which can be attributed to in situ irradiation by energetic particles. Because 26Al-rich CAIs in CV chondrites lack comparable noble gas irradiation records (Vogel et al. 2004), K\u00f6\u00f6p et al. (2018) concluded that 26Al-free hibonite-rich CAIs experienced intense energetic particle irradiation at the earliest stage of solar protoplanetary disk evolution. This conclusion seems to be consistent with our above scenario for 26Al-free CH\u2013CB CAIs. Note, however, that 10Be\/9Be ratios of 26Al-free hibonite-rich CAIs in CM chondrites tend to be in the range of those for 26Al-rich CV CAIs (Liu et al. 2009, 2010), which is inconsistent with the above scenario. Therefore, the relationship between 10Be and 26Al in the early solar system would be more complicated than we thought.","Citation Text":["Murty et al. 2007"],"Citation Start End":[[1786,1803]]}
{"Identifier":"2019MNRAS.490.1870L__Guo_et_al._2015_Instance_1","Paragraph":"In principle, we could measure the multipole moments in a given simulation by directly populating dark matter haloes in the simulation with galaxies. However, this is computationally very expensive and comes with realization noise due to the random number and phase-space positions of galaxies. Instead, we use a tabulation method (Neistein & Khochfar 2012; Reid et al. 2014; Zheng & Guo 2016) to speed up the computation dramatically and eliminate any realization noise. We first take all haloes in a given simulation to serve as tracers of central galaxies. We furthermore assign to each halo of mass M a Poisson number of satellite tracers with expectation value $3 \\times (M\/10^{13} \\, h^{-1}\\rm M_\\odot)$. This expectation value is chosen to be significantly larger than the number of satellites we typically expect in haloes of that mass (see e.g. Guo et al. 2015). We then bin all haloes and their central and satellite tracers by halo mass and whether the concentration is above or below the median. Next, we measure all cross- and autocorrelation multipole moments between all tracers in each bin. One can then show that an estimate for the galaxy number density and the multipole moments of any arbitrary galaxy\u2013halo model are given by\n(13)$$\\begin{eqnarray*}\r\n\\hat{n}_{\\rm gal} = \\sum \\limits _{i = {\\rm c}, {\\rm s}} \\sum \\limits _{k = 1}^{n_{\\rm bins}} N_{{\\rm h}, k} \\langle N_i | M_k, c_k \\rangle\r\n\\end{eqnarray*}$$and\n(14)$$\\begin{eqnarray*}\r\n\\hat{\\xi }_\\ell &=& \\hat{n}_{\\rm gal}^{-2} \\sum \\limits _{i = {\\rm c}, {\\rm s}} \\sum \\limits _{j = {\\rm c}, {\\rm s}} \\sum \\limits _{k = 1}^{n_{\\rm bins}} \\sum \\limits _{l = 1}^{n_{\\rm bins}} \\Big [ N_{{\\rm h}, k} N_{{\\rm h}, l} \\langle N_i | M_k, c_k \\rangle \\nonumber \\\\\r\n&&\\times \\,\\langle N_j | M_l, c_l \\rangle \\xi _{\\ell , kl}^{ij} \\Big ] \\, ,\r\n\\end{eqnarray*}$$respectively. In the above expression, Nh, k denotes the number of haloes in bin k and, for example, $\\xi _{\\ell , kl}^{{\\rm c}{\\rm s}}$ denotes the multipole moments between centrals in bin k and satellites in bin l. The above estimate $\\hat{\\xi }_\\ell$ approaches the expectation value of \u03be\u2113 for sufficiently small halo mass bins. We find that 100 logarithmic bins in halo mass is sufficient to adequately sample all haloes with mass $M \\gt 3.52 \\times 10^{13} (\\Omega _{\\rm m}\/0.3) \\, h^{-1}\\, \\rm M_\\odot$ (corresponding to 100 particles). With such a bin width of \u223c0.03\u2009dex, any biases in \u03be are less than $5{{\\ \\rm per\\ cent}}$ of the observational uncertainty for a BOSS CMASS-like sample (see Guo et al. 2015). The above method only works for a fixed value of the satellite radial profile parameter \u03b7. In practice, it suffices to tabulate correlation function for bins in \u03b7 of \u0394log\u2009\u03b7 = 0.1 and linearly interpolate between them.","Citation Text":["Guo et al. 2015"],"Citation Start End":[[854,869]]}
{"Identifier":"2021AandA...649A..84H__Mulders_et_al._2013_Instance_1","Paragraph":"HD 100546 is one of the closest very well studied Herbig Be stars (d = 110 \u00b1 4 pc, Gaia Collaboration 2018). It shows clear evidence of a large flared disk, and based on images obtained with the Hubble Space Telescope (HST) and ground-based high-contrast images, an elliptical structure was detected that extends up to 350\u2013380 au (Augereau et al. 2001). Furthermore, multiple-armed spiral patterns were identified as well (Grady et al. 2001; Ardila et al. 2007; Boccaletti et al. 2013; Avenhaus et al. 2014). The disk position angle measurements found in the literature range from ~ 130 to 160\u00b0 (Grady et al. 2001; Pantin et al. 2000; Ardila et al. 2007; Pani\u0107 et al. 2014; Augereau et al. 2001; Avenhaus et al. 2014), and the inclination of the disk is smaller than 50\u00b0 (Avenhaus et al. 2014). An inner dust disk extending from ~0.2 to ~ 1\u22124 au was resolved using near-IR interferometry (Benisty et al. 2010; Tatulli et al. 2011; Mulders et al. 2013; Pani\u0107 et al. 2014). The (pre-)transitional nature of HD 100546 was initially proposed by Bouwman et al. (2003) based on a spectral energy distribution (SED) analysis. The presence of a gap extending up to 10\u201315 au has been confirmed by mid-IR interferometry (Liu et al. 2003; Pani\u0107 et al. 2014), spectroscopy in the UV and near-IR (Grady et al. 2005; Brittain et al. 2009; van der Plas et al. 2009), and high-resolution polarimetric imaging in the optical and near-IR (Avenhaus et al. 2014; Quanz et al. 2015; Garufi et al. 2016; Follette et al. 2017). Recent ALMA observations reveal an asymmetric ring between ~ 20\u221240 au with largely optically thin dust emission (Pineda et al. 2019). A central compact emission is also detected, which arises from the inner central disk, which given its mass and the accretion rate onto the star, must be replenished with material from the outer disk (Pineda et al. 2019). Miley et al. (2019) presented the first detection of C18 O in this disk, which spatially coincides with the spiral arms, and derived a lower-limit on the total gas mass (around 1% of the stellar mass) and a gas-to-dust mass ratio in the disk of ~ 20 assuming interstellar medium (ISM) abundances of C18O relative to H2. HD 100546 is also one of the few cases where protoplanet candidates have been suggested (e.g. Quanz et al. 2015; Currie et al. 2015), although this is still being debated (e.g. Follette et al. 2017; Rameau et al. 2017; P\u00e9rez et al. 2020). Recently, P\u00e9rez et al. (2020) detected a compact 1.3 mm continuum dust emission source that lies in the middle of the HD 100546 cavity (0.051\u2032\u2032 from the central star). This is compatible with circumplanetary disk emission.","Citation Text":["Mulders et al. 2013"],"Citation Start End":[[931,950]]}
{"Identifier":"2021MNRAS.503.2380S__Duffy_et_al._2010_Instance_1","Paragraph":"In summary, our constraints on the elliptical galaxy structure in the context of previous observations and simulations are consistent with the following formation and evolution scenario for massive elliptical galaxies. The first stage of the formation of elliptical galaxies is through dissipational processes at z \u2273 2, when most of their present day stars are formed. The dissipation leads to contraction in the dark matter halo and the resultant total density profile observed in cosmological numerical simulations is steeper than the isothermal case (e.g. Gnedin et al. 2004; Naab et al. 2007; Duffy et al. 2010). After z \u2248 2, the growth of the elliptical galaxies is dominated by gas-poor dissipation-less mergers, explaining their growth in size without the addition of younger stellar populations (e.g. Newman et al. 2012; Nipoti et al. 2012). This growth mechanisms is consistent with the Reff\u2013fdm correlation observed in our sample. Multiple dissipation-less mergers decrease the total density profile of the galaxies to bring it close to isothermal, and increase the half-mass radius and the central dark matter fraction with decreasing redshift (Tortora et al. 2014). Furthermore, AGN feedback expands back the contracted dark matter haloes (Martizzi, Teyssier & Moore 2013; Peirani et al. 2019), which is supported by the slope of our observed $\\gamma ^{\\rm LD}_{[0.4,\\ 4]R_{\\rm eff}}$\u2013$f_{\\rm dm}^{\\rm 3D}$ distribution. Dynamical heating from accretion may also play a role in expanding the dark matter haloes in addition to the AGN feedback, however we do not find any indication either in favour or against the presence of dynamical heating in our sample. Gas-rich dissipational mergers can also happen at some point of the galaxies\u2019 growth history. However, the small intrinsic scatter in the halo response parameter (\u03c3\u03bd \u2272 0.1) around the mean \u03bc\u03bd \u223c 0 in our sample, and the ages of the stellar populations (Thomas et al. 2005), indicate that such gas-rich mergers are relatively rare, even though their contribution cannot be ruled out or confirmed conclusively given the present precision of numerical simulations and observations (see, e.g. Sonnenfeld et al. 2014; Remus et al. 2017; Xu et al. 2017, for discussion).","Citation Text":["Duffy et al. 2010"],"Citation Start End":[[597,614]]}
{"Identifier":"2018MNRAS.480.1819R__Ricci_et_al._2017a_Instance_1","Paragraph":"A relation between the photon index and the Eddington ratio has been reported by several authors over the past two decades (e.g. Brandt, Mathur & Elvis 1997; Shemmer et al. 2006, 2008; Risaliti, Young & Elvis 2009; Brightman et al. 2013, 2016; Fanali et al. 2013; Kawamuro et al. 2016), which have shown that, for increasing \u03bbEdd, the X-ray continuum tend to be steeper. Most of these works have found that the correlation\n(7)\r\n\\begin{equation*}\r\n\\Gamma =\\psi \\log \\lambda _{\\rm Edd}+ \\omega\r\n\\end{equation*}\r\nhas a slope \u03c8 \u223c 0.3 (e.g. Shemmer et al. 2008; Brightman et al. 2013), while a steeper slope (\u03c8 \u2243 0.6) was reported by Risaliti et al. (2009), who studied SDSS quasars with archival XMM\u2013Newton observations. More recently, Trakhtenbrot et al. (2017), using BASS, found instead a significantly weaker and flatter (\u03c8 \u2243 0.15) correlation when using \u0393 obtained by considering complex spectral models (see Ricci et al. 2017a for details). Interestingly, when using \u0393 obtained by fitting the spectra of unobscured AGNs with a simple power law model in the 2\u201310\u2009keV range, Trakhtenbrot et al. (2017) found a slope similar (\u03d5 = 0.30 \u00b1 0.09) to that reported by previous studies. The existence of a relation between \u0393 and \u03bbEdd has been confirmed by repeated observations of individual sources, which have shown that the photon index increases with the flux (e.g. Perola et al. 1986; Matsuoka et al. 1990; Lamer et al. 2003; Sobolewska & Papadakis 2009). Interestingly, Sobolewska & Papadakis (2009) found that \u03c8 differs from object to object, varying from \u2243 0.10 to \u2243 0.30, and that the slope for the average spectral slope versus the average Eddington ratio is \u03c8 = 0.08 \u00b1 0.02. This slope is consistent with that found for BASS by Trakhtenbrot et al. (2017), and with the value reported by Ricci et al. (2013; \u03c8 = 0.12 \u00b1 0.04) for a sample of 36 nearby AGNs, considering the average \u0393 and \u03bbEdd. The difference between the slopes found by the works reported above is likely related to the approach used for the spectral fitting (i.e. a simple power-law model or more complex models), to the energy band, and to the sample used.","Citation Text":["Ricci et al. 2017a"],"Citation Start End":[[910,928]]}
{"Identifier":"2019MNRAS.486.4671M__Zhang_et_al._2001_Instance_1","Paragraph":"CMEs are known for large-scale expulsion of magnetized plasma structures from closed magnetic field regions on the Sun. They were first detected in the coronagraphic images taken in 1971 by NASA\u2019s OSO-7 spacecraft (Tousey 1973). However, some definite inferences for the solar wind (Eddington 1910; Birkeland 1916; Biermann 1951) as well as CMEs from the Sun (Chapman & Ferraro 1931; Eddy 1974) were made decades before their formal discovery. Following OSO-7, a series of spacecraft (Skylab, Helios, P78-1 Solwind, SOHO, Coriolis, and STEREO, etc.) have observed thousands of CMEs leading to a vast literature (Munro et al. 1979; Howard et al. 1985; Gosling 1993; Hundhausen 1999; Gopalswamy et al. 2000; Schwenn 2006; Vourlidas et al. 2010; Chen 2011; Wang et al. 2011; Webb & Howard 2012; Mishra & Srivastava 2013; Mishra et al. 2017; Harrison et al. 2018). CMEs have been observed to occur often having spatial and temporal relation with solar flares, eruptive prominences (Munro et al. 1979; Webb & Hundhausen 1987; Zhang et al. 2001; Gopalswamy et al. 2003) and with helmet streamer disruptions (Dryer 1996). Unlike CMEs from the Sun, to observe stellar CMEs are challenging because the close stellar environment cannot be spatially resolved. Although stellar CMEs have not yet been directly detected in Thomson-scattered optical light from other stars, it is believed that the extreme X-ray flares observed on stars may be in conjunction with extreme stellar CMEs (Houdebine, Foing & Rodono 1990; Wheatley 1998; Leitzinger et al. 2011; Aarnio, Matt & Stassun 2012; Osten & Wolk 2015; Vida et al. 2016). Indeed, the stellar X-ray flare, helmet streamers, and prominences observed on T Tauri Stars have shown similarities with those observed on the Sun (Haisch, Antunes & Schmitt 1995; Massi et al. 2008). The CMEs and flares themselves may not be causally related, they both seem to be involved with the reconfiguration of complex magnetic field lines within the corona caused by the same underlying physical processes, e.g. magnetic reconnection (Priest & Forbes 2002; Compagnino, Romano & Zuccarello 2017). But, even for the sun, it has been noted that not all flares are accompanied by CMEs and not all CMEs by flares (Munro et al. 1979; Harrison 1995; Yashiro et al. 2008b; Wang & Zhang 2008).","Citation Text":["Zhang et al. 2001"],"Citation Start End":[[1021,1038]]}
{"Identifier":"2020MNRAS.494.5396B__Cappellari_&_Emsellem_2004_Instance_1","Paragraph":"The 1D SDSS spectrum of SDSS J124754.95-033738.6 (Fig. 1) is dominated by A-type stars and shows strong Balmer absorption lines, which is a clear signature of post starburst E+A galaxies. We used the 1D SDSS spectrum to measure the equivalent width (EW) of the H\u03b4 absorption line and found EW(H\u03b4) = 6.7\u2009\u00c5, which is above the typical 5\u2009\u2009\u00c5\u2009threshold used to select post-starburst E+A galaxies (e.g. Goto 2007; Alatalo et al. 2016). We then fitted a stellar population synthesis model using the python implementation of Penalized Pixel-Fitting stellar kinematics extraction code (pPXF; Cappellari 2012). pPXF is a public code for extraction of the stellar kinematics and stellar population from absorption line spectra of galaxies (Cappellari & Emsellem 2004). Its output includes the best-fitting stellar model, the relative contribution of stars with different ages, the stellar velocity dispersion, and the dust reddening towards the stars (assuming a Calzetti et al. 2000 extinction law). The code uses the MILES library, which contains single stellar population synthesis models that cover the entire optical wavelength range with an FWHM resolution of 2.3\u2009\u00c5 \u2009(Vazdekis et al. 2010). We used models produced with the Padova 2000 stellar isochrones assuming a Chabrier initial mass function (IMF; Chabrier 2003). The stellar ages range from 0.03 to 14\u2009Gyr, thus allowing the analysis of systems with different star formation histories. The best-fitting stellar model is marked with red in Fig. 1. Its stellar age distribution consists of two star formation episodes, with a recent short episode that started 70\u2009Myr ago and was quenched 30\u2009Myr ago, and an older long episode that started 14\u2009Gyr ago and ended 6\u2009Gyr ago. The older episode is poorly constrained and we find similarly reasonable fits when forcing the code to use templates with different ages within the range 3\u201310\u2009Gyr. According to the best-fitting model, the stellar velocity dispersion is 185\u2009km\u2009sec\u22121 and the dust reddening towards the stars is EB\u2013V = 0.188\u2009mag. The total stellar mass is $\\mathrm{M_{*}} = 10^{10.8}\\, M_{\\odot }$, which is consistent with other estimates (e.g. Chen et al. 2012). About \u223c2 per\u2009cent of this mass was formed during the recent burst.","Citation Text":["Cappellari & Emsellem 2004"],"Citation Start End":[[729,755]]}
{"Identifier":"2017MNRAS.464.3597L__Lind_et_al._2011_Instance_2","Paragraph":"It is important to emphasize that any estimate of [Na\/Fe] should be actually taken as an upper limit. Note that we can compute the theoretical response of stellar spectra to [Na\/Fe] only for stars hotter than 3500 K (see Section 3.2). As discussed in Appendix A, a linear extrapolation of Na responses to cooler temperatures would likely lead to lower [Na\/Fe] estimates, by about 0.1 dex. Moreover, as already noticed by CvD12b, atomic NaI transitions in the atmospheres of late-type stars are prone to substantial departures from LTE (Bruls, Rutten & Shchukina 1992; Gehren et al. 2006; Andrievsky al. 2007; Lind et al. 2011). In a Sun-like star, LTE calculations predict weaker lines (Allende Prieto et al. 2003) requiring corrections for the strongest lines, which can be as high as an effective change in the abundance of \u223c0.5 dex. For a lower temperature star (with Teff \u223c 4000 K), more relevant for models having old ages (as those of our sample of ETGs), we may expect NLTE corrections in the range of 0.1\u20130.2 dex (based on fig. 4 of Lind et al. 2011). Hence, we may expect that NLTE models could enhance the predicted absorption in Na indices, resulting in lower inferred values of [Na\/Fe]. However, note that our LTE-based methodology can match NaD, $\\rm Na\\,\\small {I}8190$, $\\rm Na\\,\\small {I}1.14$, and $\\rm Na\\,\\small {I}2.21$, simultaneously, suggesting that NLTE corrections should be approximately the same for all four Na lines \u2013 an important aspect to test with future models. We emphasize that although we have introduced four free-fitting parameters to match the four Na lines (the $\\alpha _{{\\rm Na}_j}$ constants; see equation 1), in practice, the values of \u03b1NaD, $\\rm \\alpha _{{\\rm Na}\\,\\small {I}1.14}$ and $\\rm \\alpha _{{\\rm Na}\\,\\small {I}2.21}$ are fully consistent with those of \u03b1-MILES and CvD12a models. Effectively, we are able to fit the four Na-sensitive line strengths of the seven X-Shooter spectra \u2013 spanning a range of age, metallicity, and [\u03b1\/Fe] \u2013 based on only one \u2018extra\u2019 free-fitting parameter (i.e. the $\\rm \\alpha _{{\\rm Na}\\,\\small {I}8190}$).","Citation Text":["Lind et al. 2011"],"Citation Start End":[[1041,1057]]}
{"Identifier":"2021ApJ...909..175Y__Buzzicotti_et_al._2018_Instance_2","Paragraph":"The filtered MHD equations read\n11\n\n\n\n\n\n\n\n12\n\n\n\n\n\nwhere we sum over repeated indices, and\n13\n\n\n\n\n\n\n\n14\n\n\n\n\n\n\n\n15\n\n\n\n\n\n\n\n16\n\n\n\n\n\ndenote the inertial (I), Maxwell (M), advective (A), and dynamo (D) subfilter-scale stresses, respectively. Despite their common origin through the electric field in the induction equation, we here treat \n\n\n\n\n\n and \n\n\n\n\n\n separately, in order to disentangle the effects of magnetic-field-line advection, encoded in \n\n\n\n\n\n, and magnetic-field-line stretching, encoded in \n\n\n\n\n\n. Usually, the magnetic subscale stress refers to the difference \n\n\n\n\n\n (Aluie 2017; Offermans et al. 2018). Equations (11) and (12) differ from expressions for the filtered MHD equations found elsewhere by an additional projection of the coupling terms. The latter ensures that the dynamics defined by Equations (11) and (12) are confined to the same finite-dimensional subspace \u03a9\u2113 of the original domain \u03a9 (Buzzicotti et al. 2018; Offermans et al. 2018). At first sight, this formulation suggests that the corresponding evolution equations for kinetic and magnetic energy feature terms are not Galilean invariant, which ought to be avoided as the measured subfilter-scale energy transfers otherwise include unphysical fluctuations (Aluie & Eyink 2009a, 2009b; Buzzicotti et al. 2018). However, the energy balance equations can be expressed in an alternative way by including terms that vanish under spatial averaging and ensure Galilean invariance of all terms (Buzzicotti et al. 2018; Offermans et al. 2018). For a statistically stationary evolution, the spatiotemporally averaged energy budget can then be written as\n17\n\n\n\n\n\n\n\n18\n\n\n\n\n\nwhere \n\n\n\n\n\n and \n\n\n\n\n\n are the filtered kinetic and magnetic dissipation rates, respectively, and \n\n\n\n\n\n are terms that convert kinetic to magnetic energy \n\n\n\n\n\n and vice versa \n\n\n\n\n\n, and\n19\n\n\n\n\n\n\n\n20\n\n\n\n\n\n\n\n21\n\n\n\n\n\n\n\n22\n\n\n\n\n\ndenote the four proper energy fluxes, in the sense that they vanish in the limit \u2113 \u2192 0, as can be seen from Equations (13)\u2013(16). If positive, the inertial and Maxwell fluxes, \n\n\n\n\n\n and \n\n\n\n\n\n transfer kinetic energy from scales larger than or equal to \u2113 to scales smaller than \u2113 and vice versa if negative, while the advective and dynamo fluxes, \n\n\n\n\n\n and \n\n\n\n\n\n, do so with magnetic energy. Note that there is no interscale energy conversion as the conversion terms \n\n\n\n\n\n and \n\n\n\n\n\n only involve filtered fields; as such they are known as resolved-scale conversion terms (Aluie 2017). The total energy flux is then given by the sum\n23\n\n\n\n\n\n\n","Citation Text":["Buzzicotti et al. 2018"],"Citation Start End":[[1266,1288]]}
{"Identifier":"2022AandA...659A.125S__Gaspari_et_al._2020_Instance_1","Paragraph":"The ability both to rapidly quench and maintain low levels of star formation in red-sequence massive galaxies has often been attributed to the impact of active galactic nuclei (AGN). Radiatively-driven winds, mechanical energy, and\/or thermal pressure release from the AGN are physical processes routinely incorporated into numerical simulations of galaxy evolution to solve the overcooling problem (e.g., Vogelsberger et al. 2014a; Schaye et al. 2015; Pillepich et al. 2018; Bassini et al. 2019; Wittor & Gaspari 2020), so that the characteristics of the observed galaxy population can be reproduced (e.g., Vogelsberger et al. 2014b; Nelson et al. 2018). Often two types of the AGN feedback mode are being distinguished, which are called the quasar mode or radiative mode (e.g., Di Matteo et al. 2005; Hopkins et al. 2008; Hopkins & Elvis 2010) and the radio mode also known as kinetic or maintenance mode (e.g., Bower et al. 2006; Croton et al. 2006; Fabian 2012; Gaspari et al. 2020). The quasar mode is thought to initially quench star formation through powerful radiatively driven winds that remove the gas from the galaxy (e.g., Nesvadba et al. 2008; Feruglio et al. 2010; Maiolino et al. 2012) and thereby suppress star formation. The radio mode is a heating mechanism affecting the gaseous halo that prevents continuous condensation of cold and warm gas, thereby quenching star formation over long timescales and maintaining a low level of star formation in red sequence galaxies (e.g., Br\u00fcggen & Kaiser 2002; McNamara et al. 2005; Gaspari et al. 2019; McDonald et al. 2021). Strong observational evidence for radio mode feedback has been collected from galaxy clusters hosting powerful radio galaxies (e.g., Fabian et al. 2006; McNamara & Nulsen 2007). The role of quasar mode feedback is much more controversial from an observational perspective. Various studies have reported negative (Ho 2005; Nandra et al. 2007; Schawinski et al. 2009; Farrah et al. 2012; Page et al. 2012; Mullaney et al. 2015; Shimizu et al. 2015; Wylezalek & Zakamska 2016; Kakkad et al. 2017; Catal\u00e1n-Torrecilla et al. 2017; Bing et al. 2019; Bluck et al. 2020; Ramos Padilla et al. 2020; Brownson et al. 2020; Smith et al. 2020), positive (Kim et al. 2006; Cresci et al. 2015a,b; Bernhard et al. 2016; Santoro et al. 2016; Maiolino et al. 2017; Koss et al. 2021), or no effect from feedback (Elbaz et al. 2011; Bongiorno et al. 2012; Harrison et al. 2012; Husemann et al. 2014; Balmaverde et al. 2016; Leung et al. 2017; Woo et al. 2017; Shangguan et al. 2018; Scholtz et al. 2020) on the star formation rates of observed AGN host galaxies.","Citation Text":["Gaspari et al. 2020"],"Citation Start End":[[966,985]]}
{"Identifier":"2020ApJ...903..105Z__Giacalone_2013_Instance_1","Paragraph":"Unveiling the transport properties of charged particles in the presence of turbulent magnetic fields represents the key to understanding the propagation and acceleration of energetic particles in many space and astrophysical plasmas. Several parameters can influence the particle transport, like the turbulence level, the turbulence anisotropy, the spectral shape, and the particle energy, and many different regimes are considered in the literature (Zimbardo 2005; Pommois et al. 2007; Shalchi 2010, 2015; Zimbardo et al. 2012; Lazarian & Yan 2014; Litvinenko et al. 2017). Besides the ballistic and diffusive regimes, many observations suggest that intermediate, so-called superdiffusive regimes for which \n\n\n\n\n\n with \u03b1 > 1, may exist. Here, r is the distance from the starting point, t is the time, and \u03b1 is the anomalous diffusion exponent. These observations range from the analysis of nonrelativistic electrons accelerated in solar energetic events (Lin 1974), to the anomalously large parallel mean free paths inferred for solar energetic protons (Reames 1999) and for termination shock accelerated ions (Giacalone 2013). On a larger, astrophysical scale, Ragot & Kirk (1997) showed, by studying the spectral index evolution of diffuse radio emission from synchrotron emitting electrons in the Coma cluster of galaxies, that this can be interpreted in terms of superdiffusive transport. The transport properties can also be deduced by studying the energetic particle profiles upstream of shock waves, where particles are thought to be accelerated. In particular, the observations of power-law energetic particle profiles upstream of shock waves in the heliosphere (Perri & Zimbardo 2007, 2008, 2009, 2015; Sugiyama & Shiota 2011; Zimbardo et al. 2012) and in the supernova remnant shock precursors (Perri et al. 2016; Perri 2018) are an indication of superdiffusive transport. Further, superdiffusive transport also influences processes like diffusive shock acceleration (Kirk et al. 1996; Ragot & Kirk 1997; Perri & Zimbardo 2012a; Zimbardo & Perri 2013; Bykov et al. 2017) and stochastic Fermi acceleration (Sioulas et al. 2020), leading to modifications of the energy spectral index and of the acceleration time. As an application of superdiffusive shock acceleration (Perri & Zimbardo 2012a), we may mention that a discrepancy is found in the Mach numbers derived by X-rays and by radio observations of galaxy clusters\u2019 merger shocks (Akamatsu et al. 2017; Colafrancesco et al. 2017; Van Weeren et al. 2017); this discrepancy could be explained by the modification of the energy spectral index implied by superdiffusive shock acceleration (Zimbardo & Perri 2017, 2018).","Citation Text":["Giacalone 2013"],"Citation Start End":[[1112,1126]]}
{"Identifier":"2017ApJ...840...71K__Eracleous_et_al._2012_Instance_1","Paragraph":"The recoiling supermassive black hole (rSMBH) carries it the broad-line region (BLR) along with it and leaves the stellar nucleus behind; it can be observable for tens of Myr as an offset AGN (Madau & Quataert 2004; Loeb 2007, Blecha et al. 2011, 2016). We therefore expect two observational characteristics: (i) the rSMBH could be observed spatially offset with respect to the stellar center of the host galaxy, and (ii) the broad emission lines could have a measurable velocity offset with respect to the systemic velocity. So far, there have been more than a dozen reports of rSMBH candidates. Some examples of rSMBH candidates are SDSS J092712.65+294344.0 (Komossa et al. 2008), SDSS J105041.35+345631.3 (Shields et al. 2009), M87 (Batcheldor et al. 2010), QSO E1821+643 (Robinson et al. 2010), CXOC J100043.1+020637 (Civano et al. 2010), CXO J122518.6+144545 (Jonker et al. 2010), a half-dozen SDSS QSOs (Eracleous et al. 2012), 10 nearby core elliptical galaxies (Lena et al. 2014), NGC 3115 (Menezes et al. 2014), 5 SDSS AGNs (Comerford et al. 2015), and 26 SDSS QSOs (Kim et al. 2016). Except for CXOC J100043.1+020637, most of the rSMBH candidates described to date either have a spatial offset or a velocity offset, but not both. The source CXOC J100043.1+020637 has two compact sources separated by \u223c2.5 kpc and has a velocity offset of \u223c1200 km s\u22121 (Civano et al. 2010). Chandra observations find that the southeastern source, which has a point-like morphology typical of a bright AGN, is responsible for the whole X-ray emission in this system (Civano et al. 2012). The northwestern source has a more extended profile in the optical band, with a scale length of \u223c0.5 kpc. Recent 3 GHz Karl G. Jansky Very Large Array (VLA) observations find that the entire observed 3 GHz radio emission can be associated with the southeastern nucleus (Novak et al. 2015). The observations favored an explanation within the rSMBH picture, but the presence of an obscured and radio-quiet SMBH in the northwestern source is not ruled out (i.e., Blecha et al. 2013; Wrobel et al. 2014).","Citation Text":["Eracleous et al. 2012"],"Citation Start End":[[910,931]]}
{"Identifier":"2019MNRAS.482.1858Y__Sharov_et_al._2017_Instance_1","Paragraph":"The non-gravitational interaction in the dark sector, precisely between dark matter and dark energy is a phenomenological concept that was originally thought to explain the different values of the time-independent cosmological constant (Wetterich 1995), but later on, such concept was found to be very useful to explain the cosmic coincidence problem (Amendola 2000; Chimento et al. 2003; Cai & Wang 2005; Hu & Ling 2006; del Campo, Herrera & Pav\u00f3n 2008, 2009). Certainly, this led to a large amount of investigations towards this direction where the dark sectors have direct interaction (Billyard & Coley 2000; Barrow & Mota 2003; Barrow & Clifton 2006; Skordis et al. 2006; Koivisto & Mota 2008b; Quartin et al. 2008; Chimento 2010; Valiviita, Maartens & Majerotto 2010; Koivisto, Mota & Zumalacarregui 2012; Thorsrud, Mota & Hervik 2012; Akrami et al. 2013; Faraoni, Dent & Saridakis 2014; Yang & Xu 2014a,b,c; Duniya, Bertacca & Maartens 2015; Pan, Bhattacharya & Chakraborty 2015; Valiviita & Palmgren 2015; Mukherjee & Banerjee 2017; Pan & Sharov 2017; Santos et al. 2017; Sharov et al. 2017; Pan, Mukherjee & Banerjee 2018; also see Pan & Chakraborty 2013; Chen et al. 2014; Pan & Chakraborty 2014; Shahalam et al. 2015, 2017; Cai, Tamanini & Yang 2017; Kumar & Nunes 2017a; Odintsov, Oikonomou & Tretyakov 2017). Such interacting scenarios have good motivation if the particle physics theory is considered because from the particle physics view, mutual interaction between any two fields is a natural phenomenon, irrespective of the nature of the fields. Although the interacting dynamics is complicated and a generalized cosmic scenario, it recovers the non-interaction cosmology as a special case. Thus, the theory of non-gravitational interaction between dark matter and dark energy is a generalized version of the non-interacting dark matter and dark energy cosmologies. Interestingly enough, the observational data at recent time found that the direct interaction between dark matter and dark energy cannot be excluded (Salvatelli et al. 2014; Kumar & Nunes 2016; Nunes, Pan & Saridakis 2016; Yang et al. 2016; Kumar & Nunes 2017b; van de Bruck, Mifsud & Morrice 2017; Yang, Banerjee & Pan 2017a; Yang, Pan & Mota 2017b; Yang, Pan & Barrow 2018). Moreover, very recently, it has been reported that the current tension on the local Hubble constant can be alleviated with the introduction of dark matter and dark energy interactions (Kumar & Nunes 2016; Di Valentino, Melchiorri & Mena 2017). Additionally, the crossing of phantom barrier has also been found to be an easy consequence of the non-gravitational interaction. Thus, the theory of interacting dark energy might be considered to be an appealing field of research and indeed a hot topic for the next generation of the astronomical surveys.","Citation Text":["Sharov et al. 2017"],"Citation Start End":[[1079,1097]]}
{"Identifier":"2018AandA...620A..80M__Coutens_et_al._2018_Instance_1","Paragraph":"The comparison with other hot corinos is not easy since we have a limited number of lines, in contrast to well-known sources that have been more extensively observed with ALMA and NOEMA, such as IRAS 16293-2422 (J\u00f8rgensen et al. 2016), NGC1333 IRAS 2A and 4A (Taquet et al. 2015; L\u00f3pez-Sepulcre et al. 2017). In general, we find a similar inventory ofCOMs but with lower abundances in B1b-S. We are going to discuss only some trends and ratios between them. For example, one similarity with hot corinos are the higher abundances of O-bearing COMs with respect to N-bearing COMs. We observe a similar trend in B1b-S, where the highest abundances are obtained for CH3OCOH, CH3OCH3, and CH3CHO, while NH2CN and NH2CHO show the lowest abundances. Cyanamide has been recently detected towards the low-mass protostars IRAS 16293-2422 and NGC 1333 IRAS2A (Coutens et al. 2018), with observed NH2CN to NH2CHO ratios of 0.2and 0.02, respectively. These values are in the range of those observed towards the molecular clouds in the Galactic centre but lower than in Orion KL (Coutens et al. 2018). We obtain in B1b-S anabundance ratio of 0.25, similar to IRAS 16293-2422. Of the three possible isomers of C2 H4O2, we have detected CH3OCOH and CH2OHCHO. The observed CH3OCOH to CH2OHCHO ratio in B1b-S, ~20, is similar to that observed in the low-mass protostars in NGC 1333 and IRAS 16293-2422 (Taquet et al. 2015; J\u00f8rgensen et al. 2012). Acetic acid (CH3COOH) is the most stable but the least abundant of the three isomers (Lattelais et al. 2010). It has been observed in IRAS 16293-2422 with a ratio with respect to glycolaldehyde of ~11 (J\u00f8rgensen et al. 2016), consistent with the 5\u201315 upper limits in B1b-S. Glycolaldehyde and its corresponding alcohol, ethylene glycol, show similar abundances in B1b-S, slightly higher for CH2OHCHO. This is in contrast to other hot cores and hot corinos (Fuente et al. 2014; J\u00f8rgensen et al. 2016; Favre et al. 2017). However, since detections of these species are based on just a few lines, in particular for glycolaldehyde, column densities may be not well constrained and it is difficult to make definite conclusions. We have detected CH3CH2OCOH for the first time in a low-mass protostar, with a relatively high abundance of 10\u221211. This detection shows that the molecular complexity is high in young hot corinos and other species with a larger number of atoms could be present. It is noteworthy that this species has not been previously observed in the well-known hot corinos IRAS 16293-2422 and NGC 133IRAS4A. We have checked in the ALMA spectral line survey of IRAS 16293-2422 (PILS, J\u00f8rgensen et al. 2016), and no clear features are seen, although the observed frequencies are different. It is possible that the high level of line confusion and wide linewidths prevent the detection of weak lines in these more evolved sources.","Citation Text":["Coutens et al. 2018"],"Citation Start End":[[849,868]]}
{"Identifier":"2021ApJ...909..175Y__Buzzicotti_et_al._2018_Instance_1","Paragraph":"The filtered MHD equations read\n11\n\n\n\n\n\n\n\n12\n\n\n\n\n\nwhere we sum over repeated indices, and\n13\n\n\n\n\n\n\n\n14\n\n\n\n\n\n\n\n15\n\n\n\n\n\n\n\n16\n\n\n\n\n\ndenote the inertial (I), Maxwell (M), advective (A), and dynamo (D) subfilter-scale stresses, respectively. Despite their common origin through the electric field in the induction equation, we here treat \n\n\n\n\n\n and \n\n\n\n\n\n separately, in order to disentangle the effects of magnetic-field-line advection, encoded in \n\n\n\n\n\n, and magnetic-field-line stretching, encoded in \n\n\n\n\n\n. Usually, the magnetic subscale stress refers to the difference \n\n\n\n\n\n (Aluie 2017; Offermans et al. 2018). Equations (11) and (12) differ from expressions for the filtered MHD equations found elsewhere by an additional projection of the coupling terms. The latter ensures that the dynamics defined by Equations (11) and (12) are confined to the same finite-dimensional subspace \u03a9\u2113 of the original domain \u03a9 (Buzzicotti et al. 2018; Offermans et al. 2018). At first sight, this formulation suggests that the corresponding evolution equations for kinetic and magnetic energy feature terms are not Galilean invariant, which ought to be avoided as the measured subfilter-scale energy transfers otherwise include unphysical fluctuations (Aluie & Eyink 2009a, 2009b; Buzzicotti et al. 2018). However, the energy balance equations can be expressed in an alternative way by including terms that vanish under spatial averaging and ensure Galilean invariance of all terms (Buzzicotti et al. 2018; Offermans et al. 2018). For a statistically stationary evolution, the spatiotemporally averaged energy budget can then be written as\n17\n\n\n\n\n\n\n\n18\n\n\n\n\n\nwhere \n\n\n\n\n\n and \n\n\n\n\n\n are the filtered kinetic and magnetic dissipation rates, respectively, and \n\n\n\n\n\n are terms that convert kinetic to magnetic energy \n\n\n\n\n\n and vice versa \n\n\n\n\n\n, and\n19\n\n\n\n\n\n\n\n20\n\n\n\n\n\n\n\n21\n\n\n\n\n\n\n\n22\n\n\n\n\n\ndenote the four proper energy fluxes, in the sense that they vanish in the limit \u2113 \u2192 0, as can be seen from Equations (13)\u2013(16). If positive, the inertial and Maxwell fluxes, \n\n\n\n\n\n and \n\n\n\n\n\n transfer kinetic energy from scales larger than or equal to \u2113 to scales smaller than \u2113 and vice versa if negative, while the advective and dynamo fluxes, \n\n\n\n\n\n and \n\n\n\n\n\n, do so with magnetic energy. Note that there is no interscale energy conversion as the conversion terms \n\n\n\n\n\n and \n\n\n\n\n\n only involve filtered fields; as such they are known as resolved-scale conversion terms (Aluie 2017). The total energy flux is then given by the sum\n23\n\n\n\n\n\n\n","Citation Text":["Buzzicotti et al. 2018"],"Citation Start End":[[913,935]]}
{"Identifier":"2018MNRAS.473.4566P__Papaderos_et_al._2006_Instance_3","Paragraph":"The young starburst inferred by the detections of high ionization emission line of He\u2009II \u03bb4686 and the blue WR bump in this and previous works (Guseva et al. 2000; Brinchmann, Kunth & Durret 2008) is confirmed by the age estimates made here for the bright and faint regions in Mrk 22 as \u223c4 and \u223c10\u2009Myr, respectively. Unlike previous works, we carried out abundance analysis for both the regions separately. We found an appreciable metallicity difference of \u223c0.5 dex between the bright region [12 + log (O\/H) \u223c 8] and the faint region [12 + log(O\/H) \u223c 7.5]. The separation between two regions is \u223c0.6\u2009kpc. Typical metallicity gradients in normal spiral galaxies have been found between \u22120.009 and \u22120.231 dex kpc\u22121, with an average gradient of \u22120.06 dex kpc\u22121 (Zaritsky, Kennicutt & Huchra 1994). The observed metallicity difference between the two regions in Mrk 22 is too large to be explained as a normal galactic metallicity gradient. The chemical composition as measured from the gas-phase metallicity [12 + log(O\/H)] shows various degree of spatial variations in different types of dwarf galaxies. For instance, shallow gradient in metallicity is seen in SBS 0335\u2212052 (Papaderos et al. 2006) while no significant variations were seen in Mrk 35 (Cair\u00f3s et al. 2007). A study on a large sample indicates that normal BCD galaxies are chemically homogeneous (Kobulnicky & Skillman 1996; Papaderos et al. 2006; Kehrig et al. 2008; Cair\u00f3s et al. 2009; P\u00e9rez-Montero & Contini 2009; P\u00e9rez-Montero et al. 2011; H\u00e4gele et al. 2011; Garc\u00eda-Benito & P\u00e9rez-Montero 2012; Lagos & Papaderos 2013). On the other hand, the metallicity of extremely metal-poor galaxies is usually not homogeneous within the galaxy, with the low metallicity seen in regions of intense star formation (Papaderos et al. 2006; Izotov & Thuan 2009; Levesque et al. 2011; S\u00e1nchez Almeida et al. 2013, 2014, 2015). However, large metallicity gradients are not common in dwarf galaxies. The simplest explanation for large metallicity difference in a single system is a recent merger of two galaxies with different metallicity. In a few cases, significantly large metallicity differences between star-forming regions in dwarf galaxies were seen and understood in terms of recent tidal interactions or mergers (L\u00f3pez-S\u00e1nchez, Esteban & Rodr\u00edguez 2004a,b; L\u00f3pez-S\u00e1nchez, Esteban & Garc\u00eda-Rojas 2006; L\u00f3pez-S\u00e1nchez & Esteban 2009, 2010). The evolution in terms of metallicities in interacting dwarf galaxies is fairly complex as it can depend on various factors such as mixing of metals with the interstellar medium (ISM), possible outflows of metals, and inflow of metal-poor gas in tidally interacting systems.","Citation Text":["Papaderos et al. 2006"],"Citation Start End":[[1770,1791]]}
{"Identifier":"2019AandA...622A.117S__Burg_et_al._2016_Instance_1","Paragraph":"Figure 1a shows for each cluster a visualization of the resulting background-subtracted color-magnitude diagram, produced with an approach similar to that of van der Burg et al. (2016). Specifically, we subtracted the residual background contamination as follows: 1) We start from a candidate cluster member sample obtained as discussed in Sect. 3.1, and a control field sample from the GOODS-S field (see Sect. 2.3) that is selected in the same manner. 2) For each galaxy in the candidate member sample, we calculate a \u201cweight\u201d that corresponds to the statistical excess of the candidate member sample over the control field density at the magnitude and colors of the given galaxy. Weights are calculated as follows (see van der Burg et al. 2016, for a more detailed description): first, all candidate member weights are initially set to 1. Then, for each galaxy in the control field sample we subtract the corresponding \u201cbackground contamination\u201d from the candidate member sample by appropriately reducing the weights of all candidate members that lie within a distance in the color (m814\u2013m140) \u2013 color (m140\u2013[3.6]) \u2013 magnitude (m140) space given by their photometric uncertainties (1\u03c3), with a minimum distance of 0.3 mag, effectively resulting in a smoothing of galaxy densities in the color-color-magnitude space. If no galaxies are found within this distance, we double the search distance, and then if necessary increase the search distance to 1.3 times the distance to the closest galaxy. This criterion allows the full subtraction of the background contamination estimated from all galaxies in the control field sample, while reducing the weights of those candidate members that are more similar to the galaxies in the control field. For each considered control-field galaxy, the weights of all selected candidate members identified with the above criterion are reduced so that the contribution of the considered field galaxy (normalized by the areas of the probed cluster region and of the control field) is removed. At the end of this procedure, the contributions of all galaxies from the field sample have been subtracted from the candidate member sample.","Citation Text":["van der Burg et al. (2016)"],"Citation Start End":[[158,184]]}
{"Identifier":"2022ApJ...938...21M__Deason_et_al._2018_Instance_1","Paragraph":"So far, only one ancient and massive intruder has been identified unambiguously, with its body now nothing but an enormous cloud of tidal wreckage scattered throughout the inner Milky Way. The existence of a vast debris structure left behind by this merger event, known today as the Gaia-Sausage\/Enceladus (GS\/E), was suggested before Gaia (see Evans 2020, for the historical development). In particular, Deason et al. (2013) argued that the rapid transition in the Galactic stellar halo structural properties at break radius of 20\u201330 kpc is likely associated with the apocentric pileup of a relatively early (8\u201310 Gyr ago) and single massive accretion event. It is hypothesized that the progenitor dwarf galaxy was massive enough for its orbit to shrink and radialize quickly as a result of a complex interplay between dynamical deceleration, host recoiling, and self-friction (Vasiliev et al. 2022). Sinking deep in the heart of the Milky Way, the dwarf sprayed the bulk of its stars in a region enclosed by its last apocenter, i.e., some \u223c30 kpc (see Deason et al. 2018). As a result, the region of the Galactic halo within this so-called break radius (see Deason et al. 2011) is inundated with the GS\/E stars, consistently showing up as the most striking substructure even in relatively small volumes around the Sun (see examples of pre-Gaia hints in, e.g., Chiba & Yoshii 1998; Brook et al. 2003; Meza et al. 2005; Nissen & Schuster 2010; Hawkins et al. 2015). Subsequently, the Gaia data made crystal clear the unusually strong radial anisotropy of the relatively metal-rich GS\/E debris and helped to reveal its dominance in the solar neighborhood (see Belokurov et al. 2018; Haywood et al. 2018b; Myeong et al. 2018b). The genesis of the GS\/E debris cloud was also unambiguously confirmed by its unique chemical fingerprint (see, e.g., Helmi et al. 2018; Mackereth et al. 2019) and the large group of globular clusters (GCs) associated with it (see, e.g., Myeong et al. 2018c, 2019; Massari et al. 2019). Close to the Sun, the GS\/E structure has been thoroughly scrutinized both kinematically and chemically (see Evans et al. 2019; Necib et al. 2019; Sahlholdt et al. 2019; Das et al. 2020; Feuillet et al. 2020, 2021; Kordopatis et al. 2020; Molaro et al. 2020; Monty et al. 2020; Aguado et al. 2021a; Buder et al. 2021a; Bonifacio et al. 2021; Carollo & Chiba 2021; Matsuno et al. 2021). Outside of the solar neighborhood, fewer studies exist; nonetheless, the global structure of the GS\/E cloud is starting to come into focus (see, e.g., Iorio & Belokurov 2019; Lancaster et al. 2019; Simion et al. 2019; Naidu et al. 2020; Balbinot & Helmi 2021; Bird et al. 2021; Iorio & Belokurov 2021; Wu et al. 2022).","Citation Text":["Deason et al. 2018"],"Citation Start End":[[1054,1072]]}
{"Identifier":"2022ApJ...928...91C__Stewart_&_Leinhardt_2012_Instance_1","Paragraph":"Our N-body planet-formation simulations utilize the well-established Mercury6 Hybrid integration package (Chambers 1999). While the resolution of our simplified numerical models is admittedly low compared to many contemporary models of planet formation (e.g., Morishima et al. 2010; Carter et al. 2015; Woo et al. 2021), this is acceptable for our purposes as the goal of our study is broadly investigate planet formation around M-dwarfs and follow the dynamical evolution of debris (i.e., collisional fragments and leftover planetesimals: Genda et al. 2012; Stewart & Leinhardt 2012; Quintana et al. 2016; Clement et al. 2019b) generated and stranded during the ultimate epoch of giant impacts in these systems. Thus, we limit the number of simulation particles employed in our initial simulation set and do not include a pebble accretion model to minimize compute time and perform a broad parameter space sweep before switching to higher-resolution models in our investigation of post-formation bombardment (Section 2.2). Moreover, as the initial conditions for planet formation around low-mass stars remain poorly understood and largely unconstrained, we begin our investigation by deviating only slightly from the accepted paradigm of rocky planet accretion in the solar system (e.g., Chambers & Wetherill 1998; Chambers 2001; Hansen 2009). An important limitation of this methodology is the absence of a gas-disk model in our simulations. If short-period planets around M-dwarfs form in situ, their growth is sufficiently rapid that it must have coincided with the gas-disk phase (unless planetesimal formation is delayed for some reason (Ogihara et al. 2015, see Section 4 for a more detailed discussion of these caveats). We plan to expound upon the results of this preliminary study in future work with more detailed models and artificial gas-disk treatments (e.g., Clement et al. 2020) specifically tuned to form specific systems of M-dwarf-hosted planets (e.g., Ormel et al. 2017; Coleman et al. 2019; Schoonenberg et al. 2019).","Citation Text":["Stewart & Leinhardt 2012"],"Citation Start End":[[559,583]]}
{"Identifier":"2022ApJ...937...62L__Xiao_et_al._2022_Instance_2","Paragraph":"However, the intrinsic effects caused by the unknown emission and acceleration mechanisms in the source could mitigate or enhance the LIV-induced time delay, which would impact the accuracy of the resulting constraints on LIV. A key challenge is then to distinguish an intrinsic time lag at the source from a delay induced by LIV. Long GRBs usually have significantly positive or negative intrinsic spectral lags and should not be used for LIV searches until reasonable progress is made on the modeling of the emission and acceleration mechanisms (Chen et al. 2005; Ukwatta et al. 2012; Bernardini et al. 2015), while short GRBs are consistent with null or negligible intrinsic spectral lag and are therefore an ideal tool to measure the LIV effect (Norris & Bonnell 2006; Bernardini et al. 2015, 2017; Xiao et al. 2022). Currently, in addition to short GRBs, active galactic nucleus (AGN) flares and gamma-ray pulsars are two other classes of astrophysical sources that have no significant intrinsic lag in general and are often used for LIV tests (Biller et al. 1999; Kaaret 1999; Aharonian et al. 2008; MAGIC Collaboration et al. 2017). It should be noted, however, that there is also evidence of intrinsic lags in some cases of AGN flares and pulsars. For example, MAGIC Collaboration during an observational campaign regarding Mkn 501 blazar found an indication of about 4 minutes time delay between the peaks at E  0.25 TeV and E > 1.2 TeV (MAGIC Collaboration et al. 2008), which may indicate a progressive acceleration of electrons in the emitting plasma blob. A robust method to study the correlations between arrival times and energy, based on a likelihood function built from the physical picture assumed for the emission, propagation, and detection of the photons was proposed by Mart\u00ednez & Errando (2009). In the case of pulsars, there are some lags if the energy range is extended too much toward low energies (e.g., radio versus TeV). No real progress on the topic of intrinsic effects will be made without accurate models for production and acceleration mechanisms for each type of source. Perennes et al. (2020) first attempted to gain knowledge on source-intrinsic spectral lags of flaring AGNs at high and very high energies and on short timescales relevant for LIV searches, using leptonic AGN flare modeling. Concerning GRBs, some ways have been proposed to reduce the impact of intrinsic effects, e.g., fitting the observed spectral lags of statistical samples of GRBs at a range of different redshifts (Ellis et al. 2006; Bernardini et al. 2017; Xiao et al. 2022), or using only a limited observer-frame energy bands range corresponding to the fixed source-frame energy bands (Wei & Wu 2017). Anyway, there is no reason to think that the low and high-energy photons should be emitted simultaneously at the source, and while detecting distinct signals at different energy channels, we have no idea which one was sent first. Previous studies usually assumed that the intrinsic time delays are either an unknown constant for all GRBs considered or scale with the photon energy E according to some power-law function (Ellis et al. 2006; Biesiada & Pi\u00f3rkowska 2009a; Zhang & Ma 2015; Wei et al. 2017a; Acciari et al. 2020; Pan et al. 2020; Du et al. 2021).","Citation Text":["Xiao et al. 2022"],"Citation Start End":[[2569,2585]]}
{"Identifier":"2022AandA...665A..91A__Kipping_2010_Instance_1","Paragraph":"We modelled the TESS photometry and HARPS RV measurements and accounted for the gravitational interactions between the adopted components of the system using a photodynamical model. The TESS transits are the most numerous of the transit observations and dominate the determination of the planet-to-star radius ratio. For simplicity, we did not model the transit observations with other instruments. The planet positions and velocities in time were obtained through an n-body integration. The sky-projected positions were used to compute the light curve (Mandel & Agol 2002) using a quadratic limb-darkening law (Manduca et al. 1977) that we parametrised following Kipping (2013). The light-time effect (Irwin 1952) was included in the model. To account for the integration time, the model was over-sampled by a factor 3 and was then binned back to match the cadence of the data points (Kipping 2010). The line-of-sight projected velocity of the star issued from the n-body integration was used to model the RV measurements. We used the n-body code REBOUND (Rein & Liu 2012) with the WHFast integrator (Rein & Tamayo 2015) and an integration step of 0.005 days, which resulted in a model error smaller than 1 ppm for the photometric model. The model was parametrised using the stellar mass and radius, planet-to-star mass ratios, planet b-to-star radius ratio, and Jacobi orbital elements (Table 3) at reference time, tref = 2458 370.418461 BJDTDB. Due to the symmetry of the problem, we fixed the longitude of the ascending node of the interior planet \u03a9B at tref to 180\u00b0, and we limited its inclination to ic  90\u00b0. We used a GP regression model with an approximate Matern kernel (celerite, Foreman-Mackey et al. 2017) for the model of the error terms of the transit light curves, and the QPC kernel (Perger et al. 2021; Ambikasaran et al. 2015) for the RVs. For the photometric dataset, we added a transit normalisation factor and a jitter parameter. For the RV, we added a systemic RV and a jitter parameter. In total, the model has 28 free parameters. We used normal priors for the stellar mass and radius from Sect. 3, non-informative sinusoidal priors for the orbital inclinations (uniform in cos i), and non-informative uniform prior distributions for the rest of the parameters. The joint posterior distribution was sampled using the emcee algorithm (Goodman & Weare 2010; Foreman-Mackey et al. 2013) with 200 walkers with starting points based on the results of Sect. 4.2. We ran 1.2\u00d7106 steps of the emcee algorithm and used the last 200 000 steps for the final inference. In Table 3, we list the prior, the median, and the 68% CI of the inferred marginal distributions of the system parameters.","Citation Text":["Kipping 2010"],"Citation Start End":[[886,898]]}
{"Identifier":"2020MNRAS.494.3790K__Wood_&_Hollerbach_2015_Instance_1","Paragraph":"The magnetic field in the crust of a neutron star is simple. In the crust, ions are fixed in a lattice, while only the electrons are mobile. These dynamics may be described based on electron magnetohydrodynamics. The evolution of the magnetic field owing to the Hall drift and Ohmic dissipation has been studied in several researchers. The magnetic fields are numerically simulated on a secular time-scale (refer Pons & Geppert 2007; Kojima & Kisaka 2012; Vigan\u00f2 et al. 2013; Gourgouliatos et al. 2013; Gourgouliatos & Cumming 2014, for axially symmetric cases). Recently, three-dimensional calculations of the evolution have been performed (Wood & Hollerbach 2015; Gourgouliatos, Wood & Hollerbach 2016; Gourgouliatos & Hollerbach 2018). The simulation shows that non-axisymmetric features such as magnetic hotspots produced by an initial perturbation persist over a long time-scale \u223c106\u2009yr. Small-scale features are likely to decay owing to the action of Ohmic dissipation, but they seem to be steadily supported by the Hall effect. This effect becomes more important as the field-strength increases. At the same time, magnetic stress causes deformations in the ion lattice. The crust responds elastically when the deviation is within a critical limit. Beyond this limit, the crust cracks or responds plastically. The effect of elastic deformation on the crustal magnetic-field evolution has been studied (Bransgrove, Levin & Beloborodov 2018). Sudden crust breaking could produce a magnetar outburst and\/or a fast radio burst (Li, Levin & Beloborodov 2016; Baiko & Chugunov 2018; Suvorov & Kokkotas 2019). Plastic flow beyond the critical point is crucial for long-term evolution (Lander 2016), and a coupled system between the flow and magnetic field can be solved numerically in a two-dimensional square box of a crust-depth size (Lander & Gourgouliatos 2019). Their simulation shows significant motion near the surface at a rate of 10\u2013100 cm yr\u22121 for a certain viscosity range. At 103\u2009yr, the path-length reaches \u223c1\u2009km, which is the size of the simulation in plane-parallel geometry. The magnetic-field structure comprises a poloidal loop field confined in a two-dimensional box. It is important to examine the effect of the plastic flow on the global features. In particular, changes in the external dipole field in 103\u2013105\u2009yr are interesting.","Citation Text":["Wood & Hollerbach 2015"],"Citation Start End":[[642,664]]}
{"Identifier":"2018ApJ...862..171W__Li_et_al._2016a_Instance_1","Paragraph":"Theoretical features of CB-SMBHs have been predicted for observations; however, observational data can be alternatively explained by complicated broad-line region (BLR) models around a single black hole. Thus, they are still elusive, though there is growing evidence for the appearance of CB-SMBHs. As one of the several signatures of CB-SMBHs, radial velocity curves of double-peaked H\u03b2 emission line have been suggested by Popovi\u0107 et al. (2000), Shen & Loeb (2010), Tsalmantza et al. (2011) and Popovi\u0107 (2012),  but the asymmetry is a better tracer than the double-peakedness in the reverberations (Shen & Loeb 2010). Candidates with this feature are NGC 4151 (Bon et al. 2012) and NGC 5548 (Li et al. 2016a, 2016b), which show opposite motions of red and blue peaks. A large sample of active galactic nuclei (AGNs) with double-peaked H\u03b2 profiles built up by Eracleous et al. (2012) has been monitored for systematic shifts of the two peaks for a couple of years (Runnoe et al. 2017). Periodical variations could be regarded as a signature of binary black holes, such as OJ 287 (Valtonen et al. 2008), PG 1320 (Graham et al. 2015), Akn 120 (Li et al. 2017), and others (Liu et al. 2015; Zheng et al. 2016; Dorn-Wallenstein et al. 2017). It usually takes at least three times the period to determine the periodicity of varying AGNs, such as, more than 10 yr to justify their periodicity if the orbital periods are 3 yr or so (Li & Wang 2018). An X-shaped radio jet has been suggested as a signature of binary black holes (Merritt & Ekers 2002; Cheung 2007), but only one could be plausible (Kharb et al. 2017). A deficit of UV emissions of spectral energy distributions is formed by the interaction between the secondary and circumbinary disks (Hayasaki et al. 2008; Schnittman 2011a, 2011b; Sesana et al. 2012; Roedig et al. 2014). However, there are alternative explanations for these phenomena, such as a processing jet as an alternative explanation of AGN periodical variabilities (since most of them are radio-loud AGNs), dust extinction for the UV deficit in Mrk 231 (Ye et al. 2016; but see Leighly et al. 2016), elliptical disks (Eracleous et al. 1995), and hot spots (e.g., Jovanovi\u0107 et al. 2010) or spiral arms (e.g., Storchi-Bergmann et al. 2017) for the double-peaked profiles of the broad H\u03b2 line. Three candidates with double-peaked profiles, Arp 102B, 3C 390.3, and 3C 332, have been excluded by long-term monitoring campaigns (Eracleous et al. 1997). No CB-SMBHs have been unambiguously identified so far.","Citation Text":["Li et al. 2016a"],"Citation Start End":[[694,709]]}
{"Identifier":"2021ApJ...907...62R__Greve_et_al._2005_Instance_1","Paragraph":"We have used the Code Investigating GALaxy Emission (CIGALE; Burgarella et al. 2005; Noll et al. 2009; Boquien et al. 2019) SED fitting package to model the full SED of ADFS-27. We have run the code on all spatially integrated photometry from Table 2 for the entire galaxy and also on the two source components ADFS-27N and S individually, using only the photometry in those bands where the emission is resolved into both sources. We ran two series of fits, with star formation histories (SFHs) limited either to approximately the age of the universe at z = 5.655 (i.e., \u22641.0 Gyr) or to an age of \u22640.2 Gyr, as is characteristic for young dusty starbursts (e.g., Greve et al. 2005; Bergvall et al. 2016). We sampled standard ranges for all main parameters within CIGALE and used a Bruzual & Charlot (2003) single stellar population and a Chabrier (2003) stellar initial mass function with a \u201cdelayed\u201d SFH (i.e., where SFR(t) \u221d t\/\u03c42 \u00d7 exp(\u2212t\/\u03c4), and the fit parameter \u03c4 is the time at which the SFR peaks; Boquien et al. 2019; see also Burgarella et al. 2020) for all fits. We fixed the power-law slope of the dust attenuation law to \u22120.7 for both the interstellar medium and birth clouds and assumed a polycyclic aromatic hydrocarbon (PAH) mass fraction (which is not directly constrained by the data) of qPAH = 3.9% (i.e., similar to dust found in the Milky Way and nearby galaxies with near-solar metallicity; e.g., Draine & Li 2007, and references therein).13\n\n13\nWe allowed for a range in metallicity in our fits, but the metallicity remains difficult to constrain directly without rest-frame optical spectroscopy.\n This is consistent with the range of qPAH = 0.47%\u20133.9% found for a sample of infrared-luminous galaxies at z = 0.5\u20134 (Magdis et al. 2012).14\n\n14\nThe uncertainty due to this assumption is fully captured by the error bars for the most relevant parameters, like Mdust.\n For ADFS-27 as a whole and for ADFS-27S (which is detected in the near-infrared bands), the best-fit parameters for both series agree within the uncertainties. We thus adopt those from the series with less constraints on the SFH in the following, and we only use those from the other series in the evaluation of the true uncertainties. For ADFS-27N (which is not detected in the near-infrared bands), the series with more stringent constraints on the SFH appears to provide more reasonable results and thus are adopted in the following.","Citation Text":["Greve et al. 2005"],"Citation Start End":[[662,679]]}
{"Identifier":"2017ApJ...850...77K__Krumholz_et_al._2012_Instance_1","Paragraph":"The Galactic Center (GC) and the Central Molecular Zone (CMZ) in particular represent an environment with conditions that are not to be found anywhere else on large scales in the Milky Way. The CMZ received its name from the presence of a large reservoir of dense (\n\n\n\n\n\n cm\u22122) and molecular gas of a few times 107 M\u2299 (Oka et al. 1998; Morris & Serabyn 1996; Ferri\u00e8re et al. 2007) and covers the central \u223c500 pc of the GC region. The large amount of molecular gas is found to be accompanied by a relatively high star formation rate (SFR) of \u223c0.1 M\u2299 yr\u22121 (Longmore et al. 2013a; Barnes et al. 2017). Star formation (SF) laws that build on the assumption of a constant gas-depletion time of 1 Gyr can fit the observed SFR (Bigiel et al. 2010; Leroy et al. 2008, 2015). However, these relations are derived from gas at much lower densities than observed in the CMZ. Density-dependent SF laws, on the other hand, strongly overpredict the SFR (Longmore et al. 2013a) at \u223c0.4 M\u2299 yr\u22121 (Kennicutt 1998; Krumholz & McKee 2005; Krumholz et al. 2012) to \u223c0.8 M\u2299 yr\u22121 Lada et al. (2010, 2012). Thus, the star formation efficiency (SFE) in the CMZ is significantly lower than expected for the observed gas densities. The complex interplay of energetic processes in the GC allows for different potential answers to this problem of low SFE, but none of these processes alone can explain the discrepancy (Kruijssen et al. 2014). A general picture of episodic starbursts in gas rings in galactic centers introduced by Krumholz & Kruijssen (2015) and Krumholz et al. (2017) might solve the SFE problem in the particular case of the CMZ and set the framework for another intriguing feature of the GC: a ring-like structure of dust and molecular gas. This structure (see Figure 1 for an overview) is projected onto an infinity (\n\n\n\n\n\n) shape that follows several arcs, the most prominent being the so-called \u201cdust ridge\u201d stretching from the massive molecular cloud G0.253+0.016 (\u201cthe Brick\u201d) to the star-forming region Sgr B2 (Lis et al. 1999). It might be continued via an \n\n\n\n\n\n arc and another loop at negative longitude ([\n\n\n\n\n\n], \n\n\n\n\n\n]) passing through the star-forming region Sgr C. The syntax \n\n\n\n\n\n denotes the four quadrants in Galactic coordinates at positive\/negative Galactic longitude (l) and latitude (b). Clouds at \n\n\n\n\n\n and \n\n\n\n\n\n are thought to be located in front of Sgr A*, i.e., the near side of the GC that is visible in silhouette against the background, whereas most of the \n\n\n\n\n\n and \n\n\n\n\n\n gas is more likely to be on the far side, i.e., behind Sgr A* (Bally et al. 2010; Molinari et al. 2011).","Citation Text":["Krumholz et al. 2012"],"Citation Start End":[[1018,1038]]}
{"Identifier":"2015ApJ...805...72M__Menten_et_al._1986_Instance_1","Paragraph":"In the interstellar medium, CH3OH  and other \u201csaturated\u201d (hydrogen-rich) molecules are primarily believed to be formed on the surface of dust grains (Tielens & Hagen 1982; Charnley et al. 1992; Watanabe & Kouchi 2002), where CH3OH and NH3  are among the most abundant mantle species present relative to H2O, as measured in both low and high-mass young stellar objects (YSOs) and cold cloud cores (Tielens & Allamandola 1987; Dartois et al. 1999; Pontoppidan et al. 2004; Gibb et al. 2004; Boogert et al. 2008; \u00d6berg et al. 2011). Notably, for CH3OH, the high abundances measured for maser sources are inconsistent with those predicted by gas-phase formation models (Menten et al. 1986; Hartquist et al. 1995). To get the CH3OH  off of the dust grains and into the gas phase in the observed large quantities then requires a mechanism to liberate the CH3OH  from the grain mantles. Proposed mechanisms include thermal desorption via heating from an embedded protostar, shocks, or cosmic rays (requiring grain temperatures \u227390 K Tielens 1995; Brown & Bolina 2007), or nonthermal desorption processes including photodesorption via far-UV photons from cosmic ray interactions (Prasad & Tarafdar 1983; D\u2019Hendecourt et al. 1985; \u00d6berg et al. 2009a), grain sputtering, wherein the ice mantles of grains are dislodged via collisions (often in shocks) with other grains, neutrals, ions, or cosmic rays (Johnson et al. 1991; Caselli et al. 1997), and finally exothermic chemical reactions on the grain surfaces Duley & Williams (1993), Roberts et al. (2007). In particular, the presence of molecules formed on grains in relatively cool and dense environments requires an efficient nonthermal desorption process (Willacy & Williams 1993; Roberts et al. 2007; \u00d6berg et al. 2009b; Caselli et al. 2012). In the CMZ, the abundance of 36 GHz CH3OH  sources has been suggested to be due to photodesorption from cosmic rays in this region Yusef-Zadeh et al. (2013a). We reconsider this in the light of our new, high-resolution observations.","Citation Text":["Menten et al. 1986"],"Citation Start End":[[666,684]]}
{"Identifier":"2016MNRAS.463.3637R__Lelli,_McGaugh_&_Schombert_2016a_Instance_1","Paragraph":"The classical framework in which theoretical galactic studies are performed these days is the \u039b cold dark matter (\u039bCDM) paradigm. However, both the cosmological constant \u039b and the CDM part of the model could also be related to a modification of gravity. On galaxy scales, the model is indeed plagued by severe problems, the most famous ones being the cusp\u2013core problem (de Blok 2010; Oman et al. 2015, but see also Read et al. 2016), the too-big-to fail problem (Boylan-Kolchin, Bullock & Kaplinghat 2011; Papastergis et al. 2015; Pawlowski et al. 2015) or the satellite planes problem (Kroupa, Theis & Boily 2005; Metz, Kroupa & Jerjen 2007; Metz, Kroupa & Libeskind 2008; Pawlowski, Pflamm-Altenburg & Kroupa 2012; Ibata et al. 2013, 2014; Pawlowski et al. 2015). There is also a more general problem linked to the finely tuned relation between the distribution of baryons and the gravitational field in galaxies, as encapsulated in various scaling relations involving a universal acceleration constant a0 \u2248 10\u221210\u2009m\u2009s\u22122, including the tight baryonic Tully\u2013Fisher relation (McGaugh et al. 2000; Lelli, McGaugh & Schombert 2016a; Papastergis, Adams & van der Hulst 2016), the diversity of shapes of rotation curves at a given maximum velocity scale (Oman et al. 2015) or the relation between the stellar and dynamical surface densities in the central regions of galaxies (Lelli et al. 2016b; Milgrom 2016), and many others (Famaey & McGaugh 2012). All this points to things happening as if the effects usually attributed to CDM on galaxy scales were actually due to a modified force law. The a priori simplest explanation for this would be that gravity is indeed effectively different in the weak field regime and accounts for the effects usually attributed to CDM. This paradigm is known as modified Newtonian dynamics (MOND), or Milgromian dynamics, suggested more than 30 years ago by Milgrom (1983). It predicted all the observed galaxy scaling relations well before they were precisely assessed by observations (Famaey & McGaugh 2012). Nevertheless, this paradigm cannot be complete, as a full theory of gravitation, also valid on cosmological scales, has not yet been found. But while successful on galaxy scales, the MOND paradigm has still been far from being fully explored even on these scales where it is currently successful. Hence, there is still potential for falsification of this paradigm in its a priori domain of validity. The main reason for this lack of exploration of all predictions of MOND on galaxy scales is its non-linear nature, and the previous lack of numerical codes devised to model galaxies in this framework.","Citation Text":["Lelli, McGaugh & Schombert 2016a"],"Citation Start End":[[1096,1128]]}
{"Identifier":"2020ApJ...892..117W__Faucher-Giguere_&_Kaspi_2006_Instance_1","Paragraph":"Over 2500 known pulsars have been included in the database of the Australia Telescope National Facility (ATNF1\n\n1\n\nhttp:\/\/www.atnf.csiro.au\/research\/pulsar\/psrcat\n\n), whose parameters can be queried and visualized, and more fascinating pulsars are waiting to be discovered (Manchester et al. 2005; Bates et al. 2012; Ng et al. 2015). Many of these stars have very stable rotation periods. Among them, the rotation periods of \u201cnormal\u201d pulsars vary from a few tens of milliseconds to a few seconds, and the derivatives of their periods are between 10\u221217 and 10\u221213 (Manchester et al. 2001; Lorimer et al. 2006); they are called cosmic \u201clighthouses\u201d because their broadband emission and steady rotation sweep out a beam of radiation across the sky. These lighthouse-like pulsars can radiate photon signals with stable periods and profiles, thus playing an important role in studying the evolution law of stars (Faucher-Giguere & Kaspi 2006), exploring the interstellar medium (ISM; Daniel et al. 2000; Bogovalov et al. 2005), detecting gravitational waves (Hobbs et al. 2009), establishing timing references (Hobbs et al. 2012) and realizing deep-space navigation (Sheikh et al. 2006; Emadzadeh & Speyer 2011). On-orbit application research on pulsars is being performed worldwide (Vikhlinin et al. 2009; Ge et al. 2012; Kaya et al. 2014). However, pulsars are far from the Earth, and compared with the receiving areas of gigantic ground-based radio observatories, that of an X-ray detector is more limited; therefore, the received photon signals are always submerged in observation noise. Meanwhile, compared with radio signals, which are more susceptible to interference from artificial signals, other forms of ground signals and photon signals from other stars, X-ray signals are almost immune to interference because of their small field of view and high energy. However, we cannot completely rule out the possibility of interference from instruments and other stars. The above factors hamper research on pulsar searches and timing observations. Therefore, it is quite important to study noise and interference mitigation methods for pulsar signals.","Citation Text":["Faucher-Giguere & Kaspi 2006"],"Citation Start End":[[907,935]]}
{"Identifier":"2017AandA...603A..24M__Silaj_et_al._2010_Instance_1","Paragraph":"Be stars are believed to rotate at ~60\u221280% of their critical velocity (Porter 1996; Rivinius et al. 2006). Earlier observations of the V850 Cen H\u03b1 line (Parkes et al. 1980; Corbet et al. 1986) detected a shell-type profile of the emission line, that is a double-peaked profile in which the central absorption extends below the stellar continuum flux. According to the seminal model by Struve (1931, yet largely accepted; see Rivinius et al. 2013, for a more recent review), shell-type profiles imply an inclination angle of i ~ 90\u00b0 (i.e. an edge-on system). During our campaign we did not observe shell-type emission lines (see Fig. 2). Therefore, based on a qualitative comparison with Fig. 1 of Rivinius et al. (2013), the profiles detected in this work support a large (60\u00b0 \u2272  i  \u2272 80\u00b0) inclination angle, yet not an edge-on view. Later, more physical models (Hummel 1994; Hanuschik 1996; Hummel & Hanuschik 1997) also support a large inclination angle. However, we should note here that line profile shapes do not always give unambiguous information about the disk inclination angle (Silaj et al. 2010; Quirrenbach et al. 1997). In any case, transient shell lines do not demand a tilt of the disk plane. For example, shell profiles need a geometrically thick medium to form. Therefore, the H\u03b1 line profiles observed in the present work may simply be due to a decrease of the disk density, with respect to previous conditions. Assuming typical values of the mass and radius of a B2 type star as M\u22c6 = 9.9 M\u2299, R\u22c6 = 5.4 R\u2299 (Townsend et al. 2004; Pasinetti-Fracassini et al. 2000; Sugizaki et al. 2015), the critical velocity of rotation is (4)\\begin{equation} V_{\\rm break}=\\sqrt{\\frac{GM_\\star}{R_\\star}}\\approx475\\,\\rm km\\,s^{-1}. \\end{equation}Vbreak=GM\u22c6R\u22c6\u2248475\u2009km\u2009s-1.With an inclination angle i = 90\u00b0 and Vsini = 330 km\u2009s-1, the resulting critical fraction is w = V\/Vbreak \u2248 0.69. On the other hand, assuming w = 0.8 and Vsini = 330 km\u2009s-1, the resulting inclination angle is i ~ 60\u00b0, which is still consistent with the inclination inferred by the H\u03b1 line profile. We therefore conclude that, at least during our observations, a more conservative range for the inclination angle of V850 Cen is 60\u00b0 \u2272  i  \u2272 90\u00b0 and that V850 Cen is likely rotating at 0.69 \u2272  w  \u2272 0.80. ","Citation Text":["Silaj et al. 2010"],"Citation Start End":[[1088,1105]]}
{"Identifier":"2022MNRAS.509.3427D__Connaughton_et_al._2017_Instance_1","Paragraph":"Unlike BH\u2013BH mergers, binary neutron star (BNS) merger mechanisms are expected to yield an optical counterparts powered by the radioactive decay of rapid neutron capture process (r-process) elements synthesized in the merger ejecta (Li & Paczy\u0144ski 1998; Metzger et al. 2010) or by the cooling of shock-heated material around the neutron star (NS; Piro & Kollmeier 2018). On 2017 August 17, the LVC detected the first and best example of a BNS merger: GW170817 (Abbott et al. 2017a). Only two seconds after its detection, a short gamma-ray burst (GRB 170817A) was detected (Connaughton et al. 2017; Goldstein et al. 2017). GRB 170817A was followed by the discovery of the optical counterpart Swope Supernova Survey 17a [SSS17a (AT 2017gfo)] \u223c11 h later in the galaxy NGC 4993 (Coulter et al. 2017) and later confirmed by others teams (Abbott et al. 2017c; Arcavi et al. 2017; Covino et al. 2017; D\u00edaz et al. 2017; Drout et al. 2017; Kilpatrick et al. 2017; Lipunov et al. 2017; McCully et al. 2017; Pian et al. 2017; Shappee et al. 2017; Smartt et al. 2017; Soares-Santos et al. 2017; Troja et al. 2017; Utsumi et al. 2017; Valenti et al. 2017). The combination of EM and GW information on the BNS can be used to constrain the mass and the radii of NS (Margalit & Metzger 2017) or their equation of state (Bauswein & Janka 2012; Annala et al. 2018). Also, observations of the optical\u2013infrared (IR) \u2018kilonova\u2019 counterpart (Chornock et al. 2017; Cowperthwaite et al. 2017; Drout et al. 2017; Kilpatrick et al. 2017; Nicholl et al. 2017; Shappee et al. 2017; Valenti et al. 2017; Villar et al. 2017) provided the first observational confirmation that NS mergers produce the majority of the r-process elements heavier than iron (Burbidge et al. 1957; Cameron 1957; Kasen et al. 2017; Pian et al. 2017; Metzger 2019). They also permit tests of theoretical kilonova models. For example, Drout et al. (2017) showed that the temperatures cooled from 10\u2009000 to 5100 K in between 12 and 36 h after the event, confirming model predictions. Also, Shappee et al. (2017) used spectra taken 11.76 and 12.72 h after the merger to show that the photosphere was expanding at \u223c0.3c.","Citation Text":["Connaughton et al. 2017"],"Citation Start End":[[573,596]]}
{"Identifier":"2016ApJ...831...63X__Neistein_et_al._1999_Instance_1","Paragraph":"The formation and evolution of S0 galaxies are very important for understanding the formation and evolution of galaxies, but they are still an open question (e.g., the recent review by D\u2019Onofrio et al. 2015). Currently, there are two possible scenarios on the origin of S0 galaxies. One is that S0 galaxies are transformed from spiral galaxies, where spirals lose their gas and star formation is rapidly quenched. The other is that S0 galaxies are intrinsically different from spiral galaxies since their formation (Kormendy & Kennicutt 2004; Barway et al. 2009; van den Bergh 2009). The transformation origin may be associated with intra-cluster medium and neighboring galaxies, via minor mergers, slow encounters, galaxy harassments (Moore et al. 1996), or tidal effects in the dense environment (Gunn & Gott 1972; Larson et al. 1980; Dressler & Sandage 1983; Mihos & Hernquist 1994; Moore et al. 1998, 1999; Neistein et al. 1999; Shioya et al. 2002). A lot of studies have discussed the environmental dependence of galaxy evolution. Generally speaking, early-type galaxies tend to be in dense environments and have low star-formation rates (SFRs; Dressler 1980; Balogh et al. 1997, 1998, 2000; Poggianti et al. 1999; Treu et al. 2003). The fraction of S0 galaxies in the field is only about 15%, while spirals are the majority (Naim et al. 1995). Within the group environment, spirals and S0s are both about 40%\u201345% (Postman & Geller 1984). Furthermore, S0s become dominant in dense environments, the fraction grows up to 60% in clusters (Dressler 1980; Postman & Geller 1984). Dressler et al. (1997), Fasano et al. (2000), and Desai et al. (2007) also found that the galaxy morphological distributions change abruptly in clusters at \n\n\n\n\n\n, about 50 \u223c 70% of spirals at high redshift (z > 0.4) are transformed into S0s, while a fraction of ellipticals, about 25%, remains nearly constant between z = 0.8 and z = 0.0. However, Wilman et al. (2009) showed that the fraction of S0s in groups is the same as in clusters but it is much higher than in the field at z = 0.4, which might suggest that S0s are formed in groups or subgroups.","Citation Text":["Neistein et al. 1999"],"Citation Start End":[[911,931]]}
{"Identifier":"2019MNRAS.487.2639R__Maslov_et_al._2015_Instance_1","Paragraph":"The early studies of the cooling of hypernuclear stars did not include the high-mass astrophysical constraint on the EoS of hypernuclear matter (Haensel & Gnedin 1994; Schaab, Balberg & Schaffner-Bielich 1998; Tsuruta et al. 2009). In Paper I we started a systematic study of the cooling of hypernuclear compact stars on the basis of modern density functionals mentioned above. Several works have appeared since then in this context: Grigorian, Voskresensky & Maslov (2018) used the \u2018nuclear medium cooling\u2019 scenario (Schaab et al. 1997; Blaschke, Grigorian & Voskresensky 2004) to account for hyperons using an EoS based on the DFT (Maslov et al. 2015), pion-mediated enhanced neutrino emissivities for the modified Urca and bremsstrahlung processes, vanishing P-wave neutron pairing, and phenomenological crust\u2013core temperature relation. The pion-mediated enhanced neutrino emissivities and vanishing P-wave neutron pairing are the key factors that distinguish their models from those in Paper I. Nevertheless, the remainder of the physical input, for example, the crust\u2013core temperature relation, is not identical to ours. Note that Paper I does not implement the minimal cooling scenario (Page et al. 2004) as it allows for the dUrca process among nucleons and hyperons, uses pair-breaking processes with dressed (in-medium) vertices, and covers a wide range of masses up to the maximum mass of Mmax\/M\u2299 \u223c 2, which implies a varying cooling behaviour for the mass hierarchy (see, in particular, the conclusion Section of Paper I). The same problem was also addressed by Negreiros et al. (2018) using hyperonic EoS based on the FSU class of relativistic density functional models which feature massive objects and are tuned to available data on hypernuclei (Tolos, Centelles & Ramos2017a,b). In these models, the main cooling agents are the (\u039b, p) and (n, p) dUrca processes which were regulated by variations of proton and neutron superfluid gaps. In particular the density range over which the proton S-wave gap is non-zero has been explored with the aim to obtain a satisfactory agreement with the data. The main difference with respect to Paper I is the absence of pairing in the hyperonic sector. Combined, the emerging new generation of cooling models of hypernuclear stars have the potential to further constraining the physics of dense nuclear matter, especially its composition.","Citation Text":["Maslov et al. 2015"],"Citation Start End":[[634,652]]}
{"Identifier":"2021ApJ...915L...8D__Bowen_et_al._2020a_Instance_1","Paragraph":"Magnetic field fluctuations in the solar wind are highly turbulent. The measured power spectral density (PSD) of the fluctuating magnetic field always exhibits power laws k\u2212\u03b1, where k is the wavenumber, and \u03b1 is the spectral index. A single spacecraft measures the PSD as a function of f\u2212\u03b1 in the frequency domain, which can be converted to the spatial domain under the Taylor hypothesis. According to the physical processes at different scales, the PSD in the solar wind can be divided into several segments, which can be fitted with different \u03b1. The inertial range, which is dominated by magnetohydrodynamic (MHD) turbulence, follows the cascade models with spectral indices \u03b1i from around 3\/2 to 5\/3 (Bruno & Carbone 2013; Chen et al. 2020). The PSDs become steepened below the ion scales (ion thermal gyroradius \u03c1i or ion inertial length di), where kinetic mechanisms should be taken into account. Sometimes a sharp transition range is observed with \u03b1t \u223c 4 (Sahraoui et al. 2010; Bowen et al. 2020a). This transition range may be caused by imbalanced turbulence (Voitenko & Keyser 2016; Meyrand et al. 2021), energy dissipation of kinetic waves (Howes et al. 2008), ion-scale coherent structures (Lion et al. 2016), or a reconnection dominated range (Mallet et al. 2017). At smaller scales, a flatter sub-ion kinetic range forms with the spectral index \u03b1k \u223c 7\/3, which can be explained as the MHD Alfv\u00e9nic turbulence developing into a type of kinetic wave turbulence, e.g., kinetic Alfv\u00e9n waves (KAWs; Schekochihin et al. 2009) or whistler waves (Cho & Lazarian 2004). Intermittency in the kinetic range could lead to an \u22128\/3 spectrum (Boldyrev & Perez 2012; Zhao et al. 2016). Ion-cyclotron-wave (ICW) turbulence could lead to a steeper \u221211\/3 spectrum (Krishan & Mahajan 2004; Galtier & Buchlin 2007; Meyrand & Galtier 2012; Schekochihin et al. 2019).The kinetic range always behaves as the KAW turbulence with the slope of \u22122.8 in the near-Earth space (Bale et al. 2005; Chen et al. 2013; Chen 2016). The spectral indices increase again beyond the electron kinetic scales in observations, indicating the conversion of turbulence energy to electrons (Sahraoui et al. 2009; Alexandrova et al. 2012; Chen et al. 2019) or transitions to a further cascade (Schekochihin et al. 2009; Chen & Boldyrev 2017). In simulations, Meyrand & Galtier (2013) obtained a \u22128\/3 spectrum at electron scales under the 3D electron\u2013MHD model.","Citation Text":["Bowen et al. 2020a"],"Citation Start End":[[984,1002]]}
{"Identifier":"2020MNRAS.495.2620R__Skrutskie_et_al._2006_Instance_1","Paragraph":"The main goal of this paper is to use the JAM method to measure the parameters MBH, (M\/L)TOT, and \u03b2z of the central region of the galaxy NGC 4546. To do this, we obtained AO assisted IFU observations from the Near-Infrared Integral Field Spectrometer (NIFS), installed on the Gemini North Telescope, and also archive photometric observations performed with the Wide Field and Planetary Camera 2 (WFPC2), installed on the Hubble Space Telescope (HST). NGC 4546 (see Fig. 1) is an SB0 galaxy (de Vaucouleurs et al. 1991) with magnitude Mk = \u221223.30 (Skrutskie et al. 2006, K band) and is located at a distance of 14 Mpc (Tully et al. 2013; 1 arcsec = 68 pc). It hosts a type 1 LINER AGN at the centre (Ricci, Steiner & Menezes 2014b). This object was also classified as a fast rotator by Emsellem et al. (2011). Gas discs in molecular, neutral, and ionized forms are all counter rotating with respect to the stellar disc (Galletta 1987; Bettoni, Galletta & Oosterloo 1991; Sage & Galletta 1994; Sarzi et al. 2006; Ricci, Steiner & Menezes 2014a). Cappellari et al. (2013a) used JAM to model the stellar kinematics of NGC 4546 using a SAURON data cube of this galaxy from the ATLAS3D project (Cappellari et al. 2011). Although their procedure was very important to determine global parameters for the stellar dynamics of NGC 4546, such as (M\/L)TOT and \u03b2z within 1 effective radius (1 Reff = 22 arcsec, Cappellari et al. 2013a), the spatial resolution of the data cubes from the ATLAS3D project is not high enough to resolve the sphere-of-influence radius of this object. Assuming $R_{\\mathrm{ soi}} \\, = \\, GM_{\\mathrm{ BH}} \/ \\sigma _e^2$, where \u03c3e is the velocity dispersion within 1 effective radius and using the MBH \u00d7 \u03c3 relation of Saglia et al. (2016), we calculate Rsoi = 18 pc (0.26 arcsec) for \u03c3e = 188 km s\u22121 (Cappellari et al. 2013a). We will show that Rsoi is resolved by our NIFS observations. Thus, this paper will complement the JAM results of Cappellari et al. (2013a) for NGC 4546, but for the vicinity of its SMBH, i.e. using AO assisted IFU data within a field of view (FOV) of 200 \u00d7 200 pc2.","Citation Text":["Skrutskie et al. 2006"],"Citation Start End":[[547,568]]}
{"Identifier":"2021MNRAS.505.5117Z__Cooray_&_Sheth_2002_Instance_1","Paragraph":"Finally, to calculate \u03a3(rp), we need to build an accurate model for \u03behm on radial scales between 0.1$\\, h^{-1}\\, {\\mathrm{Mpc}}$ and ${\\sim}100\\, h^{-1}\\, {\\mathrm{Mpc}}$. We adopt the two-component model of \u03behm developed by Zu et al. (2014; a modified version proposed by Hayashi & White 2008):\n(6)$$\\begin{eqnarray*}\r\n\\xi _{\\mathrm{hm}}(r) &=& \\left\\lbrace \\begin{array}{ll}\\xi _\\mathrm{1h} & \\quad \\mbox{if}\\ \\xi _\\mathrm{1h} \\geqslant \\xi _\\mathrm{2h},\\\\\r\n\\xi _\\mathrm{2h} & \\quad \\mbox{if}\\ \\xi _\\mathrm{1h} \\lt \\xi _\\mathrm{2h},\\\\\r\n\\end{array} \\right.\\nonumber \\\\\r\n\\xi _\\mathrm{1h} &=& \\frac{\\rho _\\mathrm{NFW}(r|M_h)}{\\rho _\\mathrm{m}} - 1, \\nonumber \\\\\r\n\\xi _\\mathrm{2h} &=& b \\,\\, \\xi _{\\mathrm{mm}}.\r\n\\end{eqnarray*}$$Here, \u03be1h and \u03be2h are the so-called \u20181-halo\u2019 and \u20182-halo\u2019 terms in the halo model (Cooray & Sheth 2002), \u03c1NFW(r|Mh, c) is the NFW density profile of halo mass Mh and concentration c, b is the average halo bias, and \u03bemm is the non-linear matter\u2013matter autocorrelation function predicted at Planck cosmology (Takahashi et al. 2012). Zu et al. (2014) found that this simple model provides an adequate description of the halo\u2013matter cross-correlation measured from simulations at the level of a few per\u2009cent on scales of our concern (i.e. below the halo radius). For massive haloes, several studies found that the average density profile in the inner region of haloes deviates from the NFW shape and the Einasto profile is more accurate (Dutton & Macci\u00f2 2014; Klypin et al. 2016). However, since we are only fitting to scales above $0.1\\, h^{-1}\\, {\\mathrm{Mpc}}$, the difference between NFW and Einasto should be negligible. We have also ignored the extra lensing effect caused by the stellar mass of the BCGs, which has negligible contribution on scales above $0.1\\, h^{-1}\\, {\\mathrm{Mpc}}$. The uncertainties of the predicted \u0394\u03a3 can be as large as $10{{\\ \\rm per\\ cent}}$ around the transition between the 1-halo and 2-halo scales ($2{-}4\\, h^{-1}\\, {\\mathrm{Mpc}}$; see the fig. 5 of Zu et al. 2014), but we are only concerned with the measurement of halo mass and concentration, which are inferred primarily from scales below the transition scale, and the large-scale bias, which is estimated from scales above $10\\, h^{-1}\\, {\\mathrm{Mpc}}$. Additional, our statistical uncertainties of weak lensing at those scales is around ${~}10{{\\ \\rm per\\ cent}}$. Therefore, the impact of relatively large uncertainties in the transitional regime on our conclusions should be negligible. For future modelling of cluster weak lensing signals of\u2009per\u2009cent level uncertainties, the transitional behaviour can be potentially improved by using the scheme proposed by Garcia et al. (2020) and emulators developed by Salcedo et al. (2020).","Citation Text":["Cooray & Sheth 2002"],"Citation Start End":[[843,862]]}
{"Identifier":"2015ApJ...808...56M__Deming_et_al._2015_Instance_2","Paragraph":"We have tested the pixel-ICA algorithm, i.e., a non-parametric method proposed by Morello et al. (2014, 2015) to detrend Spitzer\/IRAC primary transit observations, on simulated data sets. Systematics similar to the ones present in Spitzer\/IRAC data sets are obtained by combining instrumental jitter with inter- or intra-pixel sensitivity variations. A variety of jitter timeseries is used to test the pixel-ICA method with:\n\n1.\nperiodic signals with different frequencies, phasing, and shape;\n\n\n2.\nnon-stationary signals with varying amplitudes or frequencies;\n\n\n3.\nsudden change.\n\nThe detrending performances of pixel-ICA method have been compared with division by a polynomial function of the centroid, in this paper PCD method, and PLD method (Deming et al. 2015). Here we summarize the main results found:\n\n1.\nPixel-ICA algorithm can detrend non-stationary signals and sudden changes, as well as periodic signals with different frequencies and phasing, relative to the transit.\n\n\n2.\nInter-pixel effects are well-detrended with pixel-ICA method.\n\n\n3.\nEven if the instrument PSF is not entirely within the array of pixels, pixel-ICA leads to results which are consistent at \u223c1\u03c3 with the input parameters.\n\n\n4.\nIn most cases, pixel-ICA outperforms PCD method, especially if the instrument PSF is narrow, or it is not entirely within the photometric aperture.\n\n\n5.\nIntra-pixel effects are only detectable if the PSF is relatively small.\n\n\n6.\nIntra-pixel effects cannot be totally detrended by any of the three methods, but pixel-ICA, in most cases, outperforms PCD method, which is largely case-dependent. Also, pixel-ICA method provides consistent results within the error bars.\n\n\n7.\nIt is possible to fit the astrophysical signal after detrending or together with the other components. The only differences are registered if at least one of the non-transit components has a similar shape at the time of transit, in which case the first approach is preferable, but the two results were consistent within 1\u03c3.\n\n\n8.\nThe PLD method, updated to include cross-term between pixel fluctuations and the astrophysical signals, lead to very similar results than pixel-ICA, particularly if the astrophysical signal is fitted together with the other components.\n\nIn conclusion, we have found in a variety of simulated cases that pixel-ICA performs as well or better than other methods used in the literature, in particular polynomial centroid corrections and PLD (Deming et al. 2015). The main advantage of pixel-ICA over other methods relies on its purely statistical foundation without the need of imposing prior knowledge on the instrument systematics, therefore avoiding a potential source of error. The results of this paper, together with previous analyses of real Spitzer\/IRAC data sets (Morello et al. 2014, 2015), suggest that photometric precision and repeatability at the level of one part in 104 can be achieved with current infrared space instruments.","Citation Text":["Deming et al. 2015"],"Citation Start End":[[2453,2471]]}
{"Identifier":"2021MNRAS.506.5640N__Kacprzak_et_al._2010a_Instance_1","Paragraph":"We investigate the possibility of our clouds corotating within a \u2018cold-flow disc\u2019, in which material that is accreting on to the galaxy arises due to filamentary structure from the cosmic web to form an extended disc structure. As done previously by Kacprzak et al. (2010b), we investigate this by employing a simple halo model from Steidel et al. (2002) with the updated values for i and position angle (PA). In order to determine if all the low-, intermediate-, and high-ionization clouds can be explained by a corotating thick disc, we set the lagging scale height term to hv = 1000 kpc. This term is related to the location above the mid-plane where the lagging halo velocity component begins to matter and vlos becomes dominated by an exponentially declining term. Following previous studies (e.g. Steidel et al. 2002; Kacprzak et al. 2010a,b, 2011b, 2019; French & Wakker 2020; Nateghi et al. 2021) we are forcing the disc to be a thick, nearly rigid rotator, where material far above the disc plane is rotating along with the disc. This places our low- and high-ionization clouds high above the galactic disc at a high angle. The line-of-sight distance (Dlos) and line-of-sight velocity (vlos) can be found in Fig. 9. The figure shows that much of the low-, intermediate-, and high-ionization clouds (except Cloud #1) can be explained by a corotating thick disc, but the low-metallicity subsystem cannot be. Using such a large value for hv may be extreme, but it clearly illustrates the contrast between the low-metallicity subsystem and the other gas phases. In short, it is not at all possible that the low-metallicity subsystem is due to a cold-flow disc. Using a more plausible (smaller) value for h\u03bd will not affect the disparity between the low-metallicity subsystem and the other gas phases. We find that even by adopting a non-physically motivated scale height of hv = 10 kpc (adopted from Ho et al. 2017), we can still explain the metal-enriched clouds with a corotating thick disc (except Cloud #1 and Cloud #2), but we are still unable to explain the low-metallicity subsystem. This agrees with recent simulations in which high angular momentum gas accretes via a cold-flow disc, creating large corotating gaseous structures in the galaxy halo (Stewart et al. 2013, 2017). The Q0454\u2212220 system also agrees with observations of galaxy\/absorber pairs, which have been found to exhibit disc-like or accretion kinematics over a large range of impact parameters (e.g. Burchett et al. 2013; Bouch\u00e9 et al. 2016; Ho et al. 2017; Rahmani et al. 2018).","Citation Text":["Kacprzak et al. 2010a"],"Citation Start End":[[824,845]]}
{"Identifier":"2022MNRAS.516.4346G__Reboul-Salze_et_al._2021b_Instance_1","Paragraph":"The origin of the extreme magnetic fields of magnetars is still an important open question as several scenarios have been proposed, such as a convective dynamo in a fast rotating proto-neutron star (PNS) (Thompson & Duncan 1993; Raynaud et al. 2020; Raynaud, Cerd\u00e1-Dur\u00e1n & Guilet 2022), the Tayler-Spruit dynamo following fallback (Barrere et al. 2022), amplification in main-sequence stellar mergers (Schneider et al. 2019), or the fossil field scenario (Ferrario & Wickramasinghe 2006). To shed light on this question, it is important to assess the efficiency of each of the dynamo mechanisms in the conditions specific to a PNS. The impact of different ingredients on the MRI has been studied in recent years: the shear parameter (Masada et al. 2012), neutrino viscosity and drag (Guilet, M\u00fcller & Janka 2015), stable stratification (Guilet & M\u00fcller 2015; Reboul-Salze et al. 2021b), and spherical geometry (Reboul-Salze et al. 2021a). A remaining open question is the dependence on diffusion coefficients. Local disc simulations have shown that the efficiency of the MRI is strongly correlated to the magnetic Prandtl number Pm (the ratio of viscosity to resistivity) when Pm \u223c 0.1\u221216 (Fromang et al. 2007; Lesur & Longaretti 2007; Simon & Hawley 2009; Longaretti & Lesur 2010; Shi, Stone & Huang 2016; Potter & Balbus 2017) and plateaus at Pm  0.1 with an imposed external magnetic field (Meheut et al. 2015). Low and high Pm regimes are computationally challenging because a high numerical resolution is needed to resolve the small viscous (low Pm) or resistive (high Pm) scale. The regime of low Pm has attracted particular attention because it is relevant to most regions of accretion discs as well as to liquid metals (laboratory experiments and Earth core), but the opposite regime of large Pm has not been targeted specifically by previous studies of the MRI. This regime is relevant to PNSs (Thompson & Duncan 1993; Guilet et al. 2015; Lander 2021), neutron star merger remnants (Rossi, Armitage & Menou 2008), which exhibit physical conditions similar to PNSs (Guilet et al. 2017), interstellar and intergalactic media (Schekochihin et al. 2004), and to the inner parts of some accretion discs (Balbus & Henri 2008; Potter & Balbus 2014, 2017). In the latter case, the Pm dependence of the MRI may drive an instability leading to variability in the accretion rate (Potter & Balbus 2014, 2017; Kawanaka & Masada 2019).","Citation Text":["Reboul-Salze et al. 2021b"],"Citation Start End":[[859,884]]}
{"Identifier":"2018AandA...619A..13V__Saviane_et_al._2012_Instance_1","Paragraph":"The EWs were measured with the methods described in V\u00e1squez et al. (2015). As in Paper I, we used the sum of the EWs of the two strongest CaT lines (\u03bb8542, \u03bb8662) as a metallicity estimator, following the Ca\u202fII triplet method of Armandroff & Da Costa (1991). Different functions have been tested in the literature to measure the EWs of the CaT lines, depending on the metallicity regime. For metal-poor stars ([Fe\/H] \u2272 \u22120.7 dex) a Gaussian function was used with excellent results (Armandroff & Da Costa 1991), while a more general function (such as a Moffat function or the sum of a Gaussian and Lorentzian, G + L) is needed to fit the strong wings of the CaT lines observed in metal-rich stars (Rutledge et al. 1997; Cole et al. 2004). Following our previous work (Gullieuszik et al. 2009; Saviane et al. 2012) we have adopted here a G+L profile fit. To measure the EWs, each spectrum was normalised with a low-order polynomial using the IRAF continuum task, and set to the rest frame by correcting for the observed radial velocity. The two strongest CaT lines were fitted by a G+L profile using a non-linear least squares fitting routine part of the R programming language. Five clusters from the sample of Saviane et al. (2012) covering a wide metallicity range were re-reduced and analysed with our code to ensure that our EWs measurements are on the same scale as the template clusters used to define the metallicity calibration. Figure 5 shows the comparison between our EWs measurements and the line strengths measured by Saviane et al. (2012) (in both cases the sum of the two strongest lines) for the five calibration clusters. The observed scatter is consistent with the internal errors of the EW measurements, computed as in V\u00e1squez et al. (2015). The measurements show a small deviation from the unity relation, which is more evident at low metallicity. A linear fit to this trend gives the relation: \u03a3EW(S12) = 0.97 \u03a3EW(this work) + 0.21, with an rms about the fit of 0.13 \u00c5. This fit is shown in Fig. 5 as a dashed black line. For internal consistency, all EWs in this work have been adjusted to the measurement scale of Saviane et al. (2012) by using this relation. In Table 3 we provide the coordinates, radial velocities, and the sum of the equivalent widths for the cluster member stars, both measured (\u201cm\u201d) and corrected (\u201cc\u201d) to the system of Saviane et al. 2012.","Citation Text":["Saviane et al. 2012"],"Citation Start End":[[792,811]]}
{"Identifier":"2017ApJ...835..166D__Poznanski_et_al._2010_Instance_1","Paragraph":"We use the H\u03b2 \u03bb4861 velocities to standardize the SNe II because the majority of the high redshift spectra (SDSS-II and SNLS samples) are noisy and taken at early phases where the Fe ii absorption lines are not visible. Errors on the H\u03b2 velocities were obtained by measuring many times the minimum of the absorption changing the continuum fit. Both quantities are listed in Table 1. To find the best epoch to use the SCM we need the velocities for different epochs. As proposed by Hamuy (2001) and used in all SNe II cosmology works (Nugent et al. 2006; Poznanski et al. 2009; D\u2019Andrea et al. 2010; Olivares et al. 2010; Poznanski et al. 2010; Rodr\u00edguez et al. 2014; de Jaeger et al. 2015) we do an interpolation\/extrapolation using a power law of the form\n9\n\n\n\n\n\nwhere A and \u03b3 are two free parameters obtained by least-squares minimization for each individual SN and t is the epoch since the explosion. To derive the velocity error following the work done by de Jaeger et al. (2015), a Monte Carlo simulation is performed, varying randomly each velocity measurement according to the observed velocity uncertainties over more than 2000 simulations. Following Poznanski et al. (2009), we add in quadrature to the velocity uncertainty of every SN II a value of 150 km s\u22121 to account for unknown host-galaxy peculiar velocities. For the SNe II with one spectrum the same power law is used but this time with a fixed \u03b3, which is derived using only the CSP-I sample for which we have many spectra per SN and a better fit can be achieved. We find a median value of \u03b3 = \u22120.407 \u00b1 0.173. It is important to note that in the majority of other SN II cosmology works, the authors used the same power law for all of the SNe, whereas in our work the \u03b3 is different for all SNe II with more than two spectra. Additionally, in Section 6.4, we show the possibility of using a new relation between A and \u03b3 in order to derive the velocity when only one spectrum is acquired without assuming the same power-law exponent.","Citation Text":["Poznanski et al. 2010"],"Citation Start End":[[621,642]]}
{"Identifier":"2021AandA...655A.112G__Weinberger_et_al._(1999)_Instance_1","Paragraph":"HD 141569 is a Herbig star classified as a B9-A0 spectral type (Augereau & Papaloizou 2004), with an effective temperature of 9750 \u00b1 250 K, an estimated age of 7.2 \u00b1 0.02 Myr, a luminosity between 16.60 \u00b1 1.07 L\u2299 (Vioque et al. 2018) and 27.0 \u00b1 3.6 L\u2299 (Di Folco et al. 2020), a mass of 2.14 \u00b1 0.01M\u2299, and a Gaia distance of 110 \u00b1 1 pc (Arun et al. 2019)2. It is a non-flaring disk system with little mid-IR excess classified as a group-II source (Meeus et al. 2001). It is the only known pre-main sequence star characterized by a hybrid disk (Wyatt et al. 2015; P\u00e9ricaud et al. 2017; Di Folco et al. 2020), an evolutionary disk state between the protoplanetary and debris-disk regimes. Near-IR imaging spatially resolved an optically thin disk consisting of two rings located at about ~ 280 and ~ 455 au from the star Augereau et al. (1999); Weinberger et al. (1999); Biller et al. (2015). A more complex system is shown in the visible, consisting of multiple rings and outer spirals that could be explained through perturbations by two nearby (~ 7.5 arcsec) M dwarfs, or by planetary perturbations (Augereau & Papaloizou 2004; Wyatt 2005; Reche et al. 2009). Fisher et al. (2000) found a warm disk component up to 110 au at 10.8 and 18.2 \u03bcm, later confirmed by Marsh et al. (2002). The short-wavelength counterpart of this component was detected by Mawet et al. (2017) through L\u2032 imaging and ranging between 20 and 85 au. Emission at 8.6 \u03bcm was detected by Thi et al. (2014), and was interpreted as emission from polycyclic aromatic hydrocarbons (PAHs). NOEMA and ALMA observations in the millimeter range showed continuum emission equally shared between a compact (\u2272 50 au) and a smooth extended dust component (~350 au), with large millimeter grains dominating the inner regions and smaller grains in the outer ones (Di Folco et al. 2020). Finally, inner disk features were detected by SPHERE in the Y, J, H, and K bands (Perrot et al. 2016) and by Keck\/NIRC2 in the L\u2032 band (Currie et al. 2016) at physical separations of 45, 61, and 88 au. These results point out the high morphological complexity of the outer disk in the HD 141569 system.","Citation Text":["Weinberger et al. (1999)"],"Citation Start End":[[842,866]]}
{"Identifier":"2021MNRAS.500.1784D__Jaff\u00e9_et_al._2012_Instance_1","Paragraph":"Several physical processes affect galaxies inside clusters in a simultaneous way. One of these mechanisms is the ram pressure stripping (e.g. Gunn & Gott 1972; Abadi, Moore & Bower 1999; Book & Benson 2010; Steinhauser, Schindler & Springel 2016). This process can remove an important fraction of the cold gas from galaxies, resulting in the inhibition of star formation. Although this mechanism is more effective at the central regions of massive clusters, it has been reported in less massive systems (e.g. Rasmussen, Ponman & Mulchaey 2006; Jaff\u00e9 et al. 2012; Hess & Wilcots 2013). Ram pressure stripping occurs as galaxies move at high speeds through the hot ionized gas of the intracluster medium, which collides with the cold gas of the galaxies and removes it. The warm gas from the galactic halo can also be removed by the gas of the intracluster medium, a process known as starvation (e.g. Larson, Tinsley & Caldwell 1980; Balogh, Navarro & Morris 2000; McCarthy et al. 2008; Bekki 2009; Bah\u00e9 et al. 2013; Vijayaraghavan & Ricker 2015). This process can cut-off further gas cooling from the galaxy\u2019s halo gas that fuels future star formation. Kawata & Mulchaey (2008) predicted that starvation can act in galaxy groups as well. Another physical process that works on galaxies in their passage through the deep potential well of the cluster is tidal stripping (e.g. Zwicky 1951; Gnedin 2003a; Villalobos, De Lucia & Murante 2014). It can induce a central star formation burst (Byrd & Valtonen 2001), bar instabilities (\u0141okas et al. 2016), changes in the pattern of the spiral arms (Semczuk, \u0141okas & del Pino 2017), and truncate dark matter haloes (e.g. Gao et al. 2004; Limousin et al. 2009). In the outskirts of clusters, mechanisms like galaxy\u2013galaxy interaction, known as harassment, are more effective (e.g. Moore et al. 1996; Moore, Lake & Katz 1998; Gnedin 2003b; Smith et al. 2015). Most of the processes mentioned above tend to decrease or to completely suppress the star formation in galaxies. As a consequence, galaxies in clusters are typically red, early-type, with an old stellar population, and have little or none star formation at all.","Citation Text":["Jaff\u00e9 et al. 2012"],"Citation Start End":[[544,561]]}
{"Identifier":"2021ApJ...910...86R__Schmidt_et_al._2014_Instance_1","Paragraph":"One of the major endeavors of modern observational cosmology is to paint a coherent picture of the history of the universe. To this end, the final frontier remains the identification and characterization of the first sources that appeared in the universe, those which played a significant role in reionizing the intergalactic medium from a neutral state to a fully ionized one over the first billion years (corresponding to redshifts of 6 \u2272 z \u2272 12). Extragalactic surveys (of deep fields as well as lensing clusters; Grogin et al. 2011; Koekemoer et al. 2011; Bradley et al. 2012; Ellis et al. 2013; Bradley et al. 2014; Schmidt et al. 2014; Treu et al. 2015; Lotz et al. 2017; Salmon et al. 2018; Coe et al. 2019) with the Hubble Space Telescope (HST) have yielded impressive gains in the number of galaxy candidates at redshifts z = 7\u201310, with samples reaching over 1000 objects, and revolutionized our understanding of galaxy evolution therein. Complementing these observations, the spectroscopic confirmation (e.g., Finkelstein et al. 2013; Oesch et al. 2015; Zitrin et al. 2015; Roberts-Borsani et al. 2016; Hoag et al. 2017; Stark et al. 2017; Hashimoto et al. 2018) and characterization (e.g., Laporte et al. 2017; Mainali et al. 2018; Endsley et al. 2021) of over a dozen sources has seen impressive advances with ground-based spectroscopy (e.g., probing the rest-frame UV and FIR with Keck\/MOSFIRE, VLT\/X-Shooter, and ALMA), particularly for the brightest and rarest objects. For the rest-frame optical, however, the Spitzer Space Telescope has, until now, afforded the only realistic means for statistical analyses. However, the Infrared Array Camera\u2019s (IRAC) coarse spatial resolution and the limited depth probed by many surveys make robust and uncontaminated constraints on galaxy properties a challenging feat. Further advances with current facilities are challenging owing to the limited wavelength coverage of HST and the observed faintness of star-forming galaxies as one approaches redshifts of z > 10. The imminent arrival of the James Webb Space Telescope (JWST) has the potential to detect galaxies well beyond the current frontier of z \u223c 12 (e.g., Behroozi et al. 2020) thanks to the unprecedented resolution and sensitivity of its near-IR (NIR) imaging and spectroscopic capabilities, and revolutionize our current understanding of galaxy evolution.","Citation Text":["Schmidt et al. 2014"],"Citation Start End":[[621,640]]}
{"Identifier":"2016AandA...592A..54A__Hill_et_al._2012_Instance_1","Paragraph":"Regardless of whether the NGC 6334 filament will form massive stars, our ArT\u00e9MiS result that the filament inner width is within a factor of 2 of 0.1 pc has interesting implications. Our NGC 6334 study is clearly insufficient to prove that interstellar filaments have a truly universal inner width, but it shows that the finding obtained with Herschel in nearby clouds is not limited to filaments in low-mass star-forming regions. It is quite remarkable that the NGC 6334 filament has almost the same inner width as the faint subcritical filaments in Polaris (cf. Men\u2019shchikov et al. 2010; Arzoumanian et al. 2011), the marginally supercritical filaments in Musca and Taurus (Cox et al. 2016; Palmeirim et al. 2013), or the lower-mass supercritical filaments in Serpens South and Vela C (Hill et al. 2012), despite being three orders of magnitude, two orders of magnitude, and at least a factor of ~3 denser and more massive than these filaments, respectively (see Table 1). While not all of these filaments may have necessarily formed in the same way, this suggests that a common physical mechanism is responsible for setting the filament width at the formation stage and that the subsequent evolution of dense filaments, through accretion of background cloud material (cf. Heitsch 2013; Hennebelle & Andr\u00e9 2013) for example, is such that the inner width remains at least approximately conserved with time. A promising mechanism for creating dense filaments, which may be quite generic especially in massive star-forming complexes, is based on multiple episodes of large-scale supersonic compression due to interaction of expanding bubbles (Inutsuka et al. 2015). With about seven bubble-like HII regions per square degree (Russeil et al. 2013, see also Fig. 2), there is ample opportunity for this mechanism to operate in NGC 6334. More specifically, at least in projection, the NGC 6334 filament appears to be part of an arc-like structure centered on the HII region Gum 63 (see Fig. 2a), suggesting the filament may partly result from the expansion of the associated bubble. Interestingly, the background column density is one order of magnitude higher for the NGC 6334 filament than for the other filaments of Table 1, which is suggestive of a significantly stronger compression. More observational studies are needed to investigate the structure and environment of a larger number of filaments in massive star-forming regions to determine whether the characteristics of the NGC 6334 filament are generic. More theoretical work is also needed to better understand the physics controlling the width of interstellar filaments. ","Citation Text":["Hill et al. 2012"],"Citation Start End":[[787,803]]}
{"Identifier":"2021MNRAS.505.2039P__Xu,_Wang_&_Chen_2014_Instance_1","Paragraph":"The purpose of this paper is to demonstrate a complementary approach to single-dish IM experiments, capable of the statistical detection of H\u2009i field and therefore the fluctuations in the underlying matter field by measuring the H\u2009i power spectrum with interferometric observations. Interferometers have inherent advantages over single-dish measurements. Besides providing high angular resolutions, they are less sensitive to systematics which poses a major problem to the autocorrelation power. However, the smallest k-modes accessible to an interferometer is determined by the shortest baselines which may hinder probing the BAO scales. Interferometers such as CHIME (Bandura et al. 2014), TIANLAI (Xu, Wang & Chen 2014), and HIRAX (Newburgh et al. 2016) are custom designed to probe the BAO scales using the 21\u2009cm signal in the redshift range z \u223c 0.5\u20132. In this paper, we study the feasibility of detecting the cosmological 21\u2009cm signal with MeerKAT and present forecasts on the statistical measurement of H\u2009i with MeerKAT L-band (856 MHz  \u03bd  1712 MHz) observations on quasi-linear scales. We present the sensitivity estimates at z \u223c 0.27 from our newly developed simulation pipeline, which is our first attempt towards the measurement of H\u2009i power spectrum with MeerKAT. Our pipeline is based on the methods being developed for similar statistical measurement of H\u2009i from the Epoch of Reionization from a series of experiments at lower radio frequencies such as LOFAR (Van Haarlem et al. 2013), GMRT (Paciga et al. 2013), PAPER (Parsons et al. 2014), HERA (DeBoer et al. 2017), and MWA (Tingay et al. 2013). We show that with MIGHTEE (Jarvis et al. 2016; Maddox et al. 2021), one of MeerKAT\u2019s large survey projects, we can achieve constraints on the H\u2009i power spectrum at z = 0.27. There are implicit advantages with such survey projects with a specific emission line. First, it provides a one to one correspondence between observed frequency and redshift, thereby delivering a very high redshift resolution. Secondly, these are generally less time consuming compared to an optical spectroscopic galaxy survey which requires very high sensitivity to detect individual galaxies.","Citation Text":["Xu, Wang & Chen 2014"],"Citation Start End":[[701,721]]}
{"Identifier":"2021AandA...650A.164M__Davies_et_al._2012_Instance_1","Paragraph":"The GMC associated with G305 is one of the most massive and luminous clouds in the Milky Way (Fig. 1). It is located in the Galactic plane at l ~ 305\u00b0, b ~ 0\u00b0 and at a kinematic distance of 4 kpc (derived from a combinationof radio and H\u03b1 observationsby Clark & Porter (2004); Davies et al. (2012) measured its spectrophotometric distance to be 3.8 \u00b1 0.6 kpc and most recently Borissova et al. (2019) measured the Gaia DR2 average distance to be 3.7 \u00b1 1.2 kpc); this places it in the Scutum-Crux spiral arm. Given this distance, the complex has a diameter of ~ 30 pc (Clark & Porter 2004) and a molecular mass of ~6 \u00d7 105 M\u2299 (Hindson et al. 2010). The G305 complex consists of a large central cavity that has been cleared by the winds from massive stars belonging to two visible central clusters (Danks 1 and 2) and the Wolf-Rayet star (WR48a; Clark & Porter 2004; Davies et al. 2012). The cavity is surrounded by a thick layer of molecular gas (traced by CO and NH3 emission; Hindson et al. 2010, 2013). Radio continuum observations by Hindson et al. (2012) have revealed that the cavity is filled with ionized gas and identified six ultra-compact HII (UC HII) regions and also one bright rimmed cloud (BRC) at the periphery of the cavity, indicating molecular gas irradiated by UV radiation (Sugitani & Ogura 1994; Thompson et al. 2004), which may cause implosion (Bertoldi 1989) or evaporation. A number of studies havereported star formation tracers (water and methanol masers, HII regions and massive young stellar objects, MYSOs; Clark & Porter 2004; Lumsden et al. 2013; Urquhart et al. 2014; Green et al. 2009, 2012). Furthermore, Hindson et al. (2010) found the concentration of star formation tracers to be enhanced inside a clump of NH3 bearing molecular gas that faces the ionizing sources, which is consistent with the hypothesis that the star formation has been triggered. Analysis of the stellar clusters in the complex reveals them to have ages of 1.5 Myr for Danks 1 and 3 Myr for Danks 2,with the former possibly being triggered by the latter (Davies et al. 2012). Additionally, a diffuse population of evolved massive stars was also found to exist within the confines of the G305 complex that had formed around the same time as the two clusters (Leistra et al. 2005; Shara et al. 2009; Mauerhan et al. 2011; Davies et al. 2012; Faimali et al. 2012; Borissova et al. 2019).","Citation Text":["Davies et al. (2012)"],"Citation Start End":[[277,297]]}
{"Identifier":"2021ApJ...907...47L__Lee_et_al._2019_Instance_2","Paragraph":"In Figure 8, we also find small differences in the [Na, Al, O\/Fe] abundance ratios between the stars in the bright and faint RC groups, although it is not as clear as in the case of [Na, Al, O\/H] abundances. In particular, unlike Figure 7, stars in the bRC group are more enhanced in [Na\/Fe] but appear to be more depleted in [Al\/Fe] and [O\/Fe] than those in the fRC group. The mean differences are 0.053 \u00b1 0.021 dex, 0.032 \u00b1 0.018 dex, and 0.071 \u00b1 0.045 dex in [Na\/Fe], [Al\/Fe], and [O\/Fe], respectively, which are marginally significant at p-values of 0.22, 0.18, and 0.23. When the relative fraction of RC stars is taken into account (27%; see Section 4), the difference in [Na\/Fe] between the genuine RC stars would correspond to \u0394[Na\/Fe] \u223c 0.20 dex, which is comparable to that expected from our chemical evolution model for the bulge stars (\u0394[Na\/Fe] = 0.2 \u223c 0.3 dex; Kim & Lee 2018; Lee et al. 2019).10\n\n10\nThe previous study by Lee et al. (2019) noted a clear separation of the two groups according to Na abundance among bright RGB stars in the outer bulge. The apparent lack of such a distinct difference between the two groups in this study may be due to a larger uncertainty on abundances of relatively faint sample stars.\n The overall chemical patterns, however, are not identical to those observed in typical GCs, where the later-generation stars are more enhanced in [Na, Al\/Fe] and more depleted in [O, Mg\/Fe] than the first-generation stars at a given metallicity, although the trend of [Na, Al O\/Fe] between the two RCs is less clear. Figure 9 shows the comparison of stars in this study with stars in metal-rich GCs ([Fe\/H] > \u22121.0) on the Na\u2013O diagram. The stars used in this study have a different distribution from stars in GCs. Although the bRC group is slightly more enhanced in [Na\/Fe] and more depleted in [O\/Fe] than the fRC group, the [Na\/Fe] variation of RC stars is smaller than that of GC stars. This discrepancy might imply the different chemical evolution between stars in the bulge and typical GCs. We note, however, that even though we employ only metal-rich GC stars for the comparison, the majority of stars are still far more metal-poor ([Fe\/H]  \u22120.5) than stars in the bulge. Because the relatively small [Na\/Fe] variation is expected from the chemical evolution model for metal-rich bulge stars and the O-depletion is indistinct in some metal-rich GCs, such as NGC 6121 and 47 Tuc (see Kim & Lee 2018; Lee et al. 2019), the direct comparison of bulge stars with similarly metal-rich GCs on the Na\u2013O plane would require further spectroscopic observations for such GCs in the bulge.","Citation Text":["Lee et al. (2019)"],"Citation Start End":[[935,952]]}
{"Identifier":"2020AandA...638A.140W__Tricco_et_al._2016_Instance_2","Paragraph":"Using the constrained hyperbolic divergence cleaning scheme with variable cleaning speed from Tricco et al. (2016), we can keep the divergence error low in all cases. The mean normalized divergence error, \u27e8\u03f5divB\u27e9=\u27e8h|\u2207\u2005\u22c5\u2005B|\/|B|\u27e9, is typically of order 10\u22125\u2005\u2212\u200510\u22123. In Fig. 17, we show the normalized divergence error maps for several test problems. Again we see that the divergence cleaning works extremely well here, the maximum error is generally around 10\u22122. Comparing to Hopkins (2016, their Fig. 4), we find that the errors are smaller than their MFM simulations with the Dedner et al. (2002) cleaning in general, with an exception for the outskirt of the advection loop (where the magnetic field is essentially zero and thus not important for the result). The improvement is probably due to the more advanced constrained cleaning method (Tricco et al. 2016). The normalized divergence error for the \u03bc\u2004=\u200410 cloud-collapse case at the jet launching time is shown in Fig. 18. Here, the divergence cleaning still performs very well in the disc, along the jets and for the majority of the regions where the outflow interacts with the ambient gas, especially when the divergence error is compared to the total gas pressure (right panel). The result is similar to the Dedner cleaning in Hopkins (2016), although our error is somewhat larger at the tip of the jets where the gas is shocked. However, we note that the comparison is not direct in this case as the jets may develop differently. Overall, the result from cleaning is still worse than the constrained transport or constrained gradient schemes (Hopkins 2016). For SPMHD, as shown in Tricco & Price (2012), divergence errors can be reduced to machine precision (or more practically to a certain tolerance value) using cleaning, with the help of a sub-cycling routine. However, local adjustments are required to determine the number of iterations for each particle to efficiently subcycle the cleaning in the simulation. This is because certain regions are more affected than others and because divergence is spread to nearby neighbours. Conceivably, if vector potentials (Stasyszyn & Elstner 2015) could work for a wider range of problems this could be an interesting avenue as well. However, the exploration of these methods in detail is beyond the scope of this work.","Citation Text":["Tricco et al. 2016"],"Citation Start End":[[843,861]]}
{"Identifier":"2015MNRAS.452.3994M___2008_Instance_1","Paragraph":"Lower limit on \u03bc: at the highest accretion rates, to observe pulsations, the magnetic field must be at least high enough to truncate the accretion disc at or above the neutron star surface. Thus by setting rt = Rs we obtain the lower limit on the dipole moment as\n\n(5)\n\n\\begin{equation}\n\\mu _{\\rm min} = \\gamma _{\\rm B}^{-1\/2}(GM)^{1\/4}\\dot{M}_{\\rm max}^{1\/2}R_{\\rm s}^{7\/4}.\n\\end{equation}\n\nWe assume the mass accretion rate can be estimated from the bolometric luminosity as $L = GM\\dot{M}\/R_{\\rm s}$, and estimate L from the observed luminosity in the X-ray band by applying a bolometric correction factor (L = \u03f5bolLX). The typical reported values of the bolometric correction factor have a range of \u03f5bol \u223c 1\u20132 (Gilfanov et al. 1998; Galloway et al. 2002, 2008; Campana et al. 2003; Migliari & Fender 2006; Casella et al. 2008). The mass accretion rate then follows as\n\n(6)\n\n\\begin{eqnarray}\n\\dot{M} &=& 10^{16} \\mbox{ g s}^{-1} \\left( \\frac{ \\epsilon _{\\rm bol} L_{\\rm X}}{1.87 \\times 10^{36}\\mbox{ erg s}^{-1}} \\right)\\nonumber\\\\\n&&\\times\\,\\left( \\frac{M}{1.4 \\,\\mathrm{M}_{\\odot }} \\right)^{-1} \\left(\\frac{R_{\\rm s}}{10 \\mbox{ km}}\\right).\n\\end{eqnarray}\n\nExpressing the X-ray luminosity in terms of the observed X-ray flux and the assumed distance (LX = 4\u03c0d2F), we obtain the lower limit on the magnetic moment\n\n(7)\n\n\\begin{eqnarray}\n\\mu _{\\rm min} &=& 9.36 \\times 10^{24} \\mbox{ G cm}^{-3} \\left( \\frac{ \\gamma _{\\rm B} }{ 1 } \\right)^{-1\/2} \\left( \\frac{ \\epsilon _{\\rm bol} }{ 1 } \\right)^{1\/2} \\nonumber \\\\\n&&\\times \\left( \\frac{ F_{\\rm max}}{10^{-10}\\, \\mbox{erg}\\mbox{ cm}^{-2}\\mbox{ s}^{-1}} \\right)^{1\/2} \\left( \\frac{d}{10 \\mbox{ kpc}}\\right) \\nonumber \\\\\n&&\\times \\left(\\frac{M}{1.4 \\,\\mathrm{M}_{\\odot }}\\right)^{-1\/4} \\left( \\frac{R_{\\rm s}}{10 \\mbox{ km}}\\right)^{9\/4},\n\\end{eqnarray}\n\nwhere Fmax is the highest observed X-ray flux with pulsation and we adopted the boundary values of \u03b3B and \u03f5bol that give the lowest magnetic moment. For a detailed discussion of these assumptions we again refer to Section 4.","Citation Text":["Galloway et al.","2008"],"Citation Start End":[[737,752],[759,763]]}
{"Identifier":"2015AandA...584A.103S__Chamel_et_al._2011_Instance_2","Paragraph":"Douchin & Haensel (2001; DH) formulated a unified EoS for NS on the basis of the SLy4 Skyrme nuclear effective force (Chabanat et al. 1998), where some parameters of the Skyrme interaction were adjusted to reproduce the Wiringa et al. calculation of neutron matter (Wiringa et al. 1988) above saturation density. Hence, the DH EoS contains certain microscopic input. In the DH model the inner crust was treated in the CLDM approach. More recently, unified EoSs for NS have been derived by the Brussels-Montreal group (Chamel et al. 2011; Pearson et al. 2012; Fantina et al. 2013; Potekhin et al. 2013). They are based on the BSk family of Skyrme nuclear effective forces (Goriely et al. 2010). Each force is fitted to the known masses of nuclei and adjusted among other constraints to reproduce a different microscopic EoS of neutron matter with different stiffness at high density. The inner crust is treated in the extended Thomas-Fermi approach with trial nucleon density profiles including perturbatively shell corrections for protons via the Strutinsky integral method. Analytical fits of these neutron-star EoSs have been constructed in order to facilitate their inclusion in astrophysical simulations (Potekhin et al. 2013). Quantal Hartree calculations for the NS crust have been systematically performed by (Shen et al. 2011b,a). This approach uses a virial expansion at low density and a RMF effective interaction at intermediate and high densities, and the EoS of the whole NS has been tabulated for different RMF parameter sets. Also recently, a complete EoS for supernova matter has been developed within the statistical model (Hempel & Schaffner-Bielich 2010). We shall adopt here the EoS of the BSk21 model (Chamel et al. 2011; Pearson et al. 2012; Fantina et al. 2013; Potekhin et al. 2013; Goriely et al. 2010) as a representative example of contemporary EoS for the complete NS structure, and a comparison with the other EoSs of the BSk family (Chamel et al. 2011; Pearson et al. 2012; Fantina et al. 2013; Potekhin et al. 2013) and the RMF family (Shen et al. 2011b,a) shall be left for future study. ","Citation Text":["Chamel et al. 2011"],"Citation Start End":[[1723,1741]]}
{"Identifier":"2019MNRAS.485..440P__Dunne_et_al._2003_Instance_1","Paragraph":"Cassiopeia A (Cas A) is a Galactic SNR located $3.4 \\, {\\rm kpc}$ away (Reed et al. 1995), with an age of approximately $330 \\, {\\rm yr}$ (Fesen et al. 2006) and a radius of $1.7 \\, {\\rm pc}$ (Reed et al. 1995). It provides a unique laboratory to test the efficiency of dust condensation in SN ejecta, and the subsequent destruction of dust by the reverse shock, and as such has been studied extensively in the past. IR\/sub-mm observations have led to derived dust masses ranging from ${\\sim } 10^{-4} \\, {\\rm M}_\\odot$ of hot ($T \\sim 100 \\, {\\rm K}$) dust (Arendt, Dwek & Moseley 1999; Douvion, Lagage & Pantin 2001) to $2{\\text{-}}4 \\, {\\rm M}_\\odot$ of cold dust emitting at sub-mm wavelengths (Dunne et al. 2003), although this higher mass has been attributed to foreground dust emission in the ISM (Krause et al. 2004). Analyses of integrated fluxes from Spitzer and Herschel observations (Rho et al. 2008; Barlow et al. 2010; Arendt et al. 2014) found ${\\sim } 0.01 \\, {\\rm M}_\\odot$ of hot dust, with ${\\sim } 0.1 \\, {\\rm M}_\\odot$ of cold, unshocked dust present in the central regions, in agreement with simulations of the dust formation and evolution in Cas A by Nozawa et al. (2010). Dunne et al. (2009) suggested the observed polarization of the sub-mm emission could be explained by ${\\sim } 1 \\, {\\rm M}_\\odot$ of cold dust, similar to the value of $1.1 \\, {\\rm M}_\\odot$ given by Bevan et al. (2017) as the most likely mass based on the shape of emission-line profiles affected by extinction. De Looze et al. (2017) utilized spatially resolved Herschel and Spitzer observations of Cas A to fit the dust continuum emission, following the removal of line and synchrotron contamination, using a four-component model including ISM dust emission and three SNR dust temperature components. They found a large mass of unshocked cold dust in the centre of the SNR (up to $0.6 \\, {\\rm M}_\\odot$), significantly above previous estimates based on the IR\/sub-mm emission.","Citation Text":["Dunne et al. 2003"],"Citation Start End":[[699,716]]}
{"Identifier":"2019AandA...632A.129W__Feng_&_Wang_2015_Instance_1","Paragraph":"In this study, the 272 eV suprathermal electron pitch-angle distributions (PADs) measured by ACE are used. The electron PADs are obtained from the Solar Wind Electron Proton Alpha Monitor (SWEPAM) with angular and time resolutions of 9\u00b0 and 64 s respectively (McComas et al. 1998). Here we examined 16 s average magnetic field, 64 s average plasma, 1 h average O7+\/O6+ ratio, and mean Fe charge state \u27e8Fe\u27e9 data from 1998 to 2008 measured by ACE and identified 272 ICMEs in total. The ICMEs were identified by the following process: (1) We take the events in previous ICME lists of Jian et al. (2006), Chi et al. (2016), and Richardson & Cane (2004)1 as candidate ICMEs. (2) Some lists also report short-duration ( 10 h) structures as ICMEs. As the origin of these smaller-scale ICMEs and flux ropes are still debated (Feng et al. 2007; Rouillard et al. 2011; Janvier et al. 2014; Feng & Wang 2015; Wang et al. 2019), we excluded them from this study. (3) The high Fe charge states (\u27e8Fe\u27e9\u2004\u2265\u200412) and abnormally high O7+\/O6+ ratio (\u22651) are the result of flare-related heating in the corona (Lepri & Zurbuchen 2004; Reinard 2005), and therefore they are independently reliable ICME indicators (Feng & Wang 2015). If the candidate ICMEs have high Fe charge states and\/or abnormally high O7+\/O6+ ratio, they are identified as ICMEs. (4) If the candidate ICMEs have no high Fe charge states and abnormally high O7+\/O6+ ratio, we look for the following five characteristics: declining speed (apparent expansion), increasing total magnetic magnitude and helium abundance (He\/P > 0.06) (Richardson & Cane 2004), and decreasing proton temperatures and proton densities. If the candidate ICMEs have three or more of the above characteristics they are identified as ICMEs. Given that magnetic flux ropes are special field topologies characterized by bundles of helical magnetic-field lines collectively spiraling around a common axis, the essential observational properties of magnetic flux ropes should be enhanced magnetic field strength and smooth rotations (Feng et al. 2008, 2010), namely, measured enhanced magnetic field strength, the center-enhanced magnetic components, and bipolar curve magnetic components. Therefore, if an ICME was found to have enhanced magnetic field strength, both center-enhanced and bipolar field components, it was identified as MC. Among the 272 ICMEs, 101 (37.1%) events were identified as MCs. All 272 ICMEs are listed in Table A.1. The second and third columns show the start and end times, the fourth column gives the duration of the ICMEs, and the fifth column provides the types of ICMEs (MC or nonMC).","Citation Text":["Feng & Wang 2015"],"Citation Start End":[[880,896]]}
{"Identifier":"2015MNRAS.450..982K__Hardee_1979_Instance_1","Paragraph":"The traditional way of studying instabilities of non-linear dynamical systems is via linear stability analysis. In many cases, it leads to much simpler system of linearized equations, which can be solved analytically. However, in many other cases even the linearized system does not allow general solutions in terms of analytic functions for arbitrary equilibrium configuration. One way to overcome this problem is to restrict the analysis to special equilibrium configurations which allow us to simplify the linearized system of equations even further. Accordingly, most early studies of jet stability assumed simplified jet structure, including the magnetic field topology (e.g. Hardee 1979, 1982; Cohn 1983; Payne & Cohn 1985; Istomin & Pariev 1996; Begelman 1998; Lyubarskii 1999; Tomimatsu, Matsuoka & Takahashi 2001; Narayan et al. 2009). This allowed to obtain the solution in terms of Bessel and hypergeometric functions. Although very useful in many respects, this approach still cannot address the stability of jets with more complex and more realistic structure. In particular, all these equilibrium jets included surface currents, whereas Gourgouliatos et al. (2012) argued that such current sheets are likely to promote resistive jet instabilities. A jet that is free of current sheets would be free of resistive instabilities in computer simulations. As an alternative, they constructed equilibrium solutions which are current-sheet-free structures. There is an intuitive reason for expecting a magnetized jet without a current sheet to be more stable \u2013 such jets carry no volumetric net current. If the surface of a magnetized jet with current sheet is perturbed, two neighbouring perturbations could behave analogously to two parallel current carrying wires that are carrying current in the same direction. Since such wires attract, one might analogously expect the surface of such a jet to become more corrugated, i.e. the perturbations can grow. By avoiding current sheets on the surface, a current-sheet-free jet avoids this mode of destabilization. Another interesting feature of these solutions is that the jets carry zero net current and magnetic flux. In some part of the jet cross-section the poloidal electric current flows outwards and in the rest of the cross-section exactly the same amount of the current flows in the opposite direction. Thus, one does not have to worry of having a return current outside of the jet on large scales. The same applies to the poloidal magnetic field. Unfortunately, the magnetic structure of these solutions is too complex for the linearized equations to allow analytical solutions.","Citation Text":["Hardee 1979"],"Citation Start End":[[681,692]]}
{"Identifier":"2018MNRAS.473..513F__Patat_et_al._1994_Instance_1","Paragraph":"Type II supernovae (SNe) are defined by the prominent hydrogen lines in their spectra. They are believed to originate from the collapse of an iron core of massive stars (\u22738 M\u2299) that retain their hydrogen envelope. The most common sub-type, comprising \u223c70 per cent of all Type II SNe, is characterized by a phase of a roughly constant magnitude in the optical bands, hence their name type II Plateau (II-P). This plateau phase typically starts 1\u20132 weeks after the explosion and lasts for \u223c100\u2009d. Pre-explosion images have revealed that the progenitors of this class are red supergiants, in the mass range of (7\u201316 M\u2299) (Smartt 2015; for individual progenitor detections see e.g. Van Dyk, Li & Filippenko 2003a; Van Dyk, Li & Filippenko 2003b; Van Dyk et al. 2012). Type II Linear (II-L) SNe constitutes another subclass of type II SNe (e.g. Patat et al. 1994; Arcavi et al. 2012; Faran et al. 2014a,b). They are spectroscopically very similar to type II-P events (Faran et al. 2014b, see), but their light curves are declining in all bands. In both types (II-P and II-L), there is typically a sharp drop in the luminosity after \u223c100\u2009d and the luminosity starts to follow roughly the exponential decay expected from emission powered by the decay of 56Ni. The distinction between these two classes is not well defined, and studies have used different definitions for II-L SNe. However, several recent works have shown that there exists a continuum of decline rates between slow declining and fast declining SNe (Anderson et al. 2014; Faran et al. 2014a), which suggests that a separation into two different classes may be artificial. Other type II sub-classes will not be discussed here. In this paper, we focus on the light curves of type II-P and type II-L SNe, without making the distinction between the two types, and refer to them in short as Type II SNe. Our goal here is to perform a uniform analysis of the bolometric luminosity and temperature evolution of a large sample of Type II SNe and to compare our findings to theoretical models, focusing on the transition to the plateau, which takes place during the first 2 weeks.","Citation Text":["Patat et al. 1994"],"Citation Start End":[[839,856]]}
{"Identifier":"2017ApJ...839...72A__Emsellem_et_al._1994_Instance_1","Paragraph":"We fit the radial dispersion profiles of each UCD to dynamical models using the Jeans Anisotropic Models (JAM) method with the corresponding code discussed in detail in Cappellari (2008). To briefly summarize, the dynamical models are made in a series of steps making two general assumptions: (1) the velocity ellipsoid is aligned with the cylindrical coordinate system (\n\n\n\n\nR\n,\nz\n,\n\u03d5\n\n\n), (2) the anisotropy is constant. Here, the anisotropy is defined as \n\n\n\n\n\n\n\u03b2\n\n\nz\n\n\n=\n1\n\u2212\n\n\n(\n\n\n\u03c3\n\n\nz\n\n\n\n\/\n\n\n\n\u03c3\n\n\nR\n\n\n)\n\n\n2\n\n\n\n\n where \n\n\n\n\n\n\n\u03c3\n\n\nz\n\n\n\n\n is the velocity dispersion parallel to the rotation axis and \n\n\n\n\n\n\n\u03c3\n\n\nr\n\n\n\n\n is the velocity dispersion in the radial direction in the plane of the galaxy. The first step in the dynamical modeling process is to construct a three-dimensional mass model by deprojecting the two-dimensional mass model MGEs discussed in the previous section. In the self-consistent case, the luminosity and mass profile are the same. However, in our case, we used the mass profile to construct the potential and we used the light profile to calculate the observable properties of the model, both described below. The choice to parameterize the light profile with MGEs is motivated by the ease of deprojecting Gaussians and the accuracy in reproducing the surface brightness profiles (Emsellem et al. 1994; Cappellari 2002). The second step in the dynamical modeling process is to construct a gravitational potential using our mass model. This potential also contains a Gaussian to represent a supermassive BH with the axis ratio, q = 1, and width, \n\n\n\n\n\u03c3\n\u2272\n\n\nr\n\n\nmin\n\n\n\n\/\n\n3\n\n\n, where rmin is the smallest distance from the BH that needs to be accurately modeled. Although a supermassive BH can be modeled by adding a Keplerian potential, it is much simpler to model the BH as this small Gaussian (Emsellem et al. 1994). Next, the MGE formalism is applied to the solution of the axisymmetric anisotropic Jeans equations (see Section 3.1.1 of Cappellari 2008). Finally, the intrinsic quantities are integrated along the LOS and convolved with the PSF from the kinematic data to generate observables that can be compared with the radially binned dispersion profiles. Supermassive BH masses are frequently measured with dynamical models that allow for fully general distribution functions (e.g., Schwarzschild), which is important to include because of the BH mass-anisotropy degeneracy in explaining central dispersion peaks in galaxies. Since plunging radial orbits have an average radius that is far from the center of the galaxy, these orbits can raise the central dispersion without significantly enhancing the central mass density. Similarly, a supermassive BH also raises the dispersion near the center of the galaxy. Other dynamical modeling techniques break this degeneracy by fitting the full orbital distribution without assumptions about the anisotropy. However, given the quality of our kinematic data, a more sophisticated dynamical modeling technique is not feasible; we further discuss the assumptions and limitations of our modeling at the beginning of Section 5.1.","Citation Text":["Emsellem et al. 1994"],"Citation Start End":[[1308,1328]]}
{"Identifier":"2022MNRAS.509.3599T__Du_et_al._2015_Instance_1","Paragraph":"Here we report the X-ray spectral and timing analysis of the joint XMM\u2013Newton and NuSTAR observations of an IRAS 04416+1215, a nearby (z = 0.0889; Boller et al. 1992) hyper-Eddington AGN. The source is part of a XMM\u2013Newton\/NuSTAR campaign that aims to constrain the broad-band X-ray properties of eight super-Eddington AGN from the best sample of bona fide super-Eddington sources available, i.e. super-Eddington accreting massive black holes (SEAMBHs; Du et al. 2014, 2015; Wang et al. 2014) that contains exclusively objects with black hole masses estimated from reverberation mapping. In this campaign we are carrying out to study the broad-band X-ray properties of super-Eddington AGN, all the sources have new NuSTAR observations performed simultaneously with XMM\u2013Newton or Swift-X-ray Telescope (XRT). IRAS 04416+1215 has bolometric luminosity $\\log (L_{\\rm bol}\/\\rm erg\\, s^{-1})=47.55$, according to Castell\u00f3-Mor, Netzer & Kaspi (2016), and $\\log (L_{\\rm bol}\/\\rm erg\\, s^{-1})=45.52$, according to Liu et al. (2021). The former estimate is computed using, for the SED fitting procedure, the Slone & Netzer (2012) code, including the comparison of the observed SED with various combinations of disc SEDs covering the range of mass, accretion rate, spins, and taking into account the correction for intrinsic reddening and host galaxy contribution. In the latter estimate, the SED fitting is done using the more semplicistic templates from Krawczyk et al. (2013). The dimensionless accretion rate (Du et al. 2014) and black hole mass of the source are $\\log (\\dot{\\mathscr {M}})$ = $2.63^{+0.16}_{-0.67}$ and log\u2009(MBH\/M\u2299) = $6.78^{+0.31}_{-0.06}$ with the reverberation mapping technique (Du et al. 2015), respectively, where $\\dot{\\mathscr {M}}\\equiv \\dot{M}_{\\bullet }c^2\/L_{\\rm Edd}$, $\\dot{M}_{\\bullet }$ is mass accretion rates, c is speed of light, and LEdd is the Eddington luminosity. The dimensionless accretion rate is estimated by $\\dot{\\mathscr {M}}=20.1\\, \\ell _{44}^{3\/2}M_7^{-2}$ from the Shakura\u2013Sunyaev disc model (Du et al. 2015), where \u211344 is the 5100 \u00c5 luminosity in units of $10^{44}\\, {\\rm erg\\, s^{-1}}$ and $M_7=M_{\\bullet }\/10^7\\, \\mathrm{M}_{\\odot }$. This approximation is valid for $\\dot{\\mathscr {M}}\\lesssim 10^3$. To compute the Eddington ratio we assumed the bolometric luminosity value from Castell\u00f3-Mor et al. (2016), which is a better and more trustable estimate of the bolometric luminosity of the source, obtaining \u03bbEdd \u223c 472. This value is in perfect agreement with the dimensionless accretion rate from Du et al. (2014). However even assuming the luminosity from Liu et al. (2021), with which the value of the accretion rate would be \u03bbEdd \u223c 4.40, the source would remain a super-Eddington accreting AGN. IRAS 04416+1215 turned out to be the most peculiar of our sample, it is classified as NLS1 galaxy, showing narrow H\u03b2 line [full width at half-maximum (FWHM) = $1670 \\, \\rm km \\, \\rm s^{-1}$; Moran, Halpern & Helfand 1996] and very broad [O\u2009iii] (FWHM = $1150 \\, \\rm km \\, \\rm s^{-1}$; V\u00e9ron-Cetty, V\u00e9ron & Gon\u00e7alves 2001) lines, which is typically found in sources accreting at such high Eddington accretion rates (Greene & Ho 2005; Ho 2009). The source shows a photon index in the Roentgen Satellite (ROSAT) (0.1\u20132.4 keV) energy band, of \u0393 = 2.96 \u00b1 0.50 (Boller et al. 1992) and of $\\Gamma =2.46^{+0.27}_{-0.26}$ for the rest-frame >2 keV spectrum, according to Liu et al. (2021).","Citation Text":["Du et al.","2015"],"Citation Start End":[[453,462],[469,473]]}
{"Identifier":"2021MNRAS.504..444C__Belloni_et_al._2020_Instance_1","Paragraph":"The strong radio flare observed on MJD 58523 is likely to be associated with the ejection of RK1, but the mechanism capable of launching and accelerating these relativistic plasma knots is still not understood. The current picture links the launch of the jets with the transition from the HS to the SS, i.e. during the IMS, where usually ejections take place prior to strong radio flares (e.g. Fender 2006; Miller-Jones et al. 2012) and are marked by changes in the X-ray timing and spectral properties (Belloni et al. 2005). RK1 fits very well in this picture, as the inferred ejection date tej, lin + Sedov = MJD 58518.9 \u00b1 2.4 (see Section 3.3) is only 4\u2009d before the strong radio flare observed with MeerKAT, and when MAXI J1348\u2013630 was in the IMS. The presence of Type-B QPOs (tentatively linked to discrete ejections, e.g. Soleri et al. 2008; Fender et al. 2009; Miller-Jones et al. 2012; Homan et al. 2020) has been detected with NICER close to the first hard-to-soft state transition during the main outburst of MAXI J1348\u2013630 (Belloni et al. 2020; Zhang et al. 2020a). In particular, Type-B QPOs started to be detected on MJD 58522.6 (Zhang et al. 2020a). This is \u22733\u2009d after our inferred ejection date, similar to what was already observed for H1743\u2013322 (Miller-Jones et al. 2012) and MAXI J1535\u2013571 (Russell et al. 2019a). However, a detailed discussion of the NICER timing results in relation to the inferred ejection dates will be discussed in a forthcoming paper. The \u223c5 arcsec MeerKAT resolution does not allow us to resolve the two components soon after the ejection, while the ATCA observations were short (providing a non-optimal uv-coverage). As the synchrotron-emitting component is launched and expands, it is predicted to have an optically thick rising phase due to an increasing surface area, followed by an optically thin decay produced by adiabatic expansion losses (van der Laan 1966). Flares produced by discrete ejecta are observed to rise on very different time-scales (from minutes or hours to days, e.g. Brocksopp et al. 2007; Tetarenko et al. 2017; Bright et al. 2020), depending on the size and energy of the ejected component. It is possible that the flare had a fast rise, and thus we missed the optically thick phase with our coverage. Another possibility is that the flare was entirely optically thin, as already observed for some sources (e.g. Fender et al. 1997). The ATCA observation taken on MJD 58524 and the decreasing flux observed on MJD 58523 (see Section 3.2) seem to confirm that we are observing the optically thin decay of the transient jet (see Fig. 10).","Citation Text":["Belloni et al. 2020"],"Citation Start End":[[1035,1054]]}
{"Identifier":"2017MNRAS.464..635M__Dekel_et_al._2009_Instance_1","Paragraph":"The basic idea, summarized in Dekel et al. (2009), is that during VDI, the high surface density of gas and \u2018cold\u2019 young stars, \u03a3, drives the Toomre Q parameter below unity, Q \u223c \u03c3\u03a9\/(\u03c0G\u03a3) \u2272 1, where \u03c3 is the 1D velocity dispersion and \u03a9 is the angular frequency, a proxy to the epicyclic frequency \u03ba, which is related to the potential well (Toomre 1964). It has been established that under such conditions, the disc will fragment and produce large star-forming clumps. This has been shown using idealized simulations of isolated galaxies (Noguchi 1999; Gammie 2001; Immeli et al. 2004a,b; Bournaud, Elmegreen & Elmegreen 2007; Elmegreen, Bournaud & Elmegreen 2008; Bournaud & Elmegreen 2009; Hopkins et al. 2012b), as well as cosmological simulations (Agertz, Teyssier & Moore 2009; Ceverino et al. 2010; Ceverino et al. 2012; Genel et al. 2012; Mandelker et al. 2014; Oklopcic et al. 2016). The ratio of clump mass to the mass of the cold disc scales as Mc\/Md \u221d \u03b42, where \u03b4 = Md\/Mtot is the ratio of the cold disc mass to the total mass within the disc radius, which includes the bulge and dark matter halo (e.g. Dekel et al. 2009). This leads to much larger clumps at z \u223c 2 than the low-redshift giant molecular clouds (GMCs). Gravitational interactions in the perturbed disc drive turbulence causing the disc to self-regulate in a marginally stable state with Q \u2272 1 (Dekel et al. 2009; Ceverino et al. 2010; Krumholz & Burkert 2010; Cacciato, Dekel & Genel 2012; Forbes, Krumholz & Burkert 2012; Forbes et al. 2014) that can last for more than a Gyr so long as the accretion is not interrupted. Some recent works have called into question the validity of linear Toomre analysis in the context of these highly non-linear galaxies (Behrendt, Burkert & Schartmann 2015; Tamburello et al. 2015; Inoue et al. 2016) and others have suggested alternate fragmentation mechanisms related to turbulence (e.g. Hopkins 2013). However, since clump formation is largely determined by the balance between self-gravity, turbulent pressure and the centrifugal force, the largest clumps are always roughly at the Toomre scale. Larger clumps would be disrupted due to the shear and\/or tidal forces within the disc, or would not collapse in the first place due to the centrifugal force. Therefore, regardless of the full validity of linear Toomre analysis, it is plausible that the Toomre Q parameter can serve as a crude criterion for instability, possibly with a critical value that is larger than unity.","Citation Text":["Dekel et al. (2009)"],"Citation Start End":[[30,49]]}
{"Identifier":"2015MNRAS.450.3458C__Cichowolski_et_al._2001_Instance_5","Paragraph":"The kinetic energy stored in the CO shell can be estimated as $E_{\\rm kin} = 0.5\\, M_{\\rm shell}\\, V^2_{\\rm exp}$, where Vexp is the expansion velocity of the shell and Mshell is the total (molecular, atomic, and ionized) shell mass. Adopting an expansion velocity equal to half the velocity interval where the structure is observed, Vexp = 7.0 \u00b1 1.3 km\u2009s\u2212 1 , the molecular mass given in Table 1 and the atomic and ionized masses estimated by Cichowolski et al. (2001), 1450 and 3000 M\u2299, respectively, we obtain Ekin = (2.5 \u00b1 1.0) \u00d7 1049 erg, assuming a 40\u2009per\u2009cent error for the masses.. Although Cichowolski et al. (2001) concluded that WR 130 could have alone created the observed structure, it is important to note that they did not take into account the molecular mass present in the shell, which considerably increases the kinetic shell energy. Thus, we can compare now the new value obtained for Ekin with the mechanical energy deposited in the ISM by the wind of the WR star, Ew = (0.7\u20132.2) \u00d7 1050 erg (Cichowolski et al. 2001). We obtain \u03f5 = Ekin\/Ew = 0.007\u20130.5. The ratio \u03f5 measures the energy conversion efficiency in the shell, and according to evolutionary models \u03f5 \u2264 0.2 (Koo & McKee 1992). Thus, not all the possible values of \u03f5 are compatible with the scenario where the energy injected during the WR phase is enough to create the structure. In this case, the contribution of the energy injected during the O-star phase and\/or other massive stars, should be considered. As mentioned in the Introduction, WR 130 is a WNH star, and according to Smith & Conti (2008) its age would be of about 2\u20133 Myr and its initial mass of at least 60 M\u2299. A rough estimation of the energy injected by such a star during its main sequence yields Ew = (2.5\u20133.5) \u00d7 1050 erg (de Jager, Nieuwenhuijzen & van der Hucht 1988), which would be enough to create the observed structure. We have nevertheless looked for the presence of other massive stars in the region. We queried the available catalogues such as the Galactic O-Star Catalog (Ma\u00edz Apell\u00e1niz et al. 2013), the Early-Type Emission-Line Stars Catalogue (Wackerling 1970), the Catalogue of Be stars (Jaschek & Egret 1982), the H-alpha Stars in the Northern Milky Way Catalogue (Kohoutek & Wehmeyer 1997), and the Catalog of Galactic OB Stars (Reed 2003), for early-type and emission stars. No stars were found in any catalogue. The only massive star located nearby is, as mentioned by Cichowolski et al. (2001), an OB star, which has an uncertain spectral type and no distance estimate (Stock, Nassau & Stephenson 1960). It is located in projection not in the centre of the structure but on to the shell (there is a second OB star mentioned by Cichowolski et al. 2001 but its location is actually outside the structure, see fig. 1 of Cichowolski et al. 2001). Although we cannot completely rule out the possibility that the OB star may be playing a role in creating the shell structure, we think that the action of WR 130 is sufficient and most likely dominant in the region.","Citation Text":["Cichowolski et al. (2001)"],"Citation Start End":[[2436,2461]]}
{"Identifier":"2018ApJ...862....8T___2005_Instance_1","Paragraph":"We introduce our target object, MC27\/L1521F (e.g., Mizuno et al. 1994; Onishi et al. 1996, 1998, 1999, 2002; Codella et al. 1997), which is one of the densest cores among nearby low-mass star-forming regions, and it contains a very low-luminosity (Lint  0.07 L\u2299) protostar (Bourke et al. 2006; Terebey et al. 2009). Earlier studies based on single-dish observations showed both the gas temperature derived from molecular line observations (e.g., Codella et al. 1997; Tatematsu et al. 2004) and the dust temperature obtained from multiband dust continuum observations (e.g., Kirk et al. 2007; Sadavoy et al. 2018) are as cold as \u223c10 K, indicating that most of the gas and dust contents at the line of site are in the cold environment. The high-deuterium fractionation estimated by the N2D+\/N2H+ ratio also supported the cold and evolved nature of the core (Crapsi et al. 2004, 2005). On the other hand, submillimeter single-dish observations detected CO (J = 6\u20135, 7\u20136) lines, which trace the warm (\u223c30\u201370 K) and dense (\u223c105 cm\u22123) component at the center of the core (Shinnaga et al. 2009). Although the filling factor of this warm gas relative to the entire core is quite small, these facts demonstrated that multiple temperature mediums are coexisting within the core. Tokuda et al. (2014) (hereafter, Paper I), and Tokuda et al. (2016; hereafter, Paper II) observed this core with ALMA and demonstrated the dynamical nature in this system. We found a few starless high-density cores, one of which has a very high density of \u223c107 cm\u22123 (MMS-2), within a region of several hundred astronomical units around a very low-luminosity protostar (MMS-1) detected by Spitzer (Bourke et al. 2006; Terebey et al. 2009). The molecular line observation showed several cores with arc-like structures, possibly due to the dynamical gas interaction. Similar arc-like structures have also been reproduced by hydrodynamical simulations both with and without a magnetic field (Matsumoto et al. 2015b, 2017), including the different mechanisms. These complex velocity\/spatial structures indicate that in MC27\/L1521F the turbulence may play an essential role in undergoing fragmentation in the central part of the cloud core, which is different from the classic scenarios of fragmentation in massive disks (Larson 1978; Boss 2002; Machida et al. 2008). More recently, Tokuda et al. (2017; hereafter, Paper III) found that the central very low-luminosity protostar has a dynamical mass of \u223c0.2 M\u2299 and the associated disk is extremely compact with the disk radius of \u223c10 au with an indication of detachment nature from the surrounding dense environment. Paper III mainly focused on gas\/dust distributions of the central protostar (MMS-1) seen in the 0.87 mm continuum and 12CO (J = 3\u20132) observations to determine its evolutionary stage. In this paper, we investigate the properties of the surrounding gas around the central protostar to better understand the dynamical nature of this system by using the high-resolution (\u223c02) 12CO (J = 3\u20132) observations as well as 13CO (J =2\u20131) and C18O (J = 2\u20131) observations with a moderate angular resolution, \u223c1\u2033.","Citation Text":["Crapsi et al.","2005"],"Citation Start End":[[856,869],[876,880]]}
{"Identifier":"2021ApJ...913...55H__Goldman_et_al._2017_Instance_1","Paragraph":"The short-plateau SNe 2006Y, 2006ai, and 2016egz most likely come from partially stripped massive progenitors,36\n\n36\nThe lack of nebular spectra for SNe 2006Y and 2006ai remains a caveat.\n but a remaining question is their exact formation channel. If it is single-star evolution as assumed in this work, the main theoretical uncertainties are RSG wind mass-loss rates and stellar rotation (e.g., Hirschi et al. 2004; Georgy 2012; Chieffi & Limongi 2013; Meynet et al. 2015; Renzo et al. 2017). We assume no rotation and tweak the wind efficiency by hand, but it is debatable whether such high mass-loss rates are physically plausible. Observationally, there is indeed a wide range of measured RSG wind mass-loss rates (e.g., de Jager et al. 1988; van Loon et al. 2005; Mauron & Josselin 2011; Goldman et al. 2017; Beasor et al. 2020). In addition, recent observational and theoretical studies on RSGs and SNe II indicate that RSG wind mass-loss rates may be independent from metallicity (Goldman et al. 2017; Chun et al. 2018; Guti\u00e9rrez et al. 2018). Thus, it could be possible that the short-plateau SNe 2006Y, 2006ai, and 2016egz originate from single-star evolution. However, it is unlikely the case if RSG mass-loss rates are metallicity dependent (as in the main-sequence O\/B stars; e.g., Vink et al. 2000, 2001; Mokiem et al. 2007), given the estimated subsolar host metallicities (Table 2). In such a case, interacting binary evolution is more plausible, as Eldridge et al. (2017, 2018) indeed show some interacting binary products also result in short-plateau SNe. It is also important to note that any mass-loss models need to reproduce the observed populations of not only SNe II but also RSGs. For example, Neugent et al. (2020) recently show that the luminosity function of RSGs can be used to constrain their mass-loss rates. Future statistical studies with both RSG and SN II populations at various metallicities are required to distinguish the formation channels of short-plateau SNe.","Citation Text":["Goldman et al. 2017"],"Citation Start End":[[793,812]]}
{"Identifier":"2021MNRAS.502.1933S__XVI_2014_Instance_1","Paragraph":"In order to explain clustering of galaxies and its evolution, a linear bias factor, a parameter representing the difference between galaxy distribution and underlying matter distribution, is often introduced. The linear bias factor b of sgzKs at a scale of 8 h\u22121 Mpc is expressed as follows:\n(15)$$\\begin{eqnarray}\r\nb = \\frac{\\sigma _{8, \\rm sgzK}}{\\sigma _{8}(z)}.\r\n\\end{eqnarray}$$The rms fluctuation of sgzK distribution $\\sigma _{8, \\rm sgzK}$ (Peebles 1980) and that of underlying matter distribution at a given redshift \u03c38(z) can be expressed as follows:\n(16)$$\\begin{eqnarray}\r\n\\sigma _{8, \\rm sgzK} = \\sqrt{J_{2}(\\gamma) \\biggl (\\frac{r_{0}}{8\\,h^{-1} ~\\rm Mpc}\\biggr)^{\\gamma } }\r\n\\end{eqnarray}$$and\n(17)$$\\begin{eqnarray}\r\n\\sigma _{8}(z) = \\sigma _{8}(0)D(z),\r\n\\end{eqnarray}$$where\n(18)$$\\begin{eqnarray}\r\nD(z) = g(z)\/[g(0)(1+z)],\r\n\\end{eqnarray}$$(19)$$\\begin{eqnarray}\r\ng(z) = \\frac{5}{2}\\Omega _{\\mathrm{ m}}\\biggl [\\Omega _{\\mathrm{ m}}^{4\/7}-\\Omega _{\\Lambda }+\\biggl (1+\\frac{\\Omega _{\\mathrm{ m}}}{2}\\biggr)\\biggl (1+\\frac{\\Omega _{\\Lambda }}{70}\\biggr)\\biggr ]^{-1},\r\n\\end{eqnarray}$$(20)$$\\begin{eqnarray}\r\n\\Omega _{\\mathrm{ m}} = \\frac{\\Omega _{\\mathrm{ m},0}(1+z)^3}{E^{2}(z)},~~~ \\Omega _{\\Lambda } = \\frac{\\Omega _{\\Lambda ,0}}{E^{2}(z)};\r\n\\end{eqnarray}$$the present rms mass fluctuation is \u03c38(0) = 0.83 (Planck Collaboration XVI 2014) and J2(\u03b3) = 72\/[2\u03b3(3 \u2212 \u03b3)(4 \u2212 \u03b3)(6 \u2212 \u03b3)]. Fig. 8 shows the bias factor of sgzKs along with those of other populations in the literature.2 The bias of sgzKs with Ks  21.5 is larger than those of SMGs with z  2, while it is similar to those of ULIRGs at comparable redshift, EROs, DOGs with z \u223c 1.2, and SMGs with z = 2\u20133. Galaxies and their host dark matter haloes can be connected through the correlation length or the bias factor under the assumption that each dark matter halo is occupied by a single galaxy. Therefore, in order to find present-day descendants of high-redshift galaxy populations, we can use a redshift evolution of the bias for dark matter haloes hosting high-redshift galaxies. For example, Hickox et al. (2012) and Wilkinson et al. (2017) investigated the evolution of either correlation length or bias for the haloes of SMGs to find an evolutionary link between high-redshift SMGs and local galaxies. Seo et al. (2019) also used the same method to find descendants of EROs. Thus, we investigated the evolution of bias for dark matter haloes hosting sgzKs with Ks  21.5. By the ellipsoidal collapse model, Mo & White (2002) estimated the bias for dark matter haloes with mass M at a given redshift as follows:\n(21)$$\\begin{eqnarray}\r\nb (M, z) = 1 + \\frac{1}{\\delta _{\\mathrm{ c}}}~\\biggl [\\nu ^{\\prime 2}+b\\nu ^{\\prime 2(1-c)}-\\frac{\\nu ^{\\prime 2c}\/\\sqrt{a}}{\\nu ^{\\prime 2c}+b(1-c)(1-c\/2)}\\biggr ], \\\\\r\n\\end{eqnarray}$$where\n(22)$$\\begin{eqnarray}\r\n\\nu ^{\\prime } = \\sqrt{a}\\nu ~~~ \\mathrm{with} ~~~ \\nu = \\frac{\\delta _{\\mathrm{ c}}}{\\sigma (R)D(z)},\r\n\\end{eqnarray}$$a = 0.707, b = 0.5, c = 0.6, the critical overdensity is \u03b4c = 1.686 (Sheth, Mo & Tormen 2001), and \u03c3(R) is the variance of density field fluctuation.3 To trace the bias of haloes over redshift, we calculated halo mass M at a given redshift using the formalism for mass growth rate of a halo. Fakhouri, Ma & Boylan-Kolchin (2010) suggested two kinds of mass growth rates, mean growth rate $\\langle \\dot{M}\\rangle _{\\mathrm{mean}}$ and median growth rate $\\langle \\dot{M}\\rangle _{\\mathrm{median}}$, as follows:\n(23)$$\\begin{eqnarray}\r\n\\langle \\dot{M}\\rangle _{\\mathrm{mean}} = 46.1~\\mathrm{M_{\\odot }\\,yr^{-1}}\\biggl (\\frac{M}{10^{12}~\\mathrm{M_{\\odot }}}\\biggr)^{1.1} \\,\\,\\,\\,(1+1.11z)E(z)\r\n\\end{eqnarray}$$and\n(24)$$\\begin{eqnarray}\r\n\\langle \\dot{M}\\rangle _{\\mathrm{median}} = 25.3~\\mathrm{M_{\\odot }\\,yr^{-1}}\\biggl (\\frac{M}{10^{12}~\\mathrm{M_{\\odot }}}\\biggr)^{1.1} \\,\\,\\,\\,(1+1.65z)E(z),\r\n\\end{eqnarray}$$where the former is larger than the latter. We estimated each mass growth rate separately, and took their average as the mass growth rate of the halo with M at a given z. In Fig. 8, dashed lines represent the evolution of bias for haloes, which are the results considering the mass growth of haloes. On the other hand, solid lines show the evolution of bias for haloes with fixed halo mass.","Citation Text":["Planck Collaboration XVI 2014"],"Citation Start End":[[1331,1360]]}
{"Identifier":"2016ApJ...821..107G__Schwadron_et_al._2011_Instance_1","Paragraph":"We repeated the plasma pressure calculation presented by Schwadron et al. (2011) and Fuselier et al. (2012) for the new ENA energy spectrum. The results for the downwind hemisphere and for the Voyager 1 region are summarized in Table 3. The measured intensity \n\n\n\n\n\n\nj\n\n\nENA\n\n\n\n\n of neutralized hydrogen at a given energy translates into a pressure of the parent ion population in the heliosheath times the integration length along the line of sight, \u0394P \u00d7 l, in the following way:\n3\n\n\n\n\n\u0394\nP\n\u00d7\nl\n=\n\n\n\n4\n\u03c0\n\n\n3\n\n\nn\n\n\nH\n\n\n\n\n\n\n\nm\n\n\nH\n\n\nv\n\n\n\n\n\nj\n\n\nENA\n\n\n(\nE\n)\n\n\n\u03c3\n(\nE\n)\n\n\n\n\u0394\nE\n\n\n\nc\n\n\nf\n\n\n\n\n\n\n4\n\n\n\n\n\n\nc\n\n\nf\n\n\n=\n\n\n\n\n\n(\nv\n+\n\n\nu\n\n\nR\n\n\n)\n\n\n2\n\n\n\n\n\n\nv\n\n\n4\n\n\n\n\n\n(\n\n\nv\n\n\n2\n\n\n+\n4\n\n\nu\n\n\nR\n\n\n2\n\n\n+\n2\n\n\nu\n\n\nR\n\n\nv\n)\n.\n\n\nIn Equation (3), \u0394E denotes the width of the respective energy bin; for the typical radial velocity of solar wind in the flanks and the downwind hemisphere of the inner heliosheath, we assumed uR = 140 km s\u22121 as measured by Voyager 2, whereas uR = 40 km s\u22121 for the heliosheath in the Voyager 1 direction (Schwadron et al. 2011; Gloeckler & Fisk 2015). For the density of neutral hydrogen in the inner heliosheath a constant nH = 0.1 cm\u22123 is assumed (Schwadron et al. 2011; Gloeckler & Fisk 2015). The charge-exchange cross section between protons and neutral hydrogen decreases from (4 to 2) \u00d7 10\u221215 cm\u22122 for 0.015 to 1.821 keV (Lindsay & Stebbings 2005). The integration length l for ENA production in the plasma is approximately the thickness of the inner heliosheath. The part of Equation (3) without the velocity factor cf can be interpreted as stationary pressure. The total pressure or dynamic pressure is the stationary pressure times this factor. Integrating over all energy bins in Table 3, we obtain the total plasma pressure times heliosheath thickness as P \u00d7 l = 304 pdyn cm\u22122 au for the downwind hemisphere and 66 pdyn cm\u22122 au for the Voyager 1 region (1 pdyn cm\u22122 au = 0.015 N m\u22121). If we want to put these numbers into the context of other studies, we face two problems. First, the uncertainty of the total pressure is large given the upper limits in the two lowest energy bins. Second, heliosheath plasma more energetic than 2 keV will produce ENAs that cannot be detected with IBEX-Lo. We therefore used the observed median j = 0 cm\u22122 sr\u22121 s\u22121 keV\u22121 for heliospheric ENAs in the two lowest energy bins of IBEX-Lo and relied on the study by Livadiotis et al. (2013). They compared the expected plasma pressure from a kappa distribution of protons with the plasma pressure derived from IBEX-Hi energy spectra: the energy range between 0.03 and 2 keV, roughly corresponding to the IBEX-Lo range, covered more than half of the total plasma pressure predicted from a kappa distribution. The authors found a total plasma pressure of P = 2.1 pdyn cm\u22122 for all sky directions except for the ENA Ribbon. Gloeckler & Fisk (2015) presented a multi-component plasma model for the heliosheath to explain Voyager and IBEX observations. At low energies they assumed the ENA energy spectra provided by Fuselier et al. (2012). They derived a total pressure of 2.5 pdyn cm\u22122 in all three plasma regions in the nose of the heliotail (Gloeckler & Fisk 2015). Pressure contributions from the slowed solar wind, magnetic pressure, and the pressure exerted from pickup ions and anomalous cosmic rays all had to be taken into account to obtain this total pressure.","Citation Text":["Schwadron et al. (2011)"],"Citation Start End":[[57,80]]}
{"Identifier":"2019ApJ...880..119C__Tomczak_et_al._2014_Instance_1","Paragraph":"Figure 9 demonstrates that the environment plays an important role in shaping the buildup of the passive population. The mild redshift dependence in the field ratios suggests that newly quenched galaxies are being added in both the bright end and the faint end of the population at similar fractions. Indeed, the number density of the bright and faint galaxies in the field have both grown by a factor of \u223c4 from the highest redshift bin (z \u223c 1.15) to the lowest redshift bin (z \u223c 0.15). Note that the split magnitude between the bright and faint galaxies (\n\n\n\n\n\n) roughly corresponds to \n\n\n\n\n\n.24\n\n24\nThis M\/L conversion is derived using a range of SSP models, assuming zf = 0.5\u22124.0 and Z = Z\u2299.\n The mild redshift dependence in the field ratios is therefore consistent with previous studies on the stellar mass function of passive galaxies in the field, which have seen significant growth in both the high-mass end and low-mass end from z \u223c 1 to z \u223c 0 (e.g., Muzzin et al. 2013b; Tomczak et al. 2014). The growth is commonly attributed to various mass-quenching processes internal to the galaxies, such as energetic feedback from supernovae and stellar winds for low-mass galaxies (e.g., Dekel & Silk 1986; Hopkins et al. 2014) and ejective feedback from AGNs for more massive galaxies (e.g., Bower et al. 2006; Terrazas et al. 2016). Given that the field ratios do not evolve strongly with redshift, it is evident that the strong redshift dependence of the cluster ratios is a result of the high-density environment and that preferentially low (or moderate) mass galaxies are quenched by the environment. Additionally, the difference between the cluster and field ratios suggests that the quenching effects induced by the environment are clearer at low redshifts. This is consistent with the mostly mass-independent environmental quenching scenario established at low redshift (e.g., Peng et al. 2010). Because the LF of star-forming galaxies is steeper at the faint end, a population of environmentally quenched galaxies would have a relatively high faint-to-luminous ratio. Recent studies of infalling galaxies in local groups and clusters have demonstrated that the ram pressure stripping of the cold gas in the galaxies when it passes through the ICM is the dominant mechanism (e.g., Boselli et al. 2016; Fillingham et al. 2016; Fossati et al. 2018), although strangulation may also play a significant role (for galaxies with \n\n\n\n\n\n; see, e.g., Fillingham et al. 2015; Peng et al. 2015).","Citation Text":["Tomczak et al. 2014"],"Citation Start End":[[981,1000]]}
{"Identifier":"2020MNRAS.493.2472G__Sugiyama_et_al._2017_Instance_1","Paragraph":"Specific examples of IR\/maser flare correlations include the 22-GHz H2O and 6.7-GHz methanol masers associated with the ucH\u2009ii region G025.65+1.05 (Stecklum et al. 2017; Sobolev et al. 2019). The H2O maser variability in this object may be correlated with IR K-band variability in a source approximately 1500 au from the maser position (Sobolev et al. 2019). To be precise, the H2O emission appears to exhibit an anticorrelation, with bright 22-GHz episodes corresponding to times of low K band (most observations centred on 2.20 \u03bcm wavelength) flux density. This is to be expected for typical 22-GHz masers in star-forming regions, where computational models, for example Gray et al. (2016), show that strong IR emission from dust tends to destroy the 22-GHz inversion, probably by inhibiting the loss of population from the lower maser level (Strelnitskii 1981; Sobolev & Gray 2012). The 6.7-GHz methanol masers in this source have been described as showing a moderate flux rise at the time of a 22-GHz H2O maser flare (Sugiyama et al. 2017). A radiative mechanism is not the only explanation for H2O flaring in G025.65+1.05. VLBI observations identify the flaring object with a single compact spot of maser emission (Burns et al. 2020). An image from spectral channels in the blue-shifted wing of this spot shows a double structure, leading to an interpretation as a possible line-of-sight overlap of two masing objects. Preliminary modelling results for spatial overlap (Gray 2018) indicate that this mechanism can produce flares of significantly higher variability index than is typical of radiative pumping mechanisms, and the overlap interpretation therefore becomes more likely if the source is at the further of the two possible distances (2.7 and 12.5 kpc) mentioned in Burns et al. (2020). Another possible example of this overlap mechanism is a very high brightness H2O flare from Orion KL (Shimoikura et al. 2005). The 22-GHz maser emission of G025.65+1.05 has been monitored by a combination of the Simeiz, Torun, and Hartebeesthoek telescopes since 2000 (Vol\u2019vach et al. 2019). Observations were approximately monthly for early data, but on an almost daily basis since a giant flare in 2017\u20132018. During the period 2000\u20132019, it has flared strongly (peak flux density >2000\u2009Jy) three times: the peak flux density of each flare shows an increasing trend, while the flare duration and inter-flare interval have decreased. The shortest, brightest, and most recent flare (2017\u20132018) showed a double-peaked light curve, with rising sections well approximated by exponentials with characteristic times of approximately 10\u2009d. The exponential shape, backed by evidence from polarization and spectral width, has led Vol\u2019vach et al. (2019) to conclude that the flaring maser is unsaturated. The continuum flux at 870\u2009\u03bcm has generally been rising towards G025.65+1.05 since 2008, but the data are very sparse compared to the maser observations, including an interval of 8 yr with no data (Vol\u2019vach et al. 2019).","Citation Text":["Sugiyama et al. 2017"],"Citation Start End":[[1022,1042]]}
{"Identifier":"2022MNRAS.515.1795B__Yang_et_al._2021_Instance_1","Paragraph":"The most widely adopted parametrization of the observed Universe is based on the so-called \u039b cold dark matter (\u039bCDM) model (Peebles 1984), relying on the existence of cold dark matter and dark energy (\u039b) associated with a cosmological constant (Carroll 2001) in a spatially flat geometry. Predictions from this model have been found to agree with most of the observational probes such as the cosmic microwave background (CMB; e.g. Planck Collaboration 2020), the baryon acoustic oscillations (BAO; e.g. Alam et al. 2021), and the present accelerated expansion of the Hubble flow, based on the distance modulus\u2013redshift relation (the so-called Hubble\u2013Lema\u00eetre, or simply Hubble diagram) of type Ia supernovae (SNe Ia; e.g. Riess et al. 1998; Perlmutter et al. 1999), where a dominant dynamical contribution, dubbed dark energy (DE) and related to the cosmological constant, should drive such an acceleration. However, the fundamental physical origin and the properties of DE are still unknown, as the interpretation of \u039b is plagued by a severe fine-tuning issue to obtain the right amount of DE observed today. Moreover, the data sets listed above do not fully fit the evolution of DE ranging from early to late epochs (Benetti et al. 2019; Yang et al. 2021) and do not fully rule out a spatially non-flat Universe (Park & Ratra 2019; Di Valentino, Melchiorri & Silk 2020, 2021; Handley 2021; Yang et al. 2021). The latter possibility has raised a remarkable debate about the importance of properly combining CMB data to infer significant statistical interpretations from the analysis (Efstathiou & Gratton 2020; Planck Collaboration 2020) and, by extension, the importance of combining data sets that do not reveal manifest tension (Gonzalez et al. 2021; Vagnozzi, Loeb & Moresco 2021). Deviations from the spatially flat \u039bCDM model would imply important theoretical and observational consequences and a change in our current understanding of cosmic evolution (e.g. Capozziello, Benetti & Spallicci 2020). Statistically significant deviations in this directions have already been found in cosmological analyses with high-redshift probes such as Gamma-Ray Bursts (see Dainotti, Cardone & Capozziello 2008; Dainotti et al. 2011b, 2013a,b, 2015, 2017, 2020a; Dainotti, Ostrowaki & Willingale 2011a; Dainotti et al. 2020b for the standardization of these sources as cosmological candles) and quasars (QSOs) combined with SNe Ia (Risaliti & Lusso 2019; Lusso et al. 2019, 2020; Bargiacchi et al. 2021). Such a joint analysis (SNe + QSO) makes use of the observed non-linear relation between the ultraviolet and the X-ray luminosity in QSOs (e.g. Steffen et al. 2006; Just et al. 2007; Lusso et al. 2010; Lusso & Risaliti 2016; Bisogni et al. 2021; Dainotti et al. 2022) to provide an independent measurement of their distance (see e.g. Risaliti & Lusso 2015, 2019; Lusso et al. 2020, for details). The methodology is complementary to the traditional resort to type Ia SNe to estimate the cosmological parameters, yet it extends the Hubble\u2013Lema\u00eetre diagram to a redshift range currently inaccessible to SNe ($\\mathit{ z}$ = 2.4\u20137.5). Within a model where an evolution of the DE equation of state (EoS) in form w($\\mathit{ z}$) = w0 + wa \u00d7 $\\mathit{ z}$\/(1 + $\\mathit{ z}$) is assumed, the data suggest that the DE parameter is increasing with time (Risaliti & Lusso 2019; Lusso et al. 2020). Therefore, it is compelling to further study extensions of the \u039bCDM model that could produce such behaviour of DE.","Citation Text":["Yang et al. 2021"],"Citation Start End":[[1240,1256]]}
{"Identifier":"2019AandA...628A.117B__Emsellem_et_al._2007_Instance_1","Paragraph":"With the introduction of the first integral-field spectrographs (IFS, see e.g. Bacon et al. 1995), it became feasible to perform such observations in a spatially resolved manner and for larger samples of galaxies. The SAURON survey (Bacon et al. 2001; de Zeeuw et al. 2002) was one of the first projects to make extensive use of this technology. Based on their representative sample of 72 nearby early-type galaxies, the project investigated stellar and gaseous kinematics (Emsellem et al. 2004; Sarzi et al. 2006; Falc\u00f3n-Barroso et al. 2006; Ganda et al. 2006) and stellar population properties (Peletier et al. 2007; Kuntschner et al. 2010), and distinguished between fast and slow rotating early-type galaxies (Emsellem et al. 2007). Subsequently, the ATLAS3D project (Cappellari et al. 2011) continued this endeavour by further investigating kinematic properties, for instance the global specific angular momentum (Emsellem et al. 2011), based on a volume complete sample of 260 early-type galaxies. The CALIFA survey (S\u00e1nchez et al. 2012) advanced these previous studies towards a morphologically unbiased sample of 667 galaxies. In addition, it provides a unique combination of a large spatial coverage of a few effective radii and high spatial sampling (see also Garc\u00eda-Benito et al. 2015). Other surveys observe even larger samples: while the SAMI survey (Croom et al. 2012; Bryant et al. 2015) includes \u223c3000 galaxies across different environments, MaNGA (Bundy et al. 2015) will contain approximately 10 000 nearby galaxies. Most recently, these projects are being complemented by IFS studies of local galaxies in unprecedented spatial resolution (e.g. MUSE, Bacon et al. 2010). For instance, the TIMER project (Gadotti et al. 2019) analyses the central structures of 24 barred local disc galaxies in order to study the formation histories of these structures and infer constraints on the related secular evolution processes. Fornax3D (Sarzi et al. 2018; Iodice et al. 2019) investigates mostly early-type galaxies in the Fornax cluster environment, spatially covering galaxies up to four effective radii. The PHANGS survey (Leroy et al., in prep.)1 aims to connect the physics of gas and star formation with the large-scale galactic structure by complementing IFS data from MUSE with interferometric data from ALMA.","Citation Text":["Emsellem et al. 2007"],"Citation Start End":[[714,734]]}
{"Identifier":"2018ApJ...860..165T__Lada_et_al._2010_Instance_1","Paragraph":"A large \n\n\n\n\n\n survey in nearby spiral galaxies and (U)LIRGs performed by Gao & Solomon (2004a, 2004b) revealed a tight linear correlation between the infrared (IR) and the HCN luminosities for normal star-forming galaxies and starbursts. This linearity seemingly extends down to the scale of Galactic massive cores in the Milky Way and holds over a total range of luminosity of about eight orders of magnitude (Wu et al. 2005). These results imply that the dense molecular gas (i.e., \n\n\n\n\n\n) as traced by the \n\n\n\n\n\n line, rather than the total molecular gas, is the direct fuel for star formation. High-resolution simulation of dense clouds found that the HCN luminosity is related to mass of dense gas of \u2273104 cm\u22123 (Onus et al. 2018). In addition, Spitzer studies of Galactic molecular clouds also show evidence that star formation is restricted to the dense cores of GMCs (e.g., Evans 2008; Lada et al. 2010). The critical density ncrit40\n\n40\nThe critical density of rotational level j is defined as \n\n\n\n\n\n, where \n\n\n\n\n\n is the Einstein coefficient for spontaneous emission, \n\n\n\n\n\n in units of s\u22121, and \n\n\n\n\n\n is the collision rate coefficient that depends on the gas temperature, in units of cm3 s\u22121. All critical densities in this work are calculated on the assumption of Tkin = 100 K and optically thin conditions. The critical densities will decrease if the lines are optically thick.\n of rotational transitions is proportional to \u03bc2\u03bd3 (for optically thin lines at frequency \u03bd; \u03bc is the dipole moment of the molecule); therefore, molecules with high dipole moment are expected to trace high-density molecular gas (e.g., \u03bcHCN \u223c 2.98 D, \n\n\n\n\n\n D, and \u03bcCS \u223c 1.96 D vs. \u03bcCO \u223c 0.11 D; see Sch\u00f6ier et al. 2005). Subsequently, a number of studies have explored the link between molecular lines of dense gas (e.g., HCN, HCO+, and CS) and IR luminosities in different populations of galaxies (e.g., Gao et al. 2007; Papadopoulos 2007; Baan et al. 2008; Liu & Gao 2010; Wu et al. 2010; Wang et al. 2011; Garc\u00eda-Burillo et al. 2012; Zhang et al. 2014; Chen et al. 2015; Usero et al. 2015; Liu et al. 2016).","Citation Text":["Lada et al. 2010"],"Citation Start End":[[894,910]]}
{"Identifier":"2022AandA...659A.124H__Levy_et_al._2019_Instance_1","Paragraph":"From the excitation maps as shown in Fig. 9, we determine the maximum distance of AGN-ionized regions (RENLR,\u2006max) with respect to the AGN location that we can cover with our depth and area. To suppress the impact of noise on misclassification close to the borders of the demarcation lines, we require that at least six surrounding unbinned spaxels share the same AGN-ionization classification or the classification is present in the spatially binned data with its intrinsically higher S\/N. We need to exclude the LINER and the intermediate BPT regions, because those are not necessarily associated with AGN photoionization and can originate from post-AGN stars (e.g., Binette et al. 1994; Singh et al. 2013), shocked gas from stellar winds (e.g., Ho et al. 2016; L\u00f3pez-Cob\u00e1 et al. 2019), or the diffuse ionized gas inbetween star forming regions (e.g., Lacerda et al. 2018; Levy et al. 2019). We define the maximum ENLR size, RENLR,\u2006max, as the maximum projected distance of a robust AGN region to the AGN location detectable within the IFU FoV at the observational depth. In general, the maximum ENLR extent can be biased due to surface brightness dimming and observational depth. This is not a concern for the large FoV of MUSE and the narrow redshift distribution of the CARS sample. Nevertheless, we also determine two additional ENLR radii, RENLR,\u200615 and RENLR,\u200616, corresponding to the ENLR size out to a surface brightness limit of \u03a315\u2004=\u200410\u221215\u2006erg\u2006s\u22121cm\u22122\u2006arcsec\u22122(1\u2005+\u2005z)\u22124 (e.g Liu et al. 2013b, 2014; Hainline et al. 2013, 2014), or \u03a316\u2004=\u200410\u221216\u2006erg\u2006s\u22121cm\u22122\u2006arcsec\u22122(1\u2005+\u2005z)\u22124 (Chen et al. 2019b), respectively. All these measurements are listed in Table 5. The ENLR size can become unresolved for the CARS targets, in particular, at high intrinsic ENLR surface brightnesses. For such nondetections we set the upper limits for the ENLR size to be the angular size of the binned spaxels. The error on the sizes are generally driven by the spatial resolution of a given observations, so that we assume the error on all sizes to be the FWHM of seeing for each observation as reported in Table 2.","Citation Text":["Levy et al. 2019"],"Citation Start End":[[875,891]]}
{"Identifier":"2021MNRAS.508.3446H__Holdom_1986_Instance_1","Paragraph":"The general form of the velocity-dependent cross-section is given by $\\bar{\\sigma }=\\sigma _0 (v\/c)^n$, where the index n depends on different physical DM processes and c is the velocity of light in space (in natural unit $\\bar{\\sigma }=\\sigma _0 v^n$). In the case of DM with magnetic and\/or electric dipole moment, n = +2, \u22122 are considered. n = 2, 1, 0, \u22121 are applicable for scattering in presence of Yukawa potential (Buckley & Fox 2010), n = \u22124 is attributed for millicharged DM (Holdom 1986; Chun, Park & Scopel 2011). In Ref. Dvorkin, Lin & Schutz (2021), the nature of the DM\u2013baryon cross-section is discussed for a wide mass range of DM. Similar investigations are also carried out in Ref. Nadler et al. (2019), Bhoonah et al. (2018), Kovetz et al. (2018), and Mack, Beacom & Bertone (2007). In this work, the DM\u2013baryon interaction cross-section ($\\bar{\\sigma }$) is parametrized as $\\bar{\\sigma }=\\sigma _0 v^{-4}$ (Mu\u00f1oz, Kovetz & Ali-Ha\u00efmoud 2015; Barkana 2018; Mukhopadhyay et al. 2021). The term \u03c30 is the DM scalar scattering cross-section with baryons (of the type $\\alpha _q \\bar{\\chi }\\chi \\bar{q}q$ for dark DM particle \u03c7 with coupling \u03b1q). It may be mentioned in some earlier works like Bhoonah et al. (2018) and Kovetz et al. (2018), where millicharged DM is considered. But here, we assume a particle DM candidate and adopt value of $\\sigma _0 \\sim 10^{-41}\\,\\rm {cm^2}$ consistent with the scalar cross-section bound obtained from ongoing direct DM search experiments [extrapolating the allowed region for 0.1\u2009GeV \u2264m\u03c7 \u2264 3\u2009GeV from recent experiments like Aprile et al. (2018), Akerib et al. (2017), and Cui et al. (2017)] in the mass range discussed in this work. Several recent investigations on EDGES 21-cm signal also suggest the similar velocity dependence (n = \u22124) of the cross-section like Mu\u00f1oz et al. (2015), Mahdawi & Farrar (2018), and Barkana (2018). Moreover, n = \u22124 is chosen in many DM-related cases, namely, hadronically interacting DM, millicharge DM, the baryon acoustic oscillations (BAO) signal, etc.","Citation Text":["Holdom 1986"],"Citation Start End":[[486,497]]}
{"Identifier":"2015MNRAS.449.3057Z__Haiman_2013_Instance_1","Paragraph":"The required stellar population mass of \u2273104\u2009M\u2299 represents a hard limit for the detection of direct star light from z \u2273 7 Pop III star clusters at magnification \u03bc \u223c 1000 in the foreseeable future. As discussed in Section 1, large numbers of Pop III stars can form in atomic cooling haloes that manage to remain chemically pristine, a scenario that likely requires a background Lyman\u2013Werner (LW) radiation field that suppresses molecular hydrogen (H2) formation and, consequently, Pop III star formation in minihaloes. This scenario has been explored predominantly in the context of supermassive star and direct collapse black hole (Bromm & Loeb 2003; Haiman 2013) formation in regions of the Universe with an LW radiation field high enough above the cosmic mean to entirely suppress H2 formation. It thought the onset of efficient H2 cooling is required for star cluster, rather than supermassive star, formation (Regan & Haehnelt 2009), requiring a UV background intensity below J21 \u223c 103 for H2 to form in sufficient abundance but above J21 \u223c 10 to suppress star formation in minihaloes (here, J21 denotes the radiation intensity at the Lyman-limit in units of 10\u221221\u2009erg\u2009s\u22121\u2009cm\u22121\u2009Hz\u22121\u2009sr\u22121). Explicitly focusing on Pop III star cluster formation, Safranek-Shrader et al. (2012) argued that total stellar masses of 104\u2009M\u2299 may be achievable given a low enough LW background intensity, though unaccounted for internal LW feedback may be significant in reducing the star formation efficiency. The inclusion of ionizing radiation further increases the propensity for Pop III galaxies with total stellar masses \u2273104\u2009M\u2299 (Johnson et al. 2014). Related studies (Regan & Haehnelt 2009; Shang, Bryan & Haiman 2010; Latif et al. 2014; Regan, Johansson & Wise 2014; Visbal, Haiman & Bryan 2014) have routinely found central gravitationally unstable clumps with masses in excess of 104\u22125\u2009M\u2299 and central accretion rates of 0.01\u20131 M\u2299 yr\u22121 to form in chemically pristine, high-redshift atomic cooling haloes. The results of these studies additionally suggest that fragmentation, if it does occur, happens on small scales (\u22721\u201310\u2009pc) favouring the formation of a compact Pop III stellar cluster. While the limit of \u2273104\u2009M\u2299 adopted in this work seems marginally achievable for the total stellar mass of a Pop III galaxy, additional cosmological simulations are needed that focus on the possibility and characteristics of a metal-free stellar cluster forming in an atomic cooling halo.","Citation Text":["Haiman 2013"],"Citation Start End":[[651,662]]}
{"Identifier":"2017ApJ...845...86E__Soler_&_Terradas_2015_Instance_1","Paragraph":"Among the suggestedmechanisms responsible for the strong damping of the coronal loop oscillations (e.g., Ruderman & Roberts 2002; Ofman 2005, 2009; Morton & Erd\u00e9lyi 2009) resonant absorption of the MHD waves, which was established first by Ionson (1978), is a strong candidate. Several works developed this theory (e.g., Davila 1987; Sakurai et al. 1991a, 1991b; Goossens et al. 1995; Goossens & Ruderman 1995; Erd\u00e9lyi 1997; Cally & Andries 2010). The necessary condition for the resonant absorption is a continuum of Alfv\u00e9n or slow frequency across the loop (Ionson 1978; Hollweg 1984, 1987; Davila 1987; Sakurai et al. 1991a). Resonant absorption occurs when the frequency of the global MHD mode matches at least with one of the frequencies of the background Alfv\u00e9n or slow continuum at a location called he resonance point. As a result, the energy of the global MHD mode transfers to the local Alfv\u00e9n modes in a layer around the resonance point, named the resonance layer (Lee & Roberts 1986; see also Goossens et al. 2013; Soler & Terradas 2015). In the absence of dissipation mechanisms, the amplitude of the oscillations diverges at the resonance point. Dissipation is important in the resonance layer, where the oscillations make large gradients. The background Alfv\u00e9n or slow continuum can be due to the variation of the plasma density (e.g., Davila 1987; Ofman et al. 1994; Ruderman & Roberts 2002; Terradas et al. 2006; Soler & Terradas 2015), twisted magnetic field (Ebrahimi & Karami 2016), or both of them together (Karami & Bahari 2010; Giagkiozis et al. 2016). There are a variety of theoretical works related to the damping of the coronal loop oscillations based on the theory of resonant absorption of MHD waves (e.g., Goossens et al. 2002, 2009; Ruderman & Roberts 2002; Van Doorsselaere et al. 2004; Andries et al. 2005; Terradas et al. 2006; Karami et al. 2009; Karami & Bahari 2010; Soler et al. 2013; Soler & Terradas 2015; Ebrahimi & Karami 2016; Giagkiozis et al. 2016; Jung Yu & Van Doorsselaere 2016). For a good review about the theory of resonant absorption, see also Goossens et al. (2011).","Citation Text":["Soler & Terradas 2015"],"Citation Start End":[[1027,1048]]}
{"Identifier":"2021MNRAS.503.4016B__Sullivan_et_al._2019_Instance_1","Paragraph":"In Fig. 20, we show two volumes containing filaments which are almost aligned to the sky plane, as an example to illustrate the implications of this effect on the rotation measurement of such objects. The top panels show the projected maps of density, RM and \u03b5 for the four Chronos simulations. In presence of a large degree of alignment between magnetic fields and filaments, we therefore expect the magnetic field to mostly lie in the sky plane as well, with a very small line-of-sight component, thereby reducing the observable |RM| towards the observer. Current instruments (e.g. VLA and LOFAR) are able to detect values of |RM| \u2273 5\u2009rad\/m2 (e.g. Bonafede et al. 2013; O\u2019Sullivan et al. 2019; Locatelli, Vazza & Dom\u00ednguez-Fern\u00e1ndez 2018). The Figure suggests that, on one hand, clusters easily meet this requirement, while filaments would only be marginally detected, even for the runs in which the magnetic field is stronger (baseline and Z), due to the large degree of $\\boldsymbol {B}$-filament alignment, which implies low values of \u03b5, as can be seen in the third column. Therefore, the small line-of-sight component yields only little |RM|, typically below the detection threshold of present instruments, especially for runs in which the magnetization is weak already (DYN5 and CSFBH2). On the bottom panels we give, for each of the two selected areas, the median value of \u03b5 as a function of the rotation measure (in absolute value). The highest values of |RM|, mostly associated to clusters, correspond to higher values of \u03b5, although never approaching \u03b5 \u2248 1; at the lower side of |RM|, corresponding to the areas populated by filaments, lower values of \u03b5 are found, as expected from our previous considerations. The simulations are in overall agreement, except for CSFBH2, where AGN and star formation feedback introduces additional effects: although the impact of a quasi-parallel $\\boldsymbol {B}$-filament configuration is noticeable for a larger sample of objects (e.g. see Fig. 18), the bursty and random occurrence of star forming\/AGN events may strongly affect the local magnetic field topology and cause the statistic over a small volume to deviate from the expected trend.","Citation Text":["O\u2019Sullivan et al. 2019"],"Citation Start End":[[672,694]]}
{"Identifier":"2020AandA...638A.140W__Tricco_et_al._2016_Instance_1","Paragraph":"Using the constrained hyperbolic divergence cleaning scheme with variable cleaning speed from Tricco et al. (2016), we can keep the divergence error low in all cases. The mean normalized divergence error, \u27e8\u03f5divB\u27e9=\u27e8h|\u2207\u2005\u22c5\u2005B|\/|B|\u27e9, is typically of order 10\u22125\u2005\u2212\u200510\u22123. In Fig. 17, we show the normalized divergence error maps for several test problems. Again we see that the divergence cleaning works extremely well here, the maximum error is generally around 10\u22122. Comparing to Hopkins (2016, their Fig. 4), we find that the errors are smaller than their MFM simulations with the Dedner et al. (2002) cleaning in general, with an exception for the outskirt of the advection loop (where the magnetic field is essentially zero and thus not important for the result). The improvement is probably due to the more advanced constrained cleaning method (Tricco et al. 2016). The normalized divergence error for the \u03bc\u2004=\u200410 cloud-collapse case at the jet launching time is shown in Fig. 18. Here, the divergence cleaning still performs very well in the disc, along the jets and for the majority of the regions where the outflow interacts with the ambient gas, especially when the divergence error is compared to the total gas pressure (right panel). The result is similar to the Dedner cleaning in Hopkins (2016), although our error is somewhat larger at the tip of the jets where the gas is shocked. However, we note that the comparison is not direct in this case as the jets may develop differently. Overall, the result from cleaning is still worse than the constrained transport or constrained gradient schemes (Hopkins 2016). For SPMHD, as shown in Tricco & Price (2012), divergence errors can be reduced to machine precision (or more practically to a certain tolerance value) using cleaning, with the help of a sub-cycling routine. However, local adjustments are required to determine the number of iterations for each particle to efficiently subcycle the cleaning in the simulation. This is because certain regions are more affected than others and because divergence is spread to nearby neighbours. Conceivably, if vector potentials (Stasyszyn & Elstner 2015) could work for a wider range of problems this could be an interesting avenue as well. However, the exploration of these methods in detail is beyond the scope of this work.","Citation Text":["Tricco et al. (2016)"],"Citation Start End":[[94,114]]}
{"Identifier":"2019MNRAS.486.1781R__Bonning_et_al._2012_Instance_2","Paragraph":"To check for any spectral variation in the optical\/IR bands, we looked for variation in the V \u2212 J band colour against the V-band brightness. This colour variation was analysed for the epochs A, B, D, and E. During epochs A and B, the source showed a \u2018redder when brighter\u2019 (RWB) behaviour. During epoch E, a bluer when brighter behaviour was observed. During epoch D, we observed a complex behaviour. Upto a V-band brightness of around 15\u2009mag, the source showed a \u2018bluer when brighter\u2019 behaviour, but for optical brightness fainter than 15.0\u2009mag, a \u2018redder when brighter\u2019 behaviour was observed. The colour\u2013magnitude diagrams for all the four epochs are shown in Fig. 10. The spectral variations shown by the source are thus complex. From studies on the optical\u2013IR colour\u2013magnitude diagram, it is known that FSRQs in general show an RWB trend, which is attributed to them having a luminous accretion disc (Gu et al. 2006; Bonning et al. 2012). The observed optical emission is a combination of thermal blue emission from the accretion disc and non-thermal red emission from the jet. As the source gets brighter, the non-thermal emission has a more dominant contribution to the total flux, giving rise to the RWB behaviour (Bonning et al. 2012). During epochs A and B, there is a trend of the object to become RWB, irrespective of its optical brightness. The optical flares dominated by synchrotron emission processes during A and B have corresponding \u03b3-ray flares that are produced by EC processes. However, during epochs D and E, the colour variations were found to depend on the optical brightness. During the epochs when this complex spectral behaviour was noticed, the source showed an optical\/IR flare with no or a weak corresponding flare in the \u03b3-ray band. The source showed a much larger amplitude of variability in the optical\/IR bands, while in the \u03b3-ray band it was either faint or below the detection limit of Fermi. This definitely points to some complex physical changes and could be due to a combination of changes in the bulk Lorentz factor, electron energy density, and magnetic field as seen from our SED modelling of the multiband data.","Citation Text":["Bonning et al. 2012"],"Citation Start End":[[1223,1242]]}
{"Identifier":"2018MNRAS.474.1277A__Gao_&_Solomon_2004_Instance_1","Paragraph":"\nDense gas-fraction and the star-forming fraction. As in Paper I we define the dense gas-fraction, Mfrac, as the fraction of gas having density, n \u2273 103\u2009cm\u22123 and being colder than \u223c50\u2009K. Although somewhat ad hoc, this density and temperature threshold is good enough to estimate the fraction of gas likely to end up in putative star-forming clumps (e.g. Ragan et al. 2016 and other references therein). Furthermore, we also use a higher density threshold, n \u2273 104\u2009cm\u22123, to trace the fraction of strongly self-gravitating gas in a pristine cloud. In typical star-forming clouds such dense pockets of gas are traced due to emission from molecules such as HCN, HNC and HCO+ (e.g. Gao & Solomon 2004, Usero et al. 2015 and Biegel et al. 2016). The dense gas fraction, Mfrac$\\equiv \\frac{M_{\\text{thresh}}(n \\gtrsim 10^{3} \\mathrm{cm}^{-3}, 10^{4} \\mathrm{cm}^{-3}; 50 \\mathrm{K})}{M_{\\text{gas}}}$, for each choice of bimpact is shown on the upper panel of Fig. 6. Irrespective of the choice of the impact parameter and the density threshold, n, the fraction, Mfrac, increases steadily with time, but at best cycled a little under \u223c 10\u2009per\u2009cent of the gas into the dense-phase at the time calculations were terminated. Indeed, this fraction is significantly smaller in the realization where the inflows merely grazed past each other. While the fraction of gas cycled into the dense phase could possibly increase as the post-collision slab continues to accrete gas from the colliding flows, it is unlikely to exhibit a dramatic increase. This is because the NTSI destabilises the shocked slab and causes fluid layers to mix, thereby rupturing the dense pockets of gas. In any case, it is evident that grazing inflows are most inefficient in cycling gas into the dense phase. The fraction of gas cycled to densities higher than \u223c104\u2009cm\u22123 is typically an order of magnitude smaller. Evidently, star formation in a pristine cloud must be quite inefficient. We also note, irrespective of the choice of the density threshold and the impact parameter, the characteristics of the dense gas fraction look mutually similar. This is because whatever the choice of the pre-collision impact parameter, the fraction of gas cycled into the dense phase is limited by shearing interaction between layers of gas in the shocked slab. The source of shear, however, is different in the extreme scenarios of near head-on collision and one in which the gas flows merely graze past each other. In the former, shear is induced by the growth of the NTSI while in the latter, it is the obliqueness of the collision that induces a strong shearing interaction between layers of gas.","Citation Text":["Gao & Solomon 2004"],"Citation Start End":[[677,695]]}
{"Identifier":"2019AandA...627A.141W__\u00d6zel_et_al._(2016b)_Instance_1","Paragraph":"The radius estimate of the X-ray emitting area is small compared to the size of a neutron star, which typically has a radius of 10\u201314 km. The emission is therefore likely to be coming from the polar caps of PSR J1909\u20133744, as is expected from such old objects where the surface of the neutron star has cooled so that it is no longer detectable in the X-ray band. The emitting radius is somewhat smaller than the classical radius of a polar cap, \n\n\n\n\nR\npc\n\n=\n\n\n\n2\n\u03c0\nR\n\n\ncP\n\n\n\nR\n\n\n$ R_{\\mathrm{pc}}= \\sqrt{\\frac{2 \\pi R}{c P}} R $\n\n\n (e.g. Dermer & Sturner 1994), where R is the neutron star radius, c the speed of light in a vacuum and P the rotation period of the neutron star. The latter is \u223c2.9 km if we suppose a 10.6 km radius for the PSR J1909\u20133744 neutron star, the average of the best radius values for a neutron star of M\u2004=\u20041.5\u2006M\u2299 as determined by \u00d6zel et al. (2016b) who undertook a comprehensive study of 12 neutron stars. Considering the range of typical radii of neutron stars (10\u201314 km), the classical polar cap radius is between 2.7 and 4.4 km. Bogdanov et al. (2006) state that the radius determined through fitting the pulsar spectrum can be smaller than expected when the spectrum has been fitted with a single-temperature model, but that a higher signal-to-noise spectrum would require a two-temperature model, as for PSR J0030+0451, PSR J2124\u20133358 (e.g. Bogdanov et al. 2008), and PSR J0437\u20134715 (Bogdanov 2013). Fitting the spectrum with two black bodies does not improve the fit, again probably due to the low signal-to-noise ratio. Further, it is known that the radii of neutron stars are underestimated when using a black-body model as opposed to a realistic neutron star atmosphere model (e.g. Heinke et al. 2006). However, assuming the classical radius of the polar cap for this neutron star (2.9 km) in the nsatmos model provides a poor fit to the data with a C-statistic = 171.6, 33 degrees of freedom, suggesting that such a radius is too large. In fact, the nsatmos model produces a polar cap with size \u223c100\u2013400 m, slightly larger than the \u223c50\u2013300 m radius obtained with the black-body model (see Fig. 4).","Citation Text":["\u00d6zel et al. (2016b)"],"Citation Start End":[[856,875]]}
{"Identifier":"2018AandA...620A..31M__Schilke_et_al._1997_Instance_1","Paragraph":"Figure 3 (left) shows the moment zero map of the SiO (5\u20134, Eu = 31.26 K) emission integrated over the central ~3 km s\u22121 of the line covering the VLSR, between 21.4 and 24.0 km s\u22121. In contrast to Cesaroni et al. (2017), here we integrate the SiO emission over a narrower velocity range to highlight the elongated structure at near VLSR velocities. The elongated structure of the emission is clearly visible and perpendicular to the large-scale outflow (Fig. 2). Fitting the integrated SiO emission with a 2D Gaussian in the image plane indicates a PA of ~ 47 \u00b1 7\u00b0 and a major axis of ~0.28\u2033 (~600 au). The deconvolved PA is 30 \u00b1 10\u00b0 consistent with the PA of the dust continuum emission. The SiO emission is broad in velocity, ranging from \u22123.0 to 45.7 km s\u22121 considering emission above 3 \u03c3, this is broader than presented in Cesaroni et al. (2017), 4\u201339.1 km s\u22121, owing to our improved images after self calibration. In most massive YSOs, SiO is very spatially extended, typically tracing emission at shock fronts (e.g. Schilke et al. 1997) created by active jets and outflows driven by the energetic central sources (e.g. Cabrit et al. 2007; S\u00e1nchez-Monge et al. 2013b; Duarte-Cabral et al. 2014; Klaassen et al. 2015; Cunningham et al. 2016; Cesaroni et al. 2017; Beltr\u00e1n et al. 2018; Moscadelli et al. 2018). In at least one massive YSO, Orion source I, the SiO emission is rather compact, associated with SiO masers and is shown to trace a rotating disc and disc wind structure (Goddi et al. 2009; Ginsburg et al. 2018). The SiO emission from G17.64 is also relatively compact, unlike a jet. There is also compact emission from other silicon and sulphur bearing species in our observations, emission from SiS (v = 0 12\u201311, Eu = 69.95 K) and 33SO (65 \u201354, Eu = 34.67 K) are centrally peaked at the location of the continuum and SiO emission. The strong detection of SiO points to a reservoir of silicon in the gas phase, which, in addition to the detected sulphur-bearing species, suggest an association with ionised, hot (>100 K), turbulent and shocked gas (e.g. Minh et al. 2010; Minh 2016) where these species are released from the grain material.","Citation Text":["Schilke et al. 1997"],"Citation Start End":[[1021,1040]]}
{"Identifier":"2017AandA...602A...1S__Smol\u010di\u0107_&_Riechers_2011_Instance_1","Paragraph":"In recent decades, radio interferometers, such as the Karl G. Jansky Very Large Array (VLA), Australia Telescope Compact Array (ATCA), and Giant Meterwave Radio Telescope (GMRT), have surveyed fields of different sizes (ranging from tens of square arcminutes to thousands of square degrees), depths (microjansky to jansky), and multiwavelength coverage (e.g., Becker et al. 1995; Condon et al. 1998, 2003, 2012; Ciliegi et al. 1999; Georgakakis et al. 1999; Bock et al. 1999; Prandoni et al. 2001; Hopkins et al. 2003; Schinnerer et al. 2004; Bondi et al. 2003, 2007; Norris et al. 2005; Schinnerer et al. 2007, 2010; Afonso et al. 2005; Tasse et al. 2007; Smol\u010di\u0107 et al. 2008, 2014; Owen & Morrison 2008; Miller et al. 2008, 2013; Owen et al. 2009; Hales et al. 2014). These past surveys have shown that deep observations at high angular resolution (\u22721\u2033) with exquisite panchromatic coverage are critical to comprehensively study the radio properties of the main galaxy populations, avoiding cosmic variance with large area coverage (e.g., Padovani et al. 2009; Padovani 2011; Smol\u010di\u0107 et al. 2008, 2009b,a; Smol\u010di\u0107 2009; Smol\u010di\u0107 & Riechers 2011; Seymour et al. 2008; Bonzini et al. 2012, 2013). In this context, large area surveys down to unprecedented depths are planned with new and upgraded facilities (e.g., VLA, Westerbork, Australian Square Kilometre Array Pathfinder, MeerKAT, and Square Kilometre Array; e.g., Jarvis 2012; Norris et al. 2011, 2013, 2015; Prandoni & Seymour 2015). Figure 1 shows the 1\u03c3 sensitivity of each survey as a function of the area covered for past, current, and future radio continuum surveys. The VLA-Cosmic Evolution Survey (COSMOS) 3 GHz Large Project bridges the gap between past and future radio continuum surveys by covering an area as large as two square degrees down to a sensitivity reached to date only for single pointing observations. This allows for individual detections of >10 000 radio sources, further building on the already extensive radio coverage of the COSMOS field at 1.4 GHz VLA (VLA-COSMOS Large, Deep and Joint projects; Schinnerer et al. 2004, 2007, 2010), 320 MHz VLA (Smol\u010di\u0107 et al. 2014), 325 MHz and 616 MHz GMRT data (Tisani\u0107 et al., in prep.), 6 GHz VLA (Myers et al., in prep.), and the deep multiwavelength X-ray to mm photometry (Scoville et al. 2007; Koekemoer et al. 2007; Hasinger et al. 2007; Capak et al. 2007; Sanders et al. 2007; Bertoldi et al. 2007; Elvis et al. 2009; Ilbert et al. 2013; McCracken et al. 2012; Scott et al. 2008; Aretxaga et al. 2011; Smol\u010di\u0107 et al. 2012; Miettinen et al. 2015; Civano et al. 2016; Laigle et al. 2016, Capak et al., in prep.) and more than 97\u2009000 optical spectroscopic redshifts (Salvato et al., in prep.; zCOSMOS, Lilly et al. 2007, 2009; Trump et al. 2007; Prescott et al. 2006; Le F\u00e8vre et al. 2015; Aihara et al. 2011; Nagao et al., priv. comm.). This further makes the survey part of one of the richest multiwavelength data sets currently available. ","Citation Text":["Smol\u010di\u0107 & Riechers 2011"],"Citation Start End":[[1122,1145]]}
{"Identifier":"2020ApJ...893...52H__Kaastra_et_al._2008_Instance_1","Paragraph":"In this paper, our primary interest is in determination of plasma temperatures, in particular the characterization and measurement of high-temperature lines and continuum in the Chandra\/HETGS spectrum of \u03b6 Pup. If X-ray emission originates solely from embedded-wind shocks (Lucy & White 1980; Feldmeier et al. 1995), then the maximum temperature is determined by the highest relative velocities of the colliding structures, and the amount of emission is indicative of the emitting volume. Hot thermal plasmas emit strongly in bound\u2013bound lines of highly ionized states, in thermal Bremsstrahlung radiation, and in bound-free continuum emission. Bremsstrahlung and bound-free emission drop exponentially for photon energies larger than about kT of the plasma (in which k is Boltzmann\u2019s constant and T is the plasma temperature). For some plasma temperatures, bound-free emission can exceed that from thermal Bremsstrahlung at high energies (Landi 2007; Kaastra et al. 2008). However, the continuum shape is similar to that from Bremsstrahlung. The drop in the spectrum at high energies (short wavelengths) is sensitive to the hottest temperatures present in a multithermal plasma. Thermal emission lines are also very sensitive to temperature, having temperatures of peak emissivities that increase with atomic number; they typically emit most strongly over a range in temperature of about 0.3 dex (defined by the full width at half maximum of their emissivity curves). Hence, we concentrate here on the 2\u20139 \u212b region due to the relative sparseness of emission lines, the high-temperature sensitivity, and the reduced importance of wind absorption. The temperature sensitivity occurs because for the expected plasma temperatures from embedded-wind shocks (2\u201310  MK), the spectrum will drop accordingly somewhere in the 1\u201310  \u212b range. In this regime, the wind absorption is relatively unimportant since the continuum opacity (due mainly to K-shell absorption by metals) drops very steeply to short wavelengths, so we do not have to be concerned with the wind density structure to interpret the spectral shape.","Citation Text":["Kaastra et al. 2008"],"Citation Start End":[[952,971]]}
{"Identifier":"2019ApJ...883..189S__Cecil_et_al._2001_Instance_1","Paragraph":"NGC 3079 is a well-known nearby (z = 0.003723) Seyfert 2\/LINER galaxy with a nuclear starburst, which is a member of an interacting galaxy pair. It is an edge-on spiral galaxy with several regions of star formation along the disk. Its eight-shaped radio structure with a ring-like feature (ring hereafter) inside its northeast lobe (Duric & Seaquist 1988; Irwin & Saikia 2003; Wiegert et al. 2015; Irwin et al. 2019), its slow-moving (v \u223c 0.1c) parsec-scale radio jet (Irwin & Seaquist 1988; Baan & Irwin 1995; Middelberg et al. 2005), and copious amount of hot X-ray gas and emission-line filaments (Cecil et al. 2001, 2002) have been well documented. NGC 3079 is a Compton-thick AGN (Iyomoto et al. 2001) having a disk-like pseudo-bulge. The origin of the double-lobed morphology observed in NGC 3079 is not yet completely understood. Duric & Seaquist (1988) argue that the lobe material is moving at supersonic velocities, though they dismiss the possibility that these are jets that may be precessing based on the difficulty in reproducing the closed-loop morphology using these models. They suggest that these are probably winds originating from a nuclear starburst or accretion activity. X-ray emission and emission-line imaging shows that NGC 3079 hosts a superwind (Veilleux et al. 1994; Cecil et al. 2001, 2002). Superwinds are believed to be generated when the kinetic energy from supernovae and massive stellar winds are transformed into thermal energy. The heated gas then expands into the lower pressure ambient medium. The direction of the steepest pressure gradient is along the minor axis. The hot gas is easily traced by the X-ray emission, which is a more direct tracer compared to the emission-line gas, which is generated at shock fronts as a result of the interaction with the ambient medium. The emission-line ratios seen in superwinds are typical of shock-ionized gas as also seen in the case of NGC 3079 (Veilleux et al. 1994). Although these authors favored a starburst-wind origin for the radio lobes, they were not able to rule out AGN-driven winds. Several authors favored the jet origin of the lobes (Irwin & Seaquist 1988; Irwin & Saikia 2003). Irwin et al. (2017) find that NGC 3079 is one of the few edge-on galaxies in the CHANG-ES sample (Irwin et al. 2012) that shows radio lobes that stand out in polarized intensity compared to the galactic disk emission that is not as highly polarized.","Citation Text":["Cecil et al. 2001","Cecil et al. 2001"],"Citation Start End":[[601,618],[1296,1313]]}
{"Identifier":"2022ApJ...932....7Y__Marino_et_al._2021_Instance_1","Paragraph":"In summary, the different components show similar evolutionary trends for the fractional rms and characteristic frequency with the hardness ratio in the LE\/ME\/HE energy bands. The evolutionary trend can be explained under the truncated disk\/corona model (Esin et al. 1997; Done et al. 2007). Actually, there is still debate on the truncation of the accretion disk in MAXI J1820+070. Buisson et al. (2019) discovered a steady inner accretion disk measured by relativistic reflection. Kara et al. (2019) found that the reverberation time lags between the continuum-emitting corona and the irradiated accretion disk are much shorter than previously seen in truncated accretion disks, and the timescale of the reverberation lags is shortened by an order of magnitude over a period of weeks, whereas the shape of the Fe K\u03b1 emission line remains remarkably constant. Similar results are also obtained from spectral analysis. Meanwhile, there are some other studies that support the truncated accretion disk argument by either spectral analysis or timing analysis (De Marco et al. 2021; Marino et al. 2021; Zdziarski et al. 2021a, 2021b). De Marco et al. (2021) found that the frequency of the thermal reverberation lags increases steadily, and, on the other hand, the temperature of the quasi-thermal component grows as the source softens, which can be explained in terms of a decrease in the disk inner radius. Moreover, De Marco et al. (2021) measured that the values of the lag amplitude are a factor of 3 longer than those reported in Kara et al. (2019). The longer lags might not be easily reconciled with the conclusion of a disk extending close to the innermost stable circular orbit. Zdziarski et al. (2021a, 2021b) confirmed the existence of an optically thick disk of at least >10R\n\ng\n from joint spectral analysis. To sum up, so far, all arguments in favor of the nontruncated disk model can be reasonably explained under the truncated disk model. Except for the methods mentioned above, in the present paper, we confirm the truncated accretion disk model in MAXI J1820+070 from the evolutionary trend of broadband noise components with a new perspective by means of the correspondence relation between break frequency and radiation region radius. Quantitatively, we can calculate the radiation region at a different energy band of L\n2, which represents the variable emission from the outermost region (see Figure 13). We take a parameter set for a standard \u03b1-disk (\u03b1 = 0.1, M\nBH = 10M\n\u2299, scale height H\/r = 0.1) to calculate the viscous frequency at a certain radius (Kato et al. 2008). As Figure 13 shows, the characteristic frequency of the L\n2 component shows an energy dependence: the emitting region spans from \u223c34R\n\ng\n to \u223c27R\n\ng\n corresponding to 1\u2013150 keV photon energy. Like Dzie\u0142ak et al. (2021) and Kawamura et al. (2022) found using frequency-resolved spectral analysis that the L\n2 component is supposed to come from variable disk emission. However, making use of the ME and HE data from Insight-HXMT, we actually detect high-energy emission >100 keV from the L\n2 component, which cannot be attributed to a simple standard accretion disk. Considering the change of radius plotted in Figure 13, even though we can attribute the high-energy emission to the propagation of fluctuation from the disk to the hot flow, we should also expect a constant characteristic frequency for L\n2 at the high-energy band according to fluctuation propagation. The characteristic frequency of L\n2 remains unchanged below 20\u201330 keV as Kawamura et al. (2022) found the constant peak frequency for P1. However, when the photon energy is greater than 20\u201330 keV, the radiation radius of L\n2 starts to decrease to 27R\n\ng\n. This phenomenon cannot be easily interpreted by fluctuation from a disk propagating to the hot flow. Therefore, we speculate that L\n2 may originate in a warm extended variable disk region. We therefore should consider a more complicated accretion flow geometry where a standard accretion disk transits to a hot advection-dominated accretion flow geometry.","Citation Text":["Marino et al. 2021"],"Citation Start End":[[1080,1098]]}
{"Identifier":"2022MNRAS.517.4813G___2016a_Instance_1","Paragraph":"The community searches for extensions to the \u039bCDM model to resolve some of the mentioned problems. The approaches to face them are divided into two main branches: (i) assume a DE fluid with the capability to accelerate the Universe or (ii) modify General Relativity (GR) to obtain the cosmic acceleration without DE (Motta et al. 2021). This paper will focus on the second point under the formalism known as fractional calculus, which consists of a generalization of the classical integer order calculus, whose derivatives and integrals are of (real or complex) arbitrary order. These fractional operators are not local. In many cases can model real-world phenomena in a better fashion than those obtained by classical calculation. For example, it is coined fractional dynamics as a field of study in physics and mechanics investigating the behaviour of objects and systems that are characterized by power-law non-locality, power-law long-term memory or fractal properties by us and differentiation of non-integer orders, i.e. by methods in the fractional calculus (see the review by Tarasov 2013). Fractional calculus is a field with multiple applications and a great deal of research activity. Fractional quantum mechanics is employed as a tool within quantum field theory and gravity for fractional spacetime (Calcagni 2010a, b) and the fractional quantum field theory at positive temperature (Lim 2006; Lim & Eab 2019) and other applications of quantum cosmology (Moniz & Jalalzadeh 2020a, b; Jalalzadeh, da Silva & Moniz 2021; Rasouli, Jalalzadeh & Moniz 2021). Recently, the community explores the fractional calculus to tackle problems associated in cosmology (Shchigolev 2011, 2013a, b, 2016, 2021; Calcagni 2013, 2017a, 2021a, b; Calcagni & De Felice 2020; Jalalzadeh, Costa & Moniz 2022), stochastic GW background (Calcagni & Kuroyanagi 2021), luminosity distance (Calcagni et al. 2019), inflation and CMB spectrum (Calcagni, Kuroyanagi & Tsujikawa 2016; Calcagni 2017b), Fractional Action Cosmology (El-Nabulsi 2012, 2016a; Jamil, Momeni & Rashid 2012), fractional geodesic equation, complex GR, and discrete gravity El-Nabulsi (2013b), minimal couplings (El-Nabulsi 2013a), phantom (Rami 2015), Ornstein\u2013Uhlenbeck-like fractional differential equation in cosmology (El-Nabulsi 2016b), a variable Order Parameter (El-Nabulsi 2017a), wormholes in fractional action cosmology (El-Nabulsi 2017c). New metrics were considered (El-Nabulsi 2017b), as well as some DE models in emergent, logamediate, and intermediate scenarios of the universe (Debnath, Jamil & Chattopadhyay 2012; Debnath, Chattopadhyay & Jamil 2013). For instance, Shchigolev (2016, 2021) found \u03b1 = 0.926 (where \u03b1 is the order of the Riemann\u2013Liouville fractional integral). In Shchigolev (2011,2013a, b) were obtained several exact solutions for cosmological models, which differ significantly from the standard model due to the fractal nature of spacetime (Calcagni 2010a, b). Jalalzadeh et al. (2022) explore the interval 1  \u03b1  2 using Riesz\u2019s fractional derivative (that is not related to the index of Riemann\u2013Liouville fractional integral) to obtain the non-boundary and tunnelling wave functions for a closed de Sitter geometry. Another example, Barrientos, Mendoza & Padilla (2021), studies the Universe dynamics without DM and DE components by modifying the mathematical structure of Friedmann equations with fractional calculus. Another approach is calculating the value of the \u039b (due to the well-known ultraviolet divergence in the standard quantum field theory), which needs to restructure the theory using the fractional calculus (Calcagni 2021c). Giusti (2020) and Torres et al. (2020) explore Modified Newtonian Dynamics Theories and quantum cosmology in this fractional approach (Barrientos et al. 2021). Finally, notice that there are several definitions of fractional derivatives and fractional integrals, such as those of Riemann\u2013Liouville, Caputo, Riesz, Hadamard, Marchand, and Griinwald\u2013Letnikov, among other more recent ones (see Kilbas, Srivastava & Trujillo 2006; Podlubny 1998, and references therein). Even though these operators are already well studied, some of the usual features related to function differentiation fails, such as Leibniz\u2019s rule, the chain rule, and the semigroup property (Podlubny 1998; Kilbas et al. 2006).","Citation Text":["El-Nabulsi","2016a"],"Citation Start End":[[2011,2021],[2028,2033]]}
{"Identifier":"2022ApJ...925...63G__Heckman_et_al._2011_Instance_1","Paragraph":"There are four possible examples of where Ly\u03b1 absorption is taking place. In Figures 3(b) and (c) we see that the high-velocity component of the double profile seen in both H\u03b2 and [C ii] is significantly suppressed compared with the low-velocity component. For Figure 2(b) in particular, the peak in the H\u03b2 profile falls close to a strong dip in the Ly\u03b1 profile, which appears as two small wings on either side of the dip. For Figure 3(c), the feature seen in both [C ii] and H\u03b2 is largely suppressed above 6500 km s\u22121. Also, in Figures 2(d) and 3(d), the Ly\u03b1 profile is shifted blueward of the main [C ii] and H\u03b2 peaks, suggestive of absorbing gas centered at V\nhelio = 6500\u20136700 km s\u22121. Asymmetric Ly\u03b1 profiles like this are often associated with radial outflows in galaxies (e.g., Heckman et al. 2011). The COS pointing (HCG92-5) samples the gas in the so-called \u201cAGN bridge\u201d (Cluver et al. 2010), a linear H2 filament that is apparently separate from the main collisional shock in SQ. A strong shear in the velocity field of the [C ii] emission was noted by Appleton et al. (2013) in that region. The spectrum of the CO (1\u20130) emission from that direction (see Figures 2(d) and 3(d)) shows significant CO emission at the high-velocity side of the profile, which would be consistent with absorption. Finally, HCG92-7 shows, in Figure 3(e), the opposite effect. In this case, again considering the high-velocity component of the double-horned profile only, we see that the Ly\u03b1 is significantly redshifted with respect to the H\u03b2 emission, with a sharp drop in emission as one approaches the peak of the H\u03b2 (around V\nhelio = 6600\u20136700 km s\u22121). This may be another example of asymmetric absorption, with resonant scattering to the red side of the wings of the kinematics. In summary, we see that two regions in the main emission-line filament are almost free of absorption, whereas other regions show strong absorption. Even in those cases, at least some of the Ly\u03b1 emission is able to resonantly scatter and undergo many scatterings within the gas before eventually escaping into the wings of the velocity profile, where the optical depth is much lower. In Section 5.2, we estimate the escape fraction of the Ly\u03b1 photons and the number of scatterings, and we discuss the multiphase structure of the intragroup medium gas. We also remind that the differences between the Ly\u03b1 and H\u03b2 spectral profiles may not only be due to Ly\u03b1 scattering, because collisional excitation could also contribute.","Citation Text":["Heckman et al. 2011"],"Citation Start End":[[784,803]]}
{"Identifier":"2022ApJ...934..100S__Eracleous_&_Halpern_1994_Instance_1","Paragraph":"\nAccretion disks around compact objects. Compact low-luminous massive objects are expected to lurk in the CGM of galaxies. From lone low-mass stars (e.g., Helmi 2020), to black holes (BHs) arising from stellar evolution (e.g., Fender et al. 2013), including remnants of Population III stars (e.g., Madau & Rees 2001; Filho & S\u00e1nchez Almeida 2018), and primordial BHs (e.g., Carr & Hawking 1974; Clesse & Garc\u00eda-Bellido 2015). The number density of these objects is assumed to be high; for example, there are claims that primordial BHs account for all the dark matter in the universe (Clesse & Garc\u00eda-Bellido 2015). Even if this is not the case (e.g., Carr & K\u00fchnel 2020), the expectations clearly overwhelm the density required to account for all the emission signals that we observe, of the order of one clump per central galaxy (Section 4.1). When these compact objects are surrounded by accretion disks, they should emit in H\u03b1, with the rotation of the disk giving rise to two-horn H\u03b1 profiles when observed with the appropriate viewing angle (Smak 1969). This kind of double-peak H\u03b1 emission is observed in X-ray binaries (Grundstrom et al. 2007; Zamanov et al. 2013; Casares & Torres 2018; Monageng et al. 2020) and cataclysmic variables (e.g., Zolotukhin & Chilingarian 2011), with peak separations of up to hundreds of km s\u22121 (Zamanov et al. 2013). At a completely different mass scale, double-peak H\u03b1 emission is sometimes observed in the broad-line region of AGN, where it is also supposed to trace a rotating disk (e.g., Eracleous & Halpern 1994; Ho 2008). The double-peak H\u03b1 emission observed around stellar-mass compact objects comes together with continuum emission (e.g., McSwain et al. 2010; Zamanov et al. 2019). There are hints of continuum emission in our stacks of observed spectra (Figure 8), but the question arises as to whether this very faint continuum is consistent or not with an accretion disk around a compact object. A number of arguments show that the observed lack of continuum emission does not discard the accretion disk scenario. First, large H\u03b1 equivalent widths (EWs) are sometimes observed in X-ray binaries. H\u03b1 EWs larger than 100 \u00c5 are not rare during quiescence (Fender et al. 2009; Casares & Torres 2018), with maxima reaching 2000 \u00c5 (Mu\u00f1oz-Darias et al. 2016), which implies line to continuum flux ratios enough to account for our observations (Figure 8). Second, H\u03b1 photons are to be produced by recombination of H atoms photoionized by the accretion disk (e.g., Matthews et al. 2015). Under particular circumstances, the nebular emission produced by a photoionizing source can emit an H\u03b1 line with very little underlying continuum. For example, young starbursts produce H\u03b1 with an EW in excess of 103 \u00c5 (Leitherer et al. 1999), and galaxy-integrated spectra with EW of hundreds of angstrom are not uncommon (e.g., Morales-Luis et al. 2011). The main physical ingredient for the photoionization source to produce little continuum is presenting a hard spectrum with an ionization flux greatly exceeding the flux in the optical. Accretion disks can easily match this requirement since they can be extremely hot with spectra peaking in the far-UV and X-ray.","Citation Text":["Eracleous & Halpern 1994"],"Citation Start End":[[1531,1555]]}
{"Identifier":"2022MNRAS.509..180L__Gonz\u00e1lez-L\u00f3pezlira_et_al._2017_Instance_1","Paragraph":"GCs are the most easily noticeable objects in the haloes of galaxies. However, their surface density in the inner parts always outnumbers that in the haloes (see e.g. Hargis & Rhode 2014; Kartha et al. 2014). In fact, the radial surface density distribution can be well described by the S\u00e9rsic function (Sersic 1968) in its classical form:\n(2)$$\\begin{eqnarray*}\r\nN(R) = N_{\\textrm {e}}\\exp \\left[-b_{n}\\left(\\frac{R}{R_{\\textrm {e}}}\\right)^{(1\/n)}-1\\right],\r\n\\end{eqnarray*}$$where, Re is the effective radius enclosing half the population, n is the S\u00e9rsic index that controls the shape of the profile, and bn = 1.992n \u2212 0.3271. We show the radial distribution of surface number density in Fig. 13. The radial axis is plotted in units of R25. For two of our sample galaxies (M81: Perelmuter & Racine 1995, and NGC 4258: Gonz\u00e1lez-L\u00f3pezlira et al. 2017), GC searches have been carried out using the ground-based data that complement the HST data in the outer parts, which have also been shown (green empty circles). The observed and completeness-corrected values are shown in solid circles of black and magenta colours, respectively, with the fits to these data shown by solid lines of corresponding colours. The numbers inside and outside the Re are corrected by the incompleteness factors corresponding to high and low surface brightness fields, respectively. For M81, the inner most number is corrected by the completeness factor obtained from simulation on the frame that included the bulge and the nucleus. The Re values before and after the incompleteness corrections are shown by dashed vertical lines of black and magenta colours, respectively. The effect of correction is to marginally shift the Re to lower values, which is noticeable only for NGC 4258 and M51 in the figure. In Table 8, we show the values for S\u00e9rsic fit. $N_{{\\rm GC},R_{\\rm e}}$ and $N_{{\\rm GC},R_{\\rm e}}^\\prime$ are the numbers of GCs inside Re before and after correction for contaminants. To estimate the total number of GCs (see Section 5.6), we use the number of GCs within Re, ($N_{{\\rm GC},R_{\\rm e}}^\\prime$).","Citation Text":["Gonz\u00e1lez-L\u00f3pezlira et al. 2017"],"Citation Start End":[[822,852]]}
{"Identifier":"2018AandA...611A.101B__Pierog_et_al._2015_Instance_1","Paragraph":"In order to connect with cosmic ray data, we extrapolate now from the previously studied single collision zone to a population of alike GRBs with a fixed duration of \u0394T = 10 s, which implies that the emission from each GRB comes from \u0394T\u2215tv such collisions. We perform a fit of UHECR observations from the Pierre Auger Observatory, combining the modeling of interactions in the source (computed as described in the previous sections) with the propagation of cosmic rays in the extragalactic space. For the propagation, the SimProp code (Aloisio et al. 2012) has been used with Gilmore EBL (Gilmore et al. 2012) and PSB cross-section model (as defined in Alves Batista et al. 2015). We also consider the extensive air shower in the Earth\u2019s atmosphere and EPOS-LHC (Pierog et al. 2015) is assumed for UHECR-air interactions. We find that a good description of the data is obtained by distributing the GRBs as sources of cosmic rays following the star formation rate Yuksel et al. (2008), assuming a pure 28 Si at the injection. The primary spectrum is described in Eq. (10) and we use here k = 1.8 and P = 2. We fix the following parameters: source evolution, spectral index and cutoff shape at injection, and the nuclear species at the injection. In the present work, we keep this procedure as simple as possible in order to show the power of the method, while leaving a more detailed analysis for a future work. The fit is performed above 1018 eV (Mixed Composition Dip Model) and above 1019 eV (Mixed Composition Ankle Model) by using the combined spectrum (Vali\u00f1o et al. 2015) and the shower depth (Xmax) distributions (Aab et al. 2014), which contain information about the mass of the nucleus interacting with the atmosphere, with a similar procedure as used in di Matteo et al. (2015) and Aab et al. (2017). A scan over (R, L\u03b3) is performed and for each pair the normalization to the experimental flux is found. For each point of the parameter space the number of expected prompt and cosmogenic neutrino events is calculated following Baerwald et al. (2015). For the prompt neutrino flux, the exposure for muon neutrinos is calculated by summing the exposure relative to the IceCube analyses of Aartsen et al. (2015,3 yr) and that of Aartsen et al. (2017, 3 yr) for a total of 1014 GRBs that occurred in the Northern Hemisphere, and that of Aartsen et al. (2017, 5 yr) for 664 GRBs that occurred in the Southern Hemisphere, and comparing the total number of bursts of the combined sample with the assumed 667 bursts per year as in Baerwald et al. (2015). For the cosmogenic neutrino flux, the exposure for muon neutrinos is taken from Aartsen et al. (2016). The exclusion regions (90% C.L.) of the prompt and cosmogenic neutrino analyses are calculated assuming both analyses as background free.","Citation Text":["Pierog et al. 2015"],"Citation Start End":[[763,781]]}
{"Identifier":"2020MNRAS.496.2000C__Isenberg_2008_Instance_1","Paragraph":"These three dimensions in momentum space and three dimensions in real space constitute the 6D phase space. The distribution of neutrinos in this phase space is given by f(r, \u03b8, \u03d5, \u03f5, \u03bc, \u03a6). The four-current of particles in the phase space cell d\u0393com is given by\n(22)$$\\begin{eqnarray*}\r\nj^\\mu = \\int f u^\\mu \\, \\mathrm{d}\\Gamma _\\mathrm{com}.\r\n\\end{eqnarray*}$$The number \u0394N of particles in d\u0393com within some three volume is\n(23)$$\\begin{eqnarray*}\r\n\\Delta N &=& \\int \\sqrt{-g}\\, \\langle \\boldsymbol{j}, \\mathrm{d}x^\\mu \\rangle \\wedge \\mathrm{d}x^\\nu \\wedge \\mathrm{d}x^\\xi \\wedge \\mathrm{d}x^\\sigma \\nonumber \\\\\r\n&=& \\int \\sqrt{-g} j^\\mu \\, \\mathrm{d}x^\\nu \\, \\mathrm{d}x^\\xi \\, \\mathrm{d}x^\\sigma ,\r\n\\end{eqnarray*}$$where $\\langle \\boldsymbol {j},dx^\\mu \\rangle$ denotes the contraction of the four-vector $\\boldsymbol {j}$ with the volume form. This general form can be used to obtain both the number of particles within a space-like three-volume cell d3x as well as the number of particles flowing through cell interfaces per coordinate time dt (two space and one time dimensions). The number of particles within the phase space volume d\u0393com in a cell $\\mathrm{d}x^1 \\, \\mathrm{d}x^2 \\, \\mathrm{d}x^3$ is\n(24)$$\\begin{eqnarray*}\r\nN = \\iiiint \\sqrt{-g} f u^0 \\, \\mathrm{d}\\Gamma _\\mathrm{com}\\, \\mathrm{d}x^1 \\, \\mathrm{d}x^2 \\, \\mathrm{d}x^3,\r\n\\end{eqnarray*}$$and the flux across a cell interface with x1 = const. is\n(25)$$\\begin{eqnarray*}\r\nF = \\iiiint \\sqrt{-g} f u^1 \\, \\mathrm{d}\\Gamma _\\mathrm{com}\\, \\mathrm{d}x^2 \\, \\mathrm{d}x^3 \\, \\mathrm{d}t.\r\n\\end{eqnarray*}$$For a conformally flat metric (Isenberg 2008,used in coconut), with\n(26)$$\\begin{eqnarray*}\r\ng_{\\mu \\nu }= \\begin{pmatrix}-\\alpha ^2+\\beta _i \\beta ^i & \\quad \\beta _x & \\quad \\beta _y & \\quad \\beta _z \\\\\r\n\\beta _x & \\quad \\phi ^4 & \\quad 0 & \\quad 0\\\\\r\n\\beta _y & \\quad 0 & \\quad \\phi ^4 & \\quad 0\\\\\r\n\\beta _z & \\quad 0 & \\quad 0 & \\quad \\phi ^4 \\\\\r\n\\end{pmatrix},\r\n\\end{eqnarray*}$$in terms of the lapse function \u03b1, the conformal factor \u03d5, the shift vector \u03b2i, and \u03b2i = \u03d54\u03b2i, these can be simplified to\n(27)$$\\begin{eqnarray*}\r\nN = \\iiiint \\phi ^6 f \\hat{u}^0 \\, \\mathrm{d}\\Gamma _\\mathrm{com}\\, \\mathrm{d}x^1 \\, \\mathrm{d}x^2 \\, \\mathrm{d}x^3,\r\n\\end{eqnarray*}$$(28)$$\\begin{eqnarray*}\r\nF = \\iiiint \\phi ^6 f \\left(\\frac{\\alpha }{\\phi ^2}\\hat{u}^1+\\beta ^r \\hat{u}^0\\right) \\, \\mathrm{d}\\Gamma _\\mathrm{com}\\, \\mathrm{d}x^2 \\, \\mathrm{d}x^3 \\, \\mathrm{d}t.\r\n\\end{eqnarray*}$$in terms of the Eulerian velocity components $\\hat{u}^\\mu$. It is critical to recognize that because of the invariance of the distribution function we can directly use f in the comoving frame to compute the number of neutrinos within phase space cell d\u0393comv that cross a zone boundary during the time interval dt. Only the neutrino four-velocity in equation (22) must be transformed into the coordinate frame. The problem, however, is that the current j\u03bc only obeys a conservation law \u2207\u03bcj\u03bc = 0 if we define d\u0393com away from a given reference point x\u03bc as the phase space volume filled by particles from d\u0393com at x\u03bc (or in technical terms if we define d\u0393com away from x\u03bc by an exponential map). In general, the exponential map of d\u0393com(x\u03bc) from x\u03bc to another point x\u2032\u03bc will differ from the momentum space volume d\u0393com(x\u2032\u03bc) defined by the same comoving-frame momentum space coordinates, which necessitates the remapping step that we outlined in Section 2.1 and will specify in detail in the next section.","Citation Text":["Isenberg 2008"],"Citation Start End":[[1624,1637]]}
{"Identifier":"2021ApJ...906..109J__Eastwood_et_al._2009_Instance_1","Paragraph":"Reconnection rate in a laminar flow can be estimated, or defined, in terms of normal diffusion of the magnetic field by magnetic diffusivity \u03b7 on large scales or, in other words, in terms of the (root-mean-square, henceforth rms) average distance \n\n\n\n\n\n the field spreads relative to a fixed point. The related random walk of magnetic field lines was discussed by Jokipii (1966; see also Snodin et al. 2013, 2016; Servidio et al. 2014). This is of course the Taylor or normal diffusion (of particles or magnetic field) in which the rms distance between the diffusing material and a fixed point increases as \n\n\n\n\n\n with time; see, e.g., Eyink et al. (2013) and Jafari et al. (2019). For a diffusing magnetic field, \n\n\n\n\n\n. In the absence of turbulence, in a reconnection zone of width \u03b4 and length \u0394 (parallel to the local magnetic field), using the Alfv\u00e9n timescale \n\n\n\n\n\n, and using mass conservation \n\n\n\n\n\n, we recover the reconnection speed\n1\n\n\n\n\n\nThis is, of course, the well-known Sweet\u2013Parker reconnection rate (Parker 1957; Sweet 1958). Reconnection, and\/or other instabilities such as tearing modes (Furth et al. 1963), will in general generate turbulence (Eastwood et al. 2009; Jafari & Vishniac 2018a), with the implication that the laminar Sweet\u2013Parker model is far from realistic in turbulent systems such as most astrophysical fluids. In the turbulence inertial range, i.e., at scales larger than dissipative scale but smaller than the larger scales where Taylor (normal) diffusion of the magnetic field occurs, diffusing particles will undergo super-linear Richardson diffusion; \n\n\n\n\n\n which is a two-particle diffusion, i.e., d is the rms separation between any pair of particles undergoing diffusion in the inertial range. If we consider magnetic diffusion in the turbulence inertial range, we have to consider Richardson diffusion of the field, in terms of the rms distance the field spreads during the time t; see Figure 1. The eddy turnover time, in the inertial range, is of order \n\n\n\n\n\n with d being the length scale perpendicular to the mean magnetic field. Here, \n\n\n\n\n\n denotes the energy transfer rate, with turbulent velocity VT and parallel energy injection length scale \n\n\n\n\n\n. This corresponds to the Richardson diffusion; \n\n\n\n\n\n. The super-linear nature of Richardson diffusion broadens the reconnection zone and thereby enhances the reconnection rate. To see this, using mass conservation \n\n\n\n\n\n, and substituting the Alfv\u00e9n time \n\n\n\n\n\n, one arrives at the fast reconnection speed (Lazarian & Vishniac 1999; Jafari et al. 2018; Lazarian et al. 2019, 2020);\n2\n\n\n\n\n\n\n","Citation Text":["Eastwood et al. 2009"],"Citation Start End":[[1165,1185]]}
{"Identifier":"2018MNRAS.478.3890B__Heckman_et_al._2017_Instance_3","Paragraph":"Rather than AGN feedback, it is possible that the effects we are seeing are from a different process coeval or prior to the onset of AGN accretion. Several works have pointed out that AGN activity coincides with a recent starburst, with the AGN having significant accretion events at least \u223c200\u2009Myr after the starburst has occurred (Davies et al. 2007; Wild et al. 2007; Wild, Heckman & Charlot 2010; Yesuf et al. 2014) giving the neutral material time to propagate out to the impact parameters probed by COS-AGN (Heckman et al. 2017). With a sample of QSO sightlines probing the CGM around 17 low-redshift starburst and post-starburst galaxies, Heckman et al. (2017) have observed a similar signature of enhanced EWs of Ly\u2009\u03b1, Si\u2009iii, and C\u2009iv (the latter of which is not measured in our control sample) relative to a control-matched sample (matched in stellar mass and impact parameter). In the range of impact parameters and stellar masses probed by COS-AGN, the strength of our enhanced EW signature is consistent with the values probed by Heckman et al. (2017). However, the results of Heckman et al. (2017) show strong offsets in the kinematics of the gas from the host galaxy (\u2248100\u2009km s\u22121; see fig. 5 from Heckman et al. 2017), whereas the COS-AGN sightlines do not (bottom panel of Fig. 6). Assuming that the AGN activity was triggered by the starburst, a minimum delay time of 200\u2009Myr could allow for any starburst-driven winds to dissipate and kinematic offsets to no longer be present at the impact parameters of the COS-AGN sample. Although this starburst picture provides a possible explanation of our observations, we caution that starbursts are not the only astrophysical event linked to AGN accretion activity. For example, mergers that trigger the AGN (Ellison et al. 2011, 2013; Satyapal et al. 2014; Silverman et al. 2014; Goulding et al. 2018) could potentially affect the surrounding CGM gas. Past and future work focusing on the CGM of galaxy mergers can further test this result (Johnson et al. 2014; Hani et al. 2017; Bordoloi et al. in preparation).","Citation Text":["Heckman et al. (2017)"],"Citation Start End":[[1043,1064]]}
{"Identifier":"2020MNRAS.496.1051A__Rudolph_et_al._2006_Instance_1","Paragraph":"The radial distribution of S\/H ratios and the corresponding gradient are shown in panel (b) of Fig. 12. We obtain a slope of \u22120.035 \u00b1 0.006\u2009dex\u2009kpc\u22121 (very similar to the one we obtain with the ICF of ADIS20, as can be seen in Table 8), which is consistent with the one of \u22120.041 \u00b1 0.014\u2009dex\u2009kpc\u22121 estimated by Rudolph et al. (2006) using FIR lines and also very similar to the slope of our O\/H gradient. We report a dispersion around the S\/H gradient of 0.10\u2009dex, somewhat larger than the individual observational uncertainties. Recently, Fern\u00e1ndez-Mart\u00edn et al. (2017) reported a slope of \u22120.108 \u00b1 0.006\u2009dex\u2009kpc\u22121 using optical spectra for H\u2009\u2009ii regions located at RG between 5 and 17\u2009kpc. That value of the slope is considerably much steeper than our determination and other previous estimates from the literature (e.g. Shaver et al. 1983; Simpson et al. 1995; Afflerbach, Churchwell & Werner 1997; Rudolph et al. 2006). Esteban & Garc\u00eda-Rojas (2018) observed some H\u2009ii regions with a very low ionization degree \u2013 O2 +\/O  0.03 \u2013 but measurable [S\u2009iii] 6312 or 9069\u2009\u00c5 lines. For such nebulae is possible to assume that S\/H \u2248 S+\/H+ + S2 +\/H+, since the contribution of S3 +\/H+ in those objects is expected to be negligible. The objects with such properties are IC 5146, Sh\u20092-235, Sh\u20092-257, and Sh\u20092-271, whose total abundances of S\/H range from 6.70 to 7.01 but, unfortunately, cover a rather narrow range of RG \u2013 between 9.3 and 11.7\u2009kpc \u2013 and no confident gradient can be estimated with so small baseline. Inspecting Tables 3 and 4 we can see that there is a sizable group of low-ionization degree H\u2009ii regions that lack of determination of their S2 +\/H+ ratios. This is because the rather faint auroral [S\u2009iii] 6312\u2009\u00c5 line could not be detected in those objects. We plan to obtain additional optical spectra covering the bright nebular [S\u2009iii] 9069, 9532\u2009\u00c5 lines of that group of nebulae for trying to increase the number of objects with S\/H ratios determined without ICF and estimate a more precise gradient for this element.","Citation Text":["Rudolph et al. (2006)"],"Citation Start End":[[311,332]]}
{"Identifier":"2022ApJ...929...98T__Parhi_et_al._1999_Instance_1","Paragraph":"It is well known that magnetic fields stabilize the KH instability from the perspective of theoretical and idealized linear calculations such as those presented in Chandrasekhar (1961), which imply that the velocity shear must be super-Alfv\u00e9nic to lead to the generation of the KH instability. However, it is by no means clear that more complex and typical environments such as the solar corona resemble such idealized conditions. The argument for suppression of the KH instability assumes that the magnetic fields and the flow across a boundary are highly aligned. If the magnetic field is more azimuthal, the restoring tension is reduced and the flow can be more susceptible to KH instability. In the extreme situation of a magnetic field orthogonal to two tangential streams, the sheared flow is always unstable to the KH instability, just as in the hydrodynamic case, even for subsonic flows, and the magnetic field basically plays no role. If such nontangential magnetic field conditions occur in the low solar corona, as might be expected, various forms of interchange reconnection can occur along the boundaries of nominally fast and slow wind where loops and open field collide, leading to current sheet formation and magnetic fields that cross from one side of what was a tangential flow to the other. For instance, Schwadron & McComas (2021) argue that nontangential magnetic fields crossing shearing flows can be the origin of switchbacks. So the existence of switchbacks may, in their interpretation, be a manifestation of nontangential magnetic fields associated with tangential flows as well. In these regions, very complicated structures with a weak and probably quite disordered magnetic field can arise. The simple idealized KH instability of Chandrasekhar (1961) is certainly unlikely to occur, and the KH instability may resemble its gasdynamic form instead. Hence, despite previous studies that investigated the conditions for the onset of KH instability in the solar wind and in coronal plumes (e.g., Parhi et al. 1999; Velli et al. 2011) and concluded that the KH instability may arise only in super-Alfv\u00e9nic regions, at distances larger than 10 R\n\u2299 (Velli et al. 2011 found that the KH instability could become important at 0.2\u20130.3 au), the KH instability has been widely imaged in the sub-Alfv\u00e9nic low solar corona, at distances well below 10 R\n\u2299, in strong magnetic field environments associated with active regions (Yuan et al. 2019), at the flanks of coronal mass ejections (Foullon et al. 2011; Ofman & Thompson 2011; Foullon et al. 2013; M\u00f6stl et al. 2013; Nykyri & Foullon 2013), along solar prominence\/corona discontinuity layers (Berger et al. 2010; Ryutova et al. 2010; Hillier & Polito 2018; Yang et al. 2018), in solar jets (Kuridze et al. 2016; Li et al. 2018), associated with streamer wave phenomena (Feng et al. 2013), and in white-light eclipse observations (Druckm\u00fcller et al. 2014). This is because the realistic solar corona is quite unlike a nice cartoon showing well-defined large-scale magnetic field lines; rather, at least locally, it is characterized by a very complex magnetic topology. The onset of the KH instability depends on the orientation of the magnetic field relative to the tangential flows. It follows that even the sub-Alfv\u00e9nic velocity shear along the fast\/slow boundary in the solar corona can lead to KH instability growth, provided that, as expected, the streamer belt region does not have a nicely ordered magnetic field and instead possesses locally a very complicated structure with a nontangential magnetic field configuration. Although theoretically predicted to play an important role in the plasma dynamics between high- and low-speed coronal streams flowing alongside each other (Ismayilli et al. 2018), the KH instability has not yet been detected at the edges of coronal holes adjacent to equatorial streamers, where shear due to the outflow of fluids with different velocities can initiate KH vortices. This lack of evidence motivates the present work, which indeed aims at providing, for the first time, possible observational evidence for manifestations of the KH instability in the solar corona in the interaction region between fast and slow solar wind.","Citation Text":["Parhi et al. 1999"],"Citation Start End":[[2022,2039]]}
{"Identifier":"2017ApJ...844...61D__Millward_et_al._2013_Instance_2","Paragraph":"Now, we must consider a specific distribution of electrons. Ideally, with a well-constructed background-subtracted image, only recently injected coronal material should be visible in excess total brightness; hence, ideally, ne should be identically zero outside the CME. However, reality is not so kind as to clearly delineate between CME and non-CME material. The explosive eruption and ejection of coronal material may rearrange previously quiescent material. If such material is pushed into a line of sight that passes through the CME, then, due to the optically thin nature of the corona, it will be impossible to distinguish newly ejected material from newly arranged material. A further issue that must be confronted when choosing a functional form for ne is the lack of observation or direct measurement on the internal mass distribution of a CME, as well as imprecise knowledge on the morphology of the CME. For example, is a CME balloon-shaped with a circular cross section (Millward et al. 2013) or croissant-shaped with an elliptical cross section (Thernisien et al. 2006, 2009)? Are there one or more cavities in the CME, and if so, where are the cavities located, how wide are they relative to the CME, and how low is the cavity density relative to the edge density? Is the mass distributed symmetrically relative to the central axis of the CME? Does the mass distribution change with time, either by pileup at the CME leading edge or by continuous mass inflow from the lower atmosphere (Bein et al. 2013; DeForest et al. 2013a)? Nevertheless, in spite of such epistemic uncertainty on the internal mass distribution, if the shape of the CME is known\u2014either directly through polarimetry (de Koning & Pizzo 2011) or by assuming a shape (Thernisien et al. 2006; Millward et al. 2013)\u2014and hence the start and end points of the line-of-sight integral are known, one might choose a particularly simple distribution, namely, \n\n\n\n\n\n (Feng et al. 2013b). On the other hand, if both the CME mass distribution and morphology are poorly known, as was the case in the pre-STEREO era, then one might choose an even simpler distribution, that is, a point source,\n12\n\n\n\n\n\nwhere \u03b4 is the Dirac delta function and n0 is a constant. To enable direct comparison of our three-view mass determination with the two-view mass determination of CV09, we employ the point-source distribution.","Citation Text":["Millward et al. 2013"],"Citation Start End":[[1773,1793]]}
{"Identifier":"2017AandA...604A..80M__Propris_et_al._(2013)_Instance_2","Paragraph":" Using the whole sample (\\hbox{$\\bar{z}=0.40$}z\u0305 = 0.40), we find a decreasing faint end for both datasets with consistent values between HST (\u03b1 =  \u2212 0.76 \u00b1 0.07) and Subaru (\u03b1 =  \u2212 0.78 \u00b1 0.06). Separating between low-redshift (\\hbox{$\\bar{z}=0.29$}z\u0305 = 0.29) and high-redshift (\\hbox{$\\bar{z}=0.51$}z\u0305 = 0.51) samples, we find an evolution of the faint end slope of 1.7\u03c3 with HST and 2.6\u03c3 with Subaru. There is thus a mild decrease of the faint end slope (less negative \u03b1) with increasing redshift over the range (0.187 z 0.686). This evolution is in good agreement with recent papers in the literature: in particular Zenteno et al. (2016) found a decrease of the RS faint end at 2.1\u03c3 for a wider range of redshifts (0.1 z 1.13), but with ~ 80% of their clusters being in the same redshift range as ours. De Propris et al. (2013) claim that the evolution in the faint end slope has a significant contribution from surface brightness selection effects. They used HST data of differing depths on a single cluster (MS 1358.4+6254) to show that surface brightness selection effects become important above the formal magnitude limit of their data and that they affect the RS GLF at magnitudes z \u2265 24.5 for 2.7 ks HST exposures (see their Fig. 18). The faint RS for their cluster has F814W \u2212 z = 0.25, implying that the SB selection effects in their sample become important at F814W> 24.75. On the other hand, our CLASH data are significantly deeper than theirs (4.1 ks) and we limit our GLFs at F814W 24.5. Therefore, the real SB selection effects noticed in De Propris et al. (2013) should not be playing a role in our space-based results. In addition, De Propris et al. (2013) claim that previous estimates of the evolution in the RS GLF (e.g.,  De Lucia et al. 2007; Rudnick et al. 2009) were also due to SB effects. Both of those works were based on the same ground-based data with a formal magnitude limit of I = 24 or 24.5 (for the low- and high-redshift clusters, respectively) and the evolution in the GLF was seen over the faintest 2 mag. We cannot directly address the role of SB effects in the EDisCS results without detailed simulations on those data (see below for such simulations for our clusters) but the similarity between our HST and Subaru GLFs imply that the EDisCS evolution in the GLF is not dominated by SB effects. ","Citation Text":["De Propris et al. (2013)"],"Citation Start End":[[1556,1580]]}
{"Identifier":"2021MNRAS.500.2336Y__Matonick_&_Fesen_1997_Instance_1","Paragraph":"Various surveys of SNRs in our Galaxy and nearby galaxies have been carried out at radio, X-ray, Infrared (IR), and optical wavelengths. The first extragalactic SNR candidates were identified in the LMC by Mathewson & Healey (1964) and later confirmed with a combination of radio and optical techniques by Westerlund & Mathewson (1966). To date, a total of 60 SNRs have been confirmed in the LMC with an additional 14 suggested candidates (Maggi et al. 2016; Bozzetto et al. 2017; Maitra et al. 2019). However, sensitivity and resolution limitations severely reduce the effectiveness of the past and present generations of radio and X-ray searches for SNRs in galaxies beyond the Small and Large Magellanic Clouds (MCs) (Goss et al. 1980; Long et al. 1981; Cowan & Branch 1985; Matonick et al. 1997; Matonick & Fesen 1997; Millar, White & Filipovic 2012; Galvin & Filipovic 2014; Sasaki et al. 2018; Lin et al. 2020; Sasaki 2020). As a result, optical studies have produced the largest number (\u223c1200) of new extra-Magellanic SNR candidates. Optical extragalactic searches for SNRs are mainly done by using an emission line ratio criterion of the form [S\u2009ii]\/H\u2009\u03b1 > 0.4\u20130.5 (Mathewson & Clarke 1973; Dodorico, Dopita & Benvenuti 1980; Fesen 1984; Blair & Long 1997; Matonick & Fesen 1997; Dopita et al. 2010b; Lee & Lee 2014; Vu\u010deti\u0107 et al. 2019b, a, 2018; Lin et al. 2020). This criterion separates shock-ionization from photoionization in SNRs from H\u2009ii regions and Planetary Nebulae (PNe) (Frew & Parker 2010). SNR radiative shocks collisionally excite sulphur ions in the extended recombination region resulting in S+, hence the larger contribution of [S\u2009ii] accounting for an increase of the [S\u2009ii] to H\u2009\u03b1 ratio. In typical H\u2009ii regions, sulphur exists predominantly in the form of S++, yielding low [S\u2009ii] to H\u2009\u03b1 emission ratios. Ratios from narrow-band imaging are usually verified spectroscopically, since [N\u2009ii] lines at 6548 and 6584\u2009\u00c5 can contaminate the H\u2009\u03b1 images at an unknown and variable level. Spectroscopic observations of such emission nebulae also can provide other evidence of shock heating, such as strong [O\u2009i] \u03bb6300 emission, elevated [N\u2009ii] to H\u2009\u03b1 with respect to H\u2009ii regions, or high [O\u2009iii] electron temperatures, verifying the candidate as being an SNR (Blair, Kirshner & Chevalier 1981, 1982; Long et al. 1990; Smith et al. 1993; Blair & Long 1997). Although somewhat biased as an isolated criterion, this method is proven and a good way of identifying ordinary radiatively cooling SNRs in nearby galaxies. We note that young, Balmer-dominated SNRs (Chevalier, Kirshner & Raymond 1980) would be missed by this criterion.","Citation Text":["Matonick & Fesen 1997"],"Citation Start End":[[800,821]]}
{"Identifier":"2016MNRAS.461.2383P__Bienaym\u00e9_2000_Instance_1","Paragraph":"These two cases correspond to the upper and lower boundaries of the tilt term. They are usually written as\n\n(73)\n\n\\begin{equation}\n\\frac{{\\mathrm{\\partial} \\sigma _{Rz}^2\\left( {R,0} \\right)}}{{\\mathrm{\\partial} z}} = \\lambda (R)\\frac{{\\sigma _{RR}^2\\left( {R,0} \\right) - \\sigma _{zz}^2\\left( {R,0} \\right)}}{R},\n\\end{equation}\n\nwhere\u03bb(R) \u2208 [0, 1]. The factor \u03bb(R) can be derived either analytically from orbit integration or is assumed \u03bb = 0 for simplicity (e.g. van der Kruit & Freeman 1986; Lewis & Freeman 1989; Sackett & Sparke 1990). Numerical simulations (e.g. Carlberg & Innanen 1987; Bienaym\u00e9 2000) are performed to calculate explicitly the moments for a given gravitational potential. The result of the above studies is that, at the Sun's position \u03bb \u2243 0.5. In our model for the Galactic kinematics, we prefer to adopt the analytical formulation of \u03bb(R) given by Amendt & Cuddeford (1991):\n\n(74)\n\n\\begin{equation}\n\\lambda (R) = {\\left. {\\frac{{{R^2}{\\mathrm{\\partial} _{R,z,z}}\\Phi }}{{3{\\mathrm{\\partial} _R}\\Phi + R{\\mathrm{\\partial} _{R,R}}\\Phi - 4R{\\mathrm{\\partial} _{z,z}}\\Phi }}} \\right|_{\\left( {R,z = 0} \\right)}}.\n\\end{equation}\n\nThe expression (74) is null for a potential separable in cylindrical coordinates because the term $\\frac{{\\mathrm{\\partial} {\\Phi _{{\\rm {tot}}}}}}{{\\mathrm{\\partial} R\\mathrm{\\partial} {z^2}}} = 0$. In spherical coordinates \u03bb = 1, the relation (73) will be used in the following to describe the tilt of the velocity ellipsoid, obtaining for the DM halo\n\n(75)\n\n\\begin{equation}\n{\\lambda _{{\\rm DM}}}(R) = - \\frac{{{R^2}}}{{{R^2}\\left( {{q_\\Phi }^2 - 2} \\right) + 2{h_{r,{\\rm DM}}}^2\\left( {{q_\\Phi }^2 - 1} \\right)}},\n\\end{equation}\n\nand similarly, a unitary constant value for the bulge and stellar halo components is estimated. Finally, for the important contribution of the disc we simplify equation (73) as\n\n(76)\n\n\\begin{equation}\n{\\lambda _D}(R) = - \\int _0^\\infty {\\frac{{k{R^2}{J_1}(kR)h_z^{ - 1}{\\rm d}k}}{{R\\left( {k - \\frac{4}{{{h_z}}}} \\right){J_0}(kR) + 2{J_1}(kR)}}}\n\\end{equation}\n\nthat has to be included in equation (73) with a sum over all the discs components.","Citation Text":["Bienaym\u00e9 2000"],"Citation Start End":[[594,607]]}
{"Identifier":"2020MNRAS.493.3045B__Jaisawal_&_Naik_2015a_Instance_1","Paragraph":"We have used 3.0\u201375.0 keV NuSTAR data to probe any cyclotron line feature. To describe the continuum of 4U 1700\u201337, we have applied the NPEX model [cons*TBpcf*(powerlaw*npex+gaus+gaus)], following the previous work of Jaisawal & Naik (2015a). The NPEX model has been created by adding two cutoffpl models with their cutoff energies tied to each other and keeping the photon index of one to be frozen at \u20132.0. For the best fit, the \u03c72\/d.o.f is found to be 577.64\/475. The fit shows some residuals in the overall spectrum. We added a Gaussian absorption model around 39 keV (following the previous work of Jaisawal & Naik 2015a), but the best fit gives the line energy as $15.44^{+0.56}_{-0.53}$ keV with \u03c72\/d.o.f = 487.32\/472. The width and the depth of the line are found to be $5.47^{+0.90}_{-0.78}$ keV and $1.29^{+0.51}_{-0.39}$, respectively. The chance probability of the line has been computed using the ftest task in xspec. The F-test with this absorption line gives an F value  = 29.2 and a chance probability of 2.61 \u00d7 10\u221217(Table 3). If we add another Gaussian absorption line at 38.9 keV (Energy value frozen) the best-fitting \u03c72\/d.o.f is found to be 486.8\/470. This indicates that the second absorption line is not required for the fit. If we use only one Gaussian absorption line and freeze the line energy at 38.9 keV then the width and the depth of the line are found to be $4.54_{-0.87}^{+0.90}$ keV and $2.87_{-1.05}^{+1.24}$, respectively with a \u03c72\/d.o.f  = 549.2\/473. The ftest gives a chance probability of the 38.9 keV line to be 6.4 \u00d7 10\u22126. So, with NPEX model we find two valid model combinations of the data. One, the presence of a Gaussian absorption line at \u223c15 keV, two, the presence of a Gaussian absorption line at 38.9 keV. But, the presence of both lines together is not supported by the data. The 10.0\u201370.0 keV flux of the source is found to be (2.26 \u00b1 0.01) \u00d7 10\u22129 erg cm-2s-1, much lower than the value (5.6 \u00b1 0.3) \u00d7 10\u22129 erg cm-2s-1, previously reported from SUZAKU data (Jaisawal & Naik 2015a).","Citation Text":["Jaisawal & Naik (2015a)"],"Citation Start End":[[218,241]]}
{"Identifier":"2019MNRAS.482.4956T__Zhang_&_Kojima_2006_Instance_1","Paragraph":"Since pulsation can be detected from the ULX pulsars, their Alfv$\\rm \\acute{e}$n radius should be larger than the neutron star radius RA \u2265 R. This means that the neutron star dipole magnetic field cannot be arbitrarily small. This places a lower limit on the neutron star dipole magnetic field (as a function of mass accretion rate):\n(16)\r\n\\begin{eqnarray*}\r\nB_{\\rm p} \\gt B_{\\rm p,min}= 2.8\\times 10^8 \\skew4\\dot{M}_{\\rm 18}^{1\/2} \\, \\rm G.\r\n\\end{eqnarray*}\r\nDuring the calculations, a typical neutron star radius of $10\\, \\rm km$ is assumed.8 This idea is similar to that of accreting normal neutron stars ( Bildsten et al. 1997; Zhang & Kojima 2006). Accreting neutron stars in high-mass X-ray binaries are often observed as X-ray pulsars. While, the pulsation of accreting neutron stars in low-mass X-ray binaries is often not detected. This may because neutron stars in low-mass X-ray binaries have a lower dipole magnetic field. This also means that if some accreting magnetars have a low dipole magnetic field, then they may not been observed as pulsating sources. Finding pulsations is only one way to confirm the neutron star nature of ULX sources. Dall\u2019Osso et al. (2015)\u2019s lower limits on the mass accretion rate is about $5\\times 10^9 \\, \\rm G$, corresponding to an accretion rate $\\skew4\\dot{M} \\sim 38 \\times 10^{18} \\, \\rm g \\, s^{-1}$ (inset in their fig. 2). According to equation (16), the corresponding lower limit at that accretion rate is about $1.7 \\times 10^9 \\, \\rm G$. A factor of 3 difference is caused by the coefficient of 0.5 in the definition of the magnetospheric radius. The lower limits on magnetic field here (equation 16) is a function of the corresponding mass accretion rate. Although the accretion torque information is not required in obtaining equation (16), Dall\u2019Osso et al. (2015) chose a specific accretion torque and combined it with equation (16). Therefore, a specific numerical value for the magnetic field lower limit can be obtained there.","Citation Text":["Zhang & Kojima 2006"],"Citation Start End":[[632,651]]}
{"Identifier":"2019ApJ...876..135A__Guo_et_al._2013_Instance_1","Paragraph":"Left panel: observed-frame H \u2212 [3.6] color plotted vs. the observed [3.6] magnitude for our sample of BBGs, color-coded by their photometric redshift and scaled in size as a function of their stellar mass (legend shown in right panel). The CANDELS color-selected sample is plotted in gray with the size also scaled according to the masses. The 0.3\u03c3, 1\u03c3, and 2\u03c3 distributions of the mass-limited sample are shown by the light blue density regions. The black triangles represent the four massive galaxies at z > 4.5 found by Huang et al. (2011), black filled circles show the extremely red sources from Caputi et al. (2012), and black diamonds correspond to the H-band dropouts reported by Wang et al. (2016). Lines (dashed) of constant H values are also shown. The darker background indicates the regions of lower completeness for the H band (Guo et al. 2013). Individual error bars are not plotted for clarity, but the average values for our sample of BBGs are shown in the top right corner of each panel. The H-band upper limits (H \u2212 [3.6] lower limits) are also shown. Right panel: rest-frame U \u2212 V vs. V \u2212 J color\u2013color plot, where BBGs are color-coded by SFR and scaled by stellar mass. The SFR lower limits are shown in dark gray. Sources from the CANDELS color-selected sample scaled by mass are also shown. The 0.3\u03c3, 1\u03c3, and 2\u03c3 distributions of the mass-limited sample are shown by the light blue areas. The MIPS-detected galaxies are surrounded by a circle, while the X-ray detected galaxies are highlighted with an asterisk inside the symbol. The Whitaker et al. (2011) upper boundary (black wedge) separates quiescent galaxies (top left) from SFGs (bottom). The black diagonal line denotes an additional criterion proposed in our work (perpendicular to the attenuation vector) to separate bSFGs from dSFGs. The plot also includes a 1 mag attenuation vector (which assumes a Calzetti et al. 2000 law). The H-band upper limits (U \u2212 V lower limits) are also shown.","Citation Text":["Guo et al. 2013"],"Citation Start End":[[842,857]]}
{"Identifier":"2022ApJ...928L..16Y__Murase_et_al._2016_Instance_1","Paragraph":"Fast radio bursts (FRBs) are cosmological radio transients with millisecond durations. Since the first FRB (FRB 010724, the Lorimer burst) was discovered in 2007 (Lorimer et al. 2007), hundreds of FRB sources have been detected, dozens of which are repeaters (e.g., the CHIME\/FRB Collaboration et al. 2021). Recently, a Galactic FRB, FRB 200428, was detected to be associated with SGR J1935+2154 (Bochenek et al. 2020; CHIME\/FRB Collaboration et al. 2020; Mereghetti et al. 2020; Li et al. 2021a; Ridnaia et al. 2021; Tavani et al. 2021), which suggests that at least some FRBs originate from magnetars born from the core collapse of massive stars (e.g., Popov & Postnov 2013; Katz 2016; Murase et al. 2016; Beloborodov 2017; Kumar et al. 2017; Yang & Zhang 2018, 2021; Metzger et al. 2019; Lu et al. 2020; Margalit et al. 2020; Wadiasingh et al. 2020; Wang et al. 2022; Zhang 2022). However, FRB 20200120E was found to be in a globular cluster of a nearby galaxy, M81 (Bhardwaj et al. 2021; Kirsten et al. 2022). This is in tension with the scenario that invokes active magnetars with ages \u227210 kyr formed in core-collapse supernovae (Kremer et al. 2021; Lu et al. 2022) and suggests that FRBs might originate from magnetars formed in compact binary mergers (Margalit et al. 2019; Wang et al. 2020; Zhong et al. 2020; Zhao et al. 2021). Therefore, the physical origin of FRBs is still not well constrained from the data (e.g., Cordes & Chatterjee 2019; Petroff et al. 2019; Zhang 2020; Xiao et al. 2021). The growing FRB detections start to shed light on the diversity among the phenomena. The repeaters presented in the first CHIME FRB catalog have relatively larger widths and narrower bandwidths compared with one-off FRBs (Pleunis et al. 2021). The behaviors of fluence with respect to peak flux exhibit statistically significant differences between bursts with long and short durations (Li et al. 2021c). Multiple origins for the FRB population seem increasingly likely.","Citation Text":["Murase et al. 2016"],"Citation Start End":[[688,706]]}
{"Identifier":"2022MNRAS.515.3299G__Tagawa_et_al._2020a_Instance_2","Paragraph":"In active galactic nuclei (AGNs), the gaseous accretion disc around the SMBH may facilitate binary formation and mergers of stellar-mass compact objects (McKernan et al. 2014; Bellovary et al. 2016; Bartos et al. 2017; Stone, Metzger & Haiman 2017; Tagawa, Haiman & Kocsis 2020a). In this scenario, BHs form in situ in the vicinity of a GN and sink to the inner region due to mass segregation or they are delivered to these regions by infalling globular clusters (Morris 1993; Miralda-Escud\u00e9 & Gould 2000; Freitag, Amaro-Seoane & Kalogera 2006; Hopman & Alexander 2006; O\u2019Leary et al. 2009; Antonini 2014), then get captured in the disc by hydrodynamic drag as they cross the disc (e.g. Goldreich, Lithwick & Sari 2002; Bartos et al. 2017; Yang et al. 2019b; Tagawa et al. 2020a). Alternatively, some BHs may have formed in the disc itself (Levin 2007; Stone et al. 2017). Once in the disc, BHs get transported to the inner regions by exchanging angular momentum with the surrounding gas (Goldreich & Tremaine 1979). In certain regions, the BHs open an annular gap in the accretion disc and accumulate in a narrow range of radii, the so-called migration traps (Bellovary et al. 2016; Secunda et al. 2019, 2020, 2021). Bellovary et al. (2016) argue that migration traps may be expected to be close to the SMBH from \u223c20 to \u223c300 Schwarzschild radii (rS = 2GMSMBH\/c2) from the central SMBH of mass MSMBH.6 However, they may exist in slim discs near the innermost stable circular orbit (ISCO; Peng & Chen 2021) or near the boundary of a gap region if a gap opens due to a heavy stellar-mass BH or an intermediate-mass BH (McKernan et al. 2014). Dynamical encounters frequently happen in migration traps leading to the formation and subsequent merger of binary black holes (BBHs) on short time-scales (Secunda et al. 2019, 2020; Yang et al. 2019a), where the binary separation is efficiently reduced by gas dynamical friction (Escala et al. 2004; Kim & Kim 2007; Baruteau, Cuadra & Lin 2011) to the point where GW emission drives the binaries together. Alternatively, BBHs may also form and merge in the disc outside migration traps (Tagawa et al. 2020a). Because of the deep potential barrier of the SMBH, the merger remnant BH remains near the migration trap and may undergo subsequent mergers with additional BHs, which leads to high BH masses, characteristic spin properties, and possibly non-zero eccentricity identifiable via GW observations (Yang et al. 2019a; Secunda et al. 2020; Tagawa et al. 2020b, 2021a,b; Samsing et al. 2022).","Citation Text":["Tagawa et al. 2020a"],"Citation Start End":[[2128,2147]]}
{"Identifier":"2019AandA...627A.114K__Gorman_et_al._(2015)_Instance_1","Paragraph":"Feature VY may not be a dusty clump in the same sense as the other features since it has a significant flux contribution from the star. In our model, VY is modeled as dusty medium of a Gaussian density distribution centered on VY CMa and uniformly surrounding the star, as in an idealized spherical outflow. Alternatively, however, it could be a clump located in front of or behind the star and seen along the same line of sight. In our implementation, we clear the central space of dust within a radius of 90 AU (equivalent to a diameter of \n\n\n\n0\n\n.\n\u2033\n\n15\n\n\n$ 0{{\\overset{\\prime\\prime}{.}}}15 $\n\n\n) where solids would be too warm to exist heated by the cool but very luminous star. The peak surface brightness simulated with a beam of 152 mas are 34 and 140 mJy beam\u22121 in Bands 6 and 7, respectively. Of that, 13.0 and 30.1 mJy beam\u22121 comes from the blackbody radiation of the model star. O\u2019Gorman et al. (2015) noted that stellar fluxes are underestimated in a blackbody model at millimeter and submillimeter wavelengths owing to the unaccounted presence of an extended radiophotosphere. We extrapolated fluxes expected from their radiophotosphere model getting 12 mJy in Band 6 and 44 mJy in Band 7. Given the small difference with our assumed blackbody fluxes and uncertainties in the radiophotosphere model of O\u2019Gorman et al. (2015), we did not correct our model for this effect. We required 0.014 M\u2299 of dust to explain the observed fluxes which is rather large compared to most other clumps (except C). The stellar contribution partially explains the very low spectral index of 1.7 toward VY because the stellar spectrum is expected to have \u03b1\u2004\u2248\u20042 (Lipscy et al. 2005). At the modeled parameters of VY, the medium must be very thick at visual light and obscure the stellar photosphere from a direct view for a terrestrial observer, as visual observations indeed seem to indicate. The stellar photosphere and its light variations are seen mainly through light scattered in the surrounding medium with a high degree of inhomogeneity (Humphreys et al. 2005).","Citation Text":["O\u2019Gorman et al. (2015)"],"Citation Start End":[[890,912]]}
{"Identifier":"2016MNRAS.462S.323E__Zolensky_et_al._2006_Instance_1","Paragraph":"Samples of dust particles from comet 81P\/Wild 2 were returned to Earth by the Stardust mission in 2006. They were collected in aerogel with entrance velocity of 6\u2009km\u2009s\u22121 and some of them made craters in the aluminium foils separating the aerogel tiles. Analyses of 81P\/Wild 2 samples trapped in aerogel and of residues in the craters in the Al foils have shown that this comet contains mineral components similar to those found in primitive meteorites, including fragments of refractory minerals found in Ca\u2013Al-rich inclusions (CAIs) (Zolensky et al. 2006; Simon et al. 2008; Schmitz et al. 2009), and chondrule fragments (Leroux et al. 2008a; Nakamura et al. 2008). However, during the capture process, the impacting 81P\/Wild 2 dust particles were dispersed along the impact tracks and mixed with the aerogel (Leroux et al. 2008b; Stodolna et al. 2012). The mineralogy of terminal particles (at the end of the tracks) is dominated by olivine and pyroxene minerals, the number of pyroxene grains compared to that of olivine being higher (\u223c1) than that currently observed in carbonaceous chondrites, where olivine grains dominate the silicate mineralogy (number of pyroxenes\/olivines 0.1) (Zolensky et al. 2006; Tomeoka et al. 2008). Apart from these single minerals present at the end of the impact tracks, fine-grained particles mixed with melted aerogel have partially survived the impact and are present along the tracks, together with significant amounts of crystalline silica (Leroux 2012; Roskosz & Leroux 2015). The oxygen isotopic composition of the minerals exhibits 18O\/16O and 17O\/16O ratios within the range observed in carbonaceous chondrites (Nakashima et al. 2012; Ogliore et al. 2015). No evidence for 26Mg excess was found in the 81P\/Wild 2 minerals, suggesting that short-lived 26Al radionuclide was not incorporated at the time of their formation, implying a formation a few million years after the formation of the meteoritic CAIs (Ogliore et al. 2012; Nakashima et al. 2015). The abundance of isotopically anomalous presolar grains in 81P\/Wild 2 is estimated to be less than 1000 ppm (Floss et al. 2013).","Citation Text":["Zolensky et al. 2006","Zolensky et al. 2006"],"Citation Start End":[[535,555],[1189,1209]]}
{"Identifier":"2019ApJ...872..155F__Takiwaki_et_al._2016_Instance_1","Paragraph":"Intensive and extensive efforts in numerical simulations have revealed so far that no successful explosion is obtained in spherical symmetry (e.g., Liebend\u00f6rfer et al. 2001, 2005; Rampp & Janka 2002; Sumiyoshi et al. 2005), and that multidimensional effects are crucially important (Burrows et al. 2006; Bruenn et al. 2009; Marek & Janka 2009; Suwa et al. 2010; M\u00fcller et al. 2012; Takiwaki et al. 2012; Couch 2013; Couch & Ott 2013; Hanke et al. 2013; Murphy et al. 2013; Lentz et al. 2015; Melson et al. 2015; Nakamura et al. 2015; Bruenn et al. 2016; Roberts et al. 2016; O\u2019Connor & Couch 2018). Among them are rotation (Fryer & Heger 2000; Kotake et al. 2003; Thompson et al. 2005; Marek & Janka 2009; Iwakami et al. 2014a; Nakamura et al. 2014; Takiwaki et al. 2016; Summa et al. 2018), a magnetic field (Akiyama et al. 2003; Kotake et al. 2004; Yamada & Sawai 2004; Sawai et al. 2005; Obergaulinger et al. 2006, 2014, 2018; Burrows et al. 2007; Takiwaki et al. 2009; Sawai & Yamada 2014, 2016; M\u00f6sta et al. 2015), non-spherical structures of the progenitor (Couch & Ott 2013; Takahashi & Yamada 2014; Couch et al. 2015; Takahashi et al. 2016), turbulence (Murphy & Burrows 2008; Murphy & Meakin 2011; Murphy et al. 2013; Couch & Ott 2015; Mabanta & Murphy 2018), (magneto)hydrodynamical instabilities (Blondin et al. 2003; Scheck et al. 2006; Blondin & Mezzacappa 2007; Iwakami et al. 2008; Guilet et al. 2010; Wongwathanarat et al. 2010; Fern\u00e1ndez et al. 2014; Takiwaki et al. 2014; Fern\u00e1ndez 2015), general relativistic gravity (Dimmelmeier et al. 2002; Shibata & Sekiguchi 2004, 2005; Kuroda et al. 2012, 2016; M\u00fcller et al. 2012; Ott et al. 2012), and neutrino transport (Nagakura et al. 2014, 2017, 2018; Dolence et al. 2015; Pan et al. 2016). It is true that large-scale dynamical simulations have played a crucial role in recent progresses in our understanding of these ingredients, but we believe that a more phenomenological approach that employs toy models still plays an indispensable and complementary role to understand each effect more deeply.","Citation Text":["Takiwaki et al. 2016"],"Citation Start End":[[750,770]]}
{"Identifier":"2019AandA...630A..30L__H\u00e4ssig_et_al._(2015)_Instance_1","Paragraph":"The many unexpected surprises of comet 67P\/Churyumov-Gerasimenko (hereafter 67P) revealed by the historic Rosetta mission highlight the importance of observing the evolution of comets throughout their orbits. One of the surprises was the drastic heterogeneity in both the major and minor volatile species in the coma that was observed early on in the mission (H\u00e4ssig et al. 2015; Luspay-Kuti et al. 2015, hereafter ALK15). When Rosetta first arrived at comet 67P in August 2014, the Rosetta Orbiter Mass Spectrometer for Ion and Neutral Analysis\/Double Focusing Mass Spectrometer (ROSINA\/DFMS; Balsiger et al. 2007) detected large diurnal variations in the intensity profiles of various species in the coma from distances to the comet as far as 250 km. At this time, 67P was still at a distance of about 3 AU and inbound from the Sun. The intensity variations in the major and minor volatile species were found to be periodic, and were dependent on both the observing sub-spacecraft latitude and longitude (H\u00e4ssig et al. 2015; Luspay-Kuti et al. (2015)). As reported in H\u00e4ssig et al. (2015), the intensity of H2O in the coma dominated the overall signal, with maxima in the H2O signal every ~6 h, about twice during a rotation. Interestingly, however, CO2 and CO displayed a separate additional maximum when the H2O signal was near its minimum. This independent maximum in CO2 and CO only occurred at negative observing latitudes that are associated with a particular \u201cview\u201d of Rosetta at 67P, with the larger lobe blocking out the neck and head. At this time, 67P had not yet reached its first equinox (10 May 2015), and the poorly illuminated southern hemisphere was experiencing winter. In addition, the largest H2O activity was localized at the well-illuminated neck region, as also seen by the Microwave Instrument on the Rostta Orbiter (MIRO; Gulkis et al. 2015; Biver et al. 2015; Lee et al. 2015) and by the Visible InfraRed Thermal Imaging Spectrometer (VIRTIS; Bockel\u00e9e-Morvan et al. 2015; Migliorini et al. 2016). VIRTIS also measured weak H2O production in regions with low solar illumination, while CO2 was outgassing from both illuminated and non-illuminated regions pre-inbound equinox (Bockel\u00e9e-Morvan et al. 2015; Migliorini et al. 2016; Fink et al. 2016). The observed outgassing pattern of the major cometary species suggested that CO and CO2 may be sublimating from a depth below the diurnal skin depth.","Citation Text":["H\u00e4ssig et al. 2015"],"Citation Start End":[[360,378]]}
{"Identifier":"2020AandA...642A...2R__Isavnin_et_al._2013_Instance_1","Paragraph":"Lynch et al. (2010) and the review by Rouillard (2011) (see Fig. 1) paint a rather positive picture for our ability to connect RS observations and IS measurements of certain interplanetary propagating (IP) structures, particularly under solar minimum conditions. The reality, however, is more complicated. Despite the availability of continuous imaging and tracking of IP disturbances from the STEREO heliospheric imagers, CME time-of-arrivals and especially speeds at 1 AU can have errors upwards of 13 h and 200 km s\u22121, respectively, even for de-projected trajectories (Colaninno et al. 2013). Reconstructions of the CME shape from RS observations rarely agree with in situ ones (e.g., Nieves-Chinchilla et al. 2013; Wood et al. 2017). Some of the discrepancies can be attributed to limitations of the triangulation techniques (whether applied on J-maps or not) arising from the optically thin nature of the white light emission (Howard et al. 2012). Much of the remaining discrepancies, though, should arise from interaction between the structures and the ambient solar wind and\/or from internal evolution that can lead to deflections (e.g., Isavnin et al. 2013, 2014), rotations (Nieves-Chinchilla et al. 2012) or distortions (Nieves-Chinchilla et al. 2013). We note, however, that these effects are inferred rather than observed directly in the data. In terms of the 3D global morphology, there is a current debate about distortion or deformation of the internal magnetic field configuration associated to the CMEs. The elongated dark cavity shape associated with the CME flux rope, called \u201cpancaking\u201d, provides an evidence of a highly distorted global structure, probably due to the solar wind interaction (Odstrcil & Pizzo 1999; Riley & Crooker 2004; Owens 2006; Savani et al. 2010; DeForest et al. 2013; Owens et al. 2017). However, another interpretation is that it could be an effect of the projection of a 3D structure in the plane of the sky, creating the appearance of distortion (Vourlidas et al. 2011; Nieves-Chinchilla et al. 2013). In contrast, from the in situ perspective, the magnetic field strength asymmetries observed can be interpreted as a consequence of the local effect of flux rope expansion (Farrugia et al. 1993; Osherovich et al. 1993; D\u00e9moulin & Dasso 2009) or erosion (Dasso et al. 2007; Ruffenach et al. 2012), but it could also be a result of the distortion (Hidalgo et al. 2002; Berdichevsky 2013; Nieves-Chinchilla et al. 2018a). Longitudinal deflections during inner heliospheric transit are particularly hard to quantify because heliospheric imaging is taken from and restricted along the ecliptic plane. This is the reason, for example, why the \u201cgarden-hose spiral\u201d nature of solar wind streams has never been imaged directly but only inferred by de-projection techniques based on J-maps (Rouillard et al. 2008; Sheeley et al. 2008). While major progress was achieved with STEREO data, establishing the exact longitudinal extent of streamer blowouts (Vourlidas & Webb 2018), whether CMEs deflect longitudinally or merge with each other, where CIRs form or how shocks evolve require observations away from the ecliptic (Gibson et al. 2018). This is an area where Solar Orbiter imaging, particularly from Solar OrbiterHI (Howard et al. 2020) will be trailblazing. However to fully exploit this data we need to adapt our techniques to account for rapidly varying vantage points located outside the ecliptic plane.","Citation Text":["Isavnin et al. 2013"],"Citation Start End":[[1145,1164]]}
{"Identifier":"2016ApJ...833....7Y__Owen_&_Wu_2013_Instance_1","Paragraph":"We use the N-body simulation package\u2014MERCURY (Chambers 1999)\u2014to numerically investigate the effects of photo-evaporation on the dynamical evolution of planet\u2013satellite systems. We choose the Bulirsch\u2013Stoer integration algorithm, which can handle close encounter accurately. It is important in the simulations, as we will see below, that many close encounters among moons and the planet are expected to happen. Collisions among moons, the planet, and the central star are also considered in simulations and treated simply as inelastic collisions without fragmentations. Each simulation consists of a central star, a planet, and some moons orbiting around the planet. The photo-evaporation is simply modeled as a slow (adiabatic) and isotropic mass-loss process of the planet. In reality, the photo-evaporation is a very slow process on a timescale of the order of 107\u2013108 year (Owen & Wu 2013). However, it is impractical and unnecessary to perform a simulation on such a long timescale. Instead, we model the mass-loss process on a timescale of \u03c4evap, and each simulation typically lasts for several \u03c4evap. As long as the adiabatic requirement is met, i.e., the mass-loss timescale is much longer than the dynamical timescale of the system (\u03c4evap \u226b Pp, where Pp is the orbital period of the planet), one could study the dynamical effects of the mass-loss process equivalently. As we discussed in Section 3.3, the results converge if \u03c4evap > 102\u2013103 Pp, indicating the adiabatic condition is met. Therefore, in all other simulations, we set \u03c4evap = 104 Pp. Other parameters are set to represent the typical values of Kepler planets. In particular, we consider a planet\u2013satellite system orbiting a star of solar mass (M\u22c6 = M\u2299) in a circular orbit (ep = 0.0) with semimajor axis of ap = 0.1 au. The orbit has a period of \u223c10 days (typical value of Kepler planets), and it is sufficiently close to the central star to be subject to significant photo-evaporation effect (Owen & Wu 2013), which removes massive hydrogen envelopes of the planet. The planet has an initial mass of Mpi and a final mass of Mpf after photo-evaporation. In this paper, we adopt Mpi = 20 M\u2295 and Mpf = 10 M\u2295 nominally (close to the standard model adopted in Owen & Wu 2013). The mean density of the planet is set to the same as that of Neptune (1.66 g cm\u22123). The effect of changing the planetary density is discussed in Section 3.3. We performed a number of sets of simulations by considering different planet\u2013satellite configurations. Similar to the definition in MERCURY, hereafter, we define \u201csmall moons\u201d as test particles (TPs) whose mutual gravity and corresponding effects on the planet and the star are ignored, while \u201cbig moons\u201d are gravitationally important enough that their gravitational effects are fully considered. Table 1 lists the initial setups and parameters of various simulations, whose results are presented in the following subsections.","Citation Text":["Owen & Wu 2013"],"Citation Start End":[[877,891]]}
{"Identifier":"2022AandA...663A..70F__Mashonkina_et_al._2007_Instance_1","Paragraph":"Summarizing the discussion above: we have a number of suggested r-process sites, but only one of them is proven by a direct observation of the explosive event. Observations of low metallicity stars show essentially three types of patterns, a weak or limited r-process, a strong solar-type r-process, and an actinide-boosted r-process (in some publications, this is also referred to as a weak, main, and strong r-process). Whether the latter two types are produced in different sites or a result of variations within the same site (e.g., neutron star mergers) is still debated. The question is now how such observations can point back to the r-process sites, and whether it is possible to identify features which can provide additional insight. A promising approach is to look for correlations among different elements, which might directly identify the nucleosynthesis of a specific site (see e.g., Barklem et al. 2005; Fran\u00e7ois et al. 2007; Mashonkina et al. 2007; Kratz et al. 2008). Cowan et al. (2005) compared the abundances of Fe, Ge, Zr, and r-process Eu in low metallicity stars. They found a strong correlation of Ge with Fe, indicating the same nucleosynthesis origin (core-collapse supernovae), a weak correlation of Zr with Fe \u2013 indicating that other sites than regular core-collapse supernovae (without or low Fe-ejection) contribute as well \u2013, and no correlation between Eu and Fe, essentially pointing to a pure r-process origin with negligible Fe-ejection. More recent data from the SAGA and JINA databases (Suda et al. 2008; Abohalima & Frebel 2018) permit a correlation between Eu and Fe for [Eu\/Fe]  0.3, that is for stars with lower than average r-process enrichment. If interpreted in a straightforward way, this would point to a negligible Fe\/Eu ratio (in comparison to solar ratios) in the major r-process sources, while a noticeable coproduction of Fe with Eu is possible in less strong r-process sources, for example with a weak r-process. Such cases could again be identified with the limited-r entry in observations (Hansen et al. 2018).","Citation Text":["Mashonkina et al. 2007"],"Citation Start End":[[942,964]]}
{"Identifier":"2016ApJ...821..107G__Schwadron_et_al._2011_Instance_2","Paragraph":"We repeated the plasma pressure calculation presented by Schwadron et al. (2011) and Fuselier et al. (2012) for the new ENA energy spectrum. The results for the downwind hemisphere and for the Voyager 1 region are summarized in Table 3. The measured intensity \n\n\n\n\n\n\nj\n\n\nENA\n\n\n\n\n of neutralized hydrogen at a given energy translates into a pressure of the parent ion population in the heliosheath times the integration length along the line of sight, \u0394P \u00d7 l, in the following way:\n3\n\n\n\n\n\u0394\nP\n\u00d7\nl\n=\n\n\n\n4\n\u03c0\n\n\n3\n\n\nn\n\n\nH\n\n\n\n\n\n\n\nm\n\n\nH\n\n\nv\n\n\n\n\n\nj\n\n\nENA\n\n\n(\nE\n)\n\n\n\u03c3\n(\nE\n)\n\n\n\n\u0394\nE\n\n\n\nc\n\n\nf\n\n\n\n\n\n\n4\n\n\n\n\n\n\nc\n\n\nf\n\n\n=\n\n\n\n\n\n(\nv\n+\n\n\nu\n\n\nR\n\n\n)\n\n\n2\n\n\n\n\n\n\nv\n\n\n4\n\n\n\n\n\n(\n\n\nv\n\n\n2\n\n\n+\n4\n\n\nu\n\n\nR\n\n\n2\n\n\n+\n2\n\n\nu\n\n\nR\n\n\nv\n)\n.\n\n\nIn Equation (3), \u0394E denotes the width of the respective energy bin; for the typical radial velocity of solar wind in the flanks and the downwind hemisphere of the inner heliosheath, we assumed uR = 140 km s\u22121 as measured by Voyager 2, whereas uR = 40 km s\u22121 for the heliosheath in the Voyager 1 direction (Schwadron et al. 2011; Gloeckler & Fisk 2015). For the density of neutral hydrogen in the inner heliosheath a constant nH = 0.1 cm\u22123 is assumed (Schwadron et al. 2011; Gloeckler & Fisk 2015). The charge-exchange cross section between protons and neutral hydrogen decreases from (4 to 2) \u00d7 10\u221215 cm\u22122 for 0.015 to 1.821 keV (Lindsay & Stebbings 2005). The integration length l for ENA production in the plasma is approximately the thickness of the inner heliosheath. The part of Equation (3) without the velocity factor cf can be interpreted as stationary pressure. The total pressure or dynamic pressure is the stationary pressure times this factor. Integrating over all energy bins in Table 3, we obtain the total plasma pressure times heliosheath thickness as P \u00d7 l = 304 pdyn cm\u22122 au for the downwind hemisphere and 66 pdyn cm\u22122 au for the Voyager 1 region (1 pdyn cm\u22122 au = 0.015 N m\u22121). If we want to put these numbers into the context of other studies, we face two problems. First, the uncertainty of the total pressure is large given the upper limits in the two lowest energy bins. Second, heliosheath plasma more energetic than 2 keV will produce ENAs that cannot be detected with IBEX-Lo. We therefore used the observed median j = 0 cm\u22122 sr\u22121 s\u22121 keV\u22121 for heliospheric ENAs in the two lowest energy bins of IBEX-Lo and relied on the study by Livadiotis et al. (2013). They compared the expected plasma pressure from a kappa distribution of protons with the plasma pressure derived from IBEX-Hi energy spectra: the energy range between 0.03 and 2 keV, roughly corresponding to the IBEX-Lo range, covered more than half of the total plasma pressure predicted from a kappa distribution. The authors found a total plasma pressure of P = 2.1 pdyn cm\u22122 for all sky directions except for the ENA Ribbon. Gloeckler & Fisk (2015) presented a multi-component plasma model for the heliosheath to explain Voyager and IBEX observations. At low energies they assumed the ENA energy spectra provided by Fuselier et al. (2012). They derived a total pressure of 2.5 pdyn cm\u22122 in all three plasma regions in the nose of the heliotail (Gloeckler & Fisk 2015). Pressure contributions from the slowed solar wind, magnetic pressure, and the pressure exerted from pickup ions and anomalous cosmic rays all had to be taken into account to obtain this total pressure.","Citation Text":["Schwadron et al. 2011"],"Citation Start End":[[1006,1027]]}
{"Identifier":"2016AandA...592A..19C__Johansson_et_al._2012a_Instance_2","Paragraph":"Figure 24 shows that the median Re increases for increasing mass (from ~5 Kpc to ~20 Kpc) and, given a stellar mass, our passive ETGs have median sizes smaller than those of the parent sample even by ~15% (at the highest masses). Furthermore, the entire Re distribution is extended to smaller radii in the case of passive ETGs, especially for log\u2009(M\/M\u2299) \u2273 11.5. Remembering that small differences in galaxy size imply large differences in stellar mass density, the derived trends suggest that the ETGs analyzed in this work should have formed from higher density progenitors which, in the hypothesis of no coeval mergers, do not increase their mass during the evolution. On the other hand, we cannot exclude the possibility that galaxies in the parent sample have experienced more dry mergers than our massive and passive galaxies, which have increased their size across cosmic time (Naab et al. 2009; Johansson et al. 2012a). The observed trend is also in agreement with some literature studies at low redshift which also suggest that more compact galaxies contain older stellar populations than larger ones (e.g. Saracco et al. 2009; Shankar & Bernardi 2009; Williams et al. 2010; Poggianti et al. 2013; McDermid et al. 2015). In general, this is qualitatively consistent with the high zF inferred by our analysis (z> 5) since, according to the cosmological evolution of the baryonic matter density, gas was denser at these cosmic epochs. These findings also agree with recent observations of compact systems (with various level of SF) at z ~ 2\u22123, which could be identified as local massive ETGs progenitors (e.g. Daddi et al. 2004; Cattaneo et al. 2013; Finkelstein et al. 2013; Marchesini et al. 2014; Nelson et al. 2014; Williams et al. 2015). Moreover, recent models of elliptical galaxy formation also predict the formation of compact and dense progenitors at high redshifts (e.g. Johansson et al. 2012a; Naab et al. 2014 and references therein). ","Citation Text":["Johansson et al. 2012a"],"Citation Start End":[[1888,1910]]}
{"Identifier":"2019AandA...627A.130D__Broadhurst_et_al._2019_Instance_3","Paragraph":"Gravitational-wave astronomy has recently become a reality with the first detection of gravitational waves (GW hereafter) by the Laser Interferometer Gravitational-Wave Observatory (LIGO) and Virgo ground-based interferemeters. To date, eleven events have been reported by the LIGO and Virgo detectors (Abbott 2018), and this number will quickly increase to tens of events in the coming years. Some of these events may correspond to gravitationally lensed events with magnification factors ranging from a few tens to a few hundreds (Dai et al. 2017; Ng et al. 2018; Li et al. 2018a; Smith et al. 2018a,b; Broadhurst et al. 2019). Recent works have studied lensing effects in the existing LIGO\/Virgo O1 and O2 events (Hannuksela et al. 2019; Broadhurst et al. 2019), while Smith et al. (2018c) searched for candidate galaxy cluster lenses for the GW170814 event. The most likely lenses for such events would be massive galaxies or galaxy clusters (Ng et al. 2018; Dai et al. 2017; Smith et al. 2018b; Broadhurst et al. 2019). On the other extreme of the lens mass regime, compact objects with masses of a few hundreds to a few tens M\u2299 can also act as lenses (Lai et al. 2018). In this case, the geometric optics limit is not valid since the Schwarszchild radius of the lens is comparable to the wavelength of the wave. For these relatively low masses, the lensing effect has a modest impact on the average magnification, but it can introduce a frequency dependence on the magnification (see for instance, Jung & Shin 2019; Lai et al. 2018). An even smaller mass regime was considered in Christian et al. (2018) where the authors find that lenses with a mass as low as 30 M\u2299 could be detected with current experiments. They also consider future, higher-sensitivity experiments and show how they can push the limit to even smaller masses of order 1 M\u2299. These conclusions are, however, obtained assuming isolated microlenses and without accounting for the effect of the macromodel, or other nearby microlenses. In the small mass regime, microlenses such as neutron stars have been also considered as scattering sources of GWs, and it is found that a GW can be focussed at a focal point near the neutron star surface (Halder et al. 2019; Stratton & Dolan 2019).","Citation Text":["Broadhurst et al. 2019"],"Citation Start End":[[1000,1022]]}
{"Identifier":"2022ApJ...934...10L__Lynch_et_al._2008_Instance_1","Paragraph":"Classic examples of this approach to investigating solar eruptions have been the many simulations of the magnetic breakout model, which uses reconnection as the eruption mechanism (Antiochos 1998; Antiochos et al. 1999). In typical breakout simulations magnetic energy is built up in the corona by prescribed surface motions at the lower boundary, either large-scale shearing flows (e.g., Karpen et al. 2012) or small-scale randomized motions leading to helicity condensation (e.g., Dahlin et al. 2022). The normal component of the coronal field is held fixed at this lower boundary, which is presumed to represent the high-beta photosphere, and the coronal flux is assumed to be multipolar so that a separatrix surface and null point are present. Note that even a bipolar AR implies the presence of a coronal separatrix surface and null when the field of the bipole is embedded in the solar global dipole (e.g., Lynch et al. 2008). The initial conditions in most breakout simulations are taken to be the minimum-energy potential field (\nJ\n \u2261 \u2207 \u00d7 \nB\n\/\u03bc\n0 = 0 where \nB\n is the magnetic field), but the imposed boundary motions inject magnetic free energy into the corona, leading to the formation of a sheared-arcade filament channel with current sheets at the overlying separatrix surface and null between the filament flux system and neighboring flux systems. In the breakout model, magnetic reconnection (often called breakout reconnection) at the separatrix current sheet is necessary so that the sheared filament flux can expand upward sufficiently for a vertical current sheet to form below\/inside this flux system. The onset of magnetic reconnection in this vertical current sheet (often loosely described as flare reconnection) greatly accelerates the eruption, leading to the ejection of the shear and the relaxation of the corona back down to a near-potential state (e.g., Karpen et al. 2012; Wyper et al. 2017). A related but physically distinct reconnection-driven eruption model is the so-called tether cutting, which posits that with sufficient boundary shear a vertical current sheet can form spontaneously in the filament channel, without the need for breakout reconnection, leading to an eruption (e.g., Moore et al. 2001). Recent simulations using the same approach as the breakout ones just described appear to support this scenario (Jiang et al. 2021).","Citation Text":["Lynch et al. 2008"],"Citation Start End":[[913,930]]}
{"Identifier":"2019ApJ...874L..32C__Taylor_et_al._2003_Instance_1","Paragraph":"Cygnus A, at z = 0.0562, is 10 times closer than the next radio galaxy of similar radio luminosity.4\n\n4\nRadio luminosity >1045 erg s\u22121.\n The nuclear regions in Cygnus A have been observed extensively at radio through X-ray wavelengths (Carilli & Barthel 1996). The inner few arcseconds is a complex mix of optically obscuring dust clouds (Vestergaard & Barthel 1993; Whysong & Antonucci 2004; Lopez-Rodriguez et al. 2014; Merlo et al. 2014), atomic gas seen in narrow line emission (Stockton et al. 1994; Taylor et al. 2003), H i 21 cm absorption toward the inner radio jets, with a neutral atomic column density >1023 cm\u22122, depending on H i excitation temperature (Struve & Conway 2010), polarized, broad optical emission lines due to scattering by dust (Antonnuci et al. 1994; Ogle et al. 1997), and a highly absorbed hard X-ray spectrum with a total gas column density of \u223c3 \u00d7 1023 cm\u22122 (Ueno et al. 1994; Reynolds et al. 2015). VLBI radio observations at 0.05 mas resolution reveal highly collimated jets originating on scales \u223c200 times the Schwarzschild radius (Boccardi et al. 2016). Tadhunter et al. (2003), derive a black hole mass of 2.5 \u00b1 0.7 \u00d7 109 M\u2299 from HST and Keck spectroscopy of Pa-\u03b1 and [O iii], and conclude that Cygnus A contains an AGN with a bolometric luminosity of order 1046 erg s\u22121, comparable to high redshift quasars (Runnoe et al. 2012). This AGN is highly obscured in the optical due to dust along our line of sight, with Av > 50 magnitudes, based on near-IR spectroscopy (Imanishi & Ueno 2000). Studies of the mid- to far-IR spectral and polarization properties have led to a model of a clumpy, dusty torus obscuring the AGN in Cygnus A, with a radius of at least 130 pc, although these conclusions are based on spatially integrated properties; these observations did not have the spatial resolution to resolve the torus, and hence are partially contaminated by emission from the radio core-jet (Privon et al. 2012; Lopez-Rodriguez et al. 2018).","Citation Text":["Taylor et al. 2003"],"Citation Start End":[[505,523]]}
{"Identifier":"2020MNRAS.492.1061V__Fiorentino_&_Monelli_2012_Instance_1","Paragraph":"Variable stars have a long tradition of being tracers of different stellar populations. The most well known of the pulsating variable stars in dwarf galaxies are RR Lyrae stars, which unequivocally trace an old population (>10 Gyr, Smith 1995). RR Lyrae stars have been found in almost all of the satellites of the Milky Way (recent compilation in Mart\u00ednez-V\u00e1zquez et al. 2019), even in low-luminosity systems such as Segue I (Simon et al. 2011). This is a clear indication that an old population is ubiquitous among dwarf galaxies. On the other hand, anomalous Cepheids, which are pulsating stars above the horizontal branch, may trace an intermediate age population. Their progenitors are likely massive (1\u20132 M\u2299) stars that have evolved off the main sequence (Fiorentino & Monelli 2012). These stars have been commonly found in the classical satellites of the Milky Way but are rare in globular clusters (Clement et al. 2001), which contain exclusively old populations. Since they have been observed also in dwarf spheroidal (dSph) galaxies with predominantly old stellar populations like Draco and Ursa Minor (Nemec, Wehlau & Mendes de Oliveira 1988; Kinemuchi et al. 2008), other formation channels are needed for these stars. Binary evolution can also bring stars to the region of the instability strip where anomalous Cepheids live (Gautschy & Saio 2017). Thus, old progenitors are also possible. Dwarf Cepheid stars (a.k.a. SX Phoenicis and\/or \u03b4 Scuti stars) are also pulsating stars in the instability strip, but they are found below the horizontal branch (Breger 2000). Their very short periods (just a few hours) and small amplitudes make them very hard to detect, especially in distant systems. They may trace either intermediate-age populations (massive main sequence stars in the instability strip) or old populations (pulsating blue stragglers). Finally, Classical Cepheids indicate the presence of young populations (0.03\u20130.7 Gyr). They are rare in dSph galaxies, with Leo I being an exception (Fiorentino et al. 2012).","Citation Text":["Fiorentino & Monelli 2012"],"Citation Start End":[[762,787]]}
{"Identifier":"2017MNRAS.464.4495G__Griv_&_Wang_2014_Instance_1","Paragraph":"The column S in the table shows that at least two minima of line-of-sight velocity irregularity are present for all numbers m considered: the first short-wavelength minimum with the radial scale \u03bbrad \u2248 1 kpc and the second long-wavelength minimum with the scale \u03bbrad \u2248 2.5 kpc. There are three possible ways to explain these scales. In the first place, one can identify the scales as fully independent global modes. This is because for a fixed number of spiral arms m, in a self-gravitating rotating astrophysical disc configuration, there usually exists a sequence of radial wavelengths (Bertin et al. 1977; Griv & Chiueh 1998; Griv et al. 2000, 2006, Griv et al. 2012; Fujii et al. 2011; Griv & Wang 2014). The non-negative azimuthal mode number m gives the number of spiral arms, whereas the real radial wavenumber krad = 2\u03c0\/\u03bbrad determines the radial structure.4 Secondly, one can identify these scales as Fourier harmonics of a single wave (Mishurov et al. 1979, p. 151). And thirdly, the phase \u03d50 that we found is between \u221248\u00b0 and \u221270\u00b0 and\/or between 327\u00b0 and 351\u00b0 for all m. To be consistent with the current ideas as to the Galactic spiral structure (Vall\u00e9e 2008, fig. 1 therein; Xu et al. 2013, fig. 12 therein; Hou & Han 2014, fig. 5 therein; Nakanishi & Sofue 2016, fig. 7 therein), we may interpret this finding to mean that the kinematics of 223 Cepheids investigated reveals the segments of different large-scale arms. Namely, whatever spiral m models we consider, the Sun is located at the outer edge of a local spiral arm (at the outer edge of the Cygnus\u2013Orion arm; \u03d50 \u2248 \u221260\u00b0 for the first minimum) and\/or at the outer edge of an inner spiral arm (at the outer edge of the Carina\u2013Sagittarius arm, very close to the arm centre; \u03d50 \u2248 330\u00b0 for the second minimum) (for notation, see Reid et al. 2009, fig. 1 therein; Choi et al. 2014, fig. 14 therein; Sanna et al. 2014, fig. 6 therein; Camargo et al. 2015, fig. 12 therein; Sato et al. 2014, fig. 3 therein). Basharina et al. (1980) have already confirmed the Sun to be located in the local arm. Further studies are needed to clarify the issue (see also Hou & Han 2014).","Citation Text":["Griv & Wang 2014"],"Citation Start End":[[690,706]]}
{"Identifier":"2022ApJ...935..135B__Cui_et_al._2012_Instance_1","Paragraph":"Disk galaxies typically reveal out-of-equilibrium features due to incomplete equilibration. These may appear in the form of bars and spiral arms, which are large-scale perturbations in the radial and azimuthal directions, responsible for a slow, secular evolution of the disk. In the vertical direction, disks often reveal warps (Binney 1992). In the case of the Milky Way (MW) disk, which can be studied in much greater detail than any other system, recent data from astrometric and radial velocity surveys such as SEGUE (Yanny et al. 2009), RAVE (Steinmetz et al. 2006), GALAH (Bland-Hawthorn et al. 2019), LAMOST (Cui et al. 2012), and above all Gaia (Gaia Collaboration et al. 2016, 2018a, 2018b) have revealed a variety of additional vertical distortions. At large galactocentric radii (>10 kpc) this includes, among others, oscillations and corrugations (Xu et al. 2015; Sch\u00f6nrich & Dehnen 2018), and streams of stars kicked up from the disk that undergo phase mixing, sometimes referred to as \u201cfeathers\u201d (e.g., Price-Whelan et al. 2015; Thomas et al. 2019; Laporte et al. 2022). Similar oscillations and vertical asymmetries have also been reported in the solar vicinity (e.g., Widrow et al. 2012; Williams et al. 2013; Yanny & Gardner 2013; Gaia Collaboration et al. 2018b; Quillen et al. 2018; Bennett & Bovy 2019; Carrillo et al. 2019). One of the most intriguing structures is the phase-space spiral discovered by Antoja et al. (2018) and studied in more detail in subsequent studies (e.g., Bland-Hawthorn et al. 2019; Li 2021; Li & Widrow 2021; Gandhi et al. 2022). Using data from Gaia DR2 (Gaia Collaboration et al. 2018a), Antoja et al. (2018) selected \u223c900,000 stars within a narrow range of galactocentric radius and azimuthal angle centered around the Sun. When plotting the density of stars in the (z, v\n\nz\n)-plane of vertical position, z, and vertical velocity, v\n\nz\n, they noticed a faint, unexpected spiral pattern, which became more enhanced when color-coding the (z, v\n\nz\n)-\u201cpixels\u201d by the median radial or azimuthal velocities. The one-armed spiral makes two to three complete wraps, resembling a snail shell, and is interpreted as a signature of phase mixing in the vertical direction following a perturbation, which Antoja et al. (2018) estimate to have occurred between 300 and 900 Myr ago. More careful analyses in later studies (e.g., Bland-Hawthorn et al. 2019; Li 2021) have nailed down the interaction time to \u223c500 Myr ago.","Citation Text":["Cui et al. 2012"],"Citation Start End":[[617,632]]}
{"Identifier":"2017AandA...602A..82D__Ram\u00edrez_et_al._2012_Instance_1","Paragraph":"Brown dwarfs and giant exoplanets populate the same temperature range and share many physical properties, such as their molecule-dominated atmospheres and gradual cooling from ~3000 K at formation to ~100 K like the solar system gas-giant planets. Recent discoveries of very massive planets (Chauvin et al. 2005; Marois et al. 2010; Delorme et al. 2013), some possibly more massive than the 13 MJup deuterium burning mass limit, hint that planets could overlap with brown dwarfs in mass. On the other hand, the discovery of isolated L dwarfs in young clusters (Zapatero Osorio et al. 2002, 2014; Pe\u00f1a Ram\u00edrez et al. 2012), in young moving groups (Liu et al. 2013; Gagn\u00e9 et al. 2015; Gauza et al. 2015), and very cold very nearby Y dwarf objects (e.g., Kirkpatrick et al. 2012; Luhman 2014) show that very low-mass isolated brown dwarfs exist and overlap with the planetary masses. When these low-mass brown dwarfs are close enough and bright enough to be observed spectroscopically their atmospheres are much easier to study than similar exoplanets that lie near their very bright host stars. Liu et al. (2013) notably showed that the ~8 MJup brown dwarf PSO J318.5\u221222, a \u03b2-pictoris moving group member shares the spectral characteristics of the young directly imaged exoplanets, as well as atypically red late-L spectral type objects (e.g., Faherty et al. 2013; Gizis et al. 2015; Kellogg et al. 2016; Schneider et al. 2014, 2016; Bonnefoy et al. 2016). When CFBDSIR J214947.2\u2212040308.9, hereafter CFBDSIR 2149, was identified (Delorme et al. 2012), it seemed to be a candidate member of the AB Doradus young moving group and, together with the low-gravity features in its spectrum, made it a unique T-type isolated planetary-mass candidate. Another earlier-type, isolated young planetary-mass T-dwarf, SDSS J111010.01+011613.1, has been identified as a bona fide member of AB Doradus moving group (\\hbox{$149^{+51}_{-19}$}149-19+51 Myr; Bell et al. 2015) by Gagn\u00e9 et al. (2015). The late-T spectral type of CFBDSIR 2149 is typical of the coolest known directly imaged exoplanets, such as GJ 504 b or 51 Eri b (Kuzuhara et al. 2013; Macintosh et al. 2015), that the latest generation of adaptive optics systems are detecting. We therefore carried out a multi-wavelength, multi-instrument follow-up of CFBDSIR 2149 to fully characterise it and constrain its nature. ","Citation Text":["Ram\u00edrez et al. 2012"],"Citation Start End":[[601,620]]}
{"Identifier":"2018ApJ...860..153M__Yatabe_&_Yamada_2017_Instance_1","Paragraph":"FRBs have been detected in the radio band. Circular polarization of FRB 140514 was measured at a level of \n\n\n\n\n\n degree, while linear polarization is less than 10% as a 1\u03c3 upper limit (Petroff et al. 2015). We apply our model to self-consistently get both circular and linear polarization degrees of FRB 140514. The case of B = 1010 G, Te = 108 K, and \u03b3th = 104 shown as the solid line in the right panels of Figure 3 seems to be one solution to explain the polarization observations. Thus, we may consider FRB 140514 to be associated with strongly magnetized sources. A neutron star can be one candidate for FRB origin. Here, we assume that the polarized emission is pointing toward observers, as usual neutron star models assumed, but we do not include additional jet view-angle effects. We note one recent theoretical work on the polarization detection of neutron stars in the soft X-ray band (Yatabe & Yamada 2017), and the relatively low electron temperature from our analysis indicates that the radiation of FRB 140514 may not occur near the neutron star magneto-atmosphere above the neutron star surface. On the other hand, the alternative case of B = 1 G, Te = 1010 K, and \n\n\n\n\n\n in our model can obtain the circular polarization degree of 20%. However, in this case, the intrinsic linear polarization degree is about 76%. This seems to be in contrast with the observation of FRB 140514. Therefore, we are inclined to believe from our polarization analysis that FRB may have a neutron star origin. Besides FRB 140514, FRB 110523 has the linear polarization degree of \n\n\n\n\n\n, but the circular polarization degree of about 23% due to instrumental leakage was not confirmed to be an astronomical result (Masui et al. 2015). FRB 081507 has the high linear polarization degree of 80% (Ravi et al. 2016) and FRB 150215 has the linear polarization degree of \n\n\n\n\n\n (Petroff et al. 2017). If we can detect both linear and circular polarization for FRBs, we may further constrain the FRB physics.","Citation Text":["Yatabe & Yamada 2017"],"Citation Start End":[[897,917]]}
{"Identifier":"2022ApJ...930..125L__Zimbardo_&_Perri_2013_Instance_1","Paragraph":"To this we can add that Zimbardo and coworkers reported in a series of papers that spacecraft observations of SEP intensity-time profiles ahead of interplanetary shocks yielded power laws instead of the exponential profiles predicted by standard DSA. They interpreted this as a sign of superdiffusive spatial SEP transport in the shock vicinity resulting in superdiffusive shock acceleration. They predicted the superdiffusive shock acceleration to be more efficient, producing a harder accelerated power-law spectrum for SEPs, than standard DSA (Zimbardo et al. 2006; Perri & Zimbardo 2007, 2008, 2009; Zimbardo & Perri 2013; Zimbardo et al. 2017, 2020; Zimbardo & Perri 2020). Zimbardo & Perri (2020) explained such superdiffusive transport as a product of subdiffusion in pitch-angle space arising from the observation that Alfv\u00e9n-wave pitch-angle scattering times in the solar wind have a power-law distribution instead of a single value. Another possibility is that there is a strong presence of SMFRs in the vicinity of heliospheric shocks so that particle trapping effects could result in anomalous diffusive transport and shock acceleration. Evidence is increasing that, in a space plasma with a significant guide\/background magnetic field like the solar wind, SMFRs naturally form as part of a self-generated, quasi-two-dimensional (quasi-2D) MHD turbulence component that might dominate other MHD wave turbulence modes. This is supported by observations in the slow solar wind near 1 au (e.g., Matthaeus et al. 1990; Bieber et al. 1996; Greco et al. 2009; Zheng & Hu 2018), MHD simulations (e.g., Shebalin et al. 1983; Dmitruk et al. 2004), and nearly incompressible MHD turbulence transport theory (e.g., Zank & Matthaeus 1992, 1993; Zank et al. 2017, 2018, 2020). It is plausible that SMFRs can have a strong presence ahead of interplanetary shocks. Furthermore, there is evidence that the occurrence of SMFRs peaks at large-scale current sheets in the solar wind where magnetic reconnection produces additional SMFRs that can trap energetic particles (Khabarova et al. 2015, 2016; Hu et al. 2018). Many of these current sheets occur behind traveling shocks (Khabarova & Zank 2017). Recently, a theoretical investigation was launched to determine the transmission of quasi-2D SMFRs through a perpendicular shock (Zank et al. 2021). The results show a strong enhancement in the magnetic energy density of the SMFR magnetic island component behind the shock, thus confirming a strong presence of SMFRs downstream of the shock as well.","Citation Text":["Zimbardo & Perri 2013"],"Citation Start End":[[604,625]]}
{"Identifier":"2021AandA...655A...1M__Laskar_1999_Instance_1","Paragraph":"Analytical insight into the motion of the inner planets requires an appropriate dynamical modelling of the long-term evolution of their orbits. On the one hand, such a model must be consistent with the predictions of the reference numerical integrations available in literature (Laskar 1990c, 2008; Laskar et al. 2004; Laskar & Gastineau 2009) to ensure that it reproduces the dynamical features of the inner system with sufficient precision. On the other hand, the corresponding Hamiltonian should be set in a form suitable for the systematic application of canonical perturbation techniques (Hori 1966; Deprit 1969), which is essential to an unbiased analysis of such a high-dimensional dynamics. Moreover, the possibility of numerically integrating the equations of motion in an efficient way is fundamental for studying the chaotic evolution of the orbital solutions in a statistical way. Unfortunately, the construction of such a model is a delicate task. In principle, we might just consider the full N-body Hamiltonian of the Newtonian gravitational interactions among the Solar System planets, with the addition of the leading corrections coming from general relativity and the Earth\u2013Moon interaction. This already reproduces the precession frequencies of the inner orbits with a precision better than 0.01\u2032\u2032 yr\u22121 (Laskar 1999). However, this Hamiltonian is unnecessarily complicated because it includes short-time harmonics, with periods shorter than 5000 yr (e.g. Carpino et al. 1987), which are known to generate small quasi-periodic oscillations in the inner orbits, without being implied in chaos generation. The inner planets are not, indeed, involved in any relevant mean-motion resonance. At the same time, the long-term numerical integration of the corresponding equations of motion is very time consuming, the resulting solutions needing to be filtered to extract the secular trend of the orbits (Carpino et al. 1987; Nobili et al. 1989). These facts suggest considering a secular Hamiltonian to directly describe the slow movement of the planet perihelia and nodes after proper averaging over the short-time orbital motion (Laskar 1984, 1985). Secular dynamics includes the essential planet interactions responsible for chaos in the inner Solar System and allows performing the fastest long-term numerical integrations (Laskar 1988, 1989, 1994, 2008). This enabled the discovery of chaos in the inner system before the use of symplectic integration schemes (Sussman & Wisdom 1992). Unfortunately, an effective secular model for the entire Solar System has to be of high order in planet masses, principally because of the 5:2 near mean-motion resonance between Jupiter and Saturn, the so-called great inequality (Laplace 1785; Laskar 1996). A simple averaging of the N-body Hamiltonian over the planet mean longitudes, resulting in a first-order secular dynamics in planet masses, would reproduce the fundamental frequencies g5 and g6 (which dominate the perihelion precession of Jupiter and Saturn, respectively) very poorly1 (Laskar 1988). The construction of higher-order models requires the manipulation of large Poisson series and the use of sophisticated computer algebra systems. This is probably the reason why they are still not widely used, at least as a basis of extensive research.","Citation Text":["Laskar 1999"],"Citation Start End":[[1323,1334]]}
{"Identifier":"2021ApJ...919...30D__Staguhn_et_al._2014_Instance_3","Paragraph":"The first SMGs were detected using SCUBA at 850 \u03bcm (Smail et al. 1997; Barger et al. 1998; Hughes et al. 1998), which remains one of the prime wavelengths to detect these galaxies (e.g., Geach et al. 2017), thanks to a combination of available instruments, spectral window, and the negative k-correction at that wavelength. Other single-dish samples of SMGs have also been obtained at 1.1\u20131.3 mm using MAMBO (e.g., Eales et al. 2003; Bertoldi et al. 2007; Greve et al. 2008) and AzTEC (e.g., Aretxaga et al. 2011; Yun et al. 2012), at 1.4 mm\/2 mm with the SPT (Vieira et al. 2010), and at 2 mm with GISMO (Staguhn et al. 2014; Magnelli et al. 2019). Selecting SMGs from observations at longer wavelengths is thought to favor galaxies at higher redshifts (e.g., Smol\u010di\u0107 et al. 2012; Vieira et al. 2013; Staguhn et al. 2014; Magnelli et al. 2019; Hodge & da Cunha 2020), although it is difficult to compare the redshift distributions in an unbiased way (see, e.g., Zavala et al. 2014 for a discussion), and account for intrinsic variations of galaxy far-IR spectral energy distributions (SEDs). Nevertheless, the 2 mm band has been put forth as a potential candidate to detect high-redshift (z > 3) galaxies (e.g., Casey et al. 2018a, 2018b, 2019; Zavala et al. 2021). The negative k-correction is stronger at 2 mm than at 850 \u03bcm; thus, for a fixed SED, the 2 mm band should pick up more high-redshift galaxies than at 870 \u03bcm. In addition, better atmospheric transmission and larger fields of view can be achieved at 2 mm (but corresponding poorer resolution). Such an effort is currently ongoing (see Zavala et al. 2021 for first results). To understand the relationship between the populations detected at 850 \u03bcm and at 2 mm, we require a detailed characterization of the (sub)millimeter SEDs of these sources. Multiwavelength submillimeter observations are still rare, with most observations focusing on a single wavelength. Only a handful of sources observed at 2 mm have complementary shorter-wavelength detections (Staguhn et al. 2014; Magnelli et al. 2019). Thus, a more systematic multiwavelength dust continuum investigation is warranted in order to reveal the dust properties of (sub)millimeter-detected sources.","Citation Text":["Staguhn et al. 2014"],"Citation Start End":[[2019,2038]]}
{"Identifier":"2022AandA...668A..50S__Ginski_et_al._(2016)_Instance_1","Paragraph":"To test IADI, we made a model of a disk in scattered light. This model is 400 \u00d7 400 pixels, each pixel having a physical size of 1 \u00d7 1 au. Multiple rings are implemented, inspired by observations of disks (e.g., HD 97048 Ginski et al. 2016; RX J1615 de Boer et al. 2016 and TW Hydrae van Boekel et al. 2017), see panel a in Fig. 2. A specific inclination and rotation is achieved using the warpAffine function from OpenCV (Bradski 2000). Young planet-forming disks are still gas rich and dust particles are stratified due to gas pressure along the vertical axis. Therefore, flaring is implemented in the model via an offset of the rings with respect to the center of the ring (see the detailed discussion in de Boer et al. 2016). For our disk model, we are using the power-law profile for the scattering surface height H and the separation r found by Ginski et al. (2016) for the disk around HD 97048. This profile describes the flaring discovered in this source reasonably well up to a separation of ~270 au. Considering that the model disk used will have a separation of 200 au, this formula will be sufficient to simulate the disk height2. Illumination effects of the central star are implemented via a ~1\/r2 intensity dependence from the center of the disk, making the inner part of the disk brighter compared to the outer part. Moreover, the intensity also depends on the light scattering angle via the phase function. Because a physical model is beyond the scope of this work, a \u201cpseudo\u201d phase function is implemented to mimic the same asymmetries in light distribution seen in observed disks via I = cos \u03d5, where the intensity I depends on the cosine of the azimuthal angle \u03d5. This pseudo phase function depends on the azimuthal angle instead of a scattering angle on which a real phase function would depend. This makes the part of the disk facing toward the observer appear brighter than the part facing away in a fairly simple way. Lastly, the model is put through a Gaussian convolution kernel from the scipy ndimage package (Virtanen et al. 2020) to remove sharp edges and give a finite resolution to the model. The final three steps are shown in the top row of Fig. 2.","Citation Text":["Ginski et al. 2016"],"Citation Start End":[[221,239]]}
{"Identifier":"2022AandARv..30....6M__Magliocchetti_et_al._2018b_Instance_1","Paragraph":"AGN While\u2014starting from the very early works (e.g., Seldner and Peebles 1978; Longair and Seldner 1979; Hill and Lilly 1991; Peacock and Nicholson 1991; Allington-Smith et al. 1993; Zirbel 1997)\u2014virtually all the studies presented in the literature converge at indicating that radio-AGN preferentially reside within overdense structures, the three (or rather five if one also includes cross-correlation studies) methods presented in Sect. 4 provide different information on the large-scale structure behaviour of these objects. To summarize them in a few words, we might say that the method based on clustering returns more information on the very large\/cosmological (i.e., at Mpc level and beyond) scales traced by radio sources and also on the dark matter content of the regions that host them. On the other hand, very different results are obtained if one searches for structures around known radio-AGN or if one pinpoints radio-AGN within known structures. Indeed, in the first case one finds that virtually all radio-AGN are surrounded by over-densities (e.g., Venemans et al. 2007; Mayo et al. 2012; Galametz et al. 2012; Castignani et al. 2014; Wylezalek et al. 2013; Rigby et al. 2014), while the second method shows that only about 20\u201330% of them inhabit rich (group- and cluster-like) structures (e.g., Best et al. 2007; Magliocchetti and Br\u00fcggen 2007; Lin and Mohr 2007; Croft et al. 2007; Magliocchetti et al. 2018b; Croston et al. 2019). The reason for this discrepancy is not known, but it is likely related to the different redshift ranges probed by the two methods (much more local sources are considered in the second one), and\/or\u2014under the assumption of a strong correlation between radio luminosity and environmental density (e.g., Bardelli et al. 2010; Magliocchetti et al. 2018b; Croston et al. 2019; Mo et al. 2020, but see further in this section for different points of view)\u2014to the fact that generally the first method images much brighter radio sources than the second one. In any case, we note that these findings also have implications for the life-time of the radio-AGN phenomenon, and it is therefore of no surprise if works based on the different methods illustrated above find different values, ranging from \u223c60\\documentclass[12pt]{minimal}\n\t\t\t\t\\usepackage{amsmath}\n\t\t\t\t\\usepackage{wasysym}\n\t\t\t\t\\usepackage{amsfonts}\n\t\t\t\t\\usepackage{amssymb}\n\t\t\t\t\\usepackage{amsbsy}\n\t\t\t\t\\usepackage{mathrsfs}\n\t\t\t\t\\usepackage{upgreek}\n\t\t\t\t\\setlength{\\oddsidemargin}{-69pt}\n\t\t\t\t\\begin{document}$$\\sim 60$$\\end{document} Myr up to a few Gyr (e.g., Lin and Mohr 2007; Smol\u010di\u0107 et al. 2011; Hatch et al. 2014; Magliocchetti et al. 2017). We also note that recent studies based on direct LOFAR observations which\u2014we remind\u2014sample lower frequencies and therefore older emission, would tend to better agree with the high-end values provided above for the radio-active phase of an AGN (\u03c4>200\\documentclass[12pt]{minimal}\n\t\t\t\t\\usepackage{amsmath}\n\t\t\t\t\\usepackage{wasysym}\n\t\t\t\t\\usepackage{amsfonts}\n\t\t\t\t\\usepackage{amssymb}\n\t\t\t\t\\usepackage{amsbsy}\n\t\t\t\t\\usepackage{mathrsfs}\n\t\t\t\t\\usepackage{upgreek}\n\t\t\t\t\\setlength{\\oddsidemargin}{-69pt}\n\t\t\t\t\\begin{document}$$\\tau > 200$$\\end{document} Myr \u2013 Heesen et al. 2018).","Citation Text":["Magliocchetti et al. 2018b"],"Citation Start End":[[1401,1427]]}
{"Identifier":"2020MNRAS.497.5413S__Bisschop_et_al._2007_Instance_1","Paragraph":"In Section 3.2, it was noted that the long range structure observed in the reaction between the formamide and amino radicals was that of neutral hydrogenated isocyanic acid. However, instead of hydrogenation, protonation could be occurring. Isocyanic acid (HNCO) has been detected in the ISM and can be formed from simple building blocks like atomic nitrogen and CO (Snyder & Buhl 1972; Buhl, Snyder & Edrich 1973; Nourry, Zins & Krim 2015). It has been used in the study of the formation of urea in irradiation studies, where its irradiation is shown to lead to the formation of formamide and urea (Ferus et al. 2018; Raunier et al. 2003, 2004). In addition, it has been shown to be a product of the dissociation of formamide and has even been suggested to be a catalyst for the formation of H2 in the ISM (Lundell, Krajewska & R\u00e4s\u00e4nen 1998; Duvernay et al. 2005; Haupa, Tarczay & Lee 2019). Moreover, there is a clear correlation with the abundances of formamide (Bisschop et al. 2007; L\u00f3pez-Sepulcre et al. 2015). There is ample evidence of HNCO molecules functioning as an imine base, both in the gas phase and in the solid state (Dekock & Jasperse 1983; Hop et al. 1989; Hunter & Lias 1997; Hudson, Khanna & Moore 2005; Gupta et al. 2013; Bouchoux 2018; Marcelino et al. 2018). The proton affinity of HNCO is higher than that of water, so that one would expect protonation to occur even in water-based ices (Dekock & Jasperse 1983; Bouchoux 2018). Thus, in this section, calculations starting from protonated isocyanic acid are reported. Only DFT calculations were performed on these reactions, since multiconfigurational effects can be assumed to be small for these non-radical charged species. The first step that was investigated is the protonation of isocyanic acid. There are multiple options for protonation in this case: Protonation can occur on the nitrogen, carbon, or oxygen atoms. Hereby, we note that N-protonated isocyanic acid has been detected in laboratory experiments and in the ISM in the dense core L483 with a more tentative detection towards Sgr B2(N) (Hudson et al. 2005; Gupta et al. 2013; Marcelino et al. 2018). Moreover, O-protonated isocyanic acid has also been detected in laboratory experiments (Lattanzi et al. 2012). The potential energy curves for protonating on each of these atoms is shown in Fig. 7. Protonation on the nitrogen and oxygen atoms is barrierless. The well depths are 603.1 and 603.9 kJ mol\u22121, respectively, making them equally likely. Protonation on the carbon atom, however, is not barrierless, but requires an energy input of 14.4 kJ mol\u22121. It also results in a much less stable species, with a well depth of only 172.8 kJ mol\u22121, making it the least likely of the three pathways. Nitrogen protonation is the product that most naturally leads on to the formation of urea as it creates one of the amino groups found in urea. None the less, the similarity in energies for the nitrogen- and oxygen-protonated species means that the oxygen-protonated species could also be formed. The latter is the species we will consider first.","Citation Text":["Bisschop et al. 2007"],"Citation Start End":[[966,986]]}
{"Identifier":"2017AandA...601A..87C__Falcke_(1996)_Instance_3","Paragraph":"In a quasi-isothermal jet, Uj is (17)\\begin{equation} \\label{eq:U_j_quasi} U_{\\rm j} = \\zeta n_0 m_{\\rm p} c^2\\left(\\frac{\\gamma_{\\rm j}\\beta_{\\rm j}}{\\gamma_0\\beta_0}\\right)^{-\\Gamma}\\left(\\frac{z}{z_0}\\right)^{-2} \\cdot \\end{equation}Uj=\u03b6n0mpc2\u03b3j\u03b2j\u03b30\u03b20\u2212\u0393zz0-2\u00b7Substituting Eqs. (17) and (13) into Eq. (10), and assuming the jet is launched with an initial \u03b30\u03b20 equal to the sound speed (Eq. (16)), the 1D Euler equation that results is \\begin{eqnarray} \\label{eq:AGNJET_Corrected} &&\\left\\{\\gamma_{\\rm j}\\beta_{\\rm j}\\frac{\\Gamma+\\xi}{\\Gamma-1}-\\Gamma\\gamma_{\\rm j}\\beta_{\\rm j}-\\frac{\\Gamma}{\\gamma_{\\rm j}\\beta_{\\rm j}}\\right\\}\\frac{\\partial \\gamma_{\\rm j}\\beta_{\\rm j}}{\\partial z} = \\frac{2}{z}; \\\\ &&\\xi = \\frac{1}{\\zeta}\\left(\\frac{\\gamma_{\\rm j}\\beta_{\\rm j}}{\\gamma_0\\beta_0}\\right)^{\\Gamma-1}; \\qquad \\gamma_0\\beta_0=\\sqrt{\\frac{\\zeta\\Gamma(\\Gamma-1)}{1+2\\zeta\\Gamma-\\zeta\\Gamma^2}} \\cdot \\end{eqnarray}\u03b3j\u03b2j\u0393+\u03be\u0393\u22121\u2212\u0393\u03b3j\u03b2j\u2212\u0393\u03b3j\u03b2j\u2202\u03b3j\u03b2j\u2202z=2z;\u03be=1\u03b6\u03b3j\u03b2j\u03b30\u03b20\u0393\u22121;\u2001\u03b30\u03b20=\u03b6\u0393(\u0393\u22121)1+2\u03b6\u0393\u2212\u03b6\u03932\u00b7The above equation should reduce to the jet Lorentz factor profile used in Falcke (1996), Markoff et al. (2005) when \u03b6 = 1. However, it differs from Eq. (2) in Falcke (1996): (20)\\begin{equation} \\label{eq:Heino96} \\left\\{\\gamma_{\\rm j}\\beta_{\\rm j}\\frac{\\Gamma+\\xi}{\\Gamma-1}-\\frac{\\Gamma}{\\gamma_{\\rm j}\\beta_{\\rm j}}\\right\\}\\frac{\\partial \\gamma_{\\rm j}\\beta_{\\rm j}}{\\partial z} = \\frac{2}{z}; \\end{equation}\u03b3j\u03b2j\u0393+\u03be\u0393\u22121\u2212\u0393\u03b3j\u03b2j\u2202\u03b3j\u03b2j\u2202z=2z;(21)\\begin{equation} \\xi = \\left(\\gamma_{\\rm j}\\beta_{\\rm j}\\frac{\\Gamma+1}{\\Gamma(\\Gamma-1)}\\right)^{1-\\Gamma} \\cdot \\end{equation}\u03be=\u03b3j\u03b2j\u0393+1\u0393(\u0393\u22121)1\u2212\u0393\u00b7The difference between our equation and the equation in Falcke (1996) can be accounted for as follows: the \u2212 \u0393\u03b3j\u03b2j term in Eq. (18) results from a neglected \\hbox{$\\frac{\\partial}{\\partial z}(U_{\\rm j}\/n)$}\u2202\u2202z(Uj\/n) term, the difference in the exponent in \u03be results from an arithmetic error, and finally the difference in the inside of the parenthesis of \u03be terms is from setting \\hbox{$\\gamma_0\\beta_0 = \\beta_{\\rm s0}^{2}$}\u03b30\u03b20=\u03b2s02 instead of using the proper value given in Eq. (16). The difference between the solutions of Eqs. (18) and (20) are small and shown in Fig. 1. In Fig. 1, we also include solutions to the 1D Euler equations when the jet is isothermal (Tj = const., i.e., Eq. 20 with \u03be = 1) and adiabatic (Tj \u221d (\u03b3j\u03b2j)1 \u2212 \u0393z2 \u2212 2\u0393, see Eq. (25)). ","Citation Text":["Falcke (1996)"],"Citation Start End":[[1648,1661]]}
{"Identifier":"2015AandA...574A..12L__Margutti_et_al._2014_Instance_1","Paragraph":"It has been also suggested that the X-ray emission from SNe is dominated by inverse Compton (IC) scattering of photospheric optical photons by relativistic electrons accelerated by the SN shock on a timescale of weeks to a month after the explosion (see Chevalier & Fransson 2006). The X-ray luminosity strongly depends on the structure of the SN ejecta (\u03c1SN \u221d r\u2212 n), the density structure of the CSM and the relativistic electron distribution responsible for the up-scattering (Chevalier & Fransson 2006; Margutti et al. 2012). Assuming a wind-like CSM (\u03c1CSM \u221d r-2) and an ISM-like CSM (\u03c1CSM = constant), Margutti et al. (2012) developed two generalized formalisms to calculate the IC luminosity. Their formalisms are strongly sensitive to the SN bolometric luminosity, Lbol (see Appendix A of Margutti et al. 2012). Here, it is assumed electrons are accelerated according to a power-law distribution n(\u03b3) \u221d \u03b3\u2212 p with index p = 3, as suggested by the radio observations of SN shocks in Type Ib\/Ic explosions (see Soderberg et al. 2006; Margutti et al. 2014). Then, an IC luminosity in the wind-like CSM case is (see Margutti et al. 2012, their Appendix A) (1)\\begin{eqnarray} \\label{eq:1}  \\frac{\\mathrm{d} L_{\\rm{IC}}}{\\mathrm{d} \\nu} \\approx 2.1\\times10^{-2}\\left(\\frac{\\epsilon_{\\rm e}}{0.1}\\right)^{-2}\\left(\\frac{M_{\\rm{ej}}}{1.4\\,M_{\\odot}}\\right)^{-0.93}\\left(\\frac{A}{\\rm{g\\,cm^{-1}}}\\right)^{0.64}  \\nonumber \\\\\\times\\left(\\frac{E}{10^{51}\\,\\rm{erg}}\\right)^{1.29}\\left(\\frac{t}{\\rm{s}}\\right)^{-1.36}\\left(\\frac{L_{\\rm{bol}}}{\\rm{erg\\,s^{-1}}}\\right)\\,\\nu^{-1}, \\end{eqnarray}dLICd\u03bd\u22482.1\u00d710-2\u03f5e0.1-2Mej1.4\u2009M\u2299-0.93Ag\u2009cm-10.64\u00d7E1051\u2009erg1.29ts-1.36Lbolerg\u2009s-1\u2009\u03bd-1,where A = \u1e40\/ (4\u03c0\u03c5wind), and \u03f5e is the fraction of thermal energy in the shock used for the electron acceleration (~0.1 see Chevalier & Fransson 2006). Theoretical IC X-ray luminosities with different mass-loss rates (for a wind velocity of 100 km s-1) are compared to the observations in Fig. 5. Here, both the observed bolometric light curve of SN 2005hk and the predicted bolometric light curve for the \u201cN5def model\u201d (see Fig. 8 of Kromer et al. 2013) are used as input for Lbol. The \u201cN5def model\u201d was shown to provide a good fit to the observations of the SN 2005hk event (Kromer et al. 2013). Moreover, E = 1.34 \u00d7 1050 erg, and Mej = 0.37 M\u2299 (these parameters are consistent with the N5def model) are used in Eq. (1) to estimate the X-ray luminosities. Here, we note that SNe Iax have a range of luminosities, ejecta velocities, and inferred ejecta masses (Foley et al. 2013). Different bolometric luminosity light curves and ejecta masses of different SNe Iax may change the results shown in Fig. 5. ","Citation Text":["Margutti et al. 2014"],"Citation Start End":[[1037,1057]]}
{"Identifier":"2020AandA...635A.121M__Matter_et_al._(2016)_Instance_2","Paragraph":"Our model consists of four zones; an inner disk (zone 1, as in Matter et al. 2016), and three zones into which the outer disk is divided (zones 2\u20134) in order to produce the azimuthal asymmetries seen in our SPHERE observations. The radial extent of these zones were constrained from our images: zone 2 corresponds to the component inside ~14 au observed in the noncoronagraphic J-band image, zone 3 corresponds to the bright ring at 16 au, and zone 4 corresponds to the asymmetrical outer disk. The radial extent of each of these zones, as well as their disk masses, is summarized in Table 2. No gaps were introduced between the three outer disk zones, but the gap between zones 1 and 2 from the Matter et al. (2016) model was kept. We note that these radii are only loosely constrained from our data, and are determined by eye from both the coronographic and noncoronographic data. We ran a grid of models with different inclinations and position angles for zones 1\u20132 (note that we systematically use the same inclination and position angle for these two zones) and for zone 3 \u2013 that is, a total offour free parameters, keeping in mind that the relative inclination between components must be small to allow a single broad shadow to be cast (as opposed to two narrow shadow lanes). The inclination and position angle of zone 4 is set to the values provided in Sect. 3. The grid was sampled in steps of 1\u00b0 for inclination, and 2\u00b0 for PA, and later refined to 0.5\u00b0 steps for inclination and 0.5\u00b0 for PA after a good initial agreement is found between the average azimuthal profile of the model and the scattered light images. The best fitting model was picked not only based on the location of the shadows cast by the inner components on the outer disk, but also on the shape (slope) of the resulting azimuthal profile of the outer disk. There is also a degeneracy if we consider that the inclinations and PAs of the two inner components (zones 1\u20132 and zone 3) can be exchanged and produce very similar results. However, doing this would cast a shadow in a different location, and thus produce a different azimuthal profile for zone 3.","Citation Text":["Matter et al. (2016)"],"Citation Start End":[[696,716]]}
{"Identifier":"2017ApJ...850L..27V__Tanvir_et_al._2017a_Instance_1","Paragraph":"After the LVC announcement, a multi-wavelength campaign immediately started. The campaign involved the X-ray and gamma-ray satellites first, and radio\/IR\u2013optical\/TeV ground observatories at later times. The detection of a new optical transient (OT) was first announced by the One-meter, Two-hemisphere (1M2H) team discovered with the 1 m Swope telescope on August 18 01:05 UT (Coulter et al. 2017a, 2017b; Drout et al. 2017), named Swope Supernova Survey 2017a (SSS17a, now with the IAU designation AT 2017gfo). The OT was also detected independently by five other teams, the Dark Energy Camera (Allam et al. 2017; Cowperthwaite et al. 2017), the Distance Less Than 40 Mpc Survey (L. Tartaglia et al. 2017, in preparation; Valenti et al. 2017; Yang et al. 2017), Las Cumbres Observatory (Arcavi et al. 2017a, 2017b, 2017c), the Visible and Infrared Survey Telescope for Astronomy (Tanvir et al. 2017a, 2017b), and MASTER (Lipunov et al. 2017b, 2017a), REM-ROS2 (Melandri et al. 2017; Pian et al. 2017a, 2017b), Swift UVOT\/XRT (Evans et al. 2017b, 2017a), and Gemini-South (Kasliwal et al. 2017; Singer et al. 2017), and for all see also MMA17. AT 2017gfo is located at 106 from the early-type galaxy NGC 4993, at a distance of \u223c40 Mpc. A sequence of satellite high-energy observations started almost immediately, with exposures of the GW170817 LR depending on satellite position and operations. The X-ray and \u03b3-ray observations included contributions by CALET (Nakahira et al. 2017), Konus-Wind (Svinkin et al. 2017), Insight-HXMT (Li et al. 2017; Liao et al. 2017), AstroSat CZTI (Balasubramanian et al. 2017), AGILE-GRID (see below), Fermi-LAT (Kocevski et al. 2017), MAXI (Sugita et al. 2017; S. Sugita et al. 2017, in preparation), Super-AGILE (MMA17; this work), Swift X-ray Telescope (Evans et al. 2017b, 2017a), NuSTAR (Harrison et al. 2017), INTEGRAL JEM-X (Savchenko et al. 2017b; V. Savchenko et al. 2017, in preparation), and Chandra (Fong et al. 2017; Margutti et al. 2017a; Troja et al. 2017b). An X-ray counterpart detection at the OT position was reported after 9 days and confirmed after 15 days by Chandra (Fong et al. 2017; Margutti et al. 2017b; Troja et al. 2017a, 2017b). Moreover, an important detection in the radio band has been reported by VLA (Alexander et al. 2017a, 2017b).","Citation Text":["Tanvir et al. 2017a"],"Citation Start End":[[881,900]]}
{"Identifier":"2015ApJ...799..149J___2014_Instance_5","Paragraph":"With our joint analysis of stellar mass fraction and source size, we find a larger stellar mass fraction than earlier statistical studies. In Figure\u00c2 2, we compare our determination of the stellar surface density fraction to a simple theoretical model and to the best fit of a sample of lens galaxies by Oguri et\u00c2 al. (2014). The simple theoretical model is the early-type galaxy equivalent of a maximal disk model for spirals. We follow the rotation curve of a de Vaucouleurs component for the stars outward in radius until it reaches its maximum and then simply extend it as a flat rotation curve to become a singular isothermal sphere (SIS) at large radius (see details in the Appendix). The ratio of the surface mass density of the de Vaucouleurs component to the total surface mass density is shown as a dashed curve in Figure\u00c2 2. We also show as a gray band the best fit for the stellar fraction in the form of stars determined by Oguri et al (2014) in a study of a large sample of lens galaxies using strong lensing and photometry, as well as the best model using a Hernquist component for the stars and an NFW halo for the dark matter with and without adiabatic contraction, also from Oguri et\u00c2 al. (2014). We have used the average and dispersion estimates for the Einstein and effective radii available for 13 of the objects in our sample from Oguri et\u00c2 al. (2014), Sluse et\u00c2 al. (2012), Fadely et\u00c2 al. (2010), and Leh\u00c3\u00a1r et\u00c2 al. (2000; see Table\u00c2 1) as an estimate of RE\/Reff in Figure\u00c2 2. The average value and dispersion of the sample is RE\/Reff = 1.8 \u00c2\u00b1 0.8. This also averages over the different radii of the lensed images. The agreement of our estimates with the expectations of the simple theoretical model and with estimates from other studies (Oguri et\u00c2 al. 2014) is quite good. For comparison, the estimate of Pooley et\u00c2 al. (2012; using the Einstein and effective radii estimates for 10 out of 14 of their objects from Schechter et\u00c2 al. 2014) seems somewhat lower than expected at those radii. The range of stellar mass fractions from MED09 for source sizes in the range 0.3\u00e2\u0080\u009315.6 light days is also shown in Figure\u00c2 2. In this case, the discrepancy between our estimate and their reported value of \u00ce\u00b1 = 0.05 is completely due to the effect of the source size. Although accretion disk sizes are known to be smaller in X-rays, recent estimates are in the range of 0.1\u00e2\u0080\u00931 light-days, depending on the mass of the black hole (see Mosquera et\u00c2 al. 2013), and these finite sizes will increase the stellar surface densities implied by the X-ray data. Another possible origin for this discrepancy is that Pooley et\u00c2 al. (2012) use the macro model as an unmicrolensed baseline for their analysis. It is well known that simple macro models are good at reproducing the positions of images, but have difficulty reproducing the flux ratios of images due to a range of effects beyond microlensing. Recently, Schechter et\u00c2 al. (2014) found that the fundamental plane stellar mass densities have to be scaled up by a factor 1.23 in order to be compatible with microlensing in X-rays in a sample of lenses with a large overlap with that analyzed by Pooley et\u00c2 al. (2012). It is unclear how this need for more mass in stars at the position of the images found by Schechter et\u00c2 al. (2014) can be reconciled with the apparently low estimate of mass in stars at those radii by Pooley et\u00c2 al. (2012). Our estimate of the stellar mass fraction agrees better with the results of microlensing studies of individual lenses (Keeton et\u00c2 al. 2006; Kochanek et\u00c2 al. 2006; Morgan et\u00c2 al. 2008, 2012; Chartas et\u00c2 al. 2009; Pooley et\u00c2 al. 2009; Dai et\u00c2 al. 2010) that reported values in the range 8%\u00e2\u0080\u009325%, and with the estimates from strong lensing studies (see for example Jiang & Kochanek 2007; Gavazzi et\u00c2 al. 2007; Treu 2010; Auger et\u00c2 al. 2010; Treu et\u00c2 al. 2010; Leier et\u00c2 al. 2011; Oguri et\u00c2 al. 2014) which produced stellar mass fractions in the range 30%\u00e2\u0080\u009370% integrated inside the Einstein radius of the lenses.","Citation Text":["Oguri et\u00c2 al. 2014"],"Citation Start End":[[1762,1780]]}
{"Identifier":"2015ApJ...808...98L__Lagage_et_al._1996_Instance_1","Paragraph":"Cassiopeia A (Cas A), which is a remnant of an SN IIb explosion (Krause et al. 2008), is an ideal target for examining the properties of SN dust. It is quite young (\u223c330 years; Fesen et al. 2006), so the SN ejecta and freshly formed SN dust have not been significantly mixed with the circumstellar\/interstellar medium (CSM\/ISM), and it is relatively close (3.4 kpc; Reed et al. 1995), so the physical and chemical structures can be resolved. Therefore, Cas A has been a major target of infrared space missions. The first direct evidence for SN dust in Cas A came from the Infrared Space Observatory (ISO), which performed mid-infrared (MIR; 2.4\u201345 \u03bcm) spectroscopic observations toward the bright ejecta shell (Lagage et al. 1996; Arendt et al. 1999; Douvion et al. 2001). The continuum spectra had a strong bump peaking at 21 \u03bcm together with a relatively weak bump at 9.5 \u03bcm, and Douvion et al. (2001) showed that the spectra can be fitted by two dust components at different temperatures, one at 90 K and the other at 350 K, with pyroxene (MgSiO3), quartz (SiO2), and aluminum oxide (Al2O3) as major components. The composition and distribution of this \u201cwarm\u201d SN dust heated by a reverse shock have been studied in detail using the spectral mapping data of the Spitzer Space Telescope (Ennis et al. 2006; Rho et al. 2008; Arendt et al. 2014). These studies showed that the dust emission exhibits distinct spectral characteristics depending on the SN material with which it is associated. The most prominent dust emission is that with strong 9 and 21 \u03bcm bumps, which is associated with SN ejecta having strong Ar emission lines. The spectral features agree with the ISO spectra and can be reproduced by Mg protosilicate\/MgSiO3, Mg\n\n\n\n\n\nSiO\n\n\n\n\n\n, or SiO2. The dust emission associated with SN ejecta having strong Ne lines is smooth without any silicate features and can be reproduced by Al2O3 dust or C glass. There is yet another type of smooth dust emission associated with X-ray Fe emission, but it is probably from swept-up CSM (Arendt et al. 2014). Meanwhile, Nozawa et al. (2010) showed that the observed spectral energy distribution (SED) of dust emission can be fitted by a physical model in which the SNR is expanding into a dense CSM, and the MIR emission is mostly from silicate grains (e.g., MgSiO3, Mg2SiO4, and SiO2) and MgO. The estimated total mass of warm dust ranges from 0.008 to 0.054 \n\n\n\n\n\n.","Citation Text":["Lagage et al. 1996"],"Citation Start End":[[711,729]]}
{"Identifier":"2017ApJ...840...17T__Iglesias_&_Rogers_1996_Instance_1","Paragraph":"These data were compared to two different grids of stellar evolution models (see Table 1 for a summary). The first was a grid run using the Yale Rotating Evolution Code (YREC, Pinsonneault et al. 1989, with updates as discussed in van Saders & Pinsonneault 2012). We ran a grid of masses (0.6 M\u2299\u20132.6 M\u2299 in 0.1 M\u2299 increments), metallicities ([Fe\/H] = \u22122.0 to +0.6 in steps of 0.2), and \u03b1-element enhancements (+0.0, +0.2, and +0.4) to generate tracks. These models use a Grevesse & Sauval (1998) elemental mixture, with the updated OPAL equation of state (Rogers et al. 1996; Rogers & Nayfonov 2002), OPAL opacity tables (Iglesias & Rogers 1996) and gray atmospheres. Alpha-enhanced models use alpha-enhanced starting models and opacity tables, but not equations of state. Our models include semiconvection (see Kippenhahn & Weigert 1994), but do not include overshoot or the diffusion of helium or heavy elements. We assumed a Big Bang helium value of Y = 0.2485 at Z = 0 (Cyburt 2004), and did a linear fit to our solar value of Y = 0.272683 at Z = 0.018011 (so \n\n\n\n\n\n). We convert Z to metallicity assuming \n\n\n\n\n\n for a solar mixture, with a correction factor for alpha-enhanced models. We compare to other theoretical models in Section 3. In the initial part of this analysis, we use grids with a solar mixing length (B\u00f6hm-Vitense 1958) of 1.72, chosen to reproduce the solar luminosity (3.827 \u00d7 1033 erg s\u22121) and the solar radius (6.957 \u00d7 1010 cm) at the solar age (4.57 Gyrs) (Mamajek 2012). For later parts of this study, we interpolate in a grid of models with mixing lengths of 1.22, 1.72, and 2.22. For the analysis in Section 3.2, we varied the initial helium abundance and surface boundary conditions in our model grid to explore their impact on the theoretical uncertainties. We ran models with fixed helium fractions of 0.239, 0.290, and 0.330, in addition to the solar calibrated value of 0.272683. We also created models with different atmospheric surface boundary conditions (defined as the boundary pressure at \u03c4 = 2\/3): a look-up table from Kurucz (1997) models, and one from Castelli & Kurucz (2004) models, in addition to the gray atmosphere used in our main model grid.","Citation Text":["Iglesias & Rogers 1996"],"Citation Start End":[[621,643]]}
{"Identifier":"2015AandA...574A..62S__Kucera_et_al._(1998)_Instance_2","Paragraph":"In the last two decades UV, EUV, and X-ray observations from space made it possible to estimate the total mass of prominences reliably without needing to solve problems that occur when using visible light as was described in the previous paragraph. One possibility would be to estimate column mass of hydrogen and\/or helium plasma in prominences using spectral observations of UV and visible lines of hydrogen and helium (Balmer and Lyman lines) and sophisticated models in which plasma is assumed not being in the state of local thermodynamic equilibrium (NLTE models; see e.g. Labrosse et al.2010, and references therein). The problem can be a rather high complexity of such NLTE models which depend on various free parameters. Another possibility is to infer column mass and subsequently the total mass from the amount of radiation absorbed by the photoionisation in the prominence plasma at resonance continua of hydrogen and helium. This estimation of the mean column mass of prominences observed near the limb was first made by Kucera et al. (1998) using extreme-ultraviolet (EUV) observations from the Coronal Diagnostic Spectrometer (CDS; Harrison et al. 1995) on board the Solar and Heliospheric Observatory (SoHO) satellite. The advantage of the observations used by Kucera et al. (1998) was that CDS as a spectrograph observed only in spectral lines of interest, but spectrographs are able to obtain spectra from only one slit position during one exposure. Larger fields of view can be scanned by the spectrograph slit, as the CDS does, for example, but in such a case, intensities in different slit positions the scan is composed of were obtained at different times. In contrast, in filtergrams, intensities in all positions of the field of view are obtained at the same time, but the filter has transmission of certain width which can be wider than the spectral line of interest. The first estimation of the prominence mass using filtergrams was made by Golub et al. (1999) using Transition Region and Coronal Explorer (TRACE1) data. Similar studies were made by Gilbert et al. (2005, 2006) using observations of EUV Imaging Telescope (EIT; Delaboudini\u00e8re et al. 1995) on board SoHO in the 195\u2009\u00c5 channel. In those two works it was shown that it is necessary to estimate an amount of coronal emission behind and in front of the prominence (hereafter referred to as background and foreground radiation, respectively) to determine correctly the amount of absorbed radiation. The foreground radiation can be measured at the darkest place at a prominence, where it is assumed that all radiation from behind the prominence was absorbed. Then, the background radiation can be derived from the total coronal emission at the prominence location. Two ways to estimate the total emission were proposed and used. The spatial interpolative approach uses interpolation from intensities measured in the corona near a prominence. The temporal interpolative approach is suitable only for erupting prominences and it uses measurements of intensity in place of prominence after its eruption. In the work of Williams et al. (2013) these two approaches are also used to estimate the column mass of material returning to limb after prominence eruption. They used observations from the Atmospheric Imaging Assembly (AIA) instrument (Lemen et al. 2012) on board the Solar Dynamics Observatory (SDO) in several EUV coronal channels of wavelengths below 228\u2009\u00c5 (head of the resonance continuum of He\u2009ii). ","Citation Text":["Kucera et al. (1998)"],"Citation Start End":[[1277,1297]]}
{"Identifier":"2018AandA...616L...2K__Frew_et_al._(2016)_Instance_3","Paragraph":"The distances to planetary nebulae (PNe) have always faced the difficulty that nearby targets were lacking that could be reached well by direct methods. Trigonometric parallaxes have been obtained in a homogeneous long time-line campaign by the US Naval Observatory (USNO; Harris et al. 2007) and from the Hubble Space Telescope (HST; Benedict et al. 2009). Other studies (Acker et al. 1998; Smith 2015) showed that Hipparcos spacecraft parallaxes do not seem to be reliable. It was assumed that contamination by the emission of the surrounding nebulae caused these problems. Another model-independent method for distances to PNe are a cluster membership, as studied extensively by Majaess et al. (2007, 2014), and as discussed in Frew et al. (2016). In addition to these model-independent methods, a wide variety of statistical, model-dependent individual distance scales have been derived. The most frequently used of these are certainly those that are based on surface brightness versus angular sizes. They sometimes include optical depth corrections. All these methods have to be calibrated against a data set of nebulae with known distances. The older, widely used method is based on the 6 cm radio continuum flux, either using the ionized mass concept of Daub (1982) in the calibrations of Cahn et al. (1992) and Stanghellini et al. (2008), or by means of the radio continuum brightness temperature as used by van de Steene & Zijlstra (1994) and calibrated with a Galactic bulge sample. The newest model developed by Frew et al. (2016) is based on similar ideas, but makes use of the optical H\u03b1 surface brightness and a wide set of various calibrators. Moreover, they use a completely homogeneous data set for the brightness data derived earlier by themselves (Frew et al. 2013). Smith (2015) and Frew et al. (2016) described the underlying physics and assumptions for all these methods in detail. With the upcoming Gaia project (Gaia Collaboration 2016), a new era was expected to start for many classes of objects. The first step into this was described by Stanghellini et al. (2017) based on the combined Tycho + Gaia DR1 solution called TGAS (Michalik et al. 2015). With the second Data Release of Gaia (hereafter GDR2; Gaia Collaboration 2018), a complete homogeneous data set based only on Gaia measurements is available now for the first time. We present here the comparison of this new data set with common previous calibrations of PNe distances. Moreover, we compare it to the preliminary TGAS results in Stanghellini et al. (2017). Finally, we discuss possible caveats using the current GDR2.","Citation Text":["Frew et al. (2016)"],"Citation Start End":[[1803,1821]]}
{"Identifier":"2019AandA...632A.129W__Lavraud_et_al._2010_Instance_1","Paragraph":"Suprathermal electron strahls in the solar wind come from the Sun and are focused along magnetic field lines (Feldman et al. 1975; Rosenbauer et al. 1977; Pagel et al. 2005). Therefore, observations of CSE strahls within MCs can indicate that the flux rope structures are still connected with the Sun\u2019s magnetic field lines on both ends (Gosling et al. 1995; Larson et al. 1997; Shodhan et al. 2000; Feng et al. 2015, 2019). Counterstreaming suprathermal electrons can also be produced by other mechanisms, for example connection to the Earth\u2019s bowshock (Feldman et al. 1982; Stansberry et al. 1988), interplanetary shocks or corotating interaction regions (CIRs; Gosling et al. 1993; Steinberg et al. 2005; Lavraud et al. 2010), and distribution function depletions near the 90\u00b0 pitch angle (Gosling et al. 2001, 2002; Skoug et al. 2006). Among these mechanisms, the depletion CSEs are often observed on closed or open field lines within ICMEs, but the depletion CSEs are centered on and roughly symmetric to the 90\u00b0 pitch angle (Gosling et al. 2002), and can be distinguished from CSE strahls. Shodhan et al. (2000) examined the CSE strahl signatures of 52 MCs detected by a spacecraft near 1 AU and determined that approximately 87.5% of MCs exhibit CSE signatures, revealing that most MCs are still attached to the Sun at both ends at 1 AU. Gosling (1990) and Gosling et al. (1995) proposed explanations for how flux ropes arise in terms of three-dimensional reconnection close to the Sun. They illustrate how the original flux rope reconnects to form a helical field line connected to the Sun at both ends. The closed field lines of flux ropes can gradually open and occasionally disconnect from the corona when closed flux ropes expand from the Sun into the interplanetary space (Gosling et al. 1995). The proposal of Gosling et al. (1995) was confirmed by statistical results of Shodhan et al. (2000), namely most MCs exhibit CSE strahls in parts of their durations, and only 6 of the 52 MCs have no CSE. We therefore want to know whether or not nonMC ICMEs have the same CSE strahl signatures. In this study, CSE signatures from the Advanced Composition Explorer (ACE) from 1998 to 2008 are compared between ICMEs with and without MCs, and we discuss whether the CSE signatures are related to the flux-rope structures.","Citation Text":["Lavraud et al. 2010"],"Citation Start End":[[708,727]]}
{"Identifier":"2016AandA...595A..16T__Eggleton_et_al._(2006)_Instance_2","Paragraph":"We used the theoretical models by Eggleton et al. (2008), Charbonnel & Lagarde (2010), and Lagarde et al. (2012) for the comparison. They provide quantitative values representing the first dredge-up, thermohaline (TH), and thermohaline and rotation (TH+V) induced mixing. Eggleton et al. (2008) estimated the mixing speed with their formula for the diffusion coefficient and found that a range of three orders of magnitude in their free parameter can lead to the observed levels of 12C \/13C. Their predicted value of the 12C \/13C ratio for a solar-metallicity 2 M\u2299 star at the RGB tip is 17. A more recent model of the thermohaline-induced mixing by Charbonnel & Lagarde (2010) lists for the same stars a higher value of about 20. The model of Charbonnel & Lagarde (2010) of thermohaline instability induced mixing and the model by Eggleton et al. (2008) are both based on the ideas of Eggleton et al. (2006) and Ulrich (1972). It includes developments by Charbonnel & Zahn (2007). Eggleton et al. (2006) found a mean molecular weight (\u03bc) inversion in their 1 M\u2299 stellar evolution model, which occurred after the so-called luminosity bump on the RGB, when the hydrogen-burning shell reaches the chemically homogeneous part of the envelope. The \u03bc-inversion is produced by the reaction \\hbox{$^3{\\rm He(}^3{\\rm He}, 2p)^4{\\rm He}$}He(33He,2p)4He, as predicted in Ulrich (1972). It does not occur earlier because the magnitude of the \u03bc-inversion is low and negligible compared to a stabilising \u03bc-stratification. Charbonnel & Zahn (2007) computed stellar evolution models including the ideas of Kippenhahn et al. (1980), who extended Ulrich\u2019s equations to the case of a non-perfect gas. Charbonnel & Zahn (2007) also introduced a double diffusive instability (i.e. thermohaline convection) and showed its importance in the chemical evolution of red giants. This mixing connects the convective envelope with the external wing of the hydrogen-burning shell and induces surface abundance modifications in evolved stars (Charbonnel & Lagarde 2010). ","Citation Text":["Eggleton et al. (2006)"],"Citation Start End":[[982,1004]]}
{"Identifier":"2015MNRAS.449..288P__Kohri_et_al._2005_Instance_1","Paragraph":"Here, we propose a new speculative site for strong r-process nucleosynthesis in which it is formed by jets launched by an NS spiralling-in inside the core of a giant star. These jets will also explode the star. This is a rare evolutionary route and hence complies with the finding of large variations in the abundances of these elements. In this first study, we limit ourselves to present the scenario and show its viability. Most ingredients of our newly proposed scenario were studied in the past but were never put together into a coherent picture to yield a new possible site for the strong r-process. Previously studied ingredients of our proposed scenario include the CE of an NS and a giant (Thorne & Zytkow 1975; Armitage & Livio 2000; Chevalier 2012), the launching of neutron-rich gas from accretion discs around compact objects (e.g. Surman & McLaughlin 2004; Kohri et al. 2005), the launching of jets by NS accreting at a high rate (Fryer, Benz & Herant 1996) and the formation of r-process elements in jets from NS (Fryer et al. 2006). The idea that jets can explode stars under specific conditions was also studied in the past, e.g. in a CE evolution (Chevalier 2012). Another specific class of models are based on magnetic amplification by a rapidly rotating stellar core (e.g. LeBlanc & Wilson 1970; Bisnovatyi-Kogan, Popov & Samokhin 1976; Meier et al. 1976; Khokhlov et al. 1999; MacFadyen, Woosley & Heger 2001; Woosley & Janka 2005; Couch, Wheeler & Milosavljevi\u0107 2009). This magnetorotational mechanism creates bipolar outflows (jets) around the newly born NS that are able to explode the star. However, the required core's rotation rate is much larger than what stellar evolution models give, hence making most of these models applicable for only special cases. We, on the other hand, claim that all CCSNe are exploded by jets, the jittering-jets model, that also synthesize r-process elements (but not the strong r-process), and hence our newly proposed scenario is part of a unified picture we try to construct for exploding all massive stars and the synthesis of r-process elements. In the jittering-jets model, the explosion of CCSNe is powered by jittering jets launched by an intermittent accretion disc around the newly born NS (Papish & Soker 2014a,b). The intermittent accretion disc is formed by gas accreted from the convective core regions that have a stochastic angular momentum (Gilkis & Soker 2014; Papish & Soker 2014a).","Citation Text":["Kohri et al. 2005"],"Citation Start End":[[871,888]]}
{"Identifier":"2017MNRAS.470..142S__Gaspari_&_Churazov_2013_Instance_1","Paragraph":"The dynamic viscosity of a fully ionized gas in the absence of magnetic fields (see Spitzer 1962)1(2)\r\n\\begin{equation}\r\n\\mu _{\\rm Spitzer} = 2.21\\times 10^{-15}\\,\\frac{(T\\,[\\mathrm{K}])^{5\/2}A_{\\rm i}^{1\/2}}{Z^4\\ln \\Lambda }\\,\\mathrm{g\\,cm^{-1}\\,s^{-1}}\\,\r\n\\end{equation}\r\nis mainly a function of the temperature T. The parameters Ai and Z are the atomic weight and charge, respectively, of the ions. The Coulomb integral ln\u2009\u039b depends logarithmically on temperature and density, and is commonly approximated by a constant value around 40 for clusters (see Gaspari & Churazov 2013; Smith et al. 2013). Slices of the resulting viscosity are shown in middle panels of Fig. 1 for three stages of the simulated merger. In Fig. 2, radial profiles of volume averages (blue), medians (red) and interquartiles (red shaded) are plotted for the cluster produced by the merger at z = 0. The typical value of the Spitzer viscosity in the WHIM is \u03bcSpitzer \u223c 10\u2009g\u2009cm\u22121\u2009s\u22121, with a gradual increase towards the core. Around the accretion shock (R \u2248 5\u2009Mpc), the viscosity drops rapidly. It turns out that the median profile reflects this drop much better than the volume-averaged (or mass-averaged) profile because averaging a quantity that varies over several orders of magnitude is biased towards high values. Beyond the outer shocks, the values calculated with equation (2) become physically meaningless because the unshocked gas is neutral or only partially ionized. For the ICM, viscosities above 100\u2009g\u2009cm\u22121\u2009s\u22121 would be representative if it were not for the depression within 100\u2009kpc from the centre. As discussed in Schmidt et al. (2016), this is a consequence of the overpronounced cool cores in our simulations. Feedback, particularly from AGNs, would raise the core temperature and, thus, the viscosity. This is important to bear in mind for the following discussion: In any case, the viscosity following from equation (2) is large enough to significantly affect or even suppress dynamical instabilities in both the ICM and the WHIM. This was also demonstrated by Roediger et al. (2013), Roediger et al. (2015) and ZuHone et al. (2015) in the context of Kelvin\u2013Helmholtz instabilities and gas sloshing induced by minor mergers.","Citation Text":["Gaspari & Churazov 2013"],"Citation Start End":[[557,580]]}
{"Identifier":"2020MNRAS.496.1718E__Wong_et_al._2019_Instance_1","Paragraph":"(i) Our capability of reproducing the lensed images down to the noise level without fully correctly modelling the cored central mass density distribution of the lenses indicates some form of the source-position transformation (SPT; Schneider & Sluse 2014), in line with the previous findings by Unruh et al. (2017). As a consequence, our reconstructions lead to a systematic fractional error on the Hubble constant of $25_{-19}^{+37}$ per cent (in comparison to a statical error of $12_{-3}^{+6}$ per cent when the shape of the lensing potential is perfectly known). This result is in agreement with the latest analysis of Blum, Castorina & Simonovi\u0107 (2020) that shows that cored (dark matter) mass density distributions give rise to approximate MSDs, and an error on the inferred Hubble constant. The latest cosmographic analyses (see e.g. Wong et al. 2019) have attempted to break these degeneracies by including the information contained in the kinematic properties of the lens galaxies and the positions of the lensed quasar images. However, the validity of this approach has been recently debated by Kochanek (2020), who has demonstrated that departures from single power-law mass distributions are responsible for a fractional error on the Hubble constant of 30 per cent. While the cores in the simulations analysed in this paper are artefacts related to limited resolution, cored mass density distribution in real galaxies may be developed by the effect of baryonic processes (see e.g. Chan et al. 2015) or changes in the dark matter properties (Schive, Chiueh & Broadhurst 2014; Spergel & Steinhardt 2000). Moreover, similar additional complexities exist in real galaxies are related, for example, to the presence of faint discs (Hsueh et al. 2018), bars, or other (baryonic) structures (Gilman et al. 2018; Xu et al. 2013). More generally, there exist many plausible deviations from a smooth power-law distribution, such as broken power laws (see e.g. Du et al. 2020) or multiple component models (see e.g. Nightingale, Dye & Massey 2018), which can produce comparable degeneracies. Together with the findings of Blum et al. (2020), our results have important implications for the analysis of time delays and a potential solution to the H0 tension (Wong et al. 2019).","Citation Text":["Wong et al. 2019"],"Citation Start End":[[841,857]]}
{"Identifier":"2022AandA...661A..58S__Pinto_et_al._2016_Instance_1","Paragraph":"In Fig. 1, the state of the corona is shown both before the introduction of Alfv\u00e9nic perturbations (panels a\u2013c), as well as after the solar wind has reached a quasi-steady state upon the injection of Alfv\u00e9n waves of the lowest considered pump frequency of 0.3 mHz (panels d\u2013f). The solar wind density, sound speed, and radial speed are shown at these two evolutionary instances for the three flux-tube geometries considered, with expansions of fmax\u2004=\u20043, 5, and 10. The response of a polytropic solar wind (\u03b3\u2004\u20045\/3) for different magnetic flux tube expansions in the absence of a pump wave has been described in previous works (Kopp & Holzer 1976; Pinto et al. 2016). Close to the Sun (r\u2004\u2004\u223c2\u2006RS), the corona is seen to exhibit similar plasma conditions irrespectively of the chosen flux-tube geometry. The heliocentric distance to which these similar plasma conditions persist (r\u2004\u2004\u223c2\u2006RS in our study) is partly due to the choice of parameters in Eq. (3). At greater heliocentric distances, the mass density is higher for smaller values of fmax, while the opposite is observed for the sound speed and radial velocity. Upon injecting a 0.3 mHz Alfv\u00e9n wave, we see that the solar wind exhibits no visually discernible perturbations, indicating that this medium-frequency (0.3 mHz) Alfv\u00e9n wave simply acts as a direct source of momentum (via the wave pressure gradient) and energy to the solar wind. As compared to the steady state, the sound speed and density are lower and radial speed higher for all considered fmax. This behaviour of the solar wind in response to hour-scale Alfv\u00e9n waves is similarly found in other numerical simulations (Shoda et al. 2018a) and is indicative of the low dissipation of Alfv\u00e9n waves at these frequencies either by reflections in the medium or by PDI. Such hour-scale Alfv\u00e9n waves are also seen to dominate observations (Banerjee et al. 2009) as these waves do not dissipate and propagate through the medium as in Fig. 1. The acceleration of the solar wind is discussed in Sect. 4.2.","Citation Text":["Pinto et al. 2016"],"Citation Start End":[[646,663]]}
{"Identifier":"2022AandA...665A.115C__Clark_&_Steele_(2000)_Instance_1","Paragraph":"While different studies devoted to the near-IR spectra of Be stars show low-resolution data, many of them are restricted to a small sample, and some others analyse reduced spectral ranges. For instance, there are only a few studies done in the J band, and they focus mainly on a particular object or a specific spectral line or element (Mathew et al. 2012a,b; \u0160tefl et al. 2009). Also, individual spectral ranges have already been studied for large samples of Be stars. Steele & Clark (2001) presented H-band spectroscopy of Be stars with a spectral resolution of R \u2243 3000. They reported Brackett and Fe ii lines in emission and, from the analysis of the strength ratio of the higher Brackett lines to Br\u03b3, were able to distinguish early- from late-type Be stars. Later, Chojnowski et al. (2015) published high-resolution H-band spectroscopy for a great number of Be stars observed with APOGEE. They found that the Br11 emission line is formed preferentially in a circumstellar disc at an average distance of ~2.2R*, while the higher Brackett lines seem to originate in an innermost region. Several emission lines have been identified for the first time, such as C i \u03bb 1.6895 \u03bcm, which is also formed in the inner region of the discs. In a later work, Chojnowski et al. (2017) analysed the variation of the emission strength, peak intensity ratio, and peak separation. Their analysis revealed a variety of temporal variability, including the disappearance and appearance of the line emission on different timescales. In the K band, Hanson et al. (1996), Clark & Steele (2000) and Granada et al. (2010) reported Br\u03b3, Br\u03b4, and Pfund lines in emission together with lines of He i in emission or absorption, and Mg ii, Fe ii and Na i lines in emission. Clark & Steele (2000) related the infrared characteristics to the underlying properties of the stars: objects that present He i features in emission or absorption are B3 or earlier; if the star presents Mg ii in emission but no He i, it is between B2 and b4; objects with Br\u03b3 emission but no evidence of He i or Mg ii are B5 or later. Lenorzer et al. (2002b) provided an extensive atlas of early type stars, including a number of Be stars, covering the K and L bands, while Lenorzer et al. (2002a) analysed the H recombination lines of those stars and constructed a diagram of flux ratios of some selected recombination lines: Hu14\/Br\u03b1 and Hu14\/Pf\u03b3. In this diagram, the location of the objects gives information about the density of the emitting gas. After that, Mennickent et al. (2009) presented a classification scheme for Be stars based on the intensity of the L-band hydrogen-emission lines. The objects in each group fall in different regions of Lenorzer\u2019s diagram; thus, this classification scheme is probably connected to the density of the disc. Granada et al. (2010) analysed a sample of eight Be stars and classified them with Mennickent\u2019s criterion. They found that for group I objects, the equivalent widths (EW) of Br\u03b1 and Br\u03b3 lines are similar, while for stars in group II the EW(Br\u03b3) is much larger (more than five times) than the EW(Br\u03b1). Besides this, Mennickent et al. (2009) and Granada et al. (2010) reported emission lines not only in Br\u03b1, Pf\u03b3, and the Humphreys series but also in He i \u03bb4.038 \u03bcm and He i \u03bb4.041 \u03bcm. Sabogal et al. (2017) showed a sample of L-band Be-star spectra and correlated the infrared features with the optical H\u03b1 line behaviour.","Citation Text":["Clark & Steele (2000)"],"Citation Start End":[[1554,1575]]}
{"Identifier":"2020AandA...634A..99B__Rothschild_et_al._2017_Instance_1","Paragraph":"where WCRSF is the width of the cyclotron line, F\u03bd is the emergent flux including cyclotron scattering, and F0 the input continuum flux without a CRSF. In the Monte Carlo simulations performed by Schwarm et al. (2017b), only 1\u201310% of the initial photons undergo resonant scattering, that is, r\u2004\u2273\u200490%. Consequently, we would expect \u03c4CRSF\u2004\u2272\u20040.1 for the observed optical depths9. From our best-fit model of the spectra of Epoch II, the calculated residual fluxes after Eq. (10) are \u226578% for all three detected cyclotron lines (see Table 2). These values together with the fitted optical depths of \u03c4CRSF\u2004\u2264\u20040.24 (see Table 1) are in excellent agreement with the expectations from the simulations by Schwarm et al. (2017b). Although spawned photons from higher harmonics, which originate from multi-photon scattering and from the radiative de-excitation of electrons excited into higher levels, can affect the shape of the fundamental line (Isenberg et al. 1998), they can only decrease its depth. We note that residual fluxes down to r\u2004\u2273\u200436% have been observed for GX 304\u22121 (Rothschild et al. 2017), which can be obtained by, for example, a higher optical depth at the line-forming region. However, the CRSFs are not expected to absorb a significant fraction of the total broad-band continuum flux, which is the case for the model by Iyer et al. (2015). We conclude that the CRSF parameters that these latter authors found are due to a degeneracy between the continuum modeling and the CRSF modeling when unphysical values for the CRSF are allowed. In summary, the choice of a phenomenological model for the X-ray spectrum of an accreting neutron star may cause strong discrepancies relative to the theoretically expected values for the CRSF parameters. In particular, an erroneous shape or an overestimation of the X-ray flux is sometimes corrected by the introduction of further strong absorption components, which are probably not real. We stress that we do not claim that our choice of the phenomenological continuum model for 4U 0115+634 describes its true X-ray spectrum best. However, in contrast to other continua the model applied here yields CRSF parameters which are consistent with theoretical expectations. These conclusions desperately call out for a self-consistent model for both the spectral continuum and the cyclotron resonant scattering features.","Citation Text":["Rothschild et al. 2017"],"Citation Start End":[[1070,1092]]}
{"Identifier":"2017ApJ...835..286I__Pope_et_al._2006_Instance_1","Paragraph":"Despite 20 years of deep submm surveys since Smail et al. (1997), our knowledge of the upper half of the redshift distribution of SMGs remains incomplete. Early attempts to determine redshifts were conducted toward SMGs with radio counterparts because low-resolution (sub)mm images obtained with single dishes require high-resolution radio continuum maps from radio interferometers such as the Karl G. Jansky Very Large Array (VLA) in order to pinpoint source positions (Ivison et al. 1998, 2000, 2002, 2005, 2007; Smail et al. 1999; Borys et al. 2004; Pope et al. 2006; Aretxaga et al. 2011; Biggs et al. 2011; Yun et al. 2012; Umehata et al. 2014). Intensive studies of radio-bright SMGs were able to yield spectroscopic redshifts for those out to z \u223c 3 (e.g., Chapman et al. 2003, 2005). However, at that time, radio sensitivities could not detect SMGs beyond z \u223c 3, and as many as half of the SMGs lacked reliable radio counterparts (see e.g., Ivison et al. 2007; Biggs et al. 2011, cf. Lindner et al. 2011). Later attempts to determine SMG positions and redshifts using near- and mid-IR imaging could not fully overcome the bias toward lower redshifts, since the K corrections there are no more favorable than those in the radio regime, such that high-redshift sources are much fainter (e.g., Wardlow et al. 2011; Yun et al. 2012). Millimeter (mm) spectroscopic surveys toward gravitationally lensed dusty star-forming galaxies, taking advantage of their apparent ultra brightness, revealed a redshift distribution stretching out to z \u223c 5.8 (e.g., Vieira et al. 2013; Wei\u00df et al. 2013; Strandet et al. 2016). These surveys suggested a larger fraction of SMGs at z \u2273 3 than previous studies of unlensed SMGs, perhaps partly because they were selected at 1.3 mm rather than the traditional 0.8\u20131.1 mm, but also because the requirement for high magnification favors galaxies with a long line of sight. We need to reveal the intrinsic redshift distributions of unlensed SMGs in large contiguous maps to determine their abundance in the early Universe and to study the evolution of the most massive galaxies via abundance matching with other populations, and with cosmological predictions (e.g., Hayward et al. 2013; Cowley et al. 2015).","Citation Text":["Pope et al. 2006"],"Citation Start End":[[553,569]]}
{"Identifier":"2017MNRAS.469.1559G__Fraser_et_al._2015_Instance_1","Paragraph":"We can also compare to the models by Dessart, Audit & Hillier (2015), who synthesize light curves and spectra for superluminous SNe IIn with a variety of physical CSM parameters: in their fig. 20, they show the simulated bolometric light curves for their models out to +450 d. Past 200 d, most of the light curves experience a sudden and significant drop in luminosity that we did not observe for SN 2009ip, except for three (xe3m6r, xm3 and xm6). These models are distinguished from the others by having progenitor stars that experienced a more rapid mass-loss rate during the late stages of stellar evolution, and in one case, a CSM that extends to a radius 50\u2009per\u2009cent farther than the others. The total bolometric luminosities for these models are \u22733 orders of magnitude higher at late times than SN 2009ip (log\u2009Lbol \u2248 40; Fraser et al. 2015) owing to a more energetic underlying explosion driving the interaction. The synthesized early-time spectra for these three models are also qualitatively similar to those of SN 2009ip (see fig. 22 of Dessart et al. 2015). Late-time synthesized spectra are not shown for these three models, but are given for their model X, in a comparison to luminous SN IIn 2010jl (see also Section 3.4). The synthetic spectra exhibit significantly more asymmetry and have a strong, blueshifted component of the H\u2009\u03b1 line, which is even more exaggerated than that seen for SN 2010jl at \u223c200 d. Dessart et al. (2015) explain this blueshifted feature as arising from an optically thick, cool dense shell between the reverse and forward shocks in the CSM, but for SN 2009ip we would not expect to see this signature because we found that the CSM exterior to the shock front is not yet optically thin (i.e. Section 3.2.1). Although not every aspect of the models by Dessart et al. (2015) represent the physical scenario of SN 2009ip, where appropriate comparisons can be made they support a moderately energetic explosion into a large and extended mass of CSM.","Citation Text":["Fraser et al. 2015"],"Citation Start End":[[827,845]]}
{"Identifier":"2015AandA...578A..74G__Gonz\u00e1lez-Mart\u00edn_et_al._(2009b)_Instance_1","Paragraph":"The most natural explanation is that the Compton-thick nature of some low-luminosity AGNs results in an underestimation of the true LX(2\u221210 keV) of these sources. These LX(2\u221210 keV) estimates come from studies using the spectra of LINERs at energies below 10 keV. However, Compton-thick sources show the bulk of the AGN power at energies above 10 keV. The intrinsic luminosity could be 10\u221270 times higher than the estimated using only energies below 10 keV in the Compton-thick scenario (Maiolino et al. 1998). Gonz\u00e1lez-Mart\u00edn et al. (2009b) classified around 50% of their LINER sample as Compton-thick candidates. Figure 8 shows the AGN-dominated objects in our study attending to their Compton-thin (left) and Compton-thick (middle) classification. Most of the Compton-thin sources are close to the previously reported correlation for AGN. Compton-thick S2s are nicely placed along the AGN correlation found by Asmus et al. (2011). This is expected because the X-ray luminosities included for S2s are all corrected for their Compton-thick nature (most of them included in Goulding et al. 2012), either using X-ray measurements above 10 keV or assuming a factor between the observed and intrinsic X-ray luminosity for other Compton-thick AGNs (see Panessa et al. 2006, for details in this conversion factor). Compton-thick LINER candidates tend to be shifted towards X-ray luminosities lower than predicted for the AGN correlation. However, most of them are not consistent with the SB correlation either. The double arrows of Fig. 8 (middle panel) show their expected locus if the X-ray luminosity were ~10\u221270 times higher. Most of the Compton-thick LINERs can be placed in the AGN correlation if the correction is applied. The linear fit to the entire sample, once the intrinsic X-ray luminosity is corrected, results in LX(2\u221210 keV)(intrinsic) = 40 \u00d7 LX(2\u221210 keV)(observed) for Compton-thick LINERs (marked as white stars in Fig. 8, right panel), which is very close to the linear relation found for AGN. Nonetheless, four LINERs and one S2 still remain very close to the SB correlation. ","Citation Text":["Gonz\u00e1lez-Mart\u00edn et al. (2009b)"],"Citation Start End":[[511,541]]}
{"Identifier":"2015AandA...582A.117M__Aschwanden_et_al._(2002)_Instance_1","Paragraph":"A common property of observed transverse coronal loop oscillations is that they are damped quickly, usually within a few oscillation periods. It is now generally accepted that the main damping mechanism is resonant absorption (Sakurai et al. 1991; Goossens et al. 1992, 2002; Ruderman & Roberts 2002), transferring energy from the global kink mode to local azimuthal Alfv\u00e9n modes in the boundary layer of the loop structure, where the two frequencies match (for a review of theoretical results, see Goossens et al. 2011). Less frequently, nearly undamped or even growing transverse oscillations are observed (Wang et al. 2012). Decay-less low amplitude kink oscillations, which are present in loops before and after a high-amplitude flare triggered damped kink oscillations can be explained in terms of a damped linear oscillator excited by a continuous low amplitude harmonic driver (Nistic\u00f2 et al. 2013; Anfinogentov et al. 2013). Some examples of observed undamped high amplitude coronal loop oscillations can be found in Aschwanden et al. (2002) and Aschwanden & Schrijver (2011). In the latter paper, it was concluded that this could only happen if the thickness of the boundary layer (a radially inhomogeneous outer layer of the coronal loops) is much smaller than the loop radius, in order to minimize the damping due to resonant absorption. It was also evident from the observations that the loop was cooling during the particular undamped oscillation event. It seemed a rather natural explanation that the amplification due to cooling may counterbalance the damping due to resonant absorption, explaining the undamped oscillations, as it was first suggested in Ruderman (2011b). The most elaborate analytical study regarding time-dependent kink oscillations (Ruderman 2011a) considered the simultaneous effects of both damping due to resonant absorption and amplification due to plasma cooling. The conclusion was that, for typical boundary layer thicknesses, the amplification due to plasma cooling can account for the observed undamped oscillations only if the cooling happens quickly, on a timescale close to the oscillation period. However, in the paper, the effects of cooling and resonant damping have been studied under the assumption that both the characteristic cooling time and damping time are much longer than the characteristic oscillation period, when using the WKB method. Thus, it is questionable whether the derived equations remain valid for rapid cooling, i.e. for cooling times near to the oscillation periods. On the other hand, neglected non-linear behavior may change the outcome considerably, for example the Kelvin-Helmholtz instability at the tube boundary (Heyvaerts & Priest 1983; Ofman et al. 1994; Terradas et al. 2008), and the ponderomotive forces for high amplitudes (Terradas & Ofman 2004). The aim of this study is to further investigate the effects of radiative plasma cooling on the fundamental standing kink oscillation by means of numerical analysis. ","Citation Text":["Aschwanden et al. (2002)"],"Citation Start End":[[1025,1049]]}
{"Identifier":"2021ApJ...921...18K__Kushwaha_et_al._2018a_Instance_2","Paragraph":"The most unique and characteristic observational feature of blazars\u2019 highly variable broadband emission is the broad bimodal SED extending from the lowest accessible EM band, i.e., the radio, to the highest accessible, i.e., GeV-TeV \u03b3-rays. The broadband SED of all blazars can be categorized into three different spectral subclasses: low-energy-peaked (LBL\/LSP), intermediate-energy-peaked (IBL\/ISP), and high-energy-peaked (HBL\/HSP; Fossati et al. 1998; Abdo et al. 2010), based on the location of the low-energy hump. A remarkable property of each spectral subclass is the stability of the location of the two peaks despite huge variations in flux and often spectral shape. Only in a few rare instances has an appreciable shift in the location of the peaks been observed, e.g., the 1997 outburst of Mrk 501 (Pian et al. 1998; Ahnen et al. 2018) and the activity of OJ 287 from the end of 2015 to the middle of 2017 (Kushwaha et al. 2018a, 2018b). Even these two cases are remarkably different. In the case of Mrk 501, the locations of both the peaks shifted to higher energies. On the contrary, in OJ 287, a shift in the location of only the high-energy peak was observed during the 2015\u20132016 activity (Kushwaha et al. 2018a, 2019), while in 2016\u20132017 a new broadband emission component overwhelmed the overall emission, appearing as an overall shift in both the peaks as revealed in the detailed study by Kushwaha et al. (2018b). With the SED being the prime observable for exploration of the yet-debated high-energy emission mechanisms, such changes offer invaluable insights about the emission processes. For example, in Mrk 501 the shift in both peaks strongly implies the same particle distribution for the overall emission, while for OJ 287 the shift of only the high-energy peak can be reproduced by either inverse Compton scattering of the broad-line region photon field (Kushwaha et al. 2018a) or emission of hadronic origin (Oikonomou et al. 2019; Rodr\u00edguez-Ram\u00edrez et al. 2020).","Citation Text":["Kushwaha et al. 2018a"],"Citation Start End":[[1206,1227]]}
{"Identifier":"2021AandA...648A..73B__Marois_et_al._2008_Instance_2","Paragraph":"We present the photometry of the companion in Fig. 4 in a color-magnitude diagram. The corresponding numerical values are reported in Table 2. YSES 2b is consistent with a late L to early T spectral type when comparing it to colors of field brown dwarfs from the NIRSPEC Brown Dwarf Spectroscopic Survey (McLean et al. 2003, 2007), the IRTF spectral library (Rayner et al. 2009; Cushing et al. 2005), the L and T dwarf data archive (Knapp et al. 2004; Golimowski et al. 2004; Chiu et al. 2006), and the SpeX Prism Libraries (Burgasser et al. 2004, 2008, 2010; Gelino & Burgasser 2010; Burgasser 2007; Siegler et al. 2007; Reid et al. 2006; Kirkpatrick et al. 2006, 2010; Cruz et al. 2004; Burgasser & McElwain 2006; McElwain & Burgasser 2006; Sheppard & Cushing 2009; Looper et al. 2007, 2010; Muench et al. 2007; Dhital et al. 2011). Object distances were derived from Gaia EDR3 (Gaia Collaboration 2021), the Brown Dwarf Kinematics Project (Faherty et al. 2009), and the Pan-STARRS1 3\u03c0 Survey (Best et al. 2018). In color-magnitude space, YSES 2b is very close to the innermost three planets of the HR 8799 multi-planetary system (Marois et al. 2008, 2010). These three planets are classified as mid to late L type dwarfs (e.g., Greenbaum et al. 2018), which agrees well with the sequence evolution of the adjacent field brown dwarfs from L to T spectral types3. A similar spectral type in this domain, therefore, seems very likely for YSES 2b, requiring confirmation by measurements at higher spectral resolution. Whereas the masses of the spectrally similar trio of HR 8799 c, d, and e are in the range 7\u201312 MJup (Wang et al. 2018; Marois et al. 2008, 2010), it is likely that YSES 2b has an even lower mass as the system age of (13.9 \u00b1 2.3) Myr is significantly younger than the age of HR 8799, which is claimed to be member of the Columba association with an age of 30\u201350 Myr (Zuckerman et al. 2011; Bell et al. 2015). This is supported by the AMES-COND and AMES-dusty models (Allard et al. 2001; Chabrier et al. 2000) that we present in Fig. 4 for a system age of 13.9 Myr. An individual evaluation of these isochrones yielded masses from 5.3 MJup to 8.0 MJup as presented in Table 2. The uncertainties originate from the errors in the system age and planet magnitude that were propagated by a bootstrapping approach with 1000 randomly drawn samples from Gaussian distributions around both parameters. When combining the posterior distributions for the different models and filters we derived a final mass estimate of \n\n$6.3^{1.6}_{-0.9}\\,M_{\\mathrm{Jup}}$6.3\u22120.9+1.6\u2009MJup\n as the 68% confidence interval around the median of the sample. This estimate is based on broadband photometric measurements alone; further spectral coverage of the planetary SED will be important to constrain its effective temperature, luminosity, surface gravity, and mass.","Citation Text":["Marois et al. 2008"],"Citation Start End":[[1636,1654]]}
{"Identifier":"2020ApJ...898...52M__Dobbs_et_al._2014_Instance_1","Paragraph":"Idealized simulations have the advantage of carefully controlled conditions but the disadvantages that the turbulence is driven in an artificially prescribed manner to maintain a fixed overall turbulent amplitude and the processes leading to cloud formation and destruction are not followed. In reality, GMCs form due to a combination of large-scale ISM flows (including turbulence, shear, and epicyclic motion) and gravity (both stellar gravity and self-gravity) that lead to collection of material from a large volume, as mediated by thermal and magnetic pressure, and a change from the atomic to the molecular phase as the gas cools (e.g., McKee & Ostriker 2007; Dobbs et al. 2014; Chevance et al. 2020). Turbulence on scales less than the scale height of the warm\u2013cold ISM likely originates primarily due to the feedback from young stars (Elmegreen & Scalo 2004; Mac Low & Klessen 2004; McKee & Ostriker 2007),3\n\n3\nGravitational instabilities in the combined gas\u2013stellar system (e.g., Jog & Solomon 1984; Romeo 1992; Rafikov 2001; Kim & Ostriker 2007) can drive horizontal motions at very large scales, as seen in numerical simulations (e.g., Kim & Ostriker 2007; Shetty & Ostriker 2008; Agertz et al. 2009; Dobbs et al. 2011; Hopkins et al. 2012; Agertz & Kravtsov 2015, and citations within), but these motions generally do not reach supersonic amplitudes unless they are associated with gravitational collapse. In addition, turbulence at scales less than the disk scale height can be driven by spiral shocks and the magnetorotational instability, but numerical simulations show that the corresponding amplitudes are relatively modest in cold gas (e.g., Wada & Koda 2004; Piontek & Ostriker 2005, 2007; Kim et al. 2006, 2010; Dobbs & Bonnell 2007; Bonnell et al. 2013, and citations within).\n whether inherited from a GMC\u2019s formation stage or produced internally. Considering that GMCs live for at most a few turbulent crossing or freefall times (Kawamura et al. 2009; Kruijssen et al. 2019), it is not clear that internal GMC conditions can control star formation in a way that is entirely divorced from their formation and destruction processes.","Citation Text":["Dobbs et al. 2014"],"Citation Start End":[[666,683]]}
{"Identifier":"2022MNRAS.512.1814S__Pinto,_Middleton_&_Fabian_2016_Instance_1","Paragraph":"The short-term (seconds to hours) temporal properties of ULXs are quite different from those observed in Galactic X-ray binary systems and are not related to specific spectral states (e.g. Heil et al. 2009): even if the amplitude of short-term variability is larger in soft ultraluminous sources than in hard ones (Sutton et al. 2013), the variability appears to be poorly predictable, i.e. it is not found in all sources with a soft spectrum and, when detected in a source, it is not necessarily found in all the observations of that ULX. Such variability was explained with the existence of optically thick and non-uniform winds, radiatively ejected by the accretion disc, which stochastically intersect our line of sight (e.g. Middleton, Sutton & Roberts 2011; Middleton et al. 2015a). Observational evidences of these winds come from the detection of blue-shifted absorption lines in the high-quality grating spectra of some ULXs (e.g. Pinto, Middleton & Fabian 2016; Kosec et al. 2018) and by the nebulae observed in the radio (e.g. Cseh et al. 2012), optical (e.g. Pakull, Soria & Motch 2010) or X-ray band (Belfiore et al. 2020). Long-term flux variability (days to months) is also observed in most of the monitored ULXs. Such variability, in some cases, can be as high as several orders of magnitudes, implying that these sources are transients (detected at least once in the ultraluminous state, Lx > 1039 erg s\u22121, and either once at a significantly smaller luminosity or undetected below the instrumental sensitivity; e.g. Soria et al. 2012; Pintore et al. 2018a). Often the pulsating ULXs (PULXs) show large flux variations, implying they can be considered as transient ULXs. A bimodal flux distribution is sometimes observed in the long-term light curves of PULXs (e.g. Walton et al. 2015a; Motch et al. 2014). A possible explanation for this long-term behaviour may be the propeller effect: when the magnetospheric radius of the NS becomes larger than the corotation radius of the accreting matter in the disc, a centrifugal barrier prevents the accretion on to the NS, with a corresponding decrease in the X-ray flux (Illarionov & Sunyaev 1975; Tsygankov et al. 2016; Grebenev 2017). The propeller effect can act only if the accretor is an NS with a strong magnetic field, so a bimodal flux distribution can be used to identify candidate PULXs, even when pulsations are not detected (e.g. Earnshaw, Roberts & Sathyaprakash 2018; Song et al. 2020). Furthermore, long-term monitored ULXs revealed possible (super-)orbital variability (e.g. Foster, Charles & Holley-Bockelmann 2010; An, Lu & Wang 2016; F\u00fcrst et al. 2018), the origin of which is still a matter of debate. Super-orbital periods are also observed in the light curves of the PULXs, with periods of tens to hundreds days (e.g. Walton et al. 2016; Brightman et al. 2019, 2020). The detections of such long-term periodicities were only possible thanks to the flexibility and performance of the Neil Gehrels Swift Observatory (hereafter Swift; Gehrels et al. 2004).","Citation Text":["Pinto, Middleton & Fabian 2016"],"Citation Start End":[[940,970]]}
{"Identifier":"2017MNRAS.470..713M__Lutovinov_et_al._2012_Instance_1","Paragraph":"X Persei (4U 0352 + 309) is a binary stellar system at a distance of \u223c0.95 kpc (Telting et al. 1998) and consists of a slowly spinning neutron star (White, Mason & Sanford 1977, Pspin \u223c 837 s) that is accreting matter from its Be-star companion (Lyubimkov et al. 1997). The source is peculiar in that it is a persistent source and does not show Type-I outbursts as are commonly observed in other Be-\/X-ray binaries. But X-ray flares and variability in the X-ray light curve of the source have been observed (La Palombara & Mereghetti 2007; Lutovinov, Tsygankov & Chernyakova 2012). The orbital period of the binary system is long, $P_{\\text{orb}}\\sim 250 \\rm {d}$, but the orbital eccentricity is only e \u223c 0.11 (Delgado-Mart\u00ed et al. 2001). The long orbital period and the low X-ray luminosity of the source support the assumption that the neutron star in this binary is accreting quasi-spherically by capturing matter from the stellar wind of the Be-star (Shakura et al. 2012). The X-ray spectrum of X Persei is also unlike most accreting X-ray pulsars, the latter can usually be modelled with a power law and a cutoff in the range 10\u201330 keV (White, Swank & Holt 1983). Instead, X Persei has a peculiarly hard X-ray spectrum, making the source detectable at energies higher than 100 keV (Lutovinov et al. 2012). Thus, a single-component spectral model is insufficient to model the source continuum spectrum, making it necessary to include a second high-energy component to explain the emission at harder X-rays. Several authors have used several different models to explain the observed spectrum, of which we note the following works. The RXTE observations were modelled using a blackbody component at lower energies and a power-law component with an exponential cutoff at higher energies (Coburn et al. 2001). A cyclotron resonance scattering feature (CRSF) was also required to account for an absorption-like feature seen around 30 keV. The spectrum obtained from BeppoSAX observation was modelled with two power-law components with exponential cutoffs at low and high energies and did not require any CRSF (Di Salvo et al. 1998). The INTEGRAL observations used thermal Comptonization and bulk motion Comptonization to account for the observed spectrum at the low and high energies, respectively (Doroshenko et al. 2012). The INTEGRAL observations also did not require any CRSF to improve the spectral model. Thus, the presence of CRSF noted in the RXTE observations was not detected in either the BeppoSAX or the INTEGRAL observation and hence needs to be verified. There is also no clear agreement on the unusual nature of the X-ray spectrum with respect to the continuum model.","Citation Text":["Lutovinov et al. 2012"],"Citation Start End":[[1288,1309]]}
{"Identifier":"2020ApJ...889....9M__Vulcani_et_al._2018_Instance_1","Paragraph":"The lower panel of Figure 9 shows the SFR density against the H2 mass density for each pixel (black dots) and the average value found within each of the analyzed region as colored symbols. Red dashed lines are fixed depletion times, while the blue dashed line is the one derived from the 30 nearby disk galaxies of the HERACLES survey by Bigiel et al. (2011) at 1 kpc scale resolution. The red dots, corresponding to regions D1, D2, and D3 located within the disk, lie below the local galaxies relation (in blue), i.e., they show low SFE. This confirms that this galaxy is forming new stars in the disk at a very low rate (see also Vulcani et al. 2018), given that the measured values of \n\n\n\n\n\n are those typically found in galaxy bulges (Fisher et al. 2013). Lower values of \u03b1CO (5\u201310 times lower than the Milky Way value adopted here) have been found in the central part of some HERACLES galaxies (Sandstrom et al. 2013), but only in the central kiloparsec, while regions D1, D2, and D3 span a larger extent of the galaxy disk. In particular, the central kiloparsec is dominated by AGN-like line ratios (Poggianti et al. 2017a; Radovich et al. 2019), according to different indicators (Poggianti et al. 2019a), and we therefore excluded it from the calculation. Our data confirm that local conditions play an important role in determining the star formation process, as already shown in the nearby galaxy M51 (Bigiel et al. 2016), and in 29 other nearby galaxies where dense gas tracers were available (Usero et al. 2015). Ram pressure then works in JW100 by enhancing the gas density in the disk and at the same time suppressing the global SFR, resulting in long depletion times both in the tail and in the galaxy disk (as already suggested in Moretti et al. 2018a). Whether this is accompanied also by an enhancement of the molecular gas fraction, as in Nehlig et al. (2016), cannot be stated yet. Ongoing H i observations with MeerKAT will shed light on this issue. It is worth noticing, though, that in the GASP jellyfish galaxy JO206, the H i depletion time turned out to be shorter than expected (Ramatsoku et al. 2019), showing the opposite trend with respect to the molecular gas in JW100.","Citation Text":["Vulcani et al. 2018"],"Citation Start End":[[632,651]]}
{"Identifier":"2020MNRAS.498.2594S__Ahnen_et_al._2015_Instance_2","Paragraph":"In the modelling, it is assumed that the low-energy peak (from radio to optical\u2013UV) is due to synchrotron emission from ultrarelativistic electrons in the jet with an energy distribution as given by equation (2). Instead, the HE peak is due to the IC scattering of internal (SSC; Ghisellini et al. 1985; Maraschi et al. 1992; Bloom & Marscher 1996) or external photons (EIC; Sikora et al. 1994; B\u0142a\u017cejowski et al. 2000; Ghisellini & Tavecchio 2009). The IC scattering of external photons is considered, since the SEDs of FSRQS are better explained by EIC, as shown by the previous studies (e.g. Abeysekara et al. 2015; Ahnen et al. 2015; Hayashida et al. 2015b; Gasparyan et al. 2018; MAGIC Collaboration 2018), and the CD is evident in the SEDs of the considered sources (Fig. 6). Localization of the emission region in the jet is an open question and along the jet, depending on the distance from the central black hole, different photon fields can be dominant for the IC scattering (Sikora et al. 2009). In this paper, we assume that the emitting region is outside the broad-line region (BLR) where the dominant photon field is the IR emission from the dusty torus. Sikora et al. (2002) showed that in MeV blazar SEDs the shift of the peak of the HE component to lower energies is most likely due to the Comptonization of IR photons from the dusty torus. The IR radiation from the dusty torus is assumed to have a blackbody spectrum with a luminosity of LIR = 0.6\u2009Ldisc (see Ahnen et al. 2015) where Ldisc is the accretion disc luminosity, which fills a volume that for simplicity is approximated as a spherical shell with a radius of RIR = 2.5 \u00d7 1018\u2009(Ld\/1045)1\/2\u2009cm (Nenkova et al. 2008) with the energy density of $u_{\\rm IR}=0.6\\:L_{\\rm d}\/4 \\pi R_{\\rm IR}^2\\:\\delta ^2$ in the co-moving frame of the jet. In Ghisellini et al. (2011) and Marcotulli et al. (2017), the HE component in the SED of distant blazars was modelled by IC scattering of BLR reflected photons, adopting a smooth broken power-law shape of the emitting electrons. We refer the reader to these papers for details on the modelling when BLR reflected photons are considered.","Citation Text":["Ahnen et al. 2015"],"Citation Start End":[[1478,1495]]}
{"Identifier":"2020AandA...637A..59A__Massalkhi_et_al._2019_Instance_2","Paragraph":"Silicon monoxide (SiO) is predicted to be the most abundant Si-bearing molecule in the entire 1\u201310 R* range in the atmospheres of M stars. In S-type atmospheres, the calculated abundance of SiO decreases by two orders of magnitude in the 1\u20135 R* but retains a very high abundance beyond, and the same occurs in C-rich atmospheres, although in this case, the abundance drop in the 1\u20135 R* is even more pronounced (see Fig. 2; see also Ag\u00fandez & Cernicharo 2006). Observations indicate that the abundance of SiO does not differ significantly between envelopes around M-, S-, and C-type stars, although in all them the SiO abundance decreases with increasing mass-loss rate (Gonz\u00e1lez Delgado et al. 2003; Sch\u00f6ier et al. 2006; Ramstedt et al. 2009; Massalkhi et al. 2019, 2020). This decline in the SiO abundance with increasing envelope density is not a consequence of chemical equilibrium (Massalkhi et al. 2019), but has been interpreted as evidence that SiO disappears from the gas phase at high densities to be incorporated into dust grains (Gonz\u00e1lez Delgado et al. 2003; Sch\u00f6ier et al. 2006; Ramstedt et al. 2009; Massalkhi et al. 2019, 2020). It therefore appears that the gradual abundance decline calculated for SiO in the 1\u20135 R* region from stellar type in the sense M \u2192 S \u2192 C does not have a direct consequence in the SiO abundance that is injected into the expanding wind. However, this behavior predicted by chemical equilibrium probably explains why SiO masers are observed in M-type stars but not toward carbon stars (e.g., Pardo et al. 2004). Except for these details, chemical equilibrium and observations agree in the fact that SiO is one of the most abundant carriers of silicon in the atmospheres of M-, S-, and C-type stars. Calculations and observations also agree for SiS in that it is an abundant molecule regardless of the C\/O. However, observations indicate a differentiation between C- and O-rich envelopes, with SiS being on average one order of magnitude more abundant in carbon-rich sources (Sch\u00f6ier et al. 2007; Danilovich et al. 2018; Massalkhi et al. 2019, 2020). Moreover, in some oxygen-rich envelopes, the fractional abundance of SiS relative to H2 is as low as ~10\u22128, which is well below the predictions of chemical equilibrium (Danilovich et al. 2019; Massalkhi et al. 2020).","Citation Text":["Massalkhi et al. 2019"],"Citation Start End":[[886,907]]}
{"Identifier":"2022MNRAS.517.1803S__Pignatari_et_al._2013a_Instance_1","Paragraph":"Classically, the explosive \u03b1-capture at the bottom of the former convective He-shell activates a neutron burst peaking up to 1018\u201320 neutrons cm\u22123, produced by the 22Ne(\u03b1, n)25Mg reaction. This nucleosynthesis process was called the n-process (Blake & Schramm 1976; Thielemann, Arnould & Hillebrandt 1979). Much later, its signature was identified in heavy elemental abundances measured in presolar SiC-X grains such as large excesses of 58Fe, 88Sr, 96Zr, 95, 97Mo, and 138Ba (e.g. Meyer, Clayton & The 2000; Pellin et al. 2006; Pignatari et al. 2018) and in the production of the radioactive isotope 32Si (Pignatari et al. 2015). However, this scenario has been discounted as a way to explain the r-process because it could not reproduce the r-process abundance pattern in the Solar system (Blake et al. 1981). The presence of H in the He-shell during the SN explosion may lead to a suppression of the n process, where it becomes more probable to destroy 22Ne by proton capture than by the 22Ne(\u03b1, n)25Mg reaction (Pignatari et al. 2015). The production of 44Ti and 40Ca in the C\/Si zone along with the possible strong depletion of 40Ca by neutron capture predicts abundances consistent with the observation of large 44Ca\/40Ca in the C-rich grains discussed (Pignatari et al. 2013a). It is still matter of debate if in presolar grains the abundance signatures of H-ingestion coexist with the neutron-capture signature of the n-process. Indeed, the n-process activation would be compatible with the 30Si- and 32S-enhancements measured in in some putative nova grains (Liu et al. 2016; Hoppe et al. 2018), and with the 32Si- and 50Ti-enhancements found in a sample of SiC\u2013AB grains (Liu et al. 2017a). On the other hand, measurements of the Sr, Mo, and Ba isotopic abundances in SiC\u2013AB grains seem to be compatible with the pre-explosive s-process abundances in the He shell, and not with the n-process (Liu et al. 2018b). In the first case, major asymmetries triggered by the H ingestion in the pre-explosive He shell structure would be required in order to explain for the same single grain the signatures of both the n-process and of a late H-ingestion event leaving H behind in the former He shell material. While this may be a natural expectation from multidimensional simulations, it cannot be captured by 1D stellar simulations. Additionally, those abundance signatures of high neutron-density exposures could also be the effect of a local activation of the intermediate neutron-capture process (i-process; Cowan & Rose 1977), following the H ingestion event. The i-process has been proposed as a nucleosynthesis source active in massive stars at low metallicities (e.g. Roederer et al. 2016; Banerjee, Qian & Heger 2018; Clarkson, Herwig & Pignatari 2018), but it also has been proposed as a source of anomalous signatures in presolar grains for Ti isotopes (HD graphites; Jadhav et al. 2013), Ba isotopes (in some mainstream SiC grains; Liu et al. 2014) and 32Si-enrichments (in SiC AB grains; Fujiya et al. 2013). In order to verify all of these different scenarios, new abundance observations are required for more presolar grains of different types.","Citation Text":["Pignatari et al. 2013a"],"Citation Start End":[[1260,1282]]}
{"Identifier":"2021AandA...649L..15C__Woods_et_al._2002_Instance_1","Paragraph":"Since polycyclic aromatic hydrocarbons (PAHs) were first hypothesized to be carriers of the unidentified infrared bands (L\u00e9ger & Puget 1984; Allamandola et al. 1985), a great deal of effort has been devoted to understanding the chemical processes leading to the formation of these molecular species (see e.g. Joblin & Cernicharo 2018). Circumstellar envelopes around carbon-rich asymptotic giant branch (AGB) stars have been suggested as the factories of PAHs (Cherchneff et al. 1992). The detection of benzene in the carbon-rich protoplanetary nebula CRL 618 (Cernicharo et al. 2001) suggests a bottom-up approach in which the small hydrocarbons that formed during the AGB phase, such as C2H2 and C2H4, interact with the ultraviolet (UV) radiation produced by the star in its evolution to the white dwarf phase (Woods et al. 2002; Cernicharo 2004). Other hypotheses involve the processing of dust grains around evolved stars, either through UV photons (Pilleri et al. 2015) or by chemical processes (Mart\u00ednez et al. 2020). Hence, it has been surprising to see that cyanide derivatives of PAHs have been found in the cold pre-stellar core TMC-1, which is well protected against UV radiation (McGuire et al. 2018, 2021). It is unlikely that these PAH cyanides arise from a reservoir of PAHs that has existed since the early stages of the cloud since these relatively small PAHs would not have survived the diffuse cloud stage. Although McGuire et al. (2021) propose a reasonable chemical network starting with the phenyl radical that could explain the observed abundances of cyanonaphthalene, the chemical routes leading to benzene itself are still unclear. An in situ formation mechanism for benzene must involve abundant hydrocarbons containing from two to four carbon atoms. Moreover, some of these species have to permit an easy cyclization in two to three steps to have an efficient yield of benzene or phenyl radical. The propargyl radical (CH2CCH) was recently found in TMC-1 by Ag\u00fandez et al. (2021) with an abundance close to 10\u22128 relative to H2. In addition, complex hydrocarbons such as vinyl and allenyl acetylene have also been observed in very large abundances (Cernicharo et al. 2021a,b). These hydrocarbons may hold the key to the formation of initial aromatic rings, from which larger PAHs can grow.","Citation Text":["Woods et al. 2002"],"Citation Start End":[[813,830]]}
{"Identifier":"2017MNRAS.471.3699M__Cardelli_et_al._1989_Instance_1","Paragraph":"For the sources with the full JHKs photometric information (309 objects out of 507 found in the 2013 HKs catalogue), we first check the source's position in the CCD diagram. In Fig. 4(c), the solid black line represents evolutionary models, the dashed black and grey lines the locus of T-Tauri stars and the corresponding uncertainties (Meyer et al. 1997), whereas the dash\u2013dotted rectangle represents the locus of Herbig AeBe stars (Hern\u00e1ndez et al. 2005). The dotted red lines are the reddening vectors (Cardelli et al. 1989), encompassing the regions where the colours are consistent with reddened evolutionary models (region A), CTTSs (region B) and Herbig AeBe stars (region C). If the star falls in region A of the CCD, its extinction and corresponding mass are derived by de-reddening its photometry to the 1\u2009Myr isochrone in the J, (J \u2212 H) CMD. In region B, the extinction is derived by de-reddening the colours to the CTTS locus (Meyer et al. 1997). We note that the objects in region A could also be part of the CTTSs; however, the derived AV differences are typically smaller than the derived uncertainties, and for simplicity we decide to only consider the evolutionary models. The derived extinction is then used to convert the J-band photometry to the absolute J-band magnitude, which is in turn compared to the models to derive the mass. For the stars in region B, in addition to the interstellar extinction, we also correct for an excess due to the circumstellar disc or envelope, chosen randomly in the interval 0\u20130.7\u2009mag for the J band (Cieza et al. 2005). In the case that the derived mass is larger than 2\u2009M\u2299 (upper mass limit for T-Tauri stars), the procedure is repeated by de-reddening the colours to the Herbig AeBe locus (Hern\u00e1ndez et al. 2005). The intrinsic colour in this case is a randomly chosen value within the dash\u2013dotted box shown in Fig. 4(c), falling along the reddening line. The same procedure is performed if the object falls within region C. Finally, if the object falls to the left of region A, or to the right of region C, the extinction and mass cannot be derived.","Citation Text":["Cardelli et al. 1989"],"Citation Start End":[[506,526]]}
{"Identifier":"2022AandA...666A..86W__Du_et_al._2014_Instance_1","Paragraph":"A great deal of this progress is attributed to reverberation mapping (RM) campaigns since the 1980s, with the underlying principle proposed by Bahcall et al. (1972) and Blandford & McKee (1982). Photons of emission lines from structured ionized gas trace different paths to observers, leading to time lags (\u03c4) of the emission lines with respect to the ionizing photons. Such RM campaigns focusing on broad Balmer lines had detected the anticipated lags in a number of Seyfert galaxies (e.g., Peterson 1993; Peterson et al. 1998; Bentz et al. 2013; Barth et al. 2015; U et al. 2022) and quasars (e.g., Kaspi et al. 2000; Du et al. 2014, 2018a; Shen et al. 2019) over the past few decades. Along with the growing investment of observing resources and the development of analytical algorithms, the general geometry and kinematics of BLRs in some AGNs have been revealed by velocity-resolved delay analysis (e.g., Bentz et al. 2010; Denney et al. 2010; Grier et al. 2012, 2013; Du et al. 2016a; U et al. 2022), velocity-delay maps (e.g., Xiao et al. 2018; Horne et al. 2021), or dynamical modeling (e.g., Bottorff et al. 1997; Pancoast et al. 2014; Li et al. 2018; Williams et al. 2020). Many resolved BLRs have a disk-like geometry of their H\u03b2 line region (the other have inflow or outflow, or a kind of mixture of the three configurations)1. Moreover, the repeat RM observations of the same emission line and the RM results of the emission lines with different ionization in a few AGNs (e.g., NGC 5548, 3C 390.3, NGC 3783, NGC 7469, Mrk 817 etc.) approximately demonstrated the relation VFWHM\u2004\u221d\u2004\u03c4\u22121\/2, demonstrating evidence for potential of SMBHs (e.g., Peterson & Wandel 2000; Peterson et al. 2004; Lu et al. 2021), where VFWHM is the full-width-half-maximum (FWHM) of the emission lines. Considering the disk-like geometry of BLRs in some AGNs, this relation probably indicates a nearly Keplerian rotation of the disk BLRs. More recently, the GRAVITY instrument mounted in Very Large Telescope Interferometer (VLTI) spatially resolved the BLRs in several AGNs (e.g., 3C 273, NGC 3783, IRAS 09149 by GRAVITY Collaboration 2018, 2020b, 2021, respectively) and also found that their BLRs are approximately characterized by Keplerian rotating disks.","Citation Text":["Du et al. 2014"],"Citation Start End":[[620,634]]}
{"Identifier":"2022AandA...667A..68B___2020_Instance_1","Paragraph":"We combine theoretical prescriptions with realistic simulations of asteroseismic observations to construct the first magnetic detection method based on the period spacing of mixed modes. This study opens the way for an extensive search of magnetic fields inside evolved stars with solar-like oscillations. Due to the known effect of magnetic fields on mixed-mode frequencies from Bugnet et al. (2021), the magnetic signature on local measurements of the period spacing of g-modes investigated here should not get mistaken for another known physical process also leading to frequency shifts (e.g., latitudinal differential rotation, centrifugal effects, glitches, near-degeneracy effects). Such a detection of magnetic-field signatures from the measurement of the period spacing of g-modes can then be inverted to get an estimate of the radial magnetic-field strength inside RGs following Mathis et al. (2021), as it is already possible for the internal rotation (e.g., Deheuvels et al. 2017, 2020). Our methodology is well-suited for stars observed with low-inclination angles or observed pole-on, as it is based on azimuthal order m\u2004=\u20040 components to avoid additional signatures in the LS periodogram resulting from the distinct periodicity of m\u2004=\u2004\u00b11 modes. As long as the magnetic-field amplitude remains moderate, the method presented here can be applied even if the field is inclined inside the star (see Appendix F). From the statistics on the amplitudes of oscillations on the RGB (e.g. Stello et al. 2016a), intermediate-mass stars seems more likely to host strong magnetic fields than low-mass stars during the RGB. There are a few stars in the study of Gehan et al. (2020) that are observed with low-inclination angles, with a mass above 1.3 solar masses, and that present a rather low value of \u0394\u03a01 from the study of Vrard et al. (2016). The detection of magnetic fields inside these stars is out of the scope of this paper and its focus on methodology, but we are able to identify these stars as the best sample for the search of buried magnetic fields. The observational search is made possible thanks to the theoretical progress made in this study. Such a detection would be game-changing for improving the understanding of all stars of low and intermediate mass, as they are described as non-magnetic bodies in current stellar evolution models. The presence of strong internal magnetic fields inside the radiative interior of stars along their evolution would modify stellar age estimates and would thus affect many astrophysical areas \u2013 from planet habitability2 studies to galacto-archeology (Rauer et al. 2014; Chaplin et al. 2020).","Citation Text":["Deheuvels et al.","2020"],"Citation Start End":[[969,985],[992,996]]}
{"Identifier":"2018MNRAS.473.3256O__Boor_1978_Instance_1","Paragraph":"As given in equations (1)\u2013(3), \u03d5(r) and i(r) are needed when deriving the LOS model velocity at a projected sky position (x, y). Specifically, equations (2) and (3) imply that\n(4)\r\n\\begin{eqnarray}\r\nr &amp;=&amp; \\Biggl [\\biggl \\lbrace -(x-x_{C}) \\, \\sin \\phi + (y-y_{C}) \\, \\cos \\phi \\biggr \\rbrace ^2 \\nonumber \\\\\r\n&amp;&amp;+\\,\\biggl \\lbrace \\frac{(x-x_{C}) \\, \\cos \\phi \\ + (y-y_{C}) \\, \\sin \\phi }{\\cos i} \\biggr \\rbrace ^2\\Biggr ]^{\\!1\/2}.\r\n\\end{eqnarray}\r\nIf \u03d5 and i are independent of r, the latter can be directly derived from equation (4). However, if not, adequate functional forms that provide a sufficient approximation to the radial variations of \u03d5 and i should be assumed. As discussed earlier, kinematic \u03d5 and i can vary with galaxy radius due to dynamical structures in galaxies including lopsidedness, warps, bars, spiral arms and non-circular motions. The combined effect of such structures tends to result in random variations of \u03d5 and i which are not necessarily described by any specific functional form. To remove any unphysical discontinuities of \u03d5 and i and regularize their radial variations, we use the basis spline (de Boor 1978), also called the \u2018B-spline\u2019. This is a piecewise radial polynomial function of degree n where the order n is less than the number of rings in the tilted-ring model. The radial extent of the galaxy is broken up into some number of intervals where each interval has two endpoints, called \u2018breakpoints\u2019. For continuity and smoothness, these breakpoints are converted to \u2018knots\u2019 which constitute a knot vector\n(5)\r\n\\begin{equation}\r\n\\displaystyle {\\boldsymbol t} = \\{t_{0}, t_{1}, \\ldots, t_{n+k-1}\\},\r\n\\end{equation}\r\nwhere n is the number of basis splines of order k.\n(6)\r\n\\begin{eqnarray}\r\n\\displaystyle B_{m,1}(r) &amp;=&amp; \\left\\lbrace \\begin{array}{rl}1 &amp; \\qquad t_{m} \\le r < t_{m+1} \\\\\r\n0 &amp; \\qquad \\rm {otherwise} \\end{array} \\right.\r\n\\end{eqnarray}\r\n(7)\r\n\\begin{eqnarray}\r\nB_{m,k}(r) &amp;=&amp; \\frac{r-t_{m}}{t_{m+k-1} - t_{m}}B_{m,k-1}(r) \\nonumber \\\\\r\n&amp;&amp;+\\,\\,\\frac{t_{m+k} - r}{t_{m+k} - t_{m+1}} B_{m+1,k-1}(r),\r\n\\end{eqnarray}\r\nwhere m = 0, 1, \u2026, n \u2212 1. Constant, linear, quadratic\nand cubic B-splines are given by k = 1, 2, 3 and 4, respectively. The models of \u03d5 and i used in the new algorithm are given by expanding the B-spline functions as follows:\n(8)\r\n\\begin{equation}\r\n\\phi (r) = \\sum _{l=1}^{U} c^{\\phi }_{u}B^{\\phi }_{l, k}(r),\r\n\\end{equation}\r\n(9)\r\n\\begin{equation}\r\ni(r) = \\sum _{m=1}^{V} c^{i}_{v}B^{i}_{m, k}(r),\r\n\\end{equation}\r\nwhere U and V are the numbers of B-splines, and $c^{\\phi }_{m}$ and $c^{i}_{m}$ are the coefficients of the B-splines for \u03d5 and i, respectively. Similarly, the expansion velocity, vEXP, can also be modelled by the expansion of W B-spline functions,\n(10)\r\n\\begin{equation}\r\nv_{\\rm EXP}(r) = \\sum _{n=1}^{W} c^{v_{\\rm EXP}}_{w}B^{v_{\\rm EXP}}_{n,k}(r).\r\n\\end{equation}\r\n","Citation Text":["de Boor 1978"],"Citation Start End":[[1144,1156]]}
{"Identifier":"2021AandA...650A.164M__Davies_et_al._2012_Instance_4","Paragraph":"The GMC associated with G305 is one of the most massive and luminous clouds in the Milky Way (Fig. 1). It is located in the Galactic plane at l ~ 305\u00b0, b ~ 0\u00b0 and at a kinematic distance of 4 kpc (derived from a combinationof radio and H\u03b1 observationsby Clark & Porter (2004); Davies et al. (2012) measured its spectrophotometric distance to be 3.8 \u00b1 0.6 kpc and most recently Borissova et al. (2019) measured the Gaia DR2 average distance to be 3.7 \u00b1 1.2 kpc); this places it in the Scutum-Crux spiral arm. Given this distance, the complex has a diameter of ~ 30 pc (Clark & Porter 2004) and a molecular mass of ~6 \u00d7 105 M\u2299 (Hindson et al. 2010). The G305 complex consists of a large central cavity that has been cleared by the winds from massive stars belonging to two visible central clusters (Danks 1 and 2) and the Wolf-Rayet star (WR48a; Clark & Porter 2004; Davies et al. 2012). The cavity is surrounded by a thick layer of molecular gas (traced by CO and NH3 emission; Hindson et al. 2010, 2013). Radio continuum observations by Hindson et al. (2012) have revealed that the cavity is filled with ionized gas and identified six ultra-compact HII (UC HII) regions and also one bright rimmed cloud (BRC) at the periphery of the cavity, indicating molecular gas irradiated by UV radiation (Sugitani & Ogura 1994; Thompson et al. 2004), which may cause implosion (Bertoldi 1989) or evaporation. A number of studies havereported star formation tracers (water and methanol masers, HII regions and massive young stellar objects, MYSOs; Clark & Porter 2004; Lumsden et al. 2013; Urquhart et al. 2014; Green et al. 2009, 2012). Furthermore, Hindson et al. (2010) found the concentration of star formation tracers to be enhanced inside a clump of NH3 bearing molecular gas that faces the ionizing sources, which is consistent with the hypothesis that the star formation has been triggered. Analysis of the stellar clusters in the complex reveals them to have ages of 1.5 Myr for Danks 1 and 3 Myr for Danks 2,with the former possibly being triggered by the latter (Davies et al. 2012). Additionally, a diffuse population of evolved massive stars was also found to exist within the confines of the G305 complex that had formed around the same time as the two clusters (Leistra et al. 2005; Shara et al. 2009; Mauerhan et al. 2011; Davies et al. 2012; Faimali et al. 2012; Borissova et al. 2019).","Citation Text":["Davies et al. 2012"],"Citation Start End":[[2327,2345]]}
{"Identifier":"2021MNRAS.505.3514Z__Han_et_al._2002_Instance_1","Paragraph":"From the plot, we can see that there are some outliers in the sample. We will discuss these systems specifically in the following. For the system with an open circle, it is more likely to have a WD mass larger than 0.48 $\\rm M_{\\odot }$ and could have a CO WD rather than a He WD. Among these observed data, there are three systems with small orbital periods (log(Porb\/days)  \u22121.2) (open squares). For the two systems with orbital periods ($\\rm {log} (P\/\\rm {days})$) around \u20131.6, their sdB masses roughly equal to 0.3 $\\rm M_{\\odot }$ (Burdge et al. 2020). Because only these sdB progenitors with non-degenerate cores can ignite the helium when their core masses are around 0.3 $\\rm M_{\\odot }$ (Han et al. 2002, 2003). Their progenitor masses of sdB stars are likely to be larger than $2.0\\,\\,\\rm M_{\\odot }$. For these progenitors, they have a larger binding energy of envelopes leading to smaller orbital periods after CE ejection. For the system PTF1 J0823, Kupfer et al. (2017) claimed that the progenitor of the sdB star in this system is likely to have a non-degenerate core. For the three systems with minimum WD masses smaller than $0.20\\,\\,\\rm M_{\\odot }$, it is unclear how the systems formed if the WD masses are confirmed to be less than 0.2 $\\rm M_{\\odot }$. In principle, stellar objects with such a low mass could be also low-mass MS stars but the fact that no reflection effect in their light curves has been detected excludes this possibility (Maxted, Morales-Rueda & Marsh 2004; Kupfer et al. 2015). As we mentioned in the first paragraph of this section, it is possible that some WDs with masses smaller than 0.48 $\\rm M_{\\odot }$ in the sample are non-pure He WDs. For these kind of systems, the progenitor of WDs have masses larger than 2.0\u2009M\u2299 and non-degenerate He core ignition on the RGB. After their H envelopes are stripped, they can evolve into WDs with CO cores. With a binary population synthesis method, we estimate that the contribution of low mass (0.48\u2009M\u2299) non-pure He WD + sdB binaries to low mass (0.48\u2009M\u2299) WD + sdB binaries is less than $5{{\\ \\rm per\\ cent}}$.","Citation Text":["Han et al. 2002"],"Citation Start End":[[697,712]]}
{"Identifier":"2020AandA...639A..88C__Chatzistergos_et_al._2019b_Instance_3","Paragraph":"To overcome these limitations, in our previous paper (Chatzistergos et al. 2018b, Paper I, hereafter) we introduced a novel approach to process the historical and modern Ca II K observations, to perform their photometric calibration, to compensate for the intensity centre-to-limb variation (CLV, hereafter), and to account for various artefacts. By using synthetic data, we also showed that our method can perform the photometric calibration and account for image artefacts with higher accuracy than other methods presented in the literature. More importantly, we showed that, as long as the archives are consistent with each other, for example, they are centred at the same wavelength and employing the same bandwidth for the observations, the method can be used to derive accurate information on the evolution of plage areas without the need of any adjustments in the processing of the various archives (Chatzistergos et al. 2019b, Paper II, hereafter). In Paper II, we applied our method to 85 972 images from 9 Ca II K archives to derive plage areas and produce the first composite of plage areas over the entire 20th century. In particular, we analysed the Ca II K archives from the Arcetri, Kodaikanal (8-bit digitisation), McMath-Hulbert, Meudon, Mitaka, Mt Wilson, Rome\/PSPT, Schauinsland, and Wendelstein sites. Five out of the nine analysed archives were amongst the most studied and prominent ones, while the remaining archives were from less studied data sources. There are, however, many other Ca II K archives that are available and still remain largely unexplored. These archives harbour the potential to fill gaps in the available plage series as well as to address inconsistencies among the various archives and within individual archives (e.g. change in data quality, or in the measuring instrument with time). Moreover, since the work presented in Paper II, more data from various historical and modern archives became available in digital form. In particular, historical data that have been made available in the meantime include those from the latest 16-bit digitisation of the Kodaikanal archive, Catania, Coimbra, Kenwood, Kharkiv, Kyoto, Manila, Rome, Sacramento Peak, and Yerkes observatories, as well as additional data from the Meudon and Mt Wilson archives. In this light, here we present results from the most comprehensive analysis to date of historical and modern Ca II K observations taken between 1892 and 2019 from 43 different datasets for the purposes of producing a composite plage area series.","Citation Text":["Paper II"],"Citation Start End":[[1868,1876]]}
{"Identifier":"2020AandA...644A.117L__Ibarra-Medel_et_al._2016_Instance_1","Paragraph":"We made use of the Pipe3D pipeline (S\u00e1nchez et al. 2016a,b, 2018) to perform the spatially resolved stellar population analysis of the data cubes. We note that Pipe3D is based on the stellar population synthesis code FIT3D (S\u00e1nchez et al. 2016a,b), which uses the GRANADA and MILES simple stellar population (SSP) libraries in the current implementation to model and subtract the stellar spectrum and fit the emission lines. The GRANADA and MILES (gsd156) stellar library is a combination of the empirical stellar library of Vazdekis et al. (2010) and the theoretical stellar libraries of Gonz\u00e1lez Delgado et al. (2005) and Gonz\u00e1lez Delgado & Cid Fernandes (2010). The SSPs use a Salpeter (1955) initial mass function and cover 39 stellar ages (from 1 Myr to 14.1 Gyr1) and four metallicities (Z\u2004=\u20040.002, 0.008, 0.02, and 0.03). It is important to note that FIT3D performs the SSP fitting within the wavelength of 3500 \u00c5 \u20137000 \u00c5. In addition, FIT3D considers the effects on dust extinction during the stellar population synthesis using a Cardelli et al. (1989) extinction law. FIT3D provides the SSP decomposition of the modeled stellar spectra of each spatially resolved region of the galaxies. With this decomposition, we can estimate the respective SFHs and the cumulative mass temporal distributions (or mass growth histories; e.g., Ibarra-Medel et al. 2016, 2019; S\u00e1nchez et al. 2019). In order to estimate the contemporary (local and integrated) SFR values, we smoothed the SFHs within a period of 30 Myr. These quantities are independent of SFR estimations that use the H\u03b1 emission line flux. For star-forming galaxies, both estimates are very similar, showing that the SFR calculated in the last 30 Myr is related to massive young stars, contributing to the H\u03b1 emission. For retired galaxies, the SFR calculated from the H\u03b1 flux is likely overestimated because this flux is dominated by effects of post-AGB stars rather than of young stars (see, e.g., Bitsakis et al. 2019; S\u00e1nchez et al. 2019). However, in these galaxies, the SFR derived from the stellar population synthesis with Pipe3D should also be considered as an upper limit. To fit the emission lines, Pipe3D subtracts the modeled stellar spectra and fits a series of Gaussian profiles for each line (S\u00e1nchez et al. 2016b). With those fits, Pipe3D provides the total flux, dispersion, and the kinematics of each emission line. The quality of the emission line fits of Pipe3D (and other codes) was tested using MUSE data (S\u00e1nchez-Menguiano et al. 2018; Bellocchi et al. 2019; L\u00f3pez-Cob\u00e1 et al. 2020). With the use of the adjacent continuum windows (from the modeled stellar spectra) for each emission line, Pipe3D estimates the stellar continuum flux and provides the equivalent width of each emission line.","Citation Text":["Ibarra-Medel et al. 2016"],"Citation Start End":[[1337,1361]]}
{"Identifier":"2016ApJ...823L...1L___2013a_Instance_1","Paragraph":"The observed QPOs are generally attributed to oscillations of the star excited by an explosion of magnetic origin that creates the flare. The oscillating stellar surface should modulate the charge density in the magnetosphere, creating variations in the optical depth for resonant Compton scattering of the hard X-rays that accompany the flare (Timokhin et al. 2008; D\u2019Angelo & Watts 2012). In this connection, the problem of finding the oscillatory modes for a strongly magnetized neutron star has received much attention, and has proven to be a formidable problem. To make the problem tractable, most theoretical treatments of the QPO problem have assumed smooth field geometries, usually dipolar or variants (e.g., Glampedakis et al. 2006; Levin 2006, 2007; Sotani et al. 2008a, 2008b; Cerd\u00e1-Dur\u00e1n et al. 2009; Colaiuda et al. 2009; Colaiuda & Kokkotas 2011; Gabler et al. 2011, 2012, 2013a, 2013b, 2014; van Hoven & Levin 2011, 2012; Passamonti & Lander 2013). Smooth field geometries support a problematic Alfv\u00e9n continuum that couples to the discrete natural spectrum of the crust. As pointed out by Levin (2006), if energy is deposited in the crust at one of the natural frequencies of the crust, and this frequency lies within a portion of the core continuum, the energy is lost to the core continuum in less than 0.1 s as the entire core continuum is excited. The crust excitation is effectively damped through resonant absorption, a familiar process in MHD; see, e.g., Goedbleod & Poedts (2004). The problem has been addressed by assuming field geometries with gaps in the Alfv\u00e9n continuum. Under this assumption, long-lived quasi-normal modes can exist inside the gaps or near the edges of the Alfv\u00e9n continuum. van Hoven & Levin (2011) showed for a \u201cbox\u201d neutron star that introduction of a magnetic tangle breaks the Alfv\u00e9n continuum. Link & van Eysden (2015) showed that for magnetic tangling in a spherical neutron star the problematic Alfv\u00e9n continuum disappears.3\n\n3\nSotani (2015) added general relativity in the treatment of the magnetic field, and confirmed some of the results of Link & van Eysden (2015).\n They found that the star acquires discrete normal modes, and quantified the mode spacing. It is clear from these investigations that the unknown magnetic field geometry is the most important ingredient in determining the oscillation spectrum of a magnetar.","Citation Text":["Gabler et al.","2013a"],"Citation Start End":[[862,875],[888,893]]}
{"Identifier":"2022AandA...663A.127L__Guo_et_al._2019_Instance_1","Paragraph":"To understand the 3D magnetic properties of the FC observed here, we reconstructed the coronal magnetic field of the source region based on the nonlinear force-free field (NLFFF) assumption using the HMI photospheric vector magnetogram as the bottom boundary (Yan et al. 2001; Canou et al. 2009; Wiegelmann & Sakurai 2012; Yang et al. 2016b; Zhong et al. 2019; Qiu et al. 2020). We first tried the extrapolation method (Guo et al. 2010; Zhu et al. 2016) and found that it is difficult to reproduce the magnetic field comparable with the morphology of the observed filament threads. We then resorted to the magnetic flux rope (MFR) embedding method based on regularized Biot-Savart laws (RBSL; Titov et al. 2014, 2018; Guo et al. 2019). The procedure was divided into four steps. First, we took advantage of the 304 \u00c5 images at 05:24 UT to derive the path of the FC, as shown in Fig. 7a. According to previous statistics (Tandberg-Hanssen 1995; Filippov & Den 2000; Engvold 2015), the height of active region filaments ranges from 5 Mm to 30 Mm. We therefore set the height of the MFR axis to 30 Mm and the MFR minor radius to 25 Mm, assuming that the lower half of the MFR is fully filled by cool material. The 3D path of the MFR was estimated according to Guo et al. (2021a). Second, we calculated a potential field utilizing the normal magnetic field component, where the projection effect was corrected (Wiegelmann et al. 2006; Guo et al. 2017). Third, we set the physical parameters of the RBSL model including the average value of unsigned magnetic flux at the two MFR footprints (1.72\u2005\u00d7\u20051021 Mx) and the strength of the electric current following Eq. (12) in Titov et al. (2018). Finally, we inserted the MFR derived by the RBSL model into the potential field along the path of the FC and then performed a relaxation using the magnetofrictional code (Guo et al. 2016a,b). After relaxation, the force-free metric was \u03c3J\u2004=\u20040.28, and the divergence-free metric was \u27e8|fi|\u27e9\u2004=\u20041.47\u2005\u00d7\u200510\u22124, which were sufficiently small and generally acceptable according to Guo et al. (2021a)1.","Citation Text":["Guo et al. 2019"],"Citation Start End":[[718,733]]}
{"Identifier":"2017MNRAS.464..840P__Tauris_2001_Instance_1","Paragraph":"The theory of angular momentum evolution due to mass transfer via Roche lobe overflow and gravitational wave radiation is well established and leads to quantitative predictions (Landau & Lifschitz 1958; Paczy\u0144ski 1967; Faulkner 1971; Verbunt 1993). However, angular momentum losses due to emission of gravitational waves only dominate at short orbital period (Porb \u2272 2 h), while other mechanisms, such as magnetic braking, are invoked to explain the orbit evolution at longer periods. The theory at the heart of the magnetic braking model is based on the same physical principles through which magnetic stellar winds are known to decelerate the rotation of low-mass stars (e.g. Kraft 1967; Skumanich 1972; Soderblom 1983). In synchronized binaries (typically the case of LMXBs), the angular momentum loss from the companion star is continuously re-distributed to the angular momentum of the whole binary system, therefore affecting its orbital period and binary separation (e.g. Tauris 2001). Large uncertainties are involved in the extrapolation of magnetic braking from isolated low-mass stars to the case of synchronized companion stars in LMXBs (e.g. see discussion in Knigge, Baraffe & Patterson 2011). Indeed, despite the large allowed range of model parameters (possibly translating into more than one order of magnitude on efficiency of angular momentum removal), magnetic braking often fails to reproduce measured values of the orbital period derivative, at least in the case of conservative mass transfer (Di Salvo et al. 2008; Hartman et al. 2008, 2009; Hu, Chou & Chung 2008; Burderi et al. 2010; Jain, Paul & Dutta 2010; Gonz\u00e1lez Hern\u00e1ndez, Rebolo & Casares 2012, 2014). Additional mechanisms, typically involving outflows, are generally necessary to explain these discrepancies (Di Salvo et al. 2008; Hartman et al. 2008, 2009; Burderi et al. 2010). We also note that some of the best-monitored sources (i.e. EXO 0748\u2212676; Wolff et al. 2009) show erratic variations of the orbital period with amplitudes of \u0394P\/P \u223c 10\u22125 and possible trends recurring on time-scales of years-to-decades, such as those observed in e.g. Algols systems (Kreiner & Ziolkowski 1978; Hall 1989). These variations are thought to be the consequence of the torque provided by the magnetic activity of a sub-surface magnetic field in the companion star with convective envelope. The latter induces a cyclic exchange of angular momentum between the inner and outer parts of the companion star causing a change in the gravitational quadrupole moment (Applegate 1992; Applegate & Shaham 1994; Lazaridis et al. 2011).","Citation Text":["Tauris 2001"],"Citation Start End":[[979,990]]}
{"Identifier":"2018MNRAS.480.3243K__Pounds_et_al._1990_Instance_1","Paragraph":"Several observational pieces of evidence confirmed the existence of a relativistically blurred reflection component in AGNs. In particular, X-ray spectroscopy has proven to be a powerful tool to identify the reflection features in AGNs X-ray spectra, allowing us to probe the innermost regions of AGNs (e.g. Fabian et al. 2009). Additionally, this technique may provide estimates of black hole spins (e.g. Risaliti et al. 2013; Marinucci et al. 2014; Walton et al. 2014). This was achieved thanks to the high-quality spectra provided by XMM\u2013Newton (Jansen et al. 2001) and NuSTAR (Harrison et al. 2013) in the 0.3\u201380\u2009keV range. The first observational X-ray feature associated with relativistic effects was the anomalous shape of the aforementioned iron K \u03b1 line detected in type-1 AGNs, where the observer has a direct view of the central engine through the polar direction of the system (e.g. Matsuoka et al. 1990; Pounds et al. 1990). The shape of the iron line, with an extended red wing spanning over several keV was soon associated with special and general relativistic effects blurring the signal (e.g. Tanaka et al. 1995; Iwasawa et al. 1996; Nandra et al. 1997). By fitting the observed iron line with relativistic models, it became possible to determine the spin of the black hole (BH), its mass and inclination (see e.g. Miller 2007, for a review). Yet, an alternative interpretation, based on partial covering absorption, has been proposed in order to explain the apparent red wing of the Fe line and the spectral curvature at hard X-rays (e.g. Miller, Turner & Reeves 2008, 2009). The two scenarios have different advantages: on the one hand, blurred reflection is able to explain the spectral and timing properties of accreting systems for a wide range of BH mass. On the other hand, the Compton-thin to Compton-thick (and vice versa) rapid transitions that are observed in \u201cchanging-look\u201d AGNs (Matt, Guainazzi & Maiolino 2003) is suggestive that partial covering should be also taken into consideration. In addition, several occultation events, associated with both Compton-thin and Compton-thick clouds in the broad-line region (BLR), have been reported in AGNs (e.g. Risaliti et al. 2007, ; Nardini & Risaliti 2011; Sanfrutos et al. 2013; Torricelli-Ciamponi et al. 2014). Markowitz, Krumpe & Nikutta (2014) estimated the probability to observe an X-ray eclipse (of any duration between 0.2 d and 16 yr) in a given source to be in in the ranges 0.003\u20130.166 and 0.039\u20130.571, for type I and type II AGNs, respectively. By studying X-ray eclipses, it might be possible to constrain the importance of partial obscuration, together with the geometry and location of the distant obscuring clouds. This is particularly relevant since obscuration events from BLR clouds do not affect only the AGNs light curves but they show also a strong impact on their spectroscopic and polarimetric properties. Risaliti et al. (2011) investigated the effects of successive eclipses of the receding and approaching parts of the accretion disc on the shape of the iron line. In fact, due to special and general relativistic effects, obscuring various parts of the disc will result in a variability in the profile of the observed emission line, which provides a new probe of the innermost regions of the disc. In a similar fashion, Marin & Dov\u010diak (2015) explored the effects of such events on the polarimetric signal. The authors showed that eclipses induce a variability in the polarization signal due to the covering of different parts of the disc emitting a non-uniformly polarized light, mainly due to relativistic effects.","Citation Text":["Pounds et al. 1990"],"Citation Start End":[[917,935]]}
{"Identifier":"2020MNRAS.499L..31P__Zehavi_et_al._2011_Instance_1","Paragraph":"The origin of the observed colour bimodality must be explained by successful models of galaxy formation. Considerable efforts have been directed towards reproducing the observed distribution of galaxy colours using different semi-analytic models of galaxy formation (Menci et al. 2005; Cattaneo et al. 2006, 2007; Driver et al. 2006; Cameron et al. 2009; Trayford et al. 2016; Nelson et al. 2018; Correa, Schaye & Trayford 2019). A significant number of studies have also been devoted to understand the spatial distribution of red and blue galaxies using various clustering measures like two-point correlation function (Zehavi et al. 2011), three-point correlation function (Kayo et al. 2004), genus (Hoyle et al. 2002), filamentarity (Pandey & Bharadwaj 2006), and local dimension (Pandey & Sarkar 2020). The studies of mass function of red and blue galaxies (Drory et al. 2009; Taylor et al. 2015) play a crucial role in guiding the theories of galaxy formation and evolution. Any such analysis requires one to classify the red and blue galaxies using some operational definition. But colour is a subjective judgment and it is difficult to classify galaxies according to their colours in an objective manner. Also, one should remember that it is the observed bimodal distribution of galaxy colours that motivates us to classify the galaxies according to their colours. In reality, no galaxies can be regarded as truly either \u2018red\u2019 or \u2018blue\u2019 based on their colours. Strateva et al. (2001) prescribed that the red and blue galaxies can be separated by using an optimal colour separator (u \u2212 r) = 2.22. However, a substantial overlap between the two populations prohibit us to objectively define galaxies as \u2018red\u2019 and \u2018blue\u2019 based on their colours. The boundary separating the two populations remains arbitrary and currently there are no meaningful theoretical arguments favouring any specific cut over others. Applying a hard cut to define samples of red and blue galaxies is expected to introduce significant contamination in these samples. Previously, Coppa et al. (2011) used an unsupervised fuzzy partition clustering algorithm applied on the principal components of a PCA analysis to study the bimodality in galaxy colours. Another study by Norman et al. (2004) use Bayesian statistical methods instead of sharp cuts in parameter space for such classification problems.","Citation Text":["Zehavi et al. 2011"],"Citation Start End":[[620,638]]}
{"Identifier":"2019ApJ...876..132N__Hopkins_2012_Instance_1","Paragraph":"Recently, dust-obscured galaxies (DOGs, Dey et al. 2008) shed light on this issue, because SMBHs in DOGs are expected to be rapidly growing during the coevolution. DOGs are originally defined as galaxies that are bright in mid-infrared (MIR), while faint in optical. Specifically, \n\n\n\n\n\n (i.e., \n\n\n\n\n\n; Dey et al. 2008; Fiore et al. 2008). In the context of the gas-rich major-merger scenario, it is expected that the SF phase evolves into the AGN phase because the merging event leads to the active SF, while the gas accretion onto the nucleus caused by such a merger requires some time (see, e.g., Davies et al. 2007; Hopkins 2012; Matsuoka et al. 2017). Because such active galaxies are expected to be heavily surrounded by dust, DOGs potentially correspond to galaxies in the SF phase or AGN phase (Dey et al. 2008; Hopkins et al. 2008). DOGs are classified into two subclasses according to their spectral energy distribution (SED): \u201cbump DOGs\u201d and \u201cpower-law (PL) DOGs\u201d (Dey et al. 2008). The bump DOGs show a rest-frame 1.6 \u03bcm stellar bump in their SEDs, while the PL DOGs show a power-law feature on their SEDs. Therefore, it is considered that the bump DOGs correspond to galaxies in SF mode (Desai et al. 2009; Bussmann et al. 2011), while the PL DOGs correspond to galaxies in the AGN phase (Fiore et al. 2008; Bussmann et al. 2009; Melbourne et al. 2012). The fraction of PL DOGs among all DOGs increases with increasing MIR flux density (e.g., Dey et al. 2008; Toba et al. 2015), which is similar to the behavior of the luminosity dependence of the AGN fraction in ultraluminous infrared (IR) galaxies (ULIRGs; see Sanders & Mirabel 1996 for a review). The comoving number density of DOGs shows its peak at z \u223c 1\u20132 (e.g., Dey et al. 2008; Toba et al. 2017), which corresponds to the peak of SF rate density and the growth rate of SMBHs (e.g., Richards et al. 2006; Madau & Dickinson 2014). This strongly suggests that DOGs are related to the most active objects in terms of the coevolution between galaxies and SMBHs. In this sense, DOGs with a high IR luminosity potentially harbor a rapidly growing SMBH and are, therefore, important for understanding the coevolution of galaxies and SMBHs.","Citation Text":["Hopkins 2012"],"Citation Start End":[[620,632]]}
{"Identifier":"2018ApJ...853...50F__Bernard_et_al._2015b_Instance_2","Paragraph":"However, using the well-assessed new post-AGB evolutionary models, we confined the main-sequence ages of our halo sample to be mostly \u223c2\u20135 Gyr, with the oldest being \u223c6\u20138 Gyr, while the outer-disk sample are mostly \u22721\u20134 Gyr. We thus conjecture that our targets probably formed prior to the encounter with M33. Obviously, our sample represents the population that is different from the underlying, smooth, extended (and mostly metal-poor) halo component of M31 (Ibata et al. 2007, 2014), which was formed through the repeated accretion of smaller galaxies in the distant past. These bright PNe seem to resemble the younger, metal-rich population in the outer stream of M31, as revealed by HST pencil-beam pointings on the Giant Stream (Brown et al. 2006a; Bernard et al. 2015b). The metallicity of the stream fields was enriched continuously from [Fe\/H] \u223c \u22121.5 to at least solar level about 5 Gyr ago (Bernard et al. 2015b). This timeline of metal enrichment is generally consistent with the stellar ages of our metal-rich sample. N-body simulations suggested that the Giant Stream and other stream-like features in the halo are debris of a massive (\u2273109\u2013\n\n\n\n\n\n) progenitor that was recently disrupted during the course of a merger (e.g., Ibata et al. 2004; Fardal et al. 2006, 2007, 2008, 2013; Font et al. 2006; Geehan et al. 2006; Mori & Rich 2008; Sadoun et al. 2014). The extended star formation history and the broad range of metallicity (\u22121.5 \u2272 [Fe\/H] \u2272 0.2) discovered in the stream fields can be explained by a disk galaxy progenitor (Brown et al. 2006a, 2006b; Bernard et al. 2015b). If the stellar streams in M31's halo indeed have a common origin, our sample of halo PNe then probably formed through extended star formation in this possibly massive, disk-like progenitor. Moreover, some simulations predict that the remnant of the disrupted satellite resides in the NE Shelf (e.g., Fardal et al. 2008, 2013; Sadoun et al. 2014); PN17 in our sample is located in this region and might be associated with this substructure (see Section 4.4).","Citation Text":["Bernard et al. 2015b"],"Citation Start End":[[901,921]]}
{"Identifier":"2019AandA...624A..15S__Petrovich_2015_Instance_1","Paragraph":"As described in Sect. 4.1.1, the TTV inversion code includes a Hill stability criterion for two-planet systems. Although this initial check filters out the least stable systems, it cannot guarantee the long-term stability of the derived orbital architecture of a four-planet system. On one hand, there is no analytical criterion for assessing the long-term stability of multi-planet systems (>2 planets) like that for two-planet systems (Gladman 1993; Chambers et al. 1996). Chambers et al. (1996), Smith & Lissauer (2009), Lissauer et al. (2011b), and similar numerical studies have found that long-term stability of multi-planet systems typically requires the mutual separations between planets to be at least ten mutual Hill radii, which is much larger than the cautious limit of \n\n$2\\sqrt{3}$\n\n\n2\n3\n\n\n\n mutual Hill radii required by the Hill stability criterion that is implemented in our TTV inversion code (Gladman 1993; Chambers et al. 1996). Meanwhile, these criteria are only valid under the assumptions of low mutual inclinations and small eccentricities. These two restrictions are usually not well-defined; limiting values of 1\u00b0\u20132\u00b0 and e = 0.1\u20130.2, respectively, are often adopted (Petrovich 2015; MacDonald et al. 2016). Apparently, these restrictions and stability criteria are well satisfied by the nominal orbital architecture of the Kepler-411 system inverted from the TTV data. On the other hand, the Hill criterion provides no information about the long-term behavior of the system, and repeated weak interactions between planets in Hill stable orbits may still lead to ejections and\/or physical collisions; these are referred to as Lagrange unstable (Petrovich 2015). The chaotic orbits that are generated primarily by first-order (or higher-order) resonance overlap are eventually subjected to large-scale variation of the semi-major axes and hence become Lagrange unstable. Based on previous studies on the chaotic behavior induced by the first-order resonance overlap (e.g. Wisdom 1980; Duncan et al. 1989), Deck et al. 2013 supply the condition that leads to chaotic behavior in a two-planet system (see Eq. (50) in Deck et al. 2013). We apply this criterion to estimate the nominal orbital architecture extracted from the TTV data of the Kepler-411 system, and find that the orbital configuration is far away from the chaotic motion. Therefore, we conclude that the nominal orbital architecture of the Kepler-411 system extracted from measured TTVs satisfies both the Hill and Lagrange stability criteria.","Citation Text":["Petrovich 2015"],"Citation Start End":[[1194,1208]]}
{"Identifier":"2018AandA...610A..44M__Kr\u00fcger_&_Dreizler_(1992)_Instance_4","Paragraph":"The first investigations of the rotational spectra of ethyl isocyanide were carried out in 1966 by Bolton et al. (1966). The spectra of the first vibrational and torsional excited states were measured in the centimeter wave domain (Anderson & Gwinn 1968). In this initial study, the dipole moment was determined to be \u03bca = 3.79 D and \u03bcb = 1.31 D; this value is usually large for a molecule that includes a CN group. This causes dense and intense rotational spectra in the millimeter wave range and also in the submillimeter wave range up to 900 GHz (bQ lines). Anderson & Gwinn (1968) also observed some A\u2013E splittings due to the internal rotation motion of the methyl group. The most recent spectroscopic study is from Kr\u00fcger & Dreizler (1992) who reinvestigated the internal rotation measurements and also determined hyperfine coupling parameters due to the nitrogen quadrupole. As in our previous studies of ethyl cyanide isotopologs, it was not possible to observe internal rotation and hyperfine splittings due to our Doppler limited resolution. Our analysis was rather easy, starting from a prediction based on Kr\u00fcger & Dreizler (1992) parameters. First, we analyzed and fit the most intense transitions, the aRh transitions, up to 330 GHz. These transitions were shifted only a few MHz from the initial predictions. Then bR and bQ lines were searched and included in the fit up to 330 GHz. Next, all the spectra were analyzed up to 990 GHz without difficulty. For the fitting, we employed ASFIT (Kisiel 2001) and predictions were made with SPCAT (Pickett 1991). The global fits included 6 transitions from Anderson & Gwinn (1968), 29 lines from Kr\u00fcger & Dreizler (1992), and 2906 from this work. The maximal quantum numbers are J = 103 and Ka = 30. Both reductions A and S were tested. A reduction permits us to check theagreement of our new parameters set with those from Kr\u00fcger & Dreizler (1992) (Table 1). Using S reduction slightly decreases root mean square from 30.3 to28.7 kHz. The condition numbers are nearly the same: 295 and 310 for the A and S reductions, respectively.The A reduction requires 23 parameters, but 5 additional parameters are required for the S reduction (Table 2). For this reason we used the A reduction even if this molecule is close to the prolate limit with kappa = \u22120.9521. Part of the new measurements are in Table 3. Owing to its large size, the complete version of the global fit Table S1 is supplied at the CDS. The fitting files .lin (S2), .par (S3), and the prediction .cat (S4) are also available at CDS.","Citation Text":["Kr\u00fcger & Dreizler (1992)"],"Citation Start End":[[1880,1904]]}
{"Identifier":"2021ApJ...922...78X__Spitler_et_al._2014_Instance_1","Paragraph":"Fast radio bursts (FRBs) are bright, cosmological origin, and millisecond-duration bursts in radio wavelengths (Lorimer et al. 2007; Thornton et al. 2013; Bassa et al. 2017; Macquart et al. 2020). After the discovery of the first FRB (Lorimer et al. 2007), a number of dedicated facilities have been conducted to search FRBs, such as the Parkes telescope (e.g., Bhandari et al. 2018), the updated Molonglo Observatory Synthesis Telescope (e.g., Farah et al. 2018), the Australian Square Kilometre Array Pathfinder (e.g., Shannon et al. 2018), the Canadian Hydrogen Intensity Mapping Experiment (CHIME; The CHIME\/FRB Collaboration et al. 2018), the Deep Synoptic Array (Kocz et al. 2019; Ravi et al. 2019), the Green Bank Telescope (Masui et al. 2019), Arecibo (Spitler et al. 2014; Patel et al. 2018), and the Five-hundred-meter Aperture Spherical radio Telescope (FAST; Nan et al. 2011; Li et al. 2019). All these efforts result in an increasing rate of new FRB detections. Among them, more than 20 repeating FRBs have been reported. Particularly, the physical origin of the repeating FRB 20121102A was identified to be with a low-metallicity star-forming dwarf galaxy at a redshift 0.19273 (Bassa et al. 2017; Tendulkar et al. 2017). Another repeating FRB 20190523A was found to be associated with a more massive but low specific star formation rate (Ravi et al. 2019). The identification of the counterpart of the brightest radio bursts from SGR 1935+2154 as a magnetar in our Galaxy by HXMT (Li et al. 2021) and INTEGRAL (Mereghetti et al. 2020) with short-duration X-ray bursts suggests that at least a fraction of FRBs are connected with newborn magnetized neutron stars (e.g., Weltman & Walters 2020; Zhang 2021). More bursts with similar characteristics need to be detected in the future to confirm this conclusion. Recently, a new large sample with 535 FRBs was presented by CHIME\/FRB Collaboration et al. (2021) that were detected by the CHIME survey, including 61 bursts from 18 previously reported repeating sources and 474 one-off bursts. Though an increasing catalog of theories and models is developing to explain the physical nature of FRBs (e.g., see the review of Platts et al. 2019; Xiao et al. 2021), the origin of FRBs remains a mystery.","Citation Text":["Spitler et al. 2014"],"Citation Start End":[[761,780]]}
{"Identifier":"2021MNRAS.508..637S__Krause_et_al._2016_Instance_1","Paragraph":"It is now well established that intrinsic alignments exist in the Universe, and must be accounted for at some level to avoid biasing cosmological analyses based on cosmic shear and galaxy\u2013galaxy lensing. IAs have been included in cosmic shear analyses for as long as shear has been a competitive cosmological probe (Heymans et al. 2013; Dark Energy Survey Collaboration 2016; Jee et al. 2016; Troxel et al. 2018; Chang et al. 2019; Hikage et al. 2019; Hamana et al. 2020; Hildebrandt et al. 2020; Asgari et al. 2021). Only recently, however, have the lensing data been of sufficient volume to potentially incur biases due to model insufficiency (see Krause et al. 2016, and the Stage IV forecasts of Fortuna et al. 2021 and the tests in the context of DES Y3 Secco, Samuroff et al. 2021; see also Joachimi et al. 2020 for an interesting counter discussion). Developing a fuller understanding of intrinsic alignments, then, will be crucial for, arguably, the current generation of cosmological surveys, and certainly the next. This paper is one small step in this direction, providing the first detailed analysis at the level of model constraints on the best available cosmological hydrodynamic simulations. Our results, of course, come with a number of caveats. Most notably, our selection function is not intended to accurately match current or future lensing surveys. This is in part because recreating the complex redshift-dependent selection function in a real lensing sample, which would typically be based on a number of correlated observables, is a difficult task; it is also, however, a function of our aim in this study. We wish to understand the behaviour of IAs at a physical level, in order to feed into understanding IAs and model building efforts, rather than make a detailed prediction or robustness test for a particular survey. The behaviour of intrinsic alignments on small physical scales is an important topic for future investigation, and one that could conceivably be addressed using hydrodynamic simulations; indeed, due to the larger number of measureable modes, the signal to noise on small scales is relatively high. The TATT approach allows some hope of pushing to smaller scales (though not into the regime of \u223c1\u2009h\u22121\u2009Mpc, where one would need an explicit model for 1h alignment contributions). Unfortunately, a number of other poorly understood effects enter on small scales, particularly non-linear galaxy bias and baryonic physics. In order to pursue IA constraints on such scales, it is likely that one would need to consider both higher order bias terms and the interplay with the higher order IA terms.","Citation Text":["Krause et al. 2016"],"Citation Start End":[[650,668]]}
{"Identifier":"2020AandA...637A..44N__Kraus_(2018)_Instance_2","Paragraph":"Among the existing IACT systems, HESS has the largest FoV and hence provides the highest sensitivity for the diffuse \u03b3-ray flux. Its electron spectrum analysis technique could be directly used to obtain a measurement of the diffuse Galactic \u03b3-ray flux above energies of several TeV in the Galactic Ridge (|l|  30\u00b0, |b|  2\u00b0) region; see Figs. 3 and 4. A multi-year exposure of HESS could be sufficient for detection of the diffuse emission even from regions of higher Galactic latitude. This is illustrated in Fig. 5, where thick blue and red data points show mild high-energy excesses of the electron spectra derived by Kraus (2018), Kerszberg et al. (2017), Kerszberg (2017) over broken power-law models derived from the fits to lower energy data. Comparing these excesses with the level of the IceCube astrophysical neutrino flux and with the Fermi\/LAT diffuse sky flux from the region |b| > 7\u00b0 (corresponding to the data selection criterium of HESS analysis Kerszberg et al. 2017; Kerszberg 2017) we find that the overall excess flux levels are comparable to expected diffuse \u03b3-ray flux from the sky region covered by the HESS analysis (the quoted systematic error on the electron flux is \u0394log(EFE) \u2243 0.4). The overall excesses within 805 and 1186 h of HESS exposures (Kraus 2018; Kerszberg 2017) are at the levels of >4\u03c3 for the analysis of Kraus (2018) and 1.7\u03c3 for the analysis of Kerszberg (2017). A factor-of-ten longer exposure (which is potentially already available with HESS) could reveal a higher significance excess at the level of up to 5\u03c3. Such an excess is predicted in a range of theoretical models including interactions of cosmic rays injected by a nearby source (Andersen et al. 2018; Neronov et al. 2018; Bouyahiaoui et al. 2019) or decays of dark matter particles (Berezinsky et al. 1997; Feldstein et al. 2013; Esmaili & Serpico 2013; Neronov et al. 2018) or a large-scale cosmic ray halo around the Galaxy (Taylor et al. 2014; Blasi & Amato 2019).","Citation Text":["Kraus 2018"],"Citation Start End":[[1272,1282]]}
{"Identifier":"2021MNRAS.501.2112S__Ford_et_al._2014_Instance_1","Paragraph":"This work improves the efficiency of component by component modelling that has been successful in recovering the physical conditions for various individual absorbers (e.g. Churchill & Charlton 1999; Charlton et al. 2000, 2003; Ding et al. 2003a, 2003b; Zonak et al. 2004; Ding, Charlton & Churchill 2005; Masiero et al. 2005; Lynch & Charlton 2007; Misawa et al. 2008; Jones et al. 2010; Lacki & Charlton 2010; Muzahid et al. 2015; Richter et al. 2018; Rosenwasser et al. 2018). Rather than averaging over components and phases, it is possible to determine how much of the H\u2009i is associated with these different phases in order to derive separate metallicities for various clouds. Resolving the individual clouds allows us to break the degeneracy for components on the flat part of the Ly\u03b1 curve of growth, even with coverage of just saturated H\u2009i lines, and derive metallicity constraints for different parcels of gas along the line of sight. It is important to do so because different processes, e.g. outflows (Bouch\u00e9 et al. 2012; Bordoloi et al. 2014; Rubin et al. 2014; Schroetter et al. 2016), pristine accretion (Martin et al. 2012; Rubin et al. 2012; Danovich et al. 2015), recycled accretion (Ford et al. 2014), minor and major mergers (Martin et al. 2012; Angl\u00e9s-Alc\u00e1zar et al. 2017), are surely contributing to the same system, and it is expected that conditions will vary significantly along a line of sight which can span hundreds of kpc spatially (Churchill et al. 2015; Peeples et al. 2019). This will lead to a more meaningful comparison to galaxy properties. For example, Pointon et al. (2019) did not find a difference between the metallicities of absorbers found within an impact parameter of 200 kpc along the major and the minor axes of isolated galaxies. Based on cosmological hydrodynamic simulations, a larger metallicity is expected along the minor axis due to outflows and a lower metallicity along the major axis due to inflows (Peroux et al. 2020). However, an observational trend could exist, for example, for the minor axis to have some, but not all, high metallicity components, or for the minor axis to have one or more low metallicity components. Such results would be \u2018washed out\u2019 by deriving an average metallicity for all gas along a line of sight, which clearly often has multiple complex origins. For some data sets\/projects, the new analysis could be transformative, however, to make it feasible to use for large statistical studies it is important that the analysis is semi-automated and robust.","Citation Text":["Ford et al. 2014"],"Citation Start End":[[1201,1217]]}
{"Identifier":"2021ApJ...923..159M__Rubin_et_al._2019_Instance_1","Paragraph":"Sulfur-bearing molecules constitute a fundamental branch of prebiotic chemistry due to the recently gained importance of cysteine C3H7NO2S as a catalyst in the assembly of complex peptides (Foden et al. 2020). In cold, interstellar environments, cysteine is far from being detected, but other complex organic molecules (COMs) of prebiotic importance bearing S have recently been detected, highlighting ethyl mercaptan C2H5SH or thioformic acid (HC(O)SH; Kolesnikov\u00e1 et al. 2014; Rodr\u00edguez-Almeida et al. 2021). In turn, the family of compounds to which HC(O)SH belongs (thioacids) has also been pointed out as possible catalysts in prebiotic processes in a primordial Earth (Chandru et al. 2016). Particularly puzzling is the fact that the presence of these molecules on Earth could have a (at least partial) panspermic origin, owing to both a confirmed presence of sulfur-bearing species in comets (Korth et al. 1986; Krishna Swamy & Wallis 1987; Calmonte et al. 2016; Rubin et al. 2019) and the fact that the simplest thioacid (thioformic acid) has been recently identified in the interstellar medium (ISM; i.e., the giant molecular cloud G+0.693; see Rodr\u00edguez-Almeida et al. 2021). The chemical mechanism behind the formation of thioformic acid is unknown, but such a mechanism needs to explain the formation of HC(O)SH from simple, sulfur-bearing precursors. In Rodr\u00edguez-Almeida et al. (2021) several possible routes were postulated for the formation of HC(O)SH, namely,\n1\n\n\n\nCO+SH\u2192HSCO\u2192HHC(O)SH,\n\n\n\n2\n\n\n\nHCO+SH\u2192HC(O)SH,\n\n\n\n3\n\n\n\nOCS+2H\u2192HC(O)SH.\n\nIn this work, we have focused on reaction (3) on interstellar dust grains, having two main objectives in mind. First, we wish to check whether or not the reaction proceeds effectively, and, second, we wish to determine the stereoisomerism of the reaction. The study of reaction (3) sparked our interest for two main reasons. First, OCS is a moderately abundant molecule in the astronomical source where HC(O)SH was detected (Armijos-Abenda\u00f1o et al. 2014; Rodr\u00edguez-Almeida et al. 2021); second OCS is the only sulfur-bearing species positively detected on interstellar ices (Palumbo et al. 1997).","Citation Text":["Rubin et al. 2019"],"Citation Start End":[[970,987]]}
{"Identifier":"2020AandA...644A..11L__Reissl_et_al._2016_Instance_1","Paragraph":"One major issue we face with the ALMA dust polarization observations is spatial filtering by the interferometer, which removes the scales of emission that are not included in the uv -coverage of the dataset. In contrast to the statistical analysis of dust polarization performed with single dish instruments such as Planck (Planck Collaboration Int. XIX 2015; Planck Collaboration Int. XX 2015; Planck Collaboration XII 2020), BLASTPOL in Fissel et al. (2016), and SCUPOL (Poidevin et al. 2013), our analysis of interferometric data requires us to characterize how the filtering alters the polarization quantities we use in our statistics. With this aim, we use a set of synthetically observed non-ideal MHD simulations computed with RAMSES (Teyssier 2002; Fromang et al. 2006) that follow the gravitational collapse of cores whose range of initial mass and turbulence reproduce the main characteristics of the sources from our sample. The set consists of six simulations of collapsing cores (with total masses of 30, 60, and 100 M\u2299). We perform radiative transfer on these models using the POLARIS code (Reissl et al. 2016), which produces the Stokes I, Q, and U maps and assumeseither that a constant fraction of the dust grains are perfectly aligned everywhere (perfect alignment, known as \u201cPA\u201d hereafter) or that paramagnetic grains are aligned via radiative torques, known as \u201cRATs\u201d hereafter (e.g., Lazarian 2007). We note that the hypothesis of perfect alignment is not physical, and we do not aim to reproduce or interpret the polarized dust emission from Class 0 envelopes as resulting from perfect alignment. However, while we recognize that an hypothesis of perfect dust alignment is not a physical model but a phenomenological one, it has been suggested that the properties of dust polarization at the larger scales of the diffuse ISM (especially the results of \n\n${{\\mathcal{S}}\\times{\\mathcal{P}_{\\rm{frac}}}}$S\u00d7Pfrac\n) can be explained and reproduced with perfect alignment (Planck Collaboration XII 2020; Seifried et al. 2020). In the first part of our discussion we aim to compare our results with those obtained at larger scales, and thus perfect alignment remains an interesting point of comparison with RATs, and is a useful benchmark to compare how different physical models of grain alignment affect the statistical properties of the polarized emission. In addition, a case where the grains are perfectly aligned is only taking into account the source-specific geometrical effects governing the resulting polarization maps, and thus is useful to understand where alignment drops or is suppressed. We present all the details of the simulations and the radiative transfer calculations in Appendix D. In order to produce realistic synthetic observations to compare with the ALMA datasets, we use the CASA simulator (with the typical ALMA uv -coverage of these observations) to implement the effects of interferometric filtering and atmospheric noise on the POLARIS synthetic emission maps.","Citation Text":["Reissl et al. 2016"],"Citation Start End":[[1105,1123]]}
{"Identifier":"2021ApJ...913..144D__Goncharov_et_al._2014_Instance_1","Paragraph":"Whistler precursors are a manifestation of dispersive radiation (Tidman & Northrop 1968; Fairfield 1974; Hoppe 1981; Mellott & Greenstadt 1984), one mechanism through which collisionless shocks may transform bulk flow kinetic energy into other forms of energy. Several other mechanisms have been proposed, including wave\u2013particle interactions (Sagdeev 1966; Coroniti 1970a; Gary 1981), particle reflection (Edmiston & Kennel 1984; Kennel et al. 1985; Kennel 1987; Bale et al. 2005; Su et al. 2012), and macroscopic quasi-static field effects (Scudder et al. 1986a, 1986b, 1986c). At low Mf, theory suggests that dispersive radiation and wave\u2013particle interactions dominate (Kennel et al. 1985). Because this study focuses on low Mach number shocks (Mf \u2272 3), particle reflection and macroscopic field effects are not investigated in detail. Whistler precursors (Tidman & Northrop 1968; Fairfield 1974) are often observed and can dissipate energy through several processes, including generation of higher frequency waves by creating electron temperature anisotropies or current-driven instabilities (Gary 1981; Hull et al. 2012), acceleration of halo electrons and thermal ions (Wilson III et al. 2012; Chen et al. 2018), and deflection and modulation of core particles (Goncharov et al. 2014). Wave\u2013particle interactions dissipate energy through anomalous resistivity, shorthand for the energy and momentum exchange between the electromagnetic fields and particles (Sagdeev 1966; Coroniti 1970a; Gary 1981; Papadopoulos 1985; Wilson III et al. 2007, 2010, 2012, 2014a, 2014b; Breneman et al. 2013). Goodrich et al. (2018) showed that ion acoustic waves may be indicators or facilitators of momentum transfer between reflected and incident ion populations, thus linking these waves to the transformation of bulk flow kinetic energy. In addition, work by Wang et al. (2020) shows how reflected ions interacting with incident ions generate Debye-scale electrostatic fluctuations, starting as ion acoustic waves that trap ions and decay into electrostatic solitary waves. Krasnoselskikh et al. (2013) provides a review of the quasi-perpendicular bow shock and the inferred dissipation mechanisms based on Cluster observations. Some of the types of waves observed near shocks include magnetosonic whistler mode waves (Fairfield 1974; Coroniti et al. 1982), solitary waves (Bale et al. 1998), ion acoustic waves (Fredricks et al. 1968; Gary et al. 1975; Rodriguez & Gurnett 1975; Gurnett et al. 1979a, 1979b; Formisano & Torbert 1982; Fuselier & Gurnett 1984; Balikhin et al. 2005; Hull et al. 2006), Langmuir waves (Filbert & Kellogg 1979; Kellogg 2003, and references therein; Pulupa & Bale 2008), and electron cyclotron harmonic waves and waves associated with the electron cyclotron drift instability (ECDI) (Wilson III et al. 2010, 2014a, 2014b; Breneman et al. 2013; Goodrich et al. 2018).","Citation Text":["Goncharov et al. 2014"],"Citation Start End":[[1269,1290]]}
{"Identifier":"2016AandA...591A..13V__Vall\u00e9e_et_al._1986_Instance_1","Paragraph":"The first direct proof of the existence of magnetic fields in large-scale extragalactic environments, i.e., galaxy clusters, dates back to the 1970s with the discovery of extended, diffuse, central synchrotron sources called radio halos (see, e.g., Feretti et al. 2012 for a review). Later, indirect evidence of the existence of intracluster magnetic fields has been given by several statistical studies on the effect of the Faraday rotation on the radio signal from background galaxies or galaxies embedded in galaxy clusters (Lawler & Dennison 1982; Vall\u00e9e et al. 1986; Clarke et al. 2001; Johnston-Hollitt 2003; Clarke 2004; Johnston-Hollitt & Ekers 2004). On scales up to a few Mpc from the nearest galaxy cluster, possibly along filaments, only a few diffuse synchrotron sources have been reported (Harris et al. 1993; Bagchi et al. 2002; Kronberg et al. 2007; Giovannini et al. 2013, 2015). Magnetic fields with strengths on the order of 10-15\u2009G in voids might be indicated by \u03b3-ray observations (see Neronov & Vovk 2010; Tavecchio et al. 2010; Takahashi et al. 2012, 2013; but see Broderick et al. 2014a,b for alternative possibilities). Nevertheless, up to now, a robust confirmed detection of magnetic fields on scales that are much larger than clusters is not available. Stasyszyn et al. (2010) and Akahori et al. (2014a) investigated the possibility of statistically measuring Faraday rotation from intergalactic magnetic fields with present observations, showing that only the Square Kilometre Array (SKA) and its pathfinders are likely to succeed in this respect. By comparing the observations with single-scale magnetic field simulations, Pshirkov et al. (2015) infer an upper limit of 1.2\u2009nG for extragalactic large-scale magnetic fields, while the Planck Collaboration XIX (2016) derived a more stringent upper limit for primordial large-scale magnetic fields of B 0.67\u2009nG from the analysis of the Cosmic Microwave Background (CMB) power spectra and the effect on the ionization history (but see also Takahashi et al. 2005; Ichiki et al. 2006). ","Citation Text":["Vall\u00e9e et al. 1986"],"Citation Start End":[[552,570]]}
{"Identifier":"2020ApJ...904...89N___2011_Instance_1","Paragraph":"However, the determination of the faint side of the QLF at z \u223c 5 is important to understand the overall picture of SMBH growth in the early universe. Previous observations suggest that the number density of luminous quasars increased from the early universe to z \u223c 2 and then decreased to the current universe (e.g., Richards et al. 2006; Croom et al. 2009). It is of particular interest to study possible luminosity dependences of the quasar number-density evolution (luminosity-dependent density evolution; LDDE). Recent optical surveys of high-redshift quasars have reported that low-luminosity quasars show peak number density at lower redshifts than do high-luminosity quasars (e.g., Croom et al. 2009; Ikeda et al. 2011, 2012; Niida et al. 2016). Since the quasar luminosity at a given Eddington ratio corresponds to MBH, the reported LDDE trend is sometimes called the downsizing evolution. The same trend has been found by X-ray surveys for all AGN populations, including obscured (type-2) ones (e.g., Ueda et al. (2003, 2014), Hasinger et al. (2005), Aird et al. (2015), Miyaji et al. (2015); see also Enoki et al. (2014) and Shirakata et al. (2019) and references therein for theoretical works on the AGN downsizing evolution). Note that the downsizing evolution was originally proposed to describe the redshift evolution of the galaxy mass function (e.g., Cowie et al. 1996; Neistein et al. 2006; Fontanot et al. 2009). Therefore, the AGN downsizing evolution, if present, may provide a significant insight into the galaxy\u2013SMBH coevolution. However, the number density of low-luminosity quasars at high redshifts has been quite uncertain, which has prevented us from understanding the whole picture of the quasar LDDE. Ueda et al. (2014) noted a possibility that the number ratio of total (type-1 and type-2) low-luminosity AGNs to high-luminosity ones may increase from z \u223c 3 to z \u223c 5, which is referred to as \u201cup-down sizing\u201d (see also Glikman et al. 2010, 2011; Giallongo et al. 2015). We need large samples of high-z low-luminosity AGNs to examine such scenarios.","Citation Text":["Glikman et al.","2011"],"Citation Start End":[[1949,1963],[1970,1974]]}
{"Identifier":"2016MNRAS.461..666K__Joshi_et_al._2011_Instance_1","Paragraph":"C-statistic (e.g. Jang & Miller 1997) is the most commonly used and the one-way analysis of variance (ANOVA; de Diego 2010) the most powerful test for verifying the presence of variability in a DLC. However, we did not employ either of these tests because, de Diego (2010) has questioned the validity of the C-test by arguing that the C-statistics does not have a Gaussian distribution and the commonly used critical value of 2.567 is too conservative. On the other hand, the ANOVA test requires a rather large number of data points in the DLC, so as to have several points within each sub-group used for the analysis. This is not feasible for our DLCs which typically have no more than about 30\u201345 data points. Therefore, we have instead used the F-test which is based on the ratio of variances, F = variance(observed)\/variance(expected) (de Diego 2010; Villforth, Koekemoer & Grogin 2010), with its two versions : (i) the standard F-test (hereafter F\u03b7-test, Goyal et al. 2012) and (ii) scaled F-test (hereafter F\u03ba-test, Joshi et al. 2011). The F\u03ba-test is preferred when the magnitude difference between the object and comparison stars is large (Joshi et al. 2011). Onward Paper II, we have only been using the F\u03b7-test because our objects are generally quite comparable in brightness to their available comparison stars. An additional gain from the use of the F\u03b7-test is that we can directly compare our INOV results with those deduced for other major AGN classes (Goyal et al. 2013). An important point to keep in mind while applying the statistical tests is that the photometric errors on individual data points in a given DLC, as returned by the algorithms in the iraf and daophot softwares are normally underestimated by the factor \u03b7 which ranges between 1.3 and 1.75, as estimated in independent studies (e.g. Gopal-Krishna, Sagar & Wiita 1995; Garcia et al. 1999; Sagar et al. 2004; Stalin et al. 2004a; Bachev, Strigachev & Semkov 2005). Recently, using a large sample, Goyal et al. (2013) estimated the best-fitting value of \u03b7 to be 1.5, which is adopted here. Thus, the F\u03b7 statistics can be expressed as\n\n\n\\begin{equation*}\nF_{1}^{\\eta } = \\frac{\\sigma ^{2}_{({\\rm q-s1})}}{ \\eta ^2 \\langle \\sigma _{{\\rm q-s1}}^2 \\rangle }, \\hspace{5.69046pt} F_{2}^{\\eta } = \\frac{\\sigma ^{2}_{({\\rm q-s2})}}{ \\eta ^2 \\langle \\sigma _{{\\rm q-s2}}^2 \\rangle }, \\hspace{5.69046pt} F_{{\\rm s1-s2}}^{\\eta } = \\frac{\\sigma ^{2}_{({\\rm s1-s2})}}{ \\eta ^2 \\langle \\sigma _{{\\rm s1-s2}}^2 \\rangle },\\end{equation*}\n\n where $\\sigma ^{2}_{({\\rm q-s1})}$, $\\sigma ^{2}_{({\\rm q-s2})}$ and $\\sigma ^{2}_{({\\rm s1-s2})}$ are the variances of the \u2018quasar\u2013star1\u2019, \u2018quasar\u2013star2\u2019 and \u2018star1\u2013star2\u2019 DLCs and $\\langle \\sigma _{{\\rm q-s1}}^2 \\rangle =\\sum _{\\boldsymbol {i}=0}^{N}\\sigma ^2_{i,{\\rm err}}({\\rm q-s1})\/N$, $\\langle \\sigma _{{\\rm q-s2}}^2 \\rangle$ and $\\langle \\sigma _{{\\rm s1-s2}}^2 \\rangle$ are the mean square (formal) rms errors of the individual data points in the \u2018quasar\u2013star1\u2019, \u2018quasar\u2013star2\u2019 and \u2018star1\u2013star2\u2019 DLCs, respectively. \u03b7 is the scaling factor (= 1.5) as mentioned above.","Citation Text":["Joshi et al. 2011"],"Citation Start End":[[1022,1039]]}
{"Identifier":"2022AandA...658A.167G__Bromley_et_al._2016_Instance_1","Paragraph":"Asymptotic giant branch (AGB) stars are a major contributor to the global dust budget in galaxies (H\u00f6fner & Olofsson 2018). Owing to their refractory nature, alumina (stoichiometric formula Al2O3) is a promising candidate to represent the first dust condensate in oxygen-rich AGB stars. Related alumina clusters are thought to initiate dustformation in these environments and are often referred to as \u2018seed particles\u2019 (Gail & Sedlmayr 2013). However, the sizes and compositions of these aluminium oxide clusters are not well characterised. In this study, we investigate a range of Al:O stoichiometries in order to review these predictions and to construct realistic models of these initial dust seeds. The emergence of a specific condensate is predicted by its condensation temperature (Tielens 2005) and depends on the thermal stability of the solid, as well as on the gas density and its composition. Usually, the evaluation of the stability of the likely condensates is based on macroscopic bulk properties such as the vapour pressure, which is a measure of the volatility of a substance. Hence, the most refractory condensate is expected to have the lowest vapour pressure. Corundum (\u03b1-alumina), corresponding to the most stable crystalline bulk form of alumina, fulfils this condition (Gail et al. 2013). The growth and size distribution of dust grains is commonly described by classical nucleation theory (CNT). However, the applicabilityof CNT in an expanding circumstellar envelope has been questioned (Donn & Nuth 1985; Goumans & Bromley 2012; Bromley et al. 2016; Gobrecht et al. 2017). In particular, the concept of vapour pressures and the universal assumption of thermodynamic equilibrium (TE) are in contradiction with the synthesis and growth of dust grains in highly dynamical AGB atmospheres. Moreover, in CNT, the properties of small solids are derived from the (crystalline)bulk material. However, the properties of nano-sized clusters often differ significantly from those of bulk analogues. The constraints associated with extremely small sizes lead to clusters with non-crystalline structures, whose characteristics (e.g. energy, geometry, bond lengths and angles, atomic coordination) differ substantially from those of the bulk material (Bromley & Zwijnenburg 2016). In particular, the energetic stability of such nano-clusters is typically higher than that of clusters with structures directly obtained from \u2018top-down\u2019 cuts from the parent bulk crystalline material, which represent metastable, or even unstable, configurations (Lamiel-Garcia et al. 2017). In addition, the concept of surface free energy (or tension), which is fundamental in CNT, is not applicable to small clusters, where it is difficult to differentiate between surface and bulk. Surface energies can only be applied to clusters with fairly large sizes (e.g. facetted bulk cut clusters). We understand nucleation as the formation and growth of stable seed nuclei (i.e. clusters) from prevalent gas-phase molecules, whose abundance varies in time (i.e. is often not in equilibrium). Therefore, a cluster is intermediate in size between a molecule and a bulk solid.","Citation Text":["Bromley et al. 2016"],"Citation Start End":[[1553,1572]]}
{"Identifier":"2019MNRAS.486.5558S__Shultz_et_al._2018_Instance_1","Paragraph":"One possible, conventional, explanation for an apparently accelerating rotational period may be the light-time effect due to the orbit of a binary companion. However, none of the magnetic B-type stars in which this phenomenon has been detected, including the present star, are known to be in binary systems (Shatsky & Tokovinin (2002) conducted an NIR search for visual companions, and found no evidence of a companion in the case of HD\u2009142990). The change in period \u0394P due to the light-time effect should correspond to a change in radial velocity \u0394RV = c\u0394P\/P, where c is speed of light (e.g. Pigulski & Boratyn 1992). Fig. 6 shows the least-squares deconvolution (LSD) profiles extracted from the ESPaDOnS data set with a line mask using all metallic lines in the spectrum (for details, see Shultz et al. 2018). No bulk RV variability is apparent. Measuring the RV is complicated by the spectroscopic variability introduced by chemical spots, which in addition to equivalent width changes also introduce RV variations coherent with the rotation phase due to changes in the star\u2019s centre of gravity. As a result, only measurements performed on observations obtained close to the same rotation phase can be compared. The ESPaDOnS data contains two observations obtained close to phase 0.6, separated by about 3 yr (one on 14\/06\/2014, the second on 14\/05\/2017, with a difference in phase of 0.02 cycles when phased using the variable ephemeris). Measuring the centres of gravity of the LSD profiles extracted from these observations yields a difference in RV of 1 km\u2009s\u22121, comparable to the measurement uncertainty. The RV change expected over 3 yr if $\\dot{P}$ is due to orbital motion is about 3 km\u2009s\u22121, so this test must be considered inconclusive. However, a change of \u221220 s over the 30 yr of observations should have led to \u0394RV = 71 km\u2009s\u22121; it is unlikely that such a large change in RV would have been missed. The Pulkovo Compilation of Radial Velocities (Gontcharov 2006) give RV = \u221212 \u00b1 3 km\u2009s\u22121, consistent with RVs measured from ESPaDOnS data (which have a mean and standard deviation of \u22124 and 5 km\u2009s\u22121), suggesting that the RV has been stable over a time span of at least a decade.","Citation Text":["Shultz et al. 2018"],"Citation Start End":[[792,810]]}
{"Identifier":"2015MNRAS.452.2731S__Stroe_et_al._2013_Instance_3","Paragraph":"The H\u2009\u03b1 studies of Umeda et al. (2004) and Stroe et al. (2014a, 2015) are tracing instantaneous (averaged over 10 Myr) SF and little is known about SF on longer time-scales and the reservoir of gas that would enable future SF. An excellent test case for studying the gas content of galaxies within merging clusters with shocks is CIZA J2242.8+5301 (Kocevski et al. 2007). For this particular cluster unfortunately, its location in the Galactic plane, prohibits studies of the rest-frame UV or FIR tracing SF on longer time-scales, as the emission is dominated by Milky Way dust. However, the rich multiwavelength data available for the cluster give us an unprecedented detailed view on the interaction of their shock systems with the member galaxies. CIZA J2242.8+5301 is an extremely massive (M200 \u223c 2 \u00d7 1015\u2009M\u2299; Dawson et al. 2015; Jee et al. 2015) and X-ray disturbed cluster (Akamatsu & Kawahara 2013; Ogrean et al. 2013, 2014) which most likely resulted from a head-on collision of two, equal-mass systems (van Weeren et al. 2011; Dawson et al. 2015). The cluster merger induced relatively strong shocks, which travelled through the ICM, accelerated particles to produce relics towards the north and south of the cluster (van Weeren et al. 2010; Stroe et al. 2013). There is evidence for a few additional smaller shock fronts throughout the cluster volume (Stroe et al. 2013; Ogrean et al. 2014). Of particular interest is the northern relic, which earned the cluster the nickname \u2018Sausage\u2019. The relic, tracing a shock of Mach number M \u223c 3 (Stroe et al. 2014c), is detected over a spatial extent of \u223c1.5 Mpc in length and up to \u223c150 kpc in width and over a wide radio frequency range (150 MHz\u201316 GHz; Stroe et al. 2013, 2014b). There is evidence that the merger and the shocks shape the evolution of cluster galaxies. The radio jets are bent into a head\u2013tail morphology aligned with the merger axis of the cluster. This is probably ram pressure caused by the relative motion of galaxies with respect to the ICM (Stroe et al. 2013). The cluster was also found to host a high fraction of H\u2009\u03b1 emitting galaxies (Stroe et al. 2014a, 2015). The cluster galaxies not only exhibit increased SF and AGN activity compared to their field counterparts, but are also more massive, more metal rich and show evidence for outflows likely driven by SNe (Sobral et al. 2015). Stroe et al. (2015) and Sobral et al. (2015) suggest that these relative massive galaxies (stellar masses of up to \u223c1010.0\u201310.7 M\u2299) retained the metal-rich gas, which was triggered to collapse into dense star-forming clouds by the passage of the shocks, travelling at speeds up to \u223c2500 km s\u22121 (Stroe et al. 2014c), in line with simulations by Roediger et al. (2014).","Citation Text":["Stroe et al. 2013"],"Citation Start End":[[1706,1723]]}
{"Identifier":"2020MNRAS.494.3627A__Ellison_et_al._2019_Instance_1","Paragraph":"Star formation and supermassive resblack hole (SMBH) growth are two important processes in galaxies that influence their evolution throughout cosmic history. However, we do not yet understand why the global rates of star formation (e.g. Hopkins & Beacom 2006; Madau & Dickinson 2014; Driver et al. 2018) and SMBH growth (e.g. Ueda et al. 2003; Shankar, Weinberg & Miralda-Escud\u00e9 2009) both peaked at z \u2248 2, and then declined by an order magnitude to this epoch. It is clear that a ready supply of cold (Tk \u226a 104\u2009K) gas is important; in the nearby Universe the surface densities of star formation and neutral gas, particularly the molecular component, are strongly correlated (e.g. Schmidt 1959; Kennicutt 1998; Bigiel et al. 2008), and likewise radiatively efficient active galactic nuclei (AGNs) are predominantly hosted by star-forming galaxies with a central young stellar population and therefore ample cold gas reservoirs (e.g. Kauffmann et al. 2003, 2007; Kauffmann & Heckman 2009; LaMassa et al. 2013; Ellison et al. 2019). Determining how the neutral interstellar medium in galaxies has evolved over the history of the Universe is therefore a key component in understanding their evolution. Much of our knowledge of the global content of neutral gas in galaxies comes from observing hydrogen gas, the most common element in the Universe. In its neutral atomic (H\u2009i) phase, hydrogen is readily detectable in nearby galaxies via 21-cm emission at radio wavelengths (see Giovanelli & Haynes 2016 for a review) or, at cosmological distances, through Lyman\u2009\u03b1 (n = 1\u20132) absorption in the ultraviolet and visible bands (see Wolfe, Gawiser & Prochaska 2005). The total H\u2009i mass density shows comparatively less evolution over cosmological time-scales than that of star formation, decreasing by at most two-fold since z \u2248 2 (e.g. Zwaan et al. 2005; Martin et al. 2010; Braun 2012; Noterdaeme et al. 2012; Zafar et al. 2013; Crighton et al. 2015; S\u00e1nchez-Ram\u00edrez et al. 2016; Bird, Garnett & Ho 2017; Rhee et al. 2018), suggesting that much of the neutral gas content of galaxies is replenished over these time-scales.","Citation Text":["Ellison et al. 2019"],"Citation Start End":[[1009,1028]]}
{"Identifier":"2017ApJ...850...20G__Dexheimer_&_Schramm_2008_Instance_1","Paragraph":"The observation of massive neutron stars Demorest et al. (2010), Antoniadis et al. (2013) indicates that the EoS of nuclear matter must be very stiff in the regime of high densities and low temperatures. The degree of stiffness in the nuclear matter EoS is directly related to the repulsive interaction among particles at high densities, as well as to the particle content in the core of the stars. In particular, it has been extensively discussed in the literature whether exotic degrees of freedom might populate the core of neutron stars. On the one hand, it is more energetically favorable for the system to populate new degrees of freedom, such as hyperons (Dexheimer & Schramm 2008; Ishizuka et al. 2008; Bednarek et al. 2012; Fukukawa et al. 2015; Gomes et al. 2015; Maslov et al. 2015; Oertel et al. 2015; Lonardoni et al. 2015, 2016); Biswal et al. 2016; Burgio & Zappal\u00e0 2016; Chatterjee & Vidana 2016; Mishra et al. 2016; Vida\u00f1a 2016; Yamamoto et al. 2016; Tolos et al. 2017); Torres et al. 2017), delta isobars (Fong et al. 2010;Schurhoff et al. 2010; Drago et al. 2014; 2016; Cai et al. 2015; Zhu et al. 2016), and meson condensates (Ellis et al. 1995; Menezes et al. 2005; Takahashi 2007; Ohnishi et al. 2009; Alford et al. 2010; Fernandez et al. 2010; Mesquita et al. 2010; Mishra et al. 2010; Lim et al. 2014; Muto et al. 2015), in order to lower its Fermi energy (starting at about two times the saturation density). On the other hand, the EoS softening due to the appearance of exotica might turn some nuclear models incompatible with observational data, in particular with the recently measured massive neutron stars. One possible way to overcome this puzzle is the introduction of an extra repulsion in the YY interaction Schaffner & Mishustin (1996), Bombaci (2016), allowing models with hyperons to be able to reproduce massive stars (Dexheimer & Schramm 2008; Bednarek et al. 2012; Weissenborn et al. 2012; Banik et al. 2014; Bhowmick et al. 2014; Gusakov et al. 2014; Lopes & Menezes 2014; van Dalen et al. 2014; Yamamoto et al. 2014; Gomes et al. 2015). Another possible solution is the introduction of a deconfinement phase transition at high densities Bombaci (2016), with a stiff EoS for quark matter, usually associated with quark vector interactions (see Kl\u00e4hn et al. 2013 and references therein).","Citation Text":["Dexheimer & Schramm 2008","Dexheimer & Schramm 2008"],"Citation Start End":[[663,687],[1857,1881]]}
{"Identifier":"2016AandA...593A.108B__Kormendy_1977_Instance_1","Paragraph":"Our models for the density distribution in the thin and thick stellar disks of the Galaxy are based on the structural disk parameters presented by PJL. These authors performed a star counts model of the Galaxy using near-infrared data of the 2MASS survey (Skrutskie et al. 2006) with lines of sight covering the entire sky and including the Galactic plane. The authors explored the parameter space and estimated its optimal values with statistical methods such as the Markov chain Monte Carlo (MCMC; Gilks et al. (1996)) and the nested sampling (NS) algorithm (Skilling 2004). Using a modified exponential law, PJL modelled the radial profile of the density of each subcomponent of the stellar disk based on the Galactic model of L\u00e9pine & Leroy (2000). Such profile is equivalent to the Freeman Type II disk brightness profile, which contains a depletion in the centre with respect to a pure exponential law (Freeman 1970; Kormendy 1977). The stellar surface densities \u03a3d\u2605 for the thin and thick disks can then be written as (1)\\begin{equation} \\label{eq:sigma_thin_thick} \\Sigma_{\\mathrm{d}_{\\bigstar,\\,i}}(R)=\\Sigma_{0\\mathrm{d}_{\\bigstar,\\,i}}\\,\\exp\\left[-\\,\\frac{(R-R_{0})}{R_{\\mathrm{d}_{\\,i}}}-R_{\\mathrm{ch}_{\\,i}}\\left(\\frac{1}{R}-\\frac{1}{R_{0}}\\right)\\right], \\end{equation}\u03a3d\u2605,\u2009i(R)=\u03a30d\u2605,\u2009i\u2009exp\u2212\u2009(R\u2212R0)Rd\u2009i\u2212Rch\u2009i1R\u22121R0,where \u03a30d\u2605 corresponds to the local disk stellar surface density (at R = R0), Rd is the radial scale length, and Rch is the radial length of the \u201ccentral hole\u201d in the density of each stellar disk i subcomponent (i = thin, thick). The hypothesis that the Galactic disk is hollow in its centre has been justified by some models that use observational data at infrared bands to describe the inner structure of the Galaxy (e.g. Freudenreich 1998; L\u00e9pine & Leroy 2000; L\u00f3pez-Corredoira et al. 2004; Picaud & Robin 2004). In the particular case of the PJL model, only the thin disk needs a density depression in its inner part; in contrast, the thick disk can be described by a simple radial exponential decay, i.e.  Rch thick = 0. ","Citation Text":["Kormendy 1977"],"Citation Start End":[[923,936]]}
{"Identifier":"2016MNRAS.461.4176H__Jaffe_&_Kaiser_1995_Instance_2","Paragraph":"The Bayesian approach is indeed robust and optimal, within the context mentioned above. It focuses on the reconstruction of the LSS within the framework of the standard model of cosmology and for a given data base. However, most previous studies have not addressed the question how consistent is the assumed cosmological model with the observed data. This is not a trivial issue \u2013 the model needs to agree with the data so as to provide a solid foundation for the WF\/CRs construction. Ideally, one should have started with establishing the agreement of the model with the data and only then reconstruct the LSS in the manner described above. However, history does not always proceeds in a linear fashion. The aim of the paper is to amend that situation and establish the likelihood of peculiar velocities data bases given the standard model of cosmology. The relevant methodology is straightforward and well established. One needs to calculate the likelihood function of the data given the model \u2013 namely the probability of the occurrence of the data within the framework of the assumed model (Jaffe & Kaiser 1995; Zaroubi et al. 1995; Hoffman 2001; Press et al. 2007). The likelihood function establishes the goodness-of-fit (GoF) of the data by the model. In the cosmological case and under the assumption of the linear regime, where the velocity field constitutes a random Gaussian vector field and the observational errors are normally distributed, the likelihood analysis amounts to calculating a \u03c72 statistics. This approach was indeed applied to velocity data bases (Jaffe & Kaiser 1995; Zaroubi et al. 1997, 2001). The application of the likelihood analysis to actual velocity data bases suffers however from one major drawback. The gravitational dynamics of structure formation induces non-linear contributions to the velocities of galaxies. These non-linear corrections render the parameter estimation and GoF analysis to be rather uncertain. The remedy to the problem involves the filtering of small scales to give linearized data. The likelihood analysis can then be safely applied to the linearized data. Here, we suggest such a small-scales filtering procedure and study the extent by which Cosmicflows-2 (CF2; Tully et al. 2013) is compatible with the \u039b cold dark matter (\u039bCDM) standard model of cosmology.","Citation Text":["Jaffe & Kaiser 1995"],"Citation Start End":[[1574,1593]]}
{"Identifier":"2022MNRAS.509.6091H___2020a_Instance_2","Paragraph":"Galactic winds have been ubiquitously observed in galaxies at both low and high redshifts, and they are critical to galaxy formation and evolution. Simulations calibrated to match these observations predict that a large amount of galactic material is ejected as a wind before reaccreting to either form stars or be ejected once again (Oppenheimer et al. 2010; Angl\u00e9s-Alc\u00e1zar et al. 2017). Current cosmological hydrodynamic simulations of galaxy formation employ a variety of subgrid models (e.g. Springel & Hernquist 2003; Oppenheimer & Dav\u00e9 2006; Stinson et al. 2006; Dalla Vecchia & Schaye 2008; Agertz et al. 2013; Schaye et al. 2015; Dav\u00e9, Thompson & Hopkins 2016; Tremmel et al. 2017; Pillepich et al. 2018; Dav\u00e9 et al. 2019; Huang et al. 2020a) that artificially launch galactic winds, but the results are sensitive to numerical resolution and the exact subgrid model employed (Huang et al. 2019, 2020a). Simulations without these subgrid wind models (e.g. Hopkins et al. 2018; Kim & Ostriker 2015; Martizzi et al. 2016) allow winds to occur \u2018naturally\u2019, but these simulations may not resolve the scales necessary to resolve the important known physical processes (Scannapieco & Br\u00fcggen 2015; Br\u00fcggen & Scannapieco 2016; Schneider & Robertson 2017; McCourt et al. 2018; Huang et al. 2020b). Hence, modelling galactic winds accurately remains a theoretical challenge for even the most refined high-resolution simulations of galaxies (see Naab & Ostriker 2017, for a review). Even if one were able to accurately model the formation of galactic winds, the subsequent propagation in galactic haloes depends on a complicated interplay of many physical processes that occur on a wide range of physical scales that cannot be simultaneously resolved in a single simulation. For example, to robustly model the propagation and disintegration of moving clouds in various situations requires cloud-crushing simulations with at least sub-parsec scale resolution (Schneider & Robertson 2017; McCourt et al. 2018), which is orders of magnitudes below the resolution limits of cosmological simulations. Furthermore, most cosmological hydrodynamic simulations concentrate their resolution in the dense, star-forming regions of galaxies and thus have lower resolution in the circumgalactic medium (CGM, but see Hummels et al. 2019; Mandelker et al. 2019; Peeples et al. 2019; Suresh et al. 2019; van de Voort et al. 2019). To date, cosmological simulations do not include physically motivated subgrid models for galactic wind evolution, which are required to capture these small-scale physical processes.","Citation Text":["Huang et al.","2020a"],"Citation Start End":[[884,896],[903,908]]}
{"Identifier":"2020AandA...633A..34C__Carlberg_et_al._2010_Instance_1","Paragraph":"Some studies use HIPPARCOS or Gaia data to determine the evolutionary status of field LiRG and show that these objects tend to accumulate close to the RGB bump, the clump, and the early-AGB (e.g. Charbonnel & Balachandran 2000; Kumar et al. 2011; Smiljanic et al. 2018; Deepak 2019), which is in agreement with open cluster studies (e.g. Delgado Mena et al. 2016). Other works report, however, that LiRG can be randomly located in the HRD (Jasniewicz et al. 1999; Smith et al. 1999; Monaco et al. 2011; Lebzelter et al. 2012; Martell & Shetrone 2013; Casey et al. 2016). The distinction is crucial to understanding the processes that may provide an explanation for the phenomenon, such as fresh Li production by internal mixing processes (Sackmann & Boothroyd 1999; Palacios et al. 2001; Guandalini et al. 2009; Strassmeier et al. 2015; Cassisi et al. 2016), prompt mass loss events (de La Reza et al. 1996, 1997), Li accretion during engulfment of planets or planetesimals (Alexander 1967; Siess & Livio 1999; Carlberg et al. 2010; Aguilera-G\u00f3mez et al. 2016a,b; Delgado Mena et al. 2016), tidal interactions between binary stars (Casey et al. 2019), or a combination of these mechanisms (Denissenkov & Weiss 2000; Denissenkov & Herwig 2004). However, since the evolution tracks of evolved stars all converge to the same area of the CMD, the definitive determination of the actual evolution status of LiRG requires asteroseismology to probe their internal structure and disentangle RGB from clump stars. As of today, very few LiRG have been observed with CoRoT and Kepler. The majority seems to be located in the core-He burning clump (Silva Aguirre et al. 2014; Bharat Kumar et al. 2018; Casey et al. 2016; Smiljanic et al. 2018; Singh et al. 2019), with the others being at the RGB bump or higher on the first ascent giant branch (Jofr\u00e9 et al. 2015; Casey et al. 2019). In a recent study using LAMOST spectra to derive both the Li abundance and asteroseismic classification, Casey et al. (2019) showed that \u223c80% of their large sample of low-mass LiRG (2330 objects) probably have helium burning cores. They find that LiRG are more frequent at higher metallicity.","Citation Text":["Carlberg et al. 2010"],"Citation Start End":[[1011,1031]]}
{"Identifier":"2019MNRAS.487.3776P__Yu_&_Liu_2018_Instance_1","Paragraph":"Age\u2013velocity dispersion relations have been found in many resolved (Nordstr\u00f6m et al. 2004; Grieves et al. 2018) and simulated (Martig, Minchev & Flynn 2014) studies. They primarily conclude that stars born earlier have a higher velocity dispersion compared to those formed later. The underlying physical cause of this relation, however, has evaded a general consensus in the literature. Some scenarios broadly claim that the increase in velocity dispersion is a dynamical effect of internal interactions that build up over time, with a number of different specific mechanisms being proposed as the culprit. Since the older stars have been experiencing these interactions for a longer period of time, they should therefore show the largest increase in velocity dispersion. Evidence in favour of such internal mechanisms has come from both observations (Yu & Liu 2018) and simulations (Saha, Tseng & Taam 2010; Aumer, Binney & Sch\u00f6nrich 2016; Grand et al. 2016). An alternative explanation is that in the early Universe, conditions were generally more chaotic (Wisnioski et al. 2015), so that any stars born at that time were more likely to have higher velocity dispersion. As conditions gradually settled over cosmic time, stars were being born in progressively lower-dispersion conditions. A number of studies have identified the conditions at birth as the dominant effect in determining a population\u2019s present-day velocity dispersion, again both from resolved observations (Leaman et al. 2017) and simulations (Bird et al. 2013; Ma et al. 2017). Finally, the comparison between observations and simulations in Pinna et al. (2018) has identified the underlying complexity and inherent degeneracy in discriminating between these scenarios. They claim that many of the effects described above likely play a role to a varying degree and that the imprint of some mechanisms fade over time, further complicating any attempt to constrain the physical cause of disc heating.","Citation Text":["Yu & Liu 2018"],"Citation Start End":[[852,865]]}
{"Identifier":"2015ApJ...813..103M__Koss_et_al._2012_Instance_2","Paragraph":"The stochastic accretion of gas and galaxy merger-driven gas inflows are both known triggers of supermassive black hole (SMBH) growth and nuclear activity, but the relative contributions of each is still unclear. Simulations of galaxy mergers show that they drive gas to the centers of merger-remnant galaxies (e.g., Springel et al. 2005; Hopkins & Hernquist 2009), predicting that merger-driven SMBH mass growth occurs when the black hole nears the center of the merger remnant. Observations have shown that the AGN fraction does increase with decreasing distance between two merging galaxies (Ellison et al. 2011; Koss et al. 2012; Ellison et al. 2013), but this has not been well tested at the very centers of merger-remnant galaxies because of the observational difficulty of detecting and resolving two AGNs with separations 10 kpc. This is known as the \u201cdual AGN\u201d phase.4\n\n4\nThe separation scale expected for dual AGNs is between 0.1 and 10 kpc. The SMBHs in a merger stay at these separations for a few hundred megayears before evolving into a gravitationally bound, parsec-scale separation binary AGN system (Begelman et al. 1980).\n Hundreds of AGN pairs with >10 kpc separations have been discovered (Myers et al. 2008; Hennawi et al. 2010; Liu et al. 2011). However, there are only a few confirmed kiloparsec-scale dual AGNs (Junkkarinen et al. 2001; Komossa et al. 2003; Hudson et al. 2006; Rodriguez et al. 2006; Bianchi et al. 2008; Fu et al. 2011b; Koss et al. 2011, 2012; Mazzarella et al. 2012; Liu et al. 2013; Comerford et al. 2015). Dual AGNs are an intermediate evolutionary stage between first encounter and final coalescence of two merging gas-rich galaxies (e.g., Comerford et al. 2009; Liu et al. 2012), in which strong tidal interactions are more likely to influence the nuclear accretion and star formation in both galaxies (Barnes & Hernquist 1996). Indeed, galaxy merger simulations and observations clearly show that the dual AGN phase is the critical stage when SMBH growth and star formation activity are the most vigorous (e.g., Koss et al. 2012; Van Wassenhove et al. 2012; Blecha et al. 2013).","Citation Text":["Koss et al.","2012"],"Citation Start End":[[1463,1474],[1481,1485]]}
{"Identifier":"2017AandA...606A.113G__Carollo_et_al._(2013)_Instance_1","Paragraph":"In this context, massive (\u2133>1011\u2009M\u2299) PGs (MPGs) deserve particular attention. These systems are expected to evolve mainly through (dry) mergers (e.g. Hopkins et al. 2009; De Lucia & Blaizot 2007). If this is the case, in this mass range we should detect a stronger signal of the size-growth with respect to a lower mass range. So far, because MPGs are extremely rare, there have been very few studies that have investigated the combined evolution of the number density and of the age of MPGs as a function of their compactness (Carollo et al. 2013; Fagioli et al. 2016). Carollo et al. (2013) found that the number density of massive quiescent and elliptical galaxies with Re 2.5\u2009kpc decreases by about 30% from z~1 to z~0.2 and that their U\u2212V colours are consistent with passive evolution. They concluded that the driving mechanism for the average size-growth of the whole population is the appearance at later epochs of larger quiescent galaxies. More recently, Fagioli et al. (2016, hereafter F16) analysed the spectroscopic properties of ~500 MPGs (defined as galaxies with absent or very weak emission lines and no MIPS detections) at 0.2 z 0.8 in the zCOSMOS-bright 20 K catalogue (Lilly et al. 2007). From the analysis of stacked spectra of small and large MPGs, they dated the stellar content of these groups and found that the two sub-populations have similar ages. The authors concluded that, in this mass regime, the size growth of individual galaxies through dry mergers is the most likely explanation for the increase in the mean effective radius of the whole population. A recent analysis by Zahid et al. (2016) on the physical properties of compact post starburst galaxies at 0.2z0.8 with \u2133>1011\u2009M\u2299 provides new insights. On the basis of both their number density and of their ages, which have been found to be 1 Gyr, the authors suggest that this class of objects are the progenitors of compact quiescent galaxies. They conclude that a substantial fraction of dense quiescent galaxies at z0.8 are newly formed. ","Citation Text":["Carollo et al. 2013"],"Citation Start End":[[528,547]]}
{"Identifier":"2022MNRAS.512.4893Z__Bourne_et_al._2017_Instance_1","Paragraph":"Submillimetre galaxies (SMGs) are ultraluminous dusty star-forming galaxies (SFGs) with the vast majority of radiation energy in the far-infrared (FIR) and submillimetre bands (Smail, Ivison & Blain 1997; Barger et al. 1998; Hughes et al. 1998; Micha\u0142owski et al. 2012; Casey, Narayanan & Cooray 2014). They are extreme starbursts with high star formation rate (SFR) of $\\sim 10^2\\, -\\, 10^3$\u2009M\u2299\u2009yr\u22121 and high stellar mass of \u223c1010\u201311\u2009M\u2299, located preferentially at z \u223c 2\u20133 (Chapman et al. 2005; Wardlow et al. 2011; Smol\u010di\u0107 et al. 2012; Simpson et al. 2014; Brisbin et al. 2017; Danielson et al. 2017; Micha\u0142owski et al. 2017; Smith et al. 2017; Hodge & da Cunha 2020, and references therein). The physical properties of SMGs derived from detailed studies of individual sources and large sky area submillimetre surveys suggest that SMGs represent an early evolutionary phase of all local ellipticals (Smail et al. 2002; Swinbank et al. 2006; Fu et al. 2013; Toft et al. 2014; Ikarashi et al. 2015; Miettinen et al. 2017; An et al. 2019; Gullberg et al. 2019; Dudzevi\u010di\u016bt\u0117 et al. 2020; Rennehan et al. 2020). Although SMGs are a rare population (\u223c400\u2009deg\u22122 down to S850 = 4\u2009mJy; Simpson et al. 2019; Shim et al. 2020), they contribute a significant fraction (\u223c20\u2009per\u2009cent) of the cosmic SFR density at z > 2 (Bourne et al. 2017; Koprowski et al. 2017; Zavala et al. 2021). Their extreme SFRs are correlated with higher gas fractions compared to normal SFGs at the same epoch (Bothwell et al. 2013; Scoville et al. 2016; Decarli et al. 2016; Tacconi et al. 2018). The large amount gas is thought to be supplied by gas infall via cold streams from surrounding gas reservoirs (Narayanan et al. 2015; Ginolfi et al. 2017). The extreme starbursts of this population are partially triggered by galaxy major mergers (Tacconi et al. 2008; Engel et al. 2010, but see Dav\u00e9 et al. 2010; Narayanan et al. 2015; McAlpine et al. 2019), while a diversity of morphologies unveiled from the rest-frame ultraviolet (UV) and optical imaging indicate galaxy interactions and disc instabilities to be important mechanisms for enhancing star formation in SMGs (Swinbank et al. 2010; Kartaltepe et al. 2012; Chen et al. 2015), as well as AGN activities (Chapman et al. 2005; Wang et al. 2013). Because the rarity and high SFR of SMGs are sensitive to the physical processes governing galaxy formation (e.g. star formation, stellar and AGN feedback, gas infall, metal enrichment, and galaxy merging\/interactions), the SMG population is used to constrain cosmological models. It remains challenging to reproduce the SMG population with high SFRs matching observations at high z (Dav\u00e9 et al. 2010; McAlpine et al. 2019; Hayward et al. 2021; Lovell et al. 2021).","Citation Text":["Bourne et al. 2017"],"Citation Start End":[[1308,1326]]}
{"Identifier":"2015ApJ...805....4Z___2006b_Instance_1","Paragraph":"In addition to successful and failed eruptions, there are also partial filament eruptions (Gilbert et al. 2007; Liu et al. 2007). After examining 54 H\u03b1 prominence activities, Gilbert et al. (2000) found that a majority of the eruptive prominences show a separation of escaping material from the bulk of the prominence; the latter initially lifted away from and then fell back to the solar surface. To explain the partial filament eruptions, the authors proposed a cartoon model in which magnetic reconnection occurs inside an inverse-polarity flux rope, leading to the separation of the escaping portion of the prominence and the formation of a second X-type neutral line in the upper portion of the prominence. The inner splitting and subsequent partial prominence eruption was also observed by Shen et al. (2012). Gilbert et al. (2001) interpreted an active prominence with the process of vertical reconnection between an inverse-polarity flux rope and an underlying magnetic arcade. Liu et al. (2008) reported a partial filament eruption characterized by a quasi-static, slow phase and a rapid kinking phase showing a bifurcation of the filament. The separation of the filament, the extreme-ultraviolet (EUV) brightening at the separation location, and the surviving sigmoidal structure provide convincing evidence that magnetic reconnection occurs within the body of a filament (Tripathi et al. 2013). Gibson & Fan (2006a, 2006b) carried out three-dimensional (3D) numerical simulations to model the partial expulsion of an MFR. After multiple reconnections at current sheets that form during the eruption, the rope breaks in an upper, escaping rope and a lower, surviving rope. The \u201cpartially expelled flux rope\u201d model has been justified observationally (Tripathi et al. 2009). Tripathi et al. (2006) observed a distinct coronal downflow following a curved path at the speed of 150 km s\u22121 during a CME-associated prominence eruption. Their observation provides support for the pinching off of field lines drawn-out by the erupting prominences, and the contraction of the arcade formed by the reconnection. Tripathi et al. (2007) reported similar multithermal downflow at the speed of \u223c380 km s\u22121, starting at the cusp-shaped structures where magnetic reconnection occurred inside the erupting flux rope that led to its bifurcation. Liu et al. (2012) studied a flare-associated partial eruption of a double-decker filament. Cheng et al. (2014b) found that a stable double-decker MFR system existed for hours prior to the eruption on 12 July 2012. After entering the domain of instability, the high-lying MFR impulsively erupted to generate a fast CME and GOES X1.4 class flare, whereas the low-lying MFR remained behind and continuously maintained the sigmoidicity of the active region (AR). From previous literature, we can conclude that magnetic reconnection and the release of free energy is involved in most of the partial filament eruptions. However, the exact mechanism of partial eruptions, which is of great importance to understanding the origin of solar eruptions and forecasting space weather, remains unclear and controversial.","Citation Text":["Gibson & Fan","2006b"],"Citation Start End":[[1406,1418],[1427,1432]]}
{"Identifier":"2017MNRAS.471...80S__Kalugina_et_al._2014_Instance_1","Paragraph":"The electronic computations were performed using the molpro (Molpro 2015, http:\/\/www.molpro.net) package. In a preliminary work, we used the complete active space self-consistent field (Knowles & Werner 1985; Werner & Knowles 1985) to examine the electronic wavefunction of the HNCO\u2013He complex. These computations showed that this wavefunction is dominantly described by a unique electron configuration (with a weight \u22650.93) over the grid used for the generation of this PES. This justifies hence the use of monoconfigurational ab initio methods. Accordingly, we applied the recently established methodology by Hochlaf and co-workers for mapping multidimensional PESs of weakly bound molecular systems with high accuracy and relatively low computational cost (Lique, Klos & Hochlaf 2010; Halvick et al. 2011; Ajili et al. 2013; Mathivon, Linguerri & Hochlaf 2013; Kalugina et al. 2014; Mogren Al Mogren et al. 2014). Briefly, these electronic computations were carried out using the explicitly correlated coupled cluster method with single, double and perturbative treatment of triple excitations (CCSD(T)-F12) (Adler, Knizia & Werner 2007; Knizia, Adler & Werner 2009) in connection with the augmented correlation-consistent aug-cc-pVTZ basis set of Dunning and co-workers (Dunning 1989; Kendall, Dunning & Harrison 1992). In addition, molpro default choices for the density fitting and resolution of identity basis sets have been applied (Yousaf & Peterson 2008). Benchmarks by Hochlaf and co-workers (Lique et al. 2010; Halvick et al. 2011; Ajili et al. 2013; Mathivon et al. 2013; Kalugina et al. 2014; Mogren Al Mogren et al. 2014) showed that results obtained from this highly correlated approach are close to those deduced using standard coupled cluster techniques extrapolated to the complete basis set (CBS) limit, whereas a strong reduction in CPU time and disc occupancy are observed. For illustration, we performed CCSD(T)\/aug-cc-pVXZ calculations (X = D, T, Q, 5) on the HNCO\u2014He cluster. Then the energies were extrapolated to the CBS limit. The comparison between CCSD(T)-F12\/aug-cc-pVTZ and CBS calculations is given in Table 1. It shows that the CCSD(T)-F12\/aug-cc-pVTZ results are off by 4 per cent (at the maximum) with those deduced from CBS extrapolation. We compare also our results with those done using the CCSD(T)\/aug-cc-pV5Z approach. We can clearly see that the CCSD(T)-F12\/aug-cc-pVTZ approach offers a good agreement with the CCSD(T)\/aug-cc-pV5Z calculations with a very reduced computational cost.","Citation Text":["Kalugina et al. 2014"],"Citation Start End":[[864,884]]}
{"Identifier":"2021AandA...650A.155Z__Oh_et_al._2012_Instance_1","Paragraph":"Many factors can affect the prevalence of AGN activity. One important question is how gas is brought down to the galaxy center to fuel supermassive black holes (SMBHs). In the literature, two kinds of mechanisms are proposed. One is the internal secular evolution process. The torque induced by non-axisymmetric galactic structures can drive slow and significant inflow (Kormendy & Kennicutt 2004; Hopkins & Quataert 2011; Sellwood 2014; Fanali et al. 2015). The galactic bar is one of the most prominent non-axisymmetric structures and it exists in about 40% of spiral galaxies (Oh et al. 2012). In addition, there is evidence demonstrating that bars can enhance star formation in the central regions of galaxies (e.g. Oh et al. 2012; Chown et al. 2019). However, the question of whether galactic bars can significantly affect AGN activity is still under debate (Arsenault 1989; Mulchaey & Regan 1997; Oh et al. 2012; Galloway et al. 2015; Goulding et al. 2017; Alonso et al. 2018). Other mechanisms, such as galaxy merger and interaction, are also expected to displace the angular momentum of the gas and transport the gas inward (e.g. Hopkins et al. 2006; Di Matteo et al. 2008; Bhowmick et al. 2020). Similarly to studies of secular evolution, observational evidence for this scenario is also mixed. Some studies have found significant environmental dependence of AGN activity (e.g. Koulouridis et al. 2006; Koss et al. 2010; Ellison et al. 2011; Sabater et al. 2013; Khabiboulline et al. 2014; Lackner et al. 2014; Satyapal et al. 2014; Hong et al. 2015; Kocevski et al. 2015; Goulding et al. 2018; Gao et al. 2020), while others have found no or only weak environmental effects (e.g. Grogin et al. 2005; Li et al. 2006a, 2008; Pierce et al. 2007; Ellison et al. 2008; Gabor et al. 2009; Darg et al. 2010; Wang & Li 2019; Man et al. 2019). The contradictory results may be caused by the difference in AGN selection criterion, observational bias, control sample, and environmental indicator used. As we show below, understanding the environmental effects on AGNs also requires knowledge about the evolutionary status of their host galaxies, as it can help us to better understand how to construct control samples and to adopt appropriate environmental indicators.","Citation Text":["Oh et al. 2012"],"Citation Start End":[[580,594]]}
{"Identifier":"2021MNRAS.506.1438K__Schneider_et_al._2012_Instance_1","Paragraph":"The oldest and most metal-poor stars contain a fossil record of the star formation processes in the early Universe. These first stars would have been composed solely of hydrogen, helium, and trace amounts of lithium (Steigman 2007; Cyburt et al. 2016); without metals to efficiently cool the gas, large Jeans masses, and thereby very massive stars ($M_* \\gtrsim 100 \\, \\mathrm{M}_{\\odot }$) are predicted to have formed (Silk 1983; Tegmark et al. 1997; Bromm, Coppi & Larson 1999; Abel, Bryan & Norman 2002; Yoshida et al. 2006). In addition, improved gas fragmentation models (e.g. Clark et al. 2011; Schneider et al. 2012) and the discovery of very low-mass ultrametal-poor stars ([Fe\/H]\u2009\u2009\u22124; Keller et al. 2014; Starkenburg et al. 2017b; Schlaufman, Thompson & Casey 2018; Nordlander et al. 2019) have suggested that lower mass, long-lived stars may have also formed in these pristine environments (e.g. Nakamura & Umemura 2001; Wise et al. 2012). Notably, if \u22640.8 M\u2299 stars were to form, they could still exist today on the main sequence, and are expected to have undergone very little surface chemical evolution over their lifetimes. Detailed studies could provide invaluable constraints on (1) the nucleosynthetic yields from the first stars and the earliest supernovae (e.g. Heger & Woosley 2010; Ishigaki et al. 2014; Tominaga, Iwamoto & Nomoto 2014; Clarkson, Herwig & Pignatari 2018; Nordlander et al. 2019), (2) the physical conditions in the high-redshift universe (where stars are too faint to be observed individually; Tegmark et al. 1997; Freeman & Bland-Hawthorn 2002; Cooke & Madau 2014; Hartwig et al. 2018; Salvadori et al. 2019), (3) the primordial initial mass function (e.g. Greif et al. 2012; Susa, Hasegawa & Tominaga 2014), and (4) early Galactic chemical evolution processes (see Sneden, Cowan & Gallino 2008; Tolstoy, Hill & Tosi 2009; Roederer et al. 2014; Yoon et al. 2016; Wanajo et al. 2018; Kobayashi, Karakas & Lugaro 2020, and references therein).","Citation Text":["Schneider et al. 2012"],"Citation Start End":[[602,623]]}
{"Identifier":"2021ApJ...908...95H__Sargent_et_al._2012_Instance_1","Paragraph":"Star-forming galaxies at redshifts z \u223c 1\u20133 probe the cosmic epoch when most of the stellar mass assembly in the universe took place (Madau & Dickinson 2014, and references therein). A better understanding of star formation (SF) during this epoch is therefore imperative to understand SF across cosmic time. Locally, less than 5% of the galaxy population has a star formation rate (SFR) that is significantly higher than the empirical main sequence for star-forming galaxies, i.e., the tight correlation (\u223c0.3 dex) between the SFR and stellar mass, M\u22c6 (Brinchmann et al. 2004; Elbaz et al. 2007, 2011; Noeske et al. 2007; Goto et al. 2011; Rodighiero et al. 2011; Sargent et al. 2012; Whitaker et al. 2012, 2014; Salmon et al. 2015). These often-called starburst galaxies, with an IR luminosity LIR \u223c (0.1\u20135) \u00d7 1012 L\u2299 (e.g., Sanders & Mirabel 1996; Downes & Solomon 1998), become increasingly more common at high z. In fact, (sub)millimeter number counts reveal that galaxies with LIR > 1012\u201313 L\u2299, at z > 0.5, are many hundreds of times more likely to exist than in the local universe (Blain et al. 2002; Chapman et al. 2005; Berta et al. 2011; Magnelli et al. 2011; B\u00e9thermin et al. 2012; Magnelli et al. 2013; Casey et al. 2013, 2014; Geach et al. 2013; Simpson et al. 2014; Strandet et al. 2016; Brisbin et al. 2017). Meanwhile, the cosmic molecular gas density also peaks at z \u223c 1\u20133 (Decarli et al. 2014, 2016a, 2016b, 2019; Walter et al. 2014; Lentati et al. 2015; Pavesi et al. 2018; Liu et al. 2019; Riechers et al. 2019). This suggests a strong link between molecular gas and SF. Rest-frame far-IR (FIR) measurements of spectral lines and thermal dust continuum emission have been used to investigate the cooling and heating processes of the interstellar medium (ISM) in star-forming galaxies; however, the physical conditions at high z are still, in general, poorly investigated (Popesso et al. 2012; Bothwell et al. 2013; Carilli & Walter 2013; Genzel et al. 2013; Yang et al. 2017; Tacconi et al. 2018, 2020; Aravena et al. 2020; Birkin et al. 2020; Boogaard et al. 2020; Lenki\u0107 et al. 2020).","Citation Text":["Sargent et al. 2012"],"Citation Start End":[[663,682]]}
{"Identifier":"2019MNRAS.487.3702O__Owen_et_al._2011b_Instance_2","Paragraph":"The photoevaporation model successfully explains the \u2018two-time-scale\u2019 nature of protoplanetary disc evolution, where the inner regions of protoplanetary discs appear to evolve slowly on Myr time-scales, before dispersing on a much more rapid time-scale (e.g. Kenyon & Hartmann 1995; Ercolano, Clarke & Hall 2011; Koepferl et al. 2013; Ercolano et al. 2014). Furthermore, slow-moving (\u223c5\u201310 km s\u22121) ionized winds are observed to be occurring in many nearby discs hosting young stars (e.g. Hartigan, Edwards & Ghandour 1995; Pascucci & Sterzik 2009; Rigliaco et al. 2013) and are consistent with the photoevaporation model (Alexander 2008; Ercolano & Owen 2010; Pascucci et al. 2011; Owen, Scaife & Ercolano 2013a; Ercolano & Owen 2016). The photoevaporation model can also explain a large fraction of observed \u2018transition discs\u2019 (e.g. Owen & Clarke 2012; Espaillat et al. 2014), specifically those with holes \u227210 au and accretion rates \u227210\u22129 M\u2299 yr\u22121 (Owen et al. 2011b) and even those with larger holes and higher accretion rates in more recent models that incorporate CO depletion in the outer disc (Ercolano, Weber & Owen 2018). Transition discs are protoplanetary discs with evidence for a large hole or cavity in their discs (e.g. Espaillat et al. 2014), but they are known to be a heterogeneous class of objects (e.g. Owen & Clarke 2012; van der Marel et al. 2016) and their origins are not always clear. However, a specific prediction of the standard photoevaporation scenario is that there should be a large number of transition discs with hole sizes \u227310 au but that are no longer accreting. This final long-lived stage of disc dispersal gives rise to transition discs which have lifetimes between 105 and 106 yr, but remain optically thick \u2013 \u2018relic discs\u2019 (Owen et al. 2011b). The long disc lifetimes emerge from the simple fact that discs store most of their mass at large radii, but photoevaporative clearing proceeds from the inside out, so it will always take longer to remove the larger disc mass that resides at larger distance. While several discs satisfy this criterion (Dong et al. 2017), the number of observed non-accreting transition discs with large holes falls far below the theoretical expectations (Owen et al. 2011b; Owen & Clarke 2012). Studies by Cieza et al. (2013) and Hardy et al. (2015) showed that optically thick relic discs are rare and many non-accreting stars that show evidence for a circumstellar disc are more consistent with young, radially optically thin, debris discs.","Citation Text":["Owen et al. 2011b"],"Citation Start End":[[2222,2239]]}
{"Identifier":"2022MNRAS.513.4645S__Stathopoulos_et_al._2021_Instance_1","Paragraph":"The broad-band SED of blazars, in a log(\u03bdF\u03bd) versus log(\u03bd) representation, shows two prominent broad components, one (low-energy component) peaking form far-infrared frequencies to X-ray energies and another (high energy component) peaking at MeV\/GeV energies. The peak of the low-energy component (\u03bds) is used to further classify blazars as high synchrotron peaked BL Lacs (HBL when \u03bds > 1015 Hz), intermediate synchrotron peaked BL Lacs (IBL when 1014  \u03bds  1015 Hz), or low synchrotron peaked BL Lacs (LBL when \u03bds  1014 Hz) objects (Padovani & Giommi 1995; Abdo et al. 2010). Sometimes the synchrotron peak can reach energies as high as \u223c1 keV, (\u223c2 \u00d7 1017 Hz) or beyond, showing what is considered to be extreme behaviour, even for these highly peculiar sources (e.g. Giommi, Menna & Padovani 1999; Costamante et al. 2001; Biteau et al. 2020). Such a high synchrotron peak was first observed during a flare of Mkn 501 (Pian et al. 1998), and subsequently in many other objects (e.g. Costamante et al. 2018; Sahakyan 2020a). Independently of the location of the peak, the low-energy part of the SED is generally interpreted as synchrotron emission from the relativistic electrons in the jet. A proton synchrotron origin of the high energy end of this component during X-ray flares has also been considered (Mastichiadis & Petropoulou 2021; Stathopoulos et al. 2021). The nature of the high energy (HE; >100 MeV) component is instead still under debate. Within one-zone leptonic scenarios, the second component originates from inverse Compton scattering of the synchrotron photons (SSC) by the electron population producing the low-energy component (Ghisellini, Maraschi & Treves 1985\n; Maraschi, Ghisellini & Celotti 1992; Bloom & Marscher 1996). Depending on the location of the emission region, the photons external to the jet [e.g. photons from the disc, or those reprocessed from the broad-line region (BLR) or those from the infrared torus] can up-scatter, producing the second component [external inverse Compton (EIC); Sikora, Begelman & Rees 1994; B\u0142a\u017cejowski et al. 2000; Ghisellini & Tavecchio 2009\n]. On the other hand, the HE component can be also produced from the interaction of relativistic protons either from their synchrotron emission (M\u00fccke & Protheroe 2001) or from the secondary particles from pion decay (Mannheim & Biermann 1989; Mannheim 1993; M\u00fccke & Protheroe 2001; M\u00fccke et al. 2003; B\u00f6ttcher et al. 2013). Recently, after associating TXS 0506+056 with the IceCube-170922A neutrino event (IceCube Collaboration 2018a, b; Padovani et al. 2018) the lepto-hadronic scenarios, when both electrons and protons contribute to the HE emission, have become more attractive. These models also predict very high energy (VHE; >100 GeV) neutrinos observable by the IceCube detector (Ansoldi et al. 2018; Keivani et al. 2018; Murase, Oikonomou & Petropoulou 2018; Padovani et al. 2018; Sahakyan 2018; Cerruti et al. 2019; Gao et al. 2019; Righi, Tavecchio & Pacciani 2019; Sahakyan 2019; Gasparyan, B\u00e9gu\u00e9 & Sahakyan 2022).","Citation Text":["Stathopoulos et al. 2021"],"Citation Start End":[[1341,1365]]}
{"Identifier":"2016MNRAS.458.3181C__Trujillo_et_al._2011_Instance_2","Paragraph":"To explain the observed evolution, the physical processes invoked have to result in a large growth in size but not in stellar mass, nor drastic increase in the star formation rate. Most plausible candidates are mass-loss driven adiabatic expansion (\u2018puffing-up\u2019) (e.g. Fan et al. 2008, 2010; Ragone-Figueroa & Granato 2011) and dry mergers scenarios (e.g. Bezanson et al. 2009; Naab, Johansson & Ostriker 2009; Trujillo, Ferreras & de La Rosa 2011). In the former scenario, galaxies experience a mass-loss from wind driven by active galactic nuclei (AGNs) or supernovae feedback, which lead to an expansion in size due to a change in the gravitational potential. In the latter, mergers either major involving merging with another galaxy of comparable mass, or minor that involves accretion of low mass companions, have to be dry to keep the low star formation rate (Trujillo et al. 2011). Nevertheless, major mergers are not compatible with the observed growth in mass function in clusters as well as the observed major merger rates since z \u223c 1 (e.g. Nipoti, Londrillo & Ciotti 2003; Bundy et al. 2009). On the other hand, minor mergers are able to produce an efficient size growth (see e.g. Trujillo et al. 2011; Shankar et al. 2013). The rates of minor mergers are roughly enough to account for the size evolution only up to z \u2272 1 Newman et al. (2012), at z \u223c 2 additional mechanisms are required (e.g. AGN feedback-driven star formation Ishibashi, Fabian & Canning 2013). In addition, the effect of continual quenched galaxies on to the red sequence as well as morphological mixing (known as the \u2018progenitor bias\u2019) further complicates the situation (e.g van Dokkum & Franx 2001). Processes that are specific in clusters such as harassment, strangulation and ram-pressure stripping (e.g. Treu et al. 2003; Moran et al. 2007) might play an important role in quenching and morphologically transforming galaxies. Several studies have already shown that the progenitor bias has a non-negligible effect on the size evolution (e.g. Saglia et al. 2010; Valentinuzzi et al. 2010b; Carollo et al. 2013; Poggianti et al. 2013; Beifiori et al. 2014; Delaye et al. 2014; Belli, Newman & Ellis 2015; Shankar et al. 2015).","Citation Text":["Trujillo et al. 2011"],"Citation Start End":[[1192,1212]]}
{"Identifier":"2021ApJ...915..128Y__C\u00f4t\u00e9_et_al._2019a_Instance_1","Paragraph":"These two isotopes are made by the rapid neutron-capture (r) process, and typical estimates for the time interval at which r-process nucleosynthetic events that are believed to enrich a parcel of gas in the Galaxy range between 200 and 500 Myr (Hotokezaka et al. 2015; Tsujimoto et al. 2017; Bartos & Marka 2019; C\u00f4t\u00e9 et al. 2021). Therefore the case of 247Cm\/129I is the best example of Regime 3 because \u03c4eq = 5085 Myr (Table 2) is much longer than \u03b3, while each \u03c4 (\u224322.5 and 22.6 Myr, respectively) is much shorter than \u03b3. The ratios to the long-lived or stable references isotopes, 247Cm\/235U and 129I\/127I, allow us to derive a TLE, for example, for the specific \u03b3 value of 316 Myr considered in Table 3, and derive typical production ratios of 1.35 for 129I\/127I and 0.3 for 247Cm\/235U. While our TLE values are not perfectly compatible with each other, the more detailed analysis shown by C\u00f4t\u00e9 et al. (2021) demonstrates that there is compatibility for TLE in the range between 100\u2013200 Myr, depending on the exact choice of the K parameter (C\u00f4t\u00e9 et al. 2019a), \u03b3, and the production ratios. The short mean lives of 247Cm and 129I ensure that there is no memory from previous events, while the long \u03c4eq of 247Cm\/129I instead ensures that this ratio did not change significantly during TLE and has a high probability to be within the 10% of the production ratio. Therefore the production ratio of the last r-process event that polluted the ESS material can be accurately determined directly from the ESS ratio. If we assume that the last event produced a 247Cm\/232Th ratio similar to the average predicted by Goriely & Janka (2016) and assume a solar ratio for 127I\/232Th, then we find an inconsistency between the numbers in the last two columns of Table 3. The back-decayed value is more than five times lower than the assumed production ratio, which indicates a weaker production of the actinides from this last event with respect to the production ratios that we are using here. The number in the last column therefore represents a unique constraint on the nature of the astrophysical sites of the r process in the Galaxy at the time of the formation of the Sun and needs to be compared directly to different possible astrophysical and nuclear models (C\u00f4t\u00e9 et al. 2021).","Citation Text":["C\u00f4t\u00e9 et al. 2019a"],"Citation Start End":[[1047,1064]]}
{"Identifier":"2021MNRAS.507.5567D__Urquhart_&_Soria_2016_Instance_1","Paragraph":"The source shows some preferred tracks in its movement along the HID. Dips occur as the source is in the SUL regime, either from SUL2 or from SUL3 regions; ingress and egress times populate the D3 and D2 spectra, whereas the D1 spectrum comes from the time segments characterized by the lowest count rate (deep dip). We noted that during the longest dips, the source occasionally switched to harder flaring episodes in D2 state. Most dips show the following pattern: SUL2\/SUL3 \u2192 D3 \u2192 D2 \u2192 D1 \u2192 D2 \u2192 D3 \u2192 SUL2\/SUL3. The passage from the SSUL to SUL state is only occasionally observed. We selected seven regions on the HID and extracted the corresponding spectra. For the normal branch spectra, the continuum emission is well fitted with two thermal components: a soft blackbody and a disc multicolour blackbody. The blackbody emission represents the bulk of the emitted power in each spectrum. It is likely due to strong reprocessing in an optically thick environment formed at Rsph, where disc inflow is mainly inflated by internal radiation pressure. This picture is consistent with the observed low temperature, large radius, and super-Eddington luminosity (see Shen et al. 2015; Soria & Kong 2016; Urquhart & Soria 2016; Guo et al. 2019). The hotter component, which we fitted using a diskbb, dominates the emission above 2\u2009keV. Although its origin is still debated, it might come from internal hard X-ray emission, which has been inefficiently reprocessed, or simply scattered along our line of sight. Alternatively, it can be continuum emission (bremsshtrahlung and\/or Comptonization) from an extended optically thin plasma where the emission lines are produced, or a tail of the blackbody emission which has been Compton upscattered in a coronal environment around the photosphere. In addition to the continuum, we added multiple emission and absorption lines derived by the combined averaged PN \/ RGS analysis to mimic the emitting and absorbing plasmas found in Paper I. Their shifts suggest different Doppler motions in several states as a consequence of a velocity field which changes depending on the launching site and on the geometry of the system. This seems supported by correlations among the parameters of the emission lines and the underneath hard X-ray flux (see Table 3). However, given the limited energy resolution of the EPIC we are not able to distinguish between a varying ionization state of the plasma, the effect of a different line broadening for the different ionized species, a complex absorption\/emission pattern. A thorough study of the lines is left to a dedicated forthcoming study. For consistency the same spectral model has been applied also to the dipping branch spectra, although the physical conditions in and out of the dips might be different.","Citation Text":["Urquhart & Soria 2016"],"Citation Start End":[[1202,1223]]}
{"Identifier":"2015AandA...577A..46D__Carilli_&_Walter_2013_Instance_1","Paragraph":"The basically linear correlation between L\u2032CO[5\u20134] and LIR is different from findings for low-J CO, where considering all galaxies together yields a trend more consistent with a slope of 1.7 (e.g., see the classical work of Solomon & van der Bout 2005) or, when distinguishing SB and MS galaxies, separate relations with slopes of 1.2 and different normalizations (e.g., Sargent et al. 2014). This suggests that CO[5\u20134] could behave as a tracer of the star-forming gas, meaning that the gas is dense enough to be directly related to the amount of stars formed, in a way that is insensitive of whether the star formation occurs in a quiescent disk or inside a violent merger. If this is the case, then CO[5\u20134] could be used as a more efficient and accessible tracer of the dense gas than standard molecules such as HCN or HCO, which also scale basically linearly with LIR, but are intrinsically 10\u201320 times fainter and harder to observe (Gao & Solomon 2004a,b; Gao et al. 2007)5. At the same time, given its fairly high frequency, CO[5\u20134] can be observed with reasonably high, subarcsec spatial resolution with both ALMA and IRAM PdBI. We caution, however, that the critical density of CO[5\u20134] is an order of magnitude lower than that of HCN (Carilli & Walter 2013). Strictly speaking, these are not equivalent as tracers. In addition, CO[5\u20134] has an excitation temperature of ~83 K, which is much higher than low-J HCN transitions. As such, the low-J lines of these molecules are clearer tracers of density because they do not require the medium to be both warm and dense; we will still need HCN[1\u20130] or other dense gas tracers at [1\u20130] to obtain the total dense gas including the much colder component that will be missed by high-J CO lines. We do not know, moreover, whether Eq. (4) manifests a direct connection between CO[5\u20134] and LIR or if both CO[5\u20134] and LIR are driven in similar ways by some other parameter, in which case their spatial distribution inside galaxies might differ. Evidence from local galaxies suggests that CO[5\u20134] emission with Tkin> 50 K, as favored by our LVG modeling, might not be excited by the photodissociation region, but more likely by mechanical heating from shocks (e.g., Rosenberg et al. 2014; Meijerink et al. 2013), which most likely are driven by winds connected to the SF regions. Future comparisons of high spatial resolution maps of CO[5\u20134] with HCN, for example, are required to clarify whether this is a viable approach or not, that is, to clarify if CO[5\u20134] can be used as a high-fidelity tracer of the star-forming gas. ","Citation Text":["Carilli & Walter 2013"],"Citation Start End":[[1242,1263]]}
{"Identifier":"2017MNRAS.471.2848L__Sunyaev_&_Titarchuk_1980_Instance_1","Paragraph":"For LLAGNs, there are two competing components for the emission at optical\/UV wavebands. One component is the thermal emission from the outer truncated SSD, which peaks in optical or UV bands, depending on the location of the truncated radius. The other component is the Compton scattering of synchrotron emission from the inner hot accretion flow (Manmoto et al. 1997; Yuan & Narayan 2014). Since our sample is limited to sources whose \u03bb  10\u22123, we argue that the emission at 2500 \u00c5 mainly comes from the inner hot accretion flow.3 Interestingly, for hot accretion flows around supermassive BHs (the LLAGN case), the optical\/UV emission usually is the first Compton up-scattered bump (Manmoto et al. 1997; Yuan & Narayan 2014). The emission from such a process has a positive correlation with $\\dot{m}$, as $\\dot{m}$ will mainly determine the optical depth (surface density), a quantity that controls the probability of Compton scattering (Sunyaev & Titarchuk 1980; Dermer, Liang & Canfield 1991). On the other hand, the first Compton bump is even more sensitive to the energy of electrons (a.k.a. the electron temperature), as it determines the location (in wavebands) of the peak. For a given $\\dot{m}$, larger \u03b1 means lower density, or equivalently lower radiative cooling to the electrons. Consequently, the electrons will be more energetic at higher temperatures (Xie & Yuan 2012). The dependence of emission at optical\/UV band on the parameter \u03b1 relies on detailed numerical calculations, where Manmoto et al. (1997) reported that the emission at 2500 \u00c5 has a positive correlation with \u03b1. For the above reasons, we simply write LUV as\n(6)\r\n\\begin{equation}\r\nL_{\\rm UV}\\sim (\\dot{m} \\alpha )^b,\r\n\\end{equation}\r\nwhere b \u223c 2. The estimation of b comes from the fact that the first-order Compton scattering (responsible for the optical\/UV emission) depends on the product of the seed photon flux (synchrotron emission, $\\propto \\dot{m}^{0.5-1}$; cf. Mahadevan 1997; Yuan et al. 2015; note that electron temperature also depends on $\\dot{m}$ for the expression in Mahadevan 1997) and optical depth in a vertical direction ($\\propto \\dot{m}^{1-2}$).","Citation Text":["Sunyaev & Titarchuk 1980"],"Citation Start End":[[940,964]]}
{"Identifier":"2018AandA...620A..99G__Shappee_&_Stanek_(2011)_Instance_1","Paragraph":"It is obviously of fundamental importance in this process to provide a very accurate absolute calibration of the Cepheid PL relation. This can be achieved in different ways, using Cepheids in our own Galaxy that have accurate parallax measurements, or using extragalactic Cepheids for which the distances of their host galaxies have been very accurately determined with some Cepheid-independent method. A critical aspect of the calibration of the PL relation is a precise determination of its possible dependence on the metallicity of the Cepheid variables. Without an accurate knowledge of this \u201cmetallicity effect\u201d a distance measurement to a galaxy accurate to 1% with a Cepheid PL relation is clearly not possible. In the past a lot of work was done with a variety of methods to determine the metallicity dependence of the Cepheid PL relation. Early work includes the studies of Gould (1994), Kennicutt et al. (1998), and Sakai et al. (2004). More recent studies on the metallicity effect are those of Shappee & Stanek (2011), Mager et al. (2013), and Fausnaugh et al. (2015), among others. A very detailed compilation of metallicity effect determinations prior to 2008 can be found in Table 1 of Romaniello et al. (2008). These studies seem to indicate that the metallicity effect in the near-infrared JHK bands is small, perhaps even vanishing, while there is a significant effect in optical and mid-infrared photometric bands (e.g., Freedman & Madore 2011). While most studies yielded a negative sign of the metallicity effect in optical bands meaning that more metal-rich Cepheids are intrinsically brighter than their more metal-poor counterparts of the same pulsation period, the work of Romaniello et al. (2008) has yielded the opposite sign for the effect in the V band, meaning that metal-poor Cepheids are intrinsically brighter than their more metal-rich counterparts of the same pulsation period. Theoretical studies (e.g., Caputo et al. 2000; Bono et al. 2008) appear to support Romaniello\u2019s results, but the uncertainty on these determinations of the metallicity effect from pulsation theory still seems to be rather substantial. In the most recent work on the subject, Wielg\u00f3rski et al. (2017), using the extremely well-established Cepheid PL relations in the Magellanic Clouds in optical and near-infrared bands combined with the accurate distance determinations to the LMC (Pietrzy\u0144ski et al. 2013) and SMC (Graczyk et al. 2014) from late-type eclipsing binaries, found a metallicity effect compatible with zero in all bands. As a conclusion, there is still no consensus about the true effect of metallicity on Cepheid absolute magnitudes in different spectral regions. It is especially important to obtain a truly accurate determination of the metallicity effect in the near-infrared bands since these have been used, due to their much lower sensitivity to extinction, in the space-based work on H0 with the Hubble Space Telescope over the last two decades, and will be used in the near future with the James Webb Space Telescope.","Citation Text":["Shappee & Stanek (2011)"],"Citation Start End":[[1006,1029]]}
{"Identifier":"2022MNRAS.516..731B__Wang_et_al._2022_Instance_1","Paragraph":"With the results from Gaia and the large number of stars collected from dedicated surveys, a plethora of Galactic mass estimates have recently been made using using different methods and data samples. We summarize these most recent results in Fig. 8. Grouping together the works applying a distribution function method (Callingham et al. 2019; Eadie & Juri\u0107 2019; Posti & Helmi 2019; Vasiliev 2019; Li et al. 2020; Deason et al. 2021; Hattori et al. 2021; Correa Magnus & Vasiliev 2022; Shen et al. 2022; Slizewski et al. 2022; Wang et al. 2022) the weighted mean mass estimate is \u223c0.93 \u00b1 0.047 \u00d7 1012 M\u2299. The methods using high velocity stars with full 6D phase-space information have been used to estimate the mass of the Milky Way (Monari et al. 2018; Deason et al. 2019; Grand et al. 2019; Koppelman & Helmi 2021; Necib & Lin 2022) have a weighted mean mass of \u223c0.93 \u00b1 0.099 \u00d7 1012 M\u2299. Two recent works using either the Sagittarius or Magellanic Steams (Craig et al. 2021; Vasiliev et al. 2021) give a weighted mean mass of \u223c0.98 \u00b1 0.15 \u00d7 1012 M\u2299. The lightest mass measurements have come from rotation curve based methods (de Salas et al. 2019; Eilers et al. 2019; Ablimit et al. 2020; Cautun et al. 2020; Karukes et al. 2020; Jiao et al. 2021) with a weighted mean of \u223c0.76 \u00b1 0.023 \u00d7 1012 M\u2299. The works which use 6D satellite phenomenology compared to Milky Way-type simulations (Patel et al. 2018; Villanueva-Domingo et al. 2021) infer a weighted mean Galactic mass of \u223c0.97 \u00b1 0.0091 \u00d7 1012 M\u2299 and Rodriguez Wimberly et al. (2022) find their results are consistent with a Galactic mass in the range of 1\u20131.2 \u00d7 1012 M\u2299. Zaritsky et al. (2020) apply the timing argument to distant Milky Way halo stars to derive a lower limit to the Milky Way mass of 0.91 \u00d7 1012 M\u2299. Within the uncertainty, our BHB mass estimates are in agreement with Zaritsky et al. (2020), although our KG mass estimates and a few of the recent mass estimates which we plot in Fig. 8 are lower than this limit.","Citation Text":["Wang et al. 2022"],"Citation Start End":[[528,544]]}
{"Identifier":"2019ApJ...875...91D___2014_Instance_1","Paragraph":"Detailed IBEX results also showed that (1) the \u223c0.5\u20136 keV ENA spectral indices in the GDF are ordered in latitude, with those observed near mid- and low latitudes between \u00b145\u00b0 being significantly steeper compared with those observed at higher latitudes (Funsten et al. 2009b; McComas et al. 2009b), and (2) the corresponding ENA fluxes exhibit temporal variations on an annual scale that reflect the decreasing output of the SW, which resulted in a general decrease and flattening of the SW dynamic pressure during 2007\u20132010 (McComas et al. 2012, 2014). Desai et al. (2015, 2016; see also Dayeh et al. 2012) studied the energy dependence and latitudinal ordering of the \u223c0.5\u20136 keV ENA spectral indices \u03b3 in the GDF between 2009 and 2013 and showed that the latitudinal (\u03b8) dependence of \u03b3 is energy dependent and well represented by the cosine function\n1\n\n\n\n\n\nwith different offsets, amplitudes, and phase angles. Desai et al. (2016) further reported that the general decrease in the >1 keV ena fluxes from 2009\u20132011 to 2012\u20132013 occurred at all latitudes and that this decrease was not accompanied by corresponding changes in the spectral indices. Indeed, over the first five years of the IBEX mission, the ENA spectral indices in the GDF exhibited a strong and persistent energy-dependent, latitudinal ordering that reflected the solar minimum latitudinal profile of the fast and slow SW speeds in the inner heliosphere (e.g., Funsten et al. 2009b; Dayeh et al. 2011; Livadiotis et al. 2011; Schwadron et al. 2014; Desai et al. 2015, 2016; Zirnstein et al. 2017). Based on the reconstructed latitudinal and temporal profiles of SW parameters from 2011 to 2014 (Sok\u00f3\u0142 et al. 2013, 2015), i.e., during the ascending phase of solar cycle 24 and following the predictions of McComas et al. (2012, 2014), Desai et al. (2016) further suggested that the GDF in 2014\u20132017 would either stabilize or increase, and that the latitudinal ordering of the corresponding spectral indices would most likely be disrupted.","Citation Text":["McComas et al.","2014"],"Citation Start End":[[526,540],[547,551]]}
{"Identifier":"2021MNRAS.500.1784D__Larson,_Tinsley_&_Caldwell_1980_Instance_1","Paragraph":"Several physical processes affect galaxies inside clusters in a simultaneous way. One of these mechanisms is the ram pressure stripping (e.g. Gunn & Gott 1972; Abadi, Moore & Bower 1999; Book & Benson 2010; Steinhauser, Schindler & Springel 2016). This process can remove an important fraction of the cold gas from galaxies, resulting in the inhibition of star formation. Although this mechanism is more effective at the central regions of massive clusters, it has been reported in less massive systems (e.g. Rasmussen, Ponman & Mulchaey 2006; Jaff\u00e9 et al. 2012; Hess & Wilcots 2013). Ram pressure stripping occurs as galaxies move at high speeds through the hot ionized gas of the intracluster medium, which collides with the cold gas of the galaxies and removes it. The warm gas from the galactic halo can also be removed by the gas of the intracluster medium, a process known as starvation (e.g. Larson, Tinsley & Caldwell 1980; Balogh, Navarro & Morris 2000; McCarthy et al. 2008; Bekki 2009; Bah\u00e9 et al. 2013; Vijayaraghavan & Ricker 2015). This process can cut-off further gas cooling from the galaxy\u2019s halo gas that fuels future star formation. Kawata & Mulchaey (2008) predicted that starvation can act in galaxy groups as well. Another physical process that works on galaxies in their passage through the deep potential well of the cluster is tidal stripping (e.g. Zwicky 1951; Gnedin 2003a; Villalobos, De Lucia & Murante 2014). It can induce a central star formation burst (Byrd & Valtonen 2001), bar instabilities (\u0141okas et al. 2016), changes in the pattern of the spiral arms (Semczuk, \u0141okas & del Pino 2017), and truncate dark matter haloes (e.g. Gao et al. 2004; Limousin et al. 2009). In the outskirts of clusters, mechanisms like galaxy\u2013galaxy interaction, known as harassment, are more effective (e.g. Moore et al. 1996; Moore, Lake & Katz 1998; Gnedin 2003b; Smith et al. 2015). Most of the processes mentioned above tend to decrease or to completely suppress the star formation in galaxies. As a consequence, galaxies in clusters are typically red, early-type, with an old stellar population, and have little or none star formation at all.","Citation Text":["Larson, Tinsley & Caldwell 1980"],"Citation Start End":[[899,930]]}
{"Identifier":"2015ApJ...801..103G___2014_Instance_1","Paragraph":"Within the framework of the fireball shock model, Pe'er et\u00c2 al. (2007) proposed a method to infer central engine parameters using observed data. With the measured temperature and flux of an identified thermal component in the spectrum, along with a flux ratio between thermal and non-thermal components, one may infer the size of the jet at the base of the outflow, r0, and the dimensionless entropy of the outflow,  (which is also the bulk Lorentz factor of the outflow, if the photosphere radius is greater than the fireball coasting radius). Some authors have applied this method to some Fermi GRBs (Iyyani et\u00c2 al. 2013; Preece et\u00c2 al. 2014; Ghirlanda et\u00c2 al. 2013). The derived central engine parameters are sometimes ad hoc or inconsistent. For instance, the analyses for GRB 110721A (Iyyani et\u00c2 al. 2013) and for GRB 130427A (Preece et\u00c2 al. 2014) led to a curious conclusion that the bulk Lorentz factor of the outflow of different layers are decreasing with time. This would lead to no, or at most very inefficient, internal shock emission. Yet both bursts have dominant non-thermal emission. More curiously, the data of GRB 110721A (Iyyani et\u00c2 al. 2013) require that r0 is rapidly varying with time by 2\u00e2\u0080\u00933 orders of magnitudes. This is hard to imagine given the well-believed paradigm of the GRB central engine: If the engine is naked, the size of the engine (a hyper-accreting black hole or a millisecond magnetar) is around r0 \u00e2\u0088\u00bc 107\u00e2\u0080\u0089cm; if an extended envelope of a collapsar progenitor is considered, the fireball may be \u00e2\u0080\u009cre-born,\u00e2\u0080\u009d with r0 \u00e2\u0088\u00bc R*\u00ce\u00b8j \u00e2\u0088\u00bc 109R*, 10\u00ce\u00b8j, \u00e2\u0088\u00921\u00e2\u0080\u0089cm (where R* is the size of the progenitor star, and \u00ce\u00b8j is the jet opening angle). If one considers the depletion of the envelope, r0 should decrease with time. However, Iyyani et\u00c2 al. (2013) showed that r0 increases from 106\u00e2\u0080\u0089cm to 108\u00e2\u0080\u0089cm early on, and then decreases mildly after 2\u00c2 s. These absurd conclusions suggest that the starting point of the analysis, i.e., the assumption of a pure fireball model, might not be valid. It is interesting to see whether a hybrid ejecta photosphere model may solve the problem. Incidentally, Ghirlanda et\u00c2 al. (2013) analyzed another burst, GRB 100507, using the fireball framework (Pe'er et\u00c2 al. 2007), but found that the derived r0 remains constant and reasonable. The jet composition of that burst may be more close to a fireball. It would be interesting to see whether a general theoretical framework can be established, which may be reduced to the standard fireball framework when \u00cf\u00830  1.","Citation Text":["Preece et\u00c2 al. 2014"],"Citation Start End":[[624,643]]}
{"Identifier":"2022AandA...659A..34C__Roth_et_al._2016_Instance_1","Paragraph":"Guillochon et al. (2014) suggest that the broad lines seen in TDEs are produced in a BLR structure (analogous to the one seen in AGNs) which lies above and below the forming accretion disk. Furthermore, they comment that if this is true, it would be reasonable to expect that these two astrophysical phenomena (i.e., AGNs and TDEs) should have many similarities in terms of velocity structures, as well as the components of the structure that play a key role in the production of the radiated light and the emergent emission lines. If their prediction is valid, it is not surprising that TDEs show the observed time lags pointed out in this work, as time lags (caused by the response of line luminosities to variations in the continuum) are commonly observed in AGNs. It has been suggested that the line emitting region must be stratified, where Helium is closer to the black hole than Hydrogen (Guillochon et al. 2014; Roth et al. 2016). Regular AGNs are well known to be stratified as well with He\u202fII, He\u202fI, H\u03b2 and H\u03b1 being emitted from closest to furthest from the black hole (e.g., see Clavel et al. 1991; Peterson & Wandel 1999). To test this, we investigated how lagged the Helium lines are compared to Hydrogen (see Sects. 5.1.1 and 5.1.2). We chose He\u202fI 5876 \u00c5 as this line is more isolated and easier to measure than He\u202fII. We find a difference, depending on the spectroscopic type: H TDEs show larger lags in H\u03b1 than He\u202fI. On the other hand, for N\u202fIII Bowen TDEs, He\u202fI has similar lag values to H\u03b1. The average (lower-limit) lag value for H\u03b1 is 40 days for the H TDEs and 12 days for the N\u202fIII Bowen TDEs (without including ASASSN-15oi because it does not belong in either of these categories as it is a He TDE with weak Balmer lines in some epochs). The average (lower-limit) lag value for He\u202fI is 15.48 light days for N\u202fIII Bowen TDEs. For the H TDEs, only AT2018zr shows a clear lag in He\u202fI since for ASASSN-14ae and LSQ12dyw we can only place upper-limits). These results are visualized in Fig. 18 where we plot the time lags against the blackbody radii values at peak for each TDE. The H\u03b1 lags (filled symbols) increase with increasing RBB (the dashed line is a fit to H\u03b1 and the dot-dashed one is the same without fitting the three lower-limit points). On the other hand, He\u202fI (empty symbols) does not show such a correlation and the lags remain small (mostly upper-limits) even for the larger radii. This picture can be understood by a combination of two factors: (i) N\u202fIII Bowen TDEs have smaller radii than H TDEs (van Velzen et al. 2021; ii) Helium lies deeper in the photosphere than Hydrogen (Roth et al. 2016). The ionization energy of neutral Hydrogen is 13.6 eV while the one of neutral Helium is 24.6 eV hence Helium would need a larger temperature in order to get ionized, hinting that it is indeed lying closer to the black hole than Hydrogen \u2013 consistent with a stratified photosphere. We find here, however, that these differences are minimized in the more compact Bowen TDEs. Guillochon et al. (2014) focus their study on PS1-10jh, the only TDE in our sample that shows no Hydrogen at all, but only strong He\u202fII. They suggest that the absence of Hydrogen lines indicates that the accretion disk had not yet extended to the distances required to produce the Hydrogen lines. Their calculations imply that He\u202fII 4686 \u00c5 should be produced at a distance of \u223c2.1 light days from the SMBH (see their Fig. 7). Such small lags are beyond the precision that can be attained with the present data, as daily spectroscopic observations would be required.","Citation Text":["Roth et al. 2016","Roth et al. 2016"],"Citation Start End":[[920,936],[2615,2631]]}
{"Identifier":"2016MNRAS.461..248S__Munari_et_al._2013_Instance_2","Paragraph":"In Sif\u00f3n et al. (2013), we used the \u03c3\u2013M200 scaling relation of Evrard et al. (2008) to estimate dynamical masses. As discussed in Section 1, the scaling relation of Evrard et al. (2008) was calibrated from a suite of N-body simulations using DM particles to estimate velocity dispersions. However, the galaxies, from which velocity measurements are made in reality do not sample the same velocity distribution as the DM particles. They feel dynamical friction and some are tidally disrupted, which distorts their velocity distribution and biases their dispersion (e.g. Carlberg 1994; Col\u00edn et al. 2000). Recent high-resolution hydrodynamical simulations of \u2018zoomed\u2019 cosmological haloes have shown that there is a significant difference between the velocity distributions of DM particles and galaxies themselves; whether galaxies (i.e. overdensities of stars in hydrodynamical simulations) or DM subhaloes are used makes comparatively little difference (Munari et al. 2013). Results from state-of-the art numerical simulations depend on the exact definition of a galaxy and the member selection applied, but the current consensus is that galaxies are biased high (i.e. at a given mass the velocity dispersion of galaxies or subhaloes is larger than that of DM particles) by 5\u201310 per cent with respect to DM particles (Lau et al. 2010; Munari et al. 2013; Wu et al. 2013), translating into a positive 15\u201320 per cent bias in dynamical masses when using DM particles. This is illustrated in Fig. 5: DM particles are not significantly impacted by either dynamical friction or baryonic physics; therefore, the scaling relations for DM particles are essentially the same for all simulations. In contrast, DM subhaloes are affected by baryons in such a way that including baryonic feedback (most importantly feedback from active galactic nuclei \u2013 AGN, but also from cooling and star formation) makes their velocity dispersions much more similar to those of simulated galaxies. This means we can rely on our analysis of the previous section, based on DM subhaloes, to correct the velocity dispersions measured for ACT clusters, and then estimate dynamical masses using predictions obtained either from galaxies or subhaloes. The difference between the Saro et al. (2013) and Munari et al. (2013) galaxy scaling relations depends on the details of the semi-analytic and hydrodynamical implementations used in Saro et al. (2013) and Munari et al. (2013), respectively. The different cosmologies used in the Millenium simulation (in particular, \u03c38 = 0.9; Springel et al. 2005) by Saro et al. (2013) and the simulations by (Munari et al. 2013, \u03c38 = 0.8) may also play a role.","Citation Text":["Munari et al. 2013"],"Citation Start End":[[1334,1352]]}
{"Identifier":"2020ApJ...899..147F__Tsai_et_al._2017_Instance_1","Paragraph":"The C\/O ratio varies across exoplanets\u2019 host star populations (Delgado Mena et al. 2010; Brewer & Fischer 2016; Brewer et al. 2017), and this variation is likely to be reflected in the composition of exoplanet atmospheres, assuming that they are formed with the same materials as their stars. Moreover, various processes in the protoplanetary disks and the planet formation process can affect the exoplanet compositions and have a significant impact on the final C\/O ratio (\u00d6berg et al. 2011; Mordasini et al. 2016; Espinoza et al. 2017; Madhusudhan et al. 2017). For these reasons, it is necessary to consider the effects of the C\/O ratio on the atmospheric chemistry and the formation of aerosols. Numerous studies have been performed using chemical models (Madhusudhan 2012; Moses et al. 2013; Venot et al. 2015; Tsai et al. 2017; Heng & Lyons 2016; Goyal et al. 2018; Drummond et al. 2019), but corresponding laboratory experiments are still largely nonexistent. Laboratory investigations can provide essential insight into the effects of the C\/O ratio on the atmospheric photochemistry and the formation of aerosols. In a previous work, we performed the first laboratory experiments dedicated to the study of the chemistry in hot Jupiter atmospheres (Fleury et al. 2019). This work focused on the chemistry in atmospheres with T > 1000 K and a C\/O ratio of 1 (representing C enhancement compared to the solar value of 0.54), because chemical models predict that the abundances of hydrocarbon and nitrile species increase by several orders of magnitude in these atmospheres compared to atmospheres with a low C\/O ratio (Venot et al. 2015). Therefore, they can be considered as better candidates for the formation of complex organic molecules with longer carbon chains. This first study revealed that photochemical aerosols could be produced at temperatures as high as 1500 K and that water could be efficiently formed through photochemical channels. In the present work, we performed new experiments to study the chemistry in hot Jupiter atmospheres at similar temperatures (1173\u20131473 K) but with lower C\/O ratios. We used a gas mixture of H2, H2O, and CO that represents the simplest plausible atmosphere for a hot Jupiter with a C\/O ratio  1. This new study, compared with our previous work, allows us to assess the evolution of the chemistry in hot Jupiter atmospheres as a function of the C\/O ratio and atmospheric composition.","Citation Text":["Tsai et al. 2017"],"Citation Start End":[[816,832]]}
{"Identifier":"2017AandA...606A..94D__Lis_et_al._(2015)_Instance_1","Paragraph":"Also, by studying different ratios we can get information about the masses of the progenitors enriching the ISM at the formation time of our stars. The work by Travaglio et al. (1999) showed that the best progenitors for reproducing the r-process contribution to the enrichment of the Galaxy are SNe II from stars with masses 8\u201310\u2009M\u2299. On the other hand, more massive SNe II of M\u2009>\u200915\u2009M\u2299 enriched the ISM with oxygen at earlier times since those massive stars evolve faster. As a consequence we can observe that the ratios of r-process elements with respect to oxygen are negative for low metallicities. In Fig. 14 we show the ratios between Eu-Y-Ba and O using the abundances of the oxygen line at 6158 \u00c5 derived for the same sample by Bertran de Lis et al. (2015). Since we do not have very metal-poor stars we cannot observe the behaviour of Y and Ba at very low metallicities, where they are thought to be mainly produced by the r- and not the s-process. We can see how [Eu\/O] has a less steep decline towards lower metallicities when compared to Ba and Y. This occurs because Eu is a pure r-process, whereas Y and Ba at [Fe\/H]\u2009~\u2009\u20131\u2009dex are mainly produced by AGB stars, which evolve more slowly than the progenitors of Eu and present an even longer delay with respect to the more massive progenitors of oxygen. At this point it is also interesting to compare our heavy elements with Mg, another \u03b1 element, using the rederived abundances in this work. In Fig. 15 we can see the same decreasing trends towards lower metallicities for Y and Ba but less steep than when comparing oxygen, while [Eu\/Mg] is mostly flat. This might be explained by increasing O\/Mg yields for higher mass SNe progenitors (e.g. Woosley & Weaver 1995; McWilliam et al. 2008, and references therein). Thus, the production of oxygen would start earlier in the Galaxy producing higher [O\/Mg] at lower [Fe\/H]. Moreover, the [Eu\/Mg] is mostly flat suggesting that these two elements receive an important contribution from SNe progenitors of similar masses but less massive than oxygen progenitors, as explained above. However, the study of McWilliam et al. (2008) discarded the possibility of increasing O\/Mg yields for higher mass SNe progenitors since it would imply a metallicity-dependent initial mass function (IMF), that is, an increase in the fraction of low-mass SNe at higher [Fe\/H]. Instead, they proposed that a metallicity-dependent modulation of the SNe O\/Mg ratio can perfectly explain the behaviour of this ratio in the disc and in the bulge. In Fig. 15 it is also interesting to see the well-defined separation of thick disc and h\u03b1mr stars with respect to the thin disc group for [Y\/Mg] and [Ba\/Mg]. ","Citation Text":["Bertran de Lis et al. (2015)"],"Citation Start End":[[736,764]]}
{"Identifier":"2018AandA...619A.152A__Wolff_et_al._(2002)_Instance_1","Paragraph":"Since our review paper 36 years ago (Allard & Kielkopf 1982), considerable progress in unified line broadening theory and in computational technology now enables us to calculate neutral atom spectra given the potential energies and radiative transition moments for relevant states of the radiating atom interacting with other atoms in its environment. Although our unified theory has been developed in Allard et al. (1999), and a detailed discussion is presented there, we provide an overview Sect. 2 for its use in this context. Because the underpinning atomic physics is understood, theoretical models may be drastically improved compared to previous work such as Koester & Wolff (2000) and Wolff et al. (2002) by including a more complete representation of the interactions, especially in the region of atomic separations that determine the line wings. Nevertheless, the main limitation remains a lack of precision in theoretical fundamental atomic interaction data suitable for such spectroscopic calculations. In a field where most work focuses on lower precision for use in chemistry and reaction kinetics, it is fortunate that Mg\u2013He and Mg+\u2013He molecular data are now very well studied and have accurate asymptotic energies available. The potentials we use are described in detail in Sect. 3. In Sect. 4.1 we examine the far wings quantitatively for the MgHe and Mg+He lines, which both contribute to the ultraviolet spectra. We present line profiles obtained over the full range of temperatures from 4000 to 12 000 K for helium densities varying from 1021 to 1022 cm\u22123. In Sect. 4.2 we study the relative contribution of the two resonance lines in their far wings and how they can contribute to the line blanketing. At sufficiently low densities of perturbers, the symmetric center of a spectral line is Lorentzian and can be defined by two line parameters, the width and the shift of the main line. The impact approximation determines the asymptotic behavior of the unified line shape autocorrelation function. The Lorentzian width can be readily extracted, and is presented in Sect. 4.3. Finally, we report the study of WD 2216\u2013657 in Sect. 5.","Citation Text":["Wolff et al. (2002)"],"Citation Start End":[[693,712]]}
{"Identifier":"2020MNRAS.497.2057L__Marco_2016_Instance_1","Paragraph":"The CE has been followed by many analytical (Rasio & Shapiro 1991; Iben & Livio 1993; Han, Podsiadlowski & Eggleton 1994; Terman, Taam & Hernquist 1994; Reg\u0151s & Tout 1995; Armitage & Livio 2000; Sandquist, Taam & Burkert 2000; Papish, Soker & Bukay 2015; MacLeod et al. 2017b), as well as numerical studies (Taam, Bodenheimer & Ostriker 1978; Meyer & Meyer-Hofmeister 1979; Bodenheimer & Taam 1984; Livio & Soker 1988; Rasio & Livio 1996; Sandquist et al. 1998; De Marco et al. 2003; Lombardi et al. 2006; Ricker & Taam 2008; Taam & Ricker 2010; De Marco et al. 2011; Passy et al. 2011, 2012; Ricker & Taam 2012; Nandez, Ivanova & Lombardi 2014; MacLeod & Ramirez-Ruiz 2015a,b; Ivanova & Nandez 2016; Kuruwita, Staff & De Marco 2016; Nandez & Ivanova 2016; Pejcha, Metzger & Tomida 2016; Staff et al. 2016; Ohlmann et al. 2016a,b; Hillel, Schreier & Soker 2017; Iaconi et al. 2017; Murguia-Berthier et al. 2017; Ohlmann et al. 2017; MacLeod et al. 2017a; Chamandy et al. 2018; Grichener, Sabach & Soker 2018; Fragos et al. 2019; Grichener & Soker 2019; Iaconi & De Marco 2019; Iaconi et al. 2019; Prust & Chang 2019; Chamandy et al. 2019a, b; Reichardt et al. 2019, 2020; Glanz & Perets 2020; Prust 2020), nevertheless, many open questions remain unanswered. The dominant processes in the evolution and termination of the CE, for example, are not yet fully understood, and many mechanisms have been proposed [e.g. the \u03b1CE \u2212 \u03bb formalism (van den Heuvel 1976), gamma formalism (Nelemans et al. 2000), recombination energy (Nandez, Ivanova & Lombardi 2015), binding energy (Iaconi et al. 2018), tidal forces (Iben & Livio 1993), internal energy (Han et al. 1994), magnetic fields (Reg\u0151s & Tout 1995), accretion (Voss & Tauris 2003; Soker 2004), nuclear energy (Ivanova & Podsiadlowski 2003) dust-driven winds, Glanz & Perets (2018), and enthalpy (Ivanova & Chaichenets 2011) amongst others], none of which have been able to fully describe the evolution and termination of the CE phase (see for example Soker 2013; Iaconi et al. 2018). The launching of a jet from a compact object (CO) within the CE is also a possible alternative mechanism in the evolution and termination of this phase (Armitage & Livio 2000; Soker 2004; Chevalier 2012; Shiber, Schreier & Soker 2016; Soker & Gilkis 2018; Gilkis, Soker & Kashi 2019; Schreier, Hillel & Soker 2019; Soker, Grichener & Gilkis 2019; Jones 2020).","Citation Text":["Kuruwita, Staff & De Marco 2016"],"Citation Start End":[[701,732]]}
{"Identifier":"2019AandA...621A..27F__DeGraf_et_al._2017_Instance_2","Paragraph":"It is difficult to isolate the impact of mass and environment on the rate and timing of quenching. Mass quenching is more important at earlier times in the evolution of galaxies and may be more important in denser regions (e.g., Peng et al. 2010; Muzzin et al. 2012; Lee et al. 2015; Darvish et al. 2016, 2018; Kawinwanichakij et al. 2017). And in the case of powerful radio galaxies, which lie in over-dense environments, both gas-rich and gas-poor mergers likely play an important role in both the growth of the stellar mass and the black holes. Volonteri et al. (2015b) suggest that in the merger phase, the AGN dominates the bolometric luminosity but the accretion can be very stochastic (see also Gabor & Bournaud 2013; DeGraf et al. 2017). It appears that the galaxies in our sample with the highest star-formation rates all host very powerful AGN, and are potentially all advanced mergers, consistent with this picture. In fact, PKS 0529\u2212549, which has one of the highest SFRs of all the galaxies in our sample, has a modest gas fraction of about 15%, a high star-formation efficiency (SFR\/molecular gas mass), and has been transforming its gas into stars rapidly (Man et al., in prep.). The star formation efficiencies in the other radio galaxies with high SFRs also appear extreme (10\u2013100 Gyr\u22121; Man et al., in prep.). But of course, that does not explain our results in themselves. Dubois et al. (2015), in a study using numerical simulations of the relative growth of SMBHs and their host galaxies, found that star formation may regulate the black hole accretion rate. During the most rapid, gas-rich phase of the growth of massive galaxies, it may be that a larger fraction of the gas in the ISM is not available to fuel the SMBHs, but is consumed via star formation (see DeGraf et al. 2017). As the gas fractions decline, the relative power of the AGN compared to that of the star formation increases, resulting in an increased star formation efficiency. Concomitantly, the increased star formation rate can then disperse the dense gas making it easier for the jets to drive vigorous and efficient outflows (Nesvadba et al. 2006, 2017).","Citation Text":["DeGraf et al. 2017"],"Citation Start End":[[1784,1802]]}
{"Identifier":"2022MNRAS.509.1504M__Timokhin_&_Arons_2013_Instance_1","Paragraph":"Energy flows induced into magnetically dominated relativistic magnetospheres of compact objects are commonly modelled by numerical simulations in the force-free electrodynamics (FFE) limit. Fueled by the track record of observations in the era of multimessenger astrophysics, current targets for such simulations include the magnetospheres of rapidly spinning black holes, spiraling neutron stars, magnetars, and pulsars. The tenuous, magnetically dominated atmosphere (magnetosphere) of pulsars is an active field of scientific interest. They fascinate both observers (e.g. Lorimer et al. 1995; Ransom et al. 2005; Abdo et al. 2013; Jankowski et al. 2018) and theorists (e.g. Kennel & Coroniti 1984; Lyubarskii 1996; Contopoulos, Kazanas & Fendt 1999; Goodwin et al. 2004; Timokhin 2006; Timokhin & Arons 2013; Contopoulos 2019; P\u00e9tri 2020). With the remarkable progress in scientific computing, their rotating magnetosphere has captured designers of numerical methods that integrate FFE and magnetohydrodynamics (MHD) with ever-improving accuracy (e.g. Komissarov 2006; Spitkovsky 2006; Tchekhovskoy, Spitkovsky & Li 2013; Parfrey, Spitkovsky & Beloborodov 2017; Carrasco & Shibata 2020). Recently, particle-in-cell (PIC) simulations were able to resolve a broad range of scale separations and allow for unprecedented insight into the microphysics of pulsar magnetospheres across the global scale (Cerutti et al. 2015; Philippov, Spitkovsky & Cerutti 2015a; Kalapotharakos et al. 2018; Philippov & Spitkovsky 2018; Gu\u00e9pin, Cerutti & Kotera 2020). In this fascinating flurry of outcomes, only few references scrutinized whether the results from ideal plasma simulations are the best possible model for the pulsar magnetosphere that contains an inherently non-ideal region, namely the equatorial current sheet (ECS) beyond the closed zone (Contopoulos 2016; Contopoulos 2019; Contopoulos, P\u00e9tri & Stefanou 2020). Here, we study with rigorous technical depth how this non-ideal region can affect the global dynamics of the force-free aligned rotator magnetosphere \u2013 effectively serving as a blueprint for force-free magnetospheres of other compact objects.","Citation Text":["Timokhin & Arons 2013"],"Citation Start End":[[789,810]]}
{"Identifier":"2015MNRAS.450...53H__Mu\u00f1oz_et_al._2014_Instance_1","Paragraph":"Using moving meshes helps reduce the angular momentum errors from advection in grid codes. We have run >200 iterations of this test problem using the public version of fvmhd3d, systematically varying choices like the mesh regularization scheme, mesh \u2018drifting\u2019 (whether to use a strictly Lagrangian drift, or locally smoothed velocity, or regularized drift), initial mesh geometry, and boundary conditions. In both fvmhd3d and more limited tests with arepo, we find that running in the \u2018simplest\u2019 initial configuration (an initial Cartesian mesh with outflow boundary conditions, with the default mesh regularization scheme used for all other test problems shown here), the disc goes unstable and the angular momentum evolution tends to be corrupted within a few orbits (similar to the fixed-grid cases). Unfortunately, some significant errors in angular momentum evolution are difficult to avoid in moving-mesh codes, as has been discussed extensively in e.g. Duffell & MacFadyen (2012), Ivanova et al. (2013), Mocz et al. (2014) and Mu\u00f1oz et al. (2014). In a shearing disc, if the cells adapt in a truly Lagrangian manner, then they are inevitably deformed into a highly sheared\/irregular shape (Mu\u00f1oz et al. 2014). This leads to other errors. As soon as they become non-spherical (or more accurately fail to be radially symmetric about their own cell centre of mass), then mass advection between cells necessarily leads to additional angular momentum errors (indeed, the angular momentum of an irregular cell cannot be defined exactly but only to the same order of integration accuracy as the local velocity gradient estimator). This is akin to the errors in our MFV method. More importantly, if some regularization procedure is used to keep the cell shapes \u2018regular\u2019 (as is necessary in any moving-mesh code used for this problem), then the regularization means the cells cannot move entirely with the fluid and the gas must be advected over the cells. This re-introduces some of the same (more serious) errors we saw with stationary-grid methods (specifically, see Ivanova et al. 2013, equation 53). This means that the results for moving meshes are quite sensitive to choices like the mesh \u2018stiffness\u2019, regularization procedure, and in particular the choice of boundary conditions for the mesh-generating points (since the rigid Voronoi volume partition can lead to a \u2018mesh tension\u2019 effect, whereby regularization-induced distortions in the central regions propagate outwards \u2018through\u2019 the mesh; Springel 2010). So there are ways to improve the situation on this problem \u2013 for this reason, we do not show a single \u2018standard\u2019 moving-mesh result, because significantly different results are obtained if we make just small changes to the mesh-regularization procedure in each code. However, like with AMR codes, the most effective methods for eliminating angular momentum errors in moving meshes generally depend on knowing the problem geometry ahead of time. For example, Duffell & MacFadyen (2012) design a moving grid which is a series of cylindrical shells free to rotate independently about a shared axis (the disco code); Mu\u00f1oz et al. (2013) use a carefully chosen initial grid configuration with a specially designed boundary condition designed to prevent inward propagation of \u2018mesh deformation\u2019; these help considerably, but must be fine-tuned to the exact disc configuration.","Citation Text":["Mu\u00f1oz et al. (2014)"],"Citation Start End":[[1035,1054]]}
{"Identifier":"2015ApJ...804..130C___2013_Instance_3","Paragraph":"We have developed the simplest spherical void lens model based on the recently developed embedded lens theory. We have assumed a uniform mass profile for the void, compensated by a thin bounding shell. The infinitesimally thin bounding shell was chosen for convenience (Maeda & Sato 1983a, 1983b). To investigate other void profiles such as a non-uniform void interior or a finite-thin bounding ridge (Colberg et al. 2005; Lavaux & Wandelt 2012; Pan et al. 2012; Sutter et al. 2012; Hamaus et al. 2014; Kantowski et al. 2015) is straightforward; one has only to evaluate the Fermat potential of Equation (1) or equivalently the potential part of the time delay of Equation (4). It is also possible to build embedded void lens models with non-spherically symmetric density profiles given that the lowest-order embedded lens theory is applicable to any distributed lens (Kantowski et al. 2013). It is well accepted by the lensing community that small overdensities attract light, whereas small underdensities repel light. This fact can be rigorously proved using general relativistic perturbation theory (Sachs & Wolfe 1967) assuming \n\n\n\n\n\n. However, the repulsive nature of lensing by a large and deep underdense region (i.e., cosmic voids) as described by the rigorously derived but simply implemented embedded lens formalism did not appear until Kantowski et al. (2013). In the case of large density contrasts, i.e., \n\n\n\n\n\n approaching its lower bound \u22121 for cosmic voids, the repulsive lens equation follows naturally from the embedded lensing theory. This theory is based on Swiss cheese models (Einstein & Straus 1945), which are exact solutions of Einstein\u2019s field equations containing inhomogeneities with large density contrasts (Chen et al. 2010, 2011, 2015; Kantowski et al. 2010, 2012, 2013). The void-lensing community takes void repulsive lensing as granted (e.g., Amendola et al. 1999; Das & Spergel 2009), whereas the galaxy\/cluster strong-lensing community has ignored embedding effects, i.e., the repulsive lensing caused by the large underdense regions surrounding the central overdense lens. Besides correctly predicting repulsive lensing by cosmic voids, our Fermat potential formulation can be used to compute the void-lensing time delay effects, including the ISW effect caused by voids; see Equation (5).","Citation Text":["Kantowski et al.","2013"],"Citation Start End":[[1767,1783],[1796,1800]]}
{"Identifier":"2021AandA...646A..62M__Shan_et_al._2014_Instance_1","Paragraph":"So far, these surveys have mostly relied on two-point estimators for their cosmological analysis, for example the shear two-point correlation functions (\u03b3-2PCF, e.g., Kilbinger et al. 2013; Troxel et al. 2018; Hikage et al. 2019; Asgari et al. 2021). These estimators are inherited from cosmic microwave background (CMB) analyses, which probe the matter distribution of the early Universe. However, cosmic shear probes the recent Universe, where the matter distribution is more complex due to the nonlinear accretion of structures that creates non-Gaussian features in the matter field on small scales (e.g., Codis et al. 2015; Barthelemy et al. 2020). Two-point statistics fail to capture this non-Gaussian information and thus yield an incomplete description of the matter distribution at low redshift. To close this gap, the community has recently started to explore non-Gaussian cosmic shear estimators: for example weak-lensing peaks (e.g., Kruse & Schneider 1999, 2000; Dietrich & Hartlap 2010; Kratochvil et al. 2010; Fan et al. 2010; Yang et al. 2011; Maturi et al. 2011; Hamana et al. 2012; Hilbert et al. 2012; Marian et al. 2012, 2013; Shan et al. 2014, 2018; Lin & Kilbinger 2015; Martinet et al. 2015, 2018; Liu et al. 2015a,b; Kacprzak et al. 2016; Petri et al. 2016; Zorrilla Matilla et al. 2016; Giocoli et al. 2018; Peel et al. 2018; Davies et al. 2019; Fong et al. 2019; Li et al. 2019; Weiss et al. 2019; Yuan et al. 2019; Coulton et al. 2020; Ajani et al. 2020; Z\u00fcrcher et al. 2021), Minkowski functionals (e.g., Kratochvil et al. 2012; Petri et al. 2015; Vicinanza et al. 2019; Parroni et al. 2020; Z\u00fcrcher et al. 2021), higher-order moments (e.g., Van Waerbeke et al. 2013; Petri et al. 2015; Peel et al. 2018; Vicinanza et al. 2018; Chang et al. 2018; Gatti et al. 2020), three-point statistics (e.g., Schneider & Lombardi 2003; Takada & Jain 2003, 2004; Semboloni et al. 2011; Fu et al. 2014), density split statistics (e.g., Friedrich et al. 2018; Gruen et al. 2018; Burger et al. 2020), persistent homology (e.g., Heydenreich et al. 2021), scattering transform (e.g., Cheng et al. 2020), and machine learning (e.g., Merten et al. 2019; Ribli et al. 2019; Peel et al. 2019; Shirasaki et al. 2019; Fluri et al. 2019). Although these new statistics have shown a great potential, they have not yet been fully generalized to data analyses, mostly because they need a large number of computationally expensive N-body simulations to calibrate their dependence on cosmology, while this dependence is accurately predicted by theoretical models in the case of two-point estimators.","Citation Text":["Shan et al. 2014"],"Citation Start End":[[1147,1163]]}
{"Identifier":"2019AandA...624A..15S__Petrovich_2015_Instance_2","Paragraph":"As described in Sect. 4.1.1, the TTV inversion code includes a Hill stability criterion for two-planet systems. Although this initial check filters out the least stable systems, it cannot guarantee the long-term stability of the derived orbital architecture of a four-planet system. On one hand, there is no analytical criterion for assessing the long-term stability of multi-planet systems (>2 planets) like that for two-planet systems (Gladman 1993; Chambers et al. 1996). Chambers et al. (1996), Smith & Lissauer (2009), Lissauer et al. (2011b), and similar numerical studies have found that long-term stability of multi-planet systems typically requires the mutual separations between planets to be at least ten mutual Hill radii, which is much larger than the cautious limit of \n\n$2\\sqrt{3}$\n\n\n2\n3\n\n\n\n mutual Hill radii required by the Hill stability criterion that is implemented in our TTV inversion code (Gladman 1993; Chambers et al. 1996). Meanwhile, these criteria are only valid under the assumptions of low mutual inclinations and small eccentricities. These two restrictions are usually not well-defined; limiting values of 1\u00b0\u20132\u00b0 and e = 0.1\u20130.2, respectively, are often adopted (Petrovich 2015; MacDonald et al. 2016). Apparently, these restrictions and stability criteria are well satisfied by the nominal orbital architecture of the Kepler-411 system inverted from the TTV data. On the other hand, the Hill criterion provides no information about the long-term behavior of the system, and repeated weak interactions between planets in Hill stable orbits may still lead to ejections and\/or physical collisions; these are referred to as Lagrange unstable (Petrovich 2015). The chaotic orbits that are generated primarily by first-order (or higher-order) resonance overlap are eventually subjected to large-scale variation of the semi-major axes and hence become Lagrange unstable. Based on previous studies on the chaotic behavior induced by the first-order resonance overlap (e.g. Wisdom 1980; Duncan et al. 1989), Deck et al. 2013 supply the condition that leads to chaotic behavior in a two-planet system (see Eq. (50) in Deck et al. 2013). We apply this criterion to estimate the nominal orbital architecture extracted from the TTV data of the Kepler-411 system, and find that the orbital configuration is far away from the chaotic motion. Therefore, we conclude that the nominal orbital architecture of the Kepler-411 system extracted from measured TTVs satisfies both the Hill and Lagrange stability criteria.","Citation Text":["Petrovich 2015"],"Citation Start End":[[1671,1685]]}
{"Identifier":"2019MNRAS.488.4638L__Drabek-Maunder_et_al._2016_Instance_3","Paragraph":"In Fig. 10, we plot the variation of the ratio of the outflow contribution to the FWHM and turbulent energy. The ratio of the outflow contribution  = 1 \u2013 \u2018non-outflow contribution\u2019\/\u2018all contributions\u2019. We observe that the outflow has a contribution in the FWHM: about 20 per\u2009cent in the local region at the H\u2009ii region (non-outflow contribution is about 81 per\u2009cent) and about 10 per\u2009cent even in the clumps. According to Eturb  = (3\/16 ln\u20092)Mcloud \u00d7 FWHM2, outflow has a contribution in the turbulent energy up to 35 per\u2009cent in the local region at the H\u2009ii region (1 \u2212 0.812). It has a contribution of at least 15 per\u2009cent in the clump at early stages of massive star formation, which is lower than that reported in previous studies (e.g. Bally 2016; Drabek-Maunder et al. 2016). The outflow contribution decreases with time once the outflow action stops. This indicates that the outflows do not have a significant cumulative impact on the turbulent levels during the occurrence of several outflow actions. Thus, the outflow energy contribution to turbulent energy increases insignificantly with the evolutionary stages. Our results suggest that the outflow energy is large enough to maintain the turbulent energy in the clumps and that the outflow has some (not significant) effect on the turbulent energy. However, there is a better correlation between the outflow energy and turbulent energy (see Fig. 5). Therefore, we could not determine if the outflow significantly contributes to the turbulent energy in the clumps. This is consistent with the study conducted by Maud et al. (2015). They also reported that there is a better correlation between the outflow energy and turbulent energy, but the core turbulence is not driven by the local input from the outflows. However, Drabek-Maunder et al. (2016) and Yang et al. (2018) reported that there is no correlation between the turbulent and outflow energies. Urquhart et al. (2018) found that the clump mass and evolutionary stage are uncorrelated. For similar mass of massive star, it is likely that we can observe the obvious difference of turbulent energy between clump without and with outflow. However, for statistics, the mass parameter of turbulent energy is less constrained for each evolutionary stage. All these findings imply that the outflow action has some impact on the local environment and cloud itself, but the contribution from outflow does not mainly drive turbulence. This observation is consistent with several other studies that suggest that turbulence is mostly driven by large-scale mechanisms (Ossenkopf & Mac Low 2002; Brunt, Heyer & Mac Low 2009; Padoan et al. 2009; Arce et al. 2010; Mottram & Brunt 2012; Plunkett et al. 2015; Drabek-Maunder et al. 2016).","Citation Text":["Drabek-Maunder et al. 2016"],"Citation Start End":[[2711,2737]]}
{"Identifier":"2015MNRAS.447.2753T__Katz,_Hernquist_&_Weinberg_1992_Instance_1","Paragraph":"The formation of dark matter haloes has been studied extensively using numerical dark matter only simulations (e.g. Springel et al. 2005b; Fosalba et al. 2008; Boylan-Kolchin et al. 2009; Teyssier et al. 2009; Klypin, Trujillo-Gomez & Primack 2011).1 Extending the insight from dark matter only simulations to include a theory of galaxy formation requires a method to link the formation of dark matter haloes to observable galaxy properties. The most direct method available is hydrodynamical simulations, which model the co-evolution of dark matter and baryons (e.g. Katz, Hernquist & Weinberg 1992; Katz, Weinberg & Hernquist 1996; Weinberg, Hernquist & Katz 1997; Murali et al. 2002; Springel & Hernquist 2003b; Kere\u0161 et al. 2005; Ocvirk, Pichon & Teyssier 2008; Crain et al. 2009; Croft et al. 2009; Oppenheimer et al. 2010; Schaye et al. 2010; Vogelsberger et al. 2012). By directly including hydrodynamics in structure formation simulations, one can probe the thermal state and column density distribution of the intergalactic medium (IGM) via the Lyman \u03b1 forest (Cen et al. 1994; Zhang, Anninos & Norman 1995; Hernquist et al. 1996; Theuns et al. 1998), the phase structure and heavy element composition of the circumgalactic medium and IGM (Aguirre et al. 2001; Cen & Fang 2006; van de Voort & Schaye 2012; Shen et al. 2013), and the nature and rates of accretion of dark matter and baryons into galaxies (Kere\u0161 et al. 2005; van de Voort et al. 2011; Nelson et al. 2013). Since these simulations can follow the dynamics of both the dark matter and baryons down to small spatial scales, predictions can be made about their internal structure including the distribution of gas (Kere\u0161 et al. 2012; Torrey et al. 2012) and the formation of stellar discs and bulges (Abadi et al. 2003; Governato et al. 2004; Agertz, Teyssier & Moore 2011; Sales et al. 2012; Marinacci, Pakmor & Springel 2014). As a result of this detailed information, simulations are a valuable tool for interpreting observational data and placing observed galaxies into a more complete evolution based cosmological context.","Citation Text":["Katz, Hernquist & Weinberg 1992"],"Citation Start End":[[568,599]]}
{"Identifier":"2021MNRAS.504..146V__Vink_&_Gr\u00e4fener_2012_Instance_2","Paragraph":"The direct detection of the first gravitational waves from the merger of two heavy black holes (BHs) in GW\u2009150914 confirmed one of the toughest predictions of Einstein\u2019s theory of general relativity. But while satisfying the world of physics in general, for astrophysics this was only the beginning: many were surprised by the large BH masses of, respectively, 36 and 29\u2009\u2009M\u2299 (Abbott et al. 2016), showcasing how the new field of multimessenger astrophysics had just re-opened the field of stellar evolution in a spectacular fashion. Stellar mass BHs had previously been revealed by their interaction in binary systems (Orosz et al. 2011), but the maximum stellar BH mass in our Milky Way is not higher than roughly 15\u201320\u2009\u2009M\u2299 (Belczynski et al. 2010). While we know that very massive stars (VMS) above 100\u2009\u2009M\u2299 exist (Crowther et al. 2010; Vink et al. 2015), this mass is significantly diminished via stellar winds already during core hydrogen (H) burning (Vink & Gr\u00e4fener 2012). The heavy nature of the BH, as measured by LIGO\/VIRGO therefore supported the assumption that the gravitational wave event occurred in a part of the Universe still pristine in its enrichment with heavy elements (\u2018metallicity (Z)\u2019), lowering stellar wind mass-loss (Vink, de Koter & Lamers 2001; Vink & de Koter 2005). A low-Z solution was widely accepted until the announcement of a 70\u2009\u2009M\u2299 BH in LB-1 (Liu et al. 2019), spurring stellar evolution theorists to avoid heavy mass-loss in the Milky Way (Belczynski et al. 2020; Groh et al. 2020), either by arbitrarily lowering the mass-loss rates of VMS \u2013 seemingly contradicting VMS mass-loss calibrations (Vink & Gr\u00e4fener 2012) \u2013 or by invoking the presence of a strong dipolar surface magnetic field that could quench the wind (Petit et al. 2017). While such magnetic fields in some 5\u201310 per\u2009cent of massive OB stars do indeed exist, no B-fields have yet been detected in VMS (Bagnulo et al. 2020). The problem of the formation of a $70\\, \\mathrm{ M}_\\odot$ BH in a solar metallicity environment apparently resolved itself when the spectral signatures of LB-1 were re-interpreted (Abdul-Masih et al. 2020; El-Badry & Quataert 2020).","Citation Text":["Vink & Gr\u00e4fener 2012"],"Citation Start End":[[1633,1653]]}
{"Identifier":"2016ApJ...817....9K__Chen_et_al._2006_Instance_1","Paragraph":"To perform a statistical analysis of the average quiescent fraction of satellites around our sample of massive galaxies, we use a statistical background subtraction technique (e.g., Kauffmann et al. 2010; Tal et al. 2012; Wang & White 2012; Kawinwanichakij et al. 2014). We detect objects within fixed apertures centered on our central galaxies and satisfying Equation (2). These apertures include physically associated galaxies as well as chance alignments of foreground and background galaxies. We estimate and correct for the contamination due to chance alignments by placing random apertures across the field. We adapt this procedure by restricting the placement of the random apertures to regions near the centrals, as demonstrated by Chen et al. (2006). This accounts for the bias due to contaminating galaxies that are physically associated with the centrals but are not satellites (i.e., the 2-halo term of the correlation function; see Chen et al. 2006).15\n\n15\nThe contaminating galaxies that are physically associated with the central galaxies in our sample are expected to have marginally different properties than truly random field galaxies due to the fact that they exist in biased regions of the universe. There may be an additional effect due to large-scale 2-halo conformity. If 2-halo conformity exists, our procedure effectively corrects for it.\n We therefore place the random apertures within annuli with inner and outer radii equal to 1 and 3 cMpc from each central galaxy for the UDS and UltraVISTA. Parenthetically, our tests showed that the restriction on the location of the background apertures has only a small effect on the conformity signal. Relative to apertures that are placed randomly through the field, this correction increases the quiescent fractions of background galaxies by 0.4%\u201310%. For the smaller ZFOURGE fields, placing the random apertures within annuli is too restrictive, and for this survey we randomly place the apertures across the fields. We do note that the ZFOURGE fields are small enough that even these randomly placed apertures trace the same large-scale environment as the centrals. Additionally, we find that when we restrict the background apertures to be \n\n\n\n\n\n cMpc from the centrals, it changes the measured quenching efficiencies (see Section 4 below) by 10%, and none of our conclusions would be changed.","Citation Text":["Chen et al. (2006)"],"Citation Start End":[[740,758]]}
{"Identifier":"2022MNRAS.516.3381J__Jones_2001_Instance_1","Paragraph":"Studying the dynamical properties of rotating neutron stars is a domain which brings out various interesting features when one assumes a perfect fluid. It is known that the centrifugal force of a rotating star counters gravitational force and hence one can expect massive stars to be fast rotors, at least in the initial stages of the stellar evolution. As a result of rotation a star may experience damping due to unstable oscillations such as the r-modes. The r-modes are one of many pulsating modes that exist in neutron stars and are characterized by the Coriolis force acting as the restoring force (Andersson 1998). The r-modes are unstable to emission of gravitational radiation (GR) by the Chandrashekhar-Friedman-Schutz (CFS) mechanism (Chandrasekhar 1970; Friedman & Schutz 1978). It was shown in Andersson (1998) that the r-modes are unstable for all rotating perfect fluid stars irrespective of their frequency. Dissipative effects such as shear and bulk viscosities work towards suppressing GR driven instabilities and has been studied by various authors over the past few years (Lindblom, Owen & Morsink 1998; Jones 2001; Lindblom & Owen 2002; van Dalen & Dieperink 2004; Drago, Lavagno & Pagliara 2005; Nayyar & Owen 2006; Jaikumar, Rupak & Steiner 2008; Jha, Mishra & Sreekanth 2010; Ofengeim et al. 2019) under various considerations. If the GR time-scale is shorter than the damping time-scale due to such dissipative processes, then the r-mode will be unstable and a rapidly rotating neutron star could lose a significant fraction of its rotational energy through GR. At higher temperatures (T > 109 K), the dominant dissipation is due to bulk viscosity, which arises due to density and pressure perturbations, a consequence of the star being driven out of equilibrium by oscillations. The system tries to restore equilibrium through various internal processes. In the case of r-modes, since the typical frequencies are of the order of the rotational frequencies of the stars, the reactions that dominate are the weak processes. Within these weak processes, although the modified Urca processes involving leptons are important, it has been shown that non-leptonic processes involving hyperons contribute more significantly towards bulk viscosity at temperatures lower than a few times 109 K (Lindblom & Owen 2002). Our goal here is to investigate the same using a chiral model calibrated to reproduce the desired nuclear matter properties, in particular the density content of the nuclear symmetry energy at both low and high densities.","Citation Text":["Jones 2001"],"Citation Start End":[[1124,1134]]}
{"Identifier":"2017AandA...604A..80M__Propris_et_al._(2013)_Instance_1","Paragraph":" Using the whole sample (\\hbox{$\\bar{z}=0.40$}z\u0305 = 0.40), we find a decreasing faint end for both datasets with consistent values between HST (\u03b1 =  \u2212 0.76 \u00b1 0.07) and Subaru (\u03b1 =  \u2212 0.78 \u00b1 0.06). Separating between low-redshift (\\hbox{$\\bar{z}=0.29$}z\u0305 = 0.29) and high-redshift (\\hbox{$\\bar{z}=0.51$}z\u0305 = 0.51) samples, we find an evolution of the faint end slope of 1.7\u03c3 with HST and 2.6\u03c3 with Subaru. There is thus a mild decrease of the faint end slope (less negative \u03b1) with increasing redshift over the range (0.187 z 0.686). This evolution is in good agreement with recent papers in the literature: in particular Zenteno et al. (2016) found a decrease of the RS faint end at 2.1\u03c3 for a wider range of redshifts (0.1 z 1.13), but with ~ 80% of their clusters being in the same redshift range as ours. De Propris et al. (2013) claim that the evolution in the faint end slope has a significant contribution from surface brightness selection effects. They used HST data of differing depths on a single cluster (MS 1358.4+6254) to show that surface brightness selection effects become important above the formal magnitude limit of their data and that they affect the RS GLF at magnitudes z \u2265 24.5 for 2.7 ks HST exposures (see their Fig. 18). The faint RS for their cluster has F814W \u2212 z = 0.25, implying that the SB selection effects in their sample become important at F814W> 24.75. On the other hand, our CLASH data are significantly deeper than theirs (4.1 ks) and we limit our GLFs at F814W 24.5. Therefore, the real SB selection effects noticed in De Propris et al. (2013) should not be playing a role in our space-based results. In addition, De Propris et al. (2013) claim that previous estimates of the evolution in the RS GLF (e.g.,  De Lucia et al. 2007; Rudnick et al. 2009) were also due to SB effects. Both of those works were based on the same ground-based data with a formal magnitude limit of I = 24 or 24.5 (for the low- and high-redshift clusters, respectively) and the evolution in the GLF was seen over the faintest 2 mag. We cannot directly address the role of SB effects in the EDisCS results without detailed simulations on those data (see below for such simulations for our clusters) but the similarity between our HST and Subaru GLFs imply that the EDisCS evolution in the GLF is not dominated by SB effects. ","Citation Text":["De Propris et al. (2013)"],"Citation Start End":[[807,831]]}
{"Identifier":"2021ApJ...923..169R__Rhea_et_al._2020b_Instance_1","Paragraph":"To measure the emission lines\u2019 flux within a spectrum, it is essential to fit the lines using the proper model with parameters that allow us to represent the intensities in all spectral elements. Estimating the number of components of each underlying emission line (i.e., multiple lines with multiple velocity components as would be expected from merging galaxy systems) is also crucial to extract meaningful information from the lines. Standard methodologies require fitting both single and double component models and computing their Bayes factors, or approximating it using a proxy such as the Akaike (or Bayesian) information criterion (AIC; Rhea et al. 2020b). However, these methods are highly reliant on the accuracy of the fits (e.g., Kiesepp\u00e4 1997; Pooley & Marion 2018). Several promising new methods have been proposed utilizing machine-learning algorithms to solve this problem (Hampton et al. 2017; Keown et al. 2019; Rhea et al. 2020a). In particular, CLOVER, developed in Keown et al. (2019), uses a convolutional neural network (CNN) to classify high-resolution radio emission lines as having either single or double underlying components. In this paper, we expand upon this methodology for medium resolution, Fourier Transform IFU spectra taken by the Spectrom\u00e8tre Imageur \u00e0 Transform\u00e9e de Fourier pour l\u2019Etude en Long et en Large de raies d\u2019Emission (SITELLE) instrument at the Canada\u2013France\u2013Hawai\u2019i Telescope (Baril et al. 2016). Each SITELLE observation contains approximately 4 million pixels and, thus, produces 4 million spectra of a given resolution set by the principal investigator (1  R  10, 000; Drissen et al. 2014; Martin & Drissen 2017; Drissen et al. 2019). The instrumental line function of SITELLE is described as a sinc model convolved with a Gaussian to represent intrinsic line broadening, which requires special care during the fitting process (Martin et al. 2016). Therefore, the development of machine-learning applications for SITELLE demands special treatment of the underlying emission profiles.","Citation Text":["Rhea et al. 2020b"],"Citation Start End":[[646,663]]}
{"Identifier":"2017ApJ...837...88B__Ribeiro_et_al._2010_Instance_1","Paragraph":"Similarly, there is general agreement about the astrophysical effects responsible for velocity segregation between more and less luminous galaxies. This effect can be explained by physical processes such as dynamical friction (also called dynamical breaking) that cause transfers of kinetic energy from larger galaxies to smaller galaxies, as well as gravitational interactions that can convert bulk kinetic energy into internal kinetic energy via, for example, accelerating and ejecting individual stars during galaxy mergers (Sarazin 1986; Kashlinsky 1987; Biviano et al. 1992; Mahajan et al. 2011). It is mildly puzzling that the velocity segregation effect that we see in our data for the brightest and faintest galaxies does not appear in some recent simulations; we are not the first to detect this signal in bright cluster member galaxies, (Chincarini & Rood 1977; Biviano et al. 1992; Mohr et al. 1996; Goto 2005; Ribeiro et al. 2010; Old et al. 2013; Barsanti et al. 2016), and the effect is seen in some simulations (e.g., Lau et al. 2010; Saro et al. 2013). We are wary of over-interpreting this discrepancy, because of the difference in our bright galaxy sample and the 50 brightest members approach used by Gifford et al. (2013), but it is possible that the discrepancy is related to the effects of missing baryonic\/gas physics that are not captured in dark matter only simulations. We elect not to speculate too much about the pros and cons of various methods for the treatment of halos in simulated clusters, and how mock galaxies are placed into those halos, though it is not unreasonable to suppose that different prescriptions could generate significantly different results. Certainly the measurement we present here could represent a useful benchmark test to be reproduced in future simulations. Looking forward, we believe it is important that both observers and simulators strive to find useful quantities that can be measured in both data and simulations; it is with this in mind that we analyze galaxies in our sample in terms of relative luminosity, scaled by the characteristic magnitude, m*.","Citation Text":["Ribeiro et al. 2010"],"Citation Start End":[[922,941]]}
{"Identifier":"2015ApJ...809L..20M__Kraft_1967_Instance_1","Paragraph":"Several models have been advanced to explain the obliquity dichotomy, but so far none can account for all the relevant observations. Interpretations based solely on differences in the stellar properties (e.g., Rogers et al. 2012) cannot address the inferred dependence of the obliquity on Mp and the detection of highly misaligned planets around cool stars. In view of the fact that \n\n\n\n\n\n corresponds to the temperature above which the size of the outer convective zone of a main-sequence (MS) F star shrinks rapidly, it was suggested that close-in planets in both cool and hot stars are initially distributed over the entire angular range (0\u00b0\u2013180\u00b0), but that only in cool stars where a substantial convective envelope is present can a sufficiently massive close-in planet (\n\n\n\n\n\n) realign the star through tidal interaction (e.g., Winn et al. 2010; Albrecht et al. 2012). The host stars are also subject to magnetic braking, which declines strongly above the same transition temperature (corresponding to the break in the Kraft curve; Kraft 1967), and it was argued (Dawson 2014) that this, rather than the tidal dissipation efficiency, is the main factor underlying the difference in obliquity properties between cool and hot stars. Both variants of this scenario, however, face the conundrum that an HJ undergoing equilibrium tidal interaction with its host star would spiral in and be ingested on a timescale that is comparable to the realignment time. A possible way out of this difficulty is to appeal to the {10} component of the dynamical tide, which could in principle significantly reduce the alignment time without affecting the ingestion time (Lai 2012).1\n\n1\nAn alternative possibility, that only the outer convective layer (Winn et al. 2010)\u2014or even just a part of it (Dawson 2014)\u2014partakes in the realignment process, is hard to justify on either theoretical or observational grounds.\n However, even though this model can be used to account for individual systems (Valsecchi & Rasio 2014), it remains unclear whether it can explain the overall \u03bb distribution of HJs and the manifested difference between cool and hot stars (Rogers & Lin 2013; Xue et al. 2014). Furthermore, even the basic tidal interaction interpretation of the obliquity dichotomy has now been called into question by the results of Mazeh et al. (2015), who analyzed the rotational photometric modulations of a large sample of Kepler sources and inferred that (1) the conclusion that planets around cool stars are well aligned, and those around hot stars are not, is general and not restricted just to HJs; and (2) the low obliquity of planets around cool stars extends to orbital periods Porb that are a factor of \u223c10 larger than the maximum value (\n\n\n\n\n\n days) for robust tidal interaction between an HJ and a \n\n\n\n\n\n Gyr old G or F star.","Citation Text":["Kraft 1967"],"Citation Start End":[[1037,1047]]}
{"Identifier":"2022ApJ...926...85S__Ehrenreich_et_al._2020_Instance_3","Paragraph":"As in Flowers et al. (2019), to compare our model transmission spectra directly against the Ehrenreich et al. (2020) results, we must calculate the transmission spectra as a function of orbital phase throughout the duration of transit. To account for orbital phase dependencies, we apply the following procedure:1.Account for phase-dependent backlighting of the planet (i.e., stellar limb-darkening effects). At different points of its transit, the planet will occult regions of its host star of varying brightness. Furthermore, at a fixed orbital phase, different regions of the planet\u2019s limb will be backlit by varying intensities of stellar light. Similar to Flowers et al. (2019), we calculate the normalized stellar intensity at the center of each cell of the 2D projected planetary grid produced by our GCM at each modeled orbital phase of the planet. We use the quadratic limb-darkening coefficients reported by Ehrenreich et al. (2020) to establish the stellar center-to-limb intensity profile, and we take into account the 89.\u00b0623 orbital inclination of WASP-76b (Ehrenreich et al. 2020) to determine where the planet resides on the stellar disk as a function of its orbital phase. We make the assumption of constant impact parameter b over the course of transit.\n9\n\n\n9\nIn reality, a planet on an inclined orbit will not have a constant b over the entire duration of transit; rather, the planet\u2019s distance from the stellar equator will be decreased at ingress and egress, reaching its maximum at center of transit. Our tests reveal that, for WASP-76b, the relative error induced by the constant b assumption is on the order of 4% in distance, which results in a change on the order of 1 m s\u22121 at the blueshift level (see Section 3.1). Hence, our b treatment is justified. This procedure allows us to calculate a backlighting factor f, which ranges from 0 to 1, effectively replacing the constant I\n\n\u03bb,0 from Equation (3) with a variable \n\n\n\nI\u03bb,0\u00d7f(\u03b8\u2032,z,\u03c6)\n\n, for a given orbital phase \u03c6 and 2D projected polar angle \n\n\n\n\u03b8\u2032\n\n.2.Account for the decreasing of the continuum by interpolating a light curve produced by the batman code (Kreidberg 2015). Step 1 ensures that less light is transmitted through the planet\u2019s atmosphere than a uniform stellar disk would emit. Step 2 further enforces that the inner, optically thick core of the planet is simulated crossing a limb-darkened star, as opposed to a star of uniform brightness.3.Account for the planet\u2019s rotation over the course of transit. Because the planet is continually rotating as it travels across the face of its host star, we must transform the GCM coordinate system so that the correct observer-facing hemisphere is modeled at each instance during transit. For simplicity, we assume zero obliquity, which allows us to calculate the coordinate transform simply by assigning a linear offset to each planetary longitude; i.e., \u03d5\nrotated = \u03d5 + \u03c6.\n","Citation Text":["Ehrenreich et al. 2020"],"Citation Start End":[[1073,1095]]}
{"Identifier":"2020AandA...639A..20K__Johnston_et_al._2019_Instance_1","Paragraph":"The above analysis addresses a number of open questions. Firstly, it is necessary to ask whether impulsive heating events, such as Ohmic dissipation associated with magnetic field braiding can trigger coronal rain formation in the first place. The evolution of loop L1 shows that this is indeed possible, as the duration of the impulsive heating events occuring in the loop is 100\u2212200 s, similar to the loop radiative cooling timescale of \u223c100 s. An impulsive heating event of significant amplitude and with duration that is comparable to the radiative cooling timescale is sufficient to cause sufficient chromospheric evaporation into the loop for the loop to become thermally unstable (Johnston et al. 2019). Previous studies that addressed prominence formation with using a time-variable coronal heating term found that formation of condensations still occurs if 1. the heating is intermittent, but the frequency of the individual heating events exceeds the inverse of the radiative cooling timescale (Karpen & Antiochos 2008) and 2. if the heating ceases during the condensation formation, the condensation still continues to grow (Xia et al. 2011). It should also be noted that such impulsive events can also act as a perturbation that trigger the catastrophic cooling in a marginally stable plasma. We noted differences between condensations formed in loops L1 and L2, in particular that the heating driving the condensation formation in L1 has an impulsive character compared to the gradual heating of L2. Also, the cooling timescale of the condensation formed in L1 is much shorter than that of the condensation formed along L2, suggesting different conditions at coronal heights. Secondly, we also addressed the issue of the apparent coronal rain formation along open field lines and shown that flows of evaporated plasma can lead to sufficient density increase to trigger radiative instability in both closed and open magnetic field configurations. Finally, our work also addresses a possible mechanism responsible for destruction of coronal rain condensations before they reach the solar surface. A reconnection event occurring in the upper part of the leg of loop 3 leads to the condensation being heated to very high temperatures and subsequently destroyed as a result. This suggests that the lifetime of the condensations can be dependent on the frequency of such reconnection events in shorter loops. Also, it might possibly explain why coronal rain is not observed in short low-lying loops close to the active region core that are heated to very high temperatures. Even if the condensations do form in such loops (which is significantly less likely as it is easier to trigger thermal instability in longer loops; see e.g. M\u00fcller et al. 2004, 2005), they are most likely short-lived.","Citation Text":["Johnston et al. 2019"],"Citation Start End":[[688,708]]}
{"Identifier":"2020ApJ...899L...1Z__Santoliquido_et_al._2020_Instance_1","Paragraph":"This work focuses on the formation of systems with high mass ratios and component masses in the LMG through canonical isolated binary evolution. Many other channels have been proposed for producing the compact binary mergers observed by LIGO\u2013Virgo. Dynamical formation in dense stellar clusters such as globular clusters preferentially produces compact binaries with similar masses (e.g., Sigurdsson & Hernquist 1993), and thus the formation of NSBHs and other compact binaries with highly asymmetric masses is predicted to be rare (Clausen et al. 2013; Arca Sedda et al. 2020; Ye et al. 2020). While hierarchical mergers of NSs have been proposed as a means of populating the LMG (Gupta et al. 2020), this scenario is unlikely since heavier BHs dominate the dynamical interactions in clusters (e.g., Samsing & Hotokezaka 2020; Ye et al. 2020). The formation of compact binaries with highly asymmetric masses may be more prevalent in young star clusters (e.g., di Carlo et al. 2019; Rastello et al. 2020; Santoliquido et al. 2020), but Fragione & Banerjee (2020) finds the merger rate of NSBH systems in young massive and open clusters to be three orders of magnitude lower, similar to the predictions from old globular clusters. Other formation mechanisms have been explored for forming highly asymmetric compact binary mergers and mergers with components in the LMG, such as hierarchical systems in the galactic field (e.g., Antonini et al. 2017; Silsbee & Tremaine 2017; Fragione & Kocsis 2019; Safarzadeh et al. 2019; Fragione et al. 2020), hierarchical systems in galactic nuclei with a supermassive BH as the outer perturber (e.g., Antonini & Perets 2012; Petrovich & Antonini 2017; Hoang et al. 2018; Fragione et al. 2019; Stephan et al. 2019), and in disks around supermassive BHs in active galactic nuclei (e.g., McKernan et al. 2019; Yang et al. 2019). However, the rates and formation properties from these channels are uncertain. Nevertheless, a full picture of compact-binary mergers will require consideration of all these channels and investigation of how physical prescriptions (such as the connection between the underlying SN mechanism and remnant mass) jointly affect population properties, rates, and branching ratios across these channels (e.g., Stevenson et al. 2017; Talbot & Thrane 2017; Vitale et al. 2017; Zevin et al. 2017; Arca Sedda et al. 2020). The identification of bona fide NSBH systems and other compact-binary mergers with highly asymmetric masses will further constrain the relative contribution of various formation channels and the underlying physics of these formation pathways.","Citation Text":["Santoliquido et al. 2020"],"Citation Start End":[[1005,1029]]}
{"Identifier":"2022ApJ...930..160S__Tremonti_et_al._2004_Instance_1","Paragraph":"Recent years have seen a burgeoning in the volume of spatially resolved spectroscopic data from large surveys such as CALIFA (S\u00e1nchez et al. 2012), SAMI (Croom et al. 2012), and MaNGA (Bundy et al. 2015). These data have stimulated interest in the relationship between the gas-phase oxygen abundances, stellar mass, star formation rate (SFR), and gas content of galaxies on kiloparsec scales (e.g., Barrera-Ballesteros et al. 2016, 2018; Mingozzi et al. 2020; Teklu et al. 2020; Wang & Lilly 2021). By analogy to the global mass\u2013metallicity relation (Tremonti et al. 2004), the oxygen abundance has been shown to be sensitive to the local stellar mass surface density (\u03a3*), tracing the integrated star formation history on local scales in a galaxy, as well as the presence of outflows (e.g., Barrera-Ballesteros et al. 2018) and inflows (Lian et al. 2019; Schaefer et al. 2019). Since early observations of this local relationship between \u03a3* and 12+log(O\/H) (Moran et al. 2012; Rosales-Ortega et al. 2012), there has been a growing consensus that the global mass\u2013metallicity relation can be explained as arising on local scales. That is, the accumulation of chemical elements on kiloparsec scales can be seen as a reflection of the buildup of stellar mass locally within galaxies over their evolutionary history rather than on global scales (see, e.g., S\u00e1nchez 2020; S\u00e1nchez et al. 2021). This is reinforced by recent observations that gas-phase chemical abundances in galaxies are correlated on \u223ckiloparsec scales (S\u00e1nchez et al. 2015; Kreckel et al. 2020; Li et al. 2021), implying that their chemical enrichment proceeds by the local injection and diffusion of metals into the interstellar medium (ISM) of galaxies. Nevertheless, some observations have shown that the local gas-phase metallicity is also related to the total stellar mass of galaxies (Gao et al. 2018). This is likely due to the greater depth of the gravitational potential well of more massive galaxies, which makes the expulsion of metals through feedback-driven outflows more difficult.","Citation Text":["Tremonti et al. 2004"],"Citation Start End":[[551,571]]}
{"Identifier":"2020MNRAS.492.3509B__Wiaux,_Puy_&_Vandergheynst_2010_Instance_1","Paragraph":"It is to be mentioned here that most of the imaging techniques in RI literature work under the assumption of absence of any calibration errors or availability of pre-calibrated data. In the more practical scenario of working in the presence of DDE calibration errors, these terms need to be accounted for within the measurement model, such as done in Bhatnagar et al. (2013) in the case when these terms are known. Particularly for joint full polarization imaging, we can generalize the model proposed in Birdi et al. (2018b) and employ it to solve inverse problem (7). The corresponding imaging strategy leverages the CS theory. The main idea is to obtain an estimation of the sought images by solving a minimization problem of the form (Birdi et al. 2018a, b)\n(9)$$\\begin{eqnarray}\r\n\\underset{^{{{\\boldsymbol{\\sf {S}}}}}}{\\mathrm{minimize}}\\,\\,\\,\\,r ({\\boldsymbol{\\sf {S}}}) \\, \\, \\text{subject to} \\, \\, \\Vert \\Phi ({\\boldsymbol{\\sf {S}}}) - \\boldsymbol {y} \\Vert _2 \\leqslant \\epsilon ,\r\n\\end{eqnarray}$$where \u03f5 is related to the norm of the additive noise, \u2016 \u00b7 \u20162 denotes the \u21132 norm of its argument vector. The term $r({\\boldsymbol{\\sf {S}}})$ is the regularization term, injecting prior information in the reconstruction process, whereas the data fidelity is ensured by the constraint $\\Vert \\Phi ({\\boldsymbol{\\sf {S}}}) - \\boldsymbol {y} \\Vert _2 \\leqslant \\epsilon$, implying that the residual is situated within an \u21132 ball whose radius is determined by \u03f5. In the current work, due to technical assumptions related to the proposed joint imaging and calibration algorithm (Chouzenoux et al. 2016), we propose to solve an unconstrained version of problem (9), given by\n(10)$$\\begin{eqnarray}\r\n\\underset{^{{{\\boldsymbol{\\sf {S}}}}}}{\\mathrm{minimize}}\\,\\,\\,\\,\\bar{h} \\big ({\\boldsymbol{\\sf {S}}} \\big) + r({\\boldsymbol{\\sf {S}}}),\r\n\\end{eqnarray}$$where $\\bar{h}({\\boldsymbol{\\sf {S}}})$ is the data fidelity term. Using this formulation, and under the assumption that the additive noise is i.i.d. Gaussian, $\\bar{h}({\\boldsymbol{\\sf {S}}})$ is given by a least-squares criterion\n(11)$$\\begin{eqnarray}\r\n(\\forall {\\boldsymbol{\\sf {S}}} \\in {\\mathbb {R}} ^{2 \\times 2 N}) \\, \\, \\bar{h}({\\boldsymbol{\\sf {S}}}) = \\frac{1}{2}\\Vert \\Phi ({\\boldsymbol{\\sf {S}}}) - \\boldsymbol {y} \\Vert _2^2.\r\n\\end{eqnarray}$$This formulation has been introduced in Repetti et al. (2017) in the particular case of joint Stokes I imaging and calibration. In order to incorporate various polarimetric imaging specific prior informations in the regularization function, Birdi et al. (2018b) proposed to use a hybrid regularization term of the form\n(12)$$\\begin{eqnarray}\r\nr({\\boldsymbol{\\sf {S}}}) = g({\\boldsymbol{\\sf {S}}}) + r^\\prime ({\\boldsymbol{\\sf {S}}}),\r\n\\end{eqnarray}$$where the function $g:{\\mathbb {R}} ^{2 \\times 2N} \\rightarrow ]-\\infty , +\\infty ]$ imposes sparsity of the underlying Stokes images and the function $r^\\prime :{\\mathbb {R}} ^{2 \\times 2N} \\rightarrow ]-\\infty , +\\infty ]$ constrains the domains of the argument images as per some physical constraints. In particular, the first term is inspired by the CS framework which proposes to exploit the sparsity of the target images in some dictionary ${\\boldsymbol{\\Psi }} \\in {\\mathbb {C}} ^{N \\times J}$ (Cand\u00e8s et al. 2006; Donoho 2006). The choice of this dictionary is dependent on the images under scrutiny. To give an intuitive idea, images which are already sparse, for instance point sources images, are well represented by dirac basis taking ${\\boldsymbol{\\Psi }}$ to be identity. Otherwise, sparsity can be imposed in some other domain, such as TV regularization for piece-wise constant images (Rudin et al. 1992; Wiaux, Puy & Vandergheynst 2010; Akiyama et al. 2017), wavelet basis for smooth images (Mallat 2009), etc. In particular for astronomical images, a collection of wavelet bases is shown to be a good candidate for sparsifying dictionary ${\\boldsymbol{\\Psi }}$ (Carrillo et al. 2012, 2014), both for Stokes I imaging (Onose et al. 2016, 2017) and specifically for polarization imaging in Birdi et al. (2018a, b). Formally, for any matrix $\\widetilde{{\\boldsymbol{{{\\sf S}}}}} \\in {\\mathbb {R}} ^{N \\times 4}$, the sparse representation of the sought images in a chosen dictionary ${\\boldsymbol{\\Psi }}$ is given by ${\\boldsymbol{\\Psi }} ^\\dagger \\widetilde{{\\boldsymbol{{{\\sf S}}}}}$. The regularization function is then defined to impose sparsity of this term. This can be achieved by using the \u21130 pseudo-norm, counting the number of non-zero components of its argument (Donoho et al. 1995). A more common approach is to use \u21131 norm which overcomes the problem of non-convexity of the \u21130 norm (Chen, Donoho & Saunders 2001). An even better approximation of the \u21130 norm is provided by the (re)weighted \u21131 norm, which tends to diminish the magnitude dependence of the \u21131 norm (Cand\u00e8s et al. 2008; Carrillo et al. 2012). The (re)weighted \u21131 norm is given by\n(13)$$\\begin{eqnarray}\r\ng({\\boldsymbol{\\sf {S}}}) = \\Vert {\\boldsymbol{\\Psi }} ^\\dagger \\mathcal {R} \\big ({\\boldsymbol{\\sf {S}}} \\big) \\Vert _{{\\boldsymbol{\\sf {W}}},1} = \\displaystyle \\sum _{i=1}^4 \\sum _{j=1}^J {\\mathsf {W}} _{j,i} \\big |[{\\boldsymbol{\\Psi }} ^\\dagger \\mathcal {R} \\big ({\\boldsymbol{\\sf {S}}} \\big) ]_{j,i} \\big | \\, ,\r\n\\end{eqnarray}$$where the subscripts j and i in the notation [ \u00b7 ]j,i stand respectively for the row and column indices of the argument matrix, and $\\mathcal {R}:{\\mathbb {R}} ^{2 \\times 2 N} \\rightarrow {\\mathbb {R}} ^{N \\times 4}$ is the operator consisting in placing the four Stokes images contained in the matrix ${\\boldsymbol{\\sf {S}}}$ in four columns, whereas its adjoint $\\mathcal {R}^\\dagger :{\\mathbb {R}} ^{N \\times 4} \\rightarrow {\\mathbb {R}} ^{2 \\times 2 N}$ does the contrary, i.e. stores the four images in the rows of a 2 \u00d7 2 block matrix. In equation (13), ${\\boldsymbol{\\sf {W}}} \\in {\\mathbb {R}} _+^{J \\times 4}$ is the weighting matrix. In the case when this matrix is chosen to be identity, the \u21131 regularization term is obtained. Another useful piece of information, which can be used to regularize the problem is the polarization constraint on the recovered Stokes images (Birdi et al. 2018b). This constraint comes from a physical point of view, that the polarized intensity should not exceed the total intensity and mathematically, can be described by the following set:\n(14)$$\\begin{eqnarray}\r\n\\mathbb {P} = \\big \\lbrace \\widetilde{{\\boldsymbol{{{\\sf S}}}}} = \\mathcal {R}\\big ({\\boldsymbol{\\sf {S}}} \\big) \\in {\\mathbb {R}} ^{N \\times 4} \\, \\big | \\, (\\forall n \\in \\lbrace 1,\\!\\!\\!&amp;&amp;\\ldots ,N\\rbrace) \\nonumber \\\\\r\n&amp;&amp; - \\, \\widetilde{{\\mathsf {S}}}_{n,1} + \\Vert \\widetilde{{\\boldsymbol{{{\\sf S}}}}}_{n, 2:4} \\Vert _2 \\leqslant 0 \\big \\rbrace ,\r\n\\end{eqnarray}$$where, for every n \u2208 {1, \u2026, N}, $\\widetilde{{\\mathsf {S}}}_{n,1}$ and $\\Vert \\widetilde{{\\boldsymbol{{{\\sf S}}}}}_{n, 2:4} \\Vert _2$ respectively represent the total intensity and the polarized intensity. As can be noticed, this constraint also imposes implicitly the positivity of the total intensity image (Stokes I). The enforcement of this constraint amounts to incorporation of the indicator function of the set $\\mathbb {P}$ in the regularization function, i.e.\n(15)$$\\begin{eqnarray}\r\nr^\\prime ({\\boldsymbol{\\sf {S}}}) = \\iota _{\\mathbb {P}} ({\\boldsymbol{\\sf {S}}}).\r\n\\end{eqnarray}$$The indicator function of any such set $\\mathbb {P}$ is defined as\n(16)$$\\begin{eqnarray}\r\n\\iota _{\\mathbb {P}}({\\boldsymbol{\\sf {S}}}) = \\left\\lbrace \\begin{array}{@{}l@{\\quad }l@{}}0, \\quad \\quad \\text{if} \\, {\\boldsymbol{\\sf {S}}} \\in \\mathbb {P} \\, , \\\\\r\n+\\infty , \\quad \\text{otherwise}. \\end{array}\\right.\r\n\\end{eqnarray}$$","Citation Text":["Wiaux, Puy & Vandergheynst 2010"],"Citation Start End":[[3683,3714]]}
{"Identifier":"2020ApJ...905...45S__Frenklach_&_Feigelson_1989_Instance_1","Paragraph":"In the study presented here, two solid-phase experiments were run for 3 hours each, to produce cosmic grain analogs from methane (5%) and acetylene (5%) precursors, respectively, seeded in a cold supersonic expansion of argon. Different concentrations of methane were tested in COSmIC, and the 5% concentration was chosen to produce sufficient solid material for analysis in our experimental setup. Methane is a one carbon-atom alkane known to be present in the CSE of carbon-rich stars (Van de Sande & Millar 2019). Acetylene is a two-carbon alkene and the most abundant molecule in the inner region of the CSE where cosmic grains are formed, after H2 and CO (Fonfria et al. 2008), and is a known precursor of aromatic compounds (Frenklach & Feigelson 1989). For each experiment, chemistry was induced by the plasma discharge through fragmentation and ionization of the molecular precursor and plasma-generated chemical reactions between the precursor, its atomic and molecular fragments, radicals, and ions, resulting in the production of larger molecules and eventually solid particles. The resulting carbon grains, analogs of cosmic dust, were deposited on 3 mm diameter transmission electron microscopy (TEM) copper grids with Ultrathin Carbon Film on Lacey Carbon Support Film (Ted Pella, 400 mesh) placed 3 cm away from the slit, far enough away to avoid strong disturbances of the expansion. Ultra-high purity (UHP; 99.9998%) Ar and CH4 gas cylinders were used. The protocol for the addition of C2H2 to the gas mixture using a cold trap has been described in Sciamma-O\u2019Brien et al. (2014). After deposition, the samples were collected under inert argon atmosphere using a custom-made glove box connected to the chamber (see Figure 1), and placed in plastic containers sealed with parafilm to minimize exposure to air before the ex situ analysis of the sample at the SEM facility located at the NASA Ames Advanced Studies Laboratories (Sciamma-O\u2019Brien et al. 2017). The SEM technique was used to characterize the cosmic grain analogs produced in COSmIC. By using the same relative concentration for methane and acetylene (5%), only one experimental parameter (the nature of the precursor) was changed while maintaining all other experimental parameters (energy, temperature, precursor density, and reaction time) identical, thus allowing us to investigate the effect of the molecular precursors (CH4 versus C2H2) on the grain size and morphology under controlled conditions.","Citation Text":["Frenklach & Feigelson 1989"],"Citation Start End":[[731,757]]}
{"Identifier":"2018ApJ...866...15N__Ludwig_&_Steffen_2016_Instance_1","Paragraph":"We find that it is possible to infer \n\n\n\n\n\n and \n\n\n\n\n\n, at the precision of spectroscopy and relatively imprecise \n\n\n\n\n\n and \n\n\n\n\n\n for red-giant stars. We attempted to infer the [Fe\/H]; this label is available from the apogee spectroscopy for our stars. However, this label failed and, on inspection, no pixels correlated with [Fe\/H]. Therefore, contrary to the findings by Corsaro et al. (2017), we find that there is no information with respect to [Fe\/H] in the granulation signal from the Kepler multiepoch photometry. We note that corrections to scaling relations between \n\n\n\n\n\n and fundamental stellar parameters include both \n\n\n\n\n\n and [Fe\/H] (White et al. 2011; Guggenberger et al. 2016; Sharma et al. 2016). Furthermore, Viani et al. (2017) showed, in the case of \n\n\n\n\n\n, a dependence on mean molecular weight. While we do not find the signature in the ACF amplitude, this does indicate that a [Fe\/H] dependence might be expected, as was also suggested by 3D hydrodynamical simulations of convection (Collet et al. 2007; Ludwig & Steffen 2016). For our proposed methodology, the model should be applied to test data that is assumed to be derived from the underlying population as the training data. Nevertheless, to assess the impact of stellar metallicity on the inference of our labels, we performed a test where we divided stars into two groups around the mean metallicity of the sample. We created a training set of stars with [Fe\/H] > 0 and a test set with [Fe\/H]  0 dex. These stars broadly cover the same \n\n\n\n\n\n, \n\n\n\n\n\n, and \n\n\n\n\n\n and \n\n\n\n\n\n ranges, although the means of these labels are shifted. We find that the \n\n\n\n\n\n label is impacted at test time when training on a sample of stars with metallicities not covered in the test sample. The inferred \n\n\n\n\n\n is biased to be about 85 K too hot, for training on the metal-rich stars and testing on the more metal-poor stars and 85 K too cold, for the reverse (although with a similar precision as before). This indicates that the [Fe\/H] of the star is constraining as to the scale of the \n\n\n\n\n\n, even if we cannot learn this information from the data. The inference of the other three labels is not similarly affected by drawing the test and training set from different metallicity distributions. It is interesting that metallicity affects the \n\n\n\n\n\n and not the other labels. This means that temperature manifests itself differently in the spectra of stars with low and high metallicities. When we do not include stars with low metallicity in the training sample, we cannot accurately predict the temperatures of low metallicity stars during test, and vice versa, as temperature-induced variations in the spectrum must depend on metallicity. The fact that this is not true for other parameters implies that variations in seismic parameters produce variations in the ACF that are independent of metallicity; we therefore do not need to know the metallicity of our stars to infer these parameters to the precision captured using this technique. There is a theoretical expectation that the asteroseismic observables will vary with stellar metallicity (e.g., White et al. 2011; Guggenberger et al. 2016; Kallinger et al. 2018; Viani et al. 2017). Following on from this, there is also opportunity to investigate this in more detail with a data-driven approach.","Citation Text":["Ludwig & Steffen 2016"],"Citation Start End":[[1030,1051]]}
{"Identifier":"2021AandA...651A.103N__Stixrude_(2014)_Instance_1","Paragraph":"Radial velocity measurements (Anglada-Escud\u00e9 et al. 2016) can only give us a minimum planetary mass, as long as no additional information is available to constrain the inclination of the star\u2013planet system with respect to our observational plane. No transit of the planet has been observed so far, leading to the conclusion that the inclination is likely below 90\u00b0, and therefore the actual mass of the planet should be larger than 1.27 M\u2295. Anglada et al. (2017) proposed the existence of a dust belt at an orbit of about 30 AU from the star, which, if confirmed, would suggest an inclination of the system of 45\u00b0 with respectto the plane of the sky. Since confirmation has not been possible to date, we assumed four different inclinations spanning the wide range of likely inclinations and hence actual planet masses: 90, 60, 45, and 30\u00b0. The corresponding planet masses would then be 1.27, 1.47, 1.8, and 2.54 M\u2295 (see also Table 2). The table lists in addition the predicted planet and core radii determined with our interior structure model (Noack et al. 2017) as well as the assumed initial temperature at the core\u2013mantle boundary after planet accretion and solidification based on Stixrude (2014) and Noack & Lasbleis (2020). Both the iron content and the planet mass have a strong effect on the planetary structure and temperature profile. The resulting initial temperature profile in mantle and core, as well as the density, gravitational acceleration, pressure, and electrical conductivity profiles are shown in Fig. 1 for all investigated planet cases. The thermodynamic parameters (density as well as heat capacity and thermal expansivity of the material) were calculated following Noack et al. (2017) and using equations ofstate from Stixrude & Lithgow-Bertelloni (2011) for the mantle and Bouchet et al. (2013) for the metal core (here composed only of pure iron). The electrical conductivity profile is taken to be Earth-like from Xu et al. (2000). Thermal conductivity was calculated with the parametrisations derived in Tosi et al. (2013). For simplicity, the profiles for the different mantle properties were fixed during the evolution and did not take into account the influence of local melting or evolving chemical heterogeneity. The initial temperature profile, which is used in our interior evolution studies (described in Sect. 2.4), resembles the temperature profile depicted in Fig. 1. The temperature at the core\u2013mantle boundary, however, and therefore the adiabatic temperature in the core, was scaled with the core\u2013mantle boundary (CMB) temperature jump, which was varied in our study to investigate its effect on outgassing efficiency. Furthermore, we did not assume the existence of a primary atmosphere (see discussion in Sect. 4.1). During the evolution of Proxima Cen b, changes in atmospheric composition and mass can affect the surface temperature. However, since we did not couple our outgassing model to an atmospheric evolution model, surface temperatures remain constant throughout the evolution.","Citation Text":["Stixrude (2014)"],"Citation Start End":[[1186,1201]]}
{"Identifier":"2022AandA...659A..85F__Jeffries_et_al._2014_Instance_1","Paragraph":"We apply our analysis to five open clusters of ages between \u223c10 and 100 Myr that were observed within GES (25 Ori, Gamma Vel, NGC 2547, NGC 2451 B, and NGC 2516). These clusters were selected because they cover the age interval in which the effect of radius inflation could be significant, allowing us to investigate how it evolves with age. The 25 Ori cluster is a group of PMS stars that was discovered by Brice\u00f1o et al. (2005) in the Orion OB1a association, with an estimated age of 6\u221213 Myr (Downes et al. 2014; Brice\u00f1o et al. 2019; Kos et al. 2019; Zari et al. 2019); a dispersed, kinematically distinct population was also found in the region using data from the Gaia Second Data Release (DR2; e.g. Zari et al. 2019). Because only a few stars of the secondary population were observed by GES, only the main cluster is considered here. Gamma Vel (age \u223c10\u221220 Myr, Jeffries et al. 2014, 2017) and NGC 2547 (35\u2005\u00b1\u20053 Myr, Jeffries & Oliveira 2005) are both located in the Vela OB2 association at a relative separation of \u223c2\u00b0. Both clusters host two kinematically distinct populations (Jeffries et al. 2014; Sacco et al. 2015); the two Gamma Vel populations (Gamma Vel A and B) are also separated by \u223c38 pc along the line of sight (Franciosini et al. 2018). NGC 2451 is a double cluster composed of two open clusters of similar age (30\u221240 Myr, Randich et al. 2018) located at different distances along the same line of sight (R\u00f6ser & Bastian 1994; Platais et al. 1996). The GES observations cover the background cluster NGC 2451 B and only a few selected regions of the closer and more dispersed NGC 2451 A. For this reason, we considered only NGC 2451 B here. Finally, NGC 2516 is the oldest cluster in our sample, with an age of \u223c100\u2005\u2212\u2005140 Myr (e.g. Lyra et al. 2006; Randich et al. 2018), so that most of its members are already close to or at their main-sequence position. All clusters have solar or slightly subsolar metallicities (Biazzo et al. 2011; Jacobson et al. 2016; Spina et al. 2017).","Citation Text":["Jeffries et al. 2014"],"Citation Start End":[[868,888]]}
{"Identifier":"2018ApJ...853L...6A__Rezzolla_et_al._2011_Instance_1","Paragraph":"Before the detection of GW170817, many associated EM counterparts had already been proposed for NS\u2013NS\/BH mergers, and their relative brightness is essentially determined by the properties of the post-merger central remnant object. In general, the merger remnant for an NS\u2013BH merger would be a BH. But an NS\u2013NS merger could lead to a BH or a supramassive NS, depending on the total mass of the NS\u2013NS system and the NS equation of state (Dai et al. 2006; Gao & Fan 2006; Zhang 2013; Lasky et al. 2014; Gao et al. 2016). It is generally believed that, for both cases, sGRBs and their afterglow emission are expected as one of the major EM counterparts of NS\u2013NS\/BH mergers (Eichler et al. 1989; Narayan et al. 1992; Berger 2014). On the other hand, for both cases, the merger remnant would be surrounded by a mildly isotropic, sub-relativistic ejecta (which is composed of the tidally ripped and dynamically launched materials during the merger and the matter launched from the neutrino-driven wind from the accretion disk or neutron star surface; Rezzolla et al. 2011; Bauswein et al. 2013; Hotokezaka et al. 2013; Lei et al. 2013; Rosswog et al. 2013; Fern\u00e1ndez et al. 2015; Song & Liu 2017). These ejecta are mostly composed of neutron-rich materials, and the radioactivity of these materials and the decay of r-process nuclei would heat the ejecta and then power an optical\/IR transient (Li & Paczy\u0144ski 1998; Metzger et al. 2010). When the merger remnant is a BH, r-process related radioactivity would be the only heating source, so that the luminosity of the optical\/IR transient would be roughly \u223c103 times that of the nova luminosity (Metzger et al. 2010). But if the merger remnant is a supramassive NS, its magnetic dipole radiation could serve as an additional heating source to the ejecta (Gao et al. 2013; Zhang 2013), which could easily exceed the r-process power, so that the thermal emission from the ejecta would be significantly enhanced (Yu et al. 2013; Metzger & Piro 2014). The luminosity of the optical transient, in this case, would be systematically brighter by more than one order of magnitude than the r-process dominating cases (Gao et al. 2017). Since the thermal emissions from the ejecta are essentially isotropic and also non-relativistic or mildly relativistic (due to the heavy mass loading), they therefore can be detected from any direction if the flux is high enough (see Metzger 2017a for a review).","Citation Text":["Rezzolla et al. 2011"],"Citation Start End":[[1044,1064]]}
{"Identifier":"2017MNRAS.464..635M__Dekel_et_al._2009_Instance_2","Paragraph":"The basic idea, summarized in Dekel et al. (2009), is that during VDI, the high surface density of gas and \u2018cold\u2019 young stars, \u03a3, drives the Toomre Q parameter below unity, Q \u223c \u03c3\u03a9\/(\u03c0G\u03a3) \u2272 1, where \u03c3 is the 1D velocity dispersion and \u03a9 is the angular frequency, a proxy to the epicyclic frequency \u03ba, which is related to the potential well (Toomre 1964). It has been established that under such conditions, the disc will fragment and produce large star-forming clumps. This has been shown using idealized simulations of isolated galaxies (Noguchi 1999; Gammie 2001; Immeli et al. 2004a,b; Bournaud, Elmegreen & Elmegreen 2007; Elmegreen, Bournaud & Elmegreen 2008; Bournaud & Elmegreen 2009; Hopkins et al. 2012b), as well as cosmological simulations (Agertz, Teyssier & Moore 2009; Ceverino et al. 2010; Ceverino et al. 2012; Genel et al. 2012; Mandelker et al. 2014; Oklopcic et al. 2016). The ratio of clump mass to the mass of the cold disc scales as Mc\/Md \u221d \u03b42, where \u03b4 = Md\/Mtot is the ratio of the cold disc mass to the total mass within the disc radius, which includes the bulge and dark matter halo (e.g. Dekel et al. 2009). This leads to much larger clumps at z \u223c 2 than the low-redshift giant molecular clouds (GMCs). Gravitational interactions in the perturbed disc drive turbulence causing the disc to self-regulate in a marginally stable state with Q \u2272 1 (Dekel et al. 2009; Ceverino et al. 2010; Krumholz & Burkert 2010; Cacciato, Dekel & Genel 2012; Forbes, Krumholz & Burkert 2012; Forbes et al. 2014) that can last for more than a Gyr so long as the accretion is not interrupted. Some recent works have called into question the validity of linear Toomre analysis in the context of these highly non-linear galaxies (Behrendt, Burkert & Schartmann 2015; Tamburello et al. 2015; Inoue et al. 2016) and others have suggested alternate fragmentation mechanisms related to turbulence (e.g. Hopkins 2013). However, since clump formation is largely determined by the balance between self-gravity, turbulent pressure and the centrifugal force, the largest clumps are always roughly at the Toomre scale. Larger clumps would be disrupted due to the shear and\/or tidal forces within the disc, or would not collapse in the first place due to the centrifugal force. Therefore, regardless of the full validity of linear Toomre analysis, it is plausible that the Toomre Q parameter can serve as a crude criterion for instability, possibly with a critical value that is larger than unity.","Citation Text":["Dekel et al. 2009"],"Citation Start End":[[1112,1129]]}
{"Identifier":"2018AandA...619A..13V__Saviane_et_al._2012_Instance_3","Paragraph":"The EWs were measured with the methods described in V\u00e1squez et al. (2015). As in Paper I, we used the sum of the EWs of the two strongest CaT lines (\u03bb8542, \u03bb8662) as a metallicity estimator, following the Ca\u202fII triplet method of Armandroff & Da Costa (1991). Different functions have been tested in the literature to measure the EWs of the CaT lines, depending on the metallicity regime. For metal-poor stars ([Fe\/H] \u2272 \u22120.7 dex) a Gaussian function was used with excellent results (Armandroff & Da Costa 1991), while a more general function (such as a Moffat function or the sum of a Gaussian and Lorentzian, G + L) is needed to fit the strong wings of the CaT lines observed in metal-rich stars (Rutledge et al. 1997; Cole et al. 2004). Following our previous work (Gullieuszik et al. 2009; Saviane et al. 2012) we have adopted here a G+L profile fit. To measure the EWs, each spectrum was normalised with a low-order polynomial using the IRAF continuum task, and set to the rest frame by correcting for the observed radial velocity. The two strongest CaT lines were fitted by a G+L profile using a non-linear least squares fitting routine part of the R programming language. Five clusters from the sample of Saviane et al. (2012) covering a wide metallicity range were re-reduced and analysed with our code to ensure that our EWs measurements are on the same scale as the template clusters used to define the metallicity calibration. Figure 5 shows the comparison between our EWs measurements and the line strengths measured by Saviane et al. (2012) (in both cases the sum of the two strongest lines) for the five calibration clusters. The observed scatter is consistent with the internal errors of the EW measurements, computed as in V\u00e1squez et al. (2015). The measurements show a small deviation from the unity relation, which is more evident at low metallicity. A linear fit to this trend gives the relation: \u03a3EW(S12) = 0.97 \u03a3EW(this work) + 0.21, with an rms about the fit of 0.13 \u00c5. This fit is shown in Fig. 5 as a dashed black line. For internal consistency, all EWs in this work have been adjusted to the measurement scale of Saviane et al. (2012) by using this relation. In Table 3 we provide the coordinates, radial velocities, and the sum of the equivalent widths for the cluster member stars, both measured (\u201cm\u201d) and corrected (\u201cc\u201d) to the system of Saviane et al. 2012.","Citation Text":["Saviane et al. (2012)"],"Citation Start End":[[1530,1551]]}
{"Identifier":"2016ApJ...831..131L__Wang_et_al._2007_Instance_1","Paragraph":"The relationship established in our study between the occurrence rates of major flares and FHCMEs and the annual average latitude of sunspot groups may indicate the role of long-distance magnetic connectivity in the dependence of occurrence rates on the SC phase. In particular, the annual average latitude of sunspot groups determines the average length of trans-equatorial loops (TLs). A TL is a coronal loop connecting two separate ARs located in the opposite hemispheres (e.g., Tsuneta 1996; Pevtsov 2000). It has been found to be a common phenomenon, with up to one-third of all ARs exhibiting TLs in soft X-ray images (Pevtsov 2000). Therefore, our results may be a statistical indication that large-scale magnetic connectivities highlighted by TLs are important for producing solar activity such as major flares and FHCMEs. There have been many studies on the topology of the large-scale solar magnetic field structure (e.g., Babcock 1961; Hansen & Hansen 1975) and the relationship between the occurrence of solar flares \/ CMEs and large-scale magnetic connectivity (e.g., Antiochos et al. 1999; Khan & Hudson 2000; Wang & Sheeley 2003; Wang et al. 2007; Zhou et al. 2007; Fu & Welsch 2016). In particular, Dalla et al. (2007) established that the flaring rate (weakly) increases with the appearance of a new AR within 12\u00b0 of the pre-existing flaring AR. Also, it has been reported that the interaction of ARs increases as sunspots move toward the equator (McIntosh & Leamon 2014; McIntosh et al. 2014). Such interactions may be related to the reconnection of magnetic fields that belong to different ARs, and the subsequent formation of large-scale structures such as TLs, which play an important role in triggering flares and CMEs (Khan & Hudson 2000). Moon et al. (2002) showed that sympathetic flares are more frequent in TLs than in loops connecting two ARs in the same hemisphere. Chen et al. (2006) examined the relation between the number of soft X-ray TLs and that of ARs each year from 1991 to 2001. They found good correlations of the TL numbers with SC indices, with the increase in TL number during the descending phase of SC 22 in comparison to other phases of SC 22. The effect of the dependence of flaring activity on AR separation and connectivity may be understood in terms of the interaction of magnetic energy between these ARs (Fu & Welsch 2016).","Citation Text":["Wang et al. 2007"],"Citation Start End":[[1145,1161]]}
{"Identifier":"2019AandA...625A.116V__Tagger_&_Varni\u00e8re_2006_Instance_1","Paragraph":"From analytical and semi-analytical work (see for example Lovelace & Romanova 2014) we know that the local conditions where the RWI grows are what trigger the growth of one mode over another and in turn the strengths of each mode directly impacts the PDS. In the following section we focus on only one spin (a\u2004=\u20040.995) and position (rc\u2004=\u20042.7rg) where the RWI is triggered but look at different local conditions (see Eq. 3 for the parametrization), meaning the (unstable) zone where the RWI is active will be different. We can then see how the local condition affect the PDS, especially focusing on which modes are selected and how they relate to one another when present. We saw in previous studies of the RWI (Tagger & Varni\u00e8re 2006; Vincent et al. 2013; Casse et al. 2017; Casse & Varniere 2018) that every simulation triggers a lot of modes, with up to m\u2004=\u200419 being detectable in Fourier space; though lower modes tend to be a lot stronger, especially once the instability reaches saturation. Because of the timescale near the last stable orbit of a black-hole of ten solar masses, we mostly observe the saturated state of the instability in a standard observation. We have therefore first performed several simulations with the same aforementioned corotation radius rc\u2004=\u20042.7rg but with various \u03b5 and \u03b4 parameters leading to different geometries of the unstable zone. A typical GRHD simulation example and the related ray-tracing synthetic observation can be found in Fig. 1. Once a simulation has reached the saturation stage, namely when the amplitude of the fluctuations remains approximately constant, we stop the simulation and proceed to use GYOTO in order to compute the light curve of the simulation as perceived by a remote observer. We stress here that the combination of the GR-AMRVAC code and the GYOTO code encompasses all general relativistic effects at work upon both the disc and its radiative emission, here providing reliable synthetic observations.","Citation Text":["Tagger & Varni\u00e8re 2006"],"Citation Start End":[[711,733]]}
{"Identifier":"2018MNRAS.475.3065G__Feger_et_al._2012_Instance_1","Paragraph":"Substantial effort has been made by research groups in this field to improve the fibre link performance by employing non-circular fibres, double scramblers, fibre shakers, and fibre stretchers (Pepe et al. 2000; Chazelas et al. 2010; Perruchot et al. 2011; Avila 2012; Chazelas, Pepe & Wildi 2012; Bouchy et al. 2013; Reynolds & Kost 2014; Halverson et al. 2015a; Spronck et al. 2015; Ishizuka et al. 2016; Sutherland et al. 2016). Recent test results on the performance of MM optical fibres with non-circular cores show improved scrambling performance over that of standard circular MM fibres (Chazelas et al. 2012; Feger et al. 2012; Sturmer et al. 2016). However, a native non-circular fibre has also been shown to fail to provide the high level of stability and homogeneity in both the near-field and far-field of the fibre that is required for the next generation of PRV instruments operating below the 1 m s\u22121 radial velocity precision level. Much work has been done to improve the degree of stability and homogeneity that the fibre link can deliver by investigating and testing various combinations of fibre core geometries combined with a double scrambler. Optical double scramblers placed in the fibre link act to exchange the near and far-field light distributions of the fibre by means of a lens rely (Barnes & MacQueen 2010). This method has been shown to improve the level of scrambling in both the near and far-fields of the fibre output (Halverson et al. 2015a; Seifahrt, Sturmer & Bean 2016). Although these methods improve the fibre scrambling they do not necessarily reduce modal noise. Modal noise mitigation\/suppression requires a time variable correction which is typically achieved by means of shaking and\/or stretching the fibre (Reynolds & Kost 2014; Roy et al. 2014; Halverson et al. 2015a). In this paper, we present a new fibre technology (multicore fibre and photonic lanterns) that has the possibility to provide improved fibre scrambling and modal noise characteristics.","Citation Text":["Feger et al. 2012"],"Citation Start End":[[617,634]]}
{"Identifier":"2017MNRAS.468.2590S__Fassnacht_et_al._2002_Instance_1","Paragraph":"We initiated the H0LiCOW (H0 Lenses in COSMOGRAIL's Wellspring) program with the aim of measuring the Hubble constant with better than 3.5\u2009per\u2009cent precision and accuracy (in most background cosmological models), through a sample of five time-delay lenses. We obtain the key ingredients to each of the lenses through observational follow-ups and novel analysis techniques. In particular, we have high-quality lensed quasar light curves, primarily obtained via optical monitoring by the COSMOGRAIL (COSmological MOnitoring of GRAvItational Lenses; e.g. Courbin et al. 2005; Vuissoz et al. 2008; Courbin et al. 2011; Tewes et al. 2013b) and Kochanek et al. (2006) teams but also via radio-wavelength monitoring (Fassnacht et al. 2002). COSMOGRAIL has been monitoring more than 20 lensed quasars for more than a decade. The unprecedented quality of the light curves combined with new curve-shifting algorithms (Tewes, Courbin & Meylan 2013a) lead to time delays with typically \u223c3 per cent accuracy (Fassnacht et al. 2002; Courbin et al. 2011; Tewes et al. 2013b). In addition, we obtain HST imaging that reveal the \u2018Einstein ring\u2019 of the lens systems in high resolution, and develop state-of-the-art lens modelling techniques (Suyu et al. 2009; Suyu & Halkola 2010; Suyu et al. 2012b) and kinematic modelling methods (Auger et al. 2010; Sonnenfeld et al. 2012) to obtain the lens mass distribution with a few percent uncertainty (e.g. Suyu et al. 2013, 2014). We further obtain wide-field imaging and spectroscopy to characterize the environment of the field, as well as the spectroscopy of the lens galaxy to obtain the stellar velocity dispersion. The exquisite follow-up data set that we have acquired allow us not only to constrain cosmology but also to study lens galaxy and source properties for understanding galaxy evolution, including the dark matter distribution in galaxies, the stellar initial mass function of galaxies and the co-evolution between supermassive black holes and their host galaxies.","Citation Text":["Fassnacht et al. 2002"],"Citation Start End":[[710,731]]}
{"Identifier":"2021ApJ...920..139M__the_1997_Instance_1","Paragraph":"Cepheus X-4 was discovered as a transient source using the X-ray telescope of the OSO-7 satellite during 1972 June\u2013July (Ulmer et al. 1973). Ginga observed the source during the 1988 March outburst and detected a spin period of 66.25 s for its neutron star for the first time (Koyama et al. 1991). Spectroscopic studies from the same Ginga observations led to the detection of a cyclotron resonance scattering feature corresponding to a centroid energy at 30.5 \u00b1 0.4 keV (Mihara et al. 1991). The ROSAT observations during the 1993 June outburst refined the source coordinates and also determined its pulsar spin period (Schulz et al. 1995). The observations by BATSE during 1993 June\u2013July and during the subsequent outburst of 1997 June\u2013July, which was also followed by RXTE, determined the pulse characteristics of Cepheus X-4. Additionally, a possible range of its orbital period from 23 to 147.3 days was suggested using RXTE data (Wilson et al. 1999). From observed characteristic features and outburst activities, it was predicted that Cepheus X-4 could possibly have a massive early-type Be star with its circumstellar disk as a companion, which was thought to be the most likely cause of its long outburst of about 40 days, as seen for other Be binaries. Optical observations of Cepheus X-4 subsequently confirmed it as a Be binary system and estimated its location at a distance of 3.8 \u00b1 0.6 kpc (Bonnet-Bidaud & Mouchet 2005). But this distance estimate was later challenged by Riquelme et al. (2012), who proposed a distance of either 7.9 kpc or 5.9 kpc according to whether the stellar type of the companion is a B1 or B2 star, respectively. The distance of Cepheus X-4 was later reported by Gaia to be \n\n\n\n\n\n\n10.2\n\n\n\u2212\n1.6\n\n\n+\n2.2\n\n\n\n\n kpc (Malacaria et al. 2020). Luminosity dependent changes in the pulse profile of Cepheus X-4 were studied during the declining phase of the 1997 outburst (Mukerjee & Agrawal et al. 2000) by combining observations by the Indian X-ray Astronomy Experiment (Agrawal et al. 1996) and RXTE (Rothschild et al. 1998). The RXTE observed another outburst in 2002 and re-established its cyclotron resonance feature corresponding to a centroid energy at 30.7 \u00b1 1.8 keV (as established earlier by Ginga), which did not show a significant dependence on X-ray luminosity, although the continuum became harder with increasing source luminosity (McBride et al. 2007). The source went into outburst again in 2014 and was observed with the Nuclear Spectroscopic Telescope Array (NuSTAR; Harrison et al. 2013) and Suzaku (Mitsuda et al. 2007). The Suzaku observation of the 2014 outburst overlapped with the second observation with NuSTAR on 2014 July 1\u20132, and detected an additional absorption feature at \u224845 keV in the phase resolved spectra of the pulsar, which was identified as the first harmonic of the fundamental cyclotron line detected at \u224828 keV (Jaiswal & Naik 2015). The source spectra obtained from the NuSTAR observations in 2014 were well fitted by a Fermi\u2013Dirac cutoff (FD-cutoff) model along with an iron emission line, and a cyclotron absorption feature was clearly detected in both of the observations at \n\n\n\n\n\n\n30.39\n\n\n\u2212\n0.14\n\n\n+\n0.17\n\n\n\n\n keV and \n\n\n\n\n\n\n29.42\n\n\n\u2212\n0.24\n\n\n+\n0.27\n\n\n\n\n keV, respectively. Although, the averaged source luminosity differed by a factor of about 3 between these two observations, it only showed a marginal variation in its centroid energy (Furst et al. 2015). Using the same NuSTAR observations of the 2014 outburst, Vybornov et al. (2017) reported that the spectrum of Cepheus X-4 showed two cyclotron resonance scattering features, with the fundamental line at \u224830 keV and its harmonic at \u224855 keV. They also showed that the energy of the fundamental cyclotron absorption feature increases and the continuum becomes harder with increasing X-ray luminosity. Pulse phase resolved spectroscopic studies of Cepheus X-4 were conducted by Bhargava et al. (2019) at these two different intensities of the source using the same two NuSTAR observations of the 2014 outburst. It was found that the observed cyclotron line profile of Cepheus X-4 had an asymmetric shape in the phase averaged spectrum. However, for phase resolved spectra, a single symmetric cyclotron profile fitted the data well. The spectral continuum and the parameters of the cyclotron line showed some variations with respect to the pulse phase only within a limited pulse phase (Bhargava et al. 2019).","Citation Text":["Mukerjee & Agrawal et al. 2000"],"Citation Start End":[[1904,1934]]}
{"Identifier":"2021MNRAS.500.2928E__Uzdensky_2004_Instance_1","Paragraph":"The diffusion time-scale of the field lines in the disc is comparable to the viscous time-scale which is much longer than the interaction time-scale of the field lines and the inner disc, tint = |\u03a9K \u2212 \u03a9*|\u22121 (Fromang & Stone 2009), where \u03a9* is the angular velocity of the neutron star. The field lines cannot slip through the disc. Theoretical studies and numerical simulations show that the field lines interacting with the inner disc in a narrow boundary inflate and open up within the interaction time-scale (Aly 1985; Lovelace, Romanova & Bisnovatyi-Kogan 1995; Hayashi, Shibata & Matsumoto 1996; Miller & Stone 1997; Uzdensky, K\u00f6nigl & Litwin 2002; Uzdensky 2004). If the system is in the SP phase, the matter flowing in to the boundary at the innermost region of the disc is expelled from the system along the open field lines. The open lines reconnect on the dynamical time-scale, $\\Omega _{\\mathrm{K}}^{-1}$, and continue to apply torque on the matter until they open up again (Lovelace, Romanova & Bisnovatyi-Kogan 1999; Ustyugova et al. 2006). The simulations show that the field lines outside the interaction boundary are decoupled from the disc. The magnetosphere inside the boundary is the region where the closed field lines and the plasma can rotate together at radius rm \u2243 rin. These results indicate that a strong-propeller mechanism could be sustained if the field lines expel the matter from the inner boundary at the same rate as that of the mass-flow from the outer disc. Furthermore, for this strong-propeller phase to be steady, the matter should be accelerated to speeds greater than the escape speed, \u03c5esc, within tint. Ertan (2017) showed that the maximum radius at which the strong-propeller condition is satisfied is much smaller than rA, while the accretion rates estimated for the WP\/SP transition in this model seem to be in agreement with the transition properties of tMSPs (Ertan 2017, 2018). The inner boundary in the spin-up phase is also continuously evacuated, as in the case of strong-propeller, but now due to accretion on to the star. At rin the field lines decelerate and bring the matter into co-rotation within tint in the spin-up phase. In other words, the same equation could determine the inner disc radius in both the strong propeller phase (for rin > rco) and the spin-up phase (for rin rco). However, this is not the whole story, since the conditions in the spin-up (SU) phase are rather different, in particular rA enters the picture (Sections 2.1\u20132.4).","Citation Text":["Uzdensky 2004"],"Citation Start End":[[653,666]]}
{"Identifier":"2022AandA...664A..63X__Teodoro_&_Fraternali_2015_Instance_1","Paragraph":"We note that there has recently been some debate about potential observational effects (beam-smearing effect) on the dynamical state of (field) high-z galaxies (e.g., Di Teodoro et al. 2016; Kohandel et al. 2020). More specifically, when studying the gas kinematics with low spectral and spatial resolution data, the unresolved rotations within the PSF can artificially increase the value of \u03c30 and decrease the value of Vmax, leading to a severe systematic underestimation of the Vmax\/\u03c30 ratio. The beam-smearing effect has been widely explored in local galaxies (e.g., see Fig. 6 in Di Teodoro & Fraternali 2015) and high-z simulated galaxies (e.g., Kohandel et al. 2020). Subsequently, given the high spatial and spectral resolution of ALMA observations, which have been able to reach subkpc spatial resolution in high-z galaxies, the beam-smearing effect is less significant. Compared to previous data, ALMA allows for a more robust way to measure intrinsic velocity dispersions. However, in our case, comparing two cluster SBs (with a spatial resolution of \u223c0.3\u2033 and a channel width of 30 km s\u22121) with other field SBs at similar redshifts and similar resolutions (with spatial resolutions of 0.1\u2033\u20130.7\u2033 and channel widths of 20\u201350 km s\u22121, except for one with a channel width of \u223c100 km s\u22121), the \u03c30 values of the two cluster SBs are about three times lower than that of those field SBs (blue triangles in Fig. 6; Calistro Rivera et al. 2018; Barro et al. 2017; Swinbank et al. 2011; Tadaki et al. 2017). In addition, even at the same spatial and spectroscopical resolution, the \u03c30 values of the two cluster SBs are still more than three times lower than those of the two cluster MSs (green stars in Fig. 6). Furthermore, the simulations have confirmed the robustness of \u03c30 and Vmax for the two SBs and two MSs in our measurements (see Sect. 3.6). Therefore, we argue that the higher Vmax\/\u03c30 and lower \u03c30 of the two cluster SBs with respect to field galaxies is a real physical difference, rather than being driven by observational effects.","Citation Text":["Di Teodoro & Fraternali 2015"],"Citation Start End":[[585,613]]}
{"Identifier":"2022MNRAS.515.5135H__Saito_&_Gary_2007_Instance_1","Paragraph":"In the present work, we neglect the processes that generate the halo, but this topic deserves some review. Notably, the apparent growth of the halo at the expense of the anti-sunward suprathermal \u2018strahl\u2019 population may imply that the halo is locally formed in the inner heliosphere by scattered strahl electrons (e.g. Maksimovic et al. 2005; \u0160tver\u00e1k et al. 2009). This has led to significant theoretical development, focused on the resonant interaction of electrons with the whistler and fast-magnetosonic whistler (FM\/W) modes (e.g. Vocks et al. 2005; Saito & Gary 2007; Vasko et al. 2019; Verscharen et al. 2019b; Micera et al. 2021; Zenteno-Quinteros, Vi\u00f1as & Moya 2021; Tang, Zank & Kolobov 2022; Vo, Lysak & Cattell 2022). Observations have struggled to confirm these theories. Notably, whistlers are practically absent (occurrence rate 0.1\u2009per\u2009cent) during PSP perihelion passes (Cattell et al. 2022). Additionally, the eVDFs sampled by Helios and PSP are stable with respect to the oblique FM\/W mode (Jeong et al. 2022a). Theoretical calculations show that at r \u2272 1\u2009AU, the strahl is stable to whistler fluctuations (Horaites et al. 2018b; Schroeder et al. 2021) and should be unaffected by whistler turbulence in the inner heliosphere (Boldyrev & Horaites 2019). High-resolution measurements of the strahl at 1\u2009AU confirm that \u2018anomalous diffusion\u2019, e.g. from whistler waves, is not required to explain the strahl angular widths at resolvable energies \u2272300\u2009eV Horaites et al. (2018a), Horaites, Boldyrev & Medvedev (2019). Similar results were found from simulations at distances r \u2272 20RS (Jeong et al. 2022b), which showed that near the corona the strahl is adequately described by a combination of Coulomb collisions and expansion effects. This all suggests that a mechanism besides local wave particle scattering may account for the halo\u2019s presence in the inner heliosphere. Such theories have been proposed (e.g. Leubner 2004; Lichko et al. 2017; Che et al. 2019; Horaites et al. 2019; Scudder 2019), though no consensus has emerged.","Citation Text":["Saito & Gary 2007"],"Citation Start End":[[554,571]]}
{"Identifier":"2018MNRAS.480.5113M__Feretti_&_Giovannini_1996_Instance_1","Paragraph":"In this paper we extend previous work based on cosmological simulations by analysing the general magnetic field properties and the diffuse radio halo emission in galaxy clusters in the IllustrisTNG project, a set of cosmological magnetohydrodynamics simulations run with the moving-mesh code arepo (Springel 2010) that include a comprehensive module for galaxy formation physics. The main and novel aspect of our work is the analysis of the diffuse radio emission resulting from radio haloes in galaxy clusters (Feretti & Giovannini 1996; Murgia et al. 2009; Vacca et al. 2011; Feretti et al. 2012). We investigate radio emission from clusters by a detailed comparison with observations, trying to match the current observational constraints and to make predictions for the upcoming radio surveys that will be performed with the new generation of radio instruments such as SKA and LOFAR. The analysis of simulated radio haloes gives us a complementary view on the spatial extent and energy content of magnetic fields in galaxy clusters, since the radio emission is proportional to their strength. As such, the study of radio halo scaling relations (Giovannini et al. 2009; Cassano et al. 2013; Zandanel, Pfrommer & Prada 2014) with the total X-ray power and halo mass may yield important information about the amplification mechanisms of magnetic fields in clusters and the level of turbulence in the ICM. The modelling of radio emission makes it also possible to study the transport of charged particles and their re-acceleration to relativistic speeds, and it constrains the probability of detecting extended radio-emitting structures in a statistical sample of realistic simulated clusters. The comparison of the simulated radio emission with actual observations might also be employed as a useful check for the implementation of the galaxy formation physics modules used to perform the simulations, although our modelling of relativistic particles is rather preliminary and might have a non-negligible impact on the results.","Citation Text":["Feretti & Giovannini 1996"],"Citation Start End":[[512,537]]}
{"Identifier":"2016AandA...595A..83E__Visser_et_al._2015_Instance_1","Paragraph":"The choice of these initial abundances is motivated by the following two scenarios about the history of the midplane material. The first scenario is that the material going into protostellar systems is inherited from the cloud out of which the protoplanetary disk collapsed and formed. This scenario is denoted \u201cinheritance\u201d, and it implies that the material has the same composition as found in dark clouds, especially their ices (see Marboeuf et al. 2014; Mumma & Charnley 2011). The second scenario is the case where the material coming from the dark cloud experiences heating events from the protostar (i.e., accretion bursts or regular stellar irradiation). These heating events are assumed to alter the chemistry in disks significantly (see, e.g., Visser et al. 2015). In the extreme case, the chemistry is reset, meaning that the molecules are assumed to be dissociated into atoms out to R = 30 AU, which can then reform molecules and solids in a condensation sequence, as traditionally assumed for the inner solar nebula (e.g. Grossman 1972). Hence, the scenario considering atomic initial abundances is denoted \u201creset\u201d. Early chemical models of protoplanetary disks often assumed a set atomic initial abundances (e.g., Willacy et al. 1998; Aikawa et al. 1999; Semenov et al. 2004; Vasyunin et al. 2008; Walsh et al. 2010); however, these early models also did not typically include a comprehensive grain-surface network and focussed solely on the gas-phase chemistry. Some recent physico-chemical models do include surface formation of e.g., H2O ice, via the Eley-Rideal mechanism (e.g., Kamp et al. 2013; Bruderer et al. 2014; Helling et al. 2014). Early models that did include grain-surface chemistry (e.g. Willacy 2007; Walsh et al. 2010) were limited to simple atom-addition stemming from Tielens & Hagen (1982) and Hasegawa & Herbst (1993). Here, a more comprehensive grain-surface network is used which includes radical-radical recombination, atom addition, and also ice processing (Garrod et al. 2008). This work differs from earlier protoplanetary disk models in that we directly compare and quantify the effects of the inheritance versus reset scenarios for a single disk model using a comprehensive gas-grain chemical network. ","Citation Text":["Visser et al. 2015"],"Citation Start End":[[754,772]]}
{"Identifier":"2018MNRAS.481.1726G__Robin_et_al._2012_Instance_1","Paragraph":"A pragmatic solution to these problems is to generate and analyse synthetic Milky Way catalogues cast in the observational frame of the survey (Bahcall & Soneira 1980; Robin & Creze 1986; Bienayme, Robin & Creze 1987). \u2018Mock catalogues\u2019 of this general type were first used in cosmology in the 2000s (e.g. Cole et al. 2005) and have now become an essential tool for the design and analysis of large galaxy and quasar surveys. Realistic mock catalogues provide assessments of an instrument\u2019s capabilities and biases, tests of statistical modelling techniques applied to realistic representations of observational data, and detailed comparisons between theoretical predictions and observations. Perhaps one of the best known recent attempts is the Besan\u00e7on model (Robin et al. 2003), which provides a disc (or set of discs) with a set of coeval and isothermal (single velocity dispersion) stellar populations assumed to be in equilibrium, with analytically specified distributions of density, metallicity, and age. This has been the basis of the Gaia Universe Model (GUMS; Robin et al. 2012). However, these models are not dynamically consistent and oversimplify the structure of the Galaxy, particularly the stellar halo that is modelled as a smooth component. An important advance was made by Sharma et al. (2011), who developed the galaxia code for creating mock stellar catalogues either analytically or from phase-space sampling of hybrid semi-analytic-N-body simulations to represent stellar haloes in a cosmological context (Bullock & Johnston 2005; Cooper et al. 2010). Rybizki et al. (2018) have developed a mock catalogue designed specifically for Gaia DR2 based on galaxia. Building on the method of Sharma et al. (2011), Lowing et al. (2015) developed a technique to distribute synthetic stars sampled from a cosmological N-body simulation in such a way as to preserve the phase-space properties of their parent stellar populations. In a separate method, Hunt et al. (2015) introduced the snapdragons code that generates a mock catalogue taking into account Gaia errors and extinction and demonstrated the resulting observable kinematics of stars around a spiral arm in an idealized smoothed particle hydrodynamic simulation set-up in isolation.","Citation Text":["Robin et al. 2012"],"Citation Start End":[[1071,1088]]}
{"Identifier":"2017MNRAS.464.2120S__Taylor_et_al._2012_Instance_1","Paragraph":"Hu, Holz & Vale (2007b) proposed the idea of using the ratio of CMB convergence to galaxy lensing convergence as a way to measure the distance ratio (distance to surface of last scattering relative to the distance to the source galaxy sample used to estimate the galaxy lensing) and hence constrain the geometry, \u03a9k and the equation of state of dark energy. The ratio is defined as\n\n(40)\n\n\\begin{equation}\n\\mathcal {R}(z_{\\rm l})=\\frac{\\kappa (z_{\\rm l},z_{\\ast })}{\\kappa (z_{\\rm l},z_{\\rm s})}=\\frac{\\Sigma _{\\rm c}(z_{\\rm l},z_{\\rm s})}{\\Sigma _{\\rm c}(z_{\\rm l},z_{\\ast })}.\n\\end{equation}\n\nSimilar distance ratio tests have also been proposed using galaxy or galaxy cluster lensing alone, in both strong lensing (e.g. Link & Pierce 1998; Golse, Kneib & Soucail 2002) and weak lensing regimes (e.g. Jain & Taylor 2003; Bernstein & Jain 2004). Several studies have already measured the distance ratios (e.g. Taylor et al. 2012; Diego et al. 2015; Kitching et al. 2015; Caminha et al. 2016, and references therein), though they are afflicted by several systematics such as, uncertainties in modelling cluster profiles and cosmic variance in the case of multiple strong lens systems, and photometric redshift uncertainties as well as imaging systematics that cause a redshift-dependent shear calibration in the case of weak lensing. The small redshift baseline also limits the cosmological applications of these measurements using optical weak lensing alone (see discussion in Hu et al. 2007b; Weinberg et al. 2013). Using CMB lensing in cosmographic measurements is advantageous in several ways. First, the source redshift for the CMB (redshift of surface of last scattering) is well known, so one of the two redshift slices being compared has no redshift uncertainty. The long redshift baseline between CMB and galaxy lensing sources also improves the sensitivity of $\\mathcal {R}$ to cosmological parameters (Hu et al. 2007b). However, using CMB lensing with galaxy lensing makes $\\mathcal {R}$ become more sensitive to some of the systematics in galaxy lensing (for example, multiplicative bias), and $\\mathcal {R}$ can also be used as test for the presence of such systematics.","Citation Text":["Taylor et al. 2012"],"Citation Start End":[[911,929]]}
{"Identifier":"2020ApJ...895...81R__Andrews_&_Thompson_2011_Instance_1","Paragraph":"The SFR surface densities of \u03a3SFR = (750 \u00b1 440) and (1800 \u00b1 700) M\u2299 yr\u22121 kpc\u22122 found above for GN10 and the central region of AzTEC-3 are even higher than the values found for some other compact starbursts like ADFS-27 (z = 5.66; 430 \u00b1 90 M\u2299 yr\u22121 kpc\u22122) and HFLS3 (z = 6.34; 480 \u00b1 30 M\u2299 yr\u22121 kpc\u22122; Riechers et al. 2013, 2017). The source-averaged value of >(500 \u00b1 160) M\u2299 yr\u22121 kpc\u22122 for AzTEC-3 is comparable to these sources. Like these systems, GN10 and AzTEC-3 thus show the properties expected for so-called \u201cmaximum starbursts\u201d. At face value, the peak \u03a3SFR in AzTEC-3 may slightly exceed the expected Eddington limit for starburst disks that are supported by radiation pressure (e.g., Thompson et al. 2005; Andrews & Thompson 2011), but it is potentially consistent under the assumption of a more complex source geometry. On the other hand, the high \u03a3SFR value of GN10 could be boosted by an obscured AGN contribution to the dust heating. GN10 exhibits strong 0.5\u20138 keV X-ray emission.43\n\n43\nAzTEC-3 is not detected at X-ray wavelengths.\n Using Equation (15) of Ranalli et al. (2003), its observed (absorption-corrected)44\n\n44\nFrom fitting a Galactic absorption plus power-law model, an effective photon index \u0393 = \n\n\n\n\n\n was found for GN10, but it was not possible to simultaneously fit the absorbing column NH and \u0393 due to the limited photon counts. The data are consistent with no intrinsic absorption within the uncertainties, such that the absorption-corrected LX could be considered an upper limit (Laird et al. 2010) or, alternatively, a lower limit, in the case that it is heavily absorbed.\n rest-frame 2\u201310 keV X-ray luminosity of LX = 5.6(12.5) \u00d7 1042 erg s\u22121 (Laird et al. 2010) corresponds to an SFRX of 1100(2500) M\u2299 yr\u22121. Given its SFRIR = \n\n\n\n\n\n M\u2299 yr\u22121, its LX remains consistent with intense star formation, but a contribution from a modestly luminous obscured AGN cannot be ruled out, depending on the (relatively uncertain) absorption correction required. This would also be consistent with a possible excess mid-infrared emission due to an obscured AGN, which may be favored by some of the SED fits.","Citation Text":["Andrews & Thompson 2011"],"Citation Start End":[[714,737]]}
{"Identifier":"2019MNRAS.482.5222T__Bergeron_et_al._2011_Instance_1","Paragraph":"The atmospheric parameters of white dwarfs, the effective temperature (Teff) and surface gravity (log\u2009g), play a fundamental role in the study of post-main-sequence stellar evolution (see e.g., Kalirai, Marigo & Tremblay 2014; Rolland, Bergeron & Fontaine 2018), evolved planetary systems (Koester, G\u00e4nsicke & Farihi 2014), Galactic formation history (Tremblay et al. 2014), and the calibration of instruments (Bohlin, Gordon & Tremblay 2014). For decades, the most precise method to determine Teff and log\u2009g has been the spectroscopic fitting of the Balmer lines in hydrogen-atmosphere DA white dwarfs (Bergeron, Saffer & Liebert 1992; Finley, Koester & Basri 1997), and the He\u2009i lines in helium-dominated DB stars (Beauchamp et al. 1999; Voss et al. 2007; Bergeron et al. 2011). In contrast, trigonometric parallax measurements of stellar remnants or their companions can also be used to characterize Teff and stellar radii from photometric analyses (Koester, Schulz & Weidemann 1979; Bergeron, Leggett & Ruiz 2001). The surface gravity can then be inferred by using the white dwarf mass\u2013radius relation (see e.g., Fontaine, Brassard & Bergeron 2001), with a small dependence on the assumed internal stratification (Romero et al. 2012). The main advantage of the photometric technique is that it can be used to fit cool DC white dwarfs with no optical transitions as well as metal- or carbon-rich remnants where spectroscopic log\u2009g determinations are difficult (Dufour, Bergeron & Fontaine 2005; Dufour et al. 2007b). One shortcoming is that for $T_{\\rm eff} \\gtrapprox 12\\, 000$ K, optical colours of all stellar remnants get rather insensitive to the temperature, requiring very precise photometry for that determination (Carrasco et al. 2014). Until now, the main limitation of the photometric analyses has been that precise parallax measurements were available for only hundreds of white dwarfs (B\u00e9dard, Bergeron & Fontaine 2017), compared to the \u2248 30\u2009000 objects (Kleinman et al. 2013; Kepler et al. 2015, 2016; Gentile Fusillo et al. 2018) with spectroscopy from the Sloan Digital Sky Survey (SDSS; Abazajian et al. 2009).","Citation Text":["Bergeron et al. 2011"],"Citation Start End":[[758,778]]}
{"Identifier":"2021ApJ...923L..22A__Rosado_et_al._2015_Instance_1","Paragraph":"Pulsar timing experiments (Sazhin 1978; Detweiler 1979) allow us to explore the low-frequency (\u223c1\u2013100 nHz) part of the gravitational-wave (GW) spectrum. By measuring deviations from the expected arrival times of radio pulses from an array of millisecond pulsars, we can search for a variety of GW signals and their sources. The most promising sources in the nanohertz part of the GW spectrum are supermassive binary black holes (SMBHBs) that form via the mergers of massive galaxies. Orbiting SMBHBs produce a stochastic GW background (GWB; Lommen & Backer 2001; Jaffe & Backer 2003; Volonteri et al. 2003; Wyithe & Loeb 2003; Enoki et al. 2004; Sesana et al. 2008; McWilliams et al. 2012; Sesana 2013; Ravi et al. 2015; Rosado et al. 2015; Kelley et al. 2016; Sesana et al. 2016; Dvorkin & Barausse 2017; Kelley et al. 2017; Bonetti et al. 2018; Ryu et al. 2018), individual periodic signals or continuous waves (CWs; Sesana et al. 2009; Sesana & Vecchio 2010; Mingarelli et al. 2012; Roedig & Sesana 2012; Ravi et al. 2012, 2015; Rosado et al. 2015; Schutz & Ma 2016; Mingarelli et al. 2017; Kelley et al. 2018), and transient GW bursts (van Haasteren & Levin 2010; Cordes & Jenet 2012; Ravi et al. 2015; Madison et al. 2017; Islo et al. 2019; B\u00e9csy & Cornish 2021). We expect to detect the GWB first, followed by detection of individual SMBHBs (Siemens et al. 2013; Rosado et al. 2015; Taylor et al. 2016; Mingarelli et al. 2017) that stand out above the GWB. Detection of GWs from SMBHBs will yield insights into galaxy mergers and evolution not possible through any other means. Other potential sources in the nanohertz band include cosmic strings (Damour & Vilenkin 2000, 2001; Berezinsky et al. 2004; Damour & Vilenkin 2005; Siemens et al. 2006, 2007; \u00d6lmez et al. 2010; Sanidas et al. 2013; Blanco-Pillado et al. 2018; Chang & Cui 2021; Ghayour et al. 2021; Gorghetto et al. 2021; Wu et al. 2021a; Blanco-Pillado et al. 2021; Lin 2021; Chiang & Lu 2021; Lazarides et al. 2021; Chakrabortty et al. 2021; Ellis & Lewicki 2021), phase transitions in the early universe (Witten 1984; Caprini et al. 2010; Addazi et al. 2021; Arzoumanian et al. 2021; Di Bari et al.2021; Borah et al. 2021; Nakai et al. 2021; Brandenburg et al.2021; Neronov et al. 2021), and relic GWs from inflation (Starobinski\u01d0 1979; Allen 1988; Lazarides et al. 2021; Ashoorioon et al. 2021; Yi & Zhu 2021; Li et al. 2021; Poletti 2021; Vagnozzi 2021; Sharma 2021), all of which would provide unique insights into high-energy and early-universe physics.","Citation Text":["Rosado et al. 2015"],"Citation Start End":[[721,739]]}
{"Identifier":"2020MNRAS.493.2452C__Tomsick_et_al._2000_Instance_1","Paragraph":"Compact objects, such as, neutron stars and black holes (BHs) are the end products of massive stars. The gravitational pull of BHs is so high that no particle or radiation can escape from them. They both can be detected only by the electromagnetic radiation emitted by the accreted matter falling on them. Most of the black hole candidates (BHCs) in our Galaxy are in close binaries with companion stars which act as donors. Wind material from the companion star or matter accreting via Roche lobe overflow falls towards the black hole due to its intense gravitational pull and starts swirling around it to form accretion disc due to the presence of angular momentum. The observed electromagnetic radiation from a black hole binary is from the accretion flow itself. Some of the black hole candidates in low-mass X-ray binaries are transients in nature, which occasionally undergo outbursts. The increased X-ray flux in these transient X-ray binaries shows variability in the temporal and spectral states. Several works have already been done on this and many papers are available in the literature (see e.g. Tomsick et al. 2000; McClintock & Remillard 2006; Debnath et al. 2008; Nandi et al. 2012; Debnath, Chakrabarti & Nandi 2013; Rao 2013) to explain the variation of spectral and temporal properties of these objects during their active X-ray outburst phases. It is also reported by some authors that these objects transit through several spectral states (e.g. hard state (HS), hard-intermediate state (HIMS), soft-intermediate state (SIMS), soft state (SS)) during their active outburst phases (see Debnath et al. 2013 and references therein). Depending upon the observed spectral states, Debnath et al. (2017) classified these transient BH binaries into two types: classical or type-I (all four states are observable) and harder or type-II (softer states are missing). Low- and high-frequency quasi-periodic oscillations (QPOs) are generally observable in the power density spectra (PDS) of these sources (see Remillard & McClintock 2006 for a review). Low-frequency QPOs are commonly observed in hard and intermediate spectral states. Generally, it has been observed that during the rising HS and HIMS, frequency of these QPOs monotonically increases with time (day) and during the declining phase, the opposite nature is observed. In SIMS, these QPOs are observed sporadically.","Citation Text":["Tomsick et al. 2000"],"Citation Start End":[[1109,1128]]}
{"Identifier":"2022MNRAS.511.1121M__Reig_&_Nespoli_2013_Instance_2","Paragraph":"Critical luminosity (Lcrit) is the luminosity above which a state transition from subcritical to supercritical takes place. The subcritical state (LX Lcrit) is known to be the low luminosity state whereas the supercritical state is high luminosity state (LX > Lcrit) (Becker et al. 2012). The critical luminosity is crucial to determine whether the radiation pressure of the emitting plasma is capable of decelerating the accretion flow (Basko & Sunyaev 1976; Becker et al. 2012). The luminosity during the 2020 giant outburst reached a record high, which was significantly higher than the critical luminosity (Reig & Nespoli 2013). The source entered a supercritical regime from a subcritical regime during the outburst. In the supercritical regime, radiation pressure is high enough to stop the accreting matter at a distance above the neutron star, forming a radiation-dominated shock (Basko & Sunyaev 1976; Becker et al. 2012). For the subcritical regime, accreted material reaches the neutron star surface through nuclear collisions with atmospheric protons or through Coulomb collision with thermal electrons (Harding 1994). These accretion regimes can also be probed by noting changes in the cyclotron line energies, pulse profiles, and changes in the spectral shape (Parmar, White, & Stella 1989; Reig & Nespoli 2013). During the transition from the subcritical to the supercritical regime, sources show two different branches in their hardness\u2013intensity diagram (HID) which are known as horizontal branch (HB) and diagonal branch (DB) (Reig & Nespoli 2013). The HB implies the low-luminosity state of the source, which is represented by spectral changes and high X-ray variability. The DB corresponds to the high-luminosity state that appears when the X-ray luminosity is above the critical limit. The classification HB and DB depends on HID patterns that the source follows. The HB pattern is generally observed in the subcritical regime and the DB pattern is observed in the supercritical regime (Reig & Nespoli 2013).","Citation Text":["Reig & Nespoli 2013"],"Citation Start End":[[1305,1324]]}
{"Identifier":"2017AandA...606A..50D__Mennella_et_al._(1998)_Instance_1","Paragraph":"With this study, our group continues the effort to investigate the optical properties of cosmic dust analogues in the mid infrared (MIR) to the millimeter domain as a function of temperature, undertaken 20 years ago by different groups (we refer to Demyk et al. 2013, for the details of the studied samples). Briefly, Agladze et al. (1996) were the first to study relevant interstellar silicate dust analogues in the temperature range from 1.2 K to 30 K and in the wavelength range from 700 \u03bcm to 2.9 mm. Mennella et al. (1998) studied amorphous carbon samples and silicate samples in the 24 \u03bcm\u22122 mm spectral domain and 24\u2212300 K temperature range. Boudet et al. (2005) investigated silica and silicate samples in the 10\u2212300 K temperature range and in the spectral region 100\u22121000 \u03bcm. The work by Coupeaud et al. (2011) was focussed on pure sol-gel Mg-rich silicates, of composition close to enstatite (MgSiO3) and forsterite (Mg2SiO4), amorphous and crystalline, whose spectra were recorded in the 100\u22121000 \u03bcm spectral range and from 10 K to 300 K. These studies have brought important results about the spectroscopic characteristics and behavior of interstellar dust analogues in the FIR at varying temperature. They show that the MAC of the amorphous analogues (silicates and carbonaceous matter) increases with the grain temperature and that its spectral shape cannot be approximated with a single power law in the form \u03bb\u2212 \u03b2. The dependence of the MAC on the temperature, which is not observed in crystalline samples, is related to the amorphous nature and to the amount of defects in the structure of the material (Coupeaud et al. 2011). A physical model was proposed by Meny et al. (2007) to explain these experimental results. This model, named the TLS model, is based on a description of the amorphous structure of the material in terms of a temperature independent disordered charge distribution and of a collection of atomic configurations modeled as two-level systems and sensitive to the temperature. ","Citation Text":["Mennella et al. (1998)"],"Citation Start End":[[505,527]]}
{"Identifier":"2019AandA...625A.114J__Tacconi_et_al._2018_Instance_2","Paragraph":"Although most galaxies have an implied SFR that scatters within a factor two around the MS, some do show a significantly higher SFR. Those objects also exhibit a higher gas content, shorter gas depletion times (e.g., Genzel et al. 2015; Schinnerer et al. 2016; Tacconi et al. 2013, 2018), and higher dust temperatures (e.g., Magnelli et al. 2014). Likewise, the stellar-light radial distribution is different in these two galaxy populations; while MS galaxies are closely approximated by exponential disks (e.g., Bremer et al. 2018), those above (and below) it exhibit a higher central mass concentration (e.g., Wuyts et al. 2011). Based on this dichotomy and the parametrization of the MS over cosmic time, a scenario has been proposed to explain the evolutionary path of galaxies along the MS. Since the normalization of the MS, the gas fraction of galaxies, and cosmic molecular gas density decrease from z\u2004\u223c\u20042.5 to 0 at a similar pace (e.g., Speagle et al. 2014; Decarli et al. 2016; Tacconi et al. 2018), it is thought that MS galaxies evolved through a steady mode of star formation, possibly regulated by the accretion of cool gas from the intergalactic medium (e.g., Dekel et al. 2009; Kere\u0161 et al. 2009; Dav\u00e9 et al. 2010; Hodge et al. 2012; Romano-D\u00edaz et al. 2014, 2017; Feng et al. 2015; Angl\u00e9s-Alc\u00e1zar et al. 2017). From theoretical predictions, the scatter of the MS could thus be explained as the result of a fluctuating gas inflow rate that is different in each galaxy (e.g., Tacchella et al. 2016; Mitra et al. 2017). In this context, a galaxy enhances its SFR and moves towards the upper envelope of the MS due to gas compaction. As the gas is depleted, the SFR decreases and the galaxy falls below the MS. As long as a SFG is replenished with fresh gas within a timescale shorter than its depletion time, it will be confined within the MS (Tacchella et al. 2016). On the other hand, the enhanced star formation efficiency of galaxies above the MS has been linked to mergers (e.g., Walter et al. 2009; Narayanan et al. 2010; Hayward et al. 2011; Alaghband-Zadeh et al. 2012; Riechers et al. 2013, 2014) and instability episodes in gas-rich disks (particularly at high redshift; e.g., Dav\u00e9 et al. 2010; Hodge et al. 2012; Wang et al. 2019).","Citation Text":["Tacconi et al. 2018"],"Citation Start End":[[988,1007]]}
{"Identifier":"2017MNRAS.466.3961S__Bogd\u00e1n_et_al._2013a_Instance_1","Paragraph":"Other important components of the X-ray sky are X-ray binaries, the hot gas present in our own Galaxy and the extragalactic hot gas. Amongst these X-ray sources, the large hot gas reservoir (Tvir \u223c 107\u2009K) filling the space between the galaxies in the clusters, known as the ICM, has been observed in its X-ray emission for a long time (Reichert et al. 1981; Jones & Forman 1984; Branduardi-Raymont et al. 1985; Oukbir, Bartlett & Blanchard 1997; Diego et al. 2003a; Diego, Silk & Sliwa 2003b; Cavagnolo et al. 2009; Hurier et al. 2015). However, the hot gas (Tvir > 106\u2009K) present in the form of circumgalactic medium (CGM) in massive galaxies (Mh \u223c 1012\u20131013\u2009h\u22121\u2009M\u2299; Birnboim & Dekel 2003; Kere\u0161 et al. 2005; Singh et al. 2015) is less explored in X-rays due to its fainter X-ray emission. Some of the recent observations and studies (Grcevich & Putman 2009; Anderson & Bregman 2011; Dai et al. 2012; Putman, Peek & Joung 2012; Anderson, Bregman & Dai 2013; Bogd\u00e1n et al. 2013a,b; Gatto et al. 2013) indicate that the CGM can account for a good fraction of the baryons in these galaxies. The X-ray emission from the CGM, therefore, is a promising tool to put strong constraints on the distribution and energetics of the gas (Singh et al. 2016) with eROSITA. At energies above 2 keV, the extragalactic point sources like AGN completely dominate the X-ray sky (Lehmann et al. 2001; Kim et al. 2007). Even below 2 keV, where the X-ray emission from the hot gas in the ICM and CGM is significant, the major contribution to the observed X-ray sky comes from AGNs (So\u0142tan 2007). Therefore, studying the X-ray emission from the AGN is crucial to understand the origin and evolution of the AGN as well as to extract the X-ray signal from the subdominant components. We thus also compute the angular power spectrum of the unresolved AGNs that are expected to contribute to the diffuse X-ray background of eROSITA and contaminate the angular power spectrum due the ICM\/CGM in the 0.5\u20132 keV X-ray band.","Citation Text":["Bogd\u00e1n et al. 2013a"],"Citation Start End":[[959,978]]}
{"Identifier":"2019MNRAS.488.5029H__Stacey_et_al._2010_Instance_4","Paragraph":"For the first time, we detected [C\u2009ii]\u2009158-\u03bcm emission from a GRB host galaxy at z > 2. This is the second detection of [C\u2009ii]\u2009158-\u03bcm emission among known GRB host galaxies, following GRB 980425 (Micha\u0142owski et al. 2016). The [C\u2009ii]\u2009158-\u03bcm fine structure line is the dominant cooling line of the cool interstellar medium, arising from photodissociation regions (PDR) on molecular cloud surfaces. It is one of the brightest emission lines from star-forming galaxies from FIR to metre wavelengths, almost unaffected by dust extinction. [C\u2009ii]\u2009158-\u03bcm luminosity, L[C\u2009II], has been discussed as an indicator of SFR (e.g. Stacey et al. 2010). If L[C\u2009II] scales linearly with SFR, the ratio to FIR luminosity, L[C\u2009II]\/LFIR, is expected to be constant, since LFIR is a linear function of SFR (e.g. Kennicutt 1998a). However, LC\u2009II\/LFIR is not constant, but declines with increasing LFIR, known as the \u2018[C\u2009ii] deficit\u2019 (e.g. Luhman et al. 1998, 2003; Malhotra et al. 2001; Sargsyan et al. 2012; D\u00edaz-Santos et al. 2013, 2017; Spilker et al. 2016). The [C\u2009ii] deficit persists when including high-z galaxies (e.g. Stacey et al. 2010; Wang et al. 2013; Rawle et al. 2014). In Fig. 5, we compare the [C\u2009ii] deficit in the GRB 080207 host and other star-forming galaxies. Two GRB hosts are shown by stars: GRB 080207 (orange star) and 980425 (blue star). The comparison sample is compiled from the literature up to z \u223c 3 (Malhotra et al. 2001; Cormier et al. 2010, 2014; Ivison et al. 2010; Stacey et al. 2010; Sargsyan et al. 2012; Farrah et al. 2013; Magdis et al. 2014; Brisbin et al. 2015; Gullberg et al. 2015; Schaerer et al. 2015). Active galactic nuclei are separated from star-forming galaxies based on either (i) the explicit description in the literature or (ii) EWPAH\u20096.2\u03bcm  0.1 (Sargsyan et al. 2012). As reported by previous studies (e.g. Maiolino et al. 2009; Stacey et al. 2010), high-z galaxies are located at a different place from local galaxies in the L[C\u2009II]\/LFIR\u2013LFIR plane.","Citation Text":["Stacey et al. 2010"],"Citation Start End":[[1863,1881]]}
{"Identifier":"2017ApJ...847...42D__Purcell_et_al._2011_Instance_2","Paragraph":"The detailed kinematic reconstruction of the Sgr tidal debris by Law & Majewski (2010) used an initial total mass of \n\n\n\n\n\n M\u2299 for the Sgr satellite. However, several studies point to a Sgr remnant mass significantly exceeding that value. Ibata et al. (1997) and Ibata & Lewis (1998) estimate lower bounds of 109 M\u2299 for the mass of the dwarf today. With the discovery of previously unseen branches of the stream, the total luminosity budget of the progenitor galaxy is now believed to be on the order of 108 L\u2299 (Niederste-Ostholt et al. 2010). As a result, recent studies have shifted to using dark-matter halo masses as large as 1011 M\u2299 (Purcell et al. 2011; G\u00f2mez et al. 2015) based on halo abundance matching arguments. Such high values are comparable to the mass of the LMC progenitor (see e.g., Jethwa et al. 2016; Pe\u00f1arrubia et al. 2016) and imply a mass ratio relative to the MW on the order of 1:10. However, unlike the Magellanic Clouds, which may be on their first passage near the MW (Besla et al. 2007), Sgr is known to have experienced multiple close passages in the past. If true, such high Sgr progenitor masses would have important implications for the formation and evolution of the MW disk (e.g., Purcell et al. 2011; G\u00f2mez et al. 2013; D\u2019Onghia et al. 2016). Because the dynamical friction force is proportional to the square of the satellite mass, we expect drag to play a much more important role in slowing down the Sgr satellite and bringing it to closer galactocentric distances. Here, we perform an exploration across orbital angular momentum parameter space analogous to Section 3.1, this time using a Sgr progenitor mass of 6\u00d71010 M\u2299 according to the recent estimates of Gibbons et al. (2017). We consider two possibilities: a \u201cslow sinking\u201d scenario in which, as in Section 3, Sgr crosses the MW virial radius at \n\n\n\n\n\n (approximately 8 Gyr ago), and a \u201crapid sinking\u201d scenario, in which we examine a first infall around \n\n\n\n\n\n, about 4 Gyr ago.","Citation Text":["Purcell et al. 2011"],"Citation Start End":[[1215,1234]]}
{"Identifier":"2019ApJ...887..118C__Dud\u00edk_et_al._2016_Instance_1","Paragraph":"For many hot channel eruption events, their exact origin and initiation process still remain under debate. Some observations supported that before the hot channel eruption, a corresponding MFR may already exist (e.g., Liu et al. 2016; Wang et al. 2016; Yang et al. 2019) and slow photospheric flows might be important for its gradual formation (e.g., Hou et al. 2018; Yan et al. 2018; Vasantharaju et al. 2019); while others argued that they can also be newly built up via rapid coronal reconnection during the eruption (Cheng et al. 2011; Song et al. 2014; Palacios et al. 2015). For their loss-of-equilibrium, some researchers found that the loss-of-equilibrium of the hot channel is initially facilitated by the breakout type (e.g., Shen et al. 2012; Chen et al. 2016; Mitra et al. 2018) or tether-cutting type (e.g., Chen et al. 2014, 2018; Joshi et al. 2015; Dud\u00edk et al. 2016) preflare reconnection, but others believed the loss-of-equilibrium of the hot channel may be directly triggered by the ideal MHD instabilities (e.g., Liu et al. 2007; Bi et al. 2015; Vemareddy et al. 2017). Note that Duan et al. (2019) recently carried out a survey on the initiation mechanism of all the major solar flares, either eruptive or confined, from 2011 to 2017. Statistically, their study results indicate that magnetic reconnection and ideal MHD instabilities seemingly play an equal role in the triggering of major flares and eruptions. In the present study, our NLFFF extrapolation and topology analysis shows that a seed MFR existed within an HFT configuration prior to the later solar eruption. Near the HFT, two sets of sheared arcades serve as an enveloping field holding the seed MFR, and have rooted their opposite footpoints at canceling fluxes. Combining these results with imaging observations, we suggest that the preflare 3D reconnection is most likely first triggered at the HFT in the precursor phase due to the preceding flux cancellation (Savcheva et al. 2012a), and the newborn hot channel should originate from the preexisting seed MFR. Through preflare 3D reconnection, the seed MFR was rapidly heated up to 10 Mk manifesting as a newborn hot channel, and its spatial size also underwent an enlarging process at the same time (as revealed by Figures 2 and 5). As a result, the enlarged hot channel may readily lift into the abovementioned torus-unstable domain (around 12 Mm above the photosphere), and soon enter its exponential accelerated eruption.","Citation Text":["Dud\u00edk et al. 2016"],"Citation Start End":[[864,881]]}
{"Identifier":"2022ApJ...926...85S__Ehrenreich_et_al._2020_Instance_1","Paragraph":"As in Flowers et al. (2019), to compare our model transmission spectra directly against the Ehrenreich et al. (2020) results, we must calculate the transmission spectra as a function of orbital phase throughout the duration of transit. To account for orbital phase dependencies, we apply the following procedure:1.Account for phase-dependent backlighting of the planet (i.e., stellar limb-darkening effects). At different points of its transit, the planet will occult regions of its host star of varying brightness. Furthermore, at a fixed orbital phase, different regions of the planet\u2019s limb will be backlit by varying intensities of stellar light. Similar to Flowers et al. (2019), we calculate the normalized stellar intensity at the center of each cell of the 2D projected planetary grid produced by our GCM at each modeled orbital phase of the planet. We use the quadratic limb-darkening coefficients reported by Ehrenreich et al. (2020) to establish the stellar center-to-limb intensity profile, and we take into account the 89.\u00b0623 orbital inclination of WASP-76b (Ehrenreich et al. 2020) to determine where the planet resides on the stellar disk as a function of its orbital phase. We make the assumption of constant impact parameter b over the course of transit.\n9\n\n\n9\nIn reality, a planet on an inclined orbit will not have a constant b over the entire duration of transit; rather, the planet\u2019s distance from the stellar equator will be decreased at ingress and egress, reaching its maximum at center of transit. Our tests reveal that, for WASP-76b, the relative error induced by the constant b assumption is on the order of 4% in distance, which results in a change on the order of 1 m s\u22121 at the blueshift level (see Section 3.1). Hence, our b treatment is justified. This procedure allows us to calculate a backlighting factor f, which ranges from 0 to 1, effectively replacing the constant I\n\n\u03bb,0 from Equation (3) with a variable \n\n\n\nI\u03bb,0\u00d7f(\u03b8\u2032,z,\u03c6)\n\n, for a given orbital phase \u03c6 and 2D projected polar angle \n\n\n\n\u03b8\u2032\n\n.2.Account for the decreasing of the continuum by interpolating a light curve produced by the batman code (Kreidberg 2015). Step 1 ensures that less light is transmitted through the planet\u2019s atmosphere than a uniform stellar disk would emit. Step 2 further enforces that the inner, optically thick core of the planet is simulated crossing a limb-darkened star, as opposed to a star of uniform brightness.3.Account for the planet\u2019s rotation over the course of transit. Because the planet is continually rotating as it travels across the face of its host star, we must transform the GCM coordinate system so that the correct observer-facing hemisphere is modeled at each instance during transit. For simplicity, we assume zero obliquity, which allows us to calculate the coordinate transform simply by assigning a linear offset to each planetary longitude; i.e., \u03d5\nrotated = \u03d5 + \u03c6.\n","Citation Text":["Ehrenreich et al. (2020)"],"Citation Start End":[[92,116]]}
{"Identifier":"2019MNRAS.487.4571B__Krause_2005_Instance_1","Paragraph":"The origin of ultra-high energy cosmic rays (UHECR) is uncertain, and many possible sources have been proposed. Radio galaxies have long been considered a likely source of UHECR because of their high power, large size, and longevity (e.g. Hardcastle et al. 2009; O\u2019Sullivan, Reville & Taylor 2009; Wykes et al. 2013; Eichmann et al. 2018). In Matthews et al. (2018, 2019), we made the case from observations and theory that UHECR may be accelerated by shocks in the lobes of radio galaxies. Our numerical simulations demonstrated the presence of shocks in concentrated, approximately annular, flows centred on the jet axis, which we refer to as hydromagnetic flux tubes, or more simply as \u2018flux tubes\u2019, that emerge from the high pressure hot-spots at the end of relativistic jets. Back-flows in radio lobes have been discussed by many authors (for example Norman et al. 1982; Falle 1991; Scheuer 1995; Saxton et al. 2002, Krause 2005; Keppens et al. 2008; Mignone et al. 2010; Cielo et al. 2014; Tchekhovskoy & Bromberg 2016). The flow expands into the lobes where the pressure is lower. As discussed in Matthews et al. (2019), from application of the Bernoulli equation, the flow becomes supersonic before being slowed down in one or more shocks as it progresses deeper into the lobe as depicted in an idealized model in Fig. 1. Our simulations of radio jets support this picture by demonstrating the presence of a succession of shocks in flux tubes with Mach numbers of a few and flow velocities of the order of c\/4. With these velocities, shocks are well suited to UHECR acceleration since fully relativistic shocks pose severe problems for acceleration to the highest energies (Kirk & Reville 2010; Lemoine & Pelletier 2010; Reville & Bell 2014; Bell et al. 2018). Further discussion of our model and the observational context can be found in Matthews et al. (2018, 2019). Simulations show that flux tubes are usually more contorted than the idealized flux tubes in Fig. 1. However, the flux tube need only be long enough to contain the CR diffusion scaleheights upstream and downstream of a shock, and curvature should not invalidate the model since magnetic field lines threading the tube should cause CR to diffuse predominantly parallel to the tube. Cosmic ray (CR) acceleration by diffusive shock acceleration (Krymskii 1977; Axford, Leer & Skadron 1977; Bell 1978a,b, Blandford & Ostriker 1978) has been studied at length, both theoretically and observationally, in the context of supernova remnants (SNR). The maximum energy of CR accelerated by SNR is constrained by various factors. The most general constraint is set by the physical size R of the shock, leading to a maximum CR energy of ZuBR (Hillas 1984), where u is the flow velocity, B is the magnetic field, and the CR has charge Ze. ZuBR is the characteristic maximum CR energy in eV if u, B & R are in SI units. A related constraint is set by the need for the shock to persist for a time longer than the time taken for CR to be accelerated (Lagage & Cesarsky 1983a,b). For a quasi-parallel shock (magnetic field parallel to or at an angle significantly smaller than \u03c0\/2 to the shock normal) the Lagage & Cesarsky limit is equivalent to the Hillas limit if the shock persists for a flow time R\/u. For perpendicular shocks, the acceleration is very rapid and the maximum CR energy is set by the Hillas limit (Jokipii 1982, 1987). In this work, we assume that the hydrodynamic flow through a flux tube is in steady state on the scale of the time taken for CR acceleration, in which case the limit on the maximum CR energy is better understood through the spatial limit of Hillas than the time-scale analysis of Lagage & Cesarsky.","Citation Text":["Krause 2005"],"Citation Start End":[[922,933]]}
{"Identifier":"2021MNRAS.508...79J___2004_Instance_1","Paragraph":"The physical parameters governing the MS are still being investigated, although some basic inferences have been made. The main parameters that describe quasars as accreting black holes are the black hole mass (MBH, ranging from 106$\\rm M_{\\rm \\odot }$ to 109.5$\\rm M_{\\rm \\odot }$; Kormendy & Richstone 1995), the accretion luminosity, Eddington ratio (LBol\/LEdd, Boroson & Green 1992; Sulentic et al. 2000b; Marziani et al. 2001, 2003b), and the black hole spin (e.g. Wang et al. 2014). The MBH can be estimated by assuming that the gas motions are predominantly Keplerian around the black hole, and by applying the virial theorem for a system whose mass is entirely concentrated in the center of gravity. Reverberation mapping provides a measurement of the radial distance rBLR of the line emitting gas from the central black hole (Peterson 1993, 2004). The so-called \u2018virial mass\u2019 have been computed for large samples of quasars employing several different emission lines, various measures of line width, and exploiting correlation between the emitting region radius and luminosity (see Marziani & Sulentic 2012; Shen 2013; Popovi\u0107 2020, for reviews). An independent method relies on the scaling law between MBH and the stellar velocity dispersion of the galaxy bulge (\u03c3\u22c6; Kormendy & Ho 2013). Recently, spectropolarimetric observations of the Hydrogen Balmer lines have allowed to measure the intrinsic line FWHM due to a Keplerian velocity field and to compute the black hole mass independently from orientation (Afanasiev & Popovi\u0107 2015; Afanasiev, Popovi\u0107 & Shapovalova 2019, and references therein). All methods to compute black hole mass in active galactic nucleus (AGN) are subject to caveats and suffer considerable uncertainties (e.g. Dalla Bont\u00e0 et al. 2020, and references therein). None the less, the LBol\/LEdd, which is proportional to the luminosity-to-black hole mass ratio (L\/MBH), has been revealed to be a fundamental driver of the E1 MS, closely related to several observational parameters (Marziani et al. 2001, 2003b; Kuraszkiewicz et al. 2004; Shen & Ho 2014; Panda et al. 2018; Panda, Marziani & Czerny 2019). With decreasing Eddington ratio, source properties change from the ones of Pop. A to the ones of Pop. B, which tend to show broader Balmer emission lines, weaker RFeII,1 and more asymmetric Balmer line profiles (Sulentic et al. 2000b; Marziani et al. 2003b; Shen & Ho 2014).","Citation Text":["Peterson","2004"],"Citation Start End":[[834,842],[849,853]]}
{"Identifier":"2022MNRAS.517.4119T__Matheson_et_al._2012_Instance_1","Paragraph":"SN 2011fe, discovered a mere \u224811 h after explosion (Nugent et al. 2011) by the Palomar Transient Facility (PTF; Law et al. 2009), is the brightest SN Ia since the advent of modern astronomical detectors. Located at just $d_L \\approx 6.5~\\rm {Mpc}$ (e.g. Shappee & Stanek 2011; Beaton et al. 2019), SN 2011fe exploded in a region of M101 uncontaminated by intervening dust (Patat et al. 2013) providing an ideal location for testing SN Ia progenitor and explosion models. The early detection allowed extensive follow-up observations across the electromagnetic spectrum (e.g. Matheson et al. 2012; Parrent et al. 2012; Pereira et al. 2013; Hsiao et al. 2013; Johansson, Amanullah & Goobar 2013; Tsvetkov et al. 2013; Munari et al. 2013; Mazzali et al. 2014; Zhang et al. 2016) and provided direct constraints on the radius of the exploding star (Nugent et al. 2011; Bloom et al. 2012). Stringent non-detections in radio (Chomiuk et al. 2012; Horesh et al. 2012; Kundu et al. 2017) and X-ray (Horesh et al. 2012; Margutti et al. 2012) observations exclude nearby CSM at high significance. Early ultraviolet (UV) photometry did not show any evidence for the ejecta encountering a nearby companion star (Brown et al. 2012) and nebular spectra lacked the Balmer emission lines from material ablated off the donor star by the ejecta impact (Shappee et al. 2013b; Lundqvist et al. 2015; Tucker et al. 2022). Pre-explosion imaging excludes the presence of a RG or He donor star (Li et al. 2011) and disfavor an accreting WD in the \u223c105 yr prior to explosion (Graur, Maoz & Shara 2014). Multi-epoch spectropolarimetry reveal consistently-low continuum polarization suggestive of a symmetric ejecta distribution with evidence for minor oblateness (Milne et al. 2017). Finally, nebular-phase observations at \u2273 1 year after maximum light allows a direct view to the inner ejecta and provides unique constraints on the explosion conditions (McClelland et al. 2013; Kerzendorf et al. 2014, 2017; Mazzali et al. 2015; Graham et al. 2015; Taubenberger et al. 2015; Dimitriadis et al. 2017; Friesen et al. 2017; Shappee et al. 2017; Tucker et al. 2022). SN 2011fe is one of the best-studied astronomical objects in the past decade and remains a key benchmark for any SN Ia theory or model.","Citation Text":["Matheson et al. 2012"],"Citation Start End":[[574,594]]}
{"Identifier":"2021MNRAS.502.3761L__Demircan_&_Kahraman_1991_Instance_1","Paragraph":"A star having highly eccentric and compact orbit around a black hole experiences tidal distortion. The tidal radii of the black hole\u2013star system are expressed as (Shiokawa et al. 2015; Coughlin et al. 2017)\n(7)$$\\begin{eqnarray*}\r\nr_t = (\\frac{M_{\\rm BH}}{M})^{1\/3} R,\r\n\\end{eqnarray*}$$where MBH is the mass of the black hole, M is the mass of the star, and R is the radius of the star. Stars orbiting at a distance smaller than the tidal radius are disrupted by tidal force of the black hole. The additional contribution to perihelion shift arising from tidal effect in terms of pericentre distance is given by (Will 2008)\n(8)$$\\begin{eqnarray*}\r\n\\delta \\phi _{\\rm prec}^{\\rm tidal} = \\frac{30 \\pi }{(1+e)^5} k_2 (1+ \\frac{3e^2}{2} + \\frac{e^4}{8})(\\frac{M_{\\rm BH}}{M}) (\\frac{{R}^5}{ r_{\\rm p}^5}),\r\n\\end{eqnarray*}$$k2 in equations (8) is the Love number of a star that measures its deformability and susceptibility of its shape to change in response to a tidal potential (Poisson & Will 2014). The tidal Love number (kl) of degree l for polytropic stars has been studied by Yip & Leung (2017) through perturbative expansion of kl around the polytropic indices n = 0 and n  = 1. These are obtained for Newtonian stars whose stiffness is governed by polytropic index (n). In the general framework of Yip & Leung (2017), the Lane Emden equation is solved perturbatively by applying Scaled Delta Expansion method to develop perturbation series of the multipolar tidal Love number about two exactly solvable polytropic indices (n = 0, n  = 1). These perturbative expansions are then used to make a two-point Pade approximation of the Love number for n \u03f5 [0, 3] which is given as\n(9)$$\\begin{eqnarray*}\r\nk_l({\\it n}) = (5 - {\\it n})^3 \\frac{a_1 + {\\it n}a_2 + {\\it n}^2 a_3 + {\\it n}^3 a_4}{1000 + {\\it n}a_5 + {\\it n}^2 a_6}.\r\n\\end{eqnarray*}$$The factor (5 \u2212 n) is empirically inserted to ensure that the Love number reduces to zero for n  = 5 that corresponds to infinite configuration. Here, ais are analytically determined coefficients of the perturbation series and these are taken from table 1 of Yip & Leung (2017). We consider polytropic models, n = 1 and n  = 3, for evaluating the Love number and obtain k2 (n = 1)  = 0.26 k2 (n = 3)  = 0.014. These values are then used to evaluate tidal contribution to the periastron shift. For the tidal effect, we also have to take into account the mass and radii of the stars. Structure of the nuclear star cluster in the vicinity of Sgr A* was reported by Do et al. (2010). A general consensus is that the cluster has preponderance of young main sequence stars (\u223c 4 \u2013 8Myr) over late-type stars (\u223c1 \u2013 10Gyr). Presence of OB main-sequence stars was also reported by Allen, David & Hyland (2010), Ghez et al. (2003a), and Do et al. (2013). Here, we consider only the main-sequence stars for the mass\u2013radius relation that is required for calculation of tidal contribution to periastron shift (see equation 8). Study of empirical mass\u2013radius relation (MRR) of stars began to appear in literature after late twentieth century (Gimenez & Zamorano 1985; Harmanec 1988; Demircan & Kahraman 1991; Malkov 2007). The stellar radii are related to luminosity through Stefan\u2013Boltzmann law which in turn is dependent on stellar masses. Unfortunately, there was no final consensus on what functions properly represent mass\u2013luminosity relation (MLR) and MRR. It had been challenging to have a single MRR for main-sequence stars till the publication of Eker et al. (2015) in which the MLR and MRR for Galactic nearby main-sequence stars were reported. The MLR was revised with the help of power law L \u221d M\u03b1 for a sample of galactic main-sequence stars with mass and radii accurate up to \u22643 per\u2009cent and luminosities accurate up to \u226430  per\u2009cent. But there was no derivation of MRR to estimate radii for a given stellar mass. It was found that for LMSs (M  1M\u2299) stellar radii have narrow distribution with mass but for massive stars (M > 1M\u2299) the distribution was quite wide. Thus, there was no unique function to represent MRR. Eker et al. (2018) updated the data of 509 main-sequence stars in the solar neighbourhood and obtained interrelated MLR, MRR, and mass-effective temperature relation (MTR) in the mass range 0.179 \u2264 (M\/M\u2299) \u2264 31. A single MRR was found for the LMSs with 0.179 \u2264 (M\/M\u2299) \u2264 1.5. An MTR function was found for massive stars with 1.5 \u2264 (M\/M\u2299) \u2264 31. The MLR was established for six mass ranges encompassing ultra-low-mass (0.179 \u2264 (M\/M\u2299) \u2264 0.45) and very high-mass (7 \u2264 (M\/M\u2299) \u2264 31) stars. Thus, by using Stefan\u2013Boltzmann law (L  = 4$\\pi R^2 \\sigma T_{\\rm eff}^4$), the MRR can be obtained for the massive stars with 1.5 \u2264 (M\/M\u2299) \u2264 31. It is true that there are effects of metallicities of the stars on the MLR and hence finally on MRR (Eker et al. 2018). But due to absence of detailed model of stellar evolution and stellar atmosphere of the GC stars with different metallicities, we have assumed that the young GC stars have nearly the solar metallicities (see Lu et al. 2013 and references therein) and hence obey the interrelated MLR, MTR, and MRR functions of the main-sequence stars in the solar neighbourhood. The MRR for the LMSs [0.179 \u2264 (M\/M\u2299) \u2264 1.5] is (Eker et al. 2018)\n(10)$$\\begin{eqnarray*}\r\n\\frac{R}{{\\rm R}_{\\odot }} \\cong [0.438(\\frac{M}{{\\rm M}_{\\odot }})^2 + 0.479 (\\frac{M}{{\\rm M}_{\\odot }}) + 0.075].\r\n\\end{eqnarray*}$$For massive stars (1.5 \u2264 (M\/M\u2299) \u2264 31), we use the interconnection between MLR, MTR, and MRR presented by Eker et al. (2018) as follows. The Stefan\u2013Boltzmann law\n(11)$$\\begin{eqnarray*}\r\n{\\rm log}(\\frac{L}{{\\rm L}_{\\odot }}) = 2 {\\rm log}(\\frac{R}{{\\rm R}_{\\odot })} + 4{\\rm log}(\\frac{T_{\\rm eff}}{T_{{\\rm eff}(\\odot)}})\r\n\\end{eqnarray*}$$is used in conjunction with the M\u2013Teff relation (MTR) for 1.5 \u2264 (M\/M\u2299) \u2264 31\n(12)$$\\begin{eqnarray*}\r\n{\\rm log}(\\frac{T_{\\rm eff}}{1K})\\simeq - 0.170({\\rm log}(\\frac{M}{{\\rm M}_{\\odot }}))^2 + 0.888 {\\rm log}(\\frac{M}{M_{\\odot }}) + 3.671\r\n\\end{eqnarray*}$$and the MLR for the mass range 7 \u2264 (M\/M\u2299) \u2264 31,\n(13)$$\\begin{eqnarray*}\r\n{{\\rm log}(\\frac{L}{{\\rm L}_{\\odot }})\\simeq 2.865 {\\rm log}(\\frac{M}{{\\rm M}_{\\odot }}) + 1.105}.\r\n\\end{eqnarray*}$$The result is the following MRR for massive stars,\n(14)$$\\begin{eqnarray*}\r\n{\\rm log}\\Big (\\frac{R}{{\\rm R}_{\\odot }}\\Big) \\cong [-0.3435 {\\rm log}\\Big (\\frac{M}{{\\rm M}_{\\odot }}\\Big) + 0.34 {{\\rm log}\\Big (\\frac{M}{{\\rm M}_{\\odot }}\\Big) }^2 + 0.7365]. \\nonumber \\\\\r\n\\end{eqnarray*}$$We study the tidal contribution to periastron shift for both the low-mass (0.179 \u2264 (M\/M\u2299) \u2264 1.5) and high-mass (7 \u2264 (M\/M\u2299) \u2264 31) categories of young main-sequence stars near Sgr A*. For low-mass category, the contribution is studied for M\/M\u2299 = 0.18, 1.0, and 1.5, whereas for high-mass category it is estimated for M\/M\u2299 = 10, 20, and 30. Equation (8) can be written as\n(15)$$\\begin{eqnarray*}\r\n\\delta \\phi _{\\rm prec}^{\\rm tidal} = \\frac{30 \\pi }{(1+e)^5} k_2 (1+ \\frac{3e^2}{2} + \\frac{e^4}{8})(\\frac{M_{\\rm BH}}{M}) (\\frac{{{\\rm R}_{\\odot }}}{ r_p})^5 (\\frac{R}{{\\rm R}_{\\odot }})^5.\r\n\\end{eqnarray*}$$Equations (10) and (14) have been used in equation (15) to estimate the tidal contributions to periastron shift arising from LMS and HMS. The actual astrometric size of the perihelion shift at periapsis is given by \u03b8  = ($\\delta \\phi ^{\\rm total}_{\\rm prec}$) ($\\frac{ r_{p} \\sin i}{D}$), where D is the distance to the GC and i is the inclination of the orbital plane of the star with respect to the observer. Therefore, the actual rate of periastron shift is expressed as\n(16)$$\\begin{eqnarray*}\r\n\\dot{\\theta }_{\\rm prec} = \\frac{(\\delta \\phi)^{\\rm total}_{\\rm prec} r_{p} \\sin i}{\\rm PD}.\r\n\\end{eqnarray*}$$Here, P is the orbital period for the stars, P  = ${ 2\\pi a^{\\frac{3}{2}}}\/{\\sqrt{GM_{\\rm BH}}}$  = ${2\\pi r_p^{\\frac{3}{2}}}\/{\\sqrt{GM_{\\rm BH}}(1-e)^{\\frac{3}{2}}}$.","Citation Text":["Demircan & Kahraman 1991"],"Citation Start End":[[3113,3137]]}
{"Identifier":"2018AandA...619A..59S__Owocki_et_al._1988_Instance_1","Paragraph":"Information about fundamental parameters of stars \u2013 like their mass, luminosity, surface temperature and chemical composition \u2013 comes primarily from matching observations to synthetic spectra computed using models of stellar atmospheres. For massive stars with hot surfaces, scattering and absorption in spectral lines further transfer momentum from the star\u2019s intense radiation field to the plasma, and so provides a force that overcomes gravity and drives a wind-outflow directly from the stellar surface (Castor et al. 1975). These starlight-powered winds are very strong and fast, and can dramatically affect the star\u2019s atmospheric structure (review by Puls et al. 2008) as well as the evolution of its mass and luminosity, chemical surface abundances, rotational velocity, and nuclear burning life-times (review by Smith 2014). Atmospheric models of such hot, massive stars must thus generally be constructed using a unified, or global, approach, wherein the basic structural equations for the quasi-static photosphere and the outflowing stellar wind are solved simultaneously (Gabler et al. 1989). In addition, the expanding atmospheres of these stars are characterized by their large departures from local thermodynamic equilibrium (LTE), meaning the full number density rate equations (typically reduced to statistical equilibrium, and often simply called non-LTE or NLTE) must be solved to obtain the atmospheric radiation field and the excitation and ionization balance. As such, quite intricate numerical solution techniques are normally required to compute synthetic observables, like spectral lines and energy distributions, for these objects (for details, see book by Hubeny & Mihalas 2014). Over the past decades, much effort has been devoted toward constructing such global NLTE, steady-state model atmospheres of hot stars with winds; several numerical computer codes meanwhile exist on the market, for example CMFGEN (Hillier & Miller 1998), POWR (Gr\u00e4fener et al. 2002; Sander et al. 2015), PHOENIX (Hauschildt 1992), WM-BASIC (Pauldrach et al. 2001), and the subject of this paper, FASTWIND (Santolaya-Rey et al. 1997; Puls et al. 2005; Rivero Gonz\u00e1lez et al. 2011; Carneiro et al. 2016). FASTWIND is routinely applied for both photospheric and wind analyses of hot stars, and used for detailed studies of individual objects as well as in large spectroscopic surveys (like within the recent VLT-FLAMES survey of massive stars in the Tarantula giant star-forming region in the Large Magellanic Cloud, Evans et al. 2011). A critical component of all these codes regards their practical treatment of the stellar wind; traditionally this has been to assume a parametrized steady-state and smooth outflow, without any clumps or shocks. However, it has been known for quite many years now, that these line-radiation driven winds are in fact inhomogeneous and highly structured on small spatial scales (see overviews in Puls et al. 2008, 2015; Hamann et al. 2008; Sundqvist et al. 2012b). Such wind clumping arises naturally from the strong line-deshadowing instability, the LDI, a fundamental and inherent property of line driving (e.g., Owocki & Rybicki 1984, 1985). Radiation-hydrodynamic, time-dependent wind models (Owocki et al. 1988; Feldmeier et al. 1997; Owocki & Puls 1999; Dessart & Owocki 2003; Sundqvist & Owocki 2013, 2015; Sundqvist et al. 2018) following the non-linear evolution of this LDI show a characteristic two-component-like structure consisting of spatially small and dense clumps separated by large regions of very rarified material, accompanied by strong thermal shocks and a highly non-monotonic velocity field. Such clumpy winds then affect both the atmospheric structure and the radiative transfer needed to derive synthetic observables; as just one example of this, neglecting clumping typically leads to observationally inferred mass-loss rates that might differ by more than an order of magnitude for the same star, depending on which spectral diagnostic is used to estimate this mass loss (Fullerton et al. 2006). Global model atmospheres nowadays normally account for such wind inhomogeneities by simply assuming a two-component medium consisting of overdense, optically thin clumps of a certain volume filling factor, following a smooth, parametrized velocity law, and an inter-clump medium that is effectively void (e.g., Hillier 1991; Puls et al. 2006). However, if clumps become optically thick, it leads to an additional leakage of light \u2013 not accounted for in the filling factor approach \u2013 through porous channels in between the clumps. Such porosity can occur either spatially (e.g., Feldmeier et al. 2003; Owocki et al. 2004; Sundqvist et al. 2012a), or for spectral lines in velocity-space due to Doppler shifts in the rapidly accelerating wind (sometimes thus called velocity-porosity, or \u201cvorosity\u201d, Owocki 2008). Regarding spatial porosity, several studies over the past years have focused on examining potential effects on the bound-free absorption of X-ray photons by the bulk wind (e.g., Oskinova et al. 2006; Owocki & Cohen 2006; Sundqvist et al. 2012a; Leutenegger et al. 2013; Herv\u00e9 et al. 2013). Regarding velocity-space porosity, similar studies (Oskinova et al. 2007; Hillier 2008; Sundqvist et al. 2010, 2011, 2014; \u0160urlan et al. 2012, 2013) have shown that clumps indeed very easily become optically thick in especially the strong UV wind-lines of hot stars (the so-called P-Cygni lines), and that the associated additional leakage of line-photons leads to weaker line profiles than predicted by smooth or volume filling factor models1. But constructing realistic, multi-dimensional ab-initio radiation-hydrodynamic wind simulations that account naturally for (time-dependent) spatial and velocity-field porosity is an extremely challenging and time-consuming task (Sundqvist et al. 2018). Thus there has also been a big need for developing simplified, parameterized models that can be more routinely applied to diagnostic work on samples of hot stars with winds. Building on their prior studies Sundqvist et al. (2010; 2011; 2012a; 2014, hereafter SPO14) developed and benchmarked such a method, using effective quantities to simulate the reduction in opacity associated with optically thick clumps. In contrast to some other models mentioned above, this \u201ceffective opacity\u201d approach has the great advantage that it can be quite readily implemented into the already existing (time-independent) global NLTE atmosphere models discussed above.","Citation Text":["Owocki et al. 1988"],"Citation Start End":[[3233,3251]]}
{"Identifier":"2016AandA...588A..26H__hand,_Bouwens_et_al._(2009)_Instance_1","Paragraph":"Various studies have investigated the color-magnitude relation at high redshifts (z\u2009\u2273\u20092), and the results are inconclusive. Reddy et al. (2008) found weak-to-no correlation between dust or UV color (related as shown in Fig. 4) and R mag for UV-selected galaxies at 1.5\u2009\u2272\u2009z\u2009\u2272\u20092.6. Similarly, Heinis et al. (2013) through their investigation of far-infrared\/dust properties of UV selected galaxies at z\u2009~\u20091.5 found that the average UV slope is mostly independent of the UV luminosity. On the other hand, Bouwens et al. (2009) studied the relation between the UV slope and MUV for a sample of U-band dropout galaxies at z\u2009\u2243\u20092.5 and found a positive color-magnitude correlation. At higher redshifts (z\u2009>\u20093), similar disagreements have been observed for samples based on the Lyman break color selection, in the sense that, some studies find no (or very weak) trends between \u03b2 and UV magnitude (e.g., Ono et al. 2010; Finkelstein et al. 2012; Dunlop et al. 2012; Castellano et al. 2012), while other studies find a strong color-magnitude relation (e.g., Wilkins et al. 2011; Bouwens et al. 2012; Rogers et al. 2014; Bouwens et al. 2014). Bouwens et al. (2014) also argue that this relation could be described by a double power-law with a different slope for fainter magnitudes. The differences in these various studies could be because of many systematic and\/or physical reasons, such as selection of galaxy samples (especially at lower redshifts), UV magnitude and\/or \u03b2 measurements, intrinsic scatter in \u03b2, dynamic range in UV luminosity. Our study relies on a UV selected sample with spectroscopic redshifts but it lacks a significant dynamic range in UV luminosity as seen in Fig. 5. A detailed study of a uniformly selected spectroscopic sample covering a wider range in UV luminosity is needed to understand the true nature of the color-magnitude relation at high redshifts. At least for ~L\u2217 (~0.5L\u2217 \u2013 3.0L\u2217) galaxies in a uniformly selected VUDS sample at 2\u2009\u2009z\u2009\u20092.5 there appears to be no correlation between \u03b2 and M1500. ","Citation Text":["Bouwens et al. (2009)"],"Citation Start End":[[502,523]]}
{"Identifier":"2021MNRAS.503.5367B__Kaspi_&_Beloborodov_2017_Instance_1","Paragraph":"Radio searches for single pulses and average pulse emission by folding the data based on the ephemeris derived from the X-ray observations (e.g. Israel et al. 2016) were carried out using the NC, GMRT (Surnis et al. 2016), and Arecibo (Younes et al. 2017b), on time-scales of days to a few weeks after the outbursts in 2014\u20132016. Neither telescope reported a detection, and the GMRT quotes upper limits of 0.4 and 0.2\u2009mJy at 326.5 and 610\u2009MHz, respectively, assuming a 10 per cent duty cycle and an 8\u03c3 detection threshold. The Arecibo observations were at 6.7 and 1.4\u2009GHz and they report upper limits of 14 and 7\u2009\u03bcJy, respectively, having used a 10\u03c3 threshold and a 20 per cent duty cycle. However, regular monitoring is supported by the most recent detection of just three pulses with CHIME (Good & CHIME\/FRB Collaboration 2020) and the subsequent detection of single pulses and average pulse emission with FAST (Zhu, Wang & Zhou 2020), which also appears to have been transient in nature. It is also unclear how long it takes before the radio emission turns on in these systems. Sometimes it is relatively rapid (e.g. Levin et al. 2019) and other times it can be quite some time after the X-ray outbursts (see Kaspi & Beloborodov 2017, for more details) and can apparently also vary for a given source too. Usually, when a magnetar turns on in the radio, it emits radio pulses during most rotations and stays on, with some significant variability in flux density and pulse shape, over a time-scale of months to years. In the case of SGR\u2009J1935+2154, the Kirsten et al. (2020a) campaign lasted for 2 months with a total observing time of 522 h, while our campaign lasted for 1.5 months for a total integration time of \u223c110 h. Between these two overlapping campaigns, only two radio pulses were seen by Kirsten et al. (2020a). These results along with the the very recent detection of bright radio single pulses and pulsed radio emission from the magnetar (Good & CHIME\/FRB Collaboration 2020; Zhu et al. 2020) suggest that the change in the magnetosphere configuration in SGR\u2009J1935+2154 to enable pulsed radio emission might be more gradual or sporadic than other magnetars like XTE J1810\u2013197 and Swift J1818.0\u20131607 where bright radio pulses were seen in the immediate aftermath of the X-ray bursting behaviour. A possible change in the magnetosphere is also supported by X-ray pulsar variations (Younes et al. 2020b). This diversity of emission characteristics emphasizes the need to regularly monitor magnetars in their active states.","Citation Text":["Kaspi & Beloborodov 2017"],"Citation Start End":[[1212,1236]]}
{"Identifier":"2020AandA...635A.185P__Aharonian_et_al._2006_Instance_1","Paragraph":"Two additional interesting sources for a comparison are the radio galaxies M 87 (z\u2004=\u20040.0044) and 3C 84 (also called NGC 1275, z\u2004=\u20040.0176). Both sources have a viewing angle similar to NGC 3894: 15\u00b0\u2005 \u03b8\u2004 \u200425\u00b0 (Mertens et al. 2016) and \u03b8\u2004 \u200418\u00b0 (Giovannini et al. 2018), respectively. The velocities of the jet components are about 0.2c in 3C 84 (Suzuki et al. 2012) and in the inner part of M 87 (Mertens et al. 2016), similar to what we find for NGC 3894. The three sources have a comparable photon index (\u0393 \u223c 2.1). Although it is less distant and more extended (\u223c100 kpc), M 87 shows a \u03b3-ray luminosity of 6\u22127\u2005\u00d7\u20051041 erg s\u22121, similar to NGC 3894. Different is the case of 3C 84, which has two-sided compact (parsec scale) jets and a distance similar to NGC 3894, but it presents a \u03b3-ray luminosity more than 200 times higher (2\u2005\u00d7\u20051044 erg s\u22121). Unlike NGC 3894, for which no observations with Cherenkov telescopes have been performed, M 87 and 3C 84 are also detected at very high energy (VHE, E > 100 GeV; Aharonian et al. 2006; Aleksi\u0107 et al. 2012). In particular, at VHE, M 87 displayed strong variability on timescales as short as one day, but no unique signature of the region responsible for the VHE flares has been identified (Abramowski et al. 2012). Despite the long-term monitoring of M 87 during 2012\u22122015 in a low-activity state, the production site of \u03b3-rays remains unclear (MAGIC Collaboration 2020). However, the correlation observed between the radio and X-ray activities is a strong indication that most often the emitting region is close to the core in this source. In the case of 3C 84, several works on radio, X-ray, and \u03b3-ray variability suggest that short-term and long-term variability may be produced in different regions of the source. In particular, short-term variability seems related to the injection of fresh particles that are accelerated in a shock in the core region, whereas long-term variability is more likely connected with the jet structure (e.g., Fukazawa et al. 2016; Hodgson et al. 2018). Similarly to M 87 and 3C 84, long-term radio and X-ray monitoring of NGC 3894 may provide important information about the origin of the \u03b3-ray emission.","Citation Text":["Aharonian et al. 2006"],"Citation Start End":[[1006,1027]]}
{"Identifier":"2022AandA...660A.108Z__Alencar_et_al._2018_Instance_1","Paragraph":"In comparison with studies focused on other T Tauri systems, the period analysis of the H\u03b1 lines for this target often do not exhibit a single and highly significant peak. As the H\u03b1 line is expected to form not only at the accretion spot, but also in the accretion funnel and in the stellar winds, the variability of this line is also influenced by the variations of the circumstellar environment. However, the rotation period might appear with varying strengths in the periodograms. It is instructive to compare our findings with similar analyses of individual targets. In the LkCa 15 system, the strength of the emission at the line center is modulated by stellar rotation, and this effect is clearly displayed in the periodogram (Alencar et al. 2018). On the other hand, the inverse P Cygni profile of the H\u03b1 line of the HQ Tau system displays modulation on the timescale of the stellar rotation not only in the line center, but all the way to the red wing (Pouilly et al. 2020). Another interesting system to mention is V2129 Oph. When analyzing this target, Sousa et al. (2021) compared their results with the periodograms for the same system as was analyzed by Alencar et al. (2012), and they found some differences on timescales of a decade, similarly to our results. Sousa et al. (2021) discussed a few possible explanations for the observed phenomena, such as variations in the magnetic field strength, which would result in variations in the magnetospheric truncation radius, or latitudinal differential rotation, or a more complex magnetic structure with two major funnel flows originating at different radii in the inner disk. However, in the case of V2129 Oph, a periodicity longer than the stellar rotation period is present at the line center for a time of almost a decade. The likely cause of the periodicity is a structure that is present beyond the corotation radius that is stable for almost a decade. In the case of CR Cha, the rotation period has been stable for decades, but the line profiles (and the photometry as well) also vary with longer periods, likely because of variations in the circumstellar environment.","Citation Text":["Alencar et al. 2018"],"Citation Start End":[[733,752]]}
{"Identifier":"2018AandA...616A.139G__Marrone_&_Rao_(2008)_Instance_1","Paragraph":"Observations of the polarized dust emission of nine low-mass protostars at 0.87 mm were obtained using the SMA (Projects 2013A-S034 and 2013B-S027, PI: A. Maury) in the compact and subcompact configuration. To increase our statistics, we also included SMA observations from three additional sources from Perseus (NGC 1333 IRAS4A and IRAS4B) and Ophiuchus (IRAS16293) observed in 2004 and 2006 (Projects 2004-142 and 2006-09-A026, PI: R. Rao; Project 2005-09-S061, PI: D.P. Marrone). The observations of NGC 1333 IRAS4A and IRAS16293 are presented in Girart et al. (2006) and Rao et al. (2009), respectively. Marrone (2006) and Marrone & Rao (2008) provide a detailed description of the SMA polarimeter system, but we provide a few details on the SMA and the polarization design below. The SMA has eight antennas. Each optical path is equipped with a quarter-wave plate (QWP), an optical element that adds a 90\u00b0 phase delay between orthogonal linear polarizations and is used to convert the linear into circular polarization. The antennas are switched between polarizations (QWP are rotated at various angles) in a coordinated temporal sequence to sample the various combinations of circular polarizations on each baseline. The 230, 345, and 400 band receivers are installed in all eight SMA antennas. Polarization can be measured in single-receiver polarization mode and in dual-receiver mode when two receivers with orthogonal linear polarizations are tuned simultaneously. In this dual-receiver mode, all correlations (the parallel-polarized RR and LL and the cross-polarized RL and LR; with R and L for right circular and left circular, respectively) are measured at the same time. Both polarization modes were used in our observations. This campaign was used to partly commission the dual-receiver full polarization mode for the SMA. A fraction of the data was lost during this period owing to issues with the correlator software. Frequent observations of various calibrators were interspersed to ensure that such issues were detected as early as possible to minimize data loss.","Citation Text":["Marrone & Rao (2008)"],"Citation Start End":[[627,647]]}
{"Identifier":"2021AandA...656A..18H__Strangeway_&_Crawford_1993_Instance_1","Paragraph":"Since the 1960s, Venus has been investigated by various flyby missions (Mariner 2, 5 and 10, Galileo, and Cassini), landers (Venera 11\u201314), atmospheric probes, and orbiter missions (Pioneer Venus, Venus Express, and Akatsuki), out of these, Mariner 10 had the most similar flyby trajectory to Solar Orbiter and could study magnetic field fluctuations far downtail (Lepping & Behannon 1978). Nevertheless, the nature of the high frequency plasma waves at Venus has been mainly studied by the Pioneer Venus Orbiter (PVO), as it was the only orbiter that carried a plasma wave instrument, the Orbiter Electric Field Detector (OEFD; Scarf et al. 1980b). Various types of propagation modes have been identified in different regions around the planet: 30 kHz Langmuir oscillations were present in the foreshock, 5.4 kHz Doppler-shifted electrostatic ion-acoustic waves were observed in the bow shock and appeared to be a prominent feature of the magnetotail boundary, and electromagnetic 100 Hz whistler waves were thought to be generated at the bow shock, propagating through the magnetosheath and eventually absorbed within the ionosphere (Scarf et al. 1979, 1980a; Strangeway 1991). However, the identification of these waves was not well determined as they were also suspected to be lower hybrid waves (Szeg\u00f6 et al. 1991), or ion acoustic waves (Strangeway & Crawford 1993; Huba 1993). This ambiguity was a direct consequence of the absence of a wave magnetic field sensor on-board PVO and the limited frequency coverage of the OEFD instrument, that had only a 4 channel electric field receiver (E field wave power at 100 Hz, 730 Hz, 5.4 kHz, 30 kHz). Furthermore, after more than forty years long search for lightning on Venus, results obtained by the different spacecraft missions are still inconclusive. The most compelling indications of some kind of atmospheric electrical activity at Venus appear to be the Venera landers\u2019 detections of electromagnetic pulses at frequencies from 10 to 80 kHz (Ksanfomaliti 1980). Their origin is questionable (Lorenz 2018); they might be assigned to interferences similarly as electric pulses found by Galileo (Gurnett et al. 1991). No lightning related signals were found during the Cassini Venus flyby in spite of the fact that the same instrument detected up to 70 sferics per second during the Earth flyby in 1999 (Gurnett et al. 2001). The whistler-mode waves at frequencies near 100 Hz recorded by Pioneer Venus and Venus Express were attributed to lightening (Russell et al. 2008). Nevertheless, a variety of other explanations were offered for these electromagnetic observations (Lorenz 2018). Up to now, not a single image from any spacecraft has been able to confirm the existence of lightning including an extensive a 5-years long search for lightning by the Akatsuki Venus orbiter (Takahashi et al. 2018). However, the modeling did not exclude electrical breakdown in the Venusian atmosphere and transient luminous events were forecasted to occur on Venus (Riousset et al. 2020).","Citation Text":["Strangeway & Crawford 1993"],"Citation Start End":[[1344,1370]]}
{"Identifier":"2020AandA...642A..24W__Nicastro_et_al._2016_Instance_1","Paragraph":"The high temperatures of these filaments result in the matter being highly ionised so that it absorbs and emits far-UV and soft X-ray photons. However, these WHIM filaments are of low column density and, therefore, produce a low signal intensity which makes them difficult to detect with the current instrumentation. In spite of this, the search for the missing baryons has continued and detections of extragalactic O\u202fVII and O\u202fVIII have been claimed (Nicastro et al. 2005) but they are often disputed (Kaastra et al. 2006; Williams et al. 2006; Rasmussen et al. 2007), of a single-line or low statistical significance (Nicastro et al. 2010; Bonamente et al. 2016), misidentifications of the WHIM (Nicastro et al. 2016), or they are local to the background source (Johnson et al. 2019). More recently, Nicastro et al. (2018), later updated by Nicastro (2018), reported the detection of two O\u202fVII absorption features at redshifts z = 0.3551 and z = 0.4339 along the sight-line towards blazar 1ES 1553+113 after a 1.75 Ms observation with the XMM-Newton reflection grating spectrometer. Johnson et al. (2019) conducted additional studies towards 1ES 1553+113, suggesting that the feature located at z = 0.4339 may not originate from the WHIM and is associated, rather, with the local environment of the blazar. A study by Kov\u00e1cs et al. (2019) quotes a 3.3\u03c3 detection of O\u202fVII towards H1821+643 from a stacked spectrum of Chandra observations with an 8 Ms total exposure. Ahoranta et al. (2020) performed spectral analysis of Chandra and XMM-Newtown observations at previously determined FUV redshifts, yielding two X-ray line candidates of Ne\u202fIX[[Inline38]He and O\u202fVIII Ly\u03b1 at a combined significance of 3.9\u03c3. In spite of the large observational efforts, only these few marginal detections have been achieved so far. The small equivalent widths of \u22480.07\u22120.42 eV (Branchini et al. 2009; Nicastro et al. 2018) from WHIM absorption features mean that high signal-to-noise ratio observations are required for their detection. This calls for an observatory with a large effective area, high energy resolution, and a low energy threshold in the soft X-ray energy band.","Citation Text":["Nicastro et al. 2016"],"Citation Start End":[[698,718]]}
{"Identifier":"2017ApJ...839...56T___2013_Instance_1","Paragraph":"Mid-infrared and millimeter polarimetric observation has so far been considered as the best method to probe the magnetic field. This is because if aspherical grains in disks become aligned with the magnetic field as is the case in the interstellar medium (ISM), the polarization vector arising from thermal emission of the aligned grains becomes perpendicular to the local magnetic field line (Cho & Lazarian 2007, henceforth CL07, Matsakos et al. 2016; Yang et al. 2016b; Bertrang et al. 2017). At mid-infrared wavelengths, Li et al. (2016) performed a polarimetric imaging observation of AB Aur using CanariCam. As a result, they detected a centrosymmetric polarization pattern, and the degree of polarization was as high as 1.5% at large radii. At millimeter wavelengths, polarimetric observations of disks have been performed (e.g., Hughes et al. 2009, 2013; Rao et al. 2014; Stephens et al. 2014; Cox et al. 2015; Kataoka et al. 2016b). Polarized emission from a circumstellar disk has been detected in the Class 0 phase (Rao et al. 2014; Cox et al. 2015). More evolved disks do not show a degree of linear polarization larger than 0.5% (Hughes et al. 2009, 2013). It should be mentioned that Stephens et al. (2014) detected polarized flux from HL Tau, which is classified as a Class I-II, with an average degree of linear polarization of 0.9%. More recently, Kataoka et al. (2016b) reported the first submillimeter polarization observation of a disk obtained with ALMA, and they clearly detected polarized flux from HD 142527. The polarization fraction at a peak position of polarized intensity was 3.26%, and the maximum polarization fraction was as high as 13.9%. The disk reveals radial polarization vectors; however, they flip by 90\u00b0 in its northeast and northwest regions. In addition, the detected polarization fraction is much larger than the stringent limit set by Hughes et al. (2009, 2013), and further polarimetric observations by ALMA will reconcile this discrepancy.","Citation Text":["Hughes et al.","2013"],"Citation Start End":[[837,850],[857,861]]}
{"Identifier":"2015ApJ...799..149J___2014_Instance_4","Paragraph":"With our joint analysis of stellar mass fraction and source size, we find a larger stellar mass fraction than earlier statistical studies. In Figure\u00c2 2, we compare our determination of the stellar surface density fraction to a simple theoretical model and to the best fit of a sample of lens galaxies by Oguri et\u00c2 al. (2014). The simple theoretical model is the early-type galaxy equivalent of a maximal disk model for spirals. We follow the rotation curve of a de Vaucouleurs component for the stars outward in radius until it reaches its maximum and then simply extend it as a flat rotation curve to become a singular isothermal sphere (SIS) at large radius (see details in the Appendix). The ratio of the surface mass density of the de Vaucouleurs component to the total surface mass density is shown as a dashed curve in Figure\u00c2 2. We also show as a gray band the best fit for the stellar fraction in the form of stars determined by Oguri et al (2014) in a study of a large sample of lens galaxies using strong lensing and photometry, as well as the best model using a Hernquist component for the stars and an NFW halo for the dark matter with and without adiabatic contraction, also from Oguri et\u00c2 al. (2014). We have used the average and dispersion estimates for the Einstein and effective radii available for 13 of the objects in our sample from Oguri et\u00c2 al. (2014), Sluse et\u00c2 al. (2012), Fadely et\u00c2 al. (2010), and Leh\u00c3\u00a1r et\u00c2 al. (2000; see Table\u00c2 1) as an estimate of RE\/Reff in Figure\u00c2 2. The average value and dispersion of the sample is RE\/Reff = 1.8 \u00c2\u00b1 0.8. This also averages over the different radii of the lensed images. The agreement of our estimates with the expectations of the simple theoretical model and with estimates from other studies (Oguri et\u00c2 al. 2014) is quite good. For comparison, the estimate of Pooley et\u00c2 al. (2012; using the Einstein and effective radii estimates for 10 out of 14 of their objects from Schechter et\u00c2 al. 2014) seems somewhat lower than expected at those radii. The range of stellar mass fractions from MED09 for source sizes in the range 0.3\u00e2\u0080\u009315.6 light days is also shown in Figure\u00c2 2. In this case, the discrepancy between our estimate and their reported value of \u00ce\u00b1 = 0.05 is completely due to the effect of the source size. Although accretion disk sizes are known to be smaller in X-rays, recent estimates are in the range of 0.1\u00e2\u0080\u00931 light-days, depending on the mass of the black hole (see Mosquera et\u00c2 al. 2013), and these finite sizes will increase the stellar surface densities implied by the X-ray data. Another possible origin for this discrepancy is that Pooley et\u00c2 al. (2012) use the macro model as an unmicrolensed baseline for their analysis. It is well known that simple macro models are good at reproducing the positions of images, but have difficulty reproducing the flux ratios of images due to a range of effects beyond microlensing. Recently, Schechter et\u00c2 al. (2014) found that the fundamental plane stellar mass densities have to be scaled up by a factor 1.23 in order to be compatible with microlensing in X-rays in a sample of lenses with a large overlap with that analyzed by Pooley et\u00c2 al. (2012). It is unclear how this need for more mass in stars at the position of the images found by Schechter et\u00c2 al. (2014) can be reconciled with the apparently low estimate of mass in stars at those radii by Pooley et\u00c2 al. (2012). Our estimate of the stellar mass fraction agrees better with the results of microlensing studies of individual lenses (Keeton et\u00c2 al. 2006; Kochanek et\u00c2 al. 2006; Morgan et\u00c2 al. 2008, 2012; Chartas et\u00c2 al. 2009; Pooley et\u00c2 al. 2009; Dai et\u00c2 al. 2010) that reported values in the range 8%\u00e2\u0080\u009325%, and with the estimates from strong lensing studies (see for example Jiang & Kochanek 2007; Gavazzi et\u00c2 al. 2007; Treu 2010; Auger et\u00c2 al. 2010; Treu et\u00c2 al. 2010; Leier et\u00c2 al. 2011; Oguri et\u00c2 al. 2014) which produced stellar mass fractions in the range 30%\u00e2\u0080\u009370% integrated inside the Einstein radius of the lenses.","Citation Text":["Oguri et\u00c2 al. (2014)"],"Citation Start End":[[1353,1373]]}
{"Identifier":"2022AandA...666A.141M__Hatch_et_al._(2011)_Instance_1","Paragraph":"In Sect. 4.2 we see that the Q.S. decreases with both the stellar mass and the local density. The Q.S. provides us with a glimpse of quenching in the past. It significantly decreases with stellar mass at the low-mass end (similarly for the cluster candidates and the field), showing that more massive galaxies are at a much later stage of quenching. This is consistent with the downsizing scenario (e.g. Cowie et al. 1996; Cattaneo et al. 2008; Webb et al. 2020) that massive galaxies form and quench earlier, and hence they are at a later, more advanced stage of evolution compared to low-mass galaxies. The mass dependence of the quenching stage is strongly affected by the local density. The trend is only significant in high-density environments. For sparse environments, the errorbars are too large for us to say anything. The scenario is consistent with Hatch et al. (2011), where faster evolution in protoclusters is found. This can also be due to the migration of galaxies from the low-density cluster outskirts to the high-density cluster cores where the quenching happens. The Q.S. also shows a stable decreasing trend with the local density, indicating that galaxies in denser environments (cluster cores) are at a later, more advanced stage of quenching. The earlier Q.S. of field galaxies also supports this scenario. This is consistent with the inside-out quenching in clusters, where galaxies in cluster cores quench early, while the galaxies in the outer parts quench progressively later, which is consistent with the scenario in Koyama et al. (2013) and Shimakawa et al. (2018). Another possibility is that during the in-falling process their gas will be stripped off, which directly leads to their quenching (Bekki et al. 2002). In this case the effect should be stronger for low-mass galaxies since their gas components are easier to strip. This environmental dependence is not strongly affected by the stellar mass, which supports the inside-out quenching that has an impact on low- and high-mass galaxies equally. In McNab et al. (2021) the authors use three types of transitional populations: green valley (GV) selected by rest-frame (NUV\u2212V) and (V\u2005\u2212\u2005J) colours; blue quiescent (BQ) selected by rest-frame (U\u2005\u2212\u2005V) and (V\u2005\u2212\u2005J) colours; and post-starburst (PSB) selected by galaxy spectra. They connect the quenching timescale to the transitional population fraction excess in clusters. They found that the lowest mass systems ( 1010.5\u2006M\u2299) have short quenching timescales ( 1 Gyr), which agrees with our result that low-mass galaxies are dominated by fast quenching ( 1 Gyr).","Citation Text":["Hatch et al. (2011)"],"Citation Start End":[[860,879]]}
{"Identifier":"2016ApJ...821...86H__Steeb_2014_Instance_1","Paragraph":"A challenge with empirical methods is that with growing numbers of input parameters, it becomes prohibitive in terms of CPU time to use all of them. It can also be counterproductive to feed the machine learning tool with all of these parameters because some of them might either be too noisy or not bring any relevant information to the specific problem being tackled. Moreover, high-dimensional problems suffer from being overfit by machine learning methods (Cawley & Talbot 2010), thereby yielding high-dimensional nonoptimal solutions. This requires subselection of relevant variables from a large N-dimensional space (with N potentially close to 1000). This task itself can also be achieved using machine learning tools. Here we present, to our knowledge, the first application to astronomy of the combination of two machine learning techniques: genetic algorithms (GA; e.g., Steeb 2014) for selecting relevant features, followed by SVMs to build a decision\u2013reward function. GAs alone have already been used in astronomy, for instance, in the study of orbits of exoplanets and the SEDs of young stellar objects (Cant\u00f3 et al. 2009), the analysis of SN Ia data (Nesseris 2011), the study of star-formation and metal-enrichment histories of a resolved stellar system (Small et al. 2013), the detection of globular clusters from Hubble Space Telescope (HST) imaging (Cavuoti et al. 2014), and photometric redshift estimation (Hogan et al. 2015). At the same time, SVMs have been extensively used to solve a number of problems, such as object classification (Zhang & Zhao 2004), the identification of red variable sources (Wo\u017aniak et al. 2004), photometric redshift estimation (Wadadekar 2005), morphological classification (Huertas-Company et al. 2011), and parameter estimation for Gaia spectrophotometry (Liu et al. 2012). We note also that the combination of GA and SVM has already been used in a number of fields, such as cancer classification (e.g., Huerta et al. 2006; Albda et al. 2007), chemistry (e.g., Fatemi & Gharaghani 2007), and bankruptcy prediction (e.g., Min et al. 2006). We do not attempt here to provide the most optimized results from this combination of methods. Rather, we present a proof of concept that shows that GA and SVM yield remarkable results when combined, as opposed to using the SVM as a standalone tool.","Citation Text":["Steeb 2014"],"Citation Start End":[[880,890]]}
{"Identifier":"2021ApJ...914L...6A__Chhiber_et_al._2018_Instance_1","Paragraph":"On 2017 November 24 the MMS orbit allowed us to collect measurements in the pristine solar wind, well outside the Earth's magnetosheath and the bow shock, for a long period (i.e., a few times longer than the typical correlation scale) of approximately 1 hour from 01:10 to 02:10 UT. Figure 1 (upper panel) displays an overview of the magnetic field measurements collected by the FIELDS instrument suite (Torbert et al. 2016) on board of MMS1 with a temporal resolution \u0394t = 128 samples s\u22121 (Russell et al. 2016). The period of interest is a typical example of slow solar wind stream (V \u223c 377 km s\u22121), with an average magnetic field \u2329B\u232a \u223c 6.6 nT and a mean plasma density \u2329n\u232a \u223c 9 cm\u22123 (Roberts et al. 2020a, 2020b). This means that the average Alfv\u00e9n speed is VA \u223c 50 km s\u22121, while the ion inertial length and gyroradius are di \u223c 76 km and \u03c1i \u223c 96 km, respectively (Chhiber et al. 2018), with the corresponding timescales \u03c4d \u223c 1.3 s and \u03c4\u03c1 \u223c 1.6 s, respectively. As reported in previous works (Bandyopadhyay et al. 2018; Chhiber et al. 2018; Roberts et al. 2020a, 2020b) this interval is characterized by two different spectral scalings: a typical inertial range \u223c \u03c45\/3 is found at large scales (i.e., \u03c4 > \u03c4b), while a steeper scaling \u223c \u03c47\/3 is found at small scales (i.e., \u03c4 \u03c4b), with \u03c4b \u223c 2.4 s (Roberts et al. 2020a). Furthermore, the magnetic field spectrum flattens near \u03c4noise \u223c 0.2 s, due to the instrumental noise floor near \u223c5 Hz (Russell et al. 2016). Finally, a decrease at shorter timescales (e.g., \u03c4 \u223c 0.1 s) is due to an anti-aliasing filter of nonphysical origin (Russell et al. 2016; Roberts et al. 2020a). Taken together, this interval is particularly suitable for testing our formalism with respect to processes of both physical and nonphysical origin. The presence of an instrumental noise floor allows us indeed to assess our formalism with respect to purely stochastic processes, while the existence of two spectral regimes (i.e., the MHD\/inertial and the kinetic\/dissipative) allows us to investigate small- versus large-scale processes and their possible coupling in a dynamical system framework.","Citation Text":["Chhiber et al. 2018"],"Citation Start End":[[865,884]]}
{"Identifier":"2016AandA...588A...2L__M\u00e4tzler_(1998)_Instance_3","Paragraph":"H2O ice on Pluto has long escaped spectroscopic detection, and based on initial New Horizons data appears to be exposed only in a number of specific locations, usually associated with red color, suggestive of water ice\/tholin mix (Grundy et al. 2015; Cook et al. 2015). Nonetheless, water ice is likely to be ubiquitous in Pluto\u2019s near subsurface, given its cosmogonical abundance, Pluto\u2019s density, and its presence on Charon\u2019s surface6. Absorption coefficients for pure water ice (kH2O) at sub-mm-to-cm wavelengths are discussed extensively by M\u00e4tzler (1998), who also provides several analytic formulations to estimate them as a function of frequency and temperature along with illustrative plots. We use the Mishima et al. (1983) formulation (see Appendix of M\u00e4tzler 1998). Its applicability is normally restricted to temperatures above 100 K, but Fig. 2 of M\u00e4tzler (1998) indicates the trend with temperature. Absorption coefficients extrapolated to 50 K (estimated as half the values at 100 K) are shown in Fig. 5. At 500 \u03bcm, our best estimate is kH2O = 0.25 cm-1, comparable to the above values for CH4 and N2 ices. The corresponding penetration length is therefore comparable to the diurnal skin depth but remains negligible compared to the seasonal skin depth, even for seasonal \u0393 = 25 MKS. According to these calculations, the seasonal layer would be probed only at a wavelength of ~4 mm and beyond. We also remark that the expression from Mishima et al. (1983) would give a penetration depth of 56 m at 2.2 cm, which is an order of magnitude larger than indicated by the laboratory measurements of Paillou et al. (2008). In addition, small concentrations of impurities can dramatically reduce the microwave transparency of water ice (e.g., Chyba et al. 1998 and references therein). Therefore, the above calculations likely indicate upper limits to the actual penetration depth of radiation in a H2O ice layer, from which we conclude that the seasonal layer is not reached at the Herschel wavelengths. ","Citation Text":["M\u00e4tzler (1998)"],"Citation Start End":[[861,875]]}
{"Identifier":"2016MNRAS.461..385B__Collings_et_al._2004_Instance_1","Paragraph":"As an additional constraint we have to make sure that our MCMax models are consistent with the observed CO snowline location. Therefore, we analyse the mid-plane temperature profile Tmid-plane(r) for all our models that match the observed SED, which we plot in Fig. 4. As mentioned above, models of the same series have the same temperature structure. We find that all of them have mid-plane temperatures between \u223c20and25\u2009K at the observed snowline radius of $R_{\\rm sl} \\approx 90$ au (Qi et al. 2015). These are well within the range of values generally assumed and observed for the freeze-out temperature of CO: The freeze-out temperature can vary between \u223c20 and \u223c30 K depending on whether CO is binding to pure CO ice or a mixture with water ice (Collings et al. 2004), which is, in turn, also dependent on the chemical history of the ice (Garrod & Pauly 2011). In general, the CO freeze-out temperature is not known unambiguously and might vary from system to system (Hersant et al. 2009; Qi et al. 2015): Qi et al. (2013a) found a freeze-out temperature of CO of 17 K from their modelling of TW Hya, whereas J\u00f8rgensen et al. (2015) obtain temperatures of about 30 K in their study of embedded protostars. Qi et al. (2011) assume a freeze-out temperature for CO of T \u2248 19\u2009K (pure CO ice) for HD 163296. However, in Qi et al. (2015), they perform a new analysis with higher resolution observational data and use a temperature in the mid-plane at Rsl of T \u2248 25\u2009K (mixed CO\/H2O ice). Thus, all our models have temperatures in the disc mid-plane at the location of the snowline radius that are well within the plausible range. Our exploration of self-consistent models confirms that all the freeze-out temperatures assumed in these previous literature references fall within the plausible range of temperatures for HD 163296. If the freeze-out temperature of CO was known unambiguously, this would, in combination with an observationally determined CO snowline location, be a powerful model discriminant and we might be able to exclude models based on this constraint. Since, however, it is unclear what exactly the relevant freeze-out temperature is, we find that all the models can match the observed snowline location of 90 au. This weak model discrimination also means that it is impossible to predict the CO snowline radius from SED model fitting alone or even from fitting the molecular line emission together with the SED (Qi et al. 2011): given the uncertainty in the sublimation temperature of CO, our viable SED fits imply predicted radii in the range \u223c40\u2013135 au as denoted by the red asterisks in Fig. 4. In general, we find that the location of the CO snowline radius does not further discriminate between models in comparison to the criterion given by the SED; however, it sets the radial location inwards of which no freeze-out is taking place in our models, and which is therefore important for the interpretation of the C18O emission. It is important to note that an SED fit does not determine the CO snowline location and that, vice versa, a CO snowline observation does not discriminate amongst possible SED models.","Citation Text":["Collings et al. 2004"],"Citation Start End":[[752,772]]}
{"Identifier":"2018ApJ...859..115W__Kang_et_al._2005_Instance_1","Paragraph":"We futher examine the distributions of the radial number fraction of satellites with different colors in different LSE in Figure 5 to study the role central galaxies play in SCA. It can be seen, in general, that the intersections of these lines for each subsample are located between \n\n\n\n\n\n to \n\n\n\n\n\n, and with a mean value of \n\n\n\n\n\n. Satellites (especially red satellites) prefer to reside in the inner part of groups and the number fraction decreases with the increasing radius. Such evolutionary behavior is consistent with \u201cstrangulation\u201d wherein cold gas that could be used for star formation is depleted as satellites are accreted (e.g., Kauffmann et al. 2004; Kang et al. 2005; Luo et al. 2016). The shapes of each curve (the maximum value, the decrease gradient) not only depend on the color of central and satellite galaxies but also depend on the LSE. In the inner part of halos, the fraction of satellites is lower in the more dense LSEs. In the outer parts, it is the opposite, namely, in knots there is a higher fraction of satellites. It must be pointed out is that, in a given environment and for a given central color, the fraction of blue satellites is always higher than red satellites in outer regions. For wall environments, about 25% \u223c 30% satellites reside at the low values of \n\n\n\n\n\n, and the fraction decreases very quickly. Additionally, in the inner parts, the fraction of satellites of red centrals (red dashed line in the top left panel of Figure 5) is \u223c5% higher than those in blue centrals (blue dashed line in the top left panel of Figure 5). The radial distributions of satellites in filaments are almost absent with the different colors of both central and satellite galaxies. However, in a knot environment, blue centrals have \u223c3% more satellites than red centrals at low values of projected radius. These radial satellite distributions with different colors in different LSE indicate that, red central galaxies have merged more satellites than blue central galaxies in the knot LSE, which means satellites with a red central galaxy may be more affected by the interaction with central galaxies and represent a better SCA, especially in the inner region.","Citation Text":["Kang et al. 2005"],"Citation Start End":[[667,683]]}
{"Identifier":"2021AandA...650A.164M__Davies_et_al._2012_Instance_2","Paragraph":"The GMC associated with G305 is one of the most massive and luminous clouds in the Milky Way (Fig. 1). It is located in the Galactic plane at l ~ 305\u00b0, b ~ 0\u00b0 and at a kinematic distance of 4 kpc (derived from a combinationof radio and H\u03b1 observationsby Clark & Porter (2004); Davies et al. (2012) measured its spectrophotometric distance to be 3.8 \u00b1 0.6 kpc and most recently Borissova et al. (2019) measured the Gaia DR2 average distance to be 3.7 \u00b1 1.2 kpc); this places it in the Scutum-Crux spiral arm. Given this distance, the complex has a diameter of ~ 30 pc (Clark & Porter 2004) and a molecular mass of ~6 \u00d7 105 M\u2299 (Hindson et al. 2010). The G305 complex consists of a large central cavity that has been cleared by the winds from massive stars belonging to two visible central clusters (Danks 1 and 2) and the Wolf-Rayet star (WR48a; Clark & Porter 2004; Davies et al. 2012). The cavity is surrounded by a thick layer of molecular gas (traced by CO and NH3 emission; Hindson et al. 2010, 2013). Radio continuum observations by Hindson et al. (2012) have revealed that the cavity is filled with ionized gas and identified six ultra-compact HII (UC HII) regions and also one bright rimmed cloud (BRC) at the periphery of the cavity, indicating molecular gas irradiated by UV radiation (Sugitani & Ogura 1994; Thompson et al. 2004), which may cause implosion (Bertoldi 1989) or evaporation. A number of studies havereported star formation tracers (water and methanol masers, HII regions and massive young stellar objects, MYSOs; Clark & Porter 2004; Lumsden et al. 2013; Urquhart et al. 2014; Green et al. 2009, 2012). Furthermore, Hindson et al. (2010) found the concentration of star formation tracers to be enhanced inside a clump of NH3 bearing molecular gas that faces the ionizing sources, which is consistent with the hypothesis that the star formation has been triggered. Analysis of the stellar clusters in the complex reveals them to have ages of 1.5 Myr for Danks 1 and 3 Myr for Danks 2,with the former possibly being triggered by the latter (Davies et al. 2012). Additionally, a diffuse population of evolved massive stars was also found to exist within the confines of the G305 complex that had formed around the same time as the two clusters (Leistra et al. 2005; Shara et al. 2009; Mauerhan et al. 2011; Davies et al. 2012; Faimali et al. 2012; Borissova et al. 2019).","Citation Text":["Davies et al. 2012"],"Citation Start End":[[865,883]]}
{"Identifier":"2018MNRAS.481..749S__Mazzotta_et_al._2004_Instance_1","Paragraph":"The mean temperature and pressure profiles, along with their associated intrinsic scatters, are also shown in Fig. 7. Overall, at larger radii, the mean temperatures are relatively flat, and we do not detect the clear decrease found in other studies (see Fig. 8; e.g. Vikhlinin et al. 2006; Leccardi & Molendi 2008; Arnaud et al. 2010; Eckert et al. 2013b; McDonald et al. 2014). Specifically, we do not see a temperature increase at small radii \u223c0.2r500 before a subsequent decrease at large radii. However, our temperature profile is compatible with the results of Leccardi & Molendi (2008) in that region, and only 1\u20132\u03c3 lower than the profiles obtained by Vikhlinin et al. (2006) and McDonald et al. (2014). While this difference is not statistically significant, it may be due to the selection effects given that the Leccardi & Molendi (2008) sample is most similar to ours in mass and redshift. Given the modest statistical significance of this difference, it may simply be due to noise fluctuations. However, inaccurate subtraction of the X-ray background may also play a role. For example, oversubtraction of the hard background would tend to bias spectroscopic X-ray temperatures low while biasing our derived temperatures high (via the reduction in X-ray-derived density). The difference could also be related to cluster physics. For instance, clumping within the ICM, which is expected to increase with radius, will tend to bias spectroscopically derived temperatures low compared to the X-ray\/SZ values derived in our analysis (e.g. Mazzotta et al. 2004; Rasia et al. 2005; Vazza et al. 2013). Elongation of the cluster along the line of sight, which will increase the SZ brightness compared to the X-ray brightness, could also artificially boost the temperatures recovered in our analysis (e.g. Cooray 2000). Another possibility for the slightly higher than expected temperatures at large radii may be biases in the SZ data. Indirect evidence for such an effect can be found by comparing the pressure profiles obtained by Sayers et al. (2013b), based solely on Bolocam data, to those obtained by Sayers et al. (2016) based on a joint analysis of Bolocam and Planck for a nearly identical set of clusters. Beyond \u22430.5r500, the latter work found a lower value for the pressure profile, with the difference increasing to a factor of \u22431.5 at r500. This implies that the Bolocam data alone may be overestimating the pressure at large radii, which would bias our temperature profiles high in that region. This can be seen in the mean pressure profile plot. However there is good agreement in the radial range of interest in this study when compared with Arnaud et al. (2010) and Planck Collaboration V (2013). In addition, the mean temperature profiles recovered for the CC and NCC subsets in the innermost shell are not different at a statistically significant level (i.e. the CC subset does not show a significant drop towards the cluster centre). This is likely a result of the coarse binning required by the SZ data. Specifically, the innermost shell extends to 0.15r500, which is generally not small enough to resolve the cool core. All of the potential biases and physical effects noted in this paragraph apply to the individual cluster temperature profiles as well, and therefore may explain differences compared to other studies (e.g. ACCEPT) that appear for some clusters.","Citation Text":["Mazzotta et al. 2004"],"Citation Start End":[[1544,1564]]}
{"Identifier":"2017AandA...604A.106J__Snaith_et_al._2014_Instance_1","Paragraph":"A third limitation, which is probably the most urgently in need of investigation, is the lack of a gaseous component both in the Milky Way disc and in the satellite(s). It is not easy to anticipate how the inclusion of gas in the simulations can impact the results. It has been shown (Moster et al. 2010; Qu et al. 2011) that the presence of a thin gaseous disc can limit the amount of disc heating: for gas fractions of the order of 20% (40%), the heating may be reduced by 25% (40%). For accretions occurred in the early history of the Galaxy, typically before redshift z ~ 1, the fractions of gas in the Milky Way disc may have been significantly higher (with fgas \u2265 50%) than those simulated in the above cited works. Moreover, in those early phases of the Galactic evolution it is not even clear how thin the gaseous disc may have been. While in nearby disc galaxies gaseous discs indeed have scale heights between a few tens and 100 pc and velocity dispersions of the order of ~10 km\u2009s-1. At higher redshift gaseous discs of approximately 1 kpc (corresponding to velocity dispersions of the order of 100 km\u2009s-1) have been observed in the continuum and in the gas component (Elmegreen & Elmegreen 2006; Epinat et al. 2012). In the N-body models cited above, the typical scale heights employed for the gaseous component are between 200 and 400 pc. It is thus crucial to quantify the heating of stellar discs and the efficiency in forming an in-situ halo also in more extreme conditions than those modeled so far, also because galactic stellar archeology seems to suggest that the Milky Way may have experienced an intense phase of star formation and significant gas turbulence at those epochs (Haywood et al. 2013; Snaith et al. 2014; Lehnert et al. 2014). If the inclusion of a gaseous component in the Milky Way disc seems inevitable for the simulation of accretion episodes that occurred at early times, the same is valid for modeling the accreted satellites. In the case of gas-rich satellites, the stellar fraction of the baryonic mass would be lower than 100%, as assumed in this paper and a certain fraction of the gaseous mass may be lost by tidal effects\/ram pressure stripping before being accreted and converted into new stars. This would naturally lead the accretions to be less \u201cstellar-rich\u201d than those simulated in the present paper, thus effectively reducing the amount of stars of extra-galactic origin that a satellite of a given mass can bring into the inner halo (i.e., inside 20 kpc) of a Milky Way-type galaxy. Finally, it is worth emphasizing that we are using idealized simulations of local mergers to quantify the response of a stellar disc to merger events that may have occurred several Gyr ago, when both the stellar masses of the Milky Way and of its satellite(s) were significantly different from those employed here. Because the properties of the Milky Way at higher z are still largely unknown (in terms of mass and effective radius, etc.) and similarly, those of the satellite galaxies, idealized simulations of local mergers are currently still one of the only viable ways to quantify the response of a stellar disc to accretion events of a given mass ratio, even for those occurring at higher z (with the limitations previously discussed, such as the lack of a gaseous component). Some N-body models similar to those analyzed in this paper have tried to mimic mergers at higher redshift, finding that the evolution and thickening of the disc, and its orbital parameters, do not depend on the redshift, but only on the mass-ratio of the merger (cf., for example, Figs 10 and 14 in Villalobos & Helmi 2008). This is robust evidence of the fact that what matters is the relative mass ratio of the interacting galaxies (see also Quinn et al. 1993; Bournaud et al. 2005, 2009; Qu et al. 2010), and the relative variation of their parameters, rather than the absolute values. As a consequence, at first approximation, local mergers can be re-scaled to also mimic mergers at higher z. This, we suggest caution to be taken due to the limitations of these models, such as, for example, the lack of gas, that can play an important role in the process. ","Citation Text":["Snaith et al. 2014"],"Citation Start End":[[1719,1737]]}
{"Identifier":"2021MNRAS.504.3316B__than_2000_Instance_6","Paragraph":"WASP-43b is the most heavily scrutinized phase curve, with four analyses of this data set already published (Stevenson et al. 2017; Mendon\u00e7a 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\u2019s phase offset and nightside temperature. Stevenson et al. (2017) initially reported only a 2\u03c3 upper limit on the nightside temperature of 650\u2009K, while all subsequent reanalyses (including ours) favour a significantly detectable nightside temperature of \u223c800\u2009K. As for the planet\u2019s phase offset, Stevenson et al. (2017) and May & Stevenson (2020) favour a larger phase offset (21 \u00b1 2\u2009\u00b0E) than Mendon\u00e7a et al. (2018) and Morello et al. (2019) (12 \u00b1 3\u2009\u00b0E and 11 \u00b1 2\u2009\u00b0E). 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 Mendon\u00e7a 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\u2009000 CPU hours (or 434\u2009d using our 12\u00d7 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 Mendon\u00e7a et al. (2018) and Morello et al. (2019). Ultimately, we cannot decide between these two discrepant offsets as the \u0394BIC between the two phase curve models for our preferred BLISS detector model is only 3.7 (insignificantly favouring the 20.4 \u00b1 3.6 offset from the v2 model). For reference, Stevenson et al. (2014b) found phase offsets ranging from roughly \u22126 to 17\u2009deg east in the Hubble\/WFC3 bandpass.","Citation Text":["Stevenson et al. (2017)"],"Citation Start End":[[1941,1964]]}
{"Identifier":"2021ApJ...922..196K__Erwin_&_Debattista_2013_Instance_1","Paragraph":"Due to their non-axisymmetry, bars can cause the gas to flow inwards (e.g., Athanassoula 1992; Regan et al. 1999; Sheth et al. 2000), thus enhancing star formation in the central regions of their host galaxies (S\u00e9rsic & Pastoriza 1967; Hawarden et al. 1986; Ho et al. 1997; Sheth et al. 2005; Coelho & Gadotti 2011; Ellison et al. 2011). As a result, barred galaxies have various stellar structures, such as nuclear rings, nuclear disks, inner rings, outer rings, and boxy peanut bulges (e.g., Combes & Gerin 1985; Athanassoula 1992; Sellwood & Wilkinson 1993; Knapen et al. 2002; Kormendy & Kennicutt 2004; Comer\u00f3n et al. 2010; Kim et al. 2012; Erwin & Debattista 2013; Buta et al. 2015; Emsellem et al. 2015, and references therein). These are dense regions in terms of stellar density. However, bars also develop sparse regions around themselves, usually producing a \u201c\u0398\u201d shaped region. These sparse regions are thought to be produced by the bar-driven secular evolution in a way that the bar sweeps out materials around it. Such sparse regions are called by different names, such as \u201clight deficit\u201d(Kim et al. 2016), \u201cstar formation desert\u201d (James & Percival 2015, 2016), and \u201cdark spacer\u201d (Buta 2017). In this respect, some barred galaxies have been found to be accompanied by a very faint disk or almost no disk around the bar (Gadotti & de Souza 2003; Gadotti 2008; \u0141okas 2021). Using H\u03b1 imaging, James et al. (2009) found that star formation is suppressed in the immediate vicinity of the bar (James & Percival 2016, 2018). The bar effectively transports gas into the central region once it is formed. Therefore, the star formation is truncated in the star formation desert. James & Percival (2016, 2018) measured the age of stellar populations and put constraints on the bar formation epoch. With zoom-in cosmological simulations, Donohoe-Keyes et al. (2019) found a truncated star formation and a lack of young stellar populations. However, they stressed that the interpretation can be complicated due to the radial migration of stars. Kim et al. (2016) found that the light deficit is more pronounced among galaxies with a longer and stronger bar. Later, Buta (2017) found that these sparse regions occur not only inside the inner ring, but also between the bar and the outer ring. He also found that the degree of the light deficit is a strong function of the bar strength, A\n2. Therefore, the bar-driven secular evolution is recorded in the light deficit and how it changes with redshift is worthy of investigation.","Citation Text":["Erwin & Debattista 2013"],"Citation Start End":[[646,669]]}
{"Identifier":"2020ApJ...892...58H__Davenport_et_al._2019_Instance_1","Paragraph":"Solar flares occur because of the explosive release of magnetic energy via the reconnection process in the solar atmosphere. Similar energetic events have been observed in other stars, especially M dwarfs like YZ CMi (dM4.5e), EVLac (dM3.5e), and AD Leo (Hawley & Pettersen 1991; Kowalski et al. 2010). The general idea is that for low-mass stars, the flare activities would be more frequent if the stellar masses are smaller. The level of flare effects depends also on the rotation periods (P) of the stars For example, the M-type dwarfs with P  10 days have been found to exhibit far more flare activity than those with P > 10 days, according to the Kepler observation (Lin et al. 2019). As P \u223c 1\/\n\n\n\n\n\n, with t being the stellar age on the basis of the gyrochronology relation (Skumanich 1972; Barnes 2010; van Saders et al. 2016), it is expected that the flare activity should be more frequent and powerful for young and fast-rotating stars (Davenport 2016; Stelzer et al. 2016; Davenport et al. 2019). Such time evolution of the stellar activities is important to the study of habitability of exoplanets orbiting around low-mass stars because of the potential impacts on the atmospheric loss and biological damage (Tarter et al. 2007; West et al. 2008; Segura et al. 2010; Armstrong et al. 2016). The large amount of high-precision photometric light curves produced by the Kepler mission (Borucki et al. 2010) is therefore ideal for statistical studies of stellar flares (Walkowicz et al. 2011). Following the discovery of superflares of solar-type stars by Maehara et al. (2015), there are increasing interests in topics related to exospace weather effects across a wide range of spectral types from A-type (Balona et al. 2015), G-type (Maehara et al. 2012, 2015; Wu et al. 2015), to M-type stars (Hawley & Pettersen 1991; Kowalski et al. 2010; Chang et al. 2017). There are two types of exoplanets of binary stars: the S-type for those in orbit around one of the stellar component, and the P-type for those orbiting around both components (Kaltenegger & Haghighipour 2013; Pilat-Lohinger et al. 2019). Their surface temperatures (and hence the condition of habitability) would be determined by the total radiative fluxes and magnetic activity of the binary systems. From this point of view, the binary systems are no less interesting because of the possible presence of the P-type or S-type habitable exoplanets.","Citation Text":["Davenport et al. 2019"],"Citation Start End":[[983,1004]]}
{"Identifier":"2017ApJ...844...61D__Millward_et_al._2013_Instance_1","Paragraph":"Now, we must consider a specific distribution of electrons. Ideally, with a well-constructed background-subtracted image, only recently injected coronal material should be visible in excess total brightness; hence, ideally, ne should be identically zero outside the CME. However, reality is not so kind as to clearly delineate between CME and non-CME material. The explosive eruption and ejection of coronal material may rearrange previously quiescent material. If such material is pushed into a line of sight that passes through the CME, then, due to the optically thin nature of the corona, it will be impossible to distinguish newly ejected material from newly arranged material. A further issue that must be confronted when choosing a functional form for ne is the lack of observation or direct measurement on the internal mass distribution of a CME, as well as imprecise knowledge on the morphology of the CME. For example, is a CME balloon-shaped with a circular cross section (Millward et al. 2013) or croissant-shaped with an elliptical cross section (Thernisien et al. 2006, 2009)? Are there one or more cavities in the CME, and if so, where are the cavities located, how wide are they relative to the CME, and how low is the cavity density relative to the edge density? Is the mass distributed symmetrically relative to the central axis of the CME? Does the mass distribution change with time, either by pileup at the CME leading edge or by continuous mass inflow from the lower atmosphere (Bein et al. 2013; DeForest et al. 2013a)? Nevertheless, in spite of such epistemic uncertainty on the internal mass distribution, if the shape of the CME is known\u2014either directly through polarimetry (de Koning & Pizzo 2011) or by assuming a shape (Thernisien et al. 2006; Millward et al. 2013)\u2014and hence the start and end points of the line-of-sight integral are known, one might choose a particularly simple distribution, namely, \n\n\n\n\n\n (Feng et al. 2013b). On the other hand, if both the CME mass distribution and morphology are poorly known, as was the case in the pre-STEREO era, then one might choose an even simpler distribution, that is, a point source,\n12\n\n\n\n\n\nwhere \u03b4 is the Dirac delta function and n0 is a constant. To enable direct comparison of our three-view mass determination with the two-view mass determination of CV09, we employ the point-source distribution.","Citation Text":["Millward et al. 2013"],"Citation Start End":[[984,1004]]}
{"Identifier":"2018MNRAS.479.5678F__Lucia_et_al._2006_Instance_1","Paragraph":"The crucial aspect we want to test in this work is whether the CR-IGIMF scenario, once embedded in a realistic galaxy formation context, can match the observed trends. To this purpose, we use mean star formation histories (SFHs) predicted by the GAEA semi-analytic model (Hirschmann et al. 2016) for $z$ = 0 galaxies in four different stellar mass bins, namely M\u22c6 \u223c 1012, 1011.5, 1010.5, and $10^{9.25} \\, {\\rm M}_\\odot$. These SFHs have been extracted from the reference GAEA realization and from runs implementing variable IMF approaches (F17 and F18). Galaxies more massive than \u223c1010.5 are typically characterized by an early peak of SFR, followed by a smooth decay. The epoch, height, and width of the peak depend on the final galaxy stellar mass, with more massive galaxies having an earlier, narrower, and higher peak (see also De Lucia et al. 2006). Lower mass galaxies have almost constant SFHs. For the mean SFH in each mass bin, we compute the mass-weighted global IMF using our CR-IGIMF grid and then the corresponding value of fdg. It is worth noting that in our estimate for fdg we only consider the $\\lt 1 \\, \\mathrm{M}_\\odot$ mass range, where the mass-weighted IMF coincides with the present-day mass function. Moreover, observational constraints are closer to luminosity-weighted IMFs. However, for very old stellar populations (i.e. massive ETGs), we do not expect this difference to change our conclusions significantly. In Fig. 3, we compare these fdg for UCR\/UMW = 1 (blue stars \u2013 for each mass bin, we plot separately the values corresponding to the mean SFHs extracted from the three GAEA realizations considered), with the observed values. We correct the observed values of fdg from LB13 to an aperture of 1 effective radius, whose properties we consider to be comparable to SAM predictions. To this aim, we assume an IMF radial profile with the same shape as that recently derived for one massive ETG by La Barbera et al. (2017). For each mass bin, the profile is rescaled, in the central region, in order to match the fdg value within the SDSS fibre aperture, as detailed in LB13. The correction, which is larger at the highest mass and negligible at the lowest mass,4 brings all fdg values from LB13 in the range from \u223c0.72 to \u223c0.8, as illustrated in Fig. 3 (see red dots with error bars, corresponding to the 18 stacked spectra of LB13). Remarkably, the aperture correction brings the observed fdg values within the range predicted by the CR-IGIMF.","Citation Text":["De Lucia et al. 2006"],"Citation Start End":[[835,855]]}
{"Identifier":"2016MNRAS.456.1195H__Tremonti_et_al._2004_Instance_1","Paragraph":"To further test the possible role of mass in driving the high emission-line velocity-widths observed in our targets, in Fig. 14 we plot W80,H\u03b1 as a function of the emission-line ratio log\u2009([N\u2009ii]\/H\u03b1). This emission-line ratio is a tracer of the metallicity of star-forming galaxies (e.g. Alloin et al. 1979; Denicol\u00f3, Terlevich & Terlevich 2002; Kewley & Dopita 2002), as well as an indicator for the source of ionizing radiation (e.g. Kewley et al. 2006; Rich et al. 2014). Furthermore, there is an observed relationship between mass and metallicity and therefore an expected relationship between this emission-line ratio and stellar mass (e.g. Lequeux et al. 1979; Tremonti et al. 2004; Maiolino et al. 2008; Stott et al. 2013). We can also make a crude prediction for the relationship between W80,H\u03b1 and mass, under the assumption that the line width is a tracer of the stellar velocity dispersion. Therefore, for a given stellar mass, we predict the position galaxies would be located in Fig. 14 by combining: (1) the observed z = 0.7 mass\u2013metallicity relation (following Maiolino et al. 2008) and (2) the observed mass\u2013velocity dispersion relationship for massive galaxies (following Bezanson, Franx & van Dokkum 2015). The majority of the star-forming galaxies appear to broadly follow the expected trend. Furthermore, the measurements from the stacked average emission-line profile (Fig. 11) are in agreement with the rough mass-driven prediction for the average mass of these galaxies [i.e. log\u2009(M\u22c6\/M\u2299) = 10.3; see Fig. 14]. In contrast, the AGN typically have higher log\u2009([N\u2009ii]\/H\u03b1) emission-line ratios and higher velocity-widths than the star-forming galaxy sample (also visible in the stacked profiles; Fig. 11) and a positive correlation is observed between these two quantities. Such a positive correlation has been shown to be a tracer of shocks and outflows in the ISM through IFS observations of AGN and star-forming galaxies (e.g. Ho et al. 2014; McElroy et al. 2015). Interestingly, the small number of star-forming galaxies with high line widths (i.e. W80 \u2273 400 km s\u22121) appear to follow the same relationship as the KASHz AGN, which may indicate that these galaxies also have a contribution from shocks and\/or host AGN that were not detected in the X-ray surveys. We clarify that some of the AGN targets have log\u2009([N\u2009ii]\/H\u03b1) emission-line ratios that could also be photoionized by H\u2009ii regions (see Fig. 14) and we will explore these ideas further when exploring the spatially resolved outflow kinematics and emission-line flux ratios of the KASHz AGN in future papers.","Citation Text":["Tremonti et al. 2004"],"Citation Start End":[[667,687]]}
{"Identifier":"2018AandA...618A..67C__Moriguchi_et_al._(2002)_Instance_1","Paragraph":"The close proximity in the sky of M 16 and M 17, two of the nearest giant HII regions of our galactic neighborhood lying at a similar distance from the Sun, naturally leads to the question of whether they are physically related, and whether they may share a common origin (Moriguchi et al. 2002; Oliveira 2008; Nishimura et al. 2017). Both giant HII regions are projected on the contour of a giant bubble-shaped structure, outlined in the distribution of HI and CO emission as first noted by Moriguchi et al. (2002). Evidence for triggered star formation in M 16 has been examined in detail by Guarcello et al. (2010), who concluded it had been induced externally, and not by the activity of its associated cluster NGC 6611. This suggests that the formation of M 16 and M 17 could have been triggered by the expansion of the bubble, powered by a previous generation of massive stars near its center, thus representing an example of triggered star formation on the scale of several tens of parsecs (Elmegreen 1998). Given the ages of the giant HII regions, the timescale of expansion of a wind-blown bubble, and the short lifetimes of massive stars, it is to be expected that the most massive members of that previous generation may have exploded as supernovae several Myr ago. The spatial dispersion of the members of the association that must have taken place progressively during its existence, combined with the distance of 2 kpc to the M 16\u2013M 17 complex and the large number of unrelated foreground and background stars in that general direction, would make it extremely difficult to identify even its currently hottest members still remaining on the main sequence. Moriguchi et al. (2002) noted the presence of O and early B stars in the area and proposed that they were part of a massive star population responsible for having caused the bubble, but a review of their properties shows them to be generally too bright to be at the distance of the bubble and the giant HII regions, and instead are more likely members of a foreground population.","Citation Text":["Moriguchi et al. 2002"],"Citation Start End":[[273,294]]}
{"Identifier":"2020MNRAS.494.5161C__Tadhunter_et_al._1998_Instance_1","Paragraph":"Cold gas kinematics of radio AGNs can tell us about the nature of feedback to their host galaxies. Both, radiative winds from central optical AGN, and radio jets can affect the gas kinematics in host galaxies. In this section, in order to understand the effect of both radiation and jet on gas kinematics, we compare the radio luminosity at 1.4\u2009GHz and [O\u2009iii] \u03bb5007 line luminosity (using the line fluxes from MPA-JHU group) for LERGs and HERGs. While radio luminosity at 1.4\u2009GHz can be used as tracer of jet power (Cavagnolo et al. 2010), [O\u2009iii] \u03bb5007 line luminosity can be used as a tracer for AGN bolometric luminosity (LaMassa et al. 2010). We find that there exists a weak but significant correlation (Kendall\u2019s tau=0.4, p=2.4 \u00d7 10\u221216) between the radio luminosities and the [O\u2009iii] \u03bb5007 line luminosities of LERGs (Fig. 9), which is consistent with earlier studies (Baum & Heckman 1989; Tadhunter et al. 1998), implying a common source of both luminosities. A similar correlation is noticed for HERGs but with higher [O\u2009iii] \u03bb5007 luminosity than LERGs for a particular value of radio luminosity, which could mean the extra energy in form of wind-driven by radiation from accretion disc to cause the turbulence in host galaxy. However, in our sample of radio AGNs, while for low radio power ($P_{1.4 \\rm GHz}\\, \\lt $ 1024.3 W\u2009Hz\u22121) LERGs and HERGs, there is almost no difference in distribution of centroid shifts in the absorption profiles relative to the systemic velocity, with all but one within \u00b1200\u2009km\u2009s\u22121, a larger number of intermediate radio power ($P_{1.4 \\rm GHz}\\, \\sim$ 1024.3\u22121026.0\u2009W\u2009Hz\u22121) LERGs show centroid velocity shifts greater than 200\u2009km\u2009s\u22121 compared to HERGs. Also the range of centroid velocity shift for LERGs is wider (\u2212479 to +356\u2009km\u2009s\u22121) than HERGs. Although, we have doubled the number of detections compared to our previous analysis (Chandola & Saikia 2017), the overall number of sources with H\u2009i absorption detections for HERGs used for kinematic analysis is still small. Hence there are large statistical uncertainties to draw a firm conclusion about H\u2009i gas kinematics in HERGs as a population.","Citation Text":["Tadhunter et al. 1998"],"Citation Start End":[[897,918]]}
{"Identifier":"2021ApJ...914..131Y__McKinney_&_Gammie_2004_Instance_1","Paragraph":"Let us first define some terminology before describing our results. \u201cOutflow\u201d means the flow with a positive radial velocity vr, including both \u201cturbulent outflow\u201d and \u201creal outflow.\u201d In the former case, the test particle will first move outward but will eventually return after moving outward for some distance. In the latter case, the test particle continues to flow outward and eventually escapes the outer boundary of the simulation domain. \u201cReal outflow\u201d consists of two components, i.e., jet and wind. The jet region is defined as the region occupied by the time-averaged magnetic field lines connected to the ergosphere (described by \n\n\n\n\n\n\nr\n\n\nerg\n\n\n\u2261\n\n\nr\n\n\ng\n\n\n+\n\n\n1\n\u2212\n\n\na\n\n\n2\n\n\n\n\ncos\n\n\n2\n\n\n\u03b8\n\n\n\n\n\nr\n\n\ng\n\n\n;\n\n\n Visser 2007) of the black hole. We note that this definition coincides with other notions of the jet based on velocity and magnetization, as seen even in early two-dimensional simulations (McKinney & Gammie 2004). Moreover, turbulence tends not to have a strong effect on these field lines, so that in both individual snapshots and a time-averaged sense, the field lines connecting to the ergosphere coincide with \u03b2  2 and \u03c3 > 1 (McKinney et al. 2012, see Figures 3 and 6). Thus, the jet region is bounded by the magnetic field line whose foot point is rooted at the black hole ergosphere with \u03b8 = 90\u00b0, i.e., the boundary between the black hole ergosphere and the accretion flow (refer to the red lines in Figure 5). In this case, all magnetic field lines in the jet region are anchored to the black hole ergosphere and thus can extract the spin energy of the black hole via the Blandford & Znajek mechanism (Blandford & Znajek 1977; hereafter BZ77) to power the jet (\u201cBZ-jet\u201d). Real outflows outside of this boundary are powered by the rotation energy of the accretion flow, and we call them wind. Note that our definition of wind adopted here is different from that adopted in some literature, where they require that the Bernoulli parameter of wind must satisfy Be > 0. We do not add this requirement because we find that for a nonsteady accretion flow, Be is not constant along the trajectories but usually increases outward, at least to the radius within which turbulence is well developed (Yuan et al. 2015). We find that even though Be is negative at a small radius, it can become positive when it propagates outward.","Citation Text":["McKinney & Gammie 2004"],"Citation Start End":[[909,931]]}
{"Identifier":"2020AandA...635A.186P__Nagy_(2018)_Instance_2","Paragraph":"The choice of \u03ba has some bearing on the value of Mej as calculated in Eq. (2). Although taken as a constant in various model light curves, the actual optical opacity varies with time (e.g. Chugai 2000; Nagy 2018), but tests have shown that that a constant opacity can be a reasonable assumption at early times (see Mazzali et al. 2000; Taddia et al. 2018). SN 2019bkc is a H\/He-deficient SN, so to select an appropriate value for \u03ba, we search the literature literature and consider only SNe of this type. For SNe Ia, where the large abundance of Fe-group elements is a source of significant opacity, Cappellaro et al. (1997) used \u03ba\u2004=\u20040.15 cm2 g\u22121, while Arnett (1982) used \u03ba\u2004=\u20040.08 cm2 g\u22121. The light curve of the broad lined SNe Ic 1997ef was modelled in Mazzali et al. (2000) using the same code as Cappellaro et al. (1997), this time a lower constant opacity of 0.08 cm2 g\u22121 was used owing to the lower Fe abundance in this SN type. The average opacity of model CC-SNe light curves was investigated by Nagy (2018), where it was found that \u03ba\u2004=\u20040.1 cm2 g\u22121 was suitable for their H\/He-deficient models. Taddia et al. (2018) demonstrated that fitting the light curves of H-deficient CC-SNe with a simple analytical model and a constant opacity \u03ba\u2004=\u20040.07 cm2 g\u22121 returned physical parameters consistent with more complex hydrodynamic light curve models. Prentice et al. (2019) also found that using \u03ba\u2004=\u20040.07 cm2 g\u22121 gave comparable ejecta masses to both photospheric phase and nebular phase spectroscopic modelling of H-deficient CC-SNe, with the exception of some gamma-ray burst SNe. From Sect. 5 we find that the outer ejecta is not rich in Fe-group elements, and there is no indication of helium. From the early-nebular spectra in Sect. 4.2 we see no indication of strong Fe\u202fII or Fe\u202fIII line emission, demonstrating that the ejecta is poor in Fe-group elements relative to normal SNe Ia and is more like H\/He-deficient CC-SNe. This justifies the use of \u03ba\u2004=\u20040.07 cm2 g\u22121, but to show the range of possible Mej, we consider the range 0.07\u20130.1 cm2 g\u22121.","Citation Text":["Nagy (2018)"],"Citation Start End":[[1005,1016]]}
{"Identifier":"2016MNRAS.463..655K__Sunyaev_&_Zeldovich_1970_Instance_1","Paragraph":"The Planck cosmic microwave background (CMB) mission (Planck Collaboration I 2011) has for the first time allowed us to retrieve all-sky maps (Hill & Spergel 2014; Planck Collaboration XXII 2015; Khatri 2016) of the Sunyaev-Zeldovich (SZ) effect (Zeldovich & Sunyaev 1969). The chief advantage of Planck over other past and current satellite and ground-based missions is its multiple frequency channels covering the low-frequency Rayleigh\u2013Jeans as well as the higher frequency Wien region of the CMB spectrum making the separation of the thermal SZ from CMB and foregrounds feasible. As the blackbody photons from the CMB interact with hot intergalactic\/intracluster medium (ICM) travelling from the last scattering surface (Peebles & Yu 1970; Sunyaev & Zeldovich 1970) to us, Compton scattering upscatters a fraction of photons to higher energy creating a distortion from the Planck spectrum $I_{\\nu }^{\\rm Planck}$ (Sunyaev & Zeldovich 1972). The change in the intensity $\\Delta I_{\\nu }=I_{\\nu }-I_{\\nu }^{\\rm Planck}$ of the CMB radiation is given by (Zeldovich & Sunyaev 1969)\n\n(1)\n\n\\begin{eqnarray}\n\\Delta I_{\\nu }=y\\frac{2 h\\nu ^3}{c^2}\\frac{x e^x}{(e^x-1)^2}\\left[x\\left(\\frac{e^x+1}{e^x-1}\\right)-4\\right],\n\\end{eqnarray}\n\nwhere $x=\\frac{h\\nu }{{{k_{\\rm B}}}T}$, T = 2.725(1 + z) is the CMB temperature at redshift z, \u03bd = \u03bd0(1 + z) is the frequency of CMB photon at redshift z, \u03bd0 is the observed frequency today (z = 0), h is the Planck constant, kB is Boltzmann constant and c is the speed of light. The amplitude of the distortion, y, is proportional to the integral of the pressure P along the line of sight,\n\n(2)\n\n\\begin{eqnarray}\ny=\\frac{{{\\sigma _{\\rm T}}}}{{{m_{\\rm e}}}c^2} \\int {{\\rm d}}s\\, {{n_{\\rm e}}}{{k_{\\rm B}}}{{T_{\\rm e}}},\n\\end{eqnarray}\n\nwhere Te and ne are the electron temperature and electron number density, respectively, in the ICM plasma, me is the mass of the electron, \u03c3T is the Thomson scattering cross-section, and s is the distance coordinate along the line of sight.","Citation Text":["Sunyaev & Zeldovich 1970"],"Citation Start End":[[744,768]]}
{"Identifier":"2020MNRAS.495.1641S__Mukherjee_et_al._2015_Instance_1","Paragraph":"As expected, the orbital parameters obtained from this fit are fully consistent with those reported in Table 1, since the long baseline of the data set guarantees to clearly disentangle orbital phase modulation from and phase variation due to the NS spin frequency changes. From the best-fitting parameter of each frequency phase component we infer the dipolar magnetic moment \u03bcF \u2243 4.1(4) \u00d7 1026 and \u03bc1h \u2243 4.1(5) \u00d7 1026 G\u2009cm3, for the fundamental and the first harmonic, respectively. Adopting the Friedman-Pandharipande-Skyrme (FPS) equation of state (see e.g. Friedman & Pandharipande 1981; Pandharipande & Ravenhall 1989) for a 1.4 M\u2299 NS, we estimate an NS radius of R = 1.14 \u00d7 106 cm, and a moment of inertia I \u2243 1.45 \u00d7 1045 g\u2009cm2. This corresponds to an equatorial magnetic field of Beq = \u03bc\/R3 = 2.8(3) \u00d7 108 \u2243 2.8(3) \u00d7 108 G. This value is consistent both with the estimation reported in Sanna et al. (2018d) from spin frequency equilibrium considerations for the X-ray pulsar, and the average magnetic field of known AMXPs (Mukherjee et al. 2015). It is noteworthy that the value of Beq obtained here is likely overestimated due to a non-accurate estimation of mass accretion rate in case of spin-down regime. In fact, as discussed by Klu\u017aniak & Rappaport (2007), in this regime the angular momentum deposited in the disc by the pulsar torques leads to a substantial additional dissipation of energy in the disc, hence, the mass accretion rate estimated from equation (17) results on an upper limit of the quantity. However, since it is not quite clear what fraction of this additional dissipation of energy is released in non-thermal process and what is kept inside the disc, it is complicated to quantify the discrepancy. We started approaching the problem considering the worst case scenario at which the whole spin-down luminosity of the pulsar $I\\omega \\dot{\\omega }$ is transferred and released by the accretion disc. With this hypothesis, we subtract the rotation luminosity from observed luminosity in equation (17) and we then fitted the frequency phase delays. We obtained that, taking into account the correction on the mass accretion rate, the NS magnetic field decreases roughly by 15 per\u2009cent, reaching the value Beq = 2.4(4) \u00d7 108 G. Both the properties of the source during the outburst and the magnetic field estimation obtained from the studies of the spin-down like behaviour of the phase delays do not seem to suggest IGR J17591\u20132342 as an exception among the AMXPs.","Citation Text":["Mukherjee et al. 2015"],"Citation Start End":[[1031,1052]]}
{"Identifier":"2018ApJ...853...84Z__Zhu_et_al._2006_Instance_1","Paragraph":"A current topic of active debate is to what extent galaxies are affected by the assembly history of their host halos. The stochasticity in the complex baryonic physics may act to erase the record of halo assembly history. If, however, the galaxy properties closely correlate with the halo formation history, this would lead to a dependence of the galaxy content on large-scale environment and a corresponding clustering signature. This effect has commonly been referred to as galaxy assembly bias both colloquially and in the literature, and we adopt this distinction here. We stress, however, that what is referred to here is the manifestation of halo assembly bias in the galaxy distribution. The predictions for galaxy assembly bias have been explored with simulated galaxies (Zhu et al. 2006; Croton et al. 2007; Reed et al. 2007; Zu et al. 2008; Zentner et al. 2014; Chaves-Montero et al. 2015; Bray et al. 2016; Romano-Diaz et al. 2017). Detecting (galaxy) assembly bias is much more challenging since halo properties are not directly observed. Observational studies of assembly bias have generally produced mixed results. There have been several suggestive detections in observations (Berlind et al. 2006; Yang et al. 2006; Tinker et al. 2008b, 2017b; Wang et al. 2008, 2013b; Cooper et al. 2010; Lacerna et al. 2014a; Hearin et al. 2015; Watson et al. 2015; Miyatake et al. 2016; Saito et al. 2016; Zentner et al. 2016; Montero-Dorta et al. 2017; Tojeiro et al. 2017), while numerous other studies indicate the impact of assembly bias to be small (Abbas & Sheth 2006; Blanton & Berlind 2007; Croton & Farrar 2008; Tinker et al. 2008a, 2011; Deason et al. 2013; Lacerna et al. 2014b; Lin et al. 2016; Vakili & Hahn 2016; Zu & Mandelbaum 2016; Dvornik et al. 2017). The situation is further complicated as various systematic effects can mimic the effects of assembly bias (e.g., Campbell et al. 2015; Zu et al. 2016; Busch & White 2017; Lacerna et al. 2017; Sin et al. 2017; Tinker et al. 2017a; Zu & Mandelbaum 2017) and the evidence for assembly bias to date remains inconclusive and controversial.","Citation Text":["Zhu et al. 2006"],"Citation Start End":[[780,795]]}
{"Identifier":"2018MNRAS.475.2787G__Leauthaud_et_al._2012_Instance_2","Paragraph":"In order to separate the role of different physical mechanisms in galaxy evolution, a number of studies have constrained stellar-to-halo mass (SHM) relations and ratios as a function of time using the abundance matching technique (e.g. Behroozi, Conroy & Wechsler 2010a; Moster et al. 2010; Behroozi, Conroy & Wechsler 2010b), the conditional luminosity function technique proposed by Yang, Mo & Van Den Bosch (2003), the halo occupation distribution (HOD) formalism (e.g. Berlind & Weinberg 2002; Kravtsov et al. 2004; Moster et al. 2010), and by combining the HOD, N-body simulations, galaxy clustering, and galaxy\u2013galaxy lensing techniques (e.g. Leauthaud et al. 2012; Coupon et al. 2015). Distinguishing the properties of central galaxies from those of satellite galaxies in studies based only on the distribution of luminosity or stellar mass is challenging (e.g. George et al. 2011). By combining several observables and techniques (e.g. HOD, galaxy\u2013galaxy lensing, and galaxy clustering) one can probe a global SHM relation for central galaxies and satellite galaxies (e.g. Leauthaud et al. 2012; Coupon et al. 2015). Coupon et al. (2015), for example, used multiwavelength data of \u223c60\u2009000 galaxies with spectroscopic redshifts in the Canada France Hawaii Telescope Lensing Survey (CFHTLenS) and Vimos Public Extragalactic Redshift Survey (VIPERS) field to constrain the relationship between central\/satellite mass and halo mass, characterizing the contributions from central and satellite galaxies in the SHM relation. In this paper, we directly identify the BGGs using their precise redshifts and estimate stellar masses using the broad-band spectral energy distribution (SED) fitting technique (Ilbert et al. 2010) as used by Coupon et al. (2015). We utilize the advantages of the X-ray selection of galaxy groups and a wealth of multiwavelength, high signal-to-noise ratio observations such as the UltraVISTA survey in the COSMOS field (Laigle, Capak & Scoville 2016) to investigate the SHM relation for the central galaxies over 9 billion years. We aim to quantify the intrinsic (lognormal) scatter in stellar mass at fixed redshift in observations and compare them to the recently implemented semi-analytic model (SAM) by Henriques et al. (2015).","Citation Text":["Leauthaud et al. 2012"],"Citation Start End":[[1081,1102]]}
{"Identifier":"2015ApJ...799...99K__Coelho_&_Gadotti_2011_Instance_1","Paragraph":"The nonaxisymmetric potential of a bar induces large-scale streaming motions in stars and gas into the central part of the galaxy (e.g., Athanassoula 1992a, 1992b; Sellwood & Wilkinson 1993). Being dissipative, the gas loses angular momentum and energy and flows inward toward the galactic center (Knapen et\u00c2 al. 1995; Regan et\u00c2 al. 1999; Sheth et\u00c2 al. 2000, 2002), accumulating in the central \u00e2\u0088\u00bc1\u00c2 kpc of galaxies (e.g., Sakamoto et\u00c2 al. 1999; Sheth et\u00c2 al. 2005). The accumulation of gas in the central parts leads to high levels of circumnuclear star formation activity (S\u00c3\u00a9rsic & Pastoriza 1965; Hawarden et\u00c2 al. 1986; Devereux 1987; Martin 1995; Ho et\u00c2 al. 1997; Sheth et\u00c2 al. 2000, 2005; Gadotti & dos Anjos 2001; Ellison et\u00c2 al. 2011; Coelho & Gadotti 2011; Wang et\u00c2 al. 2012); this circumnuclear star formation is often in the shape of nuclear rings (Knapen et\u00c2 al. 2002; Comer\u00c3\u00b3n et\u00c2 al. 2010; Kim et\u00c2 al. 2012; Seo & Kim 2013) and nuclear star clusters (B\u00c3\u00b6ker et\u00c2 al. 2002, 2004, 2011). Such star formation activities may help to create disky pseudobulges (Kormendy & Kennicutt 2004; Sheth et\u00c2 al. 2005; Athanassoula 2005; Debattista et\u00c2 al. 2006). Bars are the primary mechanism for transporting gas from a few kiloparsec scale to the central kiloparsec. However, there have been mixed answers to the question of whether the presence of a bar and active galactic nucleus (AGN) activity are connected. Some studies find weak statistical links among AGN activity and the presence of bars (e.g., Arsenault 1989; Knapen et\u00c2 al. 2000; Laine et\u00c2 al. 2002; Laurikainen et\u00c2 al. 2004a; Coelho & Gadotti 2011), whereas others find little or no link (e.g., Moles et\u00c2 al. 1995; McLeod & Rieke 1995; Mulchaey & Regan 1997; Ho et\u00c2 al. 1997; Hunt & Malkan 1999; Martini et\u00c2 al. 2003; Lee et\u00c2 al. 2012; Cisternas et\u00c2 al. 2013). While bar torques drive gas inside the bar corotation inward, they push the gas between the corotation and outer Lindblad resonance (OLR) outward (Combes 2008; Kubryk et\u00c2 al. 2013).","Citation Text":["Coelho & Gadotti 2011"],"Citation Start End":[[743,764]]}
{"Identifier":"2019ApJ...874L..32C__Stockton_et_al._1994_Instance_1","Paragraph":"Cygnus A, at z = 0.0562, is 10 times closer than the next radio galaxy of similar radio luminosity.4\n\n4\nRadio luminosity >1045 erg s\u22121.\n The nuclear regions in Cygnus A have been observed extensively at radio through X-ray wavelengths (Carilli & Barthel 1996). The inner few arcseconds is a complex mix of optically obscuring dust clouds (Vestergaard & Barthel 1993; Whysong & Antonucci 2004; Lopez-Rodriguez et al. 2014; Merlo et al. 2014), atomic gas seen in narrow line emission (Stockton et al. 1994; Taylor et al. 2003), H i 21 cm absorption toward the inner radio jets, with a neutral atomic column density >1023 cm\u22122, depending on H i excitation temperature (Struve & Conway 2010), polarized, broad optical emission lines due to scattering by dust (Antonnuci et al. 1994; Ogle et al. 1997), and a highly absorbed hard X-ray spectrum with a total gas column density of \u223c3 \u00d7 1023 cm\u22122 (Ueno et al. 1994; Reynolds et al. 2015). VLBI radio observations at 0.05 mas resolution reveal highly collimated jets originating on scales \u223c200 times the Schwarzschild radius (Boccardi et al. 2016). Tadhunter et al. (2003), derive a black hole mass of 2.5 \u00b1 0.7 \u00d7 109 M\u2299 from HST and Keck spectroscopy of Pa-\u03b1 and [O iii], and conclude that Cygnus A contains an AGN with a bolometric luminosity of order 1046 erg s\u22121, comparable to high redshift quasars (Runnoe et al. 2012). This AGN is highly obscured in the optical due to dust along our line of sight, with Av > 50 magnitudes, based on near-IR spectroscopy (Imanishi & Ueno 2000). Studies of the mid- to far-IR spectral and polarization properties have led to a model of a clumpy, dusty torus obscuring the AGN in Cygnus A, with a radius of at least 130 pc, although these conclusions are based on spatially integrated properties; these observations did not have the spatial resolution to resolve the torus, and hence are partially contaminated by emission from the radio core-jet (Privon et al. 2012; Lopez-Rodriguez et al. 2018).","Citation Text":["Stockton et al. 1994"],"Citation Start End":[[483,503]]}
{"Identifier":"2015ApJ...815..129S__Shen_et_al._2011_Instance_2","Paragraph":"The mass accretion onto the black hole is important for a better understanding of AGN evolution. The Eddington ratio, the ratio between the AGN bolometric luminosity and the Eddington luminosity (Lbol\/LEdd), provides insight into the black hole growth because the bolometric luminosity reflects the mass accretion rate. We show AGN bolometric luminosity versus black hole mass for our sample of broad-line AGNs in the different redshift bins in the left panel of Figure 5. The different X-ray surveys are shown with different symbols as labeled. The dotted reference lines indicate constant Eddington ratios of 1, 0.1, 0.01, and 0.001. Our sample of broad-line AGNs covers the black hole mass range 7.0  log MBH\/M\u2299  9.5 and the bolometric luminosity range 43  log Lbol  47 with a wide dispersion in the Eddington ratio distribution. For comparison, we show published observations in the same redshift range from the literature in the right panel of Figure 5 (Gavignaud et al. 2008; Merloni et al. 2010; Shen et al. 2011; Nobuta et al. 2012; Matsuoka et al. 2013). The SDSS quasar sample (gray points; Shen et al. 2011) is limited to the high-mass and high-luminosity regime because the SDSS detection limit corresponds to a luminosity of log Lbol \u223c 46 at z \u223c 1. Compared to the SDSS quasar sample, our sample of broad-line AGNs shows a wider dispersion in the black hole mass, AGN bolometric luminosity, and Eddington ratio distribution, consistent with previous studies on deep AGN samples (Gavignaud et al. 2008; Merloni et al. 2010; Nobuta et al. 2012; Matsuoka et al. 2013), which fill in the low-mass and low-luminosity region. The figure shows contours at the 1\u03c3 level, together with the literature data, except the SDSS quasar sample. The figure also reveals that only a small number of AGNs exceed the Eddington limit by a small amount. AGNs with similar black hole masses show a broad range of bolometric luminosities spanning about two orders of magnitude, indicating that the accretion rate of black holes is widely distributed. This suggests that the AGN cosmic downsizing phenomenon could be explained by some more-massive black holes with low accretion rates, which are relatively fainter than less-massive black holes with efficient accretion. Lusso et al. (2012) suggest that AGNs show higher Eddington ratios at higher redshift at any given MBH, and the Eddington ratio increases with bolometric luminosity. We confirm that there is a tendency for low-luminosity AGNs (log Lbol \u2272 45.5) with less-massive black holes (log MBH\/M\u2299 \u2272 8) to have lower Eddington ratios than high-luminosity AGNs (log Lbol \u2273 45.5) with massive black holes (log MBH\/M\u2299 \u2273 8), consistent with Lusso et al. (2012). It is important to note that, when comparing with results in the literature, one should take into account the different methods of spectral line fitting and correction for bolometric luminosities. Nevertheless, they show similar distributions of the accretion rate of black holes over a wide range, consistent with previous studies.","Citation Text":["Shen et al. 2011"],"Citation Start End":[[1101,1117]]}
{"Identifier":"2020ApJ...904..185O__Kratter_et_al._2008_Instance_1","Paragraph":"Recently, rotationally supported disks have been found not only in Class I sources but also in some Class 0 sources (e.g., Murillo et al. 2013; Yen et al. 2013, 2017; Ohashi et al. 2014; Tobin et al. 2015, 2016b, 2016a; Seifried et al. 2016; Aso et al. 2017; Lee et al. 2017; Okoda et al. 2018). In spite of these extensive studies, there is still controversy regarding when and how a disk structure is formed around a newly born protostar. Moreover, the disk formation process has been revealed to be much more complicated for binary and multiple cases, both in observations (Dutrey et al. 2014; Tokuda et al. 2014; Takakuwa et al. 2014, 2017; Tobin et al. 2016b, 2016a; Boehler et al. 2017; Artur de la Villarmois et al. 2018; Alves et al. 2019) and in numerical simulations (e.g., Bate & Bonnell 1997; Kratter et al. 2008; Fateeva et al. 2011; Shi et al. 2012; Ragusa et al. 2017; Satsuka et al. 2017; Price et al. 2018; Matsumoto et al. 2019). For instance, circumbinary\/circummultiple disk structures with a spiral structure as well as a circumstellar disk for each component are reported (e.g., Tobin et al. 2016a; Takakuwa et al. 2017; Artur de la Villarmois et al. 2018; Matsumoto et al. 2019; Alves et al. 2019). In addition, it is not clear how molecules are processed during the disk formation process and what kinds of molecules are finally inherited by protoplanetary disks and potentially by planets. Understanding these processes is crucial, as they will provide important constraints on the initial physical and chemical conditions for the planetary system formation study. In this context, physical and chemical structures and their mutual relation for disk-forming regions of low-mass protostellar sources have been investigated with the Atacama Large Millimeter\/submillimeter Array (ALMA; e.g., Sakai et al. 2014b, 2014a; Oya et al. 2016, 2017, 2018, 2019; Imai et al. 2016, 2019; Jacobsen et al. 2019). These studies reveal that infalling envelopes and rotationally supported disks are not smoothly connected to each other either in physical structure or in chemical composition, unlike previous expectations.","Citation Text":["Kratter et al. 2008"],"Citation Start End":[[805,824]]}
{"Identifier":"2019ApJ...881..151L__Shen_et_al._2011_Instance_1","Paragraph":"For confined flares, the key question is the factor determining their confined character. Wang & Zhang (2007) analyzed eight X-class flares and found that confined events occur closer to the magnetic center and eruptive events tend to occur close to the edge of active regions (ARs), implying that the strong external field overlying the AR core is probably the main reason for the confinement. Similar results have been found by Baumgartner et al. (2018) based on a statistical analysis of 44 flares during 2011\u20132015. Amari et al. (2018) suggested that the role of the magnetic cage crucially affects the class of eruption\u2014either confined or eruptive. To date, many studies have drawn the consistent conclusion that the decay index of the potential strapping field determines the likelihood of ejective\/confined eruptions (Green et al. 2002; Shen et al. 2011; Yang et al. 2014; Chen et al. 2015; Thalmann et al. 2015; Li et al. 2018a). Previous studies showed that the torus instability of a magnetic FR occurs when the critical decay index reaches 1.5 (Bateman 1978; Kliem & T\u00f6r\u00f6k 2006). Zuccarello et al. (2015) performed a series of numerical magnetohydrodynamics (MHD) simulations of FR eruptions and found that the critical decay index for the onset of the torus instability lies in the range 1.3\u20131.5. Another factor determining whether a flare event is CME-eruptive or not is the non-potentiality of ARs including the free magnetic energy, relative helicity, and magnetic twists (Falconer et al. 2002, 2006; Nindos & Andrews 2004; Tziotziou et al. 2012). Sun et al. (2015) suggested that AR eruptiveness is related to the relative value of magnetic non-potentiality over the constraint of a background field. However, in the statistical study of Jing et al. (2018), the unsigned twist number of a magnetic FR plays little role in differentiating between confined and ejective flares, and the decay index of the potential strapping field above the FR discriminates them well.","Citation Text":["Shen et al. 2011"],"Citation Start End":[[843,859]]}
{"Identifier":"2016MNRAS.459.1602L__Vasudevan_et_al._2009_Instance_1","Paragraph":"The bolometric luminosity is then computed as the sum of the disc luminosity and the 2\u201310 keV X-ray luminosity. The \u03bbEdd, given by \u03bbEdd = Lbol\/LEdd, as a function of the MBH is plotted on the central panel of Fig. 5, while the \u03bbEdd distribution is shown in the horizontal histogram. Fig. 12 shows the 2\u201310\u2009keV X-ray bolometric correction ($\\kappa _{\\rm 2{\\rm -}10\\,keV}=L_{\\rm bol}\/L_{\\rm 2{\\rm -}10\\,keV}$) versus Lbol (left-hand panel) and the \u03bbEdd (right-hand panel). Tight correlations between the $\\kappa _{\\rm 2{\\rm -}10\\,keV}$ and Lbol (Pearson correlation coefficient r = 0.71) as well as the \u03bbEdd (Pearson correlation coefficient r = 0.78) are found in our sample as clearly shown in Fig. 12. A similar trend, that the sources with higher bolometric luminosities have higher bolometric corrections, was found in Lusso et al. (2012). A relation between the $\\kappa _{\\rm 2{\\rm -}10\\,keV}$ and \u03bbEdd is also found in several previous works (Vasudevan & Fabian 2007, 2009; Vasudevan et al. 2009; Lusso et al. 2012). The tight $\\kappa _{\\rm 2{\\rm -}10\\,keV}{\\rm -}L_{\\rm bol}$ relation shown in Fig. 12 as well as the correlation between $\\kappa _{\\rm 2{\\rm -}10\\,keV}$ and \u03bbEdd, rather than intrinsic, can however be a natural outcome of our flux-limited sample, which will lead to a correlation between the monochromatic optical luminosity Lmon,opt and the X-ray luminosity $L_{\\rm 2{\\rm -}10\\,keV}$. We tested the significance of the correlation by means of simulations. We generated a sample which the log\u2009N \u2212 log\u2009S of both the X-ray and optical flux follows a power-law form, and has a redshift distribution similar to our AGN sample. We then applied the flux cut-off, f0.5\u201310\u2009keV > 1 \u00d7 10\u221215\u2009erg\u2009s\u22121\u2009cm\u22122 and R4681  22\u2009mag, on the simulated sources. Using the Lmon,opt and $L_{\\rm 2{\\rm -}10\\,keV}$ of the simulated data we obtain a Pearson correlation coefficient value r of \u223c0.57, which is smaller than the correlation coefficient found in the observed data (r \u223c 0.71). Thus the observed correlations in our sample should not be purely due to the selection bias. This result is also consistent with the conclusion that the strong correlation between the monochromatic UV luminosity at 2500\u00c5 and the X-ray luminosity at 2 keV is independent of redshift (Risaliti & Lusso 2015). The distributions of the Lbol and $\\kappa _{\\rm 2{\\rm -}10\\,keV}$ are shown in the left-hand panel of Fig. 12 (top panel: Lbol, right-hand panel: $\\kappa _{\\rm 2{\\rm -}10\\,keV}$). The $\\kappa _{\\rm 2{\\rm -}10\\,keV}$ found in our sample, with a mean value of \u223c18, are consistent with the values found in previous works (e.g. Vasudevan & Fabian 2007; Lusso et al. 2012).","Citation Text":["Vasudevan et al. 2009"],"Citation Start End":[[978,999]]}
{"Identifier":"2020AandA...641A..67S__Starck_et_al._(2000)_Instance_1","Paragraph":"The deconvolution of large galaxy survey images requires that the spatial variation of the point spread function (PSF) across the field of view is taken into account. The PSF field is usually estimated beforehand through parametric models and simulations as in Krist et al. (2011) or is directly estimated from the (noisy) observations of stars in the field of view (Bertin 2011; Kuijken et al. 2015; Zuntz et al. 2018; Mboula et al. 2016; Schmitz et al. 2020). Even when the PSF is known perfectly, this ill-posed deconvolution problem is challenging, in particular because of the size of the image that is to be processed. Starck et al. (2000) proposed an object-oriented deconvolution that consists of first detecting galaxies and then deconvolving each object independently, taking the PSF at the position of the center of the galaxy into account (but not the variation in the PSF field at the galaxy scale). Following this idea, Farrens et al. (2017) introduced a space-variant deconvolution approach for galaxy images that is based on two regularization strategies: using either a sparse prior in a transformed domain (Starck et al. 2015a), or trying to learn without supervision a low-dimensional subspace for galaxy representation using a low-rank prior on the recovered galaxy images. When a sufficient number of galaxies were processed jointly (more than 1000), the authors found that the low-rank approach provided significantly lower ellipticity errors than sparsity. This illustrates the importance of learning adequate representations for galaxies. To proceed in learning, supervised deep-learning techniques that use databases of galaxy images might be employed to learn complex mappings that might regularize our deconvolution problem. Deep convolutional architectures have also proved to be computationally efficient in processing large numbers of images when the model has been learned. They are therefore promising in the context of modern galaxy surveys.","Citation Text":["Starck et al. (2000)"],"Citation Start End":[[625,645]]}
{"Identifier":"2021AandA...656A..94G__Gronow_et_al._2021_Instance_1","Paragraph":"Major differences between our approach and a full re-calculation of the hydrodynamics were not expected since the changes in the 14N and 22Ne abundances at the different metallicities do not alter the energy release in the hydrodynamic simulations significantly. The situation is different for deflagrations where the buoyancy, and therefore the Rayleigh-Taylor instabilities, depend on Ye. In contrast to detonations, the propagation of a deflagration front is thus affected by the metallicity (e.g., Meakin et al. 2009). The assumption we made here is confirmed by the comparison of the models presented in Table 1. Model M2a is taken from Gronow et al. (2020). The model was calculated at zero metallicity and has a total mass of 1.05\u2006M\u2299 with a He shell of 0.07\u2006M\u2299 at He ignition. Model M10_05_1, on the other hand, has a similar mass configuration, though it was calculated at solar metallicity (Model M10_05 in Gronow et al. 2021). Model M2a_pp is the same model as Model M2a, but the postprocessing step was calculated with solar metallicity instead of zero metallicity. An inspection of the abundances of Models M2a_pp and M10_05_1 at t\u2004=\u2004100\u2006s after He detonation ignition shows that the results of the postprocessing step with varying metallicities are in reasonably good agreement with a full re-calculation of the hydrodynamic model. The maximum difference in the yields produced in the core detonation is only 10%, while the maximum difference is 50% in the He detonation (excluding 12C in both). However, differences in the yields produced in the He detonation can in part be attributed to the slightly different setups of Model M2a (and therefore Model M2a_pp) and Model M10_05_1 at the beginning of the relaxation simulation, with the differences in the total and shell masses being less than 1% (see Gronow et al. 2021 for an explanation of the difference). Generally, the contribution of the yields from the He detonation to the total nucleosynthetic yields are small compared to those of the core detonation. Our approach is thus sufficient to derive nucleosynthetic yields for GCE studies. It saves significant computational costs as additional 3D hydrodynamical simulations of the explosion do not need to be carried out. Nevertheless, there might be slight differences visible in the observables because they are sensitive to the products of the He shell detonation (H\u00f6flich et al. 1996; Nugent et al. 1997; Kromer et al. 2010).","Citation Text":["Gronow et al. 2021"],"Citation Start End":[[916,934]]}
{"Identifier":"2019AandA...630A.123K__Kohutova_&_Verwichte_2016_Instance_2","Paragraph":"The coronal rain plasma can be distinguished from the prominence material by looking at their trajectories and average speeds. The timescale on which the coronal rain forms following the heating onset is much shorter than for the quiescent scenario; in the studied event condensations appear 10 min after the reconnection event, whereas observations of quiescent rain suggests it recurs in the same loop of the order of hours (Antolin & Rouppe van der Voort 2012; Kohutova & Verwichte 2016). The period of the loop heating-condensation cycle in the quiescent scenario is equivalent to the time it takes for the sustained footpoint heating to refill the loop sufficiently with evaporated plasma to reach the thermally unstable regime, after the loop has been evacuated by the previous coronal rain event. This short timescale for coronal rain formation is likely a consequence of the heating input being much greater than in the quiescent case and of the short length of the studied loop. The 1D numerical simulations suggest that although the loop length is a contributing factor, the heating input is the main factor affecting the coronal rain formation timescale (Froment et al. 2018). The thermal instability in the case associated with magnetic reconnection is also more concentrated spatially and only a certain fraction of the loop with a cross section of around 5 Mm width becomes unstable. This implies that the heating that triggers the thermal instability is more localised and only affects a small number of field lines in the loop. As the thermal conduction acts predominantly along the magnetic field, most of the matter and energy transport occurs along the affected field lines. Comparing this to the quiescent scenario, the typical width of the loop bundles observed to undergo condensation formation is around 10\u221215 Mm (Antolin & Rouppe van der Voort 2012; Kohutova & Verwichte 2016), and in some cases, reaching 40 Mm (Auch\u00e8re et al. 2018; Froment et al. 2019).","Citation Text":["Kohutova & Verwichte 2016"],"Citation Start End":[[1874,1899]]}
{"Identifier":"2019MNRAS.484.4083H__McPherron,_Russell_&_Aubry_1973_Instance_1","Paragraph":"According to the Dst index reconstructed by Love, Hayakawa & Cliver (2019), the interval of the telegraphic glitches taking place between \u223c13:35\u2009ut and \u223c17:20\u2009ut corresponds to the storm main phase (Dst \u223c \u2212320\u2009nT to \u223c \u2212570\u2009nT). This interval also corresponds to that of the low-latitude aurorae witnessed at many points at \u00b140\u00b0 MLAT. A large-amplitude of the GIC is induced by magnetospheric and\/or ionospheric current system. One possible cause of the telegraphic disturbance is the storm-time ring current that developed between \u223c11:40\u2009ut and \u223c17:40\u2009ut. Kappenman (2004) has shown that the magnitude of GICs flowing in the Japanese power grid increases with the magnitude of the Dst index. Another possible cause is the substorm\u2013current wedge system that consists of field-aligned current and the tail current (McPherron, Russell & Aubry 1973). Downward field-aligned current is connected to the dawnside ionosphere, and upward field-aligned current is connected to the duskside ionosphere. The pair of field-aligned currents causes magnetic disturbances at mid- and low- latitudes (Pytte, Mcpherron & Kokubun 1976). The magnetic disturbance recorded at Tokyo (Fig. 5) shows a decrease in the H-component and a positive excursion of the D-component during the interval of \u223c22\u201323 LT (\u223c13\u201314\u2009ut). It is probable that Tokyo would be located at southwest of the upward field-aligned current during this interval. The D-component variation shows a few negative excursions during \u223c25:40\u201330:00\u2009LT (\u223c16:40\u201321:00\u2009ut). Although the disturbance of the H-component is unavailable during this interval, it can be speculated that Tokyo was located at southeast of the downward field-aligned current. In addition to the development of the ring current, the formation of the substorm current wedge may also cause the telegraphic disruption. By considering the above discussion, at mid-latitude, we suggest that there are two regions where GICs can be caused by field-aligned currents. One is southwest of the upward field-aligned current, and the other is southeast of the downward field-aligned current. In these regions, the H-component of the magnetic field is expected to decrease significantly because of the combination of the effects of the ring current and the current wedge. Tokyo was probably situated in such hazardous regions during the 25 September 1909 storm.","Citation Text":["McPherron, Russell & Aubry 1973"],"Citation Start End":[[813,844]]}
{"Identifier":"2018AandA...613A..76J__Kennedy_&_Kenyon_2008_Instance_1","Paragraph":"One of the most intriguing results from RV surveys is the observed scarcity of relatively close-in (a \u2272 0.5 AU) planets around post-MS stars. This observational trend has been attributed to the strong tidal torque exerted by the star as its radius grows during the giant phase. As a result, planets are expected to lose orbital angular momentum, thus moving inward until they are evaporated in the stellar atmosphere (Livio & Soker 1983; Sato et al. 2008; Villaver & Livio 2009; Kunitomo et al. 2011). On the other hand, the majority of the giant stars targeted by RV surveys are intermediate-mass stars (M\u22c6 ~ 1.5\u20133.0 M\u2299), thus they are the post-MS counterpart of A and early F main-sequence stars. Therefore, their companions should not be directly compared to those orbiting solar-type stars. Based on this analysis, known planets orbiting field giant stars are expected to be born in different conditions from those around low-mass stars. In particular, these planets are formed in more massive disks (since Md \u221d M\u22c6; Andrews et al. 2013), from which they can efficiently accrete a significant amount of gas, becoming gas giants (e.g., Kennedy & Kenyon 2008). In addition, due to the higher gas accretion rate (Muzerolle et al. 2005) and higher irradiation, these disks have shorter dissipation timescales (Currie 2009; Kennedy & Kenyon 2009) and the snow line is located at a greater distance from the central star (Kennedy & Kenyon 2008). As a consequence, these planets are most likely formed at greater orbital distances and, due to the shorter disk timescale, inward migration is halted; they thus reach their final position at a relatively large distance from the parent star. For comparison, Currie (2009) predicted that only ~1.5% of intermediate-mass stars host giant planets with a \u2272 0.5 AU, while \u22737.5% of them host at least one gas giant at a \u2273 0.5 AU. Fig. 9 shows the mass versus the orbital distance of planets detected around giant stars (log g \u2272 3.5), via RV measurements (black dots) and by the transit method (red open circles). We note that values of the RV detected systems correspond to the minimum planet mass (Mp sini). The dotted line represents a radial velocity semi-amplitude of K = 30 m s\u22121 for a 1.5 M\u2299 star, (corresponding to a 3-\u03c3 detection; e.g., Hekker et al. 2006). As can be seen, there is only one companion detected via RVs interior to 0.1 AU, and the rest of them reside at an orbital distance a \u2273 0.4 AU. As discussed above, this observational result might be explained by the engulfment of the innermost planets as the parent star evolves off the MS and becomes a giant star. However, since a similar trend is observed in less evolved subgiants whose radii have not yet reached a value where tidal interactions are strong enough to affect the orbits of their companions, Johnson et al. (2007) argued that this is probably explained by a different formation scenario between planets around low-mass stars and those formed in more massive disks. From Fig. 9 it is also evident that planets residing interior to ~0.1 AU are significantly less massive (Mp \u2272 1 MJ) than those orbiting at a greater distance. In fact, two of these transiting planets are well below the 3-\u03c3 detection threshold, thus they are not detectable via radial velocities. A similar trend is also observed in MS stars (Zucker & Mazeh 2002), which might be caused by a decrease in the type II migration speed with increasing planetary mass, i.e., d a\u2215dt \u221d M\n\n$_P^{-1}$\nP\u22121\n (Mordasini et al. 2009). This theoretical prediction naturally explains why the most massive planets are found at a \u2273 0.4 AU. On the other hand, the mass distribution of the parent stars of these two populations of planets are different. While the mean stellar mass of the RV detected planets is 1.78 M\u2299, this value is only 1.38 M\u2299 for the transiting systems and thus two distinct planet mass distributions are expected to be found. Moreover, a similar result is observed between the mass of planets orbiting subgiant and giant stars (planets around giant stars being significantly more massive than those around subgiants; see Jones et al. 2014). In fact, the mean mass of the subgiant parent stars is 1.5 M\u2299, significantly lower than giant host stars. These results provide further observational support of a different formation and migration scenario for planets at different host star mass. This result suggest that the observed lack of planet around giant stars is mainly due to the primordial distinct formation scenario proposed by Johnson et al. (Johnson et al. (2007)).","Citation Text":["Kennedy & Kenyon 2008"],"Citation Start End":[[1138,1159]]}
{"Identifier":"2021MNRAS.502.5779N__Plas_et_al._2019_Instance_1","Paragraph":"The dynamical interaction with a companion is proposed to be the driving force behind most features seen in transition discs including the formations of dust traps. For our RC-TDs, van der Marel et al. (2015a, 2016a, 2018) showed gas cavities inside dust cavities for SR21, HD135344B, DoAr44, J1604, Sz111, J16083070. van der Plas et al. (2019) utilized SPH and radiative transfer codes to show some disc features seen in HD100453 can be attributed to an undetected, low mass close companion within the disc\u2019s cavity. The study of HD100453 is further refined by Gonzalez et al. (2020) and Nealon et al. (2020), who show that while the outer disc morphology can be caused by a companion star on an inclined orbit exterior to the disc, the inner cavity can be explained by a $\\rm \\lesssim 5$ MJ planet, less massive than previously suggested (van der Plas et al. 2019), at 15\u201320\u2009au. Using similar 1D radial brightness profiles for the majority of our sources (HD135344B, RY Lup, J1608, Sz111, SR24S, and DoAr44), Pinilla et al. (2017, 2018) suggest that the clearing of the inner cavity could be attributed to the dynamical interaction with an embedded planet. For HD169142, Fedele et al. (2017) uses a thermo-chemical code to model the gas and dust distribution. Their results suggest that dynamical interaction between the disc and two giant embedded planets results in the depletion of the inner disc and creates two pressure bumps that aid in the trapping of dust at each observable ring position. It still remains unclear for the majority of transition discs if a companion is indeed responsible for the formation of the cavity, whether this companion is an embedded planet or a binary star. What is notable about our sample is that for the RC-TDs the 8.8\u2009mm and sub-mm grains share a similar cavity size ($\\rm R_{peak}$ value). This is also seen in SR24S and HD142527 (included in Fig. 4). Pinilla et al. (2019) present a planet-disc model for SR24S that predicts both mm and sub-mm grains will share a similar radial peak in the dust density. Price et al. (2018) extensively modelled HD142527 with 3D hydrodynamical simulations considering several possible orbits for the M-dwarf binary companion presented in Lacour et al. (2016). They conclude that all the observable disc features can be attributed to the tidal truncation of the disc from the binary companion. Given that our RC-TDs and HD142527 follow the one-to-one correlation shown in Fig. 4, we suggest similar physical mechanisms affecting the grain evolution for these discs may have occurred.","Citation Text":["van der Plas et al. 2019"],"Citation Start End":[[841,865]]}
{"Identifier":"2015MNRAS.450.2749G__Brinchmann_et_al._2004_Instance_2","Paragraph":"Ideally, however, one would want to go beyond the description of cosmic global history, and trace galaxy evolution on a galaxy-by-galaxy basis to understand the physical processes driving it. In this respect, great progress has been made by surveys at different redshifts that have established the existence of a strong dependence of galaxy histories on galaxy stellar mass. On average, more massive galaxies have formed their stars and completed their star formation activity at higher z than less massive galaxies (the so-called downsizing effect, Cowie et al. 1996; Gavazzi et al. 2006; De Lucia et al. 2007; S\u00e1nchez-Bl\u00e1zquez et al. 2009). The existence of relations between SFR and galaxy stellar mass (SFR\u2013Mass) and specific star formation rate and mass (sSFR = SFR\/Mass) have been established from z = 0 out to z > 2 (Brinchmann et al. 2004; Daddi et al. 2007; Noeske et al. 2007; Salim et al. 2007; Rodighiero et al. 2011; Whitaker et al. 2012; Sobral et al. 2014; Speagle et al. 2014), and many other galaxy properties have been found to be correlated with galaxy mass. Furthermore, a number of works have pointed out that galaxy properties are even more strongly correlated with a combination of galaxy mass and galaxy \u2018size\u2019, arguing for velocity dispersion (Bernardi et al. 2003; Franx et al. 2008; Smith, Lucey & Hudson 2009; Wake, van Dokkum & Franx 2012) or galaxy surface mass density (Brinchmann et al. 2004; Kauffmann et al. 2006) as principal drivers. The exact origin of these trends is still unknown, but evidence has accumulated for a dependence of galaxy stellar population ages on galaxy sizes at fixed mass (van der Wel et al. 2009; Cappellari et al. 2012; Poggianti et al. 2013), suggesting that also galaxy structure, and not just stellar mass, is relevant. In a recent paper, Omand, Balogh & Poggianti (2014) argue that the observed correlation of the quenched fraction with M\/R1.5 is related to the dominance of the bulge component with respect to the disc, suggesting it might ultimately be linked with galaxy morphology (see also Driver et al. 2013). Even the sSFR\u2013Mass relation might be due to the increase of the bulge mass fractions with galaxy stellar mass, as the ratio of SFR and stellar mass of the galaxy disc is virtually independent of total stellar mass (Abramson et al. 2014).","Citation Text":["Brinchmann et al. 2004"],"Citation Start End":[[1401,1423]]}
{"Identifier":"2022ApJ...937...62L__Biller_et_al._1999_Instance_1","Paragraph":"However, the intrinsic effects caused by the unknown emission and acceleration mechanisms in the source could mitigate or enhance the LIV-induced time delay, which would impact the accuracy of the resulting constraints on LIV. A key challenge is then to distinguish an intrinsic time lag at the source from a delay induced by LIV. Long GRBs usually have significantly positive or negative intrinsic spectral lags and should not be used for LIV searches until reasonable progress is made on the modeling of the emission and acceleration mechanisms (Chen et al. 2005; Ukwatta et al. 2012; Bernardini et al. 2015), while short GRBs are consistent with null or negligible intrinsic spectral lag and are therefore an ideal tool to measure the LIV effect (Norris & Bonnell 2006; Bernardini et al. 2015, 2017; Xiao et al. 2022). Currently, in addition to short GRBs, active galactic nucleus (AGN) flares and gamma-ray pulsars are two other classes of astrophysical sources that have no significant intrinsic lag in general and are often used for LIV tests (Biller et al. 1999; Kaaret 1999; Aharonian et al. 2008; MAGIC Collaboration et al. 2017). It should be noted, however, that there is also evidence of intrinsic lags in some cases of AGN flares and pulsars. For example, MAGIC Collaboration during an observational campaign regarding Mkn 501 blazar found an indication of about 4 minutes time delay between the peaks at E  0.25 TeV and E > 1.2 TeV (MAGIC Collaboration et al. 2008), which may indicate a progressive acceleration of electrons in the emitting plasma blob. A robust method to study the correlations between arrival times and energy, based on a likelihood function built from the physical picture assumed for the emission, propagation, and detection of the photons was proposed by Mart\u00ednez & Errando (2009). In the case of pulsars, there are some lags if the energy range is extended too much toward low energies (e.g., radio versus TeV). No real progress on the topic of intrinsic effects will be made without accurate models for production and acceleration mechanisms for each type of source. Perennes et al. (2020) first attempted to gain knowledge on source-intrinsic spectral lags of flaring AGNs at high and very high energies and on short timescales relevant for LIV searches, using leptonic AGN flare modeling. Concerning GRBs, some ways have been proposed to reduce the impact of intrinsic effects, e.g., fitting the observed spectral lags of statistical samples of GRBs at a range of different redshifts (Ellis et al. 2006; Bernardini et al. 2017; Xiao et al. 2022), or using only a limited observer-frame energy bands range corresponding to the fixed source-frame energy bands (Wei & Wu 2017). Anyway, there is no reason to think that the low and high-energy photons should be emitted simultaneously at the source, and while detecting distinct signals at different energy channels, we have no idea which one was sent first. Previous studies usually assumed that the intrinsic time delays are either an unknown constant for all GRBs considered or scale with the photon energy E according to some power-law function (Ellis et al. 2006; Biesiada & Pi\u00f3rkowska 2009a; Zhang & Ma 2015; Wei et al. 2017a; Acciari et al. 2020; Pan et al. 2020; Du et al. 2021).","Citation Text":["Biller et al. 1999"],"Citation Start End":[[1050,1068]]}
{"Identifier":"2022ApJ...929....8E__Qin_&_Shen_2017_Instance_1","Paragraph":"Diffusion plays a significant role in the transport of charged, energetic particles in the heliosphere, whether these be galactic cosmic rays (e.g., Engelbrecht 2019; Shen et al. 2019; Moloto & Engelbrecht 2020), solar energetic particles (e.g., Strauss et al. 2017; van den Berg et al. 2021), or Jovian electrons (e.g., Ferreira et al. 2001; Vogt et al. 2020). As such, the diffusion of such particles parallel and perpendicular to the heliospheric magnetic field (HMF) has attracted considerable interest over the years, with various theoretical models for these quantities being proposed, ranging from the quasilinear theory (QLT) of Jokipii (1966), to the nonlinear guiding center (NLGC) theory of Matthaeus et al. (2003) and various recent theories (see, e.g., Shalchi 2010; Ruffolo et al. 2012; Qin & Zhang 2014), which have in turn been applied in various particle transport studies (e.g., Florinski & Pogorelov 2009; Engelbrecht & Burger 2015; Qin & Shen 2017; Moloto et al. 2018; Engelbrecht & Moloto 2021). These theoretical advances have benefited from new insights into particle diffusion via two separate avenues of research, both involving numerical simulations, as particle diffusion coefficients cannot be directly observed. One avenue of research involves numerical test particle simulations of particle diffusion coefficients in the presence of synthetically generated magnetic turbulence, in the process of gaining new insights as to the behavior of these particles, in particular, turbulence scenarios corresponding to a greater or lesser degree with what is observed in the solar wind (see, e.g., Minnie et al. 2007a; Ruffolo et al. 2008; Tautz & Shalchi 2011; Dalena et al.2012; Kong et al. 2017; Mertsch 2020). The other avenue for insights into the diffusion coefficients of charged particles involves the simulation of their transport in the heliosphere. Such studies report on the transport parameters needed as inputs for transport models to ensure agreement between computed particle intensities and various in situ spacecraft observations of the same. These approaches are by no means incompatible, as results from numerical test particle simulations have been shown to be in agreement with those yielded by transport modeling studies (Tautz & Shalchi 2013). The present study follows the latter avenue of inquiry, focusing in particular on the transport of Jovian electrons. In a sense, these particles are the ideal candidates for such a study, as the transport of low-energy electrons in the inner heliosphere has long been known to be negligibly influenced by drift effects (see, e.g., Ferreira et al. 2001; Engelbrecht & Burger 2010; Engelbrecht et al. 2017). Furthermore, the relative orbital motions of the Earth and Jupiter across the Parker (1958) HMF, coupled with the relative abundance of electron observations at 1 au as well as the Jovian source, allows one to compare computed intensities with observations taken during periods of good and bad magnetic connectivity with the source. This, then, allows for the simultaneous study of the transport coefficients of these electrons both parallel and perpendicular to the HMF, which is the aim of this study.","Citation Text":["Qin & Shen 2017"],"Citation Start End":[[952,967]]}
{"Identifier":"2019AandA...628A.118B__Feruglio_et_al._2017_Instance_1","Paragraph":"Ultra-fast outflows (UFOs) of highly ionised gas observed at sub-parsec scales (Reeves et al. 2003; Tombesi et al. 2012) have been proposed as the likely origin of galaxy-wide outflows, interpreted as the result of the impact of UFOs on the ISM (King & Pounds 2015, and references therein). Furthermore, both models and observations of kiloparsec-scale outflows seem to indicate a UFO-ISM interaction in an energy-conserving regime, whereby the swept-up gas expands adiabatically. So far, the co-existence of a massive molecular outflow with a nuclear UFO has been confirmed in a handful of AGNs with LBol\u2006\u2004\u223c\u2004\u20061044\u2005\u2212\u20051046 erg s\u22121 (Tombesi et al. 2015; Feruglio et al. 2015; Longinotti et al. 2015) and in APM 08279+5255 (Feruglio et al. 2017), which is a gravitationally lensed QSO at z\u2006\u2004\u223c\u2004\u20064 with an estimated intrinsic LBol of a few times 1047 erg s\u22121 (Saturni et al. 2018). In all these sources the momentum boost (i.e. the momentum flux of the wind normalised to the AGN radiative momentum output, LBol\/c) of the UFO is \u223c1, while the momentum rate of the molecular outflow is usually \u226b1, in qualitative agreement with the theoretical predictions for an energy-conserving expansion (Faucher-Gigu\u00e8re & Quataert 2012; Costa et al. 2014). However, these results are still limited to a very small sample and suffer from large observational uncertainties, mostly due to the relatively low signal-to-noise ratio of the UFO- or outflow-related features confirmed in spectra, or to the limited spatial resolution of sub-millimetre observations. Recent studies increasing the statistics of sources with detection of molecular outflows have widened the range of measured energetics (e.g. Garc\u00eda-Burillo et al. 2014; Veilleux et al. 2017; Feruglio et al. 2017; Brusa et al. 2018; Barcos-Mu\u00f1oz et al. 2018; Fluetsch et al. 2019). These outflows are consistent with driving mechanisms alternative to the energy-conserving expansion, such as direct radiation pressure onto the host-galaxy ISM (e.g. Ishibashi & Fabian 2014; Ishibashi et al. 2018; Costa et al. 2018).","Citation Text":["Feruglio et al. 2017"],"Citation Start End":[[721,741]]}
{"Identifier":"2019MNRAS.484.1645O__Stewart_et_al._2013_Instance_1","Paragraph":"This additional heating effect would raise the temperature of the inflowing gases to a level where they cannot drive star formation effectively and begin to evaporate. We may use the virial theorem to estimate the time-scale over which filamentary inflows may halted in this manner: the virial temperature is the temperature above which gravitational collapse of gas is halted and, presumably, any heating to temperatures above this level would lead to evaporation. If assuming inflows persist over filaments of lengths of up to 50 kpc (Dekel et al. 2009; Stewart et al. 2013; Goerdt & Ceverino 2015; Dayal & Ferrara 2018) and that they are the sole driver of the star-formation activity arising at a rate of $\\mathcal {R}_{\\rm SF} \\approx 16~\\text{M}_{\\odot }~\\text{yr}^{-1}$ with a 30 per\u2009cent mass conversion efficiency (Meier, Turner & Beck 2002; Turner et al. 2015; Behroozi & Silk 2015; Sun & Furlanetto 2016) with an inflow velocity of 400\u2009km s\u22121 (from the velocity offset of the Lyman-\u03b1 line in MACS1149-JD1 \u2013 see Hashimoto et al. 2018), the steady-state mass of these inflows would be around 6.7 \u00d7 107 M\u2299. If the typical diameter of these flows is similar to the galaxy which they feed, i.e. 1 kpc, this would give a virial temperature of Tvir \u2248 5000 K (Binney & Tremaine 2008). If adopting a number density of 10 cm\u22123 for the filamentary inflows (i.e. comparable to the mean density of the ISM in the host galaxy), the estimated CR heating rate of 10\u221225erg cm\u22123 s\u22121 would suggest it would take of order only a few Myrs for the virial temperature of the inflows to be exceeded. At this point, the supply of gas to the galaxy would be strangulated, halting star-formation within a dynamical time-scale \u03c4dyn (being the time required for the system to respond to the strangulation and internal heating). Since the strangulation and ISM quenching time-scales are comparatively short, the magnetic saturation and dynamical time-scales alone specify the time-scale over which star-formation would be quenched.","Citation Text":["Stewart et al. 2013"],"Citation Start End":[[556,575]]}
{"Identifier":"2020AandA...641A.123H__Mikal-Evans_et_al._(2019)_Instance_3","Paragraph":"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\u00d7 to 30\u00d7 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\u20131.1 \u03bcm), 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 \u03bcm feature is caused by VO emission. The secondary eclipse was observed at 2 \u03bcm (Kov\u00e1cs & Kov\u00e1cs 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 \u03bcm, prompting the authors to entirely discard their previous interpretation of emission caused by VO.","Citation Text":["Mikal-Evans et al. (2019)"],"Citation Start End":[[1086,1111]]}
{"Identifier":"2021AandA...656A..63M__Paice_et_al._2019_Instance_1","Paragraph":"MAXI J1820 is a BHT that was observed for the first time in the optical band by the All-Sky Automated Search for SuperNovae ASSAS-SN (Shappee et al. 2014) on 2018 March 3 (Tucker et al. 2018), and one week later by MAXI in X-rays (Kawamuro et al. 2018). Detailed studies of the optical counterpart revealed that the system hosts a stellar-mass black hole (\u223c8.5\u2006M\u2299) accreting from a \u223c0.4\u2006M\u2299 companion star (Torres et al. 2019, 2020). Since its discovery, the source has undergone a long (approximately one year) and bright outburst, becoming at its peak the second brightest object in the X-ray sky. Due to its brightness, the system was the object of an impressive multi-wavelength observing campaign (see, e.g., Shidatsu et al. 2018; Paice et al. 2019; Hoang et al. 2019; Trushkin et al. 2018; Bright et al. 2020; Tetarenko et al. 2021) and of a large number of studies. The most recent and precise measure of the distance, determined via radio parallax, is 3.0\u2005\u00b1\u20050.3 kpc (Atri et al. 2020). Furthermore, the system shows X-ray dips (Kajava et al. 2019) but not eclipses, suggesting an inclination between 60\u00b0 and 80\u00b0. Other evidence of the high inclination of the system are provided by optical spectroscopy (Torres et al. 2019) and by the estimate of the inclination of the jet axis (Atri et al. 2020; Wood et al. 2021), which is about 63\u00b0. The orbital period of the system has also been estimated to be around 0.68 days (Patterson et al. 2018; Torres et al. 2020). In X-rays, the outburst was studied in detail in HS (see, e.g., Bharali et al. 2019; Buisson et al. 2019; Zdziarski et al. 2021), SS (see, e.g., Fabian et al. 2020) and in its final phase (see, e.g., Xu et al. 2020). The truncation of the disk during the HS is one of the most controversial aspects of the system. On the one hand, a number of spectral-timing works have pointed out that the disk reaches the ISCO already in HS, while a contracting lamppost corona is responsible for the hard X-ray emission (Kara et al. 2019; Buisson et al. 2019; You et al. 2021; Wang et al. 2021). On the other hand, a truncated disk scenario has also been proposed on the basis of further spectral-timing analyses, even on the same sets of data (Zdziarski et al. 2021; De Marco et al. 2021; Axelsson & Veledina 2021).","Citation Text":["Paice et al. 2019"],"Citation Start End":[[735,752]]}
{"Identifier":"2020MNRAS.498.2594S__Ahnen_et_al._2015_Instance_1","Paragraph":"In the modelling, it is assumed that the low-energy peak (from radio to optical\u2013UV) is due to synchrotron emission from ultrarelativistic electrons in the jet with an energy distribution as given by equation (2). Instead, the HE peak is due to the IC scattering of internal (SSC; Ghisellini et al. 1985; Maraschi et al. 1992; Bloom & Marscher 1996) or external photons (EIC; Sikora et al. 1994; B\u0142a\u017cejowski et al. 2000; Ghisellini & Tavecchio 2009). The IC scattering of external photons is considered, since the SEDs of FSRQS are better explained by EIC, as shown by the previous studies (e.g. Abeysekara et al. 2015; Ahnen et al. 2015; Hayashida et al. 2015b; Gasparyan et al. 2018; MAGIC Collaboration 2018), and the CD is evident in the SEDs of the considered sources (Fig. 6). Localization of the emission region in the jet is an open question and along the jet, depending on the distance from the central black hole, different photon fields can be dominant for the IC scattering (Sikora et al. 2009). In this paper, we assume that the emitting region is outside the broad-line region (BLR) where the dominant photon field is the IR emission from the dusty torus. Sikora et al. (2002) showed that in MeV blazar SEDs the shift of the peak of the HE component to lower energies is most likely due to the Comptonization of IR photons from the dusty torus. The IR radiation from the dusty torus is assumed to have a blackbody spectrum with a luminosity of LIR = 0.6\u2009Ldisc (see Ahnen et al. 2015) where Ldisc is the accretion disc luminosity, which fills a volume that for simplicity is approximated as a spherical shell with a radius of RIR = 2.5 \u00d7 1018\u2009(Ld\/1045)1\/2\u2009cm (Nenkova et al. 2008) with the energy density of $u_{\\rm IR}=0.6\\:L_{\\rm d}\/4 \\pi R_{\\rm IR}^2\\:\\delta ^2$ in the co-moving frame of the jet. In Ghisellini et al. (2011) and Marcotulli et al. (2017), the HE component in the SED of distant blazars was modelled by IC scattering of BLR reflected photons, adopting a smooth broken power-law shape of the emitting electrons. We refer the reader to these papers for details on the modelling when BLR reflected photons are considered.","Citation Text":["Ahnen et al. 2015"],"Citation Start End":[[619,636]]}
{"Identifier":"2020AandA...637A..44N__Kraus_(2018)_Instance_3","Paragraph":"Among the existing IACT systems, HESS has the largest FoV and hence provides the highest sensitivity for the diffuse \u03b3-ray flux. Its electron spectrum analysis technique could be directly used to obtain a measurement of the diffuse Galactic \u03b3-ray flux above energies of several TeV in the Galactic Ridge (|l|  30\u00b0, |b|  2\u00b0) region; see Figs. 3 and 4. A multi-year exposure of HESS could be sufficient for detection of the diffuse emission even from regions of higher Galactic latitude. This is illustrated in Fig. 5, where thick blue and red data points show mild high-energy excesses of the electron spectra derived by Kraus (2018), Kerszberg et al. (2017), Kerszberg (2017) over broken power-law models derived from the fits to lower energy data. Comparing these excesses with the level of the IceCube astrophysical neutrino flux and with the Fermi\/LAT diffuse sky flux from the region |b| > 7\u00b0 (corresponding to the data selection criterium of HESS analysis Kerszberg et al. 2017; Kerszberg 2017) we find that the overall excess flux levels are comparable to expected diffuse \u03b3-ray flux from the sky region covered by the HESS analysis (the quoted systematic error on the electron flux is \u0394log(EFE) \u2243 0.4). The overall excesses within 805 and 1186 h of HESS exposures (Kraus 2018; Kerszberg 2017) are at the levels of >4\u03c3 for the analysis of Kraus (2018) and 1.7\u03c3 for the analysis of Kerszberg (2017). A factor-of-ten longer exposure (which is potentially already available with HESS) could reveal a higher significance excess at the level of up to 5\u03c3. Such an excess is predicted in a range of theoretical models including interactions of cosmic rays injected by a nearby source (Andersen et al. 2018; Neronov et al. 2018; Bouyahiaoui et al. 2019) or decays of dark matter particles (Berezinsky et al. 1997; Feldstein et al. 2013; Esmaili & Serpico 2013; Neronov et al. 2018) or a large-scale cosmic ray halo around the Galaxy (Taylor et al. 2014; Blasi & Amato 2019).","Citation Text":["Kraus (2018)"],"Citation Start End":[[1345,1357]]}
{"Identifier":"2018ApJ...856...51R__Reale_2014_Instance_1","Paragraph":"Close to the end of their formation, stars are surrounded by a gas and dust disk, from which planets form. Magnetic fields are known to play a key role in the star-disk system (Johns-Krull 2014). It is believed that the inner regions of the disk are significantly ionized by the stellar radiation and that accreting material flows along magnetic channels that connect the disk to the star (Koenigl 1991). Very long and intense X-ray flares in star-forming regions might occur in such long channels (Favata et al. 2005), but this is highly debated (Getman et al. 2008). These flux tubes might resemble those observed in the solar corona and diagnosed in the stellar coronae, but on a much larger scale. On the Sun we see the so-called coronal loops on the scale of several thousand kilometers in active regions, but some faint large-scale structures can extend up to \u223c1 R\u2299 (Reale 2014). Most solar flares occur in active region loops, but the long-lasting ones can involve more and more loops aligned in arcades. The other stars are so distant that we cannot resolve the flaring regions, but it is supposed that they occur in loops and even in arcades. Whereas the duration of solar flares typically ranges from a few minutes to several hours, stellar flares can be very intense, more than the solar bolometric luminosity, and long-lasting, including longer than one day, in very active stars. Several such gigantic coronal flares have been surveyed in star-forming regions (Favata et al. 2005) and where they occur is a big question. Magnetic instabilities in flux tubes were proposed to be the origin of the flaring activity also in T Tauri stars (Birk 1998; Birk et al. 2000), and long-lasting stellar flares might be expected to involve loop arcades (Getman et al. 2008), like those on the Sun. In long-lasting solar flares, the duration is mainly due to the progressive involvement of more and similar loops, therefore duration is not directly linked to the size of the flaring structures. This might also be the case for giant stellar flares. However, if a single stellar loop were flaring, the cooling time of the confined plasma would be proportional to the loop length (Serio et al. 1991; Reale 2014), and day-long flares would correspond to giant loops, as long as they possibly connected the star with the disk (Hartmann et al. 2016). There are ways to distinguish between a pure cooling in a single loop and a decay only due to progressive reduction of the energy release in a loop arcade (Reale et al. 1997), but the explanations are contested and the uncertainties are large (Getman et al. 2008). Several studies (Favata et al. 2005; Giardino et al. 2007) find results compatible with long magnetic channels in pre-main sequence (PMS) stars, but the derivation of the loop length is based on the assumption of a flare occurring in a single loop (Reale 2007).","Citation Text":["Reale 2014"],"Citation Start End":[[873,883]]}
{"Identifier":"2018AandA...616A..34H__Mohamed_&_Podsiadlowski_(2012)_Instance_3","Paragraph":"The CO emission, tracking the bulk of the gas, reveals an almost face-on one-armed spiral, of which almost two full windings can be traced. What could be the origin of this spiral structure? As the majority of AGB stars are in binary systems, and perhaps all host planets, interaction between the outflow and a sufficiently massive and nearby companion may be the explanation of the observed CO morphology. The intricate emission features in the inner 2\u2033 of the central CO emission maps is strongly reminiscent of hydrodynamical simulations of wind\u2013binary interaction by Mastrodemos & Morris (1998) and Mohamed & Podsiadlowski (2012), where the latter authors performed tailored simulations for the Mira AB system in which the outflow of the AGB star Mira A is perturbed by the presence of its close companion Mira B. The wind\u2013binary interaction that ensues leads to what is known as wind Roche-lobe overflow (WRLOF), where the slow AGB wind is confined to the star\u2019s Roche lobe, while overflowing through the L1 Lagrange point. Gravitational interaction of the overflowing material with the companion produces an intricate feedback system where the stellar outflow material is ejected into the surrounding CSE through two distinct streams (through L2 and the stagnation point3) which combine to form an annular stream. As this stream travels outwards, it creates the larger scale spiral observed in the wind. The morphology resulting from this particular type of wind\u2013binary interaction is shown in Fig. 3 in Mohamed & Podsiadlowski (2012). In Fig. 11 we show the emission pattern seen in the central regions of the CO channel at \u03c5*. We compare this image with the bottom left panel of Fig. 3 in Mohamed & Podsiadlowski (2012), an opacity map of the interaction zone. Though the two properties that are compared differ in nature, they likely still trace the same global morphological structure. Indeed, several of the predicted morphological features can be identified in the data of EP Aqr. The bright central region with a north and southward hook-like extension are strikingly similar, as are the eastern and western crescent-shaped \u201cvoids\u201d, the overall shape, and the morphological properties of the small-scale instabilities.","Citation Text":["Mohamed & Podsiadlowski (2012)"],"Citation Start End":[[1697,1727]]}
{"Identifier":"2021MNRAS.507.5882S__Mackereth_et_al._2018_Instance_4","Paragraph":"Cosmological hydro dynamical N-body simulations offer another possibility to investigate the origin of the bimodality in the ([Fe\/H], [\u03b1\/Fe]) plane. Earlier simulations, e.g. full N-body simulations by Loebman et al. (2011), Brook et al. (2012) or hybrid simulations in which a semi-analytic chemical evolution was added on top of a cosmological simulation (Minchev, Chiappini & Martig 2013, 2014), were able to show that the thin and thick discs lie along different tracks in the ([Fe\/H], [\u03b1\/Fe]) plane, with the thick disc being old metal poor and rich in [\u03b1\/Fe] and the thin disc being young, metal-rich and poor in [\u03b1\/Fe]. They also showed that migration was important to generate the two discs. However, a clear bimodality in the ([Fe\/H], [\u03b1\/Fe]) plane was not seen. In the past few years good progress has been made to improve the spatial resolution as well as the chemical enrichment prescriptions. The bimodality has now been observed in some simulations (Grand et al. 2018; Mackereth et al. 2018; Clarke et al. 2019), and some of the simulations, in addition to the bimodality, also reproduce the basic trends of the ([Fe\/H], [\u03b1\/Fe]) distribution with radius R (Buck 2020; Vincenzo & Kobayashi 2020). Unlike analytical models, such simulations cannot be fine tuned to reproduce the Milky Way data, hence, the focus of these simulations is to qualitatively reproduce the abundance trends seen in the Milky Way, to understand how frequently do we get the bimodality and what is the mechanism for it. However, there is a lack of consensus between the different studies. Clarke et al. (2019) and Buck (2020) suggest that bimodality should be common in disc galaxies, whereas Mackereth et al. (2018) suggest that it is rare. Each simulation suggests slightly different mechanisms for the existence of the bimodality. Clarke et al. (2019) attribute bimodality to vigorous star formation in clumps at high redshift. Grand et al. (2018) suggest two distinct pathways, a centralized starbust pathway induced by mergers and a shrinking gas disc pathway. Buck (2020) suggest that after the formation of the high-[\u03b1\/Fe] sequence a gas-rich merger dilutes the metallicity of the ISM leading to the formation of the low-[\u03b1\/Fe] sequence. Mackereth et al. (2018) attribute the bimodality to unusually rapid gas accretion at earlier times, which is also characterized by a short time-scale to convert gas to stars. While some simulations clearly identify migration as key process to shape the sequences, others do not. In spite of the differences, it seems that some of the simulations (e.g. Mackereth et al. 2018; Buck 2020; Vincenzo & Kobayashi 2020) are not inconsistent with the Sch\u00f6nrich & Binney (2009a) paradigm.","Citation Text":["Mackereth et al. 2018"],"Citation Start End":[[2584,2605]]}
{"Identifier":"2022MNRAS.510.5088B__Faisst_et_al._2017_Instance_1","Paragraph":"In the expectation of a relationship between the observed IRX and the rest-frame UV slope there is the assumption that the stars and dust are well mixed, which leads to the coupling of any observed reddening in the UV to the FIR emission detected (e.g. Meurer et al. 1999; Charlot & Fall 2000; Calzetti 2001). If instead the galaxy consists of regions of significantly different obscuration, then the relationship will break down for the galaxy as a whole. Indeed, geometric effects have been put forward as an explanation of the discrepant results at z > 5 (Faisst et al. 2017; Popping, Somerville & Galametz 2017). In local starburst galaxies, the existence of an IRX\u2013\u03b2 relation and a clear morphological similarity, indicates that the rest-frame UV emission from young stars is being attenuated from dust that is tracing broadly the same star-forming regions of the galaxy (e.g. in the spiral arms; Kennicutt et al. 2003; Gil de Paz et al. 2007). In the high-redshift Universe, however, where galaxies become more turbulent and irregular (e.g. F\u00f6rster Schreiber et al. 2011; Buitrago et al. 2013; Guo et al. 2015) the expected morphology of the dust relative to the observed UV emission is not clear. Evidence for offset dust continuum emission relative to the rest-UV has been identified in several high-redshift Lyman-break galaxies (LBGs; Koprowski et al. 2016; Faisst et al. 2017; Laporte et al. 2017; Bowler et al. 2018), and similar trends have been found when comparing the [C ii] FIR line and the rest-UV continuum (e.g. Maiolino et al. 2015; Carniani et al. 2017). Although some of these offsets have been attributed to astrometric systematics (e.g. Dunlop et al. 2017) there is a growing consensus that FIR continuum and line emission are frequently physically offset as compared to the observed rest-UV emission (see Carniani et al. 2018). Whether high-redshift galaxies show large and distinct regions of obscured and unobscured star formation has implications for the use of the IRX\u2013\u03b2 relation in deriving the cosmic SFR density (e.g. Bouwens et al. 2016b), as the assumed energy balance will break down (Buat et al. 2019) and the global \u03b2 measurement will not be representative of the full source.","Citation Text":["Faisst et al. 2017"],"Citation Start End":[[559,577]]}
{"Identifier":"2021MNRAS.504.4626K__Kraljic_et_al._2020b_Instance_1","Paragraph":"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, Dav\u00e9 & 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\u2019s 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\u2019s 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 Text":["Kraljic et al. 2020b"],"Citation Start End":[[565,585]]}
{"Identifier":"2016ApJ...833..216G__Gopalswamy_et_al._2014a_Instance_2","Paragraph":"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 (\u223c2000 km s\u22121) 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 (\u223c2 km s\u22122, 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 Text":["Gopalswamy et al. 2014a"],"Citation Start End":[[1091,1114]]}
{"Identifier":"2020ApJ...900..100R__White_et_al._2019_Instance_2","Paragraph":"It is much harder to localize and track the formation of current sheets in realistic black hole accretion flows in a larger domain and for a longer period because of the effects of the more complicated global dynamics governed by the central object, and due to the turbulence induced by the MRI. Both the evolution of accretion flows and the formation of current sheets therein strongly depend on the magnetic field geometry. We model an accretion disk around a rotating black hole, varying the initial conditions to study current sheet formation in different scenarios of magnetic field geometry. In the magnetically arrested disk (MAD; Igumenshchev et al. 2003; Narayan et al. 2003) scenario, the MRI and subsequent turbulence in the inner accretion disk are suppressed due to large-scale magnetic flux (see, e.g., White et al. 2019). In axisymmetric simulations as considered here, the arrested inflow is regularly broken by frequent bursts of accretion, allowing for a macroscopic equatorial current sheet to form and break in a periodic fashion. In a full 3D setup, magnetically buoyant structures are interchanged with less-magnetized dense fluid (Igumenshchev 2008; White et al. 2019), resulting in a magnetic Rayleigh\u2013Taylor instability (Kruskal & Schwarzschild 1954) potentially sourcing interchange-type magnetic reconnection. In the Standard And Normal Evolution (SANE; Narayan et al. 2012; Sadowski et al. 2013) state, a fully turbulent accretion disk can develop due to a smaller magnetic flux (see, e.g., Porth et al. 2019), and current sheets can ubiquitously form and interact with the turbulent flow. Polarized synchrotron radiation observed by the Event Horizon Telescope (Event Horizon Telescope Collaboration et al. 2019a) can probe the field-line structure at event-horizon scales and put tighter constraints on the magnetization and address whether the accretion is in a SANE or a MAD state (Event Horizon Telescope Collaboration et al. 2019b).","Citation Text":["White et al. 2019"],"Citation Start End":[[1173,1190]]}
{"Identifier":"2018MNRAS.477.4308R___1999_Instance_1","Paragraph":"Astrophysical processes involving radiative energy transfer are calculated by the balance between compressional heating and adiabatic cooling (Blondin et al. 1990; Taam, Fu & Fryxell 1991). Analytical prescriptions for the heating and cooling rates in complex environments are only possible under certain limits. However, the increasing computer power available today is allowing to model complex astrophysical scenarios efficiently and at a relatively low cost, including the dynamical update of the microphysics and chemistry. For instance, thermal instabilities, non-axisymmetric rotating objects, the relationships between ionization, molecular states, level populations, and kinetic temperatures of low-density environments are some of the ingredients that have no analytical counterparts and that can be calculated with highly efficient numerical algorithms. A particularly interesting question is related to the accretion of angular momentum, when the flow contains a density or velocity gradient perpendicular to the flow direction (Davies & Pringle 1980; Ruffert & Anzer 1995; Ruffert 1997, 1999). In particular, Proga & Begelman (2003) performed two-dimensional (2D) axisymmetric simulations of the accretion of slowly rotating gas on to a black hole. They found that for low angular momentum, the flow is Bondi-like, while for intermediate total angular momentum, the gas with higher angular momentum does not accrete and forms a dense torus about the black hole and, as a result, the accretion rate is substantially below the Bondi rate. For even higher total angular momentum, a dense torus coupled to a centrifugal barrier sets in, which further reduces the accretion rate. Proga & Begelman (2003) have argued that rotation may explain the difference between accretion rates estimated from observations and those predicted by the Bondi flow. Moreover, Perna et al. (2003) suggested that combining rotation and magnetic fields results in accretion rates on to isolated neutron stars that are much lower than those estimated using the Bondi formula. Later on, Krumholz, McKee & Klein (2005) considered the three-dimensional (3D) accretion of gas with constant vorticity and found that the resulting flow field is highly non-axisymmetric and time-dependent and that even a small amount of vorticity can substantially change the accretion rate. Considerable progress has been made over the past 20 yr in developing advanced numerical models to understand the physics of accretion on to compact objects. A thorough review of Bondi accretion theory and related earlier numerical simulations is given by Edgar (2004).","Citation Text":["Ruffert","1999"],"Citation Start End":[[1086,1093],[1100,1104]]}
{"Identifier":"2015AandA...580A..84D__Pont_et_al._(2008)_Instance_1","Paragraph":"We used the publicly available code AROME (Bou\u00e9 et al. 2013) to analytically compute the RM RV anomaly as measured with the CCF method. This code, written in C, was given the orbital planetary and stellar parameters (the projected rotation velocity Vsin(i), the semi-major axis a, the radius of the star Rstar, expressed in solar radii units, the inclination angle i, the mutual inclination angle \u03bb, and the planet orbital period P, all taken from Triaud et al. 2009, see Table 1), and the times of observation, which it subsequently used to calculate the CCFs and the expected RV measurements from a Gaussian fit to these. The fraction of flux blocked by a planet crossing the surface of its host star depends on the limb darkening of the star. For the calculation with AROME we chose the non-linear law of Claret (2000) to match those used in the analysis of Pont et al. (2008) and Sing et al. (2011b), that is, (3)\\begin{equation} \\frac{I \\left( \\mu \\right)}{I \\left( 1 \\right)} \\ = \\ 1 - \\sum_{k=1}^{4} a_{k} \\left( 1 - \\mu^{k\/2} \\right) \\label{claret2000}. \\end{equation}I\u03bcI(1)=1\u2212\u2211k=14ak(1\u2212\u03bck\/2).For the white-light RV time series we used the quadratic limb-darkening law coefficients as in Triaud et al. (2009), who studied the RM effect for this planet using the same data set. All the limb-darkening coefficients used in our analysis are given in Table 2. Since we are interested in the dependence of the RM effect on the planet effective size, for each passband we selected the limb-darkening law corresponding to those wavelengths and calculated the anomaly for 5000 values of the planet-star radius ratio (Rp\/Rstar), between 0.105 and 0.205 with a step of 2 \u00d7 10-5. These boundaries were chosen arbitrarily with the purpose of including the measured planet-star radius ratio for this planet using the same HARPS data as Triaud et al. (2009) (0.1581 \u00b1 0.0005). All other parameters were kept fixed. We note that for the analysis we are only interested in the relative change of the Rp\/Rstar as function of wavelength, not in absolute values. Therefore it is not relevant to fit for parameters such as the projected rotation velocity Vsini, the mutual inclination angle \u03bb, the macroturbulence z, and the scaled semi-major axis a\/Rstar separately for each passband. These should be independent of wavelength and therefore can only result in a small vertical offset in the transmission spectrum, not in a change of slope. ","Citation Text":["Pont et al. (2008)"],"Citation Start End":[[861,879]]}
{"Identifier":"2021MNRAS.503.2622C___1996_Instance_1","Paragraph":"Here, we relate the observed morphology of the dust continuum emission to the physical processes taking place within star-forming galaxies around the peak of cosmic star formation. As discussed by Cheng et al. (2020), the extended dust continuum emission observed in the less-FIR luminous SHiZELS galaxies suggests a dominant component of extended, disc-wide star formation; in contrast, the emission from sub-millimetre selected galaxies appears to be dominated by a compact, nuclear starburst. SHiZELS-14 is an outlier in the sense that it has both a submillimetre bright compact core and very extended emission. Tadaki et al. (2020) show that the most compact galaxies in their sample tend to have high gas fractions (derived via $S_{870\\mu \\rm {m}}$), and argue that this reflects efficient radial gas inflows. Numerical simulations have long shown that galaxy mergers are capable of triggering tidally driven gas inflows (Hernquist 1989; Barnes & Hernquist 1991), which can cause strong nuclear starbursts (e.g. Mihos & Hernquist 1994, 1996; Hopkins et al. 2013; Moreno et al. 2015). However, observations of local galaxies such as the Antennae system demonstrate that galaxy interactions can also trigger widespread star formation that is not limited to a compact, nuclear region (Wang et al. 2004). More recent, high resolution simulations show that these observations can be explained via merger-driven injections of turbulence into the ISM: extended compression results in fragmentation into dense, star-forming gas, and spatially extended starburst activity (Renaud et al. 2014; Renaud, Bournaud & Duc 2015). Renaud et al. (2015) argue that this process is particularly important in the early and mid-stages of a galaxy merger: during the first two simulated pericentre passages, star clusters form kiloparsecs from the galactic nucleus, with the central starburst dominating only from the beginning of the final coalescence. This progression of star formation from extended to compact as the merger unfolds is also consistent with observations of local galaxies (Pan et al. 2019). The extended star formation observed in SHiZELS-14 may therefore suggest that we are viewing the short-lived mid-stages of a merger; this would be consistent with its complex, irregular morphology and dispersion-dominated $\\rm {H}\\,\\alpha$ kinematics. The similarly TIR-luminous but more compact sources within the AS2UDS samples may comprise galaxies experiencing a wider range of evolutionary stages, including some later-stage mergers.","Citation Text":["Mihos & Hernquist","1996"],"Citation Start End":[[1017,1034],[1041,1045]]}
{"Identifier":"2018ApJ...853..107R__Imanaka_&_Smith_2010_Instance_1","Paragraph":"An important reason for modeling the THS plasma is identifying the pathways of chemical growth. Tables 3 and 4 list the dominant production and loss reactions, respectively, for a set of species. The neutral versions of these species have each been identified in Titan's upper atmosphere (Waite et al. 2007; Cui et al. 2009). The values of the calculated reaction rates are listed if larger than 1 \u00d7 10\u22126 mol m\u22123 s\u22121. For the experimental conditions considered, the model predicts that neutral C2H2 production is only appreciable in the mixture that already contains an acetylene precursor through the reaction of C4H2 with C4H3+, which are themselves reaction products that, at Titan, might be derived through other chemical pathways. C2H2 is destroyed mainly by reactions with C2H2+ ions or CHx+ fragments. Neutral C2H4 exhibits large production rates for all the precursor mixtures since its formation pathway involves CH4 and the CH fragment, which are both abundant. Nitrogen incorporation is an important part of the haze chemistry of Titan (Imanaka & Smith 2010). HCN has been detected in Cassini Composite Infrared Spectrometer observations at 713 cm\u22121 (Coustenis et al. 2008). Although its abundance is seasonal, modeled concentrations of HCN nominally peak near the mesopause at around 108 cm\u22123 (Liang et al. 2007), and HCN mixing ratios are highest in the ionosphere where they are between 0.1% and 1% (H\u00e9brard et al. 2012). According to H\u00e9brard et al.'s network, the dominant pathways for HCN production in the upper atmosphere are HCNH+H\n\n\n\n\n\nHCN+H2 and N+CH3\n\n\n\n\n\nHCN+2H. In the THS plasma, various paths for nitrogen incorporation are included in the chemical network including N+ reactions with CH and CH4 leading to CN+ and HCNH+ as well as the dominant neutral mechanisms highlighted by H\u00e9brard et al. The model suggests that the neutral pathway, N+CH3, is a dominant source of HCN for each of the N2\/CH4\/CxHy mixtures. The reason is that atomic nitrogen and CH3 are abundant products from the precursor fragmentation. The HCNH+H pathway is less productive since HCNH must first be formed from the fragments. The model also suggests that once formed, HCN is stable, subject only to small rates of charge transfer with N2(X, v = 0)+ or reaction with the C2H5+ ion.","Citation Text":["Imanaka & Smith 2010"],"Citation Start End":[[1048,1068]]}
{"Identifier":"2019MNRAS.490.5478W__Winter_et_al._2018b_Instance_2","Paragraph":"A growing body of work suggests that planet formation is strongly dependent on the birth environment of the host star. Stars preferentially form in groups (Lada & Lada 2003), and in sufficiently dense environments the evolution of a PPD can be significantly influenced by neighbours (de Juan Ovelar et al. 2012). Close star\u2013disc encounters are one such environmental influence on PPDs that can result in enhanced accretion and hasten disc depletion (Clarke & Pringle 1993; Ostriker 1994; Pfalzner et al. 2005; Olczak, Pfalzner & Spurzem 2006; Bate 2018; Winter et al. 2018a; Cuello et al. 2019). However, the stellar number densities required for tidal truncation are high, and in practice few observed regions satisfy this condition (Winter et al. 2018b, 2019a). The influence of tidal truncation is therefore limited to stellar multiples, either in bound systems (Dai et al. 2015; Kurtovic et al. 2018) or during the decay of higher order multiplicity (Winter, Booth & Clarke 2018c). Since stellar multiplicity does not appear to be strongly dependent on environment (see Duch\u00eane & Kraus 2013, for a review), this suggests that encounters are not an environmental influence, but may set disc initial conditions during the early phases of cluster evolution (Bate 2018). Discs can also be externally depleted via thermal winds driven by far-ultraviolet (FUV) and extreme ultraviolet (EUV) photons from neighbouring massive stars (Johnstone, Fabian & Taylor 1998; St\u00f6rzer & Hollenbach 1999; Adams et al. 2004; Facchini, Clarke & Bisbas 2016; Haworth et al. 2018; Haworth & Clarke 2019). This process of external photoevaporation dominates over dynamical encounters in observed environments, and can deplete PPDs rapidly for many stars that are born in massive and dense clustered environments (Scally & Clarke 2001; Winter et al. 2018b). Many stars in the solar neighbourhood are born in regions where UV fields are sufficient to significantly shorten disc lifetimes (Fatuzzo & Adams 2008; Winter et al. 2018b), and the fraction of stars born in such environments may be much greater outside of this region, dependent on galactic environment (Winter et al. 2019a). From an observational perspective, Guarcello et al. (2016) report disc survival fractions that decrease with increasing FUV flux in Cygnus OB2 (see also Winter, Clarke & Rosotti 2019b), and Ansdell et al. 2017 find a correlation between the dust mass in PPDs and separation from \u03c3 Ori. However, Richert et al. (2015) find no correlation of disc fraction with distance from OB stars. Reconciling these contradictory findings may require appealing to the inefficiency of external photoevaporation at small radii within the disc, dynamical and projection effects, or the stellar age gradient apparent in many star forming regions (Getman et al. 2018).","Citation Text":["Winter et al. 2018b"],"Citation Start End":[[1815,1834]]}
{"Identifier":"2016MNRAS.461..344P__Bernardi_2009_Instance_1","Paragraph":"The fact that all massive ETGs fall into a single FP, coupled with the segregation of the highest ranked objects from lesser galaxies, entails that any non-edge-on projection of this flat surface must also lead to segregated 2D scaling laws. To investigate this issue we now use M\u22c6 to represent scale, shifting to the (logarithmic) RVM coordinate system defined by the three most basic global parameters connected by the standard plane. In this manner we set the framework for the study of two of the most firmly established empirical scaling relations of elliptical galaxies: the Kormendy-like RM relation (Shen et al. 2003) and the VM relation, a mass analogue of the classical FJR. In good agreement with observational studies of BCG\/BGGs (Bernardi et al. 2007; Lauer et al. 2007; Bernardi 2009; M\u00e9ndez-Abreu et al. 2012), we find that our simulated first-ranked galaxies lie off the standard 2D scaling relations defined by the bulk of the ETG population. In particular, as shown in Fig. 3, BGGs are larger and have lower effective velocity dispersions than ordinary ellipticals of the same stellar mass. The combination of this latter result with the fact that light appears to be similarly concentrated for BGGs than for non-BGGs (Paper I), tells us that the total mass-to-light fraction interior to Re is lower for the former than for the latter. Our numerical experiments, however, do not seem to support claims of the steepening of the R \u221d M\u03b1 relation for CGs towards values \u03b1 \u2273 1 (Lauer et al. 2007; Bernardi 2009). We find instead that our first-ranked objects, which occupy a locus moderately offset from the central axis of the observational data in Fig. 3A, are well fitted by a model with \u03b1 = 0.60 \u00b1 0.03. This value of the power-law index agrees well with the slope \u03b1 \u223c 0.65 found for non-BCG galaxies in the Sloan r-band (Shen et al. 2003; Bernardi et al. 2007), while a plain orthogonal fit to the general elliptical population data included in Fig. 3A (green dots and small grey circles) also gives \u03b1 \u223c 0.6. Approaching closer to our findings, Liu et al. (2008) obtain \u03b1 \u223c 0.9 from surface measurements up an isophotal limit of 25 r-mag arcsec\u22122, which reduces to \u223c0.8 when their magnitudes are transformed into mass in old stars, whereas they get \u03b1 \u223c 0.75 for a control sample of non-BCGs. Differences in the photometry, the waveband of the observations, sample construction (i.e. incompletenesses and selection biases), and fitting methods would help explain the lack of a closer agreement between these results about the most robust of the 2D relationships.","Citation Text":["Bernardi 2009"],"Citation Start End":[[784,797]]}
{"Identifier":"2021AandA...646A.144S__Molaro_et_al._2013a_Instance_1","Paragraph":"In general, a precise measurement of the instrumental line-spread function can only be obtained from the LFC lines since only these are truly unresolved. Unfortunately, the LFC covers only 57% of the ESPRESSO spectral range. So even if a proper modeling of the line-spread function would resolve some of the discovered discrepancies, it might only be available for a limited wavelength range. The transitions used to constrain the fine-structure constant are, however, distributed over a wide range in wavelengths. The exact location of the lines depends of course on the absorber redshift, but restricting the range to the region covered by the LFC (approximately 5000 \u00c5\u20137000 \u00c5, see Fig. 8) would render many well-studied absorption systems basically unobservable with ESPRESSO. For example, for the zabs\u2004=\u20041.69 systems along the HE 2217-1818 sightline (Molaro et al. 2013a), the Mg\u202fII lines would not be covered anymore, leaving only the Fe\u202fII lines around \u2243 2500 \u00c5 for a constraint of the fine-structure constant (compare to Fig. 1). However, these five lines have a limited spread in sensitivity with respect to \u03b1. In addition, relying only on the LFC wavelength range will make it at any redshift impossible to observe the Fe\u202fII 1608 \u00c5 line together with the other Fe\u202fII lines at \u03bbrest > 2300 \u00c5. Usually, this is a very valuable combination of transitions, since the Fe\u202fII 1608 \u00c5 line shifts in opposite direction compared to the other Fe\u202fII lines and using only transitions originating from the same ion avoids possible systematics related to different velocity structures and isotopic abundances. For many absorption systems, it is thus crucial to use the full spectral range of ESPRESSO and a restriction to the limited range covered by the currently installed laser frequency comb would substantially diminish the scientific return. Therefore, it is important that future improvements of the ESPRESSO wavelength calibration are applicable to the full spectral range of ESPRESSO.","Citation Text":["Molaro et al. 2013a"],"Citation Start End":[[855,874]]}
{"Identifier":"2015MNRAS.446.1140T__Murray_et_al._2010_Instance_1","Paragraph":"Recently, forms of feedback that are fundamentally different from SNe have been shown to be essential to galaxy formation. Murray, Quataert & Thompson (2010) analysed the dynamical effects of several forms of stellar feedback on parent molecular clouds. In their models they include momentum input from ionized gas in H ii regions, shocked stellar winds, hot gas pressure, protostellar jets and cosmic rays. Murray et al. (2010) conclude that radiation pressure (RP) on dust grains is likely to be the dominant form of feedback in star-forming galaxies. A variety of other studies have reached the same conclusions, placing the combination of RP and photoionization of gas by massive stars as the dominant mechanism for disruption of molecular clouds and internal regulation of the SF process (Indebetouw et al. 2009; Krumholz & Matzner 2009; Murray et al. 2010; Andrews & Thompson 2011; Hopkins, Quataert & Murray 2011; Lopez et al. 2011; Pellegrini, Baldwin & Ferland 2011). RP alone might also be the only mechanism that explains galactic fountains and the warm gas outflows observed in absorption in high-redshift galaxies (Murray, M\u00e9nard & Thompson 2011). In addition, recent numerical work by Krumholz & Thompson (2012) shows that radiation feedback fully accounts for the large gas velocity dispersions measured in young star clusters in the MW. There are at least three reasons why radiative feedback is an essential ingredient of the galaxy formation process. First, observations show that molecular clouds begin to disperse shortly after the O stars form and before the first SNe explode and deposit their energy into the gas (Kawamura et al. 2009). Secondly, the total energy output of a stellar cluster is dominated by radiation. The rate of radiative energy output by O and B stars is \u223c200 times larger than the average power injected by SNe and stellar winds during the lifetime of the most massive stars. Thirdly, it is difficult to explain the large gas turbulence values observed in star-forming regions without including the momentum input by radiation (Murray et al. 2010).","Citation Text":["Murray et al. (2010)"],"Citation Start End":[[408,428]]}
{"Identifier":"2019ApJ...881L..31Y__Chisham_et_al._1998_Instance_1","Paragraph":"The mirror mode is a fundamental process in magnetized plasma environments. It plays many important roles in solar physics, interplanetary, planetary, astrophysical, and laboratory plasma environments, for example, it converts electromagnetic energy and particle kinetic energy (Kivelson & Southwood 1996), excites plasma waves and instabilities (Smith et al. 1969), modulates particle distributions (Yao et al. 2018b), and is suggested as a significant source of turbulent energy (Pokhotelov et al. 2003) and serves as messengers of solar corona (Russell et al. 2008). The most common mirror mode is in space, for example, it is observed to be full of Earth magnetosheath and can exist in the Earth magnetosphere for a long time. These structures are generally suggested to be magnetohydrodynamics (MHD) scale and static in the plasma frame. They display anticorrelation between magnetic and plasma pressures. The spatiotemporal scales and three-dimensional (3D) structures, evolution processes, and other important effects (e.g., drift, finite Larmor radius effect, non-Maxwellian ion distribution, and electron temperature influence) have been extensively studied in the past half century (e.g., Hasegawa 1969; Southwood & Kivelson 1993; Chisham et al. 1998; Pokhotelov et al. 2000, 2013; Gary & Karimabadi 2006; Genot et al. 2009; Hellinger et al. 2009). The studies on the mirror mode were mostly confined to the MHD scale. Numerical simulations and theoretical studies further revealed their kinetic effects, although the fundamental theoretical framework is still at MHD scale and lacks support from observations. If the electron distribution is anisotropic (with perpendicular temperature exceeding the parallel), the electrons become unstable and would excite the electron mirror-mode instability to remove the free energy. This instability was first discovered by Basu & Coppi (1982), and subsequently detailed by Basu & Coppi (1984), Migliuolo (1986), Marchenko et al. (1988), and more recently by Gary & Karimabadi (2006), Pokhotelov et al. (2013), and Hellinger & \u0160tver\u00e1k (2018). This electron mirror-mode instability condition is similar to the MHD-scale mirror mode and can be expressed as \n\n\n\n\n\n, where \u03b2e\u22a5, Te\u22a5, and Te\u2225 denote electron perpendicular plasma beta and perpendicular and parallel electron temperature, respectively. This has been theoretically predicted using a linear calculation (Pokhotelov et al. 2000) and also found in more precise numerical simulations (for example, including the effect of off-diagonal terms in dielectric tensor elements for the mirror instability) by Noreen et al. (2017). The contribution of electrons in the mirror instability in the quasilinear regime was then also discussed by Noreen et al. (2017). Historically, the electron mirror mode has been studied in theories and numerical simulations. But so far there have been few unambiguous observations of their existence because the capability to resolve electron-scale structures has been strongly limited by the insufficient measurement resolution of satellite-borne instrumentation. Using magnetic field data obtained from the spacecraft AMPTE-IRM and Equator-S, Treumann & Baumjohann (2018) presented observational evidence for the electron mirror mode in the Earth's magnetosheath. Low-frequency whistler waves were also observed and were used to diagnose the effect of anisotropic electron temperature on the mirror mode. These observations suggest that the whistler waves were probably related to the trapped electrons for which the trapping condition was provided by the magnetic depressions caused by the electron mirror branch. The study also pointed out that high temporal resolution observations are required for the confirmation of the electron mirror mode and its responsibility for electron trapping and wave excitation. In the recent Magnetospheric Multiscale (MMS; Burch et al. 2016) studies, similar low-frequency whistler waves were observed inside mirror-mode structures by Ahmadi et al. (2018) with a detailed in-depth investigation. The electron temperature anisotropy was suggested to provide a favorable environment for the growth of whistler waves. This result indicates the possible existence of localized electrons that were trapped by the electron mirror branch. Similar observations can be found in Breuillard et al. (2018), which observed whistler waves in mirror-mode structures with perpendicular electron temperature anisotropy.","Citation Text":["Chisham et al. 1998"],"Citation Start End":[[1241,1260]]}
{"Identifier":"2019ApJ...883...76R__Vignali_et_al._2003_Instance_1","Paragraph":"Previous observations of AGN that investigate correlations between \u03b1OX and Eddington ratio have revealed some similarities with X-ray binary outbursts at high Lbol\/LEdd, but these comparisons have not been possible below the critical Lbol\/LEdd \u2272 10\u22122, where an inversion in this correlation is predicted to occur. At higher Eddington ratios of Lbol\/LEdd \u2273 10\u22122, single-epoch X-ray and UV observations of large samples of AGN have previously revealed a hardening of \u03b1OX as Lbol\/LEdd drops from \u223c1 to \u223c10\u22122 (e.g., Vignali et al. 2003; Strateva et al. 2005; Steffen et al. 2006; Just et al. 2007; Grupe et al. 2010; Jin et al. 2012; Wu et al. 2012; Trichas et al. 2013; Vagnetti et al. 2013). This correlation was also observed in multi-epoch UV\/X-ray observations of the fading of Mrk 1018 (Noda & Done 2018), which confirms this behavior in an individual AGN. However, the predicted softening of \u03b1OX below Lbol\/LEdd \u2272 10\u22122 (thus causing an inversion in the correlation between \u03b1OX and Lbol\/LEdd) has not been previously observed. This is primarily due to the difficulty of robustly measuring both \u03b1OX and Lbol\/LEdd for AGN below Lbol\/LEdd \u2272 10\u22122, for three main reasons. First, at low Eddington ratios, AGN are often dust-obscured (Fabian et al. 2008), and thus measuring their intrinsic UV luminosities (and \u03b1OX) is difficult. Second, broad emission lines often disappear in low-luminosity AGN below Lbol\/LEdd \u2272 10\u22122, making it difficult to measure MBH (and LEdd). Third, using a sample of AGN with a wide range of Eddington ratios to trace how \u03b1OX changes as a function of Lbol\/LEdd can be hampered by the \n\n\n\n\n\n scaling of the thin disk temperature with MBH at a fixed Eddington ratio. If the AGN sample has a large range in MBH, this can cause an additional scatter in \u03b1OX. Thus, we would ideally use a sample of AGN with a narrow range in MBH, but the difficulty of measuring MBH at Lbol\/LEdd \u2272 10\u22122 also hampers the construction of such a sample. In this paper, we will use a new method to bypass all of these issues, with the goal of extending this spectral comparison between X-ray binaries and AGN to Lbol\/LEdd \u2272 10\u22122.","Citation Text":["Vignali et al. 2003"],"Citation Start End":[[512,531]]}
{"Identifier":"2021ApJ...923..105L__Owens_et_al._2013_Instance_1","Paragraph":"For the purpose of this study, we adopt the general concept that the global solar wind is generally composed of relatively steady high-speed solar winds from long-lived high latitude open coronal field regions. The high-speed solar winds from the northern and southern hemispheres are separated by a layer of\u2014on average\u2014slower, denser, more complicated, and variable outflows and structures (Gosling 1996; McComas et al. 2008). This boundary layer between the polar outflows has a latitudinal extent that is roughly defined by the heliospheric current sheet (HCS) maximum excursion. In the solar wind, there also exist various anomalous episodes, as consequences of transient and\/or energetic activities that may be remote (in the corona) or heliospheric in origin. The following phenomena or transients relating to anomalous solar wind parameters have been studied and have inspired our study. (1) False magnetic field polarity reversals\u2014evoking images of switchbacks\u2014that have been argued as being interchange reconnection debris (Owens et al. 2013, 2017, 2020), and recently also observed much closer to the Sun by the Parker Solar Probe mission (Laker et al. 2021), and the origin is being debated. (2) Heliospheric magnetic field (HMF) departures from the Parker spiral to nearly radial field at 1 au (Murphy et al. 2002; Watari et al. 2005). The great majority of the initial mass function orientation statistics fall in the vicinity of Parker spiral fields for typical solar wind speeds. However, blocks of time occasionally stand out when the HMF is nearly radial. (3) Interplanetary coronal mass ejections (ICMEs), magnetic clouds (MCs), and other flux ropes, whose structures are foreign to the ambient solar wind and whose plasma differs from the background plasma, often bring anomalously low temperatures and plasma beta (Burlaga et al. 1981; Klein & Burlaga 1982; Burlaga 1988; Richardson & Cane 1995; Kilpua et al. 2009). (4) The intervals of counterstreaming suprathermal electrons (CSEs), which are often interpreted as closed loops of magnetic fields with both ends being connected to a hot electron source within the corona (Bame et al. 1981; Gosling et al. 1987, 1992). (5) Heat flux dropouts, suggesting complete disconnection from hot electron sources (McComas et al. 1989). Many studies have been devoted to these phenomena, including their relationship to the location of the HCS\/coronal streamer belt and their solar cycle variations (McComas et al. 1989; Fitzenreiter & Ogilvie 1992; Gosling et al. 1992, 1993; Shodhan et al. 2000; Murphy et al. 2002; Watari et al. 2005; Lavraud et al. 2010; Yu et al. 2014, 2016; Owens et al. 2017, 2020). Weaker solar cycles and advanced imaging and modeling capabilities have raised the profile of pseudostreamers and their boundaries as contributors to the outflows within this slow wind layer (Wang et al. 2012; Owens et al. 2013, 2014). Similarly, the density blobs that are released from the ends of the coronal streamers are another source of slow solar wind (Sheeley & Wang 2007; Rouillard et al. 2010a, 2010b) that belong to this collection of features. These have been associated with in situ observations of small flux rope-like structures in the solar wind at 1 au (Kilpua et al. 2009; Rouillard et al. 2010a, 2010b; Viall & Vourlidas 2015; Sanchez-Diaz et al. 2017; Lavraud et al. 2020). Clearly, there are many reasons and sources of evidence for the occurrence of anomalous characteristics that differ from the typical ambient solar wind. The question is whether they occur in organized ways that provide insights into their origins.","Citation Text":["Owens et al. 2013"],"Citation Start End":[[1033,1050]]}
{"Identifier":"2021MNRAS.501.5408W__Lilley_et_al._2018a_Instance_1","Paragraph":"To study this, one must represent the dynamics in coordinates that focus on the mechanism. Action-angle expansions do this very well and are useful for systems which do not evolve with time (Binney & Tremaine 2008). However, we know that galaxies do evolve in time, and this time dependence influences the evolution itself. The nature of galactic dynamics requires analysis that can accurately treat multiple time-scales. For example, interesting evolution takes place on time-scales longer than the characteristic period, largely driven by resonant exchange and phase mixing. To do this efficiently, one may follow the gravitational field in evolving systems using basis-function expansions (BFE). BFE represents the potential and density of near-equilibrium galaxies in basis functions: pairs of biorthogonal functions that satisfy the Poisson equation and are mutually orthogonal (Clutton-Brock 1972, 1973; Kalnajs 1976; Fridman & Polyachenko 1984; Hernquist & Ostriker 1992; Weinberg 1999; Lilley et al. 2018a; Lilley, Sanders & Evans 2018b). Each of these basis sets (op. cit.) share the same properties but differ in the dimensionality, node location, profile shape, and method of construction. Most of these use integral transform methods following Clutton\u2013Brock and Kalnajs. Alternatively, Weinberg (1999) solves radial part of the Poisson equation directly to best represent the underlying galaxy density profile. We adopt this method here. In effect, this windows the length scales to particular ones of interest based on the equilibrium model. We describe the relationship between these methods in the next section. Additionally useful, one can take any set of biorthogonal functions and create a new biorthogonal set of functions by a vector-space rotation. In particular, one may select the rotation by diagonalizing the variance matrix of the expansion coefficients given the particle distribution; this generates a new set of orthogonal functions that looks like the particle distribution (Petersen, Weinberg & Katz 2021). This new data-motivated basis is sometimes called empirical orthogonal functions (EOFs) after Lorenz (1956).","Citation Text":["Lilley et al. 2018a"],"Citation Start End":[[994,1013]]}
{"Identifier":"2018MNRAS.476.2689B__Woosley_&_Weaver_1995_Instance_1","Paragraph":"We only recall here the main features of this model. Heavy elements are produced through three main enrichment channels, namely by SNIa, SNII, and AGB stars, depending on their typical lifetimes and yields. To this purpose, we assume the initial mass function by Chabrier (2003) as a distribution of initial masses for the population of stars, and the mass-dependent lifetimes by Padovani & Matteucci (1993) to account for the different time-scales of stars of different masses. The abundance of the various metal species produced during the evolution of a stellar particle is then estimated by considering different sets of stellar yields for SNIa, SNII, and AGB (namely Thielemann et al. 2003, for SNIa; Woosley & Weaver 1995 combined with those by Romano et al. 2010, for SNII; Karakas 2010, for AGB stars). Once metals are produced, they are distributed from stars to surrounding gas particles by smoothing them on the SPH kernel, consistently with the other thermodynamical quantities, while no process of metal diffusion is explicitly included in the simulations. Previous studies in the literature have investigated the influence of changing the details of the metal spreading on the final enrichment pattern in simulated clusters (e.g. Tornatore et al. 2007; Wiersma et al. 2009b). From various tests presented in Tornatore et al. 2007 for the same chemical evolution model used in the current simulations, we conclude that changes in the numerical details of the metal injection scheme (e.g. number of neighbouring particles or weighting scheme for the metal diffusion to the gas) have only a minor impact on the resulting star formation history, stellar population, and chemical enrichment pattern. The differences, in particular, are smaller than the scatter from cluster to cluster. Moreover, for a given scheme of metal spreading, the final enrichment level of the simulated clusters is more strongly influenced by the physical processes included in the simulations, such as the energy feedback (see Fabjan et al. 2010; Planelles et al. 2014; Biffi et al. 2017).","Citation Text":["Woosley & Weaver 1995"],"Citation Start End":[[706,727]]}
{"Identifier":"2022ApJ...925..203J__Malkan_et_al._1998_Instance_1","Paragraph":"Besides AGN feedback, additional clues on the SMBH\u2013galaxy connection can be obtained from constraints on AGN obscuration, which can also occur on a range of physical scales (see review by Ramos Almeida & Ricci 2017). The most basic AGN unification model predicts that the obscuring medium is a small (parsec-scale) nuclear torus surrounding the active black hole (BH), and the degree of obscuration simply depends on the viewing angle (Antonucci 1993; Urry & Padovani 1995). However, there is compelling evidence for AGN obscuration being closely related to their host galaxies. For example, previous work has revealed different hosts and dust structures between Seyfert 1\u2019s and 2\u2019s down to fine angular scales, strongly disfavoring the conventional torus orientation picture (e.g., Malkan et al. 1998; Prieto et al. 2021). Furthermore, spatially resolved mid-IR or hard-X-ray studies of nearby heavily absorbed systems found that the obscuration is fully consistent with moderate- to large-scale dust lanes and features with a range of hydrogen column densities (Bianchi et al. 2007; Ar\u00e9valo et al. 2014; Bauer et al. 2015), and with AGN ionization cone collimation taking place at distances well beyond the inner torus (Prieto et al. 2014; Mezcua et al. 2016). Additionally, the observed relationship between the 9.7 \u03bcm silicate absorption depth and the inclination of the host galaxies of heavily absorbed AGNs support large-scale obscuration (Goulding et al. 2012), and an intriguing trend for intermediate redshift star-forming galaxies points to an increasing fraction of heavily absorbed AGNs with increasing galaxy-wide specific star formation rate (= SFR\/stellar mass, a tracer of galaxy gas fractions; Juneau et al. 2013). The latter suggests either a direct role of host galaxies gas in the absorption of X-rays from the AGN or a true physical link between the small-scale torus and the multiscale interstellar medium. Some studies find cases with high X-ray absorption (i.e., Compton-thick regime), which are consistent with small-scale torus absorption (e.g., Balokovi\u0107 et al. 2014; Markowitz et al. 2014; Ricci et al. 2014; Koss et al. 2017), or evidence for both nuclear and possible host galaxy contributions, where the latter might be dependent on galaxy mass and redshift (Brightman et al. 2014; Buchner et al. 2015; Buchner & Bauer 2017). Therefore, assessing AGN feedback, AGN obscuration, and the possible interplay with host galaxy properties and substructure can shed light on the SMBH\u2013galaxy connection.","Citation Text":["Malkan et al. 1998"],"Citation Start End":[[783,801]]}
{"Identifier":"2021MNRAS.502.1933S__Hayashi_et_al._2007_Instance_1","Paragraph":"The observed two-point correlation function \u03c9obs(\u03b8) is corrected for an underestimation due to a finite coverage of the observation data. The correction is conducted by adding an integral constraint (IC; Groth & Peebles 1977) as follows:\n(6)$$\\begin{eqnarray}\r\n\\omega (\\theta) = \\omega _{\\rm obs}(\\theta) + \\rm {IC}.\r\n\\end{eqnarray}$$With the assumption of a single power-law correlation function \u03c9(\u03b8) = A\u03c9\u03b8\u2212\u03b4, IC is expressed by the formalism in Roche et al. (1999) as\n(7)$$\\begin{eqnarray}\r\n\\mathrm{IC} = \\frac{\\Sigma _{i}[RR(\\theta _{i})\\omega (\\theta _{i})]}{\\Sigma _{i} RR(\\theta _{i})},\r\n\\end{eqnarray}$$where RR(\u03b8i) is the non-normalized number of random\u2013random pairs. Combined with equation (7), equation (6) becomes\n(8)$$\\begin{eqnarray}\r\n\\omega _{\\rm obs}(\\theta) = A_{\\omega }\\biggl [\\theta ^{-\\delta } - \\frac{\\Sigma _{i}[RR(\\theta _{i})\\theta _{i}^{-\\delta }]}{\\Sigma _{i} RR(\\theta _{i})}\\biggr ].\r\n\\end{eqnarray}$$The observed two-point correlation function \u03c9obs(\u03b8) is fitted by equation (8) to obtain the corrected two-point correlation function \u03c9(\u03b8). We conducted the fit by fixing a slope parameter \u03b4 to 0.8. Since \u03b4 is sensitive to the sample size, constraining the slope \u03b4 is not effective because of the small number of sgzK galaxies. The result with fixed \u03b4 is useful for direct comparison with previous studies since many studies assumed the slope parameter \u03b4 fixed to 0.8 (Kong et al. 2006; Hayashi et al. 2007; Blanc et al. 2008; Hartley et al. 2008; Ishikawa et al. 2015). The least-squares fit procedure in mpfit1 (Markwardt 2009) was used to determine A\u03c9. The two-point correlation functions for sgzKs with Ks  21.1 and 21.5 are presented in Fig. 4. Ishikawa et al. (2015) reported that the two-point correlation function of sgzKs at small angular scale (\u03b8 \u2272 0${_{.}^{\\circ}}$01) has a steeper slope compared to that at large angular scale, which represents the difference of contributions from the one-halo term (pairs of sgzKs in the same halo) and the two-halo term (pairs of sgzKs in different haloes) as predicted from the analysis by the dark matter halo model. However, such characteristic is unclear in Fig. 4 due to the small number of sgzK galaxies. The results of clustering analysis for the two magnitude-limited sgzKs are listed in Table 1. In Fig. 5, the amplitudes of two-point correlation functions for sgzKs are presented together with the results in the literature. For the estimation of the error, we considered the effect of cosmic variance. For this, we used four sub-regions presented in Fig. 1. Two-point correlation functions were estimated four times separately after excluding each of the sub-regions in turn. A standard deviation of the four two-point correlation functions can be treated as the uncertainty of two-point correlation function due to cosmic variance. It was added quadratically in the estimation of the error of two-point correlation function. As shown in Fig. 5, our result is comparable to the previous studies. In addition to the results in the literature, our result shows Ks-band magnitude dependence of the clustering strength.","Citation Text":["Hayashi et al. 2007"],"Citation Start End":[[1415,1434]]}
{"Identifier":"2022AandA...666A.102H__Schinnerer_et_al._(2000)_Instance_1","Paragraph":"NGC1068 is a nearby (D\u2004=\u200414 Mpc Bland-Hawthorn et al. 1997, 1\u2033\u2004\u223c\u200470 pc) Seyfert 2 galaxy and is considered to be the archetype of a composite AGN-starburst system. The proximity of this composite galaxy makes it an ideal laboratory for resolving the feedback from the starburst regions that are spatially distinct from the AGN activity. NGC 1068 has been extensively investigated by many single-dish and interferometric campaigns focused on the study of how its central region is fueled and the related feedback activity, using molecular line observations (e.g., Usero et al. 2004; Israel 2009; Kamenetzky et al. 2011; Hailey-Dunsheath et al. 2012; Aladro et al. 2013; Garc\u00eda-Burillo et al. 2014, 2017, 2019; Viti et al. 2014; Impellizzeri et al. 2019; Imanishi et al. 2020). The available CO Observations of NGC 1068 by Schinnerer et al. (2000) reveal the molecular gas distributing over three regions, which has also been confirmed by, for instance, Garc\u00eda-Burillo et al. (2014, 2019) and S\u00e1nchez-Garc\u00eda et al. (2022): a starburst ring (SB ring) with a radius \u223c1.5 kpc, a circumnuclear disk (CND) of radius \u223c200 pc, and a \u223c2 kpc stellar bar running north east, along PA \u223c48\u00b0 (Scoville et al. 1988), from the CND. In Garc\u00eda-Burillo et al. (2014) and Viti et al. (2014), five chemically distinct regions were found to be present within the CND: the AGN, the East Knot, West Knot, and regions to the north and south of the AGN (CND-N and CND-S) using data from the Atacama Large Millimeter\/submillimeter Array (ALMA). Viti et al. (2014) combined these ALMA data with Plateau de Bure Interferometer (PdBI) data and determined the physical and chemical properties of each region. It was found that a pronounced chemical differentiation is present across the CND and that each sub-region could be characterized by a three-phase component interstellar medium, where one of the components is comprised of shocked gas. In fact, Garc\u00eda-Burillo et al. (2010) used the PdBI to map NGC 1068 and found strong emission of SiO(2\u20131) to the east and west of CND. The SiO kinematics of the CND point to an overall rotating structure and is distorted by non-circular or non-coplanar motions, or a combination of both. The authors have concluded that this could be due to large scale shocks through cloud-cloud collisions or through a jet-ISM interaction. Such shock-related non-circular kinematics of gas was also identified by Krips et al. (2011) using several molecular ratios of CO, 13CO, HCN, and HCO+. However, due to strong CN emission that is not easily explained by shock models nor photon-dominated region (PDR) chemistry, these authors also suggested that the CND could actually be one large X-ray dominated region (XDR).","Citation Text":["Schinnerer et al. (2000)"],"Citation Start End":[[821,845]]}
{"Identifier":"2016AandA...588A..25M__Weiss_&_Ferguson_2009_Instance_2","Paragraph":"There are some indications that available models of the post-AGB and CSPN phases are not accurate enough. First, the two available grids of post-AGB models (Vassiliadis & Wood 1994; Bl\u00f6cker 1995a) do not agree with each other on the predicted timescales (Zijlstra et al. 2008). Second, consistency between the masses of white dwarfs and those of CSPNe seems to require faster evolutionary speeds than predicted by both sets of models (Gesicki et al. 2014). Third, present models of the CSPNe phase are unable to explain why the cut-off of the PNe luminosity function is constant in most galaxies (Marigo et al. 2001, 2004). Lastly, post-AGB stellar evolution models, computed with updated physics in a reduced mass range (Kitsikis 2008; Weiss & Ferguson 2009), show a strong disagreement with the previous grids. This is not a surprise since many improvements have been carried out in the field of stellar physics in recent decades. Most importantly, available grids have been computed with opacities, which are now 45 years old (Cox & Stewart 1970b,a) before the big changes introduced by the OPAL (Iglesias & Rogers 1996), and Opacity Project (Seaton 2007) redeterminations. Similarly, nuclear reaction rates, equation of states, conductive opacities, and neutrino emission rates adopted in the models date from the early eighties and even earlier. In addition, Herwig et al. (1997) showed that the existence of carbon stars at low luminosities can be explained by the addition of mixing beyond the formal convective boundaries during the thermal pulses (TP) on the AGB. Finally, Marigo (2002) showed that C-rich molecular opacities are essential to predict the correct effective temperatures once the AGB models become carbon rich (NC\/NO> 1, by number fractions). This is particularly important because of the impact of effective temperatures on the mass loss rates. While all these improvements in stellar modeling have been implemented in AGB stellar models, and very detailed and exhaustive grids and models are available (Weiss & Ferguson 2009; Cristallo et al. 2009, 2011; Ventura & Marigo 2010; Karakas 2010; Lugaro et al. 2012; Constantino et al. 2014; Doherty et al. 2015), the inclusion of these improvements in post-AGB stellar models is still missing. It is time for a recomputation of the post-AGB models in the light of all these advances. ","Citation Text":["Weiss & Ferguson 2009"],"Citation Start End":[[2029,2050]]}
{"Identifier":"2022ApJ...924...38K__Bazavov_et_al._2019_Instance_1","Paragraph":"Recently, another explosion mechanism has gained more attention in the context of both the explosion dynamics and formation process of neutron stars composed of exotic matter at their interiors (see Fischer et al. 2018; Zha et al. 2020). When the energy density becomes sufficiently high, then nuclear\u2014in general, hadronic\u2014matter undergoes a phase transition to the deconfined quark\u2013gluon plasma. Therefore, the solution of quantum chromodynamics (QCD)\u2014the theory for strong interactions with quarks and gluons as the degrees of freedom\u2014predicts a smooth crossover transition at a pseudocritical temperature in the range of 150\u2013160 MeV (Bazavov et al. 2014; Bors\u00e1nyi et al. 2014; Bazavov et al. 2019). However, this is the case only at vanishing baryon density, and hence cannot be applied to astrophysical applications of compact stellar objects such as neutron stars and supernovae. At high density, phenomenological quark matter models have long been employed in astrophysical studies, based on the common two-phase approach with separate hadronic and quark matter phases. Such phase transition construction results not only in the transition to the third family of compact stars, known as hybrid stars, as an extension to the neutron stars second family, but also in the as-yet incompletely understood question of the nature of the QCD phase transition at high baryon density (see Annala et al. 2020, and references therein). It is likely that low- and intermediate-mass NSs, having masses of \u223c1.4 M\n\u2299, can be explained solely by hadronic equations of state (EOS). More massive NSs with \u223c2 M\n\u2299, the presence of which has been observationally confirmed at high accuracy (see. Demorest et al. 2010; Antoniadis et al. 2013; Cromartie et al. 2020), feature the highest baryon densities encountered in the universe. It is a valid question to ask if a QCD transition occurs at the densities encountered at the interior of such massive pulsars (see Kurkela et al. 2014, and references therein). Furthermore, if confirmed, this is likely to affect our understanding of the core-collapse supernova phenomenology.","Citation Text":["Bazavov et al. 2019"],"Citation Start End":[[680,699]]}
{"Identifier":"2019AandA...621A..43F__Felipe_et_al._2018_Instance_1","Paragraph":"The magnetic field configuration at the deep photosphere shows a low twist in most of the umbra (Fig. 4b). Only some penumbral regions exhibit a certain twist, which is generally below 30\u00b0. On the contrary, the twist measured at the height of the Si I 10 827 \u00c5 line shows high values at the bottom right quarter, which reach values up to \u221290\u00b0 around the umbra-penumbra boundary. This twist contrasts with that found inside the umbra at the bottom left quarter, which is mainly positive. Socas-Navarro (2005) found that the magnetic twist can change the sign from the photosphere to the chromosphere. Our observations indicate that the twist can change from the deep photosphere (where it is almost negligible) to higher photospheric layers (where negative and positive twists coexist). This variation takes place over a height difference of around 130 km (Felipe et al. 2018). However, the degree of the twist may be sensitive to the choice of the center of the sunspot. In this context, it can be more meaningful to discuss the magnetic field torsion. Figure 4e shows that the bottom part of the penumbra has a slightly negative torsion, with stronger values near the umbra-penumbra boundary (in the same region where the twist from the Si I 10 827 \u00c5 line was higher). Since the magnetic field at the formation height of the Ca I 10 839 \u00c5 line is approximately pointing towards the center, a negative torsion indicates that at the formation height of the Si I 10 827 \u00c5 line the horizontal component of the magnetic field is shifted clockwise from the radial direction (see Fig. 5 for a vectorial representation of the horizontal magnetic field). A positive torsion, such as that measured at the left and bottom part of the umbra, means that the magnetic field is shifted counterclockwise. The position on the horizontal plane of the wavefront of a slow magnetoacoustic wave propagating along magnetic field lines from the deep photosphere to the high photosphere will rotate according to the direction of the torsion. As discussed in Sect. 3.1, the bottom arm of the spiral wavefront initially appears at the left side of the umbra and moves counterclockwise towards the center of the sunspot. This movement is in agreement with the positive torsion measured at that region from the analysis of the spectropolarimetric map. On the contrary, at the regions where the measured torsion is negative, we do not detect clockwise motions of the spiral wavefront. In that part we only find a radially outward movement.","Citation Text":["Felipe et al. 2018"],"Citation Start End":[[856,874]]}
{"Identifier":"2019ApJ...876L..28D__Abbott_et_al._2017_Instance_1","Paragraph":"The mergers of double neutron star systems or the neutron star-black hole binaries generate strong gravitational wave (GW) radiation as well as short-duration gamma-ray bursts (SGRBs; including the so-called long-short GRBs that have a duration longer than 2 s but are unaccompanied by supernova emission down to very stringent limits; Eichler et al. 1989; Piran 2004; Della Valle et al. 2006; Fynbo et al. 2006; Berger 2014). Before 2017, it was widely believed that the GW\/SGRB association rate is low because the SGRB outflows are highly collimated with a typical half-opening angle of \u223c0.1 rad (Clark et al. 2015; Li et al. 2016). Surprisingly, on 2017 August 17, the gamma-ray monitor on board the Fermi \u03b3-ray space telescope and INTEGRAL had successfully detected a weak-short GRB 170817A (Goldstein et al. 2017; Savchenko et al. 2017) that is spatially and temporally correlated with GW170817, the first neutron star merger event detected by Advanced LIGO\/Virgo (Abbott et al. 2017). The GW\/SGRB association has been formally established. However, considering the relatively small event distance (D \u223c 40 Mpc), the isotropic-equivalent gamma-ray radiation energy of GRB 170817A is just \u223c3 \u00d7 1046 erg, which is at least 100 times dimmer than that of the typical SGRBs. An underluminous SGRB could either result from the breakout of the mildly relativistic shock from the leading edge of the merger-driven quasi-isotropic sub-relativistic ejecta (Kasliwal et al. 2017), or be the faint prompt emission of a highly structured relativistic ejecta viewed at a large polar angle (Jin et al. 2018). The puzzling fact that GRB 170817A and the long-duration GRB 980425 (at a distance of D \u223c 36 Mpc, and that has been suggested to be the shock breakout signal; Kulkarni et al. 1998), the two close events with remarkably different progenitors, have rather similar luminosity and spectral peak energy (Wang et al. 2017), may favor the shock breakout model. It is thus unclear whether or not GW170817-like mergers are indeed the sources of the bright SGRBs. The forward shock afterglow observations of GW170817\/GRB 170817A are helpful in answering such a question. Though the \u201cearly\u201d rising X-ray and radio afterglow emission could be reproduced by a cocoon-like mildly relativistic ejecta (Kasliwal et al. 2017), the late-time afterglow data modelings strongly favor the presence of an off-axis relativistic (structured) outflow component (D\u2019Avanzo et al. 2018; Lamb et al. 2018; Mooley et al. 2018b; Yue et al. 2018). Particularly, the off-axis relativistic outflow component at a viewing angle of \u03b8v \u223c 0.35 rad has been convincingly identified\/measured in the radio image (Mooley et al. 2018a; Ghirlanda et al. 2019). Nevertheless, a direct \u201cobservational\u201d link between GW170817\/GRB 170817A and bright SGRBs is still lacking. In this Letter, we carry out statistical studies of the SGRB afterglow data and aim to establish such a connection.","Citation Text":["Abbott et al. 2017"],"Citation Start End":[[970,988]]}
{"Identifier":"2019MNRAS.487.1626Q__Fender_2006_Instance_2","Paragraph":"In the coupled ADAF-jet model, the accretion flow ADAF and the jet are connected by a defined parameter, $\\eta \\equiv \\dot{M}_{\\rm jet}\/\\dot{M}$, and $\\dot{M}_{\\rm jet}$ is input by assuming a value of, \u03b7, which is free parameter in the present model. The half-opening angle \u03d5 of the jet in the low\/hard state of NS-LMXBs is uncertain. In this paper, we fix \u03d5 = 0.1 as assumed by several other authors for modelling the SED of the BH-LMXBs (e.g. Yuan et al. 2005; Zhang et al. 2010). Observationally, the bulk Lorentz factor of the jet in the low\/hard state of X-ray binaries can be restricted in a relatively narrow range and the velocity of the jet is mildly relativistic, i.e. \u0393jet \u2272 2. More strictly, the bulk Lorentz factor is restricted to be as \u0393jet \u2272 1.67 (Gallo et al. 2003), and \u0393jet \u2272 1.2 (Fender 2006). In the internal shock model, the energy density of the internal shock increases with increasing \u0393jet, which finally will result in an increase of both the radio emission and X-ray emission (Yuan et al. 2005). However, since \u0393jet is restricted in a very narrow range by observations, we expect that a slight change of \u0393jet will result in a slight change of the jet emission. In this paper, we fix \u0393jet = 1.2 corresponding the bulk velocity of the jet 0.55c (Fender 2006). The value of \u03f5e and \u03f5B describing the fraction of the internal energy of the internal shock stored in the accelerated electrons and the magnetic field, and the index, pjet, describing the power-law distribution of the electrons in the jet after the acceleration by the shock are uncertain. Qiao & Liu (2015) tested the effect of \u03f5e and \u03f5B on the emergent spectrum of the jet in an observationally inferred range of 0.01  \u03f5e  0.1 and 0.01  \u03f5B  0.1. It was found that a change of \u03f5B in the range of 0.01\u20130.1, the emergent spectrum of the jet nearly does not change (see the right-hand panel of fig. 3 of Qiao & Liu 2015). A change of \u03f5e in the range of 0.01\u20130.1, the radio spectrum nearly does not change. However, the X-ray luminosity changes obviously by changing the value of \u03f5e from 0.01 to 0.1 (see the left-hand panel of fig. 3 of Qiao & Liu 2015). As shown in the left-hand panel of fig. 2 of Qiao & Liu (2015), the X-ray emission is completely dominated by the accretion flow (corona) in the luminous X-ray state, which is also true in this paper, i.e. the X-ray emission from the ADAF completely dominates the X-ray emission from the jet. In this paper, we fix \u03f5e = 0.04 and \u03f5B = 0.02, respectively, throughout the paper as (Qiao & Liu 2015). The value of the power-law index pjet of the electron distribution in the jet predicted by the shock acceleration is 2  pjet  3. By modelling the SEDs of three BH-LMXBs, the value of the power-law index pjet of the electron distribution is constrained to be 2.1 (Zhang et al. 2010). Meanwhile, a change of pjet in the range of 2  pjet  3 has very minor effect on the X-ray spectrum. In this paper, we fix the power-law index pjet = 2.1 throughout the paper.","Citation Text":["Fender 2006"],"Citation Start End":[[1272,1283]]}
{"Identifier":"2019AandA...625A..12G__Haqq-Misra_et_al._(2016)_Instance_1","Paragraph":"Another important sink for CO2 is the loss of CO2 by weathering of the surface. We neglected this effect in the present study as in Tosi et al. (2017). Foley & Smye (2018) study the potential impact of weathering for stagnant-lid planets. They find that for large enough planetary carbon reservoirs and radiogenic heating, weathering and outgassing can balance each other for 1000\u20135000 Myr, thus enabling habitable surface conditions. On the one hand, for large carbon reservoirs they show that weathering is supply-limited, i.e. the supply of fresh surface material is likely too small such that CO2 would accumulate in the atmosphere, which leads to a Venus-like, uninhabitable hot climate. On the other hand, for low outgassing rates limited by the carbon reservoir, the planet would likely exist in a snowball state. For the snowball state limit, they use outgassing values determined by Haqq-Misra et al. (2016) and Kadoya & Tajika (2014) as lower limits, which are equal to 10 and 100% of present Earth\u2019s outgassing, respectively; 10% of Earth\u2019s CO2 outgassing flux, as adopted by Foley & Smye (2018), corresponds to ~2.6 \u22c5 1010 kg yr\u22121. In our models the rate of CO2 outgassing depends on the initial water concentration and on the assumed redox state. For 62 ppm water in the mantle and \n\n$f_{\\mathrm{O}_2}$\n\n\n\nf\n\n\nO\n2\n\n\n\n\n\n at the IW buffer, the outgassing rate during the phase of melt production peaks at about 2.5 \u22c5109 kg yr\u22121, while for \n\n$f_{O_2}$\n\n\n\nf\n\n\nO\n2\n\n\n\n\n\n at IW+1 it is one order of magnitude higher, and hence comparable with the lower limit mentioned above. Foley & Smye (2018) discuss weathering for Earth-like water reservoirs. In our scenarios, the weathering of the surface could be limited by the smaller surface water reservoir and lower precipitation rates. Therefore, although our outgassing rates are close to the lower limits adopted by Foley & Smye (2018) that might lead to snowball states, the overall impact of weathering on our findings is not easily predictable and will require further study. Furthermore, Foley & Smye (2018) consider metamorphic outgassing as an additional source of CO2, which could increase CO2 outgassing rates. Metamorphic outgassing refers to the release of CO2 associated with the decarbonation of buried carbonated crust. On Earth, the contribution of this mechanism to the planet\u2019s long-termcarbon cycle is thought to be potentially important, but is still poorly understood (Evans 2011). Furthermore, the release of carbon due to metamorphic processes is generally associated with partial melting in subduction zones (e.g. Dasgupta & Hirschmann 2010) and deformation in active orogenic regions (e.g. Evans et al. 2008), two settings that are intimately related to plate tectonics and hence absent in a stagnant-lid planet. Whether metamorphic outgassing on such bodies can be as relevant as on a planet with plate tectonics is therefore unclear and deserves more investigation.","Citation Text":["Haqq-Misra et al. (2016)"],"Citation Start End":[[892,916]]}
{"Identifier":"2021MNRAS.501.4850R__Roe_1981_Instance_1","Paragraph":"The hydrodynamic equations (1\u20135) are solved in this paper using total variation diminishing (TVD) scheme, introduced and developed by Harten (1983). The scheme (or the modified version of it) is applicable to hydrodynamic problems and has been used extensively in relevant astrophysical applications (Ryu 1993; Ryu et al. 1995a; Ryu, Jones & Frank 1995b; Lee, Ryu & Chattopadhyay 2011; Chattopadhyay et al. 2012; Lee et al. 2016). TVD scheme is an Eulerian, second-order accurate, non-linear, finite difference scheme, which accurately captures shock. The temporal and spatial evolution of the conserved quantities \u03c1, \u03c1vi, and ${\\cal E}$ is computed using approximate Roe type Riemann solver to solve the differential equations, followed by application of a non-oscillatory first-order accurate scheme to the modified flux functions to achieve second-order accuracy (see Roe 1981; Harten 1983; Ryu 1993). Equations of motion (1\u20135) are similar to those solved in Chattopadhyay et al. (2012). In Chattopadhyay et al. (2012), the galactic outflow was powered by the radiation from the galactic disc, while being decelerated by the gravity of the galactic disc, the halo, and the bulge matter. In contrast, in this paper, the accretion disc outflow is powered by the radiative fluxes and the centrifugal force from the KD, and is decelerated by the radiative drag terms as well as the gravity of the central black hole. To solve the equations of motion (1\u20135), we considered the TVD scheme (see Chattopadhyay et al. 2012, for details) for the resolution 512 \u00d7 512. A schematic representation of the computational arrangement is presented in Fig. 2, which marks the ghost cells where the boundary conditions are implemented and also the computational domain. We employed continuous boundary condition at z = 0 boundary, and outflow boundary condition at the outer r and z boundaries (i.e. no inflow but continuous if v > 0). At r = 0, or the axis of symmetry, reflection boundary condition has been employed. The type of boundary conditions employed are also mentioned in Fig. 2. We simulate a region of 512 from the black hole, each in r and z direction, therefore the dimension of each cell is equivalent to 1. The gravity of the black hole is described by Paczy\u0144ski and Wiita potential (Paczy\u0144ski & Wiita 1980). In order to avoid the coordinate singularity on the horizon, the black hole is covered by a sink region of radius 3 around the origin, which do not affect the physics since the inner edge of KD is 3.","Citation Text":["Roe 1981"],"Citation Start End":[[871,879]]}
{"Identifier":"2022AandA...659A..10V__Kay_et_al._2017_Instance_1","Paragraph":"The CMEs are solar eruptions caused by magnetic reconnections in the star corona (Low 2001; Howard 2006). They expel a large amount of fast charged particles and a magnetic cloud that evolves into an interplanetary coronal mass ejection (ICME; Sheeley Jr. et al. 1985; Neugebauer & Goldstein 1997; Cane & Richardson 2003; Gosling 1990). If the ICME impacts the Earth, the measured SW dynamic pressure increases to 10\u2212100 nPa and the IMF intensity to 100\u2013300 nT (Gosling et al. 1991; Huttunen et al. 2002; Manchester IV et al. 2004; Schwenn et al. 2005; Riley 2012; Howard 2014; Mays et al. 2015; Kay et al. 2017; Savani et al. 2017; Salman et al. 2018; Kilpua et al. 2019; Hapgood 2019). The Disturbance Storm Time Index (Dst) indicates the magnetic activity derived from a network of near-equatorial geomagnetic observatories that measures the intensity of the globally symmetrical equatorial electrojet (the ring current), which is widely used to identify extreme SW and IMF space weather conditions (Sugiura & Chapman 1960; Loewe & Pr\u00f6lss 1997; Siscoe et al. 2006; Borovsky & Shprits, Yuri 2017). A negative Dst value means that Earth\u2019s magnetic field is weakened due to the IMF erosion, particularly during solar storms. The strongest event observed so far is the Carrington event that occurred in 1859 (Carrington 1859). An unusual large number of sunspots on the solar disk and a wide active region was registered, and an extremely fast ICME was launched from it toward the Earth. Several authors studied the Carrington event and suggested that it was a shock that traveled at about 2000 km s\u22121 (Cliver et al. 1990) that generated the strongest geomagnetic storm with Dst \u2248\u22121700 nT (Tsurutani et al. 2003). This was later revised to Dst \u2248\u2212850 nT by Siscoe et al. (2006). The most recent strongest event, called the Bastille Day event (14\u201316 July 2000), reached a Dst \u2248\u2212300 nT for an SW velocity of 1000 km s\u22121 and an IMF intensity of \u224845 nT (Rastatter et al. 2002). On the other hand, typical ICMEs impacting the Earth show an averaged plasma velocity of 350\u2013500 km s\u22121 and IMF intensities between 9\u201313 nT, leading to geomagnetic storms with Dst  \u221250 nT (Cane & Richardson 2003).","Citation Text":["Kay et al. 2017"],"Citation Start End":[[596,611]]}
{"Identifier":"2021AandA...652A.117B__Rodriguez_et_al._2016_Instance_1","Paragraph":"Ring systems are a ubiquitous feature in planetary systems \u2013 all the gas giants in the Solar System have ring systems around them of varying optical depths (see, e.g., Tiscareno 2013; Charnoz et al. 2018), and ring systems have been detected around minor planets (e.g., Chariklo; Braga-Ribas et al. 2014), so it is reasonable that exoplanets and substellar objects host ring systems as well. Long-period eclipsing binary star systems, where one star is surrounded by an extended dark disk-like structure that periodically eclipses the other component, have already been observed, such as EE Cep (Mikolajewski & Graczyk 1999), \u03f5 Aurigae (Guinan & Dewarf 2002), and TYC 2505-672-1, with a companion period of 69 yr (Lipunov et al. 2016; Rodriguez et al. 2016). A large ring-like structure around a substellar companion was proposed to explain observations from 2007 from the J1407 system (Mamajek et al. 2012). 1SWASP J140747.93-394 542.6 (V1400 Cen; hereafter called \u201cJ1407\u201d) is a young, pre-main-sequence star in the Sco-Cen OB association (Mamajek et al. 2012) with spectral type K5 IV(e) Li and is similar in size and mass to the Sun. In 2007, it displayed a complex symmetric dimming pattern of up to ~ 3 magnitudes during a 56 day eclipse. This has been attributed to the transit of a substellar companion (called \u201cJ1407 b\u201d) with a mass of 60\u2013100 MJup (Rieder & Kenworthy 2016) surrounded by an exoring system consisting of at least 37 rings and extending out to 0.6 au in radius (Kenworthy & Mamajek 2015). For these rings to show detectable transit signatures, they must be significantly misaligned with respect to the orbital plane of J1407 b (Zanazzi & Lai 2017). This potential ring system would be considerably larger than the ring system of Saturn, which is located within the planet\u2019s tidal disruption radius. The proposed rings around J1407 b would even cover a significant fraction of the companion\u2019s Hill sphere and would not be expected to be stable over gigayear timescales. If the candidate ringed companion is in a bound orbit around the star, this orbit must be moderately eccentric in order for no othereclipses to have been detected to date (Kenworthy et al. 2015), raising the possibility that there might be a second as yet undetected companion in the system that causes the implied orbital eccentricity for J1407 b. Radial velocity measurements are overwhelmed by the chromospheric noise of the star and do not place strong constraints on other substellar companions (Kenworthy et al. 2015). The transit of J1407 suggests that its orbital plane has a high inclination to our line of sight \u2013 if there are other planets inside the orbit of J1407 b, their orbits may well be coplanar with J1407 b and there is a high chance that these companions may transit J1407.","Citation Text":["Rodriguez et al. 2016"],"Citation Start End":[[735,756]]}
{"Identifier":"2015MNRAS.452.1112E__Carpano_et_al._2007b_Instance_1","Paragraph":"The X-ray luminosity of CG X-1 is variable by a factor of \u224810 (Bianchi et al. 2002; Weisskopf et al. 2004). The highest flux reported in the literature (5.2 \u00d7 10\u221212 erg cm\u22122 s\u22121 for a power-law fit or 5 \u00d7 10\u221212 erg cm\u22122 s\u22121 for an MCD fit, in the 0.5\u20138 keV band; Weisskopf et al. 2004) would imply, for a distance of 4.2 Mpc, a 0.5\u201310 keV luminosity of LX = (1.5\u20132) \u00d7 1040 erg s\u22121. If the system is Eddington-limited, the lower limit on the mass of the accreting BH is MBH \u2273 75 M\u2299 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 \u2273 1.5\u2009vw,\u200910004\u03b42 M\u2299, where MBH is the BH mass, vw,\u20091000 is the wind velocity in units of 1000\u2009km\u2009s\u22121, and \u03b4 \u2248 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\n\n(1)\n\n\\begin{equation}\nL_{\\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},\n\\end{equation}\n\nwhere \u03b7 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\u2299 yr\u22121 and vw = 1000\u2009km\u2009s\u22121 for the WR star (e.g. Crowther 2007), a = 5.8 \u00d7 1011 cm (for a 10 M\u2299 companion), MBH = 75 M\u2299, and the formation of a disc with \u03b7 = 0.1, the corresponding luminosity is LX \u2243 2 \u00d7 1040 erg s\u22121. More in general, for MBH > 10 M\u2299 and all the other things being equal, one finds LX \u2273 3 \u00d7 1039 erg s\u22121. 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\u2013BH 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 \u223c2 \u00d7 1040 erg s\u22121 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 \u00d7105 yr (Meynet & Maeder 2005).","Citation Text":["Carpano et al. 2007b"],"Citation Start End":[[857,877]]}
{"Identifier":"2019ApJ...883...73C__Engelbrecht_&_Burger_2013a_Instance_2","Paragraph":"Both Nel (2016) and Zhao et al. (2018) report that these parallel mean free paths are smaller during periods of high solar activity and larger during solar minima, with the opposite being true for the resulting perpendicular mean free paths. Zhao et al. (2018) also show that such solar cycle dependences would not be fully accounted for if only the solar cycle changes in the HMF were taken into consideration. These mean free paths also display complicated spatial dependences if turbulence quantities beyond 1 au, yielded by various turbulence transport models (see, e.g., Oughton et al. 2011; Usmanov et al. 2016; Wiengarten et al. 2016; Zank et al. 2017), taking into account observed latitudinal variations of turbulence quantities in the heliosphere (Forsyth et al. 1996; Bavassono et al. 2000a, 2000b; Erd\u00f6s & Balogh 2005), were used as inputs for the turbulence quantities (e.g., Engelbrecht & Burger 2013a, 2015b; Chhiber et al. 2017; Moloto et al. 2018). Furthermore, mean free paths such as those discussed have, when used in conjunction with turbulence-reduced drift coefficients (see, e.g., Minnie et al. 2007b; Engelbrecht et al. 2017) in 3D stochastic modulation codes, led to computed galactic intensities in reasonable agreement with spacecraft observations at Earth (see, e.g., Engelbrecht & Burger 2013a, 2013b; Qin & Shen 2017) and even for several different solar minima (Moloto et al. 2018). The necessity of taking cosmic-ray drift effects into account, combined with the complicated spatial dependences and the fact that none of the diffusion coefficients described above display a P1 rigidity dependence, implies that a convection\u2013diffusion or force-field approach would not be ideal to describe long-term modulation, and that the assumptions implicit to the effective diffusion coefficient used in these formulations are unrealistic. This latter point would call into question any conclusions drawn as to historic solar parameters from quantities such as the modulation potential.","Citation Text":["Engelbrecht & Burger 2013a"],"Citation Start End":[[1297,1323]]}
{"Identifier":"2018MNRAS.478.3890B__Heckman_et_al._2017_Instance_2","Paragraph":"Rather than AGN feedback, it is possible that the effects we are seeing are from a different process coeval or prior to the onset of AGN accretion. Several works have pointed out that AGN activity coincides with a recent starburst, with the AGN having significant accretion events at least \u223c200\u2009Myr after the starburst has occurred (Davies et al. 2007; Wild et al. 2007; Wild, Heckman & Charlot 2010; Yesuf et al. 2014) giving the neutral material time to propagate out to the impact parameters probed by COS-AGN (Heckman et al. 2017). With a sample of QSO sightlines probing the CGM around 17 low-redshift starburst and post-starburst galaxies, Heckman et al. (2017) have observed a similar signature of enhanced EWs of Ly\u2009\u03b1, Si\u2009iii, and C\u2009iv (the latter of which is not measured in our control sample) relative to a control-matched sample (matched in stellar mass and impact parameter). In the range of impact parameters and stellar masses probed by COS-AGN, the strength of our enhanced EW signature is consistent with the values probed by Heckman et al. (2017). However, the results of Heckman et al. (2017) show strong offsets in the kinematics of the gas from the host galaxy (\u2248100\u2009km s\u22121; see fig. 5 from Heckman et al. 2017), whereas the COS-AGN sightlines do not (bottom panel of Fig. 6). Assuming that the AGN activity was triggered by the starburst, a minimum delay time of 200\u2009Myr could allow for any starburst-driven winds to dissipate and kinematic offsets to no longer be present at the impact parameters of the COS-AGN sample. Although this starburst picture provides a possible explanation of our observations, we caution that starbursts are not the only astrophysical event linked to AGN accretion activity. For example, mergers that trigger the AGN (Ellison et al. 2011, 2013; Satyapal et al. 2014; Silverman et al. 2014; Goulding et al. 2018) could potentially affect the surrounding CGM gas. Past and future work focusing on the CGM of galaxy mergers can further test this result (Johnson et al. 2014; Hani et al. 2017; Bordoloi et al. in preparation).","Citation Text":["Heckman et al. (2017)"],"Citation Start End":[[646,667]]}
{"Identifier":"2021AandA...649A.115V__Pravec_et_al._(2012)_Instance_1","Paragraph":"The largest fragment in this cluster, (525) Adelaide, is the only asteroid with some physical characterisation available. Masiero et al. (2014) analysed WISE\/NEOWISE infrared observations and reported Adelaide\u2019s size D = 9.33 \u00b1 0.24 km and geometric albedo pV = 0.22 \u00b1 0.05. These results assume an absolute magnitude H = 12.4 mag in the visible band. More up-to-date absolute magnitude determinations across all standard databases (such as MPC, JPL, or AstDyS) indicate a slightly smaller value of absolute magnitude: H \u2243 12.1. The difference of about 0.3 magnitude is a characteristic scatter in this parameter reported by Pravec et al. (2012), who also noted that near H \u2243 12 mag the systematic offset of the database-reported values should be small (becoming as large as \u2212 0.5 mag for smaller objects). It is therefore possible that the true size of (525) Adelaide is slightly larger, its geometric albedo slightly higher, or a combination of both. The answer will be found when photometrically calibrated data of this asteroid are taken. Luckily, this issue is not critical for our analysis. With a safe margin we can assume Adelaide\u2019s size to be between 9 and 11.5 km, and the geometric albedo between 0.22 and 0.29. This albedo range closely matches the characteristic value in the Flora family region of the inner main belt. Nesvorn\u00fd et al. (2015) identified (525) Adelaide as a member in the Flora family despite the proper inclination value of its orbit being at the high end of the Flora family members. Nevertheless, whether Adelaide is a fragment from the Flora-family formation event, which took part more than 1 Gyr ago (e.g. Vokrouhlick\u00fd et al. 2017b), is again not a crucial issue for our study. The confirmation of this albedo value also comes from broad-band photometry of (525) Adelaide taken by the Sloan Digital Sky Survey (SDSS) project. Considering the methodology in Parker et al. (2008), we used the SDSS-based observations to infer its colour indexes a* = 0.108 and i \u2212 z = \u22120.08. Given the information in Fig. 3 of Parker et al. (2008) we concluded that these values are characteristic of S-complex asteroids. This is a prevalent spectral type in the Flora zone, and the characteristic albedo values of these objects very closely match those given for (525) Adelaide.","Citation Text":["Pravec et al. (2012)"],"Citation Start End":[[625,645]]}
{"Identifier":"2018ApJ...859..122L__Antiochos_et_al._1999_Instance_1","Paragraph":"By comparing our results with the numerical studies of Wyper et al. (2017, 2018), we propose a compound eruption model of circular-ribbon flares consisting of two sets of successively formed ribbons and eruptions of multiple filaments in a fan-spine-type magnetic configuration. A simple schematic picture is shown in Figure 6. First, the null-point reconnection between the outermost flux of the filament F1 (blue curves) and the ambient open field occurs (panel (a)), which transfers partial flux of F1 to the closed field under the far side of the dome and to the open field exterior to the near side of the dome (dashed cyan lines in panel (a)). Meanwhile, the circular ribbon CR1 and inner ribbon IR1 (orange curves) are generated due to the flow of reconnection-accelerated particles from the null along the fan plane and the spine field into the lower atmosphere. The null-point reconnection is likely of the breakout nature (Antiochos et al. 1999; Wyper et al. 2017, 2018). Then, the F1 undergoes the slipping magnetic reconnection along the QSL binding the F1, which induces the far flux of F1 reaching the null point (panel (b)). The continuous null-point reconnection causes the sequential opening of F1, and thus leads to the significant shifts of the brightenings of CR1 and IR1 (orange regions in panel (b)). In the first episode, partial flux is transferred from under the separatrix dome via the null-point reconnection, which effectively alleviates the constraint of the nearby filament F2 (green curves). Later on, the stability of F2 is disrupted and the second episode of the flare is initiated. The continuous null-point reconnection between F2 and open field is at work and generates the sequential brightenings of CR2 and IR2 (panels (c)\u2013(d)), similar to the evolution of the first episode. Meanwhile, the blowout jet is formed at the same time as filament material below the separatrix is ejected along field lines outside the separatrix.","Citation Text":["Antiochos et al. 1999"],"Citation Start End":[[933,954]]}
{"Identifier":"2018MNRAS.475.2787G__Leauthaud_et_al._2012_Instance_1","Paragraph":"In order to separate the role of different physical mechanisms in galaxy evolution, a number of studies have constrained stellar-to-halo mass (SHM) relations and ratios as a function of time using the abundance matching technique (e.g. Behroozi, Conroy & Wechsler 2010a; Moster et al. 2010; Behroozi, Conroy & Wechsler 2010b), the conditional luminosity function technique proposed by Yang, Mo & Van Den Bosch (2003), the halo occupation distribution (HOD) formalism (e.g. Berlind & Weinberg 2002; Kravtsov et al. 2004; Moster et al. 2010), and by combining the HOD, N-body simulations, galaxy clustering, and galaxy\u2013galaxy lensing techniques (e.g. Leauthaud et al. 2012; Coupon et al. 2015). Distinguishing the properties of central galaxies from those of satellite galaxies in studies based only on the distribution of luminosity or stellar mass is challenging (e.g. George et al. 2011). By combining several observables and techniques (e.g. HOD, galaxy\u2013galaxy lensing, and galaxy clustering) one can probe a global SHM relation for central galaxies and satellite galaxies (e.g. Leauthaud et al. 2012; Coupon et al. 2015). Coupon et al. (2015), for example, used multiwavelength data of \u223c60\u2009000 galaxies with spectroscopic redshifts in the Canada France Hawaii Telescope Lensing Survey (CFHTLenS) and Vimos Public Extragalactic Redshift Survey (VIPERS) field to constrain the relationship between central\/satellite mass and halo mass, characterizing the contributions from central and satellite galaxies in the SHM relation. In this paper, we directly identify the BGGs using their precise redshifts and estimate stellar masses using the broad-band spectral energy distribution (SED) fitting technique (Ilbert et al. 2010) as used by Coupon et al. (2015). We utilize the advantages of the X-ray selection of galaxy groups and a wealth of multiwavelength, high signal-to-noise ratio observations such as the UltraVISTA survey in the COSMOS field (Laigle, Capak & Scoville 2016) to investigate the SHM relation for the central galaxies over 9 billion years. We aim to quantify the intrinsic (lognormal) scatter in stellar mass at fixed redshift in observations and compare them to the recently implemented semi-analytic model (SAM) by Henriques et al. (2015).","Citation Text":["Leauthaud et al. 2012"],"Citation Start End":[[649,670]]}
{"Identifier":"2020AandA...635A.186P__Nagy_(2018)_Instance_1","Paragraph":"The choice of \u03ba has some bearing on the value of Mej as calculated in Eq. (2). Although taken as a constant in various model light curves, the actual optical opacity varies with time (e.g. Chugai 2000; Nagy 2018), but tests have shown that that a constant opacity can be a reasonable assumption at early times (see Mazzali et al. 2000; Taddia et al. 2018). SN 2019bkc is a H\/He-deficient SN, so to select an appropriate value for \u03ba, we search the literature literature and consider only SNe of this type. For SNe Ia, where the large abundance of Fe-group elements is a source of significant opacity, Cappellaro et al. (1997) used \u03ba\u2004=\u20040.15 cm2 g\u22121, while Arnett (1982) used \u03ba\u2004=\u20040.08 cm2 g\u22121. The light curve of the broad lined SNe Ic 1997ef was modelled in Mazzali et al. (2000) using the same code as Cappellaro et al. (1997), this time a lower constant opacity of 0.08 cm2 g\u22121 was used owing to the lower Fe abundance in this SN type. The average opacity of model CC-SNe light curves was investigated by Nagy (2018), where it was found that \u03ba\u2004=\u20040.1 cm2 g\u22121 was suitable for their H\/He-deficient models. Taddia et al. (2018) demonstrated that fitting the light curves of H-deficient CC-SNe with a simple analytical model and a constant opacity \u03ba\u2004=\u20040.07 cm2 g\u22121 returned physical parameters consistent with more complex hydrodynamic light curve models. Prentice et al. (2019) also found that using \u03ba\u2004=\u20040.07 cm2 g\u22121 gave comparable ejecta masses to both photospheric phase and nebular phase spectroscopic modelling of H-deficient CC-SNe, with the exception of some gamma-ray burst SNe. From Sect. 5 we find that the outer ejecta is not rich in Fe-group elements, and there is no indication of helium. From the early-nebular spectra in Sect. 4.2 we see no indication of strong Fe\u202fII or Fe\u202fIII line emission, demonstrating that the ejecta is poor in Fe-group elements relative to normal SNe Ia and is more like H\/He-deficient CC-SNe. This justifies the use of \u03ba\u2004=\u20040.07 cm2 g\u22121, but to show the range of possible Mej, we consider the range 0.07\u20130.1 cm2 g\u22121.","Citation Text":["Nagy 2018"],"Citation Start End":[[202,211]]}
{"Identifier":"2022ApJ...936..102A__Williams_et_al._2006_Instance_2","Paragraph":"As regards the modeling of BGK modes, there are two main theoretical approaches: the integral solution or BGK methodology and the differential (or Schamel) technique. In the former method (BGK), one assumes that the initial particle distribution function and the electrostatic potential profiles are known, so these are substituted into the Poisson equation and the integral equation is solved to obtain the trapped particle distribution function (Bernstein et al. 1957; Aravindakshan et al. 2018a, 2018b, and the references therein). In Schamel\u2019s approach, the form of the trapped particle distribution function and of the passing (i.e., free, nontrapped) particle distribution function is assumed and substituted in Poisson\u2019s equation, leading to a differential equation that is then solved to obtain the form of the potential (Schamel 1986; Luque & Schamel 2005, and the references therein). A distinguishing factor in the former (BGK) approach is that it involves a condition in the form of an inequality to be satisfied by the potential parameters (width and amplitude) in order for a BGK mode to be sustained. The BGK approach will be adopted in this work. The above models tacitly assume a collisionless electron-ion plasma. These assumptions are acceptable in the Earth\u2019s magnetosphere. However, as we move farther from near-Earth plasma environments, the presence of charged dust in the plasma cannot be neglected. In the case of Saturn, there are observations of streaming ions by the Cassini spacecraft (Badman et al. 2012a, 2012b). We know that these streaming ion flows can lead to the generation of ion holes. Electrostatic solitary waves have been observed in Saturn\u2019s magnetosphere (Williams et al. 2006) and in the dusty environment near its moon Enceladus (Pickett et al. 2015). Williams et al. (2006) reported observations of solitary structures in the vicinity of Saturn\u2019s magnetosphere. They detected a series of bipolar pulses and speculated that these could be either electron holes or ion holes (Williams et al. 2006). Later on, Pickett et al. (2015) observed solitary wave pulses within 10 Rs (Rs is the Saturn radius) and near Enceladus. Near the Enceladus plume, they discussed how dust impacts affected the observed solitary waves. In fact, Pickett et al. (2015) pointed out that some of the bipolar electric field pulses associated with the solitary waves observed had an inverse polarity (i.e., a positive pulse first, followed by a negative pulse in a short time period) and suggested that this might be due to either an inverse direction of propagation or to a true inverse potential pulse polarity (sign). Moreover, Farrell et al. (2017) examined the conditions that allow low-energy ions, such as those produced in the Enceladus plume, to be attracted and trapped within the sheath of negatively charged dust grains. Using particle-in-cell simulations, they showed that with dust in the system, the large electric field from the grain charge disrupts pickup and leads to ion trapping. Their simulation results also reveal that the bipolar pulses reported in the Enceladus plume by Williams et al. (2006) and Pickett et al. (2015) could most probably be ion holes. In the light of the above information, we may suggest that the formation of ion holes is highly likely in the dusty plasma of environments such as the one found in Saturn. Importantly, the instrument on board Cassini, the Radio Plasma Wave Science instrument, does not have the ability to determine the polarity of the electric fields associated with the ESWs observed with 100% certainty as it lacks the ability to perform interferometry (Williams et al. 2006).","Citation Text":["Williams et al. (2006)","Williams et al. 2006"],"Citation Start End":[[1797,1819],[2020,2040]]}
{"Identifier":"2022AandA...664A.145D__Marconi_et_al._1998_Instance_1","Paragraph":"The stellar initial mass function (IMF) of stars (i.e., the distribution of the masses of stars at their birth) is of fundamental importance in astrophysics. The IMF controls the efficiency of star formation in molecular clouds (Dib et al. 2011, 2013; Hony et al. 2015), the radiative and mechanical feedback from stars into the large-scale interstellar medium (Dib et al. 2006, 2021; Martizzi et al. 2016; Silich & Tenorio-Tagle 2017), and the dynamical and chemical evolution of galaxies (C\u00f4t\u00e9 et al. 2016). Significant efforts have been devoted to the determination of the shape of the IMF in a variety of environments, including the Galactic field (Salpeter 1955; Bochanski et al. 2010; Rybizki & Just 2015; Mor et al. 2019; Sollima 2019), the Galactic bulge (Zoccali et al. 2000; Wegg et al. 2017), and globular clusters (GCs; Da Costa & Freeman 1976; De Marchi & Paresce 1997; Covino & Ortolani 1997; Marconi et al. 1998; Paresce & De Marchi 2000; Sollima et al. 2007; Balbinot et al. 2009; Sollima & Baumgardt 2017; Cadelano et al. 2020), as well as in many young open clusters and associations (Dib 2014; Weisz et al. 2015; Maia et al. 2016; Dib et al. 2017; Jose et al. 2017; Madaan et al. 2020; Bisht et al. 2021; Damian et al. 2021; Elsanhoury et al. 2022). The existence of IMF variations in clusters has significant consequences for the galaxy-wide IMF and for galactic evolution (Dib & Basu 2018; Dib 2022). In the galactic ecosystem, GCs stand out as relics of early star formation with ages that span the range \u224811 Gyr to \u224813.2 Gyr (Salaris et al. 1997; VandenBerg et al. 2013; Pfeffer et al. 2018; Usher et al. 2019; Oliveira et al. 2020), and it is well established that GCs in the Milky Way and in the Magellanic clouds harbor two or more stellar populations (D\u2019Antona & Caloi 2008; Milone et al. 2010; Sbordone et al. 2011; Gratton et al. 2012; Cummings et al. 2014; Piotto et al. 2015; Lee 2015; Oldham & Auger 2016; Mucciarelli et al. 2016; Massari et al. 2016; Dalessandro et al. 2016; Bowman et al. 2017; Carretta & Bragaglia 2018; Latour et al. 2019; Gilligan et al. 2020; Dondoglio et al. 2021; Jang et al. 2021; D\u2019Antona et al. 2022; Kapse et al. 2022).","Citation Text":["Marconi et al. 1998"],"Citation Start End":[[907,926]]}
{"Identifier":"2020MNRAS.497.4423N__Demangeon_et_al._2018_Instance_1","Paragraph":"With a mass of Mb = 37.1 \u00b1 5.6 M\u2295 and a radius of Rb = 7.50 \u00b1 0.44 R\u2295, K2-280\u2009b joins the group of sub-Saturns planets \u2013 defined as planets having radii between 4 and 8 R\u2295 (Petigura et al. 2017) \u2013 whose masses and radii have been measured. The basic physical parameters of a sample of 23 sub-Saturns with densities measured with a precision better than 50 per cent have been presented and discussed by Petigura et al. (2017). We here extend this sample by adding K2-280\u2009b alongside 6 additional sub-Saturns that have densities measured with a precision better than 50 per cent, as described below. WASP-156\u2009b (Demangeon et al. 2018), an \u223c0.5\u2009$\\mathrm{R_\\mathrm{Jup}}$ planet with a Jupiter-like density was discovered by the ground-based SuperWASP transit survey (Pollacco et al. 2006; Smith & WASP Consortium 2014). Kepler-1656\u2009b, a dense sub-Saturn with a high eccentricity of e = 0.84 transiting a relatively bright (V =\u200911.6\u2009mag) solar-type star, was recently reported by Brady et al. (2018). Three sub-Saturns were discovered and characterized by the KESPRINT consortium, two of them in K2 campaign 3 (K2-60\u2009b, Eigm\u00fcller et al. 2017) and campaign 14 (HD\u200989345\u2009b, aka K2-234\u2009b, Van Eylen et al. 2018; Yu et al. 2018b), and HD\u2009219666\u2009b (Esposito et al. 2019) in TESS Sector 1. One sub-Saturn, GJ\u20093470\u2009b (Bonfils et al. 2012), orbiting an M1.5 dwarf was not included by Petigura et al. (2017), but we add it to the current sample, adopting the parameters from Awiphan et al. (2016). All of these new sub-Saturns, including K2-280\u2009b, reside in apparently single systems. Fig. 6 shows the mass\u2013radius and mass\u2013density diagrams for this extended sample of 30 planets. Sub-Saturns found to be in multiplanet systems are marked with green filled circles, whereas those in single systems are marked with blue filled circles. The position of K2-280\u2009b is indicated with a red-rimmed circle. Sub-Saturns whose density has been measured with a precision slightly worse than 50 per cent are marked with green and blue open circles. All the remaining transiting planets with measured radii and masses are marked with open grey circles.12 According to the Fortney, Marley & Barnes (2007)\u2019s models \u2013 also shown in the mass\u2013radius diagram (Fig. 6, upper panel) \u2013 K2-280\u2009b has a core of about 10\u201325 M\u2295, accounting for \u223c25\u201365 per cent of its total mass.","Citation Text":["Demangeon et al. 2018"],"Citation Start End":[[610,631]]}
{"Identifier":"2017ApJ...836L...4S__Scholer_&_Burgess_2007_Instance_1","Paragraph":"At quasi-perpendicular shocks, the average shock structure is dominated by a foot of reflected ions, which is upstream of the shock ramp where the major thermalization and deceleration occurs. Non-stationarity in the form of rippling of the surface or steepened whistler waves (Moullard et al. 2006; Lobzin et al. 2007) is an intrinsic feature of the shock, but this is generally manifest as minor perturbations on top of an otherwise stationary shock ramp. Simulations have predicted that if the fraction of ions reflected by the shock front becomes sufficiently high, the quasi-perpendicular shock can become periodically reforming on timescales of the ion gyroperiod. Various theories have been suggested for such non-stationarity, including self-reformation where a new shock ramp grows at the edge of the foot (Biskamp & Welter 1972a; Lemb\u00e8ge & Dawson 1987), whistler-induced reformation (Biskamp & Welter 1972b; Scholer & Burgess 2007), kinetic instabilities such as the Buneman and modified two-stream instability (e.g., Cargill & Papadopoulos 1988; Matsukiyo & Scholer 2003, 2006b; Scholer et al. 2003; Scholer & Burgess 2007; Matsumoto et al. 2013), and gradient catastrophe of nonlinear whistler waves due to steepening (Krasnoselskikh et al. 2002). However, it has not been until recently that such non-stationarity has been confirmed with in situ spacecraft observations. In a survey of Cassini shock crossings at Saturn, Sulaiman et al. (2015) found evidence of a periodically reforming shock, pulsating at a period near 0.3 of the ion gyroperiod in the unperturbed upstream medium. This period agrees with the time taken for a specularly reflected proton to gyrate across the foot and return to the main shock ramp. Sulaiman et al. (2015) also report that these periodic non-stationary shocks are primarily found in the very high Mach number regime, which gives evidence for a relation between Mach number and reformation. The main processes behind the non-stationary behavior of these very high Mach number shocks, such as the details of the ion- and electron-scale processes acting within the shock transition, remain elusive.","Citation Text":["Scholer & Burgess 2007"],"Citation Start End":[[918,940]]}
{"Identifier":"2019MNRAS.490.2071Y__Riess_et_al._2018_Instance_3","Paragraph":"\nSet II: we now focus on the observational constraints on the model parameters after the inclusion of the local measurement of H0 by Riess et al. (2018) with the previous data sets (CMB, Pantheon, and CC) in order to see how the parameters could be improved with the inclusion of this data point. Since for this present UM, the estimation of H0 from CMB alone is compatible with the local estimation of H0 by Riess et al. (2018), thus, we can safely add both the data sets to see whether we could have something interesting. Following this, we perform another couple of tests after the inclusion of R18. The observational results on the model parameters are summarized in Table 4. However, comparing the observational constraints reported in Table 3 (without R18 data) and Table 4 (with R18), one can see that the inclusion of R18 data (Riess et al. 2018) does not seem to improve the constraints on the model parameters. In fact, the estimation of the Hubble constant H0 remains almost similar to what we found in Table 3. In order to be more elaborate in this issue, we have compared the observational constraints on the model parameters before and after the inclusion of R18 to other data sets. In Figs 7 (CMB versus CMB+R18), 8 (CMB+CC versus CMB+CC+R18), 9 (CMB+Pantheon versus CMB+Pantheon+R18), and 10 (CMB+Pantheon+CC versus CMB+Pantheon+CC+R18), we have shown the comparisons which prove our claim. One can further point out that the strong correlation between the parameters \u03bc and H0 as observed in Fig. 5 still remains after the inclusion of R18 [see specifically the (\u03bc, H0) planes in Figs 7\u201310]. The physical nature of \u03bc does not alter at all. That means the correlation between H0 and \u03bc is still existing after the inclusion of R18 to the previous data sets, such as CMB, Pantheon, and CC. In addition to that since \u03bc \u2272 0.9 according to all the observational data sets, thus, the transition from past decelerating era to current accelerating era occurs to be around z \u2272 0.6, similar to what we have found with previous data sets (Table 3).","Citation Text":["Riess et al. 2018"],"Citation Start End":[[837,854]]}
{"Identifier":"2019MNRAS.490.2071Y__Riess_et_al._2018_Instance_2","Paragraph":"\nSet II: we now focus on the observational constraints on the model parameters after the inclusion of the local measurement of H0 by Riess et al. (2018) with the previous data sets (CMB, Pantheon, and CC) in order to see how the parameters could be improved with the inclusion of this data point. Since for this present UM, the estimation of H0 from CMB alone is compatible with the local estimation of H0 by Riess et al. (2018), thus, we can safely add both the data sets to see whether we could have something interesting. Following this, we perform another couple of tests after the inclusion of R18. The observational results on the model parameters are summarized in Table 4. However, comparing the observational constraints reported in Table 3 (without R18 data) and Table 4 (with R18), one can see that the inclusion of R18 data (Riess et al. 2018) does not seem to improve the constraints on the model parameters. In fact, the estimation of the Hubble constant H0 remains almost similar to what we found in Table 3. In order to be more elaborate in this issue, we have compared the observational constraints on the model parameters before and after the inclusion of R18 to other data sets. In Figs 7 (CMB versus CMB+R18), 8 (CMB+CC versus CMB+CC+R18), 9 (CMB+Pantheon versus CMB+Pantheon+R18), and 10 (CMB+Pantheon+CC versus CMB+Pantheon+CC+R18), we have shown the comparisons which prove our claim. One can further point out that the strong correlation between the parameters \u03bc and H0 as observed in Fig. 5 still remains after the inclusion of R18 [see specifically the (\u03bc, H0) planes in Figs 7\u201310]. The physical nature of \u03bc does not alter at all. That means the correlation between H0 and \u03bc is still existing after the inclusion of R18 to the previous data sets, such as CMB, Pantheon, and CC. In addition to that since \u03bc \u2272 0.9 according to all the observational data sets, thus, the transition from past decelerating era to current accelerating era occurs to be around z \u2272 0.6, similar to what we have found with previous data sets (Table 3).","Citation Text":["Riess et al. (2018)"],"Citation Start End":[[409,428]]}
{"Identifier":"2016MNRAS.463.4121T__Brouillet_et_al._2005_Instance_1","Paragraph":"Molecular line ratio diagrams for NGC 4710, NGC 5866 and a variety of other galaxies. Our data for NGC 4710 and NGC 5866 are shown as filled circles and squares, respectively, while our data for the nuclear discs and inner rings are shown in blue and red, respectively (black for the intermediate region). Upper and lower limits are represented by arrows. Other lenticular galaxies are indicated by magenta filled stars (Krips et al. 2010; Crocker et al. 2012), starburst nuclei by dark green filled stars, Seyferts by brown filled stars, spiral arm GMCs by black circles with an X (see Baan et al. 2008; table 3 in Krips et al. 2010 and references therein), NGC 6946 (starburst) GMCs by black circles with a cross (Topal et al. 2014), and M31 GMCs by turquoise circles with an X (Brouillet et al. 2005). The data for NGC 1266 (a lenticular galaxy with a molecular outflow) are shown by magenta squares with a filled star (Alatalo et al. 2011). The green shaded region in panel a indicates the typical range of 12CO(1\u20130)\/HCN(1\u20130) ratios in starbursts with LFIR > 1011 L\u2299 (see table B2 in Baan et al. 2008). The range of R11 ratios in the nuclear disc and inner ring of NGC 5866 (this work) are indicated by respectively the blue and red horizontal lines in panel b, while the typical range in spirals with LFIR  1011 L\u2299 (Paglione et al. 2001) is indicated by the pale grey shaded region. In panels a and b, Crocker et al.'s (2012) single-dish observations of NGC 4710 and NGC 5866 are shown as an open black circle and an open black square, respectively (see table 4 of Crocker et al. 2012). The HCN(1\u20130)\/HCO+ ratios for M31 GMCs (Brouillet et al. 2005) are indicated by the turquoise shaded region in panels c and d. The green shaded region in panel d indicates the typical range of 13CO(1\u20130)\/HCO+(1\u20130) ratios in the disc of M82 (starburst; Tan et al. 2011). The 12CO(1\u20130)\/HCN(1\u20130) ratios in spirals (Gao & Solomon 2004a) are indicated by the dark grey shaded region in panels a, e and f, respectively. The black solid lines in a number of panels show the 1 : 1 relation and are there to guide the eye. Similarly, the black dashed lines show a ratio of 1 in panels c, d and f.","Citation Text":["Brouillet et al. 2005"],"Citation Start End":[[781,802]]}
{"Identifier":"2019ApJ...873..111I__Scovacricchi_et_al._2017_Instance_1","Paragraph":"The two LSST observing programs are complementary in the SN samples they will provide. The main survey will obtain light curves in six bands and photometric redshifts of about 400,000 photometrically classified SNe Ia that can be used for cosmological distance measurements, with further spectroscopic follow-up of a subsample of their host galaxies. Such a sample will not only provide larger statistics for the study of the Type Ia population in the universe but also be spread across the full 18,000 deg2 LSST main survey footprint, allowing different probes of the large-scale structure of the low-redshift universe. This sample of SNe can be used as a tracer of large-scale structure by directly probing the gravitational potential of structure through inferences of their peculiar velocities (Gordon et al. 2007; Bhattacharya et al. 2011; Howlett et al. 2017), WL of SN brightnesses (Dodelson & Vallinotto 2006; Quartin et al. 2014; Macaulay et al. 2017; Scovacricchi et al. 2017), and the local bulk flow (Riess 2000; Dai et al. 2011; Turnbull et al. 2012; Feindt et al. 2013; Huterer et al. 2015), as well as low-redshift constraints on the isotropy of the universe (Antoniou & Perivolaropoulos 2010; Campanelli et al. 2011; Colin et al. 2011; Cai et al. 2013; Javanmardi et al. 2015). The rapidly sampled Deep Drilling Fields, possibly co-added over short timescales, will yield well-sampled light curves of tens of thousands of SNe to redshifts peaking around z \u223c 0.7 and reaching beyond a redshift of 1.0, limited by the systematics related to the limits of our astrophysical understanding of SN populations and relative photometric calibration. In addition to the usual use of SNe Ia to probe the redshift\u2013distance relation to high redshift, the luminosities will be magnified by lensing from foreground structure, a correlation that can be probed with these data. The ultimate promise of such SN surveys will be linked to the observing strategy employed by the LSST.","Citation Text":["Scovacricchi et al. 2017"],"Citation Start End":[[961,985]]}
{"Identifier":"2016ApJ...822...15S__Bruntt_et_al._2010_Instance_1","Paragraph":"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, \n\n\n\n\n\n\n\u03bd\n\n\nmax\n\n\n\u221d\n\n\ngT\n\n\neff\n\n\n\u2212\n1\n\n\/\n\n2\n\n\n\n\n and \n\n\n\n\n\u0394\n\u03bd\n\u221d\n\u03c1\n\n\n (Brown et al. 1991; Kjeldsen & Bedding 1995), to estimate their radius and mass (and hence age). Here, \n\n\n\n\n\n\n\u03bd\n\n\nmax\n\n\n\n\n is the frequency of maximum amplitude and \n\n\n\n\n\u0394\n\u03bd\n\n\n 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 \n\n\n\n\n\n\n\u03bd\n\n\nmax\n\n\n\n\n scaling relation has been shown to work well (Bedding 2014) and recently, Coelho et al. (2015) found the proportionality \n\n\n\n\n\n\n\u03bd\n\n\nmax\n\n\n\u221d\n\n\ngT\n\n\neff\n\n\n\u2212\n1\n\n\/\n\n2\n\n\n\n\n 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 \n\n\n\n\n\u0394\n\u03bd\n\n\n (but with additional \n\n\n\n\n\n\n\u03bd\n\n\nmax\n\n\n\n\n independent quantities) is lower compared to using both \n\n\n\n\n\u0394\n\u03bd\n\n\n and \n\n\n\n\n\n\n\u03bd\n\n\nmax\n\n\n\n\n. Based on this they suggested that a modification to the \n\n\n\n\n\n\n\u03bd\n\n\nmax\n\n\n\n\n scaling relation might be required. Theoretical modeling has suggested corrections to the \n\n\n\n\n\u0394\n\u03bd\n\n\n 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 \n\n\n\n\n\n\n\u03bd\n\n\nmax\n\n\n\n\n, Houdek et al. (1999) and Chaplin et al. (2008) had suggested theoretically that \n\n\n\n\n\n\n\u03bd\n\n\nmax\n\n\n\n\n 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 \n\n\n\n\n\n\n\u03bd\n\n\nmax\n\n\n\n\n from theory.","Citation Text":["Bruntt et al. 2010"],"Citation Start End":[[1248,1266]]}
{"Identifier":"2015AandA...579A..48W__Allard_et_al._2003_Instance_1","Paragraph":"We wish to add to the work of Zhou et al. (2014) by deriving \u1e40acc for LkCa15 b. To this end we use the upper limit on the flux from the H\u03b1, Pa\u03b3 and Pa\u03b2 lines to estimate the accretion luminosity (Lacc) in each line. The equation \u1e40acc = 1.25 (LaccR\u2217)\/(GM\u2217) is used (Gullbring et al. 1998; Hartmann et al. 1998) and the results of Alcal\u00e1 et al. (2014) are used to convert the luminosity of each line (Lline) to Lacc. Rpl is taken as 0.22 R\u2299 and 0.34 R\u2299 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 \u1e40acc in low-mass stars and BDs (Alcal\u00e1 et al. 2014; Whelan et al. 2014a), and therefore it is assumed here that the correlation between Lline and \u1e40acc for low-mass pre-main sequence stars is applicable for planetary mass objects. Studies of \u1e40acc using Lacc estimated from Lline have shown that the optimum way to estimate \u1e40acc is to use several accretion tracers found in different wavelength regimes (Rigliaco et al. 2011, 2012). In this way any spread in \u1e40acc 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\u200915 show strong variability (refer to Fig. 2) we chose to limit our calculations to the lines with the least variability namely H\u03b1, Pa\u03b3 and Pa\u03b2. Here we are specifically referring to the presence of variable absorption features. For example in E1 the He I 1.083 \u03bcm line is double peaked but in E2 and E3 it is single peaked with a deep red-shifted absorption. The H\u03b3 and H\u03b2 lines also show red-shifted absorption in E2. Additionally, H\u03b1 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 3\u03c3 upper limit on the flux, log\u2009(\u1e40acc [ H\u03b1 ] ) = \u22129.0 to \u22129.2, log\u2009(\u1e40acc [ Pa\u03b3 ] ) = \u22129.1 to \u22129.3 and log\u2009(\u1e40acc [ Pa\u03b2 ] ) = \u22128.9 to \u22129.1. The range in \u1e40acc 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 \u1e40acc for a 1\u221210 MJup planet in the range 10-9 to 10-8+M\u2299\/yr (Zhu 2015). ","Citation Text":["Allard et al. 2003"],"Citation Start End":[[564,582]]}
{"Identifier":"2022MNRAS.510.1528V__Takeda,_Hashimoto_&_Honda_2017_Instance_1","Paragraph":"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\u2019Orazi 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 Text":["Takeda, Hashimoto & Honda 2017"],"Citation Start End":[[988,1018]]}
{"Identifier":"2016AandA...595A.106C__Korista_&_Goad_2000_Instance_2","Paragraph":"The present HST-COS data were taken 20 days after the last XMM-Newton pointing (Kaastra et al. 2011) as the closing measurements of the campaign, which lasted in total about 100 days. Spectral coverage simultaneous to HST-COS was provided instead by both Chandra-LETGS (Ebrero et al. 2011) and Swift-XRT (Mehdipour et al. 2011). We used the average SED recorded by the XMM-Newton instruments 20\u201360 days before the HST-COS observation. The choice of SED is very important in the BLR modeling, as different lines respond to the continuum variations on different time scales (Korista & Goad 2000; Peterson et al. 2004). Reverberation mapping studies of Mrk 509 report that the delay of the H\u03b2 with respect to the continuum is very long (about 80 days for H\u03b2, Carone et al. 1996; Peterson et al. 2004). However, higher ionization lines respond more quickly to the continuum variations. Taking as a reference the average H\u03b2\/C\u2009iv delay ratio for NGC 5548 (Peterson et al. 2004), for which \u2013 contrary to Mrk 509 \u2013 a large set of line measurements is available, we obtain that the C\u2009iv line in Mrk 509 should respond in approximately 40 days. A similar (but shorter) time delay should apply to the Ly\u03b1 line (Korista & Goad 2000). This delay falls in the time interval covered by the XMM-Newton data. Therefore, our choice of SED should be appropriate for the modeling of at least the main UV lines. Variability of the X-ray broad lines has been reported on time scales of years (Costantini et al. 2010); however, no short-term studies are available. We expect that the X-ray broad lines should respond promptly to the continuum variations, as they may be located up to three times closer to the black hole with respect to the UV lines (C07). During the XMM-Newton campaign the flux changed by 30% at most, with a minimal change in spectral shape (Sect. 3.1). The used SED should therefore represent what the BLR gas sees for the X-ray band. However, for the optical lines the used SED might be too luminous as we observed an increase in luminosity of about 30% during the XMM-Newton campaign and, as seen above, the time-delay of the optical lines may be large. ","Citation Text":["Korista & Goad 2000"],"Citation Start End":[[1200,1219]]}
{"Identifier":"2022AandA...659L...1L__Lellouch_et_al._2013_Instance_1","Paragraph":"Without knowledge of nucleus shape and spin parameters (pole orientation and shape), a thermophysical model is pointless, and we instead adopted a NEATM (Near Earth Asteroid Thermal Model) model, used extensively for asteroids (Harris 1998) and TNOs (M\u00fcller et al. 2020, and references therein). NEATM is based on the asteroid standard thermal model (STM; Lebofsky et al. 1989) but accounts for phase angle effects; additionally, the temperature distribution is modified by an adjustable \u03b7\u22121\/4 factor, which represents the combined and opposed effects of roughness (\u03b7  1) and thermal inertia (\u03b7 > 1). For fixed surface (thermal inertia, roughness) and spin properties, \u03b7 is also a function of the subsolar temperature, and, therefore, of the heliocentric distance (e.g., Spencer et al. 1989; Lellouch et al. 2013). Given the rh = 20 au distance of our measurements (and the expected large size of 2014 UN271), we adopted a beaming factor \u03b7 = 1.175 \u00b1 0.42, based on measurements of 85 Centaurs and TNOs (Lellouch et al. 2013, 2017). We also specified a bolometric emissivity \u03f5b = 0.90 \u00b1 0.06 and a relative radio emissivity \u03f5r = \u03f5mm\/\u03f5b = 0.70 \u00b1 0.13, as inferred from combined Spitzer\/Herschel\/ALMA measurements of nine objects (Brown & Butler 2017; Lellouch et al. 2017). The lower-than-unity relative radio emissivity is interpreted as resulting from (i) the sounding of a colder dayside subsurface and (ii) the loss of outgoing thermal radiation due to volume scattering in the subsurface and\/or Fresnel reflection at the surface. The few available radio observations of cometary nuclei also generally indicate radio emissivities lower than 1, for example \u223c0.5 for Hale-Bopp (Fern\u00e1ndez 2002) and  0.8 for 8P\/Tuttle (Boissier et al. 2011). Comets are also found to have low thermal inertias (e.g.,  10,  30, and  45 MKS for 8P\/Tuttle, 22P\/Kopff, and 9P\/Tempel 1, respectively; Boissier et al. 2011; Groussin et al. 2009, 2013), consistent with a beaming factor, \u03b7, of order unity. Based on NEATM analysis of a large sample of comet nuclei observed with Spitzer at rh = 3.5\u20136 au, Fern\u00e1ndez et al. (2013) find a mean \u03b7 of 1.03 \u00b1 0.11. The large 29P\/Schwassmann-Wachmann nucleus (D\u2004=\u200465 km) has \u03b7 = 1.1 \u00b1 0.2 (Schambeau et al. 2021). These numbers are fully consistent with our choice of \u03b7. Given the values of rh, \u03b7, and \u03f5b, NEATM calculations indicate that the object\u2019s spectral index over 224\u2013242 GHz is 1.93, slightly lower than the Rayleigh-Jeans limit of 2.","Citation Text":["Lellouch et al. 2013"],"Citation Start End":[[792,812]]}
{"Identifier":"2018MNRAS.475.3419A__Davis_et_al._1999_Instance_2","Paragraph":"If we consider for the bulk density the value 4500 kg m\u22123, which is one of the highest measured in the asteroid population out of those asteroids with good quality of data (see Carry 2012), it will strengthen the hypothesis that Psyche could be an exposed metal core of a differentiated asteroid (Elkins-Tanton et al. 2017). According to the models of asteroid differentiation, the process that led to the formation of Psyche happened very early. Considering Psyche's current diameter, Deff = 226\u2009km (Shepard et al. 2017), the Psyche parent body (PPB) was supposed to be \u223c500\u2009km in diameter and have suffered severe \u2018hit-and-run\u2019 impact events capable of removing all crust and mantle, exposing the core (Elkins-Tanton et al. 2016). In addition, Psyche should have \u223c40\u2009per cent macroporosity, if we assume that it is made of blocks of iron\/nickel with a density around 7500\u2009kg\u2009m\u22123. In that case, the core itself was possibly destroyed and re-accumulated, implying a severe collisional history. When an asteroid is disrupted catastrophically, with a remaining mass \u226450\u2009per cent of the initial one, after a collision with another body, an asteroid family is formed. If the collision happened in the Main Belt, a family of asteroid fragments should be in the region of Psyche; however, no family related to Psyche has been found yet (Davis, Farinella & Marzari 1999). One possibility to solve this issue is that the potential Psyche asteroid family was created at an early time, e.g. within the first 500\u2009Myr of Solar system history (Davis et al. 1999). This would allow the family fragments to be ground down by collisional evolution and be unobservable today. The same models show that, even in this case, today there should be several surviving fragments having diameters around 20\u2009km and above the detection limit. There is a lack of primordial asteroid families in the Main Belt (Bro\u017e et al. 2013; Spoto, Milani & Kne\u017eevi\u0107 2015), very likely due to the classical methods that are used to identify them. The hierarchical clustering method (HCM) is not sensitive enough to find old and dispersed families, as it searches for asteroids forming compact groups in orbital element space (semi-major axis, eccentricity and inclination). A new approach has been proposed and implemented with success (Walsh et al. 2013; Delbo\u2019 et al. 2017), as it is able to distinguish very old families, having eccentricities and inclinations dispersed in space. Therefore the possibility of the absence of a Psyche family could be due to searching biases. However, this may be an unlikely hypothesis, because A-type asteroids that could represent mantle material (almost pure olivine) from differentiated bodies do not exist extensively in the orbital space related to Psyche, but instead are distributed randomly in the Main Belt (Davis et al. 1999; DeMeo et al. 2015). In order to study this puzzling small body further, NASA is sending a new Discovery Mission to Psyche. The main goal is to get insight into whether it is a core of a parent body and understand the procedures of differentiation, making all the above questions more valid than ever. The alternative theory is that Psyche is a planetesimal that bears primitive unmelted material (Elkins-Tanton et al. 2016).","Citation Text":["Davis et al. 1999"],"Citation Start End":[[2812,2829]]}
{"Identifier":"2021MNRAS.503.5367B__Burgay_et_al._2014_Instance_1","Paragraph":"Magnetars are possibly the most exotic objects of all the classes of neutron stars. High magnetic fields (\u22731014 G) and intense X-ray\/soft \u03b3-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\u2009J1935+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 $\\mathrm{\\mu Jy}$ (3\u03c3; Fong & Berger 2014). Subsequent observations with the Chandra X-ray telescope detected X-ray pulsations with a period of 3.24\u2009s (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 3\u03c3 upper limits of 1.2 and 4.5\u2009mJy at frequencies of 362.5 and 610\u2009MHz, 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 Text":["Burgay et al. 2014"],"Citation Start End":[[1133,1151]]}
{"Identifier":"2016ApJ...819...66D__Burns_et_al._2015_Instance_1","Paragraph":"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; Chavarr\u00eda 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\u22120) line data, Kirsanova et al. (2008) found three molecular gas components (i.e., \u221218 km s\u22121\n\n\n\n\n\n km s\u22121 (red), \u221221 km s\u22121 \n\n\n\n\n\n km s\u22121 (central), and \u221225 km s\u22121 \n\n\n\n\n\n km s\u22121 (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 Text":["Burns et al. 2015"],"Citation Start End":[[633,650]]}
{"Identifier":"2019MNRAS.489.4669S__Bigiel_et_al._2010_Instance_1","Paragraph":"In Fig. 7 we compare UGC 1378\u2019s SFR density versus gas surface density (the Schmidt\u2013Kennicutt relation) to data in the literature. The gas surface density corresponds to H\u2009i 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\u2009i 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) \u2013 total and H\u2009isurface densities. A blue line shows the best-fitting relation for the H\u2009isurface 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\u2009i 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\u2013Kennicutt relation. Similar deviations are observed in \u2018classical\u2019 LSB galaxies, Bluedisk galaxies (Roychowdhury et al. 2015), outer parts of HSB spiral galaxies (Bigiel et al. 2010), and H\u2009idiscs in early-type galaxies (Y\u0131ld\u0131z 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\u2009i rich galaxies and proposed that very high specific angular momentum in H\u2009irich 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 Text":["Bigiel et al. 2010","Bigiel et al. 2010"],"Citation Start End":[[916,934],[1862,1880]]}
{"Identifier":"2022MNRAS.516.1406S__Mapelli_2016_Instance_1","Paragraph":"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 & M\u00fcller 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 Text":["Mapelli 2016"],"Citation Start End":[[906,918]]}
{"Identifier":"2022MNRAS.512.4185A__Kitzmann_&_Heng_2018_Instance_1","Paragraph":"Our free chemistry retrieval suite results are further at odds with equilibrium chemistry due to the potential presence of an overabundance of H\u2212 in many retrievals across our suite. Across our free chemistry retrieval suite, the increasing abundance of H\u2212 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\u2212 acts like a uniform opacity source in the optical, which turns off at around 1.6\u2009$\\rm{\\mu m}$\u2009coinciding with these data points. Under equilibrium chemistry conditions, the abundances of H\u2212 retrieved in our suite would be possible only at temperatures \u22732500\u2009K\u2009(Kitzmann & Heng 2018). However, the production of appreciable quantities of H\u2212 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\u2212 by H2 dissociative electron attachment and destruction by atomic H collisional detachment can enable mixing ratios of H\u2212, 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\u2212 seen in our retrieval suite plausible outside of equilibrium conditions. While the downturn at 1.6\u2009$\\rm{\\mu m}$, which seems to drive the inclusion of H\u2212, 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\u2009$\\rm{\\mu m}$,\u2009which 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\u2212 is indeed present within the atmosphere of WASP-17b.","Citation Text":["Kitzmann & Heng 2018"],"Citation Start End":[[657,677]]}
{"Identifier":"2016AandA...587A.133G__Roy_et_al._2015_Instance_1","Paragraph":"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) \u2264 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\u2299 (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, \u201crunaway\u201d 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\u22126 (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 \n        20\u221230% (e.g. Mitra et al. 2013) or even higher (e.g. Wise & Cen 2009). ","Citation Text":["Roy et al. 2015"],"Citation Start End":[[1008,1023]]}
{"Identifier":"2015ApJ...806..168Z__Spitkovsky_2008_Instance_1","Paragraph":"It is not clear which mechanism is operating to accelerate electrons up to relativistic energies. It could be shock acceleration, stochastic acceleration or magnetic reconnection (see Zdziarski et al. 2014, for some discussions), or shock interaction in a magnetic reconnection site (Lazarian & Vishniac 1999; de Gouveia dal Pino & Lazarian 2005). The diffusive shock acceleration involving the first-order Fermi process is often regarded as the most effective scenario in jets. However, this process requires that electrons are pre-heated up to a value comparable to the energy of thermal ions. Fortunately, a quasi-Maxwellian electron distribution with strong coupling with protons has been found in PIC simulations of collisionless relativistic (Spitkovsky 2008; Sironi & Spitkovsky 2011) and non-relativistic (Riquelme & Spitkovsky 2011) shocks. For the current study, jets are relativistic or should be at least mildly relativistic. In fact, the low-energy tail of electrons with a thermal distribution has a relatively small contribution to the SEDs. On one hand, a sharp low-energy cutoff, \n\n\n\n\n\n, from which electrons can be energized, is assumed in Zdziarski et al. (2014) for the jets of black hole X-ray binaries in the low\/hard state. On the other hand, a hard injection spectrum below \n\n\n\n\n\n, which could approximate the Maxwellian distribution, is used to model a sample of blazars (e.g., Ghisellini & Tavecchio 2009; Zhang et al. 2012). Similarly to the latter, the relativistic electrons injected in the dissipation region are considered as\n4\n\n\n\n\n\nwhere \n\n\n\n\n\n is the break energy of relativistic electrons, and p and q are the energy spectral indices of the electrons below \n\n\n\n\n\n and above \n\n\n\n\n\n, respectively. The normalization constant of the electrons, \n\n\n\n\n\n, is determined by\n5\n\n\n\n\n\nwhere \n\n\n\n\n\n is the accretion power of the system via Roche lobe outflows of the companion star. Here, \n\n\n\n\n\n is the mass accretion rate. The power of relativistic electrons, \n\n\n\n\n\n, accounts for a fraction of the jet power of \n\n\n\n\n\n. We further consider \n\n\n\n\n\n to be 0.1 and \n\n\n\n\n\n an adjustable parameter. \n\n\n\n\n\n (assuming to be 1) and \n\n\n\n\n\n are minimum and maximum energies of the relativistic electron, respectively.","Citation Text":["Spitkovsky 2008"],"Citation Start End":[[749,764]]}
{"Identifier":"2018MNRAS.477.2326F__P\u00e9rez_et_al._2016_Instance_1","Paragraph":"In this work, we use APOGEE calibrated atmospheric parameters and chemical abundances from the Sloan Digital Sky Survey IV (SDSS-IV; Gunn et al. 2006; Blanton et al. 2017) Data Release 14 (DR14; Abolfathi et al. 2017). The APOGEE target selection is described in Zasowski et al. (2013) and data reduction is described in Nidever et al. (2015). The atmospheric parameters and chemical abundances are determined using the APOGEE Stellar Parameters and Chemical Abundances Pipeline (ASPCAP; Garc\u00eda P\u00e9rez et al. 2016). ASPCAP uses the code ferre (Allende Prieto et al. 2006) in combination with grids of normalized stellar synthetic spectra (see Zamora et al. 2015), as created with extensive atomic and molecular linelists (Shetrone et al. 2015), to simultaneously determine Teff, log\u2009g, [M\/H], [\u03b1\/M], [C\/M], [N\/M], and microturbulance through \u03c72 minimization. During this process, ASPCAP fits for the overall abundance of metals [M\/H] and the relative abundance of \u03b1-elements [\u03b1\/M], using O, Mg, Si, S, Ca, and Ti; the relative abundances of non-\u03b1 elements (except for C and N) are set to solar ratios when inferring [M\/H]. The measurements of both [M\/H] and [\u03b1\/M] are dominated by different elements in different temperature regimes. The abundances of the individual elements are then determined by fitting specific spectral windows containing relevant absorption lines while varying the corresponding family of elements (Holtzman et al. 2015, in preparation). The DR14 abundances have been internally calibrated to remove Teff\u2009correlations in open clusters, similar to the corrections to DR12 described in Holtzman et al. (2015); see Holtzman (in preparation) for a description of the DR14 calibrations. The APOGEE team has found that the metallicity parameter [M\/H] closely correlates with the [Fe\/H] measurement (Holtzman et al. 2015, in preparation). In this paper, we will assume that the ASPCAP measurements of [M\/H] and [\u03b1\/M] can be taken as proxies from [Fe\/H] and [\u03b1\/Fe], and therefore compared to literature results and theoretical models for these ratios.","Citation Text":["Garc\u00eda P\u00e9rez et al. 2016"],"Citation Start End":[[488,512]]}
{"Identifier":"2018AandA...611A..74R__Grady_et_al._2013_Instance_4","Paragraph":"In this context, MWC 758 (HD 36112) offers a unique environment to probe the existence of planetary companions and to explore the connection between disk structures and planet formation. MWC 758 is a young stellar object (3.5 \u00b1 2 Myr, Meeus et al. 2012) at a distance of 151\n$^{+9}_{-8}$\n\n\n\n\n151\n\n\u22129\n\n+8\n\n\n\n4\n\n\n\n\n pc (Gaia Collaboration 2016) close to the edge of the Taurus star forming region (stellar properties are given in Table 1). Measurements of resolved CO emission around the star determined the stellar mass to be 2.0 \u00b1 0.2 M\u2299 and the disk to have an inclination of 21\u00b0 \u00b1 2\u00b0 and a position angle of the semi-major axis of 65\u00b0 \u00b1 7\u00b0 (Isella et al. 2010). The mass and age estimates were based on the previously adopted hipparcos distances of 200 pc (van den Ancker et al. 1998) and 279 pc (van Leeuwen 2007). Given the revised Gaia distance, the star could be older and lighter than previously thought. In this paper, we assume a stellar mass of 1.5 \u00b10.2 M\u2299, reflecting the scaling of the dynamical mass estimate to the new Gaia distance. Based on its SED, MWC 758 has been classified as a pre-transition disk (Grady et al. 2013). Although a cavity of 55 astronomical units (au) in radius has been inferred from dust millimeter emission (Andrews et al. 2011), infrared polarized intensity observations have found no clear evidence for a cavity in scattered light (Grady et al. 2013; Benisty et al. 2015). Using Ks-band (2.15 \u03bcm) direct imaging andH-band (1.65 \u03bcm) polarimetric imaging with the High-Contrast Instrument with Adaptive Optics (HiCIAO) at the Subaru Telescope, Grady et al. (2013) detected two spiral arms and polarized light down to 0.\u2032\u2032 1 (15 au) from the star. Recent VLT Spectro-Polarimetric High-contrast Exoplanet REsearch (SPHERE) observations in the Y band (1.04 \u03bcm) have confirmed the presence of scattered light at least down to 14 au (Benisty et al. 2015). The asymmetries observed by Isella et al. (2010) in the mm-dust distribution and in CO emission suggest that the disk may be gravitationally perturbed by a low-mass companion orbiting within a radius of 23 au (assuming a distance of 151 pc). The asymmetric cm-dust distribution was shown to follow the location of the mm-dust (Marino et al. 2015a), hinting towards the hypothesis of a dust trap, which could also be created by a companion in the gap through the Rossby wave instability (e.g., Pinilla et al. 2012b). Hydrodynamical simulations of the disk indicate that the observed spirals could instead be launched by a massive planet or brown dwarf at larger separations (~ 100 au based on the revised Gaia distance, Dong et al. 2015b). The presence of stellar companions down to a mass limit of 12 MJup at 0.\u2032\u2032 25 and of planets outside 0.\u2032\u20325 (5 MJup at 0.\u2032\u2032 5, and 3 MJup at 1\u2032\u2032 , according to the BT-SETTL models; Allard et al. 2012) has been ruled out based on a combination of sparse aperture masking observations at L\u2032 band and angular differential imaging at K\u2032 and Ks bands (Grady et al. 2013).","Citation Text":["Grady et al. 2013"],"Citation Start End":[[2975,2992]]}
{"Identifier":"2018MNRAS.480..927P__Goerdt_et_al._2006_Instance_1","Paragraph":"The core\/cusp problem is a clear example of this controversy: on the one hand, cosmological dark matter only N-body simulations predict cuspy dark halo density profiles; on the other hand, the rotation curves of low surface brightness disc and gas-rich dwarf galaxies favour shallower or cored dark matter density distributions (de Blok 2010 and references therein). Also for dSphs, for which the determination of the dark matter density distribution is more difficult, there are indications that cored dark matter density profiles may be favoured with respect to cuspy profiles (Kleyna et al. 2003; Goerdt et al. 2006, Battaglia et al. 2008; Walker & Pe\u00f1arrubia 2011; Salucci et al. 2012; Amorisco, Agnello & Evans 2013; Zhu et al. 2016), though this finding is still debated (Richardson & Fairbairn 2014; Strigari, Frenk & White 2017). It must be stressed, however, that cored dark haloes in dSphs do not necessarily imply a failure of \u039bCDM: dark matter only cosmological simulations may not reliably predict the present-day dark matter distribution in dSphs because, by definition, they neglect the effects of baryons on the dark haloes. Even in a galaxy that is everywhere dark matter dominated today, baryons must have been locally dominant in the past to permit star formation. Therefore, the effect of baryon physics on the dark halo is expected to be important also in dSphs. For instance, Nipoti & Binney (2015) showed how, due to the fragmentation of a disc in cuspy dark halo, dynamical friction may cause the halo to flatten the original cusp into a core even before the formation of the first stars (see also El-Zant, Shlosman & Hoffman 2001; Mo & Mao 2004; Goerdt et al. 2010; Cole, Dehnen & Wilkinson 2011; Arca-Sedda & Capuzzo-Dolcetta 2017). Moreover, the results of hydrodynamical simulations suggest that, following star formation, supernova feedback can also help to flatten the central dark matter distribution, by expelling the gas (Navarro, Eke & Frenk 1996a; Read & Gilmore 2005) and thus inducing rapid fluctuations in the gravitational potential (Mashchenko, Couchman & Wadsley 2006, Pontzen & Governato 2012, Tollet et al. 2016).","Citation Text":["Goerdt et al. 2006"],"Citation Start End":[[600,618]]}
{"Identifier":"2021ApJ...914L..19Z__McKernan_et_al._2012_Instance_2","Paragraph":"Massive stars are believed to exist in the accretion disks of active galactic nuclei (AGNs). Such AGN stars and compact objects can be either the result of in situ formation inside the accretion disk or be captured from the nuclear star clusters around the AGNs (e.g., Artymowicz et al. 1993; Collin & Zahn 1999; Goodman 2003; Goodman & Tan 2004; Wang et al. 2011, 2012; Fabj et al. 2020; Cantiello et al. 2021). These AGN stars will end up with supernovae (SNe), which can eject heavy elements into the disk; this offers a possible explanation for the observational features of high-metallicity environments in AGN disks (e.g., Artymowicz et al. 1993; Hamann & Ferland 1999; Warner et al. 2003). Some compact objects, including white dwarfs (WDs), neutron stars (NSs), and black holes (BHs), can be thus formed within AGN disks. These compact objects can also be captured from the surrounding nuclear star clusters. The disk of an AGN provides a natural environment for stars and compact objects to accrete materials and to migrate within it (e.g., McKernan et al. 2012; Yang et al. 2020; Dittmann et al. 2021; Jermyn et al. 2021; Wang et al. 2021; Tagawa et al. 2021; Kimura et al. 2021). Some of these stars can be very massive and have high spin caused by accretion (Dittmann et al. 2021; Jermyn et al. 2021), so that they can easily produce high-spin stellar remnants. Abundant compact objects, especially with the presence of the massive BHs (M > 10M\u2299), would likely accrete, collide, and merge within the trapping orbits, and hence would grow into \u223c100 M\u2299 intermediate-mass BHs (McKernan et al. 2012; Secunda et al. 2019; Yang et al. 2019b). Some AGN stars can be tidally disrupted by these BHs that can power micro-tidal disruption events (Yang et al. 2021). The death of high-spin stars and neutron star mergers are expected to power gamma-ray burst (GRB) jets, which would be always choked by the dense atmosphere of the disks (Zhu et al. 2021a, 2021b; Perna et al. 2021a). Zhu et al. (2021a) suggested that these choked jets can produce high-energy neutrinos that may contribute a substantial fraction of the diffuse neutrino background. A candidate electromagnetic (EM) counterpart that emerged from an AGN, explained as ram pressure stripping of gas within the kicked BH hill sphere colliding with the AGN disk gas (McKernan et al. 2019), was reported by the Zwicky Transient Facility (Graham et al. 2020). This was thought to be associated with a (85 + 66)M\u2299 binary BH merger (GW190521) detected by the LIGO\/Virgo collaboration (Abbott et al. 2020). This connection provided plausible evidence of a potentially important AGN channel for compact star mergers.","Citation Text":["McKernan et al. 2012"],"Citation Start End":[[1586,1606]]}
{"Identifier":"2021MNRAS.501.3540D__Just_et_al._2012_Instance_1","Paragraph":"Many near-equal mass MBH mergers in the LISA band may occur in gaseous environments, given that the galactic mergers that lead to the eventual formation of an MBH binary often carry a fresh supply of gas into the post-merger galactic nucleus (see Mayer 2013 for a review). Likewise, \u2018gas-embedded\u2019 E\/IMRIs may frequently occur in the accretion discs expected in active galactic nuclei (AGNs) in which an MBH is surrounded by a thin, dense accretion disc. Compact objects in the nucleus can eventually align their orbits with the disc, provided a sufficient number of disc-crossing intersections (e.g. Syer, Clarke & Rees 1991; Artymowicz et al. 1993; Artymowicz, Lin & Wampler 1993; Rauch 1995; Just et al. 2012; McKernan et al. 2012; Kennedy et al. 2016; Panamarev et al. 2018; Fabj et al. 2020; MacLeod & Lin 2020), or such events may arise from in-situ star formation in the disc that leaves compact remnants (Goodman & Tan 2004; Levin 2007). Embedded stars and BHs can subsequently accrete, migrate, and merge with each other (Bellovary et al. 2016; Secunda et al. 2019, 2020; Tagawa, Haiman & Kocsis 2020a), before eventually merging with the central MBH. While the event rate for E\/IMRIs has previously been estimated in dry nuclei (e.g. Barausse, Cardoso & Pani 2015), the rate may be considerably higher when including formation pathways in accretion discs. The recent discovery of stellar-mass BH mergers by the ground based interferometer LIGO has stimulated work on such mergers that may be occurring in AGN discs, showing that they might contribute to LIGO events (McKernan et al. 2014; Bartos et al. 2017; Stone, Metzger & Haiman 2017; McKernan et al. 2018; Gr\u00f6bner et al. 2020). Subsequent work has suggested that some of the LIGO events, based on their high chirp masses, may indeed have formed via this channel \u2013 large masses are expected both because of the preferential capture of heavier BHs by the disc (Yang et al. 2019b) and because repeated mergers are common and lead to hierarchical build-up (Yang et al. 2019a). The high effective spin provides additional support for this channel (Gayathri et al. 2020) as well as events with more unequal mass ratios (e.g. GW190412; see Tagawa et al. 2020b) which have been predicted to occur naturally via dynamical interactions in a dense, gaseous environment (McKernan et al. 2020; Secunda et al. 2020; Tagawa et al. 2020a). These results suggest that BH mergers indeed occur in gas discs, and that they could constitute the building blocks of IMBHs (McKernan et al. 2012; Tagawa et al. 2020a). Recent semi-analytical estimates of in-situ star formation suggest that the accretion of embedded compact objects may also contribute substantially to the growth of MBHs (Dittmann & Miller 2020).","Citation Text":["Just et al. 2012"],"Citation Start End":[[695,711]]}
{"Identifier":"2021MNRAS.502.2815C__Nhan_et_al._2019_Instance_1","Paragraph":"The evolution of our Universe from the Cosmic Dawn (CD) and Epoch of Reionization (EoR) is not very well understood. The H\u2009i 21-cm line promises to be an excellent probe into these epochs (Furlanetto, Oh & Pierpaoli 2006a; Morales & Wyithe 2010; Pritchard & Loeb 2012). The global 21-cm experiments aim to measure the sky-averaged signature of this redshifted line, by using ground-based radio telescopes. Examples of such experiments are the Shaped Antenna measurement of the background RAdio Spectrum (SARAS; Patra et al. 2013; Singh et al. 2017), the Large-Aperture Experiment to Detect the Dark Ages (LEDA; Greenhill & Bernardi 2012), SCI-HI (Voytek et al. 2014), the Broadband Instrument for Global HydrOgen ReioNisation Signal (BIGHORNS; Sokolowski et al. 2015), and the Cosmic Twilight Polarimeter (CTP; Nhan et al. 2019). The Experiment to Detect the Global EoR Signature (EDGES; Bowman et al. 2018) has reported a possible detection of the sky-averaged H\u2009i 21-cm global signal from the CD. However, this signal has an absorption trough of about 0.5 K, which is twice the amplitude predicted by the standard model of cosmology. If confirmed, this measured signal would give us a completely new insight into the physics of the evolution of the Universe. Following this detection, several models explaining this exceedingly deep absorption trough has been explored. Barkana (2018), Fraser et al. (2018), Pospelov et al. (2018), and Slatyer & Wu (2018) had explained this with models with excess cooling from non-standard physics, including dark matter particles scattering off baryons. Fialkov, Barkana & Cohen (2018), Fraser et al. (2018), Pospelov et al. (2018), Ewall-Wice et al. (2018), and Feng & Holder (2018) have explored various possible 21-cm signals, varying the properties of the dark matter particles, in addition to varying the astrophysical parameters. Ewall-Wice et al. (2018), Feng & Holder (2018), and Fialkov & Barkana (2019) have explored models that produce an excess radio background, in addition to the cosmic radio background to explain the large amplitude of the absorption feature in the EDGES detection. This radio excess could be, for example, due to supernovae (Jana, Nath & Biermann 2019; Mirocha & Furlanetto 2019) or primordial black holes (Ewall-Wice et al. 2018).","Citation Text":["Nhan et al. 2019"],"Citation Start End":[[811,827]]}
{"Identifier":"2022ApJ...929..122W___2021_Instance_1","Paragraph":"Calculating the magnetic helicity in the corona by finite volume methods is more complicated than estimating the magnetic helicity flux through the photosphere by optical flow methods, and the significant point of the finite volume methods is how to obtain coronal magnetic fields. Cargill (2009) pointed out the difficulties in the traditional approach of using the Zeeman effect to measure coronal magnetic fields. Some researchers have attempted to explain the observations of coronal oscillations by using magnetohydrodynamic (MHD) wave models, and the magnetic field strength of coronal magnetic fields can be inferred in this process (Nakariakov et al. 1999; Tomczyk et al. 2007). Another approach involves using sophisticated numerical computations to reconstruct realistic 3D nonlinear force-free coronal magnetic field models, which need observed photospheric vector magnetograms as boundary conditions (Wiegelmann & Sakurai 2012, 2021). These coronal models include the MHD models, the magnetohydrostatics (MHS) models, the nonlinear force-free field (NLFFF) models, and the potential-field source-surface models. The force-free approach is a static model that neglects time-dependent phenomena and plasma flows, and the gradient of the plasma pressure and the gravity force in the MHS equation are also neglected as these terms are at least four orders of magnitude smaller than the magnetic pressure in the corona, so that the Lorentz force is vanishing and the electric currents are parallel\/antiparallel to the magnetic field lines (Wiegelmann et al. 2017). However, whichever numerical computation method is chosen to approximate the 3D coronal magnetic field, they all have limitations. DeRosa et al. (2015) compared the influence of spatial resolution on the NLFFF model, and they found that the output of the NLFFF model depends on the spatial resolution of the input photospheric magnetograms (bottom boundary), with the values of free energy generally tending to be higher with increasing resolution, while the values of the relative magnetic helicity vary significantly with different resolutions. They also found that the results from the more highly resolved data were more self-consistent.","Citation Text":["Wiegelmann & Sakurai","2021"],"Citation Start End":[[913,933],[940,944]]}
{"Identifier":"2015AandA...576A.110N__Deharveng_et_al._2010_Instance_1","Paragraph":"The exciting star of the H\u2009ii region has not been clearly identified so far. According to Ortiz et al. (2007), Obj1a (B0 v) and Obj2 (O9 v) are the best candidates for ionization sources of the RCW 41 region. At 5.8 \u03bcm, the region shows a cavity at its center, surrounded by PAH emission in the PhotoDissociation Region (PDR). The H\u03b1 emission is strongest in the center of the nebula (see Fig. 1), implying that the ionizing source might be embedded within the nebula. This phenomenon has been observed in many bubbles associated with H\u2009ii regions (e.g., Deharveng et al. 2010). Thus, we believe that the exciting star must be inside the 5.8 \u03bcm cavity that surrounds the H\u03b1 emission, which excludes the Obj1a source as a potential candidate. Also Obj1a is at the extreme edge of the H\u03b1 emission, which makes it unlikely to be the ionizing candidate. To identify the probable ionizing candidates of the individual H\u2009ii regions, we then used a JHKs catalog of the region inside the cavity and searched for ionizing sources in a circular area of radius ~1.15 pc (i.e., the radius of the 5.8 \u03bcm cavity) from the center of the bubble. Using the extinction laws of Rieke & Lebofsky (1985), the observed (J \u2212 H) color and MJ-spectral type calibration table of Bessel & Brett (1988), and a distance of 1.3 kpc, we selected luminous sources that have spectral types earlier than B3 MS star. We followed the same prescription as described in Samal et al. (2010) of rejecting the most likely giants and foreground sources based on (J \u00d7 [J \u2212 H]) and ([J \u2212 H] \u00d7 [H \u2212 Ks]) diagrams. After these eliminations, we are left with only one massive O-type star located close to the northwestern border of 5.8 \u03bcm cavity. This source is associated with a group of stars in its close vicinity, possibly part of an exciting cluster. The source corresponds to the \u201cObj2\u201d identified by Santos et al. (2012), whose spectroscopic spectral type agrees with the photometric spectral derived form our observations. From the above discussion, we conclude that the Obj2 source is the most likely ionizing star of the RCW\u200941 H\u2009ii region. ","Citation Text":["Deharveng et al. 2010"],"Citation Start End":[[555,576]]}
{"Identifier":"2015ApJ...808...56M__Beaulieu_et_al._2011_Instance_1","Paragraph":"The field of extrasolar planetary transits is one of the most productive and innovative subject in astrophysics in the last decade. Transit observations can be used to measure the size of planets, their orbital parameters (Seager and Mall\u00e9n-Ornelas 2003), and stellar properties (Mandel & Agol 2002; Howarth 2011), to study the atmospheres of planets (Brown 2001; Charbonneau et al. 2002; Tinetti et al. 2007), and to detect small planets (Miralda-Escud\u00e9 2002; Agol et al. 2005) and exomoons (Kipping 2009a, 2009b). In particular, the study of planetary atmospheres requires a high level of photometric precision, i.e., one part in \u223c104 in stellar flux (Brown 2001), which is comparable to the effects of current instrumental systematics and stellar activity (Berta et al. 2011; Ballerini et al. 2012), hence the necessity of testable methods for data detrending. In some cases, different assumptions, e.g., using different instrumental information or functional forms to describe them, leed to controversial results even from the same data sets; examples in the literature are Tinetti et al. (2007), Ehrenreich et al. (2007), Beaulieu et al. (2008) and D\u00e9sert et al. (2009, 2011) for the hot-Jupiter HD 189733b, and Stevenson et al. (2010), Beaulieu et al. (2011) and Knutson et al. (2011, 2014) for the warm-Neptune GJ436b. Some of these controversies are based on Spitzer\/IRAC data sets at 3.6 and 4.5 \u03bcm. The main systematic effect for these two channels is an almost regular undulation with period \u223c3000 s, so called pixel-phase effect, as it is correlated with the relative position of the source centroid with respect to a pixel center (Fazio et al. 2004; Morales-Cald\u00e9ron et al. 2006). Conventional parametric techniques correct for this effect by dividing the measured flux by a polynomial function of the coordinates of the photometric centroid; some variants may include time-dependence (e.g., Stevenson et al. 2010; Beaulieu et al. 2011). Newer techniques attempt to map the intra-pixel variability at a fine-scale level, e.g., adopting spatial weighting functions (Ballard et al. 2010; Cowan et al. 2012; Lewis et al. 2013) or interpolating grids (Stevenson et al. 2012a, 2012b). The results obtained with these methods appear to be strongly dependent on a few assumptions, e.g., the degree of the polynomial adopted, the photometric technique, the centroid determination, calibrating instrument systematics over the out-of-transit only or the whole observation (e.g., Beaulieu et al. 2011; Diamond-Lowe et al. 2014; Zellem et al. 2014). Also, the very same method, applied to different observations of the same system, often leads to significantly different results. Non-parametric methods have been proposed to guarantee a higher degree of objectivity (Carter & Winn 2009; Thatte et al. 2010; Gibson et al. 2012; Waldmann 2012, 2014; Waldmann et al. 2013). Morello et al. (2014, 2015) reanalyzed the 3.6 and 4.5 \u03bcm Spitzer\/IRAC primary transits of HD 189733b and GJ436b obtained during the cryogenic regime, so called \u201ccold Spitzer\u201d era, adopting a blind source separation technique, based on an Independent Component Analysis (ICA) of individual pixel timeseries, in this paper called \u201cpixel-ICA\u201d. The results obtained with this method are repeatable over different epochs, and a photometric precision of one part in \u223c104 in stellar flux is achieved, with no signs of significant stellar variability as suggested in the previous literature (D\u00e9sert et al. 2011; Knutson et al. 2011). The use of ICA to decorrelate the transit signals from astrophysical and instrumental noise, in spectrophotometric observations, has been proposed by Waldmann (2012, 2014) and Waldmann et al. (2013). The reason to prefer such blind detrending methods is twofold: they require very little, if any, prior knowledge of the instrument systematics and astrophysical signals, therefore they also ensure a higher degree of objectivity compared to methods based on approximate instrument systematics models. As an added value, they give stable results over several data sets, also in those cases where more conventional methods have been unsuccessful. Recently, Deming et al. (2015) proposed a different pixel-level decorrelation method (PLD) that uses pixel timeseries to correct for the pixel-phase effect, while simultaneously modeling the astrophysical signals and possible detector sensitivity variability in a parametric way. PLD has been applied to some Spitzer\/IRAC eclipses and synthetic Spitzer data, showing better performances compared to previously published detrending methods.","Citation Text":["Beaulieu et al. (2011)"],"Citation Start End":[[1242,1264]]}
{"Identifier":"2021ApJ...908..159T__Dishoeck_2009_Instance_1","Paragraph":"Complex organic molecules (or COMs) are the building blocks of life. COMs are believed to first form on the icy mantle of dust grains (e.g., Garrod et al. 2008; Jim\u00e9nez-Serra et al. 2016). Subsequently, COMs are released to the gas phase during the warm and hot phase, where icy grain mantles can be heated to high temperatures \n\n\n\n\n\n\nT\n\n\nd\n\n\n\u223c\n100\n\u2013\n300\n\nK\n\n\n (e.g., Blake et al. 1987; Brown et al. 1988; Bisschop et al. 2007). COMs are increasingly observed in the environs of young stellar objects, including hot cores\/corinos around high-mass\/low-mass protostars and protoplanetary disks (see Herbst & van Dishoeck 2009 and van Dishoeck 2014 for recent reviews). In particular, many surveys that search for COMs toward the nearest massive star-forming regions, the Orion Becklin\u2013Neugebauer\/Kleinmann\u2013Low object (BN\/KL), have been conducted with single-dish telescopes from the ground, e.g., IRAM-30 m (Tercero et al. 2010, 2011; Cernicharo et al. 2016), in space, e.g., Herschel (Bergin et al. 2010; Crockett et al. 2014), and with interferometers, e.g., IRAM Plateau de Bure (Favre et al. 2011; Peng et al. 2012; Brouillet et al. 2013), and ALMA (Tercero et al. 2015, 2018; Favre et al. 2017; Pagani et al. 2017, 2019; Peng et al. 2017, 2019). The single-dish telescopes are unable to spatially\/spectrally separate the contribution of the different sources in Orion BN\/KL because of their low resolution (i.e., lines are overlapped). All constraints from those data, therefore, are manifested as uncertainties. Conversely, the extremely high spatial resolution of interferometers naturally can separate the components, which reveals the complexity of the chemical and physical structure in Orion BN\/KL. Observations show the existence of large COMs with very strong binding energies, such as acetic acid (CH3COOH), ethyl formate (C2H5OCHO), propyl cyanide (C3H7CN), methyl acetate (CH3COOCH3), methoxymethanol (CH3OCH2OH), ethylene glycol (OHCH2CH2OH), etc. The sublimation temperature of these large COMs is higher than that of water, i.e., \n\n\n\n\n\n\nT\n\n\nsub\n\n\n\u2273\n152\n\n\n K (see Section 2.1). However, the dust temperature estimated in Orion BN\/KL is far from Tsub. This casts doubt on the popular sublimation mechanism and suggests nonthermal mechanisms that can desorb COMs within these environments.","Citation Text":["Herbst & van Dishoeck 2009"],"Citation Start End":[[597,623]]}
{"Identifier":"2022MNRAS.515.4810L__Ajello_et_al._2020_Instance_1","Paragraph":"In spite of the above considerations, we can appreciate a statistical trend of increasing spectral index (i.e. softer average spectra) while moving through the HSP, ISP, and LSP classes. This result, summarized in Table 1, is consistent with our natural expectations and it supports the reliability of spectral decomposition techniques to extract and analyse the non-thermal emission of these objects. As a consequence, we can investigate the relation existing between the non-thermal component of the optical spectra and the corresponding \u03b3-ray spectra. It is well known that the \u03b3-ray photon index correlates with the position of the synchrotron peak in frequency (e.g. Ajello et al. 2020). So far, however, the presence of thermal contributions in the optical domain has made it very difficult to perform direct comparisons of spectral indices between the optical and the \u03b3-ray energy domains. From our sample, instead, we can infer a substantial degree of linear correlation between the optical spectral index \u03b1opt and its \u03b3-ray counterpart \u03b1\u03b3. The situation, summarized in Fig. 4, is well represented by a linear relation described by\n(2)$$\\begin{eqnarray}\r\n\\alpha _{\\rm opt} = (0.862 \\pm 0.110) \\cdot \\alpha _{\\rm \\gamma} - (0.818 \\pm 0.445),\r\n\\end{eqnarray}$$with weighted correlation coefficient R = 0.743 and null-hypothesis probability p0 = 2.71 \u00d7 10\u221210 (i.e. statistical significance at $6.3\\, \\sigma$ level), for the whole sample, and by\n(3)$$\\begin{eqnarray}\r\n\\alpha _{\\rm opt} = (1.469 \\pm 0.224) \\cdot \\alpha _{\\rm \\gamma} -(1.931 \\pm 0.919),\r\n\\end{eqnarray}$$with weighted correlation coefficient R = 0.751 and null-hypothesis probability p0 = 1.54 \u00d7 10\u22128 (i.e. $5.7\\, \\sigma$), for the subset of 98 objects that were fitted with \u0394\u03b1opt \u2264 0.6 (corresponding to 3 times the systematic uncertainty estimated for the method). The observed correlations are particularly suggestive if compared to the situation that is obtained by simply fitting a power-law model to the optical data without accounting for the host component (more details in Appendix C). Another interesting result, also presented in Fig. 4, is that the above correlation results in a situation where the power-law fits of the \u03b3-ray spectra appear to be softer than the associated optical non-thermal emission by a value that lies mostly close to 1.0. This circumstance points towards a tight relationship between the population of particles that explain the non-thermal component of the optical spectra and the ones that are responsible for the \u03b3-ray emission. The existence of a similar relationship between the observed spectral indices is a characteristic prediction of SSC-based models. However, the systematic difference between the optical and \u03b3-ray energy regions requires further investigation to be properly interpreted.","Citation Text":["Ajello et al. 2020"],"Citation Start End":[[672,690]]}
{"Identifier":"2016AandA...586A..80O__Fornasier_et_al._2015_Instance_2","Paragraph":"Figure 1 shows that in the regions where activity was detected visually, i.e., Hapi, Seth, and Ma\u2019at pits have lower (8\u201313%\/100 nm) spectral slopes than the rest of the comet surface (13\u201322%\/100 nm). In addition to those places, Seth alcoves, the wall of the large Anuket alcove, around the circular features, both clustered and isolated bright features (see Thomas et al. 2015b; Auger et al. 2015; Pommerol et al. 2015b, for definitions) show similar lower spectral slopes than the rest of the surface, even though there was no visual detection of activity features rising from them at the time of the observations used in this study4. This may be because the observing geometry was not suited for their detections during the observations. In the regions we investigated, the Hapi region displays the lowest spectral slopes 8\u201311%\/100 nm (see also Fornasier et al. 2015) together with the isolated bright features (IBFs) and the clustered bright features in the Imhotep region. The locations of the bright features on the Imhotep image (image #4) are shown in Fig. B.4. According to the spectral slope values, the IBFs of Imhotep seem to be more similar to the Hapi region than the active pits of Seth and Ma\u2019at regions. Active pits, alcoves, and the large alcove of Anuket have slope values of typically 10\u201313%\/100 nm. The Ma\u2019at region, which is located on the smaller lobe (head) of the comet, displays higher spectral slope values than the Seth region, which is located on the larger (body) lobe of the comet. In the investigated regions, the highest slope values are detected in the Imhotep region (see Fig. 1d). Here it should be mentioned that the comparison of spectral slopes is performed under the assumption of no spectral reddening between the phase angles of the images we investigated, although the spectral slopes show reddening by phase as presented in Fornasier et al. (2015). Unfortunately, the previous work does not cover all the phase angles of the images we investigated, but the spectral slope variation between 35\u201354\u00b0 (Fig. 3 of  Fornasier et al. 2015) is small so that we can make this comparison. However, if we follow the linear trend of the phase reddening, for the image taken in 70.45\u00b0 phase angle (image #4), the spectral slopes would vary from 15%\/100 nm to 18%\/100 nm in the observations we used. ","Citation Text":["Fornasier et al. (2015)"],"Citation Start End":[[1868,1891]]}
{"Identifier":"2019ApJ...884..154W__Marrone_et_al._2018_Instance_1","Paragraph":"In past decades, single-dish submillimeter surveys have identified populations of massive, dusty star-forming galaxies at z > 1 (e.g., Casey et al. 2014). While these galaxies are rare even at Cosmic Noon (1  z  3) when the star formation activity in the universe peaks, their contribution to the cosmic star formation rate density (CSFRD) equals that of all optical and near-infrared selected galaxies combined (Madau & Dickinson 2014). However, at z > 3 the situation is much less certain. A tail of submillimeter-selected galaxies have been confirmed beyond z > 4, but they trace only the very tip of the star formation rate (SFR) distribution at early times (e.g., Cooray et al. 2014; Strandet et al. 2017; Marrone et al. 2018). The total contribution of dust-obscured star formation, and therefore the census of star formation in the early universe, is unknown. Despite the strongly negative k-correction allowing sources to be found to z = 10, the overwhelming majority of (sub)millimeter-selected galaxies continue to be confirmed at z  3 with Atacama Large Millimeter\/submillimeter Array (ALMA) spectroscopy (Brisbin et al. 2017; Danielson et al. 2017). While some dusty galaxies have been discovered beyond z > 5 through gravitational lensing (Spilker et al. 2016; Zavala et al. 2018b), the lensing correction and selection effects make it challenging to establish their contribution to the CSFRD. Progress is hampered by the limited sensitivity and low spatial resolution of single-dish submillimeter observations and the difficulty of associating detections with counterparts in the optical\u2013near-infrared (NIR). Ultradeep SCUBA surveys over moderate \u223c100 arcmin2 are now pushing into the range of \u201cnormal\u201d SFRs (several 100 M\u2299 yr\u22121, main-sequence galaxies; e.g., Koprowski et al. 2016; Cowie et al. 2017, 2018) and extending to z > 4, but the analysis is often limited by the ability to identify counterparts at other wavelengths and derive accurate redshifts.","Citation Text":["Marrone et al. 2018"],"Citation Start End":[[711,730]]}
{"Identifier":"2019MNRAS.490..157M__Yu_&_Tremaine_2003_Instance_1","Paragraph":"As a class, the fastest stars in our Galaxy are expected to be hypervelocity stars (HVSs). These were first theoretically predicted by Hills (1988) as the result of a three-body interaction between a binary star and the massive black hole in the Galactic Centre (GC), Sagittarius A*. Following this close encounter, a star can be ejected with a velocity \u223c1000 km\u2009s\u22121, sufficiently high to escape from the gravitational field of the MW (Kenyon et al. 2008; Brown 2015). The first HVS candidate was discovered by Brown et al. (2005); a B-type star with a velocity more than twice the Galactic escape speed at its position. Currently about \u223c20 unbound HVSs with velocities \u223c300\u2013700 km\u2009s\u22121 have been discovered by targeting young stars in the outer halo of the MW (Brown, Geller & Kenyon 2014). In addition, tens of mostly bound candidates have been found at smaller distances but uncertainties prevent the precise identification of the GC as their ejection location (e.g. Hawkins et al. 2015; Vickers, Smith & Grebel 2015; Zhang, Smith & Carlin 2016; Marchetti et al. 2017; Ziegerer et al. 2017). HVSs are predicted to be ejected from the GC with an uncertain rate around 10\u22124 yr\u22121 (Yu & Tremaine 2003; Zhang, Lu & Yu 2013), two orders of magnitude larger than the rate of ejection of RSs with comparable velocities from the stellar disc (Brown 2015). Because of their extremely high velocities, HVS trajectories span a large range of distances, from the GC to the outer halo. Thus, HVSs have been proposed as tools to study the matter distribution in our Galaxy (e.g. Gnedin et al. 2005; Sesana, Haardt & Madau 2007; Kenyon et al. 2014; Fragione & Loeb 2017; Rossi et al. 2017; Contigiani, Rossi & Marchetti 2018) and the GC environment (e.g. Zhang et al. 2013; Madigan et al. 2014), but a larger and less observationally biased sample is needed in order to break degeneracies between the GC binary content and the Galactic potential parameters (Rossi et al. 2017). Using the fact that their angular momentum should be very close to zero, HVSs have also been proposed as tools to constrain the solar position and velocity (Hattori, Valluri & Castro 2018a). Other possible alternative mechanisms leading to the acceleration of HVSs are the encounter between a single star and a massive black hole binary in the GC (e.g. Yu & Tremaine 2003; Sesana, Haardt & Madau 2006, 2008), the interaction between a globular cluster with a single or a binary massive black hole in the GC (Capuzzo-Dolcetta & Fragione 2015; Fragione & Capuzzo-Dolcetta 2016), and the tidal interaction of a dwarf galaxy near the centre of the Galaxy (Abadi et al. 2009). Another possible ejection origin for HVSs and high-velocity stars in our Galaxy is the Large Magellanic Cloud (LMC; Boubert & Evans 2016; Boubert et al. 2017; Erkal et al. 2018), orbiting the MW with a velocity \u223c380 km\u2009s\u22121 (van der Marel & Kallivayalil 2014).","Citation Text":["Yu & Tremaine 2003"],"Citation Start End":[[1180,1198]]}
{"Identifier":"2022ApJ...929..186L__Ferraro_et_al._2018b_Instance_1","Paragraph":"Our group is addressing this problem by combining a variety of complementary perspectives: (i) by constructing a new generation of high-quality star density profiles derived from star counts instead of surface brightness (see Lanzoni et al. 2007a, 2010, 2019; Miocchi et al. 2013; Pallanca et al. 2021); (ii) by investigating the population of stellar exotica (Ferraro et al. 2001, 2003, 2015, 2016; Pallanca et al. 2010, 2013,2014, 2017; Cadelano et al. 2017, 2018, 2020) and their connection with the dynamical evolution of the parent cluster (see Ferraro et al. 2009, 2012, 2018a, 2019; Lanzoni et al. 2016); (iii) by characterizing the three-dimensional (3D) global velocity space through the analysis of the velocity dispersion profile and rotation curve from resolved star spectroscopy (Lanzoni et al. 2013, 2018a, 2018b; Ferraro et al. 2018b) and proper motions (PMs; see Raso et al. 2020). The determination of GGC internal kinematics from resolved star velocities is particularly relevant and challenging. In this context we promoted the ESO-VLT Multi-Instrument Kinematic Survey (hereafter the MIKiS survey; Ferraro et al. 2018b, 2018c), a project specifically designed to characterize the kinematical properties of a sample of GGCs in different dynamical evolutionary stages from the radial velocities (RVs) of hundreds of individual stars distributed over the entire radial range of each stellar system. To this end, the survey fully exploits the spectroscopic capabilities of different instruments currently available at the ESO Very Large Telescope (VLT): originally designed to use the adaptive optics (AO) assisted integral-field spectrograph SINFONI, the multiobject integral-field spectrograph KMOS, and the multiobject fiber-fed spectrograph FLAMES\/GIRAFFE, it has been recently complemented with individual projects and an ongoing large program (PI: Ferraro) fully exploiting the remarkable performances of the AO-assisted integral-field spectrograph MUSE.","Citation Text":["Ferraro et al. 2018b"],"Citation Start End":[[828,848]]}
{"Identifier":"2020ApJ...890...89G__Bieler_et_al._2015_Instance_2","Paragraph":"Consequently, the neutral gas measured in situ in the coma of comet 67P by the ROSINA experiment (and also MIRO, VIRITS, and Alice) on board Rosetta likely originated several tens of meters beneath the primordial surface of the comet. ROSINA observations provided evidence that this comet is formed from pristine material that has not been significantly altered after its formation in the first Myr of the solar nebula stage. The high abundance of super-volatiles like CO and CO2 (Le Roy et al. 2015), the detection of argon (Balsiger et al. 2015), of molecular nitrogen (Rubin et al. 2015), of molecular oxygen (Bieler et al. 2015), of a high D\/H in HDO\/H2O and D2O\/HDO and HDS\/H2S (Altwegg et al. 2015, 2017), and of hydrogen halides (De Keyser et al. 2017; Dhooghe et al. 2017), coupled with the low density, high porosity, and homogeneity of the nucleus (P\u00e4tzold et al. 2016) and the absence of signatures of aqueous alteration (see Capaccioni et al. 2015; Davidsson et al. 2016; Quirico et al. 2016; Bardyn et al. 2017) all indicate that comet 67P formed at low temperature and did not experience any substantial global-scale heating after its formation. This suggests that 67P is representative of the solar nebula material from which the solar system had formed. This has strong implications not only for how the measurements made in cometary environments can be used to constrain the protosolar environment but also for the contribution of comets to Earth\u2019s composition. For instance, the measurement of the D\/H isotopic ratio in 67P (Altwegg et al. 2015) suggests that comets cannot be considered as the main source of water on Earth. The discovery of significant amounts of O2 in comets (Bieler et al. 2015) was not predicted by astrochemical models and challenges our understanding of the chemistry of molecular clouds and of the protosolar nebula. However, the Jupiter-family comets (JFCs; which include 67P) are a diverse groups. Indeed, even if Giotto measurements indicate that comet 1P\/Halley contains similar amounts of O2 (Rubin et al. 2015), different D\/H ratios (lower than observed for 67P and compatible with the D\/H ratio in the Earth\u2019s oceans) have been measured for other JFC comets like Hartley 2 (Balsiger et al. 2015) and 46P\/Wirtanen (Lis et al. 2019). The causes for this diversity may be already present at the formation of these comets or may result from a different evolution after their formation.","Citation Text":["Bieler et al. 2015"],"Citation Start End":[[1698,1716]]}
{"Identifier":"2019MNRAS.490..243C__Jeli\u0107_et_al._2010_Instance_1","Paragraph":"The main challenge in detecting the cosmological H\u2009i 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\u2013free 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\u2009i signal (Zaldarriaga, Furlanetto & Hernquist 2004; Bharadwaj & Ali 2005; Jeli\u0107 et al. 2008; Bernardi et al. 2009; Jeli\u0107 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\u2009i 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 \u2018smoothness\u2019 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\u2009i brightness temperature fluctuation, a cylindrical (2D) power spectrum, P(k\u22a5, k\u2225) is a useful diagnostic in terms of foreground avoidance. Spectral smoothness of foregrounds confines the majority of foreground power to low k\u2225 modes, resulting in \u2018foreground wedge\u2019. In the foreground avoidance technique, the EoR signal is searched for outside this wedge, in the so-called \u2018EoR window\u2019 (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 Text":["Jeli\u0107 et al. 2010"],"Citation Start End":[[876,893]]}
{"Identifier":"2016AandA...591A..13V__Govoni_&_Feretti_2004_Instance_1","Paragraph":"The origin and evolution of cosmic magnetism are at present poorly understood. Answering the many open questions surrounding the physics of astrophysical magnetic fields is a difficult task since magnetic fields can be significantly affected by structure and galaxy formation and evolutionary processes. Their strength can be amplified, for example, in galaxy clusters through mergers and in galaxies through large-scale dynamos, invoking differential rotation and turbulence. Insights into the origin and properties of magnetic fields in the Universe could be provided by probing them on even larger scales. Along filaments and voids of the cosmic web, turbulent intracluster gas motions have not yet enhanced the magnetic field; its strength thus still depends on the seed field intensity, in contrast to galaxy clusters, where it probably mostly reflects the present level of turbulence (see, e.g., Donnert et al. 2009; Xu et al. 2010, 2011). Intervening magnetoionic media cause a difference in the phase velocity between the left-handed and right-handed circular polarization components of the linearly polarized synchrotron radiation emitted by a background radio source (e.g., Carilli & Taylor 2002; Govoni & Feretti 2004). This effect translates into a rotation of the intrinsic polarization angle, \u03c80, (1)\\begin{equation} \\psi(\\lambda^2)=\\psi_0+\\phi\\lambda^2. \\end{equation}\u03c8(\u03bb2)=\u03c80+\u03c6\u03bb2.Following Burn (1966), the observed polarization angle, \u03c8, depends on the observation wavelength, \u03bb, through the Faraday depth, \u03c6, (2)\\begin{equation} \\phi=a_0\\int_0^{z_{\\rm s}}B_l(z)~n_{\\rm e}(z)~\\frac{\\mathrm{d}l}{\\mathrm{d}z}~\\mathrm{d}z, \\label{definition} \\end{equation}\u03c6=a0\u222b0zsBl(z)ne(z)dldzdz,where a0 depends only on fundamental constants, ne is the electron density, Bl is the component of the magnetic field along the line of sight, and zs is the redshift of the source. When the rotation is completely due to a foreground screen, the Faraday depth has the same value as the rotation measure (RM), defined by (3)\\begin{equation} {\\rm RM}\\equiv\\frac{\\partial \\psi}{\\partial \\lambda^2}\\cdot \\end{equation}RM\u2261\u2202\u03c8\u2202\u03bb2\u00b7The Faraday depth is assumed to be positive when the line-of-sight average component of the magnetic field points toward the observer, otherwise it is negative for a field with an average component pointing away from the observer. The amount of Faraday depth measured by radio observations along a given line of sight is the sum of the contributions from the Milky Way, the emitting radio source, and any other source and large-scale structure in between hosting a magnetized plasma. The investigation of these contributions and of their possible dependence on redshift is essential to discriminate among the different scenarios of magnetic field formation and evolution and, therefore, crucial for the understanding of cosmic magnetism. Sensitive observations, a good knowledge of the Galactic Faraday foreground screen, and a statistical approach that is able to properly combine all of the observational information are necessary. An all-sky map of the Galactic Faraday rotation foreground and an estimate of the overall extragalactic contribution has been derived by Oppermann et al. (2012, 2015) in the framework of \u201cInformation Field Theory\u201d (En\u00dflin et al. 2009) by assuming a correlated Galactic foreground and a completely uncorrelated extragalactic term. In this paper, we propose a new, fully Bayesian approach aiming at further disentangling the contribution intrinsic to emitting sources from the contribution due to the intergalactic environment between the source and the observer and at investigating the dependence of these contributions on redshift. ","Citation Text":["Govoni & Feretti 2004"],"Citation Start End":[[1207,1228]]}
{"Identifier":"2021MNRAS.504.1939G__Zhang_&_Yan_2011_Instance_2","Paragraph":"For the magnetic field configurations considered in this work, the polarization angle can only change exactly by \u0394\u03d5 = 90\u25cb and a gradual change of the PA is not possible. There are tantalizing hints of a 90\u25cb change in the PA in some of the GRBs, as discussed above, but the results are not yet conclusive. The result presented by Sharma et al. (2019) where the PA changes by 90\u25cb twice over the emission is again very exciting as such a change over a single pulse can only occur for the Btor field configuration. The only difficulty, according to the modelling done here, is that both 90\u25cb changes occur in the decaying tail of the pulse when high latitude emission dominates the flux. In the measurement presented by Sharma et al. (2019), the PA shows a change close to the peak of the emission. Another scenario in which a 90\u25cb PA change can be obtained includes contribution from multiple pulses and when the LOS is close to the edge of the jet, such that \u03b8obs \u2248 \u03b8j, along with a change in bulk \u0393 between the pulses which would change \u03bej = (\u0393\u03b8j)2. Alternatively, such a change in the PA can be obtained due to magnetic reconnection, e.g. in the ICMART model (Zhang & Yan 2011), where the local magnetic field orientation, which is orthogonal to the wave vector of the emitted photon, itself changes by 90\u25cb as the field lines are destroyed and reconnected in the emission region (Deng et al. 2016). To obtain a change in the PA other than \u0394\u03d5 = 90\u25cb or to get a gradually changing PA the condition for axisymmetry must be relaxed and the magnetic field configuration or orientation in the emission region must change. One possibility is that if the different pulses that contribute to the emission arise in a \u2018mini-jet\u2019 within the outflow (e.g. Shaviv & Dar 1995; Lyutikov & Blandford 2003; Kumar & Narayan 2009; Lazar, Nakar & Piran 2009; Narayan & Kumar 2009; Zhang & Yan 2011). In this case, the different directions of the mini-jets or bright patches w.r.t. the LOS (e.g. Granot & K\u00f6nigl 2003; Nakar & Oren 2004) would cause the PA to also be different between the pulses even for a field that is locally symmetric w.r.t the local radial direction (e.g. B\u22a5 or B\u2225) as well as for fields that are axisymmetric w.r.t to the centre of each mini-jet (e.g. a local Btor for each mini-jet). Finally, broadly similar result would follow from an ordered field within each mini-jet (Bord) which are incoherent between different mini-jets. Time-resolved measurement in such a case would naturally yield a time-varying PA. Alternatively, as shown by Granot & K\u00f6nigl (2003) for GRB afterglow polarization, a combination of an ordered field component (e.g. Bord) and a random field, like B\u22a5, can give rise to a time-varying PA between different pulses that, e.g. arise from internal shocks. The ordered field component here would be that advected from the central engine and the random field component can be argued to be shock-generated. Notice that the ordered field component should not be axisymmetric in order for the PA to smoothly vary.","Citation Text":["Zhang & Yan 2011"],"Citation Start End":[[1858,1874]]}
{"Identifier":"2019MNRAS.490.2071Y__Riess_et_al._2018_Instance_1","Paragraph":"\nSet II: we now focus on the observational constraints on the model parameters after the inclusion of the local measurement of H0 by Riess et al. (2018) with the previous data sets (CMB, Pantheon, and CC) in order to see how the parameters could be improved with the inclusion of this data point. Since for this present UM, the estimation of H0 from CMB alone is compatible with the local estimation of H0 by Riess et al. (2018), thus, we can safely add both the data sets to see whether we could have something interesting. Following this, we perform another couple of tests after the inclusion of R18. The observational results on the model parameters are summarized in Table 4. However, comparing the observational constraints reported in Table 3 (without R18 data) and Table 4 (with R18), one can see that the inclusion of R18 data (Riess et al. 2018) does not seem to improve the constraints on the model parameters. In fact, the estimation of the Hubble constant H0 remains almost similar to what we found in Table 3. In order to be more elaborate in this issue, we have compared the observational constraints on the model parameters before and after the inclusion of R18 to other data sets. In Figs 7 (CMB versus CMB+R18), 8 (CMB+CC versus CMB+CC+R18), 9 (CMB+Pantheon versus CMB+Pantheon+R18), and 10 (CMB+Pantheon+CC versus CMB+Pantheon+CC+R18), we have shown the comparisons which prove our claim. One can further point out that the strong correlation between the parameters \u03bc and H0 as observed in Fig. 5 still remains after the inclusion of R18 [see specifically the (\u03bc, H0) planes in Figs 7\u201310]. The physical nature of \u03bc does not alter at all. That means the correlation between H0 and \u03bc is still existing after the inclusion of R18 to the previous data sets, such as CMB, Pantheon, and CC. In addition to that since \u03bc \u2272 0.9 according to all the observational data sets, thus, the transition from past decelerating era to current accelerating era occurs to be around z \u2272 0.6, similar to what we have found with previous data sets (Table 3).","Citation Text":["Riess et al. (2018)"],"Citation Start End":[[133,152]]}
{"Identifier":"2016MNRAS.462.3441D__Namouni_1999_Instance_1","Paragraph":"In principle, Fig. 5, central panel G, shows that (469219) 2016 HO3 may have been locked in a Kozai\u2013Lidov resonance with \u03c9 librating about 270\u00b0 for nearly 100 kyr and probably more. Because of the Kozai\u2013Lidov resonance, both e (central panel E) and i (central panel F) oscillate with the same frequency but out of phase (for a more detailed view, see Fig. 4, panels E and F); when the value of e reaches its maximum the value of i is the lowest and vice versa ($\\sqrt{1 - e^2} \\cos i \\sim$ constant, see Fig. 4, panel B). During the simulated time and for the nominal orbit, 469219 reaches perihelion and aphelion the farthest possible from the ecliptic. Fig. 5, G-panels, show that for other incarnations of the orbit of 469219, different from the nominal one, \u03c9 may librate about 90\u00b0 as well during the simulated time interval. However, is this a true Kozai\u2013Lidov resonance? Namouni (1999) has shown that the secular evolution of co-orbital objects is viewed more naturally in the er\u03c9r-plane, where er and \u03c9r are the relative eccentricity and argument of perihelion computed as defined in Namouni's work (see equations 3 in Namouni 1999); these are based on the vector eccentricity and the vector inclination. Fig. 6 shows the multi-planet er\u03c9r-portrait for the nominal orbit of this object. It clearly resembles figs 13 and 19 in Namouni (1999). Asteroid 469219 librates around $\\omega _{\\rm r}=-90{^\\circ }$ for Venus, the Earth, and Jupiter. This behaviour corresponds to domain III in Namouni (1999), horseshoe-retrograde satellite orbit transitions and librations (around $\\omega _{\\rm r}=-90{^\\circ }$ or 90\u00b0). For a given cycle, the lower part corresponds to the horseshoe phase and the upper part to the quasi-satellite or retrograde satellite phase. This is not the Kozai\u2013Lidov resonance; in this case, the Kozai\u2013Lidov domain (domain II in Namouni 1999) is characterized by libration around $\\omega _{\\rm r}=0{^\\circ }$ (or 180\u00b0) which is only briefly observed at the end of the backwards integrations (see Fig. 6). The Kozai\u2013Lidov resonance is however in action at some stage in the orbits displayed in Figs 5 and 8. Our calculations show that the orbital evolution followed by 469219 is the result of the dominant secular perturbation of Jupiter as the periodic switching between co-orbital states ceases after about 8 kyr if Jupiter is removed from the calculations. Fig. 7 shows that, without Jupiter, 469219 switches between the Kozai\u2013Lidov domain and that of horseshoe-quasi-satellite orbit transitions and librations (including both \u221290\u00b0and 90\u00b0). Jupiter plays a stabilizing role in the dynamics of objects following orbits similar to that of 469219. It is not surprising that Jupiter instead of the Earth or Venus is acting as main secular perturber of 469219. Ito & Tanikawa (1999) have shown that the inner planets share the effect of the secular perturbation from Jupiter; in fact, Venus and our planet exchange angular momentum (Ito & Tanikawa 2002). In their work, these authors argue that the inner planets maintain their stability by sharing and weakening the secular perturbation from Jupiter. Tanikawa & Ito (2007) have extended this analysis to conclude that, regarding the secular perturbation from Jupiter, the terrestrial planets form a collection of loosely connected mutually dynamically dependent massive objects. The existence of such planetary grouping has direct implications on the dynamical situation studied here; if Jupiter is removed from the calculations, the overlapping secular resonances and the recurrent dynamics disappear as well.","Citation Text":["Namouni (1999)"],"Citation Start End":[[877,891]]}
{"Identifier":"2015ApJ...806....1M__Sheth_&_Tormen_1999_Instance_1","Paragraph":"The measurement of galaxy clustering statistics is one of the most powerful probes in observational cosmology (Tegmark et al. 2004; Cole et al. 2005; Percival et al. 2010; Saito et al. 2011; Zehavi et al. 2011; Reid et al. 2012, 2014; Samushia et al. 2014). However, our lack of a detailed understanding of the relationship between the distribution of galaxies and that of dark matter limits the full use of the measured amplitude of the clustering signal for constraining cosmological parameters. Seljak et al. (2005) first proposed and demonstrated that it is possible to break this degeneracy between the unknown bias between galaxies and matter and the cosmological parameters by utilizing the theoretical dependence of galaxy bias on halo mass (Mo & White 1996; Sheth & Tormen 1999). By probing the halo masses of galaxies via the weak gravitational lensing around galaxies, commonly referred to as galaxy\u2013galaxy lensing, on small scales (see also Yoo et al. 2006; Cacciato et al. 2009; Li et al. 2009; Leauthaud et al. 2012; Tinker et al. 2012; Cacciato et al. 2013, 2014; Gillis et al. 2013; More et al. 2013; Simpson et al. 2013; see also for the galaxy\u2013galaxy lensing measurements Fischer et al. 2000; Mandelbaum et al. 2006; van Uitert et al. 2011; Velander et al. 2014). An alternative method, which uses the ratio of the clustering and lensing signals, has also been proposed in order to avoid the complex astrophysics that complicates the interpretation of these observables on scales smaller than the typical halo radii (Baldauf et al. 2010). As a proof of this concept, Mandelbaum et al. (2013, hereafter RM13) used the state-of-the-art measurement of galaxy\u2013galaxy lensing signal and galaxy clustering up to large-scales (\n\n\n\n\n\n) by combining the spectroscopic and multi-color imaging galaxy catalogs from the Sloan Digital Sky Survey I\/II (SDSS I\/II; York et al. 2000; Eisenstein et al. 2001; Strauss et al. 2002). They showed that the joint analysis provides a significant improvement in the cosmological parameters \n\n\n\n\n\n and \n\n\n\n\n\n when combined with results from the WMAP7 experiment (Komatsu et al. 2011).","Citation Text":["Sheth & Tormen 1999"],"Citation Start End":[[767,786]]}
{"Identifier":"2022AandA...664L..16E__Spruit_(2002)_Instance_1","Paragraph":"This n\u2004=\u20041, CT\u2004=\u2004216 case corresponds to a new asteroseismic-calibrated version of the original TS dynamo. The parameter CT is introduced here to account for uncertainties on the adopted timescale for the damping of the azimuthal field (while the \u03b1 parameter in Fuller et al. 2019 was introduced for the saturated value of \u03c9A, hence the relation CT\u2004=\u2004\u03b13). We thus find that the damping timescale adopted for the azimuthal field in the original TS dynamo has to be increased by a factor of about 200 to correctly reproduce the asteroseismic data of evolved stars. Ideally, one would expect a value of CT closer to unity for a well-defined physical process. We however recall here that this timescale is known to be quite approximated as mentioned by Spruit (2002) who indicated that the original estimates of the dynamo process were made by neglecting all multiplying factors of order unity and that these factors could sometimes compound to rather large numbers. A first uncertainty is related to the exact value of the growth rate of the Tayler instability which seems to be somewhat smaller than the adopted value of \u03c9A (e.g., Goldstein et al. 2019). In the same way, the correction to this growth rate in the case of fast rotation (i.e., \u03a9\u2004\u226b\u2004\u03c9A) introduces another uncertainty that also seems to overestimate its value (Ib\u00e1\u00f1ez-Mej\u00eda & Braithwaite 2015). Another source of uncertainty is related to the fact that these estimates of the growth rate are only based on the fastest growing non-axisymmetric m\u2004=\u20041 Tayler mode, while the small-scale dynamo process could perfectly be dominated by modes with different values of m (see e.g., Ib\u00e1\u00f1ez-Mej\u00eda & Braithwaite 2015). In this context, the damping timescale adopted in the original TS dynamo seems to correspond more to a minimal than an exact value; it is thus interesting to be able to constrain its value from asteroseismic data and it is then not surprising to deduce a longer timescale in this way. Of course, it is difficult to speculate whether such a large increase of about two orders of magnitude is really physically motivated in the framework of the original TS dynamo and we can only await numerical simulations performed under more realistic stellar conditions to obtain some answers on this point.","Citation Text":["Spruit (2002)"],"Citation Start End":[[749,762]]}
{"Identifier":"2020AandA...636A..63C__Halfen_et_al._2011_Instance_1","Paragraph":"Among these latter complex molecules, formamide (HCONH2) is of crucial importance. Chemical reactions of molecules containing H, C, N, and O such as formamide are considered a plausible pathway for synthesis of biomolecules under prebiotic conditions (Oparin 1938). Formamide is the simplest molecule containing the peptide bond, which is known to be the basis for assembling proteins and polypeptides starting from amino acids, with a crucial role in the biotic processes of life on Earth. Formamide was first detected in the gaseous phase in two high-mass star forming regions: Orion-KL, an active star forming region, and the giant molecular cloud SgrB2 (e.g., Turner 1991; Nummelin et al. 1998; Halfen et al. 2011); later it was detected in the comet Hale Bopp (Bockelee-Morvan et al. 2000), whose chemical composition is suspected to be similar to the chemical composition of the primitive Solar Nebula. Furthermore, during landing of Philae aboard Rosetta mission, in situ mass spectrometer data inferred the presence of formamide in the comet nucleus with the highest abundance after water (Goesmann et al. 2015). In recent years, formamide was also observed in two types of low-mass star forming environment: shocked regions by protostellar jets (e.g., Codella et al. 2017) and hot corinos (Kahane et al. 2013; L\u00f3pez-Sepulcre et al. 2015; Marcelino et al. 2018; Imai et al. 2016; Oya et al. 2017; Lee et al. 2017). In high-temperature regions (> 100 K) such as hot corinos, thermal desorption is responsible for sublimation of frozen mantles into the gas phase. Moreover, outbursting young stars like V883 Ori are good new targets for looking for organic complex molecules that thermally desorb from icy mantles. The interpretation of observations can benefit from information coming from the laboratory, where it is possible to simulate the thermal desorption process and UV irradiation of formamide under simulated space conditions. Here we report two laboratory analyses. We focus both on photostability and on thermal desorption of pure formamide, and in the presence of grains, before and after UV irradiation. Section 2 presents our choice of minerals used as analog samples and the experimental setup. In Sect. 3, in situ UV irradiation of pure formamide ice and its adsorption by minerals at 63 K are investigated by infrared spectroscopy (FTIR). Section 4 presents a temperature-programmed desorption (TPD) analysis of pure formamide ice and in the presence of TiO2 dust, before and after UV irradiation.","Citation Text":["Halfen et al. 2011"],"Citation Start End":[[699,717]]}
{"Identifier":"2019AandA...622A.117S__Burg_et_al._2016_Instance_2","Paragraph":"Figure 1a shows for each cluster a visualization of the resulting background-subtracted color-magnitude diagram, produced with an approach similar to that of van der Burg et al. (2016). Specifically, we subtracted the residual background contamination as follows: 1) We start from a candidate cluster member sample obtained as discussed in Sect. 3.1, and a control field sample from the GOODS-S field (see Sect. 2.3) that is selected in the same manner. 2) For each galaxy in the candidate member sample, we calculate a \u201cweight\u201d that corresponds to the statistical excess of the candidate member sample over the control field density at the magnitude and colors of the given galaxy. Weights are calculated as follows (see van der Burg et al. 2016, for a more detailed description): first, all candidate member weights are initially set to 1. Then, for each galaxy in the control field sample we subtract the corresponding \u201cbackground contamination\u201d from the candidate member sample by appropriately reducing the weights of all candidate members that lie within a distance in the color (m814\u2013m140) \u2013 color (m140\u2013[3.6]) \u2013 magnitude (m140) space given by their photometric uncertainties (1\u03c3), with a minimum distance of 0.3 mag, effectively resulting in a smoothing of galaxy densities in the color-color-magnitude space. If no galaxies are found within this distance, we double the search distance, and then if necessary increase the search distance to 1.3 times the distance to the closest galaxy. This criterion allows the full subtraction of the background contamination estimated from all galaxies in the control field sample, while reducing the weights of those candidate members that are more similar to the galaxies in the control field. For each considered control-field galaxy, the weights of all selected candidate members identified with the above criterion are reduced so that the contribution of the considered field galaxy (normalized by the areas of the probed cluster region and of the control field) is removed. At the end of this procedure, the contributions of all galaxies from the field sample have been subtracted from the candidate member sample.","Citation Text":["van der Burg et al. 2016"],"Citation Start End":[[722,746]]}
{"Identifier":"2018MNRAS.473.4566P__Papaderos_et_al._2006_Instance_2","Paragraph":"The young starburst inferred by the detections of high ionization emission line of He\u2009II \u03bb4686 and the blue WR bump in this and previous works (Guseva et al. 2000; Brinchmann, Kunth & Durret 2008) is confirmed by the age estimates made here for the bright and faint regions in Mrk 22 as \u223c4 and \u223c10\u2009Myr, respectively. Unlike previous works, we carried out abundance analysis for both the regions separately. We found an appreciable metallicity difference of \u223c0.5 dex between the bright region [12 + log (O\/H) \u223c 8] and the faint region [12 + log(O\/H) \u223c 7.5]. The separation between two regions is \u223c0.6\u2009kpc. Typical metallicity gradients in normal spiral galaxies have been found between \u22120.009 and \u22120.231 dex kpc\u22121, with an average gradient of \u22120.06 dex kpc\u22121 (Zaritsky, Kennicutt & Huchra 1994). The observed metallicity difference between the two regions in Mrk 22 is too large to be explained as a normal galactic metallicity gradient. The chemical composition as measured from the gas-phase metallicity [12 + log(O\/H)] shows various degree of spatial variations in different types of dwarf galaxies. For instance, shallow gradient in metallicity is seen in SBS 0335\u2212052 (Papaderos et al. 2006) while no significant variations were seen in Mrk 35 (Cair\u00f3s et al. 2007). A study on a large sample indicates that normal BCD galaxies are chemically homogeneous (Kobulnicky & Skillman 1996; Papaderos et al. 2006; Kehrig et al. 2008; Cair\u00f3s et al. 2009; P\u00e9rez-Montero & Contini 2009; P\u00e9rez-Montero et al. 2011; H\u00e4gele et al. 2011; Garc\u00eda-Benito & P\u00e9rez-Montero 2012; Lagos & Papaderos 2013). On the other hand, the metallicity of extremely metal-poor galaxies is usually not homogeneous within the galaxy, with the low metallicity seen in regions of intense star formation (Papaderos et al. 2006; Izotov & Thuan 2009; Levesque et al. 2011; S\u00e1nchez Almeida et al. 2013, 2014, 2015). However, large metallicity gradients are not common in dwarf galaxies. The simplest explanation for large metallicity difference in a single system is a recent merger of two galaxies with different metallicity. In a few cases, significantly large metallicity differences between star-forming regions in dwarf galaxies were seen and understood in terms of recent tidal interactions or mergers (L\u00f3pez-S\u00e1nchez, Esteban & Rodr\u00edguez 2004a,b; L\u00f3pez-S\u00e1nchez, Esteban & Garc\u00eda-Rojas 2006; L\u00f3pez-S\u00e1nchez & Esteban 2009, 2010). The evolution in terms of metallicities in interacting dwarf galaxies is fairly complex as it can depend on various factors such as mixing of metals with the interstellar medium (ISM), possible outflows of metals, and inflow of metal-poor gas in tidally interacting systems.","Citation Text":["Papaderos et al. 2006"],"Citation Start End":[[1387,1408]]}
{"Identifier":"2021AandA...650A.203G__Grand_et_al._2018_Instance_1","Paragraph":"In the past, as well as in the more recent literature, there have been many attempts to explain the different chemical evolutionary paths of different MW components, particular those of the thin and thick discs. The outcome of these studies is that the observed different chemical evolutionary paths are related to differences in the main physical processes that drive galaxy evolution, the most significant of which are the gas accretion time-scale and the star formation efficiency and, possibly, radial migration (Larson 1972; Lynden-Bell 1975; Pagel & Edmunds 1981; Matteucci & Greggio 1986; Matteucci & Brocato 1990; Ferrini et al. 1994; Prantzos & Aubert 1995; Chiappini et al. 1997, 2001; Portinari & Chiosi 1999; Bekki & Tsujimoto 2011; Micali et al. 2013; Sahijpal 2014; Snaith et al. 2014; Grisoni et al. 2017; Grand et al. 2018; Spitoni et al. 2021). A good agreement between observations and theoretical predictions for the Galaxy is obtained by models that are based on the assumption that the disc formed via the infalling of gas (Chiosi 1980; Matteucci & Francois 1989; Chiappini et al. 1997). The formation of the different components is associated with distinct sequential main episodes of gas accretion (infall phases) that, at first, rapidly accumulates in the central regions and then, more slowly, in the more external ones, according to the so-called \u2018inside-out scenario\u2019 (Chiappini et al. 2001). In particular, the three-infall model, devised by Micali et al. (2013), is capable of reproducing the abundance patterns of the MW halo, thick and thin disc at once. In this model, the halo forms in a first gas infall episode of short timescale (0.2 Gyr) and mild star formation efficiency, \u03bd\u2004=\u20042 Gyr\u22121, lasting for about 0.4 Gyr. It is immediately followed by the thick disc formation, characterised by a somewhat longer infall timescale (1.2 Gyr), a longer duration (about 2 Gyr) and a higher star formation efficiency (\u03bd\u2004=\u200410 Gyr\u22121). Finally, star formation continues in the thin disc with a longer infall timescale (6 Gyr in the solar vicinity) and is still continuing to this day, with a star-formation efficiency of \u03bd\u2004=\u20041 Gyr\u22121. The [O\/Fe] vs. [Fe\/H] path is thus continuous across the regions populated by halo, thick, and thin disc stars. While Micali et al. (2013) described the chemical enrichment as continuous across the three different infall stages, Grisoni et al. (2017) used also an alternative scheme where the thin and thick disc components evolve separately, in a parallel approach (see also Chiappini 2009). In such a parallel approach, the disc populations are assumed to form in parallel but to proceed at different rates. The gas infall exponentially decreases with a timescale that is 0.1 Gyr and 7 Gyr, for the thick and thin disc, respectively. This alternative approach better reproduces the presence of the metal-rich \u03b1-enhanced stars in the [Mg\/Fe] vs. [Fe\/H] diagram obtained with the recent AMBRE data (Mikolaitis et al. 2017).","Citation Text":["Grand et al. 2018"],"Citation Start End":[[821,838]]}
{"Identifier":"2022ApJ...925..136C__Matsumoto_et_al._2015_Instance_1","Paragraph":"Extensive studies have attempted to probe the near-infrared EBL through different methods. Galaxy counts measures emission from resolved galaxies and the observations are extrapolated to estimate contributions from faint galaxies below the detection limit (Keenan et al. 2010; Dom\u00ednguez et al. 2011; Helgason et al. 2012; Driver et al. 2016; Saldana-Lopez et al. 2021; Koushan et al. 2021). This sets a lower bound to the integrated galaxy light (IGL) component of EBL, as EBL from diffuse emission or sources not associated with galaxies is not included. Direct measurements using absolute photometry capture all emission in the EBL (Bernstein 2007; Levenson et al. 2007; Tsumura et al. 2013; Matsumoto et al. 2015; Sano et al. 2015; Matsuura et al. 2017; Zemcov et al. 2017; Lauer et al. 2021; Sano et al. 2020), however, absolute photometry is challenging since systematic errors and foreground contamination have to be tightly controlled. Fluctuation analysis is an alternative approach that is less susceptible to foreground contamination (e.g., zodiacal light, Galactic cirrus), since foregrounds and the EBL signal have distinct spatial and spectral correlation features (Kashlinsky et al. 2005; Thompson et al. 2007; Matsumoto et al. 2011; Cooray et al. 2012; Kashlinsky et al. 2012; Zemcov et al. 2014; Mitchell-Wynne et al. 2015; Seo et al. 2015; Kim et al. 2019; Matsumoto & Tsumura 2019); nevertheless inferring the absolute intensity of the EBL from fluctuation measurements depends on model assumptions. Another limitation of absolute photometry and fluctuation measurements is the inferred EBL redshift resolution. As the measured intensity is a projection of all emission along the line of sight, one way to infer the redshift dependency of the EBL is to measure the opacity of gamma-ray photons from individual blazars. Near-infrared photons along the line of sight will interact with gamma-ray photons by pair production, and thus the redshift composition of EBL can be constrained by the absorption features in blazar spectra (e.g., Aharonian et al. 2006, 2007; MAGIC Collaboration et al. 2008; Abdo et al. 2010; Ackermann et al. 2012; H. E. S. S. Collaboration et al. 2017; Ackermann et al. 2018; Abeysekara et al. 2019; Acciari et al. 2019; Abdalla et al. 2020). This method enables a redshift tomography of the EBL by observing blazars from different distances. However, these estimates have low spectral resolution and depend on the assumption of intrinsic blazar spectra.","Citation Text":["Matsumoto et al. 2015"],"Citation Start End":[[694,715]]}
{"Identifier":"2020MNRAS.493.3045B__Jaisawal_&_Naik_2015a_Instance_2","Paragraph":"We have used 3.0\u201375.0 keV NuSTAR data to probe any cyclotron line feature. To describe the continuum of 4U 1700\u201337, we have applied the NPEX model [cons*TBpcf*(powerlaw*npex+gaus+gaus)], following the previous work of Jaisawal & Naik (2015a). The NPEX model has been created by adding two cutoffpl models with their cutoff energies tied to each other and keeping the photon index of one to be frozen at \u20132.0. For the best fit, the \u03c72\/d.o.f is found to be 577.64\/475. The fit shows some residuals in the overall spectrum. We added a Gaussian absorption model around 39 keV (following the previous work of Jaisawal & Naik 2015a), but the best fit gives the line energy as $15.44^{+0.56}_{-0.53}$ keV with \u03c72\/d.o.f = 487.32\/472. The width and the depth of the line are found to be $5.47^{+0.90}_{-0.78}$ keV and $1.29^{+0.51}_{-0.39}$, respectively. The chance probability of the line has been computed using the ftest task in xspec. The F-test with this absorption line gives an F value  = 29.2 and a chance probability of 2.61 \u00d7 10\u221217(Table 3). If we add another Gaussian absorption line at 38.9 keV (Energy value frozen) the best-fitting \u03c72\/d.o.f is found to be 486.8\/470. This indicates that the second absorption line is not required for the fit. If we use only one Gaussian absorption line and freeze the line energy at 38.9 keV then the width and the depth of the line are found to be $4.54_{-0.87}^{+0.90}$ keV and $2.87_{-1.05}^{+1.24}$, respectively with a \u03c72\/d.o.f  = 549.2\/473. The ftest gives a chance probability of the 38.9 keV line to be 6.4 \u00d7 10\u22126. So, with NPEX model we find two valid model combinations of the data. One, the presence of a Gaussian absorption line at \u223c15 keV, two, the presence of a Gaussian absorption line at 38.9 keV. But, the presence of both lines together is not supported by the data. The 10.0\u201370.0 keV flux of the source is found to be (2.26 \u00b1 0.01) \u00d7 10\u22129 erg cm-2s-1, much lower than the value (5.6 \u00b1 0.3) \u00d7 10\u22129 erg cm-2s-1, previously reported from SUZAKU data (Jaisawal & Naik 2015a).","Citation Text":["Jaisawal & Naik 2015a"],"Citation Start End":[[604,625]]}
{"Identifier":"2022MNRAS.516.3900A__Cazaux_et_al._2022_Instance_2","Paragraph":"Sudden outbursts of NH3 simultaneously with H2S detected with the ROSINA-DFMS instrument on the Rosetta S\/C point to the presence of abundant ammonium hydrosulphide in or on carbonaceous grains from comet 67P\/Churyumov-Gerasimenko. There seems to be a clear distinction between the nucleus ice, where H2S and NH3 exist independently and grains, where they desorb together. S2 is much more abundant on grains compared to water than in the ice of the comet, while S3 is found only in grain impacts. This higher abundance points to radiolysis in these grains, which means they must have been exposed to energetic particles over an extended time. While for operational reasons, S4 could not be measured close to the dust impacts, S4 was clearly identified in periods where the coma was very dusty (Calmonte et al. 2016). Longer sulphur chains very likely are refractory, not sublimating at temperatures reached in the instrument or on grains in the coma. While Sn can also be formed from pure H2S ice by photo processing (Cazaux et al. 2022), the fact that S3 is clearly related to dust and is not found in the normal nucleus ice, where H2S is quite abundant, indicates that S3 is a product of radiolysis of the ammonium salt. In addition, photo processing of H2S results not only in Sn, but also in H2S2 (Cazaux et al. 2022), a species not detected in the DFMS m\/z 66 and m\/z 65 (HS2) spectra. This exposure rules out a contemporary formation of the salt on the surface or in the interior of the comet or a formation of the salt in the mid-plane of the protoplanetary disc, while the comet accreted. A pre-stellar formation is therefore likely. The salt is semivolatile, less volatile than water and could probably have survived quite high temperatures. It seems that on these grains, acids and ammonia are all locked in salts, be it sulphur, halogens, or carbon bearing acids like HOCN. If indeed, a relatively large part of sulphur and nitrogen is therefore in a semivolatile state in these grains, then the depletion of nitrogen in comets and of sulphur in star-forming regions could probably be explained, primarily because salts escape detection unless they experience temperatures above water sublimation. With the JWST S\/C in orbit, there is hopefully the possibility to detect salts, or at least several of the acids in ices, which are supposed to be part of ammonium salt, like HOCN, H2CO, and formamide while looking for ammonium salts in star-forming regions and possibly comets.","Citation Text":["Cazaux et al. 2022"],"Citation Start End":[[1302,1320]]}
{"Identifier":"2016ApJ...829..131A__Reid_1964_Instance_1","Paragraph":"Solar energetic particle (SEP) events observed in interplanetary space present a large variability in terms of intensity, composition, and spatial and temporal extent (e.g., Kahler et al. 1999; von Rosenvinge & Cane 2006; Gopalswamy 2012). In particular, the observed SEP intensity time profiles result from both (i) the injection history of SEPs as they are released from their acceleration sites, and (ii) the transport processes undergone by the particles as they travel through the interplanetary medium from their sources to the particle detectors on board spacecraft. The transport effects pose the main difficulty in reconstructing the processes of particle release occurring at the Sun from in situ spacecraft observations. In order to correctly extract interplanetary transport effects, an effective treatment of the SEP transport processes in the heliosphere is required. Several approaches have been used over the years to characterize the SEP transport conditions (e.g., Reid 1964; Earl 1974, 1976; Roelof 1975, 1979, 2008; Hamilton 1977; Beeck et al. 1987; Kallenrode et al. 1992; Ruffolo 1995; Kocharov et al. 1998; Ruffolo et al. 1998; Laitinen et al. 2000, 2013; Dr\u00f6ge 2003; Qin et al. 2005; Wang et al. 2006; Agueda et al. 2008, and references therein). Such approaches allow us to reproduce spacecraft SEP observations and thereby infer SEP injection histories. The particle angular distributions relative to the local direction of the magnetic field (i.e., pitch-angle distributions (PADs)) have proven to be relevant in these analyses. A given intensity time profile can be fitted using a large variety of combinations of the injection profile at the Sun and interplanetary transport conditions, although this non-uniqueness can be constrained by a simultaneous analysis of the omni-directional intensity and the computed anisotropy time profiles (e.g., Schulze et al. 1977). More recently this problem has been further constrained by including the most direct form of data in the analysis, that is, the measured directional distributions together with the angular response of the detector (Agueda et al. 2008, 2009a, 2014).","Citation Text":["Reid 1964"],"Citation Start End":[[983,992]]}
{"Identifier":"2018AandA...619A..13V__Saviane_et_al._2012_Instance_2","Paragraph":"The EWs were measured with the methods described in V\u00e1squez et al. (2015). As in Paper I, we used the sum of the EWs of the two strongest CaT lines (\u03bb8542, \u03bb8662) as a metallicity estimator, following the Ca\u202fII triplet method of Armandroff & Da Costa (1991). Different functions have been tested in the literature to measure the EWs of the CaT lines, depending on the metallicity regime. For metal-poor stars ([Fe\/H] \u2272 \u22120.7 dex) a Gaussian function was used with excellent results (Armandroff & Da Costa 1991), while a more general function (such as a Moffat function or the sum of a Gaussian and Lorentzian, G + L) is needed to fit the strong wings of the CaT lines observed in metal-rich stars (Rutledge et al. 1997; Cole et al. 2004). Following our previous work (Gullieuszik et al. 2009; Saviane et al. 2012) we have adopted here a G+L profile fit. To measure the EWs, each spectrum was normalised with a low-order polynomial using the IRAF continuum task, and set to the rest frame by correcting for the observed radial velocity. The two strongest CaT lines were fitted by a G+L profile using a non-linear least squares fitting routine part of the R programming language. Five clusters from the sample of Saviane et al. (2012) covering a wide metallicity range were re-reduced and analysed with our code to ensure that our EWs measurements are on the same scale as the template clusters used to define the metallicity calibration. Figure 5 shows the comparison between our EWs measurements and the line strengths measured by Saviane et al. (2012) (in both cases the sum of the two strongest lines) for the five calibration clusters. The observed scatter is consistent with the internal errors of the EW measurements, computed as in V\u00e1squez et al. (2015). The measurements show a small deviation from the unity relation, which is more evident at low metallicity. A linear fit to this trend gives the relation: \u03a3EW(S12) = 0.97 \u03a3EW(this work) + 0.21, with an rms about the fit of 0.13 \u00c5. This fit is shown in Fig. 5 as a dashed black line. For internal consistency, all EWs in this work have been adjusted to the measurement scale of Saviane et al. (2012) by using this relation. In Table 3 we provide the coordinates, radial velocities, and the sum of the equivalent widths for the cluster member stars, both measured (\u201cm\u201d) and corrected (\u201cc\u201d) to the system of Saviane et al. 2012.","Citation Text":["Saviane et al. (2012)"],"Citation Start End":[[1210,1231]]}
{"Identifier":"2018ApJ...860...59K___2017b_Instance_1","Paragraph":"The total energy released at an SN type II explosion, \n\n\n\n\n\n, exceeds, by two orders of magnitude, the energy transferred to the ejecta, \n\n\n\n\n\n. The bulk of energy escapes with the neutrino emission. Typically, the pulsar rotation energy is small compared to the ejecta energy. However, it cannot be excluded that at some circumstances the pulsar can receive a larger fraction of the explosion energy. When it becomes comparable to the kinetic energy of the ejecta, the sizes of both the SNR and the PWN can significantly deviate from Sedov\u2019s solution. To illustrate the implication of this scenario, we considered a specific case of HESS J1825\u2212137, a very extended and bright in VHE gamma-ray PWN (Aharonian et al. 2006; Pavlov et al. 2008; Mitchell et al. 2017a, 2017b; H.E.S.S. Collaboration et al. 2018a). The X-ray and gamma-ray observations of HESS J1825\u2212137 show that the shape of the nebula deviates from the spherically symmetric geometry. Thus, the 1D model described in Section 2 has limited capability for accurate quantitative predictions. However, in the framework of the suggested scenario, the hydrodynamic processes evolve differently from the conventional situations, and the 1D model should provide a correct estimate for the energy required for the PWN\/SNR inflation. In the framework of the scenario, the pressure in the PWN is the main driving force of the SNR expansion. Independently on symmetry of the system and dimension of the used model, the PWN is nearly isobaric, thus the radius of the pulsar wind TS provides an observational constraint for the pressure in the nebula. Therefore, we adopt the following requirement: the present-day pressure in the nebula should be consistent with the pulsar wind TS located at \n\n\n\n\n\n from the pulsar (Van Etten & Romani 2011). The 1D model allows us to obtain the radii of the SNR, PWN, and pulsar wind TS as functions of three key model parameters: density of the ISM, pulsar braking index, and the initial angular velocity of the pulsar. In the framework of this model, it is possible to reproduce the three key properties of the system: \n\n\n\n\n\n, \n\n\n\n\n\n, and \n\n\n\n\n\n, provided that (i) the pulsar obtains, at its birth, a significant fraction of the explosion energy (\n\n\n\n\n\n), and (ii) the braking index is small, \n\n\n\n\n\n. The conventional Sedov-like solution implies a low density of the ISM, \n\n\n\n\n\n. In contrast, the model proposed in this paper requires a rather dense environment, \n\n\n\n\n\n. This agrees with the presence of dense molecular clouds reported in the vicinity of the source (Voisin et al. 2016). What concerns the requirement for the pulsar braking index, \n\n\n\n\n\n, is that it is small compared to the value measured for the Crab pulsar (Lyne et al. 1993), but it matches both the old (e.g., Vela; n = 1.4 \u00b1 0.2; Lyne et al. 1996) and young (e.g., J1833\u20131034; n = 1.8569 \u00b1 0.001; Roy et al. 2012) pulsars.","Citation Text":["Mitchell et al.","2017b"],"Citation Start End":[[742,757],[765,770]]}
{"Identifier":"2020ApJ...902..116M__Palumbo_et_al._2008_Instance_1","Paragraph":"Three additional products were detected using the QMS during the TPD of the processed ice analogs in Experiments 1\u22123: C2O,3\n\n3\nNo other carbon chain oxide molecules were monitored with the QMS.\n H2C2O, and H15NCO. The TPD curves of these product main mass fragments (m\/z = 40, C2O; m\/z = 14, H2C2O; and m\/z = 44, H15NCO4\n\n4\nWe assign the thermal desorption detected for the m\/z = 44 mass fragment at T > 100 K to H15NCO because it is the most stable isomer among HOC15N, H15NCO, HC15NO, and HO15NC. In previous experimental simulations where HNCO ice molecules were produced, the metastable isomers were only tentatively detected, if at all (see, e.g., Jim\u00e9nez-Escobar et al. 2014). In addition, gas-phase observations of different astrophysical environments reveal an abundance of the metastable isomers with respect to HNCO of less than 1% (Quan et al. 2010 and references therein).\n in Experiment 3 are also shown in Figure 3. The C2O IR feature at 1989 cm\u22121 (Palumbo et al. 2008) was only tentatively detected at the edge of the hatched region in the top panel of Figure 1. The H2C2O and H15NCO IR features at 2129 cm\u22121 and 2260 cm\u22121, respectively (van Broekhuizen et al. 2004; Hudson & Ferrante 2020), overlapped with different CO and carbon chain oxide IR features in the hatched region and could not be unambiguously detected. In order to confirm the assignments of the m\/z = 40, m\/z = 14, and m\/z = 44 desorptions in Figure 3, to C2O, H2C2O, and H15NCO (the thermal desorption detected at T > 100 K), respectively, we energetically processed four additional ice analogs composed of different combinations of isotopically labeled H2, CO, and N2 molecules with 2 keV electrons (Experiments 4\u22127; see Table 1),  and checked that the same thermal desorption was observed for the expected mass fragments, considering the shift in the mass of the main fragments according to their isotopic composition in each experiment. The results are shown in Figure 4, confirming the formation of these species upon electron irradiation of H2:CO:N2 ice analogs.","Citation Text":["Palumbo et al. 2008"],"Citation Start End":[[963,982]]}
{"Identifier":"2021MNRAS.503..324M__Zhao_et_al._2019_Instance_2","Paragraph":"We first determined the orbital parameters for RS Ser, V449 Per, and V1095 Her. Further, we updated the parameters for V593 Cen and MR Del. Using the formula f  = (\u03a9in \u2013 \u03a91)\/(\u03a9in \u2013 \u03a9out), we calculated the contact factors f for RS Ser, V593 Cen, and V1095 Her as 6.5 per\u2009cent, 40 per\u2009cent, and 53 per\u2009cent, respectively. RS Ser is a contact binary with a small temperature difference of 131 K and a low contact factor. For V593 Cen, we updated the orbital parameters using more complete light curves. The orbital inclination of 83\u00b0.18 is similar to the result (82\u00b0.6) obtained by Zhao et al. (2019). The temperature of the secondary component (15 284 K) is higher than the previous result of 15 099 K. However, the mass ratio of 0.6 is lower than the previous result of 1.05 (Zhao et al. 2019). More spectroscopic observations are required to confirm the mass ratio. We confirmed that V593 Cen is an early-type contact binary with a deep contact factor as well as a black hole candidate. For MR Del, we revised the absolute parameters using its full light curve and the published radial velocities, which are similar to those published previously (Zhao et al. 2019; Pribulla et al. 2009; Djura\u0161evi\u0107 et al. 2011). V1095 Her is also a contact binary with a deep contact factor of 40 per\u2009cent and a temperature difference of about 172 K. Looking over our four complete light curves, we found no evident starspot activity, and estimated variations exist over a long-term time-scale of years. V449 Per is an interesting target for detecting extra-solar and brown dwarfs using the minima timing variability of a low-mass eclipsing binary (Pribulla et al. 2012). Additional minima with higher precision are required to study its periodic variation further. Our physical parameters for RS Ser, V593 Cen, and V1095 Her are based on a light curve with a q-search-determined mass ratio. The nature of these parameters is speculative and preliminary. Radial velocities may eventually come to rescue them and provide a more definitive determination.","Citation Text":["Zhao et al. 2019"],"Citation Start End":[[776,792]]}
{"Identifier":"2016MNRAS.462.1415C__Kochiashvili_et_al._2015_Instance_1","Paragraph":"Two major limitations, often neglected, affect this type of analysis: the adoption of oversimplified models to describe the wide variety of observed galaxy SEDs and the presence of \u2018systematic\u2019 model uncertainties. This second limitation has been addressed in several studies already (e.g. Charlot, Worthey & Bressan 1996; Cervi\u00f1o, Luridiana & Castander 2000; Conroy, Gunn & White 2009; Percival & Salaris 2009; Conroy & Gunn 2010; Conroy, White & Gunn 2010a). The difficulty of precisely quantifying systematic model uncertainties has led to mainly qualitative conclusions, leaving the problem unsolved. The first limitation is easier to tackle, for example, by using more physically realistic models of galaxy SEDs and combining these with advanced statistical techniques to extract physical constraints from data. This appears as the most promising route to fully exploit the information gathered by modern photometric and spectroscopic galaxy surveys. Yet, the several tools proposed so far to interpret galaxy SEDs in terms of physical parameters do not allow one to fully exploit the high quality of modern data. For example, most existing approaches rely on the adoption of a rigid physical model (e.g. analytic, two-parameter star formation histories combined with a standard dust attenuation curve and the assumption that all stars in a galaxy have the same metallicity) to describe galaxy SEDs (e.g. Bolzonella et al. 2010; Wuyts et al. 2011; Hern\u00e1n-Caballero et al. 2013; Ilbert et al. 2013; Bauer et al. 2013; Muzzin et al. 2013; Lundgren et al. 2014; Kochiashvili et al. 2015; Mortlock et al. 2015; Kawinwanichakij et al. 2016). Even with the inclusion of superimposed bursts of star formation (e.g. Kauffmann et al. 2003; Gallazzi et al. 2005; Pozzetti et al. 2007; Gallazzi & Bell 2009; da Cunha et al. 2010), this does not allow a physically consistent description of the contributions by stars, gas and dust to the integrated emission from a galaxy, nor the inclusion of a potential AGN component (a notable exception is the approach of Pacifici et al. 2012, who incorporate star formation and chemical enrichment histories from numerical simulations of galaxy formation and emission from photoionized gas). Also, current spectral analysis tools are generally optimized to interpret either photometric or spectroscopic observations of galaxies, but not arbitrary combinations thereof. Finally, most existing tools suffer from additional limitations: many focus on the selection of \u2018best-fitting\u2019 parameters rather than on the uncertainties associated with these parameters (e.g. chi-square minimization techniques; Arnouts et al. 1999; Bolzonella, Miralles & Pell\u00f3 2000; Kriek et al. 2009); when this is not the case, the number of free parameters that can be explored is generally limited (e.g. with grid-based Bayesian techniques; da Cunha, Charlot & Elbaz 2008; Noll et al. 2009; Pacifici et al. 2012); and when more sophisticated (e.g. Markov Chain Monte Carlo, hereafter MCMC) techniques allow the exploration of more parameters, instrumental effects are generally not incorporated in the analysis (e.g. Acquaviva, Gawiser & Guaita 2011; Serra et al. 2011; Han & Han 2014).","Citation Text":["Kochiashvili et al. 2015"],"Citation Start End":[[1564,1588]]}
{"Identifier":"2021MNRAS.506.1258W__Borsa_et_al._2021_Instance_2","Paragraph":"Situated in the closest vicinity of their host stars (0.05\u2009AU) and having no counterparts in our Solar system, ultra-hot Jupiters (Arcangeli et al. 2018; Bell & Cowan 2018; Parmentier et al. 2018) are ideal testbeds for studying the impact of 3D effects on high-resolution spectra. There are two important reasons for this. First, ultra-hot Jupiters are accessible objects to observe. Their short orbital periods (1\u20132\u2009d) and hot, extended atmospheres make them perfect targets for transmission spectroscopy (Hoeijmakers et al. 2019; Von Essen et al. 2019; Ehrenreich et al. 2020; Borsa et al. 2021), emission spectroscopy (Evans et al. 2017; Arcangeli et al. 2018; Mikal-Evans et al. 2020) and phase-curve studies (Zhang et al. 2018; Bourrier et al. 2020b; Mansfield et al. 2020). Secondly, ultra-hot Jupiters display extreme variations across their atmospheres, because they are expected to become tidally locked soon after their formation (Rasio et al. 1996; Showman & Guillot 2002). As a result, their atmospheres virtually consist two different worlds: a permanently irradiated dayside and a permanently dark nightside. The scorching, cloud-free dayside (T \u2273 2500\u2009K) nearly resembles a stellar photosphere, where most molecules are dissociated1 and metals become ionized (Parmentier et al. 2018; Hoeijmakers et al. 2019). On the other hand, the nightside is substantially cooler (T \u2272 1000\u2009K) and may even serve as a stage for cloud formation (Helling et al. 2019; Ehrenreich et al. 2020). Ultra-hot Jupiters also exhibit large differences in their thermal structures: the dayside is expected to show strong thermal inversions (Haynes et al. 2015; Evans et al. 2017; Kreidberg et al. 2018; Pino et al. 2020; Yan et al. 2020), whereas nightside temperatures are expected to monotonically decrease with altitude. Furthermore, ultra-hot Jupiters feature strong winds in the order of 1\u201310\u2009km\u2009s\u22121 (Tan & Komacek 2019), which arise as a result of the continuous day-night forcing. Many observational studies have measured Doppler shifts due to winds on ultra-hot Jupiters (Casasayas-Barris et al. 2019; Bourrier et al. 2020a; Cabot et al. 2020; Ehrenreich et al. 2020; Gibson et al. 2020; Hoeijmakers et al. 2020; Nugroho et al. 2020; Stangret et al. 2020; Borsa et al. 2021; Kesseli & Snellen 2021; Rainer et al. 2021; Tabernero et al. 2021), yet inferring the underlying 3D circulation pattern is a formidable challenge.","Citation Text":["Borsa et al. 2021"],"Citation Start End":[[2254,2271]]}
{"Identifier":"2021MNRAS.503.5179N__Blanton_et_al._2004_Instance_1","Paragraph":"Here, we report on molecular gas observations of NGC\u20090708, the BCG in the low-mass galaxy cluster Abell\u2009262, itself part of the Perseus\u2013Pisces galaxy supercluster. NGC\u20090708 lies 58.3 \u00b1 5.4\u2009Mpc away (estimated using infrared surface brightness fluctuations; Jensen et al. 2003). It is a giant elliptical galaxy with a weak dust lane (Ebneter & Balick 1985; Wegner et al. 1996) and an effective radius of 33 arcsec ($\\approx \\, 9.3$\u2009kpc; Wegner et al. 2012). See Fig. 1 for an HST image of NGC\u20090708. Abell\u2009262 was identified as having an X-ray emitting ICM by Jones & Forman (1984), and Stewart et al. (1984) measured the cooling time to be 1.3 \u00d7 109\u2009yr, smaller than the age of the Universe so that the cluster is expected to form a cooling flow. The 20-cm observations of Parma et al. (1986) revealed a double-lobed, \u2018S\u2019-shaped jet and led to the classification of NGC\u20090708 as a weak Fanaroff\u2013Riley Class I radio source (Blanton et al. 2004). The top panel of Fig. 1 also has 330\u2009MHz continuum observations from Clarke et al. (2009) overlaid (blue contours) to show the shape and orientation of the large-scale jet. Analysis of Chandra observations revealed a hole or bubble within the ICM, cospatial with the eastern lobe of the jet (Blanton et al. 2004). Clarke et al. (2009) found additional 3\u20136 kpc radius cavities at differing position angles within the X-ray gas, and at a range of radial distances from the BCG (8\u201329 kpc), indicating multiple episodes of AGN activity from a precessing SMBH jet. They concluded that the total AGN emission should be capable of counteracting the cooling flow over several outbursts. Using their multifrequency observations of NGC\u20090708, Clarke et al. (2009) also calculated the radio spectral index (\u03b1) from 235 to 610\u2009MHz, finding the spectrum to be flat in the core (\u03b1 = \u22120.5), typical of new particles in a jet. They also estimated a lower limit on average outburst repetition time-scales in Abell\u2009262 to be \u03c4rep \u2265 28\u2009Myr.","Citation Text":["Blanton et al. 2004"],"Citation Start End":[[921,940]]}
{"Identifier":"2022MNRAS.510.4962G__Komissarov_&_Barkov_2009_Instance_1","Paragraph":"For the launching of a two-sided jet, the BZ power depends on the BH spin (assumed to be aligned with the disc) and on the magnetic flux threading the BH horizon at rh (e.g. Tchekhovskoy, Narayan & McKinney 2011):1(1)$$\\begin{eqnarray*}\r\nL \\approx \\frac{10^{-3}}{c} \\Phi _h^2\\Omega _h^2f(\\Omega _h)\\approx 10^{51} M_{\\, {\\rm BH},5}^2 B_{h,15}^2a_{-0.1}^2 \\, {\\rm erg}~\\, {\\rm s}^{-1},\r\n\\end{eqnarray*}$$where Qx denotes the value of the quantity Q in units of 10x times its c.g.s. units, except for Mx, which is given in units of $\\, {\\rm M_{\\odot }}$. $\\Phi _h = 4\\pi r_h^2|B_h|$ is the integrated magnetic flux on one hemisphere of the BH horizon, |Bh| is the value of the radial contravariant magnetic field on the horizon, $\\, {M_{\\rm BH}}$ and a are the BH mass and spin, \u03a9h = ac\/(2rh) is the angular velocity at the BH horizon, and f(\u03a9h) \u2248 1 + 1.38x2 \u2212 9.2x4, where $x \\equiv 0.5a\\left(1+\\sqrt{1-a^2}\\right)^{-1}$. A second necessary condition for a successful launching of a BZ jet is that the BZ jet power will be sufficient to overcome the accretion power of the infalling material along the jet path (Burrows et al. 2007; Komissarov & Barkov 2009). Prior to the jet launching, matter is free-falling on the BH quasi-spherically with a power on one hemisphere of \n(2)$$\\begin{eqnarray*}\r\n\\dot{M}c^2 = 4\\pi r_g^2 \\rho _h \\beta ^r_h c^3\\approx 2\\times 10^{51} M_{\\, {\\rm BH},5}^2\\beta ^r_h\\rho _{h,7} \\, {\\rm erg}~\\, {\\rm s}^{-1},\r\n\\end{eqnarray*}$$where \u03c1h is the mass density on the horizon at the time of the jet launching, and $\\beta ^r_h$ is the dimensionless radial velocity of the infalling material at that time, which is expected to be close to unity. Comparing the jet luminosity from equation (1) with the accretion power from equation (2), we get a necessary condition for the strength of magnetic flux on the horizon that allows a successful jet launching \n(3)$$\\begin{eqnarray*}\r\n\\Phi _h \\gtrsim \\, {\\Phi _{h,\\rm min}}\\approx 7\\times 10^{27}\\frac{\\sqrt{\\beta ^r_h\\rho _{h,7}}}{a_{-0.1}}~{\\rm G~{\\rm cm^2}},\r\n\\end{eqnarray*}$$corresponds to minimal magnetic field on the horizon $\\, {B_{h,\\rm min}}\\approx 1.4\\times 10^{15}$ G for the same normalization of parameters. Note that since a and $\\beta ^r_h$ are both of order unity, the minimal magnetic field is primarily dictated by the central density of the star, which in turn depends on the stellar mass, radius, and density profile.","Citation Text":["Komissarov & Barkov 2009"],"Citation Start End":[[1132,1156]]}
{"Identifier":"2019MNRAS.484..892R__Wolf_et_al._2009_Instance_2","Paragraph":"In the left-hand panel of Fig. 2, we compare the morphological types assigned by the STAGES collaboration for the galaxies in the whole OMEGA sample and the jellyfish candidates sample. The sample of jellyfish galaxies (JC345) is composed mainly of late-type spirals and irregulars. In the middle panel of Fig. 2, we show the distribution of SED types for both samples. Based on the SED types of the galaxies, of the 70 jellyfish galaxy candidates analysed, 66 were found to be part of the blue cloud and 4 as being dusty reds (IDs: 11633, 17155, 19108, 30604). However, contrary to what could be expected, dusty red galaxies are only a small portion of our sample of jellyfish candidates. One reason why we may not detect many dusty reds as jellyfish galaxies might be due to the fact that these galaxies, despite having relatively high SFRs (only four times lower than that in blue spirals at fixed mass, Wolf et al. 2009), have significant levels of obscuration by dust that might hamper the identification of the jellyfish signatures. Another reason for that is that we selected jellyfish galaxy candidates within a parent sample of H\u2009\u03b1-emitting galaxies that already had a low fraction of dusty red galaxies (\u224815 ${{\\ \\rm per\\ cent}}$). As these galaxies have low star formation, it is harder to perceive the morphological features of RPS. Dusty red galaxies have been previously studied in this same system (Wolf et al. 2009), and RPS was suggested to be the main mechanism acting in these galaxies (B\u00f6sch et al. 2013). While in B\u00f6sch et al. (2013) one of the main pieces of evidence suggesting the action of enhanced RPS were the existence of disturbed kinematics without disturbed morphologies, in our study we strongly base our selection on such morphological distortions. Both our jellyfish galaxy candidates and the dusty red galaxies show different characteristics that can be correlated to the effect of RPS. Nevertheless, they might be tracing different stages of the same phenomenon, where dusty red galaxies have more regular morphologies, but disturbed kinematics. Our sample of morphologically disturbed jellyfish galaxy candidates may be showing the stage where the features of RPS are the most visible and the SFRs are enhanced.","Citation Text":["Wolf et al. 2009"],"Citation Start End":[[1414,1430]]}
{"Identifier":"2022AandA...658A..78S__Seifried_et_al._2012_Instance_1","Paragraph":"Molecular outflows are a common and essential component in the formation process of low- and high-mass stars. In the past 40 yr, astronomers have mapped outflows in the whole mass range of young stellar objects (YSOs; e.g., Frank et al. 2014; Bally 2016; Anglada et al. 2018; Ray & Ferreira 2021). Magnetohydrodynamical (MHD) simulations have shown that the magnetic field plays a crucial role in the launching of molecular outflows (e.g., Pudritz & Ray 2019), more significantly so in the case of massive YSOs (e.g., Matsushita et al. 2018). Here, for instance, the presence of a magnetic field leads to the formation of early outflows. These reduce the radiation pressure, which allows the protostar mass to grow further (Banerjee & Pudritz 2007; Rosen & Krumholz 2020). In addition, the intensity of the magnetic field may influence the collimation of the outflows in massive YSOs. The outflows are well collimated for weak fields and poorly collimated for strong fields (Hennebelle et al. 2011; Seifried et al. 2012). In case of strong magnetic fields, the structure of the outflows is determined by the large-scale geometry of the magnetic field lines (Matsushita et al. 2017). Recently, Machida & Hosokawa (2020) have found a strong dependence of the evolution of outflows in massive YSOs on the initial magnetic field strength of the prestellar cloud for different accretion rates. In their 3D MHD simulations, they grouped the results into three categories: successful outflows, failed outflows, and delayed outflows. In the successful outflows, the outflows appear only when the prestellar cloud is strongly magnetized (\u03bc1 = 2,3), and after an evolution time of ~104 yr, they reach a distance from the central protostar of about 104 au. When the magnetic field is weak (\u03bc \u2265 5), we have failed and delayed outflows; even though small outflows of about 100\u20131000 au are observed in both cases, only in delayed outflows they can overcome the ram pressure and can ultimately grow. In a massive YSO, a large molecular outflow is therefore formed only if the initial magnetic field strength is B0 \u2273B0,cr = 10\u22124(Mcl\u2215100M\u2299) G, where Mcl is the cloud mass (Machida & Hosokawa 2020).","Citation Text":["Seifried et al. 2012"],"Citation Start End":[[999,1019]]}
{"Identifier":"2017MNRAS.464..183N__Biviano_&_Katgert_2004_Instance_2","Paragraph":"Other important result we reported in Section 3.3 is the reversing behaviour of red and blue galaxies with respect to velocity and groupcentric distances segregation, with redshift. Regarding velocity segregation, the preceding paragraph provides a qualitative scenario. Now, to explain the spatial segregation, we should notice that our analyses in Sections 3.2 and 3.3 take into account galaxies within 2R\/R200. One can reasonably assume that such objects at lower redshifts correspond to a mixture of descendants of galaxies at higher redshifts in the same radii and of infalling objects from outer radii. Thus, both survival and replenishment of galaxies should be expected over the time, and two important factors come into play: (i) the accretion rate of galaxies; and (ii) the orbital dependence of galaxy properties (e.g. Biviano & Katgert 2004; Iannuzzi & Dolag 2012). Indeed, regarding velocity segregation, it has also been interpreted as red and blue galaxies having different kinds of orbits, with the orbits of blue galaxies being more anisotropic than the red ones (e.g. Biviano & Katgert 2004). Recently, Biviano et al. (2016) verified that the anisotropy profile of z \u223c 1 clusters is nearly isotropic near the cluster centre, and increasingly elongated with radius. This result is consistent with a halo evolution through an initial phase of fast collapse and a subsequent slow phase of inside-out growth by accrection of field material (e.g. Lapi & Cavaliere 2009). Since the accretion rate of galaxies from the field is higher at higher redshifts (e.g. McGee et al. 2009), our sample at z \u223c 0.8 is expected to be more affected by recent infalls, which had less time to go deeper into the group potential. This could explain the development of a more marked difference between the mean groupcentric distance of red and blue galaxies (see Fig. 12). After \u223c3 Gyr, part of these infalling galaxies may reach the R  2R200 region, at z \u223c 0.4, mixing with virialized and backsplash objects, and thus presenting a less pronounced radial segregation between red and blue galaxies.","Citation Text":["Biviano & Katgert 2004"],"Citation Start End":[[1086,1108]]}
{"Identifier":"2016ApJ...831...74R__Dotter_et_al._2008_Instance_1","Paragraph":"In order to determine the fundamental parameters of DM Ori, we take a similar approach to our V409 Tau study (Rodriguez et al. 2015), namely, to infer stellar parameters using stellar evolution isochrone models. However, unlike the V409 Tau analysis, we are very limited in the available data for DM Ori. We are only able to use the temperature of \n\n\n\n\n\n = 6740 \u00b1 250 K determined from the fits to the SED, as well as the star\u2019s B- and V-band magnitudes (See Table 1). The longer wavelength photometric data for DM Ori are not useful due to contamination at these wavelengths from the disk. Assuming DM Ori is associated with the OB1c or the Orion Nebula Cluster portion of the Orion association (Bally 2008) allows us to place priors on some of the parameters. We assume uniform priors for distance (250\u2013550 pc), metallicity (\n\n\n\n\n\n dex), and age (0\u20136 Myr). We also assume an extinction of \n\n\n\n\n\n based on our best-fit SED. We use both the Dartmouth Stellar Evolution Models (Dotter et al. 2008) and the MIST stellar evolution models (Choi et al. 2016) combined with a Markov-Chain Monte Carlo (MCMC) to derive posterior distributions for the mass, age, radius, and distance of DM Ori (see Figure 4). From this analysis, we determine the most probable parameter values based on the posterior medians with errors from their 68% confidence interval uncertainties. For the Dartmouth models, we find a mass of \n\n\n\n\n\n M\u2299, an age of \n\n\n\n\n\n Myr, a radius of \n\n\n\n\n\n R\u2299, and a distance of \n\n\n\n\n\n pc. Similarly, for the MIST models, we find a mass of \n\n\n\n\n\n M\u2299, an age of \n\n\n\n\n\n Myr, a radius of \n\n\n\n\n\n R\u2299, and a distance of \n\n\n\n\n\n pc. Both sets of stellar evolution models produce consistent parameters for DM Ori, and the inferred age and distance are consistent with it being associated with the Orion star-forming region. To rule out the possibility that DM Ori is actually a foreground star in our line of sight to the Orion association, we reran our analysis with distance prior of 0\u2013550 pc. Our results are very similar to the distribution shown in Figure 4, providing strong evidence that DM Ori is not a foreground star with a distance  250 pc.","Citation Text":["Dotter et al. 2008"],"Citation Start End":[[977,995]]}
{"Identifier":"2016ApJ...823...50S__Ho_et_al._1993_Instance_1","Paragraph":"For normal single AGNs, the bolometric luminosity, \n\n\n\n\n\n, can usually be estimated from either optical or X-ray luminosities. However, for our targets, the hard X-ray luminosities are much lower than for single AGNs with the same [O iii] luminosities (Paper I). Three possible reasons are: (1) the observed X-rays are low due to high gas absorption (e.g., Bassani et al. 1999); (2) the X-rays are intrinsically weak due to a high accretion state that radiates little coronal emission (e.g., Desroches et al. 2009; Dong et al. 2012); (3) there are significant [O iii] luminosity excesses due to shocks in our galaxy mergers (e.g., Dopita & Sutherland 1995). The emission-line diagnostic ratios featuring [O iii]\/H\u03b2 and [N ii]\/H\u03b1 show that our targets are well in the regime of Seyferts, with only one displaying relatively large uncertainty (Liu et al. 2010a). It is unlikely that the emission lines of our AGNs are mainly powered by shocks, as they do not lie in the region of LINERs on the diagnostic diagrams (Ho et al. 1993). It is more likely that our targets are highly obscured rather than in a high accretion state, since the estimated Eddington ratios (Table 4; see discussion below) are not close to unity. Additionally, our targets are all detected by the Wide-field Infrared Survey Explorer (Wright et al. 2010) with \u03bbL\u03bb(12 \u03bcm) \u223c 1044 erg s\u22121. The X-ray\u2013MIR ratios of our target binary AGNs (\u22720.01) are much smaller than those of normal AGNs (\u223c0.3; Horst et al. 2008), indicating that the X-ray emission of the AGNs is absorbed and reprocessed into IR emission. As shown by Heckman et al. (2005) at low redshift, selection by narrow optical emission lines will recover most AGNs selected by hard X-rays (with the exception of BL Lac objects). On the other hand, selection by hard X-rays misses a significant fraction of the local AGN population with strong emission lines. In view of the uncertainties associated with X-ray absorption in gas-rich mergers, we estimate the bolometric luminosities of the binary AGNs from their extinction-corrected [O iii] luminosity (Paper I),\n3\n\n\n\n\n\nwith intrinsic scatter 0.4 dex (see Trump et al. 2015 and references therein). The results are listed in Table 4. We adopt extinction-corrected rather than observed [O iii] luminosity (e.g., Stern & Laor 2012) to estimate the bolometric luminosity, since the former provides a smaller scatter.","Citation Text":["Ho et al. 1993"],"Citation Start End":[[1013,1027]]}
{"Identifier":"2020ApJ...890....6S__Greisen_2015_Instance_1","Paragraph":"The Stokes V profile of maser A+B from our 2009 observations (lower panel of Figures 2 and 3) reveals an S-shaped structure that is usually taken to be a detection of the Zeeman effect (also see the discussion in Section 4 below); so does the Stokes V profile from our 2017 observations (lower panel of Figures 4 and 5). No other compact maser spots in this field show any such feature. Whenever an S-shaped feature is observed in Stokes V, the magnetic field strength is usually determined by fitting a numerical frequency derivative of the Stokes I spectrum to the Stokes V spectrum; details are in, e.g., Momjian & Sarma (2017). The Stokes V profile is usually fit simultaneously to the derivative of the I profile and a scaled replica of the I profile itself via the equation (Troland & Heiles 1982; Sault et al. 1990)\n1\n\n\n\n\n\nThe scaled replica of the I spectrum is included in the fit to account for small calibration errors in RCP versus LCP; for all results reported in this paper, a \u2272 10\u22123. The magnetic field values are contained in the fit parameter b, which is equal to z \n\n\n\n\n\n, where z is the Zeeman splitting factor, and \n\n\n\n\n\n is the line-of-sight (LOS) magnetic field strength (assuming, of course, that the signature in Stokes V is due to the magnetic field in the region; see Section 4.2). We used the AIPS task ZEMAN (Greisen 2015) to carry out the fit in Equation (1). This task allows multiple Gaussian components in I to be fitted simultaneously to V, with each Gaussian component fitted for a different b, and hence a different LOS magnetic field strength. For the three Gaussian components of maser A+B from our 2009 observations (listed in Table 2, and shown in the upper panel of Figure 2), the derivative profiles scaled by the respective values fitted for b = z \n\n\n\n\n\n are each shown in the lower panel of Figure 2; the composite profile obtained by adding together these three scaled derivative profiles is shown in the lower panel of Figure 3. Of these three components for maser A+B from our 2009 observations, only component 2 showed a significant fit, with the fitted value given by z \n\n\n\n\n\n = 152 \u00b1 12 Hz. Following customary practice in the field of Zeeman observations, we consider fits to be significant only if the ratio of fitted value to the fitted error is at the 3\u03c3 level or greater. Meanwhile, for the three Gaussian components of maser B from our 2017 observations (listed in Table 2, and shown in the upper panel of Figure 4), the derivative profiles scaled by the respective fitted values for \n\n\n\n\n\n are each shown in the lower panel of Figure 4, and the composite profile is shown in the lower panel of Figure 5. Only component 2 of maser B from our 2017 observations showed a significant fit, with the fitted value given by z \n\n\n\n\n\n = 149 \u00b1 19 Hz. Since component 2 of maser A+B from our 2009 observations and component 2 of maser B from our 2017 observations are very likely the same maser spot (see Section 4.1), this is a remarkable coincidence in \n\n\n\n\n\n over observations taken eight years apart.","Citation Text":["Greisen 2015"],"Citation Start End":[[1337,1349]]}
{"Identifier":"2021MNRAS.502..772L__Hopkins_2013_Instance_1","Paragraph":"A number of semi-analytic models (SAMs) have attempted to reproduce submm number counts (e.g. Granato et al. 2000; Fontanot et al. 2007; Somerville et al. 2012). One such model is the galform (SAM), which has been tuned to successfully reproduce the number counts of 850 \u03bcm and $\\mathrm{1.1 \\, mm}$ selected galaxies.1 However, in order to achieve this good agreement galform invokes a top-heavy initial mass function (IMF). Early versions of the model used a flat IMF above $1 \\, \\mathrm{M_{\\odot }}$, in sub-L* mergers (Baugh et al. 2005; Swinbank et al. 2008). This is required to produce sufficiently bright submm emission during frequent low-mass merger events. Later versions of the model used a more moderately top-heavy IMF in starbursts, triggered by disc instabilities rather than mergers, and found similarly good agreement with the number counts (Cowley et al. 2015, 2019; Lacey et al. 2016; Park et al. 2016). However, such IMF variability is still controversial, particularly extreme forms and any dependence on merger state (Bastian, Covey & Meyer 2010; Hopkins 2013; Krumholz 2014), and is inconsistent with the constraints on the IMF in massive star-forming galaxies that are significantly less extreme (e.g. Tacconi et al. 2008), though there is tentative evidence of a bottom-light\/top-heavy IMF in both local star-forming region analogues (Motte et al. 2018; Schneider et al. 2018) and some gravitationally lensed high-redshift starbursts (Zhang et al. 2018). Safarzadeh, Lu & Hayward (2017) showed that a variable IMF is degenerate with a number of other modelling processes in SAMs, such as the form of stellar feedback. They highlight that taking in to account dust mass allows for a good fit to the number counts without resorting to a variable IMF. Most recently, the shark SAM (Lagos et al. 2018) is able to broadly reproduce the 850 \u03bcm counts (whilst slightly overestimating the bright-end counts compared to S2CLS; Geach et al. 2017) using a fixed Chabrier (2003) IMF (Lagos et al. 2019). They attribute the good agreement to their use of physically motivated attenuation curves obtained from a self-consistent galaxy evolution model (eagle; Trayford et al. 2020).","Citation Text":["Hopkins 2013"],"Citation Start End":[[1069,1081]]}
{"Identifier":"2015ApJ...807..148G__Rees_&_M\u00e9sz\u00e1ros_1994_Instance_1","Paragraph":"The fireball model remains the most popular scenario for the gamma-ray burst (GRB) phenomenon (Cavallo & Rees 1978; Goodman 1986; Paczynski 1986; Shemi & Piran 1990; Rees & M\u00e9sz\u00e1ros 1992, 1994; M\u00e9sz\u00e1ros & Rees 1993). In this model, the GRB central engine is a stellar-mass black hole or a rapidly spinning and highly magnetized neutron star formed by either the collapse of a supermassive star (collapsar; Woosley 1993; MacFadyen & Woosley 1999; Woosley & Heger 2006) or the merger of two compact objects (Paczynski 1986; Fryer et al. 1999; Rosswog 2003). In both cases, the original explosion creates a bipolar collimated jet composed mainly of photons, electrons, positrons, and a small fraction of baryons. The relativistic explosion ejecta within the jet are not homogeneous\u2014they form multiple high density layers, which propagate at various velocities. When the fastest layers catch up with the slowest, the charged particles contained in the layers are accelerated through mildly relativistic collisionless shocks (internal shocks; Rees & M\u00e9sz\u00e1ros 1994; Kobayashi et al. 1997; Daigne & Mochkovitch 1998). The particles subsequently cool via emission processes such as synchrotron, Synchrotron Self Compton (SSC), and Inverse Compton (IC). The internal shock phase is usually associated with the so-called GRB prompt emission,17\n\n17\nSee Pe\u2019er (2015) for a recent review of GRB prompt emission.\n mainly observed in the keV\u2212MeV energy range (see, e.g., the spectral catalogs by Gruber et al. 2014; von Kienlin et al. 2014) and usually lasting from a few ms up to several tens to hundreds of seconds. As the ejecta interact with the interstellar medium they slow down via relativistic collisionless shocks (external shocks; Rees & M\u00e9sz\u00e1ros 1992; M\u00e9sz\u00e1ros & Rees 1993), accelerating charged particles, which then emit non-thermal synchrotron photons. This external shock phase is usually associated with the so-called GRB afterglow emission observed at radio wavelengths up to X-rays and in some cases even up to the GeV regime hours after the prompt phase, and days and even years for the lowest frequencies. The detailed origin of the gamma-ray emission, however, is not fully understood and many theoretical difficulties remain, such as the composition of the jet, the energy dissipation mechanisms, as well as the radiation mechanisms (e.g., Zhang 2011).","Citation Text":["Rees & M\u00e9sz\u00e1ros","1994"],"Citation Start End":[[166,181],[188,192]]}
{"Identifier":"2019ApJ...883...88B__Heinzel_&_Kleint_2014_Instance_1","Paragraph":"The GALEX NUV observations span a wide wavelength range, from 1771 to 2831 \u212b. While these data provide no spectral information within that bandpass, we rely on solar and stellar flare studies to inform the likely contributors to the flare flux. The NUV spectral region has not had as many observational constraints as the far-UV region in flare studies. Few NUV stellar flare spectra exist at all, and the few that do were obtained either for solar flares or on nearby M dwarfs. Flares observed in the UV are often associated with the more impulsive phases of solar flares, starting with early observations showing a close temporal association between UV and hard X-ray emission (Cheng et al. 1981). Welsh et al. (2006) reported on high time-resolution NUV+FUV flares seen with GALEX on nearby M dwarfs, and Hawley et al. (2007) presented high spectral-resolution NUV flare measurements of an M-dwarf flare. From these two studies, the contribution of emission lines relative to continuum emission could be determined; the main emission lines in the flare NUV spectrum were Mg ii, Fe ii, Al iii, and C iii. However, the main emission component overall was a continuum component. Recent results from solar flares observed from space (Heinzel & Kleint 2014; Kleint et al. 2016) demonstrate an NUV spectrum originating largely from Hydrogen Balmer continuum emission. The formation of the NUV emission appears to originate from an impulsive thermal and nonthermal ionization caused by the precipitation of electron beams through the chromosphere. This explains the temporal correlation with solar flare hard X-ray emission observed previously. More recently, Kowalski et al. (2019) presented accurately calibrated NUV flare spectra at high time cadence on an M dwarf, again finding a large flux enhancement due to continuum radiation. They commented that the oft-used 9000 K blackbody used to describe blue-optical stellar flare emission (Hawley et al. 2003) under-predicts the NUV continuum flare flux by a factor of two. Based on general similarities in radiative properties between solar and stellar flares studied thus far (Osten 2016), it is likely that a combination of line and continuum emission enhancements are present in the NUV flare flux from the flares being considered, but we cannot speculate about the relative contribution of one versus the other. These sources originate from different layers of the stellar atmosphere: singly and doubly ionized emission lines likely originate in the chromosphere, and the Balmer continuum emission also originates from the chromosphere. Some lines, such as Mg ii, exhibit absorption components and self-reversals, indicating optical depth effects in the atmosphere, while other lines such as Fe ii appear to be optically thin. Any hot blackbody emission might originate further in the photosphere.","Citation Text":["Heinzel & Kleint 2014"],"Citation Start End":[[1233,1254]]}
{"Identifier":"2021MNRAS.504.4626K__Kraljic_et_al._2020b_Instance_2","Paragraph":"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, Dav\u00e9 & 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\u2019s 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\u2019s 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 Text":["Kraljic et al. 2020b"],"Citation Start End":[[646,666]]}
{"Identifier":"2021MNRAS.505.4048L__Leroy_et_al._2015_Instance_1","Paragraph":"It is well known that giant molecular clouds (GMCs) are the major gas reservoirs for star formation (SF) and the sites where essentially all stars are born. Understanding the properties of GMCs is thus key to unravelling the interplay between gas and stars within galaxies. Early GMC studies were restricted to our own Milky Way (MW) and the late-type galaxies (LTGs) in our Galactic neighbourhood (e.g. Engargiola et al. 2003; Rosolowsky 2005, 2007; Rosolowsky et al. 2007; Gratier et al. 2012; Colombo et al. 2014; Wu, Sakamoto & Pan 2017; Faesi, Lada & Forbrich 2018), where GMCs have relatively uniform properties and generally follow the so-called Larson relations (between size, velocity dispersion, and luminosity; e.g. Blitz et al. 2007; Bolatto et al. 2008). However, more recent studies of other local galaxies have raised doubts on the universality of cloud properties. The cloud properties in some LTGs (such as M51 and NGC 253) vary with galactic environment and do not universally obey the usual scaling relations (e.g. Hughes et al. 2013; Leroy et al. 2015; Schruba, Kruijssen & Leroy 2019). The first study of individual GMCs in an early-type galaxy (ETG; NGC 4526) has also clearly shown that the clouds in that galaxy do not follow the usual size\u2013linewidth correlation and tend to be more luminous, denser and to have larger velocity dispersions than the GMCs in the MW and other Local Group galaxies (Utomo et al. 2015). The differences in NGC 4526 may be due to a higher interstellar radiation field (and\/or cloud extinctions), a different external pressure relative to each cloud\u2019s self-gravity, and\/or different galactic dynamics. GMCs in ETGs seem to have shorter orbital periods and be subjected to stronger shear\/tidal forces, analogous to the highly dynamic environment in the MW central molecular zone (CMZ; e.g. Dale, Kruijssen & Longmore 2019; Henshaw et al. 2019; Kruijssen et al. 2019). Although we are entering an era of large surveys of GMC populations (e.g. Sun et al. 2018), current samples of ETGs are still very limited. More studies of GMCs in varied LTGs and ETGs are thus required to provide a comprehensive census of GMC properties across different galaxy environments.","Citation Text":["Leroy et al. 2015"],"Citation Start End":[[1054,1071]]}
{"Identifier":"2022MNRAS.517.5744G__Caro_et_al._2016_Instance_2","Paragraph":"The CO photodesorption yield reaches its highest value when this ice is deposited at low temperatures (down to 7\u2009K, the lowest temperature studied experimentally) and decreases gradually at higher deposition temperatures (\u00d6berg et al. 2007; \u00d6berg et al. 2009; Mu\u00f1oz Caro et al. 2010, 2016; Sie et al. 2022). The explanation for this phenomenon motivated further research. It was found that the columnar structure of CO ice samples, grown at incidence angles larger than 45\u00b0, increases the effective ice surface exposed to UV photons and therefore the photodesorption efficiency (Gonz\u00e1lez D\u00edaz et al. 2019), but ice surface effects cannot account for the large variations observed in the photodesorption of CO ice samples deposited at different temperatures (Mu\u00f1oz Caro et al. 2016). Absorption band shifts of CO ice in the UV and IR ranges only occurred at deposition temperatures above 20\u2009K (Lasne et al. 2015; Mu\u00f1oz Caro et al. 2016), suggesting that CO ice grown at lower temperatures is amorphous below 20\u2009K in our experiments, and therefore, the decreasing photodesorption yield is not related to a transition from amorphous to crystalline ice, instead it might be associated to a different degree of molecular disorder in CO ice samples, depending on their deposition temperature. Photon energy transfer via Wannier-Mott excitons between the first photoexcited molecule in the ice and a molecule on the ice surface capable to desorb was proposed (Chen et al. 2017; McCoustra & Thrower 2018). Molecular disorder seems to enhance this energy transfer between neighbour molecules. The colour temperature variations measured at different deposition temperatures could also be the result of molecular disorder (Carrascosa et al. 2021). Urso et al. (2016), Cazaux et al. (2017), and Carrascosa et al. (2021) did not find significant changes in the desorption behaviour or the colour temperature of pure CO ice during controlled warm-up, which points to a low value of the diffusion in the ice. Finally, Sie et al. (2022) investigated the CO photodesorption yield dependence on ice thickness.","Citation Text":["Mu\u00f1oz Caro et al. 2016"],"Citation Start End":[[758,780]]}
{"Identifier":"2021MNRAS.503.5367B__Lu,_Kumar_&_Zhang_2020_Instance_1","Paragraph":"Since the discovery of the radio burst, there have been extensive follow-up observations of SGR\u2009J1935+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\u2009J1935+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\u22126127, 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 Text":["Lu, Kumar & Zhang 2020"],"Citation Start End":[[906,928]]}
{"Identifier":"2018AandA...612A..34D__Goldston_et_al._2005_Instance_1","Paragraph":"Sagittarius A* (Sgr A*) is a supermassive black hole system that allows one to observationally test the aforementioned GRMHD models of accretion flows (Goddi et al. 2017). Millimeter-Very Long Baseline Interferometry (mm-VLBI) is capable of resolving the shadow of the event horizon (Falcke et al. 2000), making this an ideal laboratory not only to tests Einstein\u2019s General Theory of Relativity (GR) but also to investigate electron acceleration in the vicinity of a black hole. Most of the radiative models for Sgr A*, which are based on post-processing GRMHD simulations, assume that electrons have a thermal, relativistic (Maxwell\u2013J\u00fcttner) distribution function, and that the proton-to-electron temperature ratio is constant across the simulation domain (Goldston et al. 2005; Noble et al. 2007; Mo\u015bcibrodzka et al. 2009; Dexter et al. 2010, 2012a; Shcherbakov et al. 2012). When the proton-to-electron temperature is constant, the disk dominates the images and spectra since most of the matter resides there. We have recently extended these radiative models by making the temperature ratios a function of the plasma \u03b2 parameter, where \n\n$\\beta = \\frac{P_{\\textrm{gas}}}{P_{\\textrm{B}}}$\n\n\n\u03b2=\n\n\nP\n\ngas\n\n\n\n\nP\nB\n\n\n\n\n\n is the ratio of gas to magnetic pressures. In these extended models, the electrons are hotter in the more magnetized plasma, which is usually outflowing from the system. The reason for this is that the previously mentioned models do not recover the flat radio spectra. The \u03b2 parameterization enforces that the disk emission is suppressed by significantly decreasing the temperature of the electrons in those regions. As a consequence of this, the jet will be the dominant source of emission. These modifications to the electron temperature model allowed us to recover some basic observational characteristics of Sgr A* (a roughly flat radio spectral slope and a size vs. wavelength relationship that is in agreement with observations) (Mo\u015bcibrodzka & Falcke 2013; Mo\u015bcibrodzka et al. 2014; Chan et al. 2015b,a; Gold et al. 2017). Our model for the electron temperatures as a function of the \u03b2 plasma parameter is now roughly confirmed with extended-GRMHD simulations that self-consistently take into account the evolution of the electron temperatures (Ressler et al. 2015, 2017). Moreover, GRMHD simulations with the new electron temperatures can naturally explain the symbiosis of disks and jets observed in many accreting black hole systems (Falcke & Biermann 1995; Mo\u015bcibrodzka et al. 2016a).","Citation Text":["Goldston et al. 2005"],"Citation Start End":[[758,778]]}
{"Identifier":"2020AandA...637A..44N__Kerszberg_et_al._2017_Instance_2","Paragraph":"Among the existing IACT systems, HESS has the largest FoV and hence provides the highest sensitivity for the diffuse \u03b3-ray flux. Its electron spectrum analysis technique could be directly used to obtain a measurement of the diffuse Galactic \u03b3-ray flux above energies of several TeV in the Galactic Ridge (|l|  30\u00b0, |b|  2\u00b0) region; see Figs. 3 and 4. A multi-year exposure of HESS could be sufficient for detection of the diffuse emission even from regions of higher Galactic latitude. This is illustrated in Fig. 5, where thick blue and red data points show mild high-energy excesses of the electron spectra derived by Kraus (2018), Kerszberg et al. (2017), Kerszberg (2017) over broken power-law models derived from the fits to lower energy data. Comparing these excesses with the level of the IceCube astrophysical neutrino flux and with the Fermi\/LAT diffuse sky flux from the region |b| > 7\u00b0 (corresponding to the data selection criterium of HESS analysis Kerszberg et al. 2017; Kerszberg 2017) we find that the overall excess flux levels are comparable to expected diffuse \u03b3-ray flux from the sky region covered by the HESS analysis (the quoted systematic error on the electron flux is \u0394log(EFE) \u2243 0.4). The overall excesses within 805 and 1186 h of HESS exposures (Kraus 2018; Kerszberg 2017) are at the levels of >4\u03c3 for the analysis of Kraus (2018) and 1.7\u03c3 for the analysis of Kerszberg (2017). A factor-of-ten longer exposure (which is potentially already available with HESS) could reveal a higher significance excess at the level of up to 5\u03c3. Such an excess is predicted in a range of theoretical models including interactions of cosmic rays injected by a nearby source (Andersen et al. 2018; Neronov et al. 2018; Bouyahiaoui et al. 2019) or decays of dark matter particles (Berezinsky et al. 1997; Feldstein et al. 2013; Esmaili & Serpico 2013; Neronov et al. 2018) or a large-scale cosmic ray halo around the Galaxy (Taylor et al. 2014; Blasi & Amato 2019).","Citation Text":["Kerszberg et al. 2017"],"Citation Start End":[[961,982]]}
{"Identifier":"2018ApJ...855L...8M__Melrose_&_Dulk_1988_Instance_1","Paragraph":"Irregular refraction (scattering) of radio waves due to density turbulence in the solar corona broadens the angular size of the background \u201cradio Sun\u201d and the compact radio sources there at any given observing frequency. But the extent of broadening and the interferometer baselines required to image the solar atmosphere in detail at radio frequencies remain a puzzle, particularly at frequencies 100 MHz (Erickson 1964; Aubier et al. 1971; Riddle 1974; McMullin & Helfer 1977; Robinson 1983; Melrose & Dulk 1988; Subramanian & Sastry 1988; Thejappa & Kundu 1992, 1994; Sastry 1994; Schmahl et al. 1994; Ramesh & Sastry 2000; Ramesh 2000; Ramesh et al. 2001; Kathiravan et al. 2002; Bastian 2004; Subramanian 2004; Thejappa & MacDowall 2008). Past\/present solar-dedicated radio imaging instruments in the above frequency range had\/have limited angular resolution (>3\u2032; Wild 1967; Labrum 1972; Kundu et al. 1983; Ramesh et al. 1998). Interestingly, observations carried out during different solar eclipses (similar to the lunar occulation technique) and the bandwidth, duration of some of the solar radio transients indicate that radiowave emitting sources with angular sizes \u22721\u2032 are present\/observable in the solar atmosphere from where radio emission at the above frequencies originate (Letfus et al. 1967; Melrose 1980; Benz & Wentzel 1981; McConnell 1983; Ramesh et al. 1999; Ramesh & Ebenezer 2001; Kathiravan et al. 2011; Ramesh et al. 2012). These are consistent with the theoretical predictions that solar radio sources of angular sizes \u227210\u2033 and \u227230\u2033 are observable at the typical frequencies of 327 MHz and 80 MHz, respectively (Riddle 1974; Subramanian & Cairns 2011). Encouraged by the above results and the availability of logistics support, we conducted interferometeric observations of the solar corona at frequencies 100 MHz on a \u2248200 km long interferometer baseline for the first time, the description of which is the subject matter of this article.","Citation Text":["Melrose & Dulk 1988"],"Citation Start End":[[494,513]]}
{"Identifier":"2019ApJ...874...86S__Moffat_1969_Instance_1","Paragraph":"For the NIR data set, we used standard IRAF11\n\n11\nIRAF is distributed by the National Optical Astronomy Observatory, which is operated by the Association of Universities for Research in Astronomy, Inc., under cooperative agreement with the National Science Foundation.\n (Tody 1986, 1993) tools to correct the raw images for flat-field and to perform the sky subtraction. A master sky was obtained by combining five sky images of a relatively empty field. For the optical data set, instead, we used images processed, flat-fielded, bias subtracted, and corrected for Charge Transfer Efficiency losses by standard HST pipelines (_flc images). The WFC3 images have also been corrected by the Pixel Area Map using the files available on the HST website. For both NIR and optical data, the photometric reduction was carried out via point-spread function (PSF) fitting techniques in each chip of each image independently, by using DAOPHOTIV (Stetson 1987). The PSF has been modeled by selecting about 200 bright and isolated stars uniformly distributed in each chip, and by using the DAOPHOTIV\/PSF routine. We allowed the PSF to vary within each chip following a cubic polynomial spatial variation. The best-fit PSF analytic models obtained for the HST F555W and F814W images, are a Moffat function with \u03b2 = 1.5 (Moffat 1969) and a Penny function (Penny 1976), respectively. The typical PSF model for the GEMINI images is a Penny function, but in some cases it was necessary to adopt a Lorentz function (Melcher & Gerth 1977) as well. The PSF models thus obtained were then applied to all the star-like sources detected at a 3\u03c3 level above the local background by using ALLSTAR. We have derived in this way the stellar instrumental magnitudes. Then, starting from the star lists thus obtained and to fill the gaps among the GSAOI chips, we created a GEMINI master star list containing all the stars measured in at least three Ks images. For HST instead, the master list was made up of all the stars measured in at least four F814W observations. As done in previous works (e.g., Dalessandro et al. 2014 and references therein), the master lists thus created for the two data sets have been used as input for ALLFRAME (Stetson 1994). The files obtained as output have been independently combined to get a GEMINI catalog containing all the stars measured in at least three J and three Ks images and an HST catalog containing stars measured in at least four F555W and four F814W exposures. For every stellar source in each catalog, different magnitude estimates have been homogenized and their mean values and standard deviations have been adopted as the star magnitudes and photometric errors in the final catalog (Ferraro et al. 1991, 1992).","Citation Text":["Moffat 1969"],"Citation Start End":[[1306,1317]]}
{"Identifier":"2016AandA...590A..19S__Benz_&_Asphaug_(1994)_Instance_1","Paragraph":"We simulated a laboratory impact experiment into a basaltic rock performed by Nakamura & Fujiwara (1991) to test the implementation of the damage model. In their experiment, a small 0.2 g nylon bullet impacts with an impact speed of 3200 m s-1 into a basaltic rocky sphere with a diameter of 6 cm. We chose to simulate an off-axis impact with an impact angle of 30 \u25e6 (see Fig. 9 for the definition of the impact angle), which corresponds to an impact parameter of 25% of the diameter. We modeled the basalt target with 523 741 particles and applied the Tillotson equation of state. Owing to the lack of parameters for nylon, we used lucite for the projectile material (values were taken from Benz & Asphaug 1994). The basalt target was modeled with the additional damage model and activation threshold strains were distributed following the Weibull distribution statistically as a preprocessing step (see Sect. 3.2 for a detailed description of the algorithm). Following Benz & Asphaug (1994), we performed several simulations with varying Weibull parameter pairs k and m. Figure 7 shows the damage pattern in the target after 50 \u03bcs for the simulation that matched the experimental outcome best. The damage in the region around the impact point is the largest and all particles close to the impact point were immediately damaged upon impact. The damage pattern is radially symmetric around the impact point. Another source of damage is found near the surface of the target. There, damage grows due to spallation stresses that occur when compressive waves get reflected at the free surface. As a result of this nature, the inner core of the target remains intact while the surface is coated with cracks. After 50 \u03bcs simulation time, the damage pattern in the target remains unchanged. We identified single fragments using a friend-of-friend algorithm in the following way: first, we removed all fully damaged particles. Then, we clustered all remaining particles that are closer than 1.01 of the original smoothing length to their neighboring particles. Fragments consisting of less than three particles were neglected. Figure 8 shows the fragment mass spectrum obtained with this friend-of-friend clustering algorithm. ","Citation Text":["Benz & Asphaug 1994","Benz & Asphaug (1994)"],"Citation Start End":[[692,711],[971,992]]}
{"Identifier":"2018ApJ...866...93L___2013b_Instance_1","Paragraph":"Interestingly, from \u223c11:33:40 UT, the PAD of the suprathermal electrons changes to cigar type, in association with a dramatic drop of electron flux (Figure 2(a)). This change of electron PAD and flux is related to a magnetic dip structure that is manifested by a conspicuous decrease of the magnetic field strength (Figure 2(b)). In contrast to the magnetic field enhancement observed at \u223c11:33:08 UT, the magnetic dip structure should arise from the local expansion of flux tubes, which are driven by the two opposite flows (see the shaded region in Figure 2(c)). The observed electron cigar distribution and the associated magnetic dip are a strong indication of the betatron cooling effect. This betatron-mediated cigar distribution has been suggested in previous studies (Fu et al. 2011b, 2012c, 2013b; Liu et al. 2017c) but never clearly observed. These observations, for the first time, show a direct link between the cigar distribution and betatron cooling. Note that in the trailing edge of the magnetic dip region, a magnetic hump structure, indicated by the sharp enhancement of magnetic field, is observed at \u223c11:33:55 UT (Figure 2(b)). This magnetic hump, reminiscent of the Earthward-propagating dipolarization-front structure typically generated by magnetic reconnection in the midtail (Fu et al. 2012d, 2012e, 2013a, 2015, 2016, 2017; Liu et al. 2013, 2018a, 2018b, 2018c; Cao et al. 2017; Peng et al. 2017; Yao et al. 2017; Chen et al. 2018), was possibly formed due to the local contraction of flux tubes driven by the ion flow with an increasing velocity in the Earthward direction (Figure 2(c)). Associated with this magnetic hump, a very weak pancake distribution of suprathermal electrons is observed (Figure 2(a)), and strong waves near the electron gyrofrequency are also observed (Figures 2(e) and (f)). These waves are whistler-mode; they indicate that the betatron acceleration inside the magnetic hump was accompanied by non-adiabatic effects (Fu et al. 2009, 2010a, 2010b, 2012a, 2012b) and is still ongoing (Fu et al. 2011a; Khotyaintsev et al. 2011; Wang et al. 2017; Yang et al. 2017).","Citation Text":["Fu et al.","2013b"],"Citation Start End":[[776,785],[800,805]]}
{"Identifier":"2021ApJ...920L..31N__Sterling_et_al._2017_Instance_2","Paragraph":"But if, as our studies in this Letter indicate, microstreams might be the result of accumulated and persistent velocity enhancements resulting from a series of switchbacks, then it could be that individual switchbacks result from coronal jets, and the microstreams are a consequence of a series of such jet-driven switchbacks occurring in close succession. Thus, this would be a modification of the idea put forth by Neugebauer (2012) whereby a series of minifilament eruptions capable of producing coronal jets could accumulate and generate a microstream. In fact, homologous jets, continuing for hours at a time, have been commonly observed (e.g., Chifor et al. 2008; Cheung et al. 2015; Panesar et al. 2016a, 2016b; Sterling et al. 2017; Joshi et al. 2017; Paraschiv & Donea 2019; Moore et al. 2021). Under the minifilament-eruption scenario, the multiple minifilament\/flux ropes would be ejected as long as the cancellation continues (Panesar et al. 2016a; Sterling et al. 2017). A swarm of such homologous jets, produced over a several-hour time period, conceivably could account for a microstream peak. Additionally, there is some recent evidence (Bale et al. 2021; Fargette et al. 2021) that switchbacks have an extent similar to the scale size of supergranules (\u223c30,000 km). Measurements of the lengths of the erupting minifilaments that produce jets range from \u223c8000 km (Sterling et al. 2015) to \u223c18,000 (Panesar et al. 2016a), and thus of similar order to (albeit somewhat smaller than) a typical supergranule diameter. Fargette et al. (2021) also found switchbacks to occur on another size scale, one that approximately corresponds to the size of photospheric granules, \u223c1000 km. Chromospheric spicules have widths of some fraction of this size, and thus their observation could be consistent with some spicules resulting from the minifilament-eruption-jet mechanism as suggested in Sterling & Moore (2016), and then those spicules making smaller-scale switchbacks as suggested in Sterling & Moore (2020).","Citation Text":["Sterling et al. 2017"],"Citation Start End":[[961,981]]}
{"Identifier":"2018ApJ...869..168D__Gosling_&_Pizzo_1999_Instance_1","Paragraph":"Figure 2 gives an overview of particle and solar wind parameters observed during a time span of 8 days before the electron event under consideration. The data seem to show no evidence for a CME and a related interplanetary shock. The solar wind speed, which had a low value of \u223c300 km s\u22121, starts to rise gradually after October 14 and reaches a maximum of \u223c650 km s\u22121 on October 19. At the same time the magnetic field strength and the proton density drop to low values. These variations are a typical signature of the presence of a CIR, i.e., a fast solar wind stream running into a slow solar wind stream and building up a compression region at a radial distance from the Sun of \u223c2 au, or between \u223c2.7 and 3 au along the field line connecting to the Sun (e.g., Gosling & Pizzo 1999, see Figure 3). The gradual increases in low-energy electron and proton fluxes (see the upper panel of Figure 2) could be caused by particle acceleration at the CIR (Giacalone et al. 2002). We found that on http:\/\/www.srl.caltech.edu\/ACE\/ASC\/DATA\/level3\/ a pair of stream interaction regions (SIRs) is listed for 2002 October 14\u201315, and that the Harvard\u2013Smithsonian Center for Astrophysics Interplanetary Shock Database (https:\/\/www.cfa.harvard.edu\/shocks\/) does not report any interplanetary shocks during 2002 October 12\u201320. We therefore conclude that the reflection of electrons in the 2002 October 20 event is likely due to\u2014depending on their pitch angle\u2014scattering or mirroring at a CIR-related compression, and not due to interaction with an interplanetary shock or propagation in a CME-related loop structure. Malandraki et al. (1997) described a technique for estimating the radial distance at which the magnetic compression arises due to stream\u2013stream interaction by mapping the observed time variation of the solar wind speed back to the solar corona. Applying this method using the observed time profile of the solar wind speed (second panel of Figure 3), we estimate the reflectivity to be located at a radial distance of 1.9\u20132.1 au, corresponding to a distance along the field line connecting to the Sun of \u223c2.4\u20132.7 au.","Citation Text":["Gosling & Pizzo 1999"],"Citation Start End":[[764,784]]}
{"Identifier":"2016ApJ...821..107G__Gloeckler_&_Fisk_2015_Instance_3","Paragraph":"We repeated the plasma pressure calculation presented by Schwadron et al. (2011) and Fuselier et al. (2012) for the new ENA energy spectrum. The results for the downwind hemisphere and for the Voyager 1 region are summarized in Table 3. The measured intensity \n\n\n\n\n\n\nj\n\n\nENA\n\n\n\n\n of neutralized hydrogen at a given energy translates into a pressure of the parent ion population in the heliosheath times the integration length along the line of sight, \u0394P \u00d7 l, in the following way:\n3\n\n\n\n\n\u0394\nP\n\u00d7\nl\n=\n\n\n\n4\n\u03c0\n\n\n3\n\n\nn\n\n\nH\n\n\n\n\n\n\n\nm\n\n\nH\n\n\nv\n\n\n\n\n\nj\n\n\nENA\n\n\n(\nE\n)\n\n\n\u03c3\n(\nE\n)\n\n\n\n\u0394\nE\n\n\n\nc\n\n\nf\n\n\n\n\n\n\n4\n\n\n\n\n\n\nc\n\n\nf\n\n\n=\n\n\n\n\n\n(\nv\n+\n\n\nu\n\n\nR\n\n\n)\n\n\n2\n\n\n\n\n\n\nv\n\n\n4\n\n\n\n\n\n(\n\n\nv\n\n\n2\n\n\n+\n4\n\n\nu\n\n\nR\n\n\n2\n\n\n+\n2\n\n\nu\n\n\nR\n\n\nv\n)\n.\n\n\nIn Equation (3), \u0394E denotes the width of the respective energy bin; for the typical radial velocity of solar wind in the flanks and the downwind hemisphere of the inner heliosheath, we assumed uR = 140 km s\u22121 as measured by Voyager 2, whereas uR = 40 km s\u22121 for the heliosheath in the Voyager 1 direction (Schwadron et al. 2011; Gloeckler & Fisk 2015). For the density of neutral hydrogen in the inner heliosheath a constant nH = 0.1 cm\u22123 is assumed (Schwadron et al. 2011; Gloeckler & Fisk 2015). The charge-exchange cross section between protons and neutral hydrogen decreases from (4 to 2) \u00d7 10\u221215 cm\u22122 for 0.015 to 1.821 keV (Lindsay & Stebbings 2005). The integration length l for ENA production in the plasma is approximately the thickness of the inner heliosheath. The part of Equation (3) without the velocity factor cf can be interpreted as stationary pressure. The total pressure or dynamic pressure is the stationary pressure times this factor. Integrating over all energy bins in Table 3, we obtain the total plasma pressure times heliosheath thickness as P \u00d7 l = 304 pdyn cm\u22122 au for the downwind hemisphere and 66 pdyn cm\u22122 au for the Voyager 1 region (1 pdyn cm\u22122 au = 0.015 N m\u22121). If we want to put these numbers into the context of other studies, we face two problems. First, the uncertainty of the total pressure is large given the upper limits in the two lowest energy bins. Second, heliosheath plasma more energetic than 2 keV will produce ENAs that cannot be detected with IBEX-Lo. We therefore used the observed median j = 0 cm\u22122 sr\u22121 s\u22121 keV\u22121 for heliospheric ENAs in the two lowest energy bins of IBEX-Lo and relied on the study by Livadiotis et al. (2013). They compared the expected plasma pressure from a kappa distribution of protons with the plasma pressure derived from IBEX-Hi energy spectra: the energy range between 0.03 and 2 keV, roughly corresponding to the IBEX-Lo range, covered more than half of the total plasma pressure predicted from a kappa distribution. The authors found a total plasma pressure of P = 2.1 pdyn cm\u22122 for all sky directions except for the ENA Ribbon. Gloeckler & Fisk (2015) presented a multi-component plasma model for the heliosheath to explain Voyager and IBEX observations. At low energies they assumed the ENA energy spectra provided by Fuselier et al. (2012). They derived a total pressure of 2.5 pdyn cm\u22122 in all three plasma regions in the nose of the heliotail (Gloeckler & Fisk 2015). Pressure contributions from the slowed solar wind, magnetic pressure, and the pressure exerted from pickup ions and anomalous cosmic rays all had to be taken into account to obtain this total pressure.","Citation Text":["Gloeckler & Fisk (2015)"],"Citation Start End":[[2813,2836]]}
{"Identifier":"2020AandA...635A.154M__Hummels_et_al._(2017)_Instance_2","Paragraph":"This paper presents a new 3D RT code, RAdiation SCattering in Astrophysical Simulations (RASCAS), which was designed to construct accurate multi-wavelength mock observations (spectra, images, or datacubes) from high-resolution simulations. RASCAS deploys a general two-step methodology (e.g. Hummels et al. 2017; Barrow et al. 2017). The first step consists in extracting all relevant information from the simulation outputs: (1) the information concerning the medium through which light will propagate, and (2) the information concerning the sources of radiation. Point (1), for example, determines the number density of H\u202fI atoms, their thermal velocity dispersion and bulk velocity, and the dust density, everywhere in a chosen volume. Fully coupled radiation-hydrodynamic simulations would naturally provide the information about ionised states, but retrieving this information in pure hydrodynamic simulations may be tricky. In such case, it is necessary to process the simulation outputs with additional software and models, typically with CLOUDY (Ferland et al. 2013), as in Hummels et al. (2017) and Barrow et al. (2017), or with independent codes which solve for H and He ionisation by propagating ionising radiation in post-processing (e.g. Li et al. 2008; Yajima et al. 2012). Point (2) is also involved to a greater or lesser extent, depending on the sources. Computing the continuum emission from star particles is relatively straightforward, using spectro-photometric models of stellar populations (e.g. Bruzual & Charlot 2003; Eldridge et al. 2017). However, computing the emission lines from the gas (e.g. in the Lyman-\u03b1 line or in other nebular lines) again requires a detailed knowledge of the ionisation and thermal state of the emitting species. This may be provided by the simulation code, as is the case for H and He with RAMSES-RT (Rosdahl et al. 2013), or it may be necessary to post-process the simulation to estimate the emissivities of the gas. This first step is very simulation- and model-dependent, and RASCAS chooses to encapsulate it in a simulation-plugin module and to implement two stand-alone pre-processing codes which generate an adaptive mesh with all the needed physical information about the gaseous medium, and the initial conditions for light emission in the form of lists of photon packets. These datasets, which could easily be generated from other simulations with any post-processing code, serve as inputs to RASCAS to perform the radiative transfer computation.","Citation Text":["Hummels et al. (2017)"],"Citation Start End":[[1082,1103]]}
{"Identifier":"2019MNRAS.484.1487E__Roca-Fabrega_et_al._2013_Instance_1","Paragraph":"Since the manifold spirals arise in a system of reference which corotates with the bar, the manifold theory in its basic form predicts that the spiral arms should have the same pattern speed as the bar. This remark seems to come in conflict with observations both in our Galaxy (as reviewed e.g. in Bland-Hawthorn & Gerhard 2016; see also Antoja et al. 2014; Junqueira et al. 2015 and references therein) and in other galaxies (e.g. Vera-Villamizar et al. 2001; Boonyasait, Patsis & Gottesman 2005; Patsis, Kaufmann & Gottesman 2009; Meidt, Rand & Merrifield 2009; Speights & Westpfahl 2012; Speights & Rooke 2016). Considering again, galactic disc simulations, the leading paradigm over the years refers to simulations showing the coexistence of multiple pattern speeds (Sellwood & Sparke 1988; Little & Carlberg 1991; Rautiainen & Salo 1999; Quillen 2003; Minchev & Quillen 2006; Dubinski, Berentzen & Shlosman 2009; Quillen et al. 2011; Minchev et al. 2012; Baba, Saitoh & Wada 2013; Roca-Fabrega et al. 2013; Font et al. 2014; Baba 2015; but see also a noticeable exception in Roca-Fabrega et al. 2013), possibly connected also to the phenomenon of nonlinear coupling of multiple disc modes (Tagger et al. 1987; Tagger & Athanassoula 1991; Sellwood & Wilkinson 1993; Masset & Tagger 1997). On the other hand, it is well known that even isolated barred galaxies undergo substantial secular evolution (see Athanassoula 2013; Binney 2013; Kormendy 2013 in the tutorial volume Falcon-Barroso & Knapen 2013). The tendency to transfer angular momentum outwards (e.g. towards the halo or across the disc, Tremaine & Weinberg 1984; Debattista & Sellwood 1998; Debattista & Sellwood 2000; Athanassoula 2002; Athanassoula & Misiriotis 2002; Athanassoula 2003; O\u0143eill & Dubinski 2003; Holley-Bockelmann, Weinberg & Katz 2005; Berentzen, Shlosman & Jogee 2006; Martinez-Valpuesta, Shlosman & Heller 2006) leads the bar to slow down and grow in size at a rate which produces non-negligible change in dynamics at time-scales comparable even to a few bar periods. This process becomes complex, and even partially reversed due to the growth of \u2018pseudo-bulges\u2019 or peanuts (Kormendy & Kennicutt 2004), caused by dynamical instabilities such as chaos or the \u2018buckling instability\u2019 (Combes & Sanders 1981; Combes et al. 1990; Pfenniger & Friedli 1991; Raha et al. 1991; Bureau & Athanassoula 1999; Martinez-Valpuesta & Shlosman 2004; Bureau & Athanassoula 2005; Debattista et al. 2006). The reduction in size of the bar by the transfer of angular momentum under constant pattern speed is discussed in Weinberg & Katz (2007). Spiral activity acts as an additional factor of outwards transfer of angular momentum (Lynden-Bell & Kalnajs 1972), while a radial re-distribution of matter can take place even under a nearly preserved distribution of angular momentum (Hohl 1971; Sellwood & Binney 2002; Avila-Reese et al. 2005). Radial migration is enhanced by the amplification of chaos due to the overlapping of resonances among the various patterns (Quillen 2003; Minchev & Quillen 2006; Quillen et al. 2011).","Citation Text":["Roca-Fabrega et al. 2013"],"Citation Start End":[[987,1011]]}
{"Identifier":"2015MNRAS.448.2798S__Ermolli_et_al._2009_Instance_1","Paragraph":"The study of the variations in chromospheric flux using Ca-K spectral line at all time-scales is important to study the solar dynamo because of relation between the Ca-K line and the magnetic fields on the Sun (Leighton 1964; Skumanich, Smythe & Frazier 1975; Sivaraman & Livingston 1982; Ortiz & Rast 2005). Short time-scale variations provide the information about the dynamics and energy transport in the chromosphere whereas the long period variations provide valuable knowledge about the dynamo process. The observed Ca-K line irradiance variation with solar cycle is due to the contribution of various features such as plages, enhanced, active and quiet network to the flux. Also, the flux modulation occurs because of varying contribution with latitude due to limb darkening. Irradiance variations in Ca-K line have been studied by many researchers using long series of spectroheliograms obtained at Kodaikanal, Mount Wilson, Sac Peak and other observatories (Ermolli et al. 2009; Foukal et al. 2009; Tlatov, Pevtsov & Singh 2009; Priyal et al. 2013, and references therein) using plage areas as a proxy to the chromospheric irradiance variation. Worden, White & Woods (1998) analysed 1400 spectroheliograms obtained at the National Solar Observatory, Sac Peak spread between the period of 1980\u20131996 to identify the different features such as plages, enhanced, active and quiet network using the intensity contrast technique to study the rotational modulation and solar cycle variations in the chromosphere. Walton, Preminger & Chapman (2003) concluded that faculae are responsible for 80\u2009per\u2009cent of the solar cycle variations whereas Ermolli, Berrilli & Florio (2003) estimated that during the ascending phase of the solar cycle (1996\u20132002) the network contributed 40\u201350\u2009per\u2009cent to the solar cycle variations. Following this work Priyal et al. (2013) digitized the spectroheliograms obtained at Kodaikanal Observatory and analysed the data for three solar cycles (1955\u20131985) and found good correlation with the Mount Wilson data and sunspot number.","Citation Text":["Ermolli et al. 2009"],"Citation Start End":[[967,986]]}
{"Identifier":"2021AandA...655A..72S___2019_Instance_1","Paragraph":"In this paper, we report on spectroscopic CH3CN, CH3OH (methanol), and dust continuum observations with the Atacama Large Millimeter\/submillimeter Array (ALMA) at 349 GHz with an angular resolution of 0\u2032\u2032.1. We exploit the CH3CN (19K\u201318K) K-ladder, with excitation energies ranging from 168 K (for K = 0) to 881 K (for K = 10), to probe, at different radii, the physical conditions in the accretion disk of an early-type young star. We targeted the star-forming region G023.01\u221200.41, at a trigonometric distance of 4.59\n\n$^{+0.38}_{-0.33}$\n\n\n\n\n\n\u22120.33\n\n+0.38\n\n\n\n kpc from the Sun (Brunthaler et al. 2009), where we recently revealed the accretion disk around a young star of 104.6 L\u2299, corresponding to a ZAMS star of 20 M\u2299 (Sanna et al. 2019, their Fig. 1); the disk was imaged by means of spectroscopic ALMA observations of both CH3CN and CH3OH lines at 0\u2032\u2032.2 resolution inthe 230 GHz band. The disk extends up to radii of 3000 au from the central star where it warps above the midplane; here, we resolve the outer disk regions in two apparent spirals projected onto the plane of the sky. We showed that molecular gas is falling in and slowly rotating with sub-Keplerian velocities down to radii of 500 au from the central star, where we measured a mass infall rate of 6 \u00d7 10\u22124 M\u2299 yr\u22121 (Sanna et al. 2019, their Fig. 5). The disk and star system drives a radio continuum jet and a molecular outflow aligned along a position angle of 57\u00b0, measured east of north (Sanna et al. 2016, their Fig. 2); their projected axis is oriented perpendicular to the disk midplane whose inclination with respect to the line-of-sight was estimated to be less than 30\u00b0 (namely, the disk is seen approximately edge-on; Sanna et al. 2014, 2019). Previously, we also measured the average gas conditions over the same extent of the whole disk, by means of Submillimeter Array (SMA) observations of the CH3CN (12K\u201311K) emission, and we estimated a kinetic temperature of 195 K and CH3CN column density of 5.1 \u00d7 1016 cm\u22122 (Sanna et al. 2014, their Fig. 2 and Table 4).","Citation Text":["Sanna et al. 2019"],"Citation Start End":[[723,740]]}
{"Identifier":"2015AandA...574A..62S__Kucera_et_al._(1998)_Instance_1","Paragraph":"In the last two decades UV, EUV, and X-ray observations from space made it possible to estimate the total mass of prominences reliably without needing to solve problems that occur when using visible light as was described in the previous paragraph. One possibility would be to estimate column mass of hydrogen and\/or helium plasma in prominences using spectral observations of UV and visible lines of hydrogen and helium (Balmer and Lyman lines) and sophisticated models in which plasma is assumed not being in the state of local thermodynamic equilibrium (NLTE models; see e.g. Labrosse et al.2010, and references therein). The problem can be a rather high complexity of such NLTE models which depend on various free parameters. Another possibility is to infer column mass and subsequently the total mass from the amount of radiation absorbed by the photoionisation in the prominence plasma at resonance continua of hydrogen and helium. This estimation of the mean column mass of prominences observed near the limb was first made by Kucera et al. (1998) using extreme-ultraviolet (EUV) observations from the Coronal Diagnostic Spectrometer (CDS; Harrison et al. 1995) on board the Solar and Heliospheric Observatory (SoHO) satellite. The advantage of the observations used by Kucera et al. (1998) was that CDS as a spectrograph observed only in spectral lines of interest, but spectrographs are able to obtain spectra from only one slit position during one exposure. Larger fields of view can be scanned by the spectrograph slit, as the CDS does, for example, but in such a case, intensities in different slit positions the scan is composed of were obtained at different times. In contrast, in filtergrams, intensities in all positions of the field of view are obtained at the same time, but the filter has transmission of certain width which can be wider than the spectral line of interest. The first estimation of the prominence mass using filtergrams was made by Golub et al. (1999) using Transition Region and Coronal Explorer (TRACE1) data. Similar studies were made by Gilbert et al. (2005, 2006) using observations of EUV Imaging Telescope (EIT; Delaboudini\u00e8re et al. 1995) on board SoHO in the 195\u2009\u00c5 channel. In those two works it was shown that it is necessary to estimate an amount of coronal emission behind and in front of the prominence (hereafter referred to as background and foreground radiation, respectively) to determine correctly the amount of absorbed radiation. The foreground radiation can be measured at the darkest place at a prominence, where it is assumed that all radiation from behind the prominence was absorbed. Then, the background radiation can be derived from the total coronal emission at the prominence location. Two ways to estimate the total emission were proposed and used. The spatial interpolative approach uses interpolation from intensities measured in the corona near a prominence. The temporal interpolative approach is suitable only for erupting prominences and it uses measurements of intensity in place of prominence after its eruption. In the work of Williams et al. (2013) these two approaches are also used to estimate the column mass of material returning to limb after prominence eruption. They used observations from the Atmospheric Imaging Assembly (AIA) instrument (Lemen et al. 2012) on board the Solar Dynamics Observatory (SDO) in several EUV coronal channels of wavelengths below 228\u2009\u00c5 (head of the resonance continuum of He\u2009ii). ","Citation Text":["Kucera et al. (1998)"],"Citation Start End":[[1034,1054]]}
{"Identifier":"2018AandA...613A...7Y__Kostogryz_et_al._(2016)_Instance_2","Paragraph":"where a is the area of the i, j pixel (constant for a regular grid), \u03bc is the anglebetween the surface normal and the line of sight to the observer, \u03d5 is the polar angle of a system with the origin at the disk center, F is the total stellar flux, and q and u are normalized Stokes parameters. We note that the relative flux is normalized to the total flux of an unspotted photosphere in our analysis. The center-to-limb variations ofintensity I(\u03bcij) and polarization P(\u03bcij) were found through trilinear interpolation using the look-up tables from Kostogryz & Berdyugina (2015) and Kostogryz et al. (2016), according to the selected wavelength, surface gravity, and temperature of a star. Kostogryz & Berdyugina (2015) calculated the tables for continuum spectra of FGK stars (Teff = 4500 K\u20136900 K, log g = 2.0\u20135.0, and wavelength range 4000\u20137000 \u212b) for the Phoenix grid of plane-parallel models, and in Kostogryz et al. (2016) similar calculations were made for a wider range of models (Teff = 4000 K\u20137000 K and log g = 1.0\u20135.5) assuming a spherical stellar atmosphere. In addition, we used unpublished data on intensity and polarization variations for both plane-parallel and spherical cooler models with temperatures down to 3000 K, obtained with the same code by Kostogryz (2016, priv. comm.). It should be noted that these data for cooler stars do not include atomic and molecular absorption lines, which can lead to overestimated polarization values, especially for blue wavelengths and for dwarfs with higher surface gravities (depending on the selected wavelength). Evidently, a proper spectrum synthesis code is needed to calculate the intrinsic polarization for these cases. However, the depolarizing effect from absorption lines may not be very significant because when the absorption line suppresses the polarization, the intensity is also reduced. In turn, this will affect the normalized polarization parameters that we calculate less strongly. Additionally, coherent scattering processes, especially in molecular bands, as seen on the Sun (e.g., Berdyugina et al. 2002) probably increase polarization further, which counteracts the effect of the absorption lines.","Citation Text":["Kostogryz et al. (2016)"],"Citation Start End":[[903,926]]}
{"Identifier":"2021ApJ...909...65K__SN_1999b_Instance_1","Paragraph":"White dwarfs (WDs) are the end state of stars with mass \u22728M\u2299. A WD attains its stable equilibrium configuration by balancing the outward force due to the degenerate electron gas with the inward force of gravity. If the WD has a binary companion, it pulls out matter from the companion, and as a result, the mass of WD increases. Once the mass of the WD reaches the Chandrasekhar mass limit (Chandrasekhar 1931; \u223c1.4M\u2299 for carbon\u2013oxygen nonrotating, nonmagnetized WDs), the pressure balance no longer sustains, and the WD bursts out to produce a Type Ia supernova (SN Ia; Choudhuri 2010). The similarity in peak luminosities of SNe Ia is used as one of the standard candles to estimate the luminosity distances for various astronomical and cosmological objects (Lieb & Yau 1987; Nomoto et al. 1997). However, recent discoveries of various under- and overluminous SNe Ia question the complete validity of considering luminosities of SNe Ia as standard candles. SNe Ia such as SN 1991bg (Filippenko et al. 1992; Mazzali et al. 1997), SN 1997cn (Turatto et al. 1998), SN 1998de (Modjaz et al. 2001), SN 1999by (Garnavich et al. 2004), and SN 2005bl (Taubenberger et al. 2008) were discovered with extremely low luminosities, which were produced from WDs with 56Ni mass content as low as \u223c0.1M\u2299 (Stritzinger et al. 2006). On the other hand, a different class of SNe Ia, such as SN 2003fg (Howell et al. 2006), SN 2006gz (Hicken et al. 2007), SN 2009dc (Yamanaka et al. 2009; Tanaka et al. 2010; Silverman et al. 2011; Taubenberger et al. 2011; Kamiya et al. 2012), SN 2007if (Scalzo et al. 2010; Yuan et al. 2010; Scalzo et al. 2012), SN 2013cv (Cao et al. 2016), and many more was discovered with an excess luminosity, with the observed mass of 56Ni as high as \u223c1.8M\u2299 (Kamiya et al. 2012), violating the Khokhlov pure detonation limit (Khokhlov et al. 1993). It was inferred that these underluminous SNe Ia were produced from WDs with a mass \u223c0.6M\u2299 (Mazzali et al. 1997; Turatto et al. 1998), while the same for overluminous SNe Ia could be \u223c2.8M\u2299 (Scalzo et al. 2010; Kamiya et al. 2012). Hence, these progenitor WDs of peculiar SNe Ia violate the Chandrasekhar mass limit: the underluminous SNe Ia were produced from sub-Chandrasekhar limiting-mass WDs (WDs burst before reaching the mass \u223c1.4M\u2299), and the overluminous SNe Ia were produced from super-Chandrasekhar limiting-mass WDs (WDs burst well above the mass \u223c1.4M\u2299). These new mass limits are important, as they may lead to modifying the standard candle.","Citation Text":["Garnavich et al. 2004"],"Citation Start End":[[1107,1128]]}
{"Identifier":"2015AandA...582A..42K__Poveda_et_al._(1967)_Instance_1","Paragraph":"Blaauw (1961) hypothesized that runaway stars are created when a massive component within a binary system explodes as a Type II supernova. The secondary is then ejected with a velocity comparable to the orbital velocity at the time of the supernova event. The explosion may not disrupt the system (Hills 1983), which will be observed as a OB-neutron star\/black-hole system and eventually become a high-mass X-ray binary. Given that D2\u2013EB is a double-lined O-type binary, it is unlikely that such a mechanism took place. Alternatively, D2\u2013EB could have been part of a triple system including a very massive component (\u226585 M\u2299) with a lifetime of less than 4.5 Myr. However, the inferred radial velocity curve lacks evidence of a low-mass third companion, unless it orbits with a high eccentricity and\/or long period. Observations of D2\u2013EB over a longer time span than the data presented here are required to investigate the possibility of a third body in the system. Poveda et al. (1967) suggested that runaway stars are dynamically ejected due to encounters of collapsing protostars in\/near the core of young clusters. Interactions in clusters that contain initial hard binaries increase the number of escapees via binary-binary interactions (Mikkola 1983). The less massive binary system reaches peculiar space velocities up to ~200 km s-1, while the more massive system travels at a speed less than ~100 km s-1 (Leonard & Duncan 1988). The binary frequency for runaways with V\u221e> 30 km s-1 is predicted to be 10% and it is striking that systems hosting components of M ~ 20 M\u2299 as the most massive members are not predicted to escape with more than 50 km s-1 (Leonard & Duncan 1990). This does not conflict with the observed lower limit for the space velocity of D2\u2013EB. A young, dense cluster core increases the possibility of dynamical ejection, with the consequence that the kinematic age of the runaways is similar to the age of the cluster (Gualandris et al. 2004). In this case, the unreasonably low estimated value for the tangential velocity of D2\u2013EB makes the possibility of the dynamical ejection through binary-binary interaction unlikely. ","Citation Text":["Poveda et al. (1967)"],"Citation Start End":[[965,985]]}
{"Identifier":"2016AandA...591A..38C__Roediger_et_al._(2011a)_Instance_2","Paragraph":"Despite the considerable scatter in both colors and color gradients (Peletier & Balcells 1996; Taylor et al. 2005; Roediger et al. 2011a), the tight correlation between color and stellar mass of the host galaxies holds true in both the region identified as intermediate and outer in Sect. 7. ETGs form a tight red sequence (see Fig. 10) for both regions and show an average inside-out gradient of ~0.1 mag (Fig. 12, bottom). LTGs form instead two different distributions: the blue cloud of the intermediate (or bulge\/bar) region becomes red (i.e., it reaches the red sequences) above 1010 M\u2299 while the outer, disk-dominated region never overlaps completely the red sequence. Still, more massive disks are redder than their lower mass counterparts but the difference between the colors of the outer and the intermediate region increases with respect to the total stellar mass. This can be possibly induced by the growth of a red and dead structure in the center of massive disks, i.e. that the central part of galaxies underwent a star formation quenching process that turned them red. Our results on the average properties of color profiles broadly agree with literature data. For example MacArthur et al. (2004) have shown that the radial profile of the average ages of the stellar populations decreases from inside out and that the steepness of the decrease is a function of morphological type. Color templates shown in Fig. 9 exhibit a radial behavior fully consistent with the average age profiles shown by MacArthur et al. (2004). We find there is also a good agreement with the (g\u2212H) color profiles published in Roediger et al. (2011a) for almost all the morphological types, although in their median profiles of early disks they find positive color gradients (and consistently positive age population gradients in their stellar population analysis, Roediger et al. 2011b) that we do not see. Roediger et al. (2011b) do not find any direct link between galaxy morphologies and the observed stellar population gradients. On the contrary, studies such as Cheung et al. (2013) and Gavazzi et al. (2015) have shown that the bar occupation fraction rises steeply above 109.5 M\u2299 (as also confirmed by the works done by  Skibba et al. 2012; Masters et al. 2012) and that, above this mass,galaxies are progressively more quenched (red) in their centers, while their disks still sustain SF and hence are blue. These studies therefore highlight that the presence of structures such as bars can indeed produce the color gradients that we observe and likely also the stellar population gradients. These authors thus conclude that a secular bar drives the quenching of the star formation in the central kiloparsecs of galaxies. Moreover massive galaxies undergo bar instability earlier than their lower mass counterparts and thus have more time to grow redder than low mass systems. Moreover, M\u00e9ndez-Abreu et al. (2012) on a study of the Virgo bar fraction have shown that this rises up to more than 50% above 1010 M\u2299, adding a further link between color\/stellar populations radial gradients that we observe and structures such as bars (Laurikainen et al. 2010). We stress, however, that disk instabilities can also rejuvinate the central stellar population by, e.g., triggering central star formation in correspondence of the ILRs of spirals or bars (an example could be the one of VCC 508 in Fig. 3a). ","Citation Text":["Roediger et al. (2011a)"],"Citation Start End":[[1618,1641]]}
{"Identifier":"2020AandA...637A..44N__Kraus_(2018)_Instance_1","Paragraph":"Among the existing IACT systems, HESS has the largest FoV and hence provides the highest sensitivity for the diffuse \u03b3-ray flux. Its electron spectrum analysis technique could be directly used to obtain a measurement of the diffuse Galactic \u03b3-ray flux above energies of several TeV in the Galactic Ridge (|l|  30\u00b0, |b|  2\u00b0) region; see Figs. 3 and 4. A multi-year exposure of HESS could be sufficient for detection of the diffuse emission even from regions of higher Galactic latitude. This is illustrated in Fig. 5, where thick blue and red data points show mild high-energy excesses of the electron spectra derived by Kraus (2018), Kerszberg et al. (2017), Kerszberg (2017) over broken power-law models derived from the fits to lower energy data. Comparing these excesses with the level of the IceCube astrophysical neutrino flux and with the Fermi\/LAT diffuse sky flux from the region |b| > 7\u00b0 (corresponding to the data selection criterium of HESS analysis Kerszberg et al. 2017; Kerszberg 2017) we find that the overall excess flux levels are comparable to expected diffuse \u03b3-ray flux from the sky region covered by the HESS analysis (the quoted systematic error on the electron flux is \u0394log(EFE) \u2243 0.4). The overall excesses within 805 and 1186 h of HESS exposures (Kraus 2018; Kerszberg 2017) are at the levels of >4\u03c3 for the analysis of Kraus (2018) and 1.7\u03c3 for the analysis of Kerszberg (2017). A factor-of-ten longer exposure (which is potentially already available with HESS) could reveal a higher significance excess at the level of up to 5\u03c3. Such an excess is predicted in a range of theoretical models including interactions of cosmic rays injected by a nearby source (Andersen et al. 2018; Neronov et al. 2018; Bouyahiaoui et al. 2019) or decays of dark matter particles (Berezinsky et al. 1997; Feldstein et al. 2013; Esmaili & Serpico 2013; Neronov et al. 2018) or a large-scale cosmic ray halo around the Galaxy (Taylor et al. 2014; Blasi & Amato 2019).","Citation Text":["Kraus (2018)"],"Citation Start End":[[620,632]]}
{"Identifier":"2018AandA...620A.122S__\u0160vanda_et_al._(2014)_Instance_2","Paragraph":"Simon & Leighton (1964) described supergranulation as a system of atmospheric currents in the photosphere. The currents form a cellular network, which is visible in Doppler maps after the reduction of other larger scale flows (e.g. differential rotation and convective blueshift). Inside the cells, the flow is radially directed from the cell centre to its boundary. The diameters of supergranules are in the range of 10 Mm up to 45 Mm (Simon & Leighton 1964; Roudier et al. 2014; Orozco Su\u00e1rez et al. 2012). Depending on the method, different values are measured (Hirzberger et al. 2008). A mean horizontal velocity of 0.4 km s\u22121 in supergranules was measured by Simon & Leighton (1964). Orozco Su\u00e1rez et al. (2012) investigated supergranular convective flows using Fourier local correlation tracking and intergranular magnetic elements. The flow velocity in supergranules increases outwards for larger distances from the centre. After reaching a maximum of 0.35 km s\u22121, the velocity decreases monotonically. By averaging over 222 976 supergranular cells and applying time-distant helioseismic inversions, \u0160vanda et al. (2014) found a symmetrical flow in supergranules, which is directed radially away from the centre of the cells to its periphery, with horizontal velocities in the range of 0.3\u20130.6 km s\u22121. Their velocity profile is found to increase with increasing distances to the supergranule centre up to a maximum value from where on a continuous decrease for larger distances starts. According to Simon & Weiss (1968), the non-stationary cells only survive their turnover time, then the flow lapses into disorder. Supergranular cells have a mean lifetime of 1.5 days, which can extend up to 4 days (Roudier et al. 2014; Hirzberger et al. 2008). De Rosa & Toomre (2004) observed the creation of supergranules due to fragmentation or merging of older cells and stated that each supergranular cell takes part in a minimum of one merging or splitting event during its lifetime. This interaction seems to be the preferred mode of evolution of the observed supergranular pattern. The mechanism of advecting magnetic field elements to the network and the similar appearance suggests a relation between the supergranular pattern and the quiet-Sun magnetic network, as was recognised by Simon & Leighton (1964). De Rosa & Toomre (2004) observed a strengthening or weakening of the network lanes due to splitting or merging of supergranules. Although there is no definite evidence of a one-to-one relation (see review by Rieutord & Rincon 2010 and references therein), in the large-scale picture, the supergranular pattern can be equated with the magnetic network. The horizontal, radially outward directed gas motion of the moat flow, which is visible in Doppler maps, resembles the characteristics of supergranular flows. In addition, magnetic lanes can form around the moat cell due to magnetic features (MMFs), which cross the moat and conglomerate, resembling the magnetic network. This relation has been described by Simon & Leighton (1964), Sheeley (1972), and Vrabec et al. (1974), for example, who proposed the sunspot to be sitting in the centre of a supergranule. Meyer et al. (1974) concluded from the observations that sunspots are related to supergranular convection, but these two flows should be distinguished because of their difference in size and occurrence by taking into account the unique relation of the moat flow to its sunspot. \u0160vanda et al. (2014) found the moat flow to be asymmetric, while flows in supergranules are symmetric, but the investigated cells approximately show the same size. In addition, the authors described the moat cell as a downflow region, while the motion in supergranules shows an upflow\u2212downflow behaviour. Various investigations have been carried out to study the differences and similarities between the moat flow and supergranules (see e.g. Sobotka & Roudier 2007 and the reviews by Solanki 2003 and Rieutord & Rincon 2010 and references therein). An overview of the characteristics of the moat flow and supergranules is given in Table 1.","Citation Text":["\u0160vanda et al. (2014)"],"Citation Start End":[[3452,3472]]}
{"Identifier":"2022MNRAS.516.2500C__Lin_et_al._2009_Instance_3","Paragraph":"Neutron star X-ray binaries are an important class of low-mass X-ray binaries to understand the radiative and dynamical configuration of the inner region of an accretion disc. Though from previous studies especially based on RXTE (Rossi X-ray Timing Explorer) data of Z sources, it was known that there must exist a corona\/comptonization region to explain the observed hardtail in their X-ray spectra but the exact location and how it changes across the intensity variation is not yet properly understood. Among the two primary categories i.e. Z and Atoll sources, Z sources emit close to the Eddington luminosity (0.5\u20131.0 LEdd; Done, Gierli\u0144ski & Kubota 2007a) and they exhibit \u2018Z\u2019 and \u2018C\u2019 shape intensity variation in the hardness intensity diagram (HID) or colour\u2013colour diagrams (CCDs; Hasinger & Van der Klis 1989; Van der Klis 2006). The Z shape variation constitutes a horizontal branch (HB) at the top, a flaring branch (FB) at the bottom, and a normal branch (NB) connecting them diagonally. These are further classified into two broad groups, namely Sco and Cyg-like sources, due to their different appearance exhibited by the HB and FB i.e. less vertical orientation of HB and a weaker FB is seen among Cyg-like sources than in Sco-like (Kuulkers et al. 1994). The hybrid source XTE J1701\u2013462 occupies a special place among NS LMXBs and is considered to be a remarkable source, as it displays all the characteristics exhibited by both Z and atoll sources (Homan et al. 2010, 2007; Lin et al. 2009). At the brightest state, the intensity variations were associated with HB, NB, and FB of Cyg-like and exhibited Sco-like variation at relatively lower brightness. During the decay phase, the variation closely resembles the soft state of an atoll source and later transits to the hard state of the atoll source just before going to the quiescent state. Many important results were noticed based on the spectral fitting of RXTE data of this source. The mass accretion rate was found to be constant along with the Z phase in Sco-like variation and different mechanisms were proposed to explain the spectral and timing variations during the Z phase variations (Lin et al. 2009). It was also found that mass accretion rate is the important driving parameter during the Z and all along with the atoll phases variation. Z sources are unique probes in the sense they provide a platform to understand the structure of accretion disc emitting close to Eddington luminosity because due to the radiation pressure the structure of the inner region of accretion is affected. The previous studies suggested that the interplay between the accretion disc and comptonization region mutually varies to produce the observed tracks in the HID. However, other physical components like a boundary layer (Popham & Sunyaev 2001) or a transition layer (TL) (Osherovich & Titarchuk 1999a, b; Titarchuk & Osherovich 1999) cannot be ruled out. The comptonization region can be in the form of a quasi-spherical cloud or it could be a base of a jet that causes the observed hard continuum in the X-ray spectrum (Migliari et al. 2007). But its association with dynamical features like various branch oscillations or band-limited noises is not known. The spectra of Z sources can also be explained by a structure known as the boundary layer over the NS surface but again, its association to the observed HBO, NBO, etc., is not properly understood (Popham & Sunyaev 2001; Gilfanov, Revnivtsev & Molkov 2003; Revnivtsev & Gilfanov 2006). Based on the detailed spectral modelling of GX 17 + 2, BL occupies a smaller area at the lower vertex (i.e. bottom of NB) in comparison to its area in other branches (Lin et al. 2012) and the comptonization dominates at the HB branch that fades away as source traverse to the FB. The inner disc radius was found to be moving towards the NS, as the Z track evolves from HB to FB. All these structural and radiative variations are found to be occurring at an almost constant mass accretion rate (Lin et al. 2009, 2012).","Citation Text":["Lin et al. 2009"],"Citation Start End":[[4006,4021]]}
{"Identifier":"2015ApJ...811L..32H__Hellinger_&_Tr\u00e1vn\u00ed\u010dek_2005_Instance_1","Paragraph":"In this Letter, we directly test the relationship between proton kinetic instabilities and plasma turbulence in the solar wind using a hybrid expanding box model that allows us to study self-consistently physical processes at ion scales. In the hybrid expanding box model, a constant solar wind radial velocity vsw is assumed. The radial distance R is then \n\n\n\n\n\n, where R0 is the initial position and \n\n\n\n\n\n is the initial value of the characteristic expansion time \n\n\n\n\n\n Transverse scales (with respect to the radial direction) of a small portion of plasma, comoving with the solar wind velocity, increase \u221d R. The expanding box uses these comoving coordinates, approximating the spherical coordinates by the Cartesian ones (Liewer et al. 2001; Hellinger & Tr\u00e1vn\u00ed\u010dek 2005). The model uses the hybrid approximation where electrons are considered as a massless, charge-neutralizing fluid and ions are described by a particle-in-cell model (Matthews 1994). Here, we use the two-dimensional (2D) version of the code, fields and moments are defined on a 2D x\u2013y grid 2048 \u00d7 2048, and periodic boundary conditions are assumed. The spatial resolution is \u0394x = \u0394y = 0.25dp0, where \n\n\n\n\n\n is the initial proton inertial length (vA0: the initial Alfv\u00e9n velocity, \u03a9p0: the initial proton gyrofrequency). There are 1024 macroparticles per cell for protons that are advanced with a time step \n\n\n\n\n\n, while the magnetic field is advanced with a smaller time step \n\n\n\n\n\n The initial ambient magnetic field is directed along the radial z-direction, perpendicular to the simulation plane \n\n\n\n\n\n, and we impose a continuous expansion in the x- and y-directions. Due to the expansion, the ambient density and the magnitude of the ambient magnetic field decrease as \n\n\n\n\n\n (the proton inertial length dp increases \u221d R; the ratio between the transverse sizes and dp remains constant; the proton gyrofrequency \u03a9p decreases as \u221dR\u22122). A small resistivity \u03b7 is used to avoid accumulation of cascading energy at grid scales; initially, we set \n\n\n\n\n\n (\u03bc0 being the magnetic permittivity of vacuum) and \u03b7 is assumed to be \n\n\n\n\n\n The simulation is initialized with an isotropic 2D spectrum of modes with random phases, linear Alfv\u00e9n polarization (\n\n\n\n\n\n), and vanishing correlation between magnetic and velocity fluctuation. These modes are in the range 0.02 \u2264 kdp \u2264 0.2 and have a flat one-dimensional (1D) power spectrum with rms fluctuations = 0.24 B0. For noninteracting zero-frequency Alfv\u00e9n waves, the linear approximation predicts \n\n\n\n\n\n (Dong et al. 2014). Protons initially have the parallel proton beta \n\n\n\n\n\n and the parallel temperature anisotropy \n\n\n\n\n\n as typical proton parameters in the solar wind in the vicinity of 1 AU (Hellinger et al. 2006; Marsch et al. 2006). Electrons are assumed to be isotropic and isothermal with \u03b2e = 0.5 at t = 0.","Citation Text":["Hellinger & Tr\u00e1vn\u00ed\u010dek 2005"],"Citation Start End":[[748,774]]}
{"Identifier":"2022MNRAS.510.3222S__Kishimoto_et_al._2007_Instance_1","Paragraph":"In order to understand the accretion process and the structure of emitting regions\/locations in AGNs, studies of broad-band spectral energy distributions and their temporal variations are necessary. The X-ray emission is produced in the innermost region or above the accretion disc by hot electrons, i.e. corona (Ramos Almeida & Ricci 2017), where ultraviolet (UV) photons are Compton upscattered (Haardt & Maraschi 1993), which results in a power-law continuum above 2 keV. The modified blackbody emission associated with the Keplerian portion of the accretion disc peaks around optical\/UV (Sanders et al. 1989; Elvis et al. 1994). Some of the emitted UV radiation is reprocessed in a structure often known as \u2018dusty torus\u2019, located around tens of pc scale environment of AGNs, and is re-emitted in infrared\/mid-infrared (MIR) (Suganuma et al. 2006; Kishimoto et al. 2007; Burtscher et al. 2013; Ramos Almeida & Ricci 2017; Leftley et al. 2018). Recent studies based on infrared observations with Very Large Telescope Interferometer suggest that the distribution of dust in the torus need not to be confined in a classical thick toroidal structure but possibly originates from two components, i.e. an equatorial thin disc and a hollow dusty cone towards the polar region (see H\u00f6nig & Kishimoto 2017; Stalevski et al. 2017, 2019). The hollow dusty cone is probably due to the puffed-up disc caused by the infrared radiation pressure, which releases the inflowing gas from the gravitational potential of the black hole (H\u00f6nig 2019). This structure is considered to be an extended polar feature probably from a dusty wind located at the edge of the NLR (L\u00f3pez-Gonzaga et al. 2014; L\u00f3pez-Gonzaga et al. 2016; Honig & Kishimoto 2017; Stalevski et al. 2017; Leftley et al. 2018). The MIR radiation is also considered to be emitted from this non-classical torus structure (Asmus et al. 2016) and the spectrum is often modelled by a warm component with a blackbody temperature around 300 K (e.g. Edelson & Malkan 1986). The orientation of the dusty molecular obscuring torus plays a key role in the characterization of the distinct classes of AGNs in the unification scheme (for more detail see, Antonucci 1993; Urry & Padovani 1995; Netzer 2015).","Citation Text":["Kishimoto et al. 2007"],"Citation Start End":[[851,872]]}
{"Identifier":"2021ApJ...920...96L__Leger_&_Puget_1984_Instance_1","Paragraph":"In addition to being an important participant in the evolution of galaxies, dust absorption and scattering can significantly alter the appearance of its host galaxy. Dust can absorb a large fraction of the total UV and optical energy emitted in a galaxy and re-radiate it in the IR. The absorbed energy can be re-radiated through several physical processes, each with different observational signatures. Dust emission in the far-IR (FIR) comes from equilibrium heating of dust grains and can be approximately described with a featureless modified blackbody function. However, rich spectral features that are commonly attributed to polycyclic aromatic hydrocarbons (PAHs) seen in the mid-IR emission (Leger & Puget 1984) arise predominantly in a nonequilibrium, single photon, process. Both the broad, featureless FIR and the spectral features of the mid-IR correlate with other galaxy properties, providing insight into the underlying physical processes operating within the galaxy. In local starburst galaxies Engelbracht et al. (2008) found that the equivalent width of the 8 \u03bcm aromatic emission complex negatively correlates with the radiation field hardness and positively correlates with metallicity. Gordon et al. (2008) further found that the former correlation is stronger than the latter in M101 H ii regions, suggesting that the variation in the strength of aromatic emissions is due to dust processing, rather than differences in dust formation. The ratio of the total infrared flux (TIR) to the far-UV (FUV) flux correlates with UV spectral slope, \u03b2 (while the correlation holds for a general definition of TIR, FUV, and \u03b2, see e.g., Meurer et al. 1995; Kong et al. 2004 for specific definitions), both across galaxies and within galaxies (Meurer et al. 1995; Calzetti et al. 2005). While the relation is tight for starburst galaxies (Meurer et al. 1999), significant scatter exists for normal star-forming galaxies (Bell 2002). Kong et al. (2004) found that higher \u03b2 values at a given IR to UV ratio in more quiescent star formation galaxies can be explained in terms of the balance between present and past averaged star formation rate (SFR). Given that \u03b2 is commonly measured from the Galaxy Evolution Explorer (GALEX) FUV and near-UV (NUV) bands and that the NUV band captures the 2175 \u212b extinction feature, variation in the 2175 \u212b feature strength can have a profound impact on \u03b2, but separating such variations from the intrinsic behavior of \u03b2 is challenging with empirical methods.","Citation Text":["Leger & Puget 1984"],"Citation Start End":[[700,718]]}
{"Identifier":"2021AandA...651L...8D__Just_et_al._2007_Instance_1","Paragraph":"As the most luminous persistent sources in the Universe, quasars are bright enough to be detected up to redshifts z\u2004>\u20047 (Mortlock et al. 2011; Banados et al. 2018; Wang et al. 2018; Yang et al. 2020). According to the currently accepted model, quasars are extremely luminous active galactic nuclei (AGNs), where the observed intense energy release are related to the accretion of a gaseous disk onto a supermassive black hole (SMBH). Quasars have a wide spectral energy distribution, which normally contains a significant emission component in the optical-UV band LUV, the so-called big blue bump, with a softening at higher energies (Sanders et al. 1989; Elvis et al. 1994; Trammell et al. 2007; Shang et al. 2011). It has long been discussed that there is a nonlinear relationship between LUV and the quasar\u2019s X-ray luminosity LX, parametrized as log(LX) = \u03b3log(LUV)+\u03b2 (Vignali et al. 2003; Strateva et al. 2005; Steffen et al. 2006; Just et al. 2007; Green et al. 2009; Young et al. 2010; Jin et al. 2012). From the theoretical point of view, this relation could be intrinsic since the UV emission is usually thought to originate from the optically thick disk surrounding the SMBH and the X-ray photons are thought to be generated through the inverse-Compton scattering of these disk UV photons by a plasma of hot relativistic electrons (the so-called corona) around the accretion disk. Such a relation is found to be independent of redshift (Lusso & Risaliti 2016), so that it could be used as a distance indicator to estimate cosmological parameters. The initial dispersion of the LUV\u2005\u2212\u2005LX relation is relatively large (\u03b4\u2004\u223c\u20040.35\u22120.4, Just et al. 2007; Young et al. 2010), but after a detailed study, Lusso & Risaliti (2016) suggest that most of the observed dispersion is not intrinsic, but it is rather due to observational effects. By gradually refining the selection technique and flux measurements, Risaliti & Lusso (2019) collected a complete sample of quasars, whose dispersion of the LUV\u2005\u2212\u2005LX relation is smaller than 0.15 dex. The sample of main quasars is composed of 1598 data points in the range from 0.036\u2004\u2004z\u2004\u20045.1. With this sample, they constructed a Hubble diagram of quasars in redshift range of 0.5\u2004\u2004z\u2004\u20045.5, which is in excellent agreement with the analogous Hubble diagram for SNIa in the redshift range of 0.5\u2004\u2004z\u2004\u20041.4. Moreover, this Hubble diagram of quasars has been studied in cosmological applications (Zheng et al. 2020, 2021). Considering that objects at the same redshift should have the same luminosity distance in any cosmology, here we first fit the model-independent cosmography formula that reflects the Hubble relation between the luminosity distance and redshif using the quasar sample, and then we obtained the distance moduli (also the luminosity distance) for GRBs at a given redshift with the best fit results.","Citation Text":["Just et al. 2007"],"Citation Start End":[[936,952]]}
{"Identifier":"2021MNRAS.507.2766S__Connor_&_Couch_2018_Instance_1","Paragraph":"Supernova explosions must be another promising source of gravitational waves. The gravitational waves from a core-collapse supernova have been mainly discussed with the numerical simulations of gravitational collapse of massive star, the bounce of central core, and subsequent accreting matter to the central compact object, i.e. protoneutron star (PNS), produced via supernova (e.g. Murphy, Ott & Burrows 2009; Cerd\u00e1-Dur\u00e1n et al. 2013; M\u00fcller, Janka & Marek 2013; Ott et al. 2013; Yakunin et al. 2015; Kuroda, Kotake & Takiwaki 2016; Andresen et al. 2017; O\u2019Connor & Couch 2018). For instance, the three-dimensional (3D) numerical simulations for the core-collapse supernova tell us the existence of gravitational wave signals related to the oscillations of PNS (or core region) (Kuroda et al. 2016; Andresen et al. 2019; Radice et al. 2019; Mezzacappa et al. 2020), where the frequency increases with time from a few hertz up to kilohertz after corebounce. This signal is originally considered to be the evidence of the Brunt\u2013V\u00e4is\u00e4l\u00e4 frequency at the PNS surface, i.e. the so-called surface gravity (g) mode (Cerd\u00e1-Dur\u00e1n et al. 2013; M\u00fcller et al. 2013), but it may be natural to consider that the signals come from the global oscillations of the PNSs because the Brunt\u2013V\u00e4is\u00e4l\u00e4 frequency is a local value and depends strongly on the definition of PNS radius. In fact, the gravitational wave signals appearing in the numerical simulations can be identified with the global oscillations, such as fundamental (f-mode) oscillations of PNS (Morozova et al. 2018; Sotani et al. 2019a; Sotani & Takiwaki 2020b, c) or the g-mode-like oscillations in the region inside the shock radius (Torres-Forn\u00e9 et al. 2018, 2019a,b). In addition to the ramp up signals, the 3D simulations tell us the existence of another gravitational wave signal with \u223c100 Hz, which seems to be associated with the standing accretion-shock instability (Kuroda et al. 2016; Andresen et al. 2017, 2019; O\u2019Connor & Couch 2018; Radice et al. 2019; Mezzacappa et al. 2020).","Citation Text":["O\u2019Connor & Couch 2018","O\u2019Connor & Couch 2018"],"Citation Start End":[[557,578],[1968,1989]]}
{"Identifier":"2015ApJ...805...82M__Chevalier_1977_Instance_1","Paragraph":"The fate of the early PWN depends on various parameters such as the spin-down power and baryon loading. If the spin-down power is high enough, some two-dimensional simulations suggest that the equatorial wind can be redirected by the anisotropic pressure, and hoop stresses lead to bipolar outflows10\n\n10\nIn this case, the (collimated) wind radius is \n\n\n\n\n\n.\n that could explain GRBs (Bucciantini et al. 2007, 2008; Komissarov & Barkov 2007). If not, we expect a quasi-spherical expanding flow embedded in the expanding stellar material (see Figure 1). Assuming a SN explosion with \n\n\n\n\n\n erg, the SN ejecta expands with its velocity \n\n\n\n\n\n and radius \n\n\n\n\n\n. The early PWN radius Rw also increases non-relativistically, which is given by (e.g., Metzger et al. 2014)\n7\n\n\n\n\n\nfor \n\n\n\n\n\n, otherwise \n\n\n\n\n\n is used. Note that we have used the ejecta density\n8\n\n\n\n\n\nwhere \u03b4 \u223c 0\u20131 is a typical value used in the literature (Kasen & Bildsten 2010; Metzger et al. 2014). The mixture of material allows us to approximate the inner density profile to be reasonably smooth and flat (Chevalier 1977; Chevalier & Fransson 1992). For demonstration, we adopt \n\n\n\n\n\n throughout this work (Kasen & Bildsten 2010; Metzger et al. 2014), and that the radiation pressure is given by \n\n\n\n\n\n. Here \n\n\n\n\n\n is the PWN volume and \n\n\n\n\n\n is the PWN expansion velocity that can be different from \n\n\n\n\n\n In general, Rw is smaller than \n\n\n\n\n\n, and both of \n\n\n\n\n\n and Rw are numerically determined in this work. Roughly speaking, \n\n\n\n\n\n becomes a good approximation for small values of P such that \n\n\n\n\n\n (implying \n\n\n\n\n\n). The ejecta velocity \n\n\n\n\n\n and radius \n\n\n\n\n\n can be determined by\n9\n\n\n\n\n\n\n\n10\n\n\n\n\n\nThe internal energy trapped in the SN ejecta, \n\n\n\n\n\n, is given by\n11\n\n\n\n\n\nwhere \n\n\n\n\n\n is the dynamical time. Since X-ray and gamma-ray emission is expected in month-to-year timescales, we only consider energy injection due to \n\n\n\n\n\n. In the early phase, as in normal SNe, heating by shocks and unstable isotopes such as \n\n\n\n\n\nNi can be relevant. In the later phase, one may assume that late interactions with circumstellar material are negligible, and injections via the \u03b2 decay of \n\n\n\n\n\nNi are irrelevant after their lifetime \n\n\n\n\n\n. Visible photons leave the SN ejecta in the escape time\n12\n\n\n\n\n\nwhere the Thomson optical depth in the ejecta is given by \n\n\n\n\n\n, which is estimated to be\n13\n\n\n\n\n\nwhere \n\n\n\n\n\n, \n\n\n\n\n\n is the mean molecular weight per electron, and mu is the atomic mass unit. See also Equation (45) below. Two of the key parameters, \n\n\n\n\n\n and \n\n\n\n\n\n, can be estimated from the SN peak emission and determination of the ejecta velocity \n\n\n\n\n\n via detailed spectroscopy. Note that the bound\u2013free or bound\u2013bound cross section is much higher at \n\n\n\n\n\n keV energies, and thermal photons are still generated at later times.","Citation Text":["Chevalier 1977"],"Citation Start End":[[1072,1086]]}
{"Identifier":"2017MNRAS.471...80S__Knizia,_Adler_&_Werner_2009_Instance_1","Paragraph":"The electronic computations were performed using the molpro (Molpro 2015, http:\/\/www.molpro.net) package. In a preliminary work, we used the complete active space self-consistent field (Knowles & Werner 1985; Werner & Knowles 1985) to examine the electronic wavefunction of the HNCO\u2013He complex. These computations showed that this wavefunction is dominantly described by a unique electron configuration (with a weight \u22650.93) over the grid used for the generation of this PES. This justifies hence the use of monoconfigurational ab initio methods. Accordingly, we applied the recently established methodology by Hochlaf and co-workers for mapping multidimensional PESs of weakly bound molecular systems with high accuracy and relatively low computational cost (Lique, Klos & Hochlaf 2010; Halvick et al. 2011; Ajili et al. 2013; Mathivon, Linguerri & Hochlaf 2013; Kalugina et al. 2014; Mogren Al Mogren et al. 2014). Briefly, these electronic computations were carried out using the explicitly correlated coupled cluster method with single, double and perturbative treatment of triple excitations (CCSD(T)-F12) (Adler, Knizia & Werner 2007; Knizia, Adler & Werner 2009) in connection with the augmented correlation-consistent aug-cc-pVTZ basis set of Dunning and co-workers (Dunning 1989; Kendall, Dunning & Harrison 1992). In addition, molpro default choices for the density fitting and resolution of identity basis sets have been applied (Yousaf & Peterson 2008). Benchmarks by Hochlaf and co-workers (Lique et al. 2010; Halvick et al. 2011; Ajili et al. 2013; Mathivon et al. 2013; Kalugina et al. 2014; Mogren Al Mogren et al. 2014) showed that results obtained from this highly correlated approach are close to those deduced using standard coupled cluster techniques extrapolated to the complete basis set (CBS) limit, whereas a strong reduction in CPU time and disc occupancy are observed. For illustration, we performed CCSD(T)\/aug-cc-pVXZ calculations (X = D, T, Q, 5) on the HNCO\u2014He cluster. Then the energies were extrapolated to the CBS limit. The comparison between CCSD(T)-F12\/aug-cc-pVTZ and CBS calculations is given in Table 1. It shows that the CCSD(T)-F12\/aug-cc-pVTZ results are off by 4 per cent (at the maximum) with those deduced from CBS extrapolation. We compare also our results with those done using the CCSD(T)\/aug-cc-pV5Z approach. We can clearly see that the CCSD(T)-F12\/aug-cc-pVTZ approach offers a good agreement with the CCSD(T)\/aug-cc-pV5Z calculations with a very reduced computational cost.","Citation Text":["Knizia, Adler & Werner 2009"],"Citation Start End":[[1141,1168]]}
{"Identifier":"2022ApJ...926..118B__Chatterjee_et_al._2010_Instance_1","Paragraph":"The results of this first exploratory work pave the way to a series of future investigations in which we will broaden the range of initial conditions and study their impact on the empirical parameters defined here. By changing the initial structural properties of the simulated clusters, we will compare the three parameters determined in stellar systems that, after one Hubble time of evolution, have reached different dynamical states. We will also include populations of primordial binaries, which are known to halt the core contraction earlier in the cluster evolution and at lower concentrations (see, e.g., Vesperini & Chernoff 1994; Trenti et al. 2007; Chatterjee et al. 2010). The main differences between the values of A\n5, P\n5, and S\n2.5 in simulations with and without primordial binaries are expected during the advanced phases of the evolution, toward CC and post-CC, when the milder contraction of the clusters with primordial binaries might lead to a different and\/or less extreme evolution of these parameters. The study presented here will also be further extended to explore the effects of different retention fractions of dark remnants (neutron stars and black holes; see, e.g., Alessandrini et al. 2016; Giersz et al. 2019; Kremer et al. 2020, 2021; Gieles et al. 2021, for some studies on the dynamical effects of dark remnants). Finally, a forthcoming paper will be dedicated to building nCRDs and determining the three parameters here defined in a sample of observed star clusters. This requires photometric observations (i) with a high enough angular resolution to resolve individual stars even in the innermost cluster regions, (ii) sampling each system at least out to 0.5 \u00d7 r\n\nh\n, and (iii) deep enough to reach a few magnitudes below the MSTO. These requirements are achieved by most HST and adaptive-optics-assisted observations currently available for many GCs, thus making the determination of the three parameters from observations relatively straightforward, although particular care is needed to deal with typical observational difficulties such as photometric incompleteness, differential reddening, and Galactic field contamination. We will discuss the relation between these parameters and other dynamical indicators (in particular, the A\n+ parameter measured from BSSs; see the Introduction), thus providing quantitative assessments of the operational ability of the three nCRD diagnostics to distinguish dynamically young GCs from systems in advanced states of dynamical evolution.","Citation Text":["Chatterjee et al. 2010"],"Citation Start End":[[660,682]]}
{"Identifier":"2016ApJ...828...83G__Guo_et_al._2012_Instance_1","Paragraph":"NLFFF models have been widely adopted to study magnetic field structures in the solar atmosphere, for instance, magnetic flux ropes (Canou et al. 2009; Su et al. 2009; Canou & Amari 2010; Cheng et al. 2010; Guo et al. 2010a, 2010b; Jing et al. 2010), magnetic null points (Zhang et al. 2012; Sun et al. 2013, 2014), and quasi-separatrix layers (Savcheva et al. 2012, 2015; Guo et al. 2013; Zhao et al. 2014; Yang et al. 2015). Most of them are modeled in Cartesian coordinates, although some models are reconstructed in spherical geometry (Su et al. 2009; Tadesse et al. 2011; Guo et al. 2012) or with tetrahedral meshes (Amari et al. 2014b). In particular, Amari et al. (2014a) has developed another code to fulfill the need for reconstructions on the scale of local active regions within a global extrapolation, using an iterative Grad\u2013Rubin scheme adapted to spherical coordinates. Such a state-of-the-art code resolves an active region in a global model with \n\n\n\n\n\n grid points. This contrasts to most earlier modeling efforts, which were limited either to a small field of view (for Cartesian coordinates) or to lower spatial resolution (for spherical coordinates) due to limited computational resources. The data volume and the computation time increase for NLFFF modeling with a large field of view and high spatial resolution simultaneously, which implies a big data problem. On the other hand, we have to use data with high spatial resolution to resolve the small-scale features of the magnetic field, such as small flux tubes and electric current channels. They are crucial for reconstructing the non-potential magnetic field. We also have to consider a field of view as large as possible to include remote magnetic field connections (DeRosa et al. 2009). Since only the bottom boundary is available at present for realistic NLFFF modeling, the magnetic field concentration should be isolated to mitigate any effects of lateral boundary conditions (Wiegelmann et al. 2006). An isolated magnetic field region has magnetic field lines originating from the bottom boundary and connecting back to the bottom. A large field of view is required to fulfill this requirement of having an isolated magnetic domain. When the field of view is large, the curvature of the Sun cannot be neglected and spherical coordinates are recommended (e.g., Guo et al. 2012; Yeates 2014).","Citation Text":["Guo et al. 2012","Guo et al. 2012"],"Citation Start End":[[577,592],[2343,2358]]}
{"Identifier":"2018ApJ...852...45W__Stern_&_Poutanen_2011_Instance_1","Paragraph":"Based on indirect evidence, there are two main candidates for the \u03b3-ray emitting region. The first one is close to the BH and the \u03b3-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 \u03b3-ray emission region through the magnification factor during the \u03b3-ray variability in some gravitationally lensed \u03b3-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 \u03b3-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 \u03b3-ray production is beyond the BLR (e.g., B\u0142a\u017cejowski 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 \u03b3-rays are produced 10\u201320 pc from the BH (e.g., Larionov et al. 2008; Sikora et al. 2008), where it was assumed that \u03b3-ray and radio emission is triggered by shocks propagating along a relativistic jet. The time delay between flares in radio and \u03b3-ray bands combined with the VLBI high-resolution observation on the size of radio core also indicated that the \u03b3-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 \u03b3-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 \u03b3-ray emitting region, which may be dominated by an accretion disk, BLR, infrared torus, and cosmic background, respectively, if the \u03b3-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 \n\n\n\n\n\n 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 Text":["Stern & Poutanen 2011"],"Citation Start End":[[915,936]]}
{"Identifier":"2016AandA...591A..91P__Halpern_et_al._2014_Instance_1","Paragraph":"MSH 11-61A (also known as G290.1-0.8) is a mixed morphology SNR detected from radio to soft X-rays (up to ~3 keV) that was formed by the core collapse of a massive progenitor star (mass \u227325 M\u2299;  Filipovi\u0107 et al. 2005; Reynoso et al. 2006; Garc\u00eda et al. 2012; Kamitsukasa et al. 2015; Auchettl et al. 2015a). Following these authors, the distance to the SNR is in the range 6\u221211 kpc; the most recently determined values converge towards 7 \u00b1 1 kpc. We adopt a distance of 7 kpc throughout the paper1. The INTEGRAL source IGR J11014-6103 is located close to MSH 11-61A and is powered by PSR J1101-6101 (Pavan et al. 2011, hereafter Paper I; Tomsick et al. 2012; Pavan et al. 2014, hereafter Paper II; Halpern et al. 2014). The pulsar shows spin-down parameters typical for pulsars of its age: a period of P = 62.8 ms and a pulse period derivative \u1e56 = (8.56 \u00b1 0.51) \u00d7 10-15 s s-1. The estimated spin-down energy is \u0116 = 1.36 \u00d7 1036 erg s-1 and the surface dipolar magnetic field is 7.4 \u00d7 1011 G (Halpern et al. 2014). Previous Chandra observations aimed at the INTEGRAL source showed that PSR J1101-6101 simultaneously powers several outflows: an X-ray and radio PWN, shaped in a narrow cone elongated towards the parent SNR, and an X-ray jet and counter-jet, both oriented nearly perpendicular to the PWN axis (Tomsick et al. 2012; Paper II). The main jet extends for nearly 5\u2032 in the sky, which corresponds to a projected length of ~11 pc, and showed a remarkable helicoidal pattern (see Paper II). Already in the data set analysed in Paper II, indications for a spatial deviation from the helical pattern of the main jet were noticed at a distance of ~50\u2033 from the pulsar. At this position the surface brightness of the jet was low, forming what looked like a gap, but its brightness profile was compatible with expectations of Doppler-deboosting in the jet-helix model. The spatial deviation was therefore not considered significant at the time because the data were hampered by the presence of CCD chip gaps, resulting in only 50% effective exposure in that region. The counter-jet was detected at 3.7\u03c3 in the Chandra image and its flux was estimated to be ~5% that of the main jet. The conical shape of the PWN in IGR J11014-6103 was ascribed to the supersonic motion of PSR J1101-6101 in the ISM (Tomsick et al. 2012). ","Citation Text":["Halpern et al. 2014"],"Citation Start End":[[698,717]]}
{"Identifier":"2021MNRAS.501.4035R__Rubin_et_al._2009_Instance_1","Paragraph":"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; H\u00e4berli 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\u2013Gerasimenko 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\u2013Gerasimenko 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\u2009cent), 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 Text":["Rubin et al. 2009"],"Citation Start End":[[678,695]]}
{"Identifier":"2019MNRAS.490..243C__Zaldarriaga,_Furlanetto_&_Hernquist_2004_Instance_1","Paragraph":"The main challenge in detecting the cosmological H\u2009i 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\u2013free 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\u2009i signal (Zaldarriaga, Furlanetto & Hernquist 2004; Bharadwaj & Ali 2005; Jeli\u0107 et al. 2008; Bernardi et al. 2009; Jeli\u0107 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\u2009i 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 \u2018smoothness\u2019 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\u2009i brightness temperature fluctuation, a cylindrical (2D) power spectrum, P(k\u22a5, k\u2225) is a useful diagnostic in terms of foreground avoidance. Spectral smoothness of foregrounds confines the majority of foreground power to low k\u2225 modes, resulting in \u2018foreground wedge\u2019. In the foreground avoidance technique, the EoR signal is searched for outside this wedge, in the so-called \u2018EoR window\u2019 (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 Text":["Zaldarriaga, Furlanetto & Hernquist 2004"],"Citation Start End":[[771,811]]}
{"Identifier":"2016MNRAS.460.4220L__Zhang,_Luo_&_Wang_2015a_Instance_1","Paragraph":"The oEA stars have pulsation features similar to classical \u03b4 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\u2009F = 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\u2009Ppul = \u22120.61log\u2009F + 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\u2009F for KIC 6220497 matches well with those for the other oEA stars with similar orbital periods, while the surface gravity log\u2009g1 = 3.78 for the primary component is clearly smaller. The period\u2013gravity 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 Text":["Zhang, Luo & Wang 2015a"],"Citation Start End":[[1293,1316]]}
{"Identifier":"2020MNRAS.493.5413K__Fryxell_et_al._2000_Instance_1","Paragraph":"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 \u03c10,7 \u2248 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 \u03c10,7 = 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 Text":["Fryxell et al. 2000"],"Citation Start End":[[690,709]]}
{"Identifier":"2021MNRAS.504.3316B__than_2000_Instance_1","Paragraph":"WASP-43b is the most heavily scrutinized phase curve, with four analyses of this data set already published (Stevenson et al. 2017; Mendon\u00e7a 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\u2019s phase offset and nightside temperature. Stevenson et al. (2017) initially reported only a 2\u03c3 upper limit on the nightside temperature of 650\u2009K, while all subsequent reanalyses (including ours) favour a significantly detectable nightside temperature of \u223c800\u2009K. As for the planet\u2019s phase offset, Stevenson et al. (2017) and May & Stevenson (2020) favour a larger phase offset (21 \u00b1 2\u2009\u00b0E) than Mendon\u00e7a et al. (2018) and Morello et al. (2019) (12 \u00b1 3\u2009\u00b0E and 11 \u00b1 2\u2009\u00b0E). 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 Mendon\u00e7a 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\u2009000 CPU hours (or 434\u2009d using our 12\u00d7 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 Mendon\u00e7a et al. (2018) and Morello et al. (2019). Ultimately, we cannot decide between these two discrepant offsets as the \u0394BIC between the two phase curve models for our preferred BLISS detector model is only 3.7 (insignificantly favouring the 20.4 \u00b1 3.6 offset from the v2 model). For reference, Stevenson et al. (2014b) found phase offsets ranging from roughly \u22126 to 17\u2009deg east in the Hubble\/WFC3 bandpass.","Citation Text":["Stevenson et al. 2017"],"Citation Start End":[[109,130]]}
{"Identifier":"2022ApJ...925..123N__Wang_&_Frenklach_1997_Instance_1","Paragraph":"Benzene (C6H6), the simplest aromatic hydrocarbon, is a molecule that has raised great interest in the astrophysical community for almost four decades. This is mainly because C6H6 is one of the main precursors of polycyclic aromatic hydrocarbons (PAHs) reported to be present in interstellar dust particles (Leger & Puget 1984; Allamandola et al. 1989; Tielens 2013 and references therein), carbonaceous chondrites (Pering & Ponnamperuma 1971; Hayatsu et al. 1977; Hahn et al. 1988), and other astrophysical environments, such as carbon-rich, high-temperature environments (circumstellar and carbon-rich protoplanetary nebulae; Buss et al. 1993; Clemett et al. 1994). Benzene rings easily produce more complex, polycyclic structures by the one-ring build-up mechanism (Simoneit & Fetzer 1996). In space, an analogous process to carbon soot formation occurring on Earth can be initiated through the completion of that first aromatic ring and may also lead to the synthesis of PAHs (Tielens & Charnley 1997). Mechanisms involving the addition of hydrocarbons, such as acetylene onto aromatic rings as well as the attachment of other aromatic rings, or hydrocarbon pyrolysis, have been proposed to characterize the growth process of PAHs (Bittner & Howard 1981; Frenklach & Feigelson 1989; Wang & Frenklach 1997; Cherchneff 2011 and references therein). PAH synthesis from shocked benzene has also been reported (Mimura 1995). PAHs are crucial materials involved in a variety of cosmochemical processes. For example, amino acids could be synthesized by aqueous alteration of precursor PAHs in carbonaceous chondrites (Shock & Schulte 1990). PAHs are also produced in laboratory-simulated planetary atmospheres of Titan and Jupiter (Sagan et al. 1993; Khare et al. 2002; Trainer et al. 2004), and results from these studies indicate that the formation of aromatic rings and polyaromatics may be, among other sources, a possible chemical pathway for the production of the atmospheric solid particles (Lebonnois et al. 2002; Wilson et al. 2003; Trainer et al. 2004). The formation and evolution of benzene in planetary environments or other solar system objects thus represents a fundamental primary stage of the PAH production and other subsequent relevant chemical and prebiotic processes (like soot formation). In this context, several works related to benzene have been devoted to better understand the physico-chemical processes of irradiated C6H6, in its gaseous and solid phases, and the derived products, by acquiring high-resolution astronomical spectra, carrying out detailed laboratory studies or developing theoretical modeling (Allamandola et al. 1989  and references therein; Callahan et al. 2013; Materese et al. 2015; Mouzay et al. 2021). Laboratory astrophysical investigations have mostly focused on performing vibrational spectroscopy of ion, electron, or UV irradiated C6H6 gas and C6H6 ice. Such investigations aim to provide data on the spectral properties of the irradiated C6H6 materials, compare them with spectra obtained from astronomical observations (e.g., observations of the interstellar medium), or to study photoprocessed benzene ices to understand the fate of benzene ices in Titan\u2019s stratosphere and help understanding the formation of aerosol analogs observed in Saturn\u2019s moon\u2019s stratosphere (Mouzay et al. 2021).","Citation Text":["Wang & Frenklach 1997"],"Citation Start End":[[1287,1308]]}
{"Identifier":"2022MNRAS.516..636S__Kurucz_1993_Instance_1","Paragraph":"We obtained atmospheric parameters using the high-resolution spectra. For our candidates with early spectral types: CoRoT\u201334, CoRoT 310204242, and CoRoT 110660135 we derived the effective temperature Teff, surface gravity log\u2009(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 \u03c72 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\u2009dex for log\u2009(g)sp, of 0.1\u2009dex for [Fe\/H], and 2\u2009 per\u2009cent 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\u201334, 35, and 36, we used the stellar mass and radius as input for the light-curve fit, which we used to constrain log\u2009(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\u2009(g)lc as prior. The best-fitting evolutionary tracks are shown in Fig. 2.","Citation Text":["Kurucz 1993"],"Citation Start End":[[718,729]]}
{"Identifier":"2017ApJ...849..140X__Zavagno_et_al._2006_Instance_1","Paragraph":"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 Text":["Zavagno et al. 2006"],"Citation Start End":[[824,843]]}
{"Identifier":"2021ApJ...923..233G__Carlton_et_al._2011_Instance_2","Paragraph":"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 \n\n\n\n100\u2032\u2032\n\n 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 \u00d7 1022 cm\u22122 (Reynolds et al. 2009), implying A\n\nV\n \u223c 23\nm\n, so radio and X-rays are the only useful observational channels. The angular expansion rate of 0.64 arcsec yr\u22121 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\u2013dominated 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 \u223c14,000 km s\u22121 confirming the large expansion proper motion, refined with a second Chandra observation (Carlton et al. 2011). The distance is still uncertain; the high column\u2014higher than the entire Galactic column along nearby sight lines\u2014suggests 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 Text":["Carlton et al. 2011"],"Citation Start End":[[1647,1666]]}
{"Identifier":"2016MNRAS.461.3982B__Merline_et_al._2002_Instance_1","Paragraph":"Many studies have been done to understand the dynamics and origin of such systems since the discovery of the first binary asteroid system, Dactyl orbiting around (243) Ida in 1993 (Chapman et al. 1995). Based on the structure of \u2018rubble pile\u2019 asteroids (a collection of gravitationally bound boulders with a distribution of size scales and very little tensile strength between them), a model for how they can disrupt due to close flybys of a planet was developed. However, close encounters with the planets proved not to be enough for creation of the current population of binary systems (Margot et al. 2002; Walsh & Richardson 2008). Another model for their formation is by increasing their spin rates due to incident and remitted solar photons, known as the Yarkovsky\u2013O'Keefe\u2013Radzievskii\u2013Paddack (YORP) effect. The YORP effect on contact binary asteroids has been studied (Bottke et al. 2002; Merline et al. 2002; Scheeres 2002; Walsh & Richardson 2006). Using a model with an ellipsoid and a sphere in a planar case, Scheeres (2007) studied fission limits (spin limit to occur a fission) and the stability of that kind of system for different initial conditions. After that, the stability of a binary system was analysed using a two-ellipsoid model (Scheeres 2009). Pravec et al. (2010) made a complete study about formation of asteroid pairs through rotation fission. Jacobson & Scheeres (2011) studied the creation of binaries and other observed near-Earth asteroid (NEA) systems, including doubly synchronous binaries, high-e binaries, ternary systems and contact binaries. That study analysed the dynamics of a binary system just after rotational fission. Using a two-ellipsoid model taking into account mutual gravitational interactions and tidal dissipation, they analysed the dynamics for different mass ratios of the system under a planar assumption. The current work follows from these results, but looks at more likely, non-planar initial configurations. This extension is significant, as non-planar cases must take into account the complete rotational motion (rotation, precession and nutation) of each body. Our results are compared with the results obtained by Jacobson & Scheeres (2011).","Citation Text":["Merline et al. 2002"],"Citation Start End":[[895,914]]}
{"Identifier":"2021MNRAS.508.5425M__Schruba,_Kruijssen_&_Leroy_2019_Instance_1","Paragraph":"In this paper, we use a legacy value data set covering the inner star-forming disc of NGC 300 taken with the VLT\/MUSE instrument (Bacon et al. 2010), consisting of a contiguous 7 arcmin \u00d7 5 arcmin mosaic (\u223c4 kpc \u00d7 3 kpc) and covering the galaxy out to galactocentric radii of about 0.45R25 (with R25 \u223c 5.33 kpc being the optical radius; Paturel et al. 2003). NGC 300 is an ideal target to study stellar feedback: it is the closest (D \u223c 2 Mpc; Dalcanton et al. 2009), non-interacting, star-forming disc galaxy that can be mapped at the necessary spatial resolution [i.e. 1 arcsec, which corresponds to \u223c10 pc, resolving individual star-forming regions and supernova remnants (SNRs)]. The large spatial coverage (to cover most of the star-forming disc) available not only in the optical with MUSE but throughout the electromagnetic spectrum (e.g. Helou et al. 2004; Westmeier, Braun & Koribalski 2011; Riener et al. 2018; Kruijssen et al. 2019; Schruba, Kruijssen & Leroy 2019) makes NGC 300 the ideal target for simultaneous resolved feedback, stellar population, and ISM studies. Closer galaxies like the Magellanic Clouds, M31, or M33 do not allow similar large-scale multiwavelength mapping due to their large angular sizes, while more distant galaxies (beyond a few Mpc) do not allow multiwavelength studies with similar spatial resolution across the optical, infrared, mm\/sub-mm, and radio. NGC 300 perfectly bridges between \u223c100 pc resolution IFU surveys of nearby galaxies like PHANGS, and upcoming (sub-)\u2009pc scale resolution IFU surveys of the Milky Way and the Magellanic Clouds like SDSS-V\/LVM (Kollmeier et al. 2017). Further, NGC 300 offers a favourable inclination of \u223c40\u00b0 (Puche, Carignan & Bosma 1990), it is actively forming stars at a rate between \u223c0.08 and \u223c0.30 M\u2299 yr\u22121 (Kang et al. 2016, and references therein), and has a well-studied population of ${\\rm H\\, \\small {II}}$ regions (e.g. Deharveng et al. 1988; Bresolin et al. 2009; Faesi et al. 2014), planetary nebulae (PNe; e.g. Soffner et al. 1996; Pe\u00f1a et al. 2012; Stasi\u0144ska et al. 2013), and SNRs (e.g. Blair & Long 1997; Millar, White & Filipovic 2012; Vu\u010deti\u0107, Arbutina & Uro\u0161evi\u0107 2015).","Citation Text":["Schruba, Kruijssen & Leroy 2019"],"Citation Start End":[[943,974]]}
{"Identifier":"2021MNRAS.502.4794N__Dullo_&_Graham_2012_Instance_2","Paragraph":"Alongside the above theoretical uncertainties in the physics of core formation in ellipticals, there have also been observational challenges. In particular, determining the size of the core has proven to be a non-trivial task. The light profiles of ellipticals are well described by the 3-parameter S\u00e9rsic profile (S\u00e9rsic 1963, 1968) over a large radial range. The most luminous ellipticals, however, show a departure from the S\u00e9rsic law in their central regions, at a radius widely known as the \u2018break\u2019 or \u2018core\u2019 radius. In these galaxies, the profiles break downward from the inward extrapolation of the outer S\u00e9rsic law. Initially the core size of a galaxy was determined by fitting the so-called \u2018Nuker-profile\u2019 (Lauer et al. 1995) to the surface brightness profile, a method that however depends sensitively on the radial fitting range and yields unreliable results when fit to surface brightness profiles with a large radial extent (e.g. Graham et al. 2003; Dullo & Graham 2012). In more recent years, it has become customary to incorporate a central flattening in the light profile by adopting a 6-parameter core-S\u00e9rsic profile (Graham et al. 2003; Trujillo et al. 2004) which provides a reliable measurement of the core size even over a large radial range (e.g. Dullo & Graham 2012, 2013, 2014). Furthermore, it has been shown that adopting a multicomponent model rather than a single core-S\u00e9rsic model over the entire radial range provides a more reliable estimate of the core size (Dullo & Graham 2014; Dullo 2019). Measured core sizes for massive ellipticals \u2013 derived in this way \u2013 are typically tens to a few hundred parsecs (e.g. Dullo & Graham 2014), while cores larger than $1{\\, \\mathrm{kpc}}$ are rare. A study by Lauer et al. (2007) considered a large sample of brightest cluster galaxies (BCGs) and found that fewer than 10 systems had a core size of $\\sim 1\\rm {kpc}$ or greater, with the largest cored system being NGC 6166 which has a core size of $\\sim 1.5{\\, \\mathrm{kpc}}$. More recently, Dullo (2019) considered the largest sample of \u2018large-core\u2019 galaxies to date, finding that only 13(7) galaxies have core sizes larger than $0.5(1){\\, \\mathrm{kpc}}$.","Citation Text":["Dullo & Graham 2012"],"Citation Start End":[[1270,1289]]}
{"Identifier":"2017MNRAS.466.3323L__Wu_et_al._2015_Instance_1","Paragraph":"Parameter  is likely different for each AGN and may depend on the accretion rate. We have 0.3 for Mrk 110 (see Fig.1). The lines parallel to the best-fitting line to Mrk 110 cannot connect two points of Mrk 493 or IRAS 04416 even considering the corresponding errors. Mrk 110, Mrk 493 and IRAS 04416 have $\\dot{\\mathscr{M}}_{f=1}=$ 5.89, 75.9 and 426.6 or $\\dot{\\mathscr{M}}_{\\rm {f}}=$ 0.49, 1.04 and 7.27, respectively. Mrk 493 and IRAS 04416 have larger accretion rates than Mrk 110 and it is possible that  depends on the accretion rate, but the present data do not allow us to test this hypothesis. There is a strong positive correlation between f and $\\dot{\\mathscr{M}}_{f=1}$, $f=6.8(\\pm 1.6)\\dot{\\mathscr{M}}_{f=1}^{0.4\\pm 0.1}$ (see Fig.3), which is consistent with f being dominated by the radiation pressure. This correlation supports the explanation of $f=5.4 r_{\\rm {BLR}}^{0.3}$ in Mrk 110 arising due to the radiation pressure from the accretion disc. Thus, the relation $f=6.8\\dot{\\mathscr{M}}_{f=1}^{0.4}$ reflects the physical essence of relations like $f=5.4 r_{\\rm {BLR}}^{0.3}$. The radiation pressure will produce more obvious effects on the BLR clouds as  increases. As the radiation pressure vanishes,  vanishes. The larger values of f  816 make MRM increase for Mrk 110. In like manner, larger f values may exist in quasars. Quasar J0100+2802 at z6.30, the most luminous quasar known at z6, has MRM(f1)1.21010M and Lbol1.621048ergs1 (Wu et al. 2015). It has LEdd  1.5  1038M\/M  1.8  1048 erg s1 for solar composition gas and Lbol\/LEdd0.9. $\\dot{\\mathscr{M}}_{f=1}= L_{\\rm {bol}}\/L_{\\rm {Edd}}\/\\eta$ (Du etal.2014), where  is the efficiency of converting rest-mass energy to radiation (Thorne1974) and in general,  is of the order of 0.1. Thus $\\dot{\\mathscr{M}}_{f=1}\\sim 9$ and f16 for J0100+2802. As f16, MRM1.91011M and Lbol\/LEdd0.06. The larger black hole mass further gives rise to the most significant challenge to the Eddington limit growth of black holes in the early Universe (Willott etal.2010; Volonteri2012). A high f value is suggested for PG 1247+267 by the ultraviolet RM of carbon lines (Trevese etal.2014). PG 1247+267 has L(1350)3.91047ergs1 and ionization stratification like low-luminosity AGNs. The broad H has a redward shift of 0.008 with respect to [Oiii] 5007 (McIntosh etal.1999b), and has vFWHM7460kms1 (McIntosh etal.1999a). If zg0.008, f  13, consistent with the high f value suggested in Trevese etal. (2014). Thus, the new method is capable of estimating f in quasars. For the Eddington ratio Lbol\/LEdd(f)Lbol\/fLEdd(f1)1, we have $\\dot{\\mathscr{M}}_{f=1}=1133$, f113 and $\\dot{\\mathscr{M}}_{f}= 10$ for 0.1. Then, we have Lbol\/LEdd113 when assuming f1. That we do not often see sources with such (apparent) Eddington ratios may suggest that the Thomson cross-section typically used in the radiation pressure calculation underestimates the coupling between the radiation and the BLR gas. On considering the line-driven radiation pressure (Castor, Abbott  Klein1975), the radiation pressure due to the gas opacity will be 103 times that due to the electron scattering opacity (Ferland etal.2009).","Citation Text":["Wu et al. 2015"],"Citation Start End":[[1459,1473]]}
{"Identifier":"2022ApJ...926..118B__Giersz_et_al._2019_Instance_1","Paragraph":"The results of this first exploratory work pave the way to a series of future investigations in which we will broaden the range of initial conditions and study their impact on the empirical parameters defined here. By changing the initial structural properties of the simulated clusters, we will compare the three parameters determined in stellar systems that, after one Hubble time of evolution, have reached different dynamical states. We will also include populations of primordial binaries, which are known to halt the core contraction earlier in the cluster evolution and at lower concentrations (see, e.g., Vesperini & Chernoff 1994; Trenti et al. 2007; Chatterjee et al. 2010). The main differences between the values of A\n5, P\n5, and S\n2.5 in simulations with and without primordial binaries are expected during the advanced phases of the evolution, toward CC and post-CC, when the milder contraction of the clusters with primordial binaries might lead to a different and\/or less extreme evolution of these parameters. The study presented here will also be further extended to explore the effects of different retention fractions of dark remnants (neutron stars and black holes; see, e.g., Alessandrini et al. 2016; Giersz et al. 2019; Kremer et al. 2020, 2021; Gieles et al. 2021, for some studies on the dynamical effects of dark remnants). Finally, a forthcoming paper will be dedicated to building nCRDs and determining the three parameters here defined in a sample of observed star clusters. This requires photometric observations (i) with a high enough angular resolution to resolve individual stars even in the innermost cluster regions, (ii) sampling each system at least out to 0.5 \u00d7 r\n\nh\n, and (iii) deep enough to reach a few magnitudes below the MSTO. These requirements are achieved by most HST and adaptive-optics-assisted observations currently available for many GCs, thus making the determination of the three parameters from observations relatively straightforward, although particular care is needed to deal with typical observational difficulties such as photometric incompleteness, differential reddening, and Galactic field contamination. We will discuss the relation between these parameters and other dynamical indicators (in particular, the A\n+ parameter measured from BSSs; see the Introduction), thus providing quantitative assessments of the operational ability of the three nCRD diagnostics to distinguish dynamically young GCs from systems in advanced states of dynamical evolution.","Citation Text":["Giersz et al. 2019"],"Citation Start End":[[1224,1242]]}
{"Identifier":"2020AandA...635A..81P__simulations,_Georgobiani_et_al._(2003)_Instance_1","Paragraph":"Furthermore, Duvall et al. (1993) noticed an inversion of the sense of asymmetry between spectrometric and photometric measurements, with line profiles in the velocity spectrum featuring more power in their low-frequency wing than in their high-frequency wing and vice-versa for line profiles in the intensity spectrum. Since intensity perturbations were expected to be proportional to velocity perturbations, one would have expected the asymmetries to be the same. Many hypotheses were posited to explain this puzzling result. Duvall et al. (1993) suggested that it was due to non-adiabatic effects lifting the proportionality relationship between the two kinds of perturbations (fluid displacement and temperature) but this hypothesis was later contradicted by Rast & Bogdan (1998). Non-adiabaticity was brought up again later on by Georgobiani et al. (2003) who suggested that the explanation resided in radiative transfer between the mode and the medium. Indeed, the observed radiation temperature corresponds to the gas temperature at local optical depth \u03c4\u2004=\u20041. But optical depth depends on opacity, which non-linearly depends on temperature. Therefore, the temperature fluctuations due to the oscillating mode entails opacity fluctuations, which in turn impacts the \u201cobserved\u201d radiation temperature. Given the non-linear nature of the \u03ba\u2005\u2212\u2005T relation, this modulation decreases the observed temperature fluctuations more significantly in the low-frequency wing of the mode than in its high-frequency wing. Since this radiative transfer does not impact the velocity measurements, this could explain the asymmetry reversal between velocity and intensity spectra. Using 3D simulations, Georgobiani et al. (2003) computed mode line profiles in both the velocity and the intensity power spectrum alternatively at mean unity optical depth and instantaneous unity optical depth. Their results indeed show that the modulation of the \u201cobserved\u201d intensity fluctuations due to radiative transfer could be significant enough to reverse the sense of mode asymmetry. One of the hypothesis enjoying the most support for asymmetry reversal, however, is based on the effect of turbulent perturbations partially correlated with the mode, which thus impact its line profile (Nigam et al. 1998; Roxburgh & Vorontsov 1997; Rast & Bogdan 1998; Kumar & Basu 1999). Indeed, a part of these perturbations is coherent with the mode and, thus, leads to interference. This interference term may be constructive or destructive, depending on the phase difference between the mode and the coherent turbulent perturbations. For frequencies at which the interference is constructive, the power spectral density is slightly elevated, whereas it drops slightly for frequencies at which it is destructive. Typically, in the vicinity of a resonant mode, the dependence of the phase difference between mode and turbulent perturbation is such that the interference term is constructive for frequencies located in one wing of the mode and destructive in the other. Therefore, as a result of this interference behaviour, one of the wings falls off more slowly and the other more rapidly, leading to mode asymmetry. It has been suggested that the degree of correlation between the turbulent perturbations and the oscillation it excites is higher in intensity than in velocity, so that it changes the sign of mode asymmetry only in the intensity spectrum. While it is widely accepted that correlated turbulent perturbations must be taken into account to explain asymmetries in the intensity spectrum, the question of whether it has a significant impact on the velocity spectrum remains an open issue (see e.g. Jefferies et al. 2003).","Citation Text":["Georgobiani et al. (2003)"],"Citation Start End":[[835,860]]}
{"Identifier":"2021AandA...655A..99D__Carigi_et_al._2005_Instance_1","Paragraph":"Another way of obtaining information about the nucleosynthesis processes involved in producing carbon is to compare it with other elements that are characterised by a well-known source of production, as in the case of oxygen. In Fig. 5, we show the variation of [C\/O] as a function of [Fe\/H], which serves as a first-order approximation to the evolution with time. To calculate the [C\/O] ratios, two oxygen abundance indicators are used independently. At subsolar metallicities, the abundance ratios with both oxygen indicators are mostly negative and show an increasing trend towards higher metallicity. This is explained by the fact that oxygen is entirely produced by SNe Type II from massive progenitors, which started to release theiryields at earlier ages in the Galaxy and, hence, at lower metallicities (e.g. Woosley & Weaver 1995). The massive stars producing carbon at low metallicities might be less massive than those producing oxygen (i.e. having a longer life), explaining a delayed contribution of carbon, hence, the negative [C\/O] ratios. Alternatively, this could be explained by increasing O\/C yields for more massive progenitors of SNeII. Once metallicity starts to increase, low- and intermediate-mass stars release carbon and massive stars start to eject more carbon than oxygen (Carigi et al. 2005). The [C\/O] ratio seems to have a constant rise towards higher metallicities when using the forbidden oxygen line. However, in the case when the O\u202fI 6158 \u212b line is employed, we do observe that the maximum in [C\/O] takes places close to solar metallicity to then become flat or decrease. This suggests that low-mass stars mostly contribute to carbon around solar metallicity, whereas at super-solar metallicities, massive stars produce carbon together with oxygen, thereby flattening or even decreasing the [C\/O] ratio. This trend is in agreement with the metallicity dependent yields from Carigi et al. (2005), which provide higher carbon as [Fe\/H] increases from massive stars (i.e. also increasing the O production) but lower carbon from low and intermediate mass stars as [Fe\/H] increases (i.e. less production of C). The turning point of increased relative production of carbon from massive stars takes place at A(O) ~ 8.7 dex (see Fig. 2 of Carigi et al. 2005) which equals to [O\/H] ~ 0.0 dex. This observed behaviour of [C\/O] is in contrast to the steady increase of [C\/O] up to [Fe\/H] ~ 0.3 dex found, for example, by Franchini et al. (2021). Nevertheless, the general trend we find when using the [O\u202fI ] 6300 \u212b line is similar to the reported by Franchini et al. (2021), who use also that oxygen indicator. All thick-disk stars present negative [C\/O] ratios and when using the oxygen line at 6158 \u212b thin-disk stars with [Fe\/H] \u2272 \u20130.2 have [C\/O]  0 as well. Thick-disk stars and low-metallicity thin-disk stars at the same metallicity have similar [C\/O] ratios, meaning that the balance between different production sites for oxygen and carbon is the same among both populations, despite [C\/Fe] and [O\/Fe] being systematically higher for thick-disk stars at a given metallicity.","Citation Text":["Carigi et al. 2005"],"Citation Start End":[[1301,1319]]}
{"Identifier":"2015AandA...576A...5C__Watanabe_et_al._2004_Instance_1","Paragraph":"In contrast, experimental studies based on irradiation of ices show that the second scenario is likely. Such studies show that glycolaldehyde, ethylene glycol, and methyl formate can be synthesized by irradiation of pure or mixed methanol (CH3OH) ices (Hudson & Moore 2000; \u00d6berg et al. 2009). Interestingly, the (CH2OH)2\/CH2OHCHO abundance ratio is found to be dependent on the initial ice composition as well as on the ice temperature during the UV irradiation. The CH3OH:CO ratio in the ices is a key parameter: for irradiated 20 K ices a composition of pure CH3OH leads to a (CH2OH)2\/CH2OHCHO ratio higher than 10, while a CH3OH:CO 1:10 ice mixture produces a (CH2OH)2\/CH2OHCHO ratio lower than 0.25 (\u00d6berg et al. 2009). The difference found between IRAS 16293 and IRAS2A could then be related to a different grain mantle composition in the two sources. If the CH3OH:CO ratio in the grain mantles of IRAS2A was higher than in IRAS 16293, a higher (CH2OH)2\/CH2OHCHO abundance ratio would be expected according to the laboratory results. In fact, the CH3OH gas-phase abundance in the inner envelope is found to be higher in IRAS2A (~4 \u00d7 10-7, J\u00f8rgensen et al. 2005b) than in IRAS 16293 (~1 \u00d7 10-7, Sch\u00f6ier et al. 2002), while the CO abundance is relatively similar (~(2\u20133)\u2009\u00d7\u200910-5, J\u00f8rgensen et al. 2002; Sch\u00f6ier et al. 2002). This could consequently be the result of the desorption of ices with a higher CH3OH:CO ratio in IRAS2A than IRAS 16293. The question then arises of how CH3OH can be more efficiently produced on grains in IRAS2A than in IRAS 16293. Several scenarios are possible: i) the initial conditions may play an important role in the CH3OH:CO ratio. In particular, experiments and simulations show that the efficiency of CH3OH formation through CO hydrogenation on the grains is dependent on temperature, ice composition (CO:H2O), and time (Watanabe et al. 2004; Fuchs et al. 2009); ii) the collapse timescale was longer in IRAS2A than in IRAS 16293, enabling the formation of more CH3OH; iii) the H2 density in the prestellar envelope of IRAS2A was lower than that of IRAS 16293. A less dense environment would lead to a higher atomic H density and consequently to a higher efficiency of CO hydrogenation. This was proposed by Maret et al. (2004) and Bottinelli et al. (2007) to explain the anticorrelation found between the inner abundances of H2CO and CH3OH and the ratios of submillimeter to bolometric luminosity (Lsmm\/Lbol) of different low-mass protostars. The Lsmm\/Lbol parameter is interpreted as an indication of different initial conditions, rather than an evolutionary parameter in this context (Maret et al. 2004). The Lsmm\/Lbol ratios of IRAS2A (~0.005, Karska et al. 2013) and IRAS 16293 (~0.019, Froebrich 2005) are consistent with this hypothesis. The current H2 density profiles of these two sources also agree with this scenario if they keep the memory of the prestellar conditions. The density derived in the outer envelope of IRAS2A with a power-law model (J\u00f8rgensen et al. 2002) is lower than the density derived in IRAS 16293 by Crimier et al. (2010), whether it be for a Shu-like model or a power-law model, while the temperature profiles are relatively similar (see Fig. 5). Along the same lines, Hudson et al. (2005) showed with proton irradiation experiments that glycolaldehyde is more sensitive to radiation damage than ethylene glycol. Irradiation would be more important in less dense envelopes, which would also be consistent with a less dense prestellar envelope in IRAS2A. A recent experiment by Fedoseev et al. (2015) shows that these two species can also be synthesized by surface hydrogenations of CO molecules in dense molecular cloud conditions. They do not directly form from CH3OH, but the results of this experiment show that similarly to CH3OH, which results from successive hydrogenations of CO, ethylene glycol forms by two successive hydrogenations of glycolaldehyde. This consequently agrees with the proposed scenario. ","Citation Text":["Watanabe et al. 2004"],"Citation Start End":[[1858,1878]]}
{"Identifier":"2016MNRAS.455.4426V__Mauche_&_Gorenstein_1986_Instance_1","Paragraph":"From the best-fitting values of the normalization at a fixed energy, we can obtain the ratios of the hydrogen column densities for different dust layers, such as:\n\n(16)\n\n\\begin{eqnarray}\n\\frac{N_{\\rm H,12}}{N_{\\rm H, 5}} \\approx \\frac{\\mathcal {F}_{\\rm X}^{(12)}(E)}{\\mathcal {F}^{(5)}_{\\rm X}(E)} \\frac{\\left(1-\\langle x \\rangle _{12} \\right){\\langle x \\rangle _{12}}}{\\left(1-x_5\\right){x_5}},\n\\end{eqnarray}\n\nwhere we assumed that the number of grains per hydrogen atom is the same in all clouds and dropped by the factor of Cd, i, since this is not sensitive to the grain composition (Mauche & Gorenstein 1986) (see also Section 2). In the above, NH, 12 represents the average column density of dust clouds 1 and 2. As already mentioned in Section 4.1 and later in Section 4.3, these are most probably one dust cloud of finite width, whose average column density we are probing. Having derived the ratios at four different energies, we calculated the weighted means, namely NH, 12\/NH, 5 = 0.19 \u00b1 0.03, NH, 12\/NH, 34 = 0.059 \u00b1 0.006, and NH, 34\/NH, 5 = 3.5 \u00b1 0.6. Thus, the intermediate dust clouds have the highest column density. This is 3.5 times higher than that of the fifth dust cloud, which may also explain the approximately equal intensities of rings 3\u20134 and 5, given their different distances from the observer. Interestingly, the column density of the clouds located closest to V404 Cyg is \u223c6\u2009per\u2009cent NH, 34. These findings are in rough qualitative agreement with the extinction histogram shown in Fig. 5, where the lowest extinction values are found at distances \u22732 kpc, i.e. where the dust layers 1 and 2 are located. Quantitatively, we find that the ratios NH, 34\/NH, 5 and NH, 34\/NH, 12 as derived from the extinction maps (Sale et al. 2014) are by a factor of \u223c3 smaller than our estimates, while the NH, 12\/NH, 5 ratio is the same (within uncertainty values). This implies that our analysis is overestimating the NH, 34 column density by a factor of \u223c3. A simple explanation for this difference is that the dust-to-hydrogen ratio is not the same for all clouds as originally assumed in our calculations (the Ai\/Aj term has dropped from the l.h.s. of equation 16). In any case, the ratio of the normalizations $\\mathcal {F}_{\\rm X}^{(i)}$ can be directly related to the ratio of dust column densities NH, iAi without resorting to any additional assumptions.","Citation Text":["Mauche & Gorenstein 1986"],"Citation Start End":[[589,613]]}
{"Identifier":"2017MNRAS.464L..26F__O'Sullivan_et_al._2001_Instance_1","Paragraph":"The diffuse hot gas X-ray luminosities in the 0.3\u20138 keV band are taken from the work of KF15. They have carefully removed the contribution from discrete sources such as low-mass X-ray binaries (Fabbiano 2006) to the total X-ray luminosity, leaving the diffuse gas contribution LX, Gas. A correction to bolometric would increase the X-ray luminosities by 0.08 dex on average. Most of the X-ray data come from Chandra observations. However, for some high-mass galaxies, the X-ray emission is particularly extended (e.g. NGC 4374, 4486, 4649, 5846) and in those cases ROSAT data from O\u2018Sullivan et al. (2001), corrected to the Chandra energy band, are used. Although the contribution from discrete sources in the ROSAT data cannot be subtracted as accurately as it can for Chandra data, their contribution is only about 1\u2009per\u2009cent of the diffuse gas luminosity for these high-mass galaxies (see O'Sullivan et al. 2001). For further details, see KF15. Here we make a very small correction to the KF15 LX,Gas luminosities for the distances used in the SLUGGS survey (Brodie et al. 2014). The KF15 compilation did not include several galaxies that appear in the Alabi et al. (in preparation) study. Here we also include the X-ray luminosities for NGC 720, NGC 1316 and NGC 3115 from Boroson et al. (2011), and for NGC 5128 from KF13. Su et al. (2014) conducted a detailed XMM and Chandra study of NGC 1400. As well as the X-ray emission centred on NGC 1400, they detected an enhanced region of X-rays to the NE of the galaxy that they associated with stripped gas. Here we use the X-ray luminosity centred on NGC 1400 with a small adjustment to our X-ray band and distance, and assume an uncertainty of 20\u2009per\u2009cent. We note that the X-ray luminosity would double if the enhanced region were also included. Two galaxies in Alabi et al. (in preparation) but not included here are NGC 2974 (not observed by Chandra) and NGC 4474 (the Chandra observation was only 5 ks).","Citation Text":["O\u2018Sullivan et al. (2001)"],"Citation Start End":[[581,605]]}
{"Identifier":"2015AandA...584A.103S__Chamel_et_al._2011_Instance_1","Paragraph":"Douchin & Haensel (2001; DH) formulated a unified EoS for NS on the basis of the SLy4 Skyrme nuclear effective force (Chabanat et al. 1998), where some parameters of the Skyrme interaction were adjusted to reproduce the Wiringa et al. calculation of neutron matter (Wiringa et al. 1988) above saturation density. Hence, the DH EoS contains certain microscopic input. In the DH model the inner crust was treated in the CLDM approach. More recently, unified EoSs for NS have been derived by the Brussels-Montreal group (Chamel et al. 2011; Pearson et al. 2012; Fantina et al. 2013; Potekhin et al. 2013). They are based on the BSk family of Skyrme nuclear effective forces (Goriely et al. 2010). Each force is fitted to the known masses of nuclei and adjusted among other constraints to reproduce a different microscopic EoS of neutron matter with different stiffness at high density. The inner crust is treated in the extended Thomas-Fermi approach with trial nucleon density profiles including perturbatively shell corrections for protons via the Strutinsky integral method. Analytical fits of these neutron-star EoSs have been constructed in order to facilitate their inclusion in astrophysical simulations (Potekhin et al. 2013). Quantal Hartree calculations for the NS crust have been systematically performed by (Shen et al. 2011b,a). This approach uses a virial expansion at low density and a RMF effective interaction at intermediate and high densities, and the EoS of the whole NS has been tabulated for different RMF parameter sets. Also recently, a complete EoS for supernova matter has been developed within the statistical model (Hempel & Schaffner-Bielich 2010). We shall adopt here the EoS of the BSk21 model (Chamel et al. 2011; Pearson et al. 2012; Fantina et al. 2013; Potekhin et al. 2013; Goriely et al. 2010) as a representative example of contemporary EoS for the complete NS structure, and a comparison with the other EoSs of the BSk family (Chamel et al. 2011; Pearson et al. 2012; Fantina et al. 2013; Potekhin et al. 2013) and the RMF family (Shen et al. 2011b,a) shall be left for future study. ","Citation Text":["Chamel et al. 2011"],"Citation Start End":[[518,536]]}
{"Identifier":"2016MNRAS.462.3441D__Namouni_1999_Instance_3","Paragraph":"In principle, Fig. 5, central panel G, shows that (469219) 2016 HO3 may have been locked in a Kozai\u2013Lidov resonance with \u03c9 librating about 270\u00b0 for nearly 100 kyr and probably more. Because of the Kozai\u2013Lidov resonance, both e (central panel E) and i (central panel F) oscillate with the same frequency but out of phase (for a more detailed view, see Fig. 4, panels E and F); when the value of e reaches its maximum the value of i is the lowest and vice versa ($\\sqrt{1 - e^2} \\cos i \\sim$ constant, see Fig. 4, panel B). During the simulated time and for the nominal orbit, 469219 reaches perihelion and aphelion the farthest possible from the ecliptic. Fig. 5, G-panels, show that for other incarnations of the orbit of 469219, different from the nominal one, \u03c9 may librate about 90\u00b0 as well during the simulated time interval. However, is this a true Kozai\u2013Lidov resonance? Namouni (1999) has shown that the secular evolution of co-orbital objects is viewed more naturally in the er\u03c9r-plane, where er and \u03c9r are the relative eccentricity and argument of perihelion computed as defined in Namouni's work (see equations 3 in Namouni 1999); these are based on the vector eccentricity and the vector inclination. Fig. 6 shows the multi-planet er\u03c9r-portrait for the nominal orbit of this object. It clearly resembles figs 13 and 19 in Namouni (1999). Asteroid 469219 librates around $\\omega _{\\rm r}=-90{^\\circ }$ for Venus, the Earth, and Jupiter. This behaviour corresponds to domain III in Namouni (1999), horseshoe-retrograde satellite orbit transitions and librations (around $\\omega _{\\rm r}=-90{^\\circ }$ or 90\u00b0). For a given cycle, the lower part corresponds to the horseshoe phase and the upper part to the quasi-satellite or retrograde satellite phase. This is not the Kozai\u2013Lidov resonance; in this case, the Kozai\u2013Lidov domain (domain II in Namouni 1999) is characterized by libration around $\\omega _{\\rm r}=0{^\\circ }$ (or 180\u00b0) which is only briefly observed at the end of the backwards integrations (see Fig. 6). The Kozai\u2013Lidov resonance is however in action at some stage in the orbits displayed in Figs 5 and 8. Our calculations show that the orbital evolution followed by 469219 is the result of the dominant secular perturbation of Jupiter as the periodic switching between co-orbital states ceases after about 8 kyr if Jupiter is removed from the calculations. Fig. 7 shows that, without Jupiter, 469219 switches between the Kozai\u2013Lidov domain and that of horseshoe-quasi-satellite orbit transitions and librations (including both \u221290\u00b0and 90\u00b0). Jupiter plays a stabilizing role in the dynamics of objects following orbits similar to that of 469219. It is not surprising that Jupiter instead of the Earth or Venus is acting as main secular perturber of 469219. Ito & Tanikawa (1999) have shown that the inner planets share the effect of the secular perturbation from Jupiter; in fact, Venus and our planet exchange angular momentum (Ito & Tanikawa 2002). In their work, these authors argue that the inner planets maintain their stability by sharing and weakening the secular perturbation from Jupiter. Tanikawa & Ito (2007) have extended this analysis to conclude that, regarding the secular perturbation from Jupiter, the terrestrial planets form a collection of loosely connected mutually dynamically dependent massive objects. The existence of such planetary grouping has direct implications on the dynamical situation studied here; if Jupiter is removed from the calculations, the overlapping secular resonances and the recurrent dynamics disappear as well.","Citation Text":["Namouni (1999)"],"Citation Start End":[[1333,1347]]}
{"Identifier":"2018AandA...616A..34H__Mohamed_&_Podsiadlowski_(2012)_Instance_1","Paragraph":"The CO emission, tracking the bulk of the gas, reveals an almost face-on one-armed spiral, of which almost two full windings can be traced. What could be the origin of this spiral structure? As the majority of AGB stars are in binary systems, and perhaps all host planets, interaction between the outflow and a sufficiently massive and nearby companion may be the explanation of the observed CO morphology. The intricate emission features in the inner 2\u2033 of the central CO emission maps is strongly reminiscent of hydrodynamical simulations of wind\u2013binary interaction by Mastrodemos & Morris (1998) and Mohamed & Podsiadlowski (2012), where the latter authors performed tailored simulations for the Mira AB system in which the outflow of the AGB star Mira A is perturbed by the presence of its close companion Mira B. The wind\u2013binary interaction that ensues leads to what is known as wind Roche-lobe overflow (WRLOF), where the slow AGB wind is confined to the star\u2019s Roche lobe, while overflowing through the L1 Lagrange point. Gravitational interaction of the overflowing material with the companion produces an intricate feedback system where the stellar outflow material is ejected into the surrounding CSE through two distinct streams (through L2 and the stagnation point3) which combine to form an annular stream. As this stream travels outwards, it creates the larger scale spiral observed in the wind. The morphology resulting from this particular type of wind\u2013binary interaction is shown in Fig. 3 in Mohamed & Podsiadlowski (2012). In Fig. 11 we show the emission pattern seen in the central regions of the CO channel at \u03c5*. We compare this image with the bottom left panel of Fig. 3 in Mohamed & Podsiadlowski (2012), an opacity map of the interaction zone. Though the two properties that are compared differ in nature, they likely still trace the same global morphological structure. Indeed, several of the predicted morphological features can be identified in the data of EP Aqr. The bright central region with a north and southward hook-like extension are strikingly similar, as are the eastern and western crescent-shaped \u201cvoids\u201d, the overall shape, and the morphological properties of the small-scale instabilities.","Citation Text":["Mohamed & Podsiadlowski (2012)"],"Citation Start End":[[603,633]]}
{"Identifier":"2020AandA...642A.231L__Maltby_et_al._2001_Instance_1","Paragraph":"We wanted to discuss the observed periods at sunspot. A period of nearly five minutes can be detected in Doppler velocities of CO 3-2 R14 and 7-6 R67 lines, which is consistent with previous findings in white light images or continuum spectrum (Beckers & Schultz 1972; Lites 1988; Nagashima et al. 2007; Yuan et al. 2014; Su et al. 2016), and might be considered the solar p-mode waves in the photosphere (Thomas 1985; Bogdan 2000; Solanki 2003). While a period of roughly three minutes is found in the Doppler velocities of the CO 3-2 R14 and Mg\u202fII k lines, which agrees closely with the previous observational results in UV\u2013infrared lines or images at the sunspot umbra (e.g., Solanki et al. 1996; Bogdan 2000; Fludra 2001; Maltby et al. 2001; Centeno et al. 2008; Tian et al. 2014; Khomenko & Collados 2015; Yang et al. 2017). They are explained as the resonant modes of sunspot oscillations (Uexkuell et al. 1983; Thomas 1984; Gurman 1987; Khomenko & Collados 2015). On the other hand, the CO 3-2 R14 line is believed to provide the information in the upper photosphere or the temperature minimum region (Uitenbroek 2000; Ayres 2002; Ayres et al. 2006). Moreover, a time delay of about two minutes is measured between the CO 3-2 R14 and Mg\u202fII k lines, as shown in Fig. 9. So, the three-minute oscillation at the sunspot umbra could come from the upper photosphere or the temperature minimum region and then propagate to the chromosphere, supporting the interpretation that propagating waves above the sunspots originate from the lower solar atmosphere (De Moortel et al. 2002; O\u2019Shea et al. 2002; Brynildsen et al. 2004; Khomenko & Collados 2015; Krishna Prasad et al. 2015). In other words, the three-minute oscillation can be regarded as the upwardly propagating slow magnetoacoustic waves (e.g., Su et al. 2016; Chae et al. 2019; Cho & Chae 2020; Feng et al. 2020b). Finally, Chae et al. (2017) found that the three-minute oscillation in the light bridge or umbral dots of a sunspot originates from the photosphere (see aslo Cho et al. 2019), which is consistent with our results and further suggests that the CYRA data is reliable.","Citation Text":["Maltby et al. 2001"],"Citation Start End":[[726,744]]}
{"Identifier":"2017MNRAS.464.3597L__Lind_et_al._2011_Instance_1","Paragraph":"It is important to emphasize that any estimate of [Na\/Fe] should be actually taken as an upper limit. Note that we can compute the theoretical response of stellar spectra to [Na\/Fe] only for stars hotter than 3500 K (see Section 3.2). As discussed in Appendix A, a linear extrapolation of Na responses to cooler temperatures would likely lead to lower [Na\/Fe] estimates, by about 0.1 dex. Moreover, as already noticed by CvD12b, atomic NaI transitions in the atmospheres of late-type stars are prone to substantial departures from LTE (Bruls, Rutten & Shchukina 1992; Gehren et al. 2006; Andrievsky al. 2007; Lind et al. 2011). In a Sun-like star, LTE calculations predict weaker lines (Allende Prieto et al. 2003) requiring corrections for the strongest lines, which can be as high as an effective change in the abundance of \u223c0.5 dex. For a lower temperature star (with Teff \u223c 4000 K), more relevant for models having old ages (as those of our sample of ETGs), we may expect NLTE corrections in the range of 0.1\u20130.2 dex (based on fig. 4 of Lind et al. 2011). Hence, we may expect that NLTE models could enhance the predicted absorption in Na indices, resulting in lower inferred values of [Na\/Fe]. However, note that our LTE-based methodology can match NaD, $\\rm Na\\,\\small {I}8190$, $\\rm Na\\,\\small {I}1.14$, and $\\rm Na\\,\\small {I}2.21$, simultaneously, suggesting that NLTE corrections should be approximately the same for all four Na lines \u2013 an important aspect to test with future models. We emphasize that although we have introduced four free-fitting parameters to match the four Na lines (the $\\alpha _{{\\rm Na}_j}$ constants; see equation 1), in practice, the values of \u03b1NaD, $\\rm \\alpha _{{\\rm Na}\\,\\small {I}1.14}$ and $\\rm \\alpha _{{\\rm Na}\\,\\small {I}2.21}$ are fully consistent with those of \u03b1-MILES and CvD12a models. Effectively, we are able to fit the four Na-sensitive line strengths of the seven X-Shooter spectra \u2013 spanning a range of age, metallicity, and [\u03b1\/Fe] \u2013 based on only one \u2018extra\u2019 free-fitting parameter (i.e. the $\\rm \\alpha _{{\\rm Na}\\,\\small {I}8190}$).","Citation Text":["Lind et al. 2011"],"Citation Start End":[[609,625]]}
{"Identifier":"2015AandA...576A.130F__Salvato_et_al._2011_Instance_1","Paragraph":"We performed an analysis on simulated XMM maps of point sources, presented in Brunner et al. (2008) for similarly large XMM exposures in the Lockman Hole, detecting no extended emission in the simulated maps containing the detected by XMM point sources. The higher sensitivity of Chandra towards the detection of point sources allows us to make a statistical assessment of the effect of sub-threshold (for XMM) AGNs toward the detection of extended emission. Performance of XMM observations was accompanied by deepening the Chandra data within one year from each other, which makes Chandra maps suitable for XMM point source contamination analysis, limiting the effect of AGN long-term variability (Salvato et al. 2011; Paolillo et al. 2004). We have computed the variation of unresolved point source flux on the detection scales for XMM, using the Chandra image, masking out the sources detected in the XMM analysis. The constructed Chandra flux map has been further smoothed with a Gaussian of 16\u2032\u2032 width, approximating the effects of the XMM PSF. In the map, the uniform distribution of the faintest point sources results in nearly constant emission, which we subtract following the procedure for local cosmic background estimates for XMM, while bright sources and clustered sources make an enhancement. We find the contamination by point sources unresolved by XMM to the flux of identified extended sources is below the 5% level of the extended source\u2019s flux. The highest peaks in the contamination map are associated with stand-alone sources near the (XMM) detection threshold, which by chance happened not to coincide with any of the detected groups and would contribute 30% to the faintest group flux. The importance of these sources is even higher in shallow surveys (Mirkazemi et al. 2015), to a degree requiring matched detection thresholds between point-like and extended sources, effectively removing faint extended sources from consideration. The importance of point source removal in XMM data is mentioned also in other cluster publications (Hilton et al. 2010; Pierre et al. 2012). ","Citation Text":["Salvato et al. 2011"],"Citation Start End":[[699,718]]}
{"Identifier":"2019AandA...624L...5P__Wang_et_al._(2011)_Instance_1","Paragraph":"One remarkable feature present in the central region of Orion is an explosive event that occurred 550\u2005\u00b1\u200525 years ago (J. Bally, priv. comm.) and was revealed by the three runaway stars BN, n, and I (G\u00f3mez et al. 2005; Rodr\u00edguez et al. 2005, 2017), and by the CO and H2 fingers (e.g., Allen & Burton 1993; Zapata et al. 2009; Nissen et al. 2012; Youngblood et al. 2016; Bally et al. 2017). However, Luhman et al. (2017) showed that the object n is no longer a runaway member because its real proper motion is much lower than previously estimated; but conversely another object, named x, is moving away at high speed from the same explosion center. J. Bally (priv. comm.) confirms the fast proper motion of x and the absence of movement of n from ground-based H2 images 14 years apart. X is further out having passed our 20% beam coupling mark (see Fig. 1 of Paper I), and therefore does not appear in the figures presented in this Letter. Zapata et al. (2011) and Orozco-Aguilera et al. (2017) in their follow-up work with ALMA proposed that the hot core (HC) is externally heated despite its high temperature, and that the heating source could be the nearby explosion. Similarly, Blake et al. (1987), Wang et al. (2011), and Favre et al. (2011) advocated that the Compact Ridge is also externally heated, although the heating source should not be the same since we presented evidence in Paper I that the Compact Ridge has not yet been affected by the impact of the explosion. A possibility could be the outflow from source I hitting the Compact Ridge (Liu et al. 2002). In Paper I, we also presented evidence for an interaction between the explosive event and the main components of the Orion KL region including the HC, several infrared (IR) components (Rieke et al. 1973), and methyl formate (CH3OCHO; hereafter MF) peaks (Favre et al. 2011). We showed that the IRc6\/MF5 and IRc20\/MF4 sources, west of the explosion center, display emission lines of various species having only red wings, while sources on the east and south sides display emission lines having only blue wings. We also confirmed that excited emission lines are found preferentially surrounding the explosion center and that complex organic molecules (COMs) rich in oxygen (O-COMs) do not occupy the same volumes as CN rich COMs (CN-COMs). We identified the ethylene glycol peak (EGP) to be coincident with a hollow sphere of material, which we interpreted to have originated from the impact of a \u201cbullet\u201d launched from the explosion center (Favre et al. 2017; Wright & Plambeck 2017). We also proposed that the Compact Ridge (MF1) is sufficiently far away from the rest of the KL region to have not yet been perturbed by the explosion, the evidence being the absence of asymmetric emission line wings and the narrowness of the lines themselves (\u223c1 km s\u22121). In this Letter, we study further the interaction of the explosion blowout with the surrounding gas and dense sources.","Citation Text":["Wang et al. (2011)"],"Citation Start End":[[1201,1219]]}
{"Identifier":"2016AandA...593A..22R__Shibuya_et_al._2015_Instance_3","Paragraph":"Although it is a simple concept, obtaining galaxy sizes is not an easy task and is subject to a number of assumptions. The most common way to derive galaxy sizes is by performing light-profile fitting assuming a given shape of the surface brightness profile using a \u03c72 minimization (e.g. Simard et al. 1999; Peng et al. 2002; Ravindranath et al. 2004; Daddi et al. 2005; Ravindranath et al. 2006; Trujillo et al. 2006; Akiyama et al. 2008; Franx et al. 2008; Tasca et al. 2009; Cassata et al. 2010, 2013; Williams et al. 2010; Mosleh et al. 2011; Huang et al. 2013; Ono et al. 2013; Stott et al. 2013; Morishita et al. 2014; van der Wel et al. 2014; Straatman et al. 2015; Shibuya et al. 2015). Another method assumes circular or elliptical apertures around a predefined galactic center and computes the size enclosing a given percentage of the total galaxy flux (e.g. Ferguson et al. 2004; Bouwens et al. 2004; Hathi et al. 2008; Oesch et al. 2010; Ichikawa et al. 2012; Curtis-Lake et al. 2016). A third approach, involving counting the number of pixels belonging to the galaxy to derive its size, was also explored in Law et al. (2007). Studies of galaxy sizes at z> 2 became possible with the deep imaging obtained with HST. The first reports on size evolution found that galaxy sizes as observed in the UV rest-frame were becoming smaller at the highest redshifts (Bouwens et al. 2003, 2004; Ferguson et al. 2004). We have now access to the size evolution up to z ~ 10 from the deepest HST imaging data (e.g., Hathi et al. 2008; Jiang et al. 2013; Ono et al. 2013; Kawamata et al. 2015; Holwerda et al. 2015; Shibuya et al. 2015). With the multiwavelength and near-infrared coverage of CANDELS (Grogin et al. 2011; Koekemoer et al. 2011) optical rest-frame measurements are reported up to z ~ 3 for a large collection of galaxies in diverse populations (e.g. Bruce et al. 2012; van der Wel et al. 2014; Morishita et al. 2014). At z ~ 2 the size of star-forming galaxies (SFGs) is, to first order, independent of the observed rest-frame bands (Shibuya et al. 2015). It is generally accepted that galaxy sizes tend to decrease with increasing redshift (e.g. Bouwens et al. 2003, 2004; Ferguson et al. 2004; Mosleh et al. 2012) and that galaxy sizes depend on stellar mass (e.g. Franx et al. 2008; van der Wel et al. 2014; Morishita et al. 2014) and luminosity (e.g. Grazian et al. 2012; Huang et al. 2013). However, some results point to a scenario consistent with no size evolution as seen in UV rest-frame from HST data (Law et al. 2007; Curtis-Lake et al. 2016) and, at a fixed stellar mass, from optical rest-frame ground-based data (Ichikawa et al. 2012; Stott et al. 2013). ","Citation Text":["Shibuya et al. 2015"],"Citation Start End":[[2048,2067]]}
{"Identifier":"2022AandA...662A..99K__Avenhaus_et_al._2018_Instance_1","Paragraph":"In recent years, the quest to observe embedded accreting planets as they form has come a long way as a result of combined observations that are covering the ultraviolet to the millimeter wavelength ranges. One instrument in particular has been playing a major role in this, due to its diffraction-limited and high-contrast observations: the Spectro-Polarimetric High contrast imager for Exoplanets REsearch (SPHERE; Beuzit et al. 2019), which covers the visible and near-infrared (VIS\/NIR) wavelength ranges. Together with other instruments, it has revealed a large abundance of substructures that are now known to be commonly present in protoplanetary disks (PPDs) ranging from gaps and rings to spiral arms, cavities, and various asymmetric features (e.g., Garufi et al. 2016, 2018; Avenhaus et al. 2018; Andrews et al. 2018, DSHARP). These features are often interpreted as the result of embedded planets that interact with their natal circumstellar disk (CSD; e.g., Wolf & D\u2019Angelo 2005; Fouchet et al. 2010; Ruge et al. 2013; Ober et al. 2015; Dong et al. 2016). However, other origins have been proposed as well (e.g., Flock et al. 2015; Zhang et al. 2015; Ruge et al. 2016; Gonzalez et al. 2017; Suriano et al. 2017, 2019; Dullemond & Penzlin 2018). Despite the ambiguity of their particular origins, the properties of the potential planets that may produce these features have been studied extensively, especially using hydrodynamics simulations (e.g., Bae et al. 2019; Toci et al. 2020; Calcino et al. 2020). Dong & Fung (2017), for instance, analyzed observed gap features in order to determine the masses of planets that may have caused them by using 2D and 3D hydrodynamics simulations with 3D radiative transfer simulations of five scattered light observations of CSDs around Herbig Ae\/Be and T Tauri stars: HD 97048 (Ginski et al. 2016), TW Hya (van Boekel et al. 2017), HD 169142 (Momose et al. 2015), LkCa15 (Thalmann et al. 2016), and RX J1615 (de Boer et al. 2016). By assuming an \u03b1-viscosity model (Shakura & Sunyaev 1973) with \u03b1visc = 10\u22123 and single gap-opening planets as origins, they deduced that the corresponding planetary masses are all of typical low-mass giant planets between about 0.1 and 1 MJ.","Citation Text":["Avenhaus et al. 2018"],"Citation Start End":[[785,805]]}
{"Identifier":"2015ApJ...801..112L___1990_Instance_1","Paragraph":"Both Greco et\u00c2 al. (2009) and Servidio et\u00c2 al. (2011) show that a remarkable similarity exists between the PDF of the field magnitude from 2D MHD simulations with a guide field and the PDF from ACE observations in the supersonic solar wind near Earth. This suggests that the common occurrence of strong magnetic field discontinuities in the solar wind (e.g., Li 2008) might be attributed to the ubiquitous presence of inertial-scale quasi-2D coherent magnetic field structures or flux ropes interspersed with current sheets and reconnection x-points. Additional arguments in support of this view come from theoretical considerations (e.g., Zank & Matthaeus 1992, 1993) and observations (Matthaeus et\u00c2 al. 1990; Bieber et\u00c2 al. 1996) that MHD turbulence in the supersonic solar wind can be thought of comprising to lowest order a dominant 2D component and a minor parallel propagating Alfv\u00c3\u00a9n wave component. Furthermore, recent observations in the supersonic solar wind suggest the common occurrence of magnetic islands with different sizes in the inertial range together with evidence for island merging near the heliospheric current sheet (Khabarova et\u00c2 al. 2014). A statistical solar wind data analysis by Osman et\u00c2 al. (2014) also supports the idea that strong magnetic field discontinuities have a larger likelihood of being associated with current sheets and x-point reconnection outflows and have a higher possibility that a connection exists between current sheets and such outflows. With today's higher resolution measurement capability, identification of x-point reconnection outflows in the solar wind has become quite common (e.g., Gosling 2010). Gosling (2010) found that most reconnection events are observed in the slow supersonic solar wind (\u00e2\u0088\u00bc2\u00e2\u0080\u00933 events per day near Earth), while in the fast supersonic solar wind much less are found (\u00e2\u0088\u00bc0.6 events per day). Reconnection events are also regularly observed in association with coronal mass ejections (Gosling 2010) where inertial-scale flux-rope merging was reported at the current sheet behind a coronal mass ejection close to the Sun by Song et\u00c2 al. (2012), and current sheet reconnection was observed at the turbulent leading edge of an interplanetary coronal mass ejection near Earth by Chain & Munoz (2011), for example.","Citation Text":["Matthaeus et\u00c2 al. 1990"],"Citation Start End":[[687,709]]}
{"Identifier":"2020MNRAS.498..689C__Matrajt_et_al._2013_Instance_1","Paragraph":"Fig. 1 clearly shows that in increasing fluence the processes of dissociation are greater, and thus new molecular species come to light. We identify the new species by comparing the peak positions with similar works on acetone ices. New observed peaks were: a flat peak in around 3 \u03bcm that could be associated to amorphous water, but since several other alcohols have a similar peak and water is not a highly expected acetone radiation product (see Section 3), we could not make any strong assignment to this peak. Some new peaks in the range of 2000\u20133500 cm\u22121, which we assign to isopropanol, ketene, and carbon monoxide according to the work of Hudson (2018). Other two peaks in the region 2000\u2013600 cm\u22121 we assign to CH4 and -CH=CH- bending (Hudson 2018; Matrajt et al. 2013). The most sharp and visible peak is the CO2 one. CO2 could be a contaminant, however the CO2 peak is not present in our control experiments, such as monitoring the ice sample before irradiation to confirm that without irradiation sample remains unaltered, or also irradiation of the bare substrate. Also, the CO2 peak increases with fluence showing this is a molecule being formed due to irradiation. Due the vacuum chamber pressure to be 10\u22129 mbar, some residual gas could contaminate the sample, but this is accounted for our error bars estimates. The H2CCO peak can be overlapped with a CO one since both are expected radiation products and have a peak near 2128 cm\u22121. However, as stated in Hudson (2018), the shift of this peak to higher wavenumbers with increasing fluence indicates this peak is most due to formation of ketene. In comparison with the work of Hudson (2018), he observed peaks for CH4, CO, H2CCO, and C3H8O. Here, we observed and confirm a peak for CO2 as well other similar peaks presented in the work of Hudson (2018) for isopropanol and ketene. Photo-irradiation of pure acetone ice with X-rays was also performed in Almeida et al. (2014), where authors employed a photo-stimulated desorption (PSID) and took the positive PSID spectra to analyse acetone fragmentation. They have thus detected lots of cations as fragments, such as HCO+, CO+, H+, C2CH$_3\\, ^+$, CO$_2\\, ^+$, H2CO+, etc. Tables 1 and 2 present a compilation of peak positions (in wavenumbers and wavelength), our corresponding assignments, band strengths and references used to determine molecular assignments. Question mark in Table 2 is a peak probably related to several different molecules being formed in the ice as explained before. Note that ketene band strength was taken from Berg & Ewing (1991) and it was also used as reference value in other recent works (Hudson & Loeffler 2013; Bergner, \u00d6berg & Rajappan 2019). No reference band strength for C3H8O were found.","Citation Text":["Matrajt et al. 2013"],"Citation Start End":[[757,776]]}
{"Identifier":"2018AandA...614A..23K__Tr\u00fcmper_et_al._2013_Instance_1","Paragraph":"The beaming functions of the polar and fan beams are given by f sinm\u03d5 and p cosk\u03d5, respectively. Here \u03d5 is the angle between the magnetic field axis and the photon propagation and is thus a function of the angle between the magnetic and rotation axis (\u03b8), the angle of the NS rotation axis, the observing angle (i), and the NS spin phase. The exponent values m and k can be as low as 1, but in the general case of the accretion column emission, they can have significantly higher values (Tr\u00fcmper et al. 2013). We constructed various pulse profiles for different combinations of fan- and polar-beam emission patterns with different relative intensities of the fan and polar beams (i.e., FPolar\u2215FFan ranging from ~0.01 to ~30). We note that these values exceed any realistic configuration between the fan beam and the reflected polar beam. More specifically, since the polar beam is a result of reflection of the fan beam, it is only for a very limited range of observing angles that its contribution will (seemingly) exceed that of the fan beam, and thus a value of FPolar\u2215FFan \u2273 2, is unlikely. On the other hand, the increasing height of the emitting region, which would result in a smaller fraction of the fan beam being reflected by the NS surface, is limited to \u2272 10 km (Mushtukov et al. 2015), thus limiting the FPolar\u2215FFan ratio to values greater than 0.1 (see, e.g., Fig. 2 of Poutanen et al. 2013). Despite these physical limitations, we have extended our estimates to these exaggerated values in order to better illustrate the contribution of each component (fan and polar) to the PF. Moreover, a value of FPolar\u2215FFan exceeding 10 qualitatively describes the pencil-beam regime, which is expected at lower accretion rates. We estimated the PF for a range of observer angles and for different combinations of angles between the magnetic and rotation axis. An indicative result for a 45\u00b0 angle between the NS magnetic field and rotation axis and for an rG\u2215rNS = 0.25 is plotted inFig. 12.","Citation Text":["Tr\u00fcmper et al. 2013"],"Citation Start End":[[488,507]]}
{"Identifier":"2020AandA...637A..59A__Massalkhi_et_al._2019_Instance_4","Paragraph":"Silicon monoxide (SiO) is predicted to be the most abundant Si-bearing molecule in the entire 1\u201310 R* range in the atmospheres of M stars. In S-type atmospheres, the calculated abundance of SiO decreases by two orders of magnitude in the 1\u20135 R* but retains a very high abundance beyond, and the same occurs in C-rich atmospheres, although in this case, the abundance drop in the 1\u20135 R* is even more pronounced (see Fig. 2; see also Ag\u00fandez & Cernicharo 2006). Observations indicate that the abundance of SiO does not differ significantly between envelopes around M-, S-, and C-type stars, although in all them the SiO abundance decreases with increasing mass-loss rate (Gonz\u00e1lez Delgado et al. 2003; Sch\u00f6ier et al. 2006; Ramstedt et al. 2009; Massalkhi et al. 2019, 2020). This decline in the SiO abundance with increasing envelope density is not a consequence of chemical equilibrium (Massalkhi et al. 2019), but has been interpreted as evidence that SiO disappears from the gas phase at high densities to be incorporated into dust grains (Gonz\u00e1lez Delgado et al. 2003; Sch\u00f6ier et al. 2006; Ramstedt et al. 2009; Massalkhi et al. 2019, 2020). It therefore appears that the gradual abundance decline calculated for SiO in the 1\u20135 R* region from stellar type in the sense M \u2192 S \u2192 C does not have a direct consequence in the SiO abundance that is injected into the expanding wind. However, this behavior predicted by chemical equilibrium probably explains why SiO masers are observed in M-type stars but not toward carbon stars (e.g., Pardo et al. 2004). Except for these details, chemical equilibrium and observations agree in the fact that SiO is one of the most abundant carriers of silicon in the atmospheres of M-, S-, and C-type stars. Calculations and observations also agree for SiS in that it is an abundant molecule regardless of the C\/O. However, observations indicate a differentiation between C- and O-rich envelopes, with SiS being on average one order of magnitude more abundant in carbon-rich sources (Sch\u00f6ier et al. 2007; Danilovich et al. 2018; Massalkhi et al. 2019, 2020). Moreover, in some oxygen-rich envelopes, the fractional abundance of SiS relative to H2 is as low as ~10\u22128, which is well below the predictions of chemical equilibrium (Danilovich et al. 2019; Massalkhi et al. 2020).","Citation Text":["Massalkhi et al. 2019"],"Citation Start End":[[2061,2082]]}
{"Identifier":"2015ApJ...806....1M__Ahn_et_al._2012_Instance_2","Paragraph":"For the clustering measurements, we use the sample of galaxies compiled in Data Release 11 (DR11) of the SDSS-III project. The SDSS-III is a spectroscopic investigation of galaxies and quasars selected from the imaging data obtained by the SDSS (York et al. 2000) I\/II covering about 11,000 deg2 (Abazajian et al. 2009) using the dedicated 2.5 m SDSS Telescope (Gunn et al. 2006). The imaging employed a drift-scan mosaic CCD camera (Gunn et al. 1998) with five photometric bands (\n\n\n\n\n\n and z; Fukugita et al. 1996; Smith et al. 2002; Doi et al. 2010). The SDSS-III (Eisenstein et al. 2011) BOSS project (Ahn et al. 2012; Dawson et al. 2013) obtained additional imaging data of about 3000 deg2 (Aihara et al. 2011). The imaging data was processed by a series of pipelines (Lupton et al. 2001; Pier et al. 2003; Padmanabhan et al. 2008) and corrected for Galactic extinction (Schlegel et al. 1998) to obtain a reliable photometric catalog. This catalog was used as an input to select targets for spectroscopy (Dawson et al. 2013) for conducting the BOSS survey (Ahn et al. 2012) with the SDSS spectrographs (Smee et al. 2013). Targets are assigned to tiles of diameter 3\u00b0 using an adaptive tiling algorithm designed to maximize the number of targets that can be successfully observed (Blanton et al. 2003). The resulting data were processed by an automated pipeline which performs spectral classification, redshift determination, and various parameter measurements, e.g., the stellar-mass measurements from a number of different stellar population synthesis codes which utilize the photometry and redshifts of the individual galaxies (Bolton et al. 2012). In addition to the galaxies targeted by the BOSS project, we also use galaxies that pass the target selection but have already been observed as part of the SDSS-I\/II project (legacy galaxies). These legacy galaxies are subsampled in each sector so that they obey the same completeness as that of the CMASS sample (Anderson et al. 2014).","Citation Text":["Ahn et al. 2012"],"Citation Start End":[[1062,1077]]}
{"Identifier":"2019MNRAS.490.5722W__Remus_et_al._2013_Instance_1","Paragraph":"Since it is infeasible to observe the evolution of individual galaxies over time, theoretical approaches focusing on understanding the formation of ETGs have made use of numerical simulations to trace the evolution of individual galaxies. Through zoom-in and cosmological simulations, a consensus has emerged between these simulations that the formation of ETGs proceeds through two phases, where galaxies first go through dissipative gas-rich wet mergers followed by in situ star formation bursts at redshifts above z \u2248 2, and then evolve towards low redshift through non-dissipative gas-poor dry mergers (Naab et al. 2007; Guo & White 2008; Hopkins et al. 2009; Nipoti et al. 2009b; Nipoti, Treu & Bolton 2009a; Oser et al. 2010; Johansson, Naab & Ostriker 2012; Moster, Naab & White 2013; Remus et al. 2013; Furlong et al. 2015; Wellons et al. 2015, 2016; Rodriguez-Gomez et al. 2016). However, regarding the redshift evolution of ETGs\u2019 total power-law density slopes, no consensus has been reached neither among different cosmological hydrodynamic simulations nor between simulations and observations, despite the many advances in cosmological simulations (Vogelsberger et al. 2019a). While the Magneticum pathfinder simulation (Remus et al. 2017) and the Illustris simulations (Xu et al. 2017) produce shallower total density profile with time, the Horizon-AGN simulations (Peirani et al. 2019) produce steeper total density profile with time, in better agreement with the redshift evolution trend found in observations. However, the latter simulation has smaller slope values compared to the former two, which are closer to the observed slope values due to different implementation of feedback models, etc. Apart from cosmological simulations, dedicated zoom-in simulations (Johansson, Naab & Burkert 2009; Johansson et al. 2012; Remus et al. 2013) have revealed that dry mergers that dominate the passive evolution of ETGs below z \u2248 2 could make the total density profile shallower than isothermal (Hilz et al. 2012; Hilz, Naab & Ostriker 2013; Remus et al. 2017). The inclusion of wet mergers is also crucial for reconciling the simulated redshift evolution trend of the slope with strong-lensing observations (Sonnenfeld, Nipoti & Treu 2014).","Citation Text":["Remus et al. 2013"],"Citation Start End":[[792,809]]}
{"Identifier":"2020ApJ...894...61P__Yuan_et_al._2009_Instance_1","Paragraph":"On the basis of its unusual morphology and steep spectrum, we suggest that R1 traces the synchrotron radiation from a wind launched from the accretion flow around the SMBH. Known as the ADAF wind, such an energetic outflow is a generic prediction by theories and numerical simulations of hot accretion flows in LLAGNs (Blandford & Begelman 1999; Yuan et al. 2012a, 2012b, 2015). Originating from the corona region of the hot accretion flow, the ADAF wind would have a wide opening angle compared to the highly collimated jet, thus more naturally matching the observed radio morphology. Relativistic electrons responsible for the synchrotron may be produced in the corona of ADAF, where frequent magnetic reconnection is expected to occur (Yuan et al. 2009). Alternatively, relativistic electrons may be produced via the outward propagating wind shocking the circumnuclear medium over a region of 10 pc, corresponding to \u223c4 \u00d7 107Rs (here \n\n\n\n\n\n is the Schwarzschild radius of the putative SMBH). To our knowledge, there has been no dedicated theoretical effort to predict the observational signatures of an ADAF wind on relevant physical scales. Here we shall provide a simple estimate of the energetics related to the radio emission. First, we notice that the radiative efficiency of an ADAF (\n\n\n\n\n\n) at highly sub-Eddington rates scales roughly with the net mass accretion rate (i.e., the mass inflow rate into the event horizon) as \n\n\n\n\n\n (Xie & Yuan 2012), where Lbol is the bolometric luminosity of the SMBH and \n\n\n\n\n\n is the Eddington accretion rate. Assuming that the 2\u201310 keV luminosity of X1 (\n\n\n\n\n\n) is 10% of the bolometric luminosity, and adopting a canonical value of \u03b4 = 0.5, the fractional viscous heating of electrons in the ADAF,11\n\n11\nSee Xie & Yuan (2012) for discussions about this parameter; adopting a smaller value of \u03b4 would result in a higher \n\n\n\n\n\n.\n we may infer \n\n\n\n\n\n. Numerical simulations show that the ADAF wind carries a kinetic energy of \n\n\n\n\n\n (Yuan et al. 2015), which is then \n\n\n\n\n\n according to the above estimate. This value is sufficient to account for the observed synchrotron radiation, given that typically 1% of the kinetic energy goes to shock-accelerating relativistic electrons (Nims et al. 2015). The ADAF wind may also have produced soft X-ray emission upon shock-heating the circumnuclear medium. We note that the observed 0.5\u20132 keV luminosity from X1 (\n\n\n\n\n\n; Y15) is compatible with the estimated wind kinetic energy.","Citation Text":["Yuan et al. 2009"],"Citation Start End":[[739,755]]}
{"Identifier":"2016MNRAS.455..739K__Gendreau,_Arzoumanian_&_Okajima_2012_Instance_1","Paragraph":"Even if non-accreting millisecond pulsars are unaffected by \u03b1-oscillations, some NSs in LMXBs (accreting millisecond pulsars), which are attached to sharper peaks centred at higher temperatures, still can be unstable with respect to \u03b1-oscillations and thus should exhibit them. To check this possibility, we need more X-ray observations with high temporal resolution that can be achieved by future missions such as LOFT (Feroci et al. 2012), NICER (Gendreau, Arzoumanian & Okajima 2012) and SRG (Merloni et al. 2012). Another opportunity to detect \u03b1-oscillations can be associated with the recently discovered transitional millisecond pulsars, which are switching between rotation- and accretion-powered states [currently, three of such objects are known: PSR J1023+0038 (\u03bd \u2248 592.4 Hz, Archibald et al. 2009; Stappers et al. 2014), IGR J18245\u22122452 (\u03bd \u2248 254.3 Hz, Papitto et al. 2013), and XSS J12270\u22124859 (\u03bd \u2248 593.0 Hz, Roy et al. 2015)]. In a rotation-powered state these NSs are observed as radio pulsars, giving a chance to study their timing properties precisely. The accretion-powered state guarantees that these pulsars cannot be too cold, and thus likely to be affected by r-mode instability and \u03b1-oscillations. For example, the redshifted surface temperature of PSR J1023+0038 can be as large as 5 \u00d7 105 K (Homer et al. 2006; Bogdanov et al. 2011). It corresponds to the internal temperature T\u221e = (1.5\u20133) \u00d7 107 K, implying that PSR J1023+0038 can be on the stability peak. Its timing behaviour was reported as \u2018complex\u2019 by Archibald et al. (2013) and can be influenced by \u03b1-oscillations. The reliable registration of \u03b1-oscillations would confirm the crucial role of the resonance interaction of oscillation modes in the evolution of NSs in LMXBs (Gusakov et al. 2014a,b). The measurements of \u03b1-oscillation parameters (such as the period of oscillations, $\\widehat{P}$) would provide us with a new powerful tool to constrain the properties of superdense matter.","Citation Text":["Gendreau, Arzoumanian & Okajima 2012"],"Citation Start End":[[449,485]]}
{"Identifier":"2022AandA...665A.118F__Rempel_2014_Instance_1","Paragraph":"Despite significant advances, the solar wind remains challenging to model (and importantly forecast) due to the large range of scales that need to be incorporated. It is well known that energy is injected into the corona at all scales, from magnetohydrodynamic (MHD) waves (Nutto et al. 2012; Van Doorsselaere et al. 2020) to flares and coronal mass ejections with global extent (Aschwanden et al. 2017; Green et al. 2018; Wyper et al. 2018). In addition to the range of spatial scales, heating events take place on a variety of timescales (Hollweg 1973; Viall & Klimchuk 2017). Putting aside large-scale eruptions so as to focus on the quasi-steady input of energy into the corona, it is possible to distinguish three broad magnetic configurations of the solar atmosphere for modelling purposes. These are the quiet Sun (Danilovic et al. 2010; Rempel 2014), active or enhanced regions (Carlsson et al. 2016; Chen et al. 2021), and coronal holes (W\u00f3jcik et al. 2019). In each of these configurations, energy is channelled from the convection into the low corona. Coronal holes are the dominant source of the solar wind in the heliosphere (Cranmer et al. 2017; Stansby et al. 2021), typically producing the fast solar wind (McComas et al. 2008; Ebert et al. 2009; Macneil et al. 2020a; Wang 2020). The magnetic field configuration of a coronal hole is relatively simple, compared with the quiet Sun and active regions, given that the field is principally open to the solar wind (Lowder et al. 2017; Hofmeister et al. 2019). However, there are still a range of dynamic processes taking place, such as the braiding of magnetic field lines (Wedemeyer-B\u00f6hm et al. 2012; Wedemeyer et al. 2013; Huang et al. 2018) and the emergence of new magnetic flux (Murray et al. 2009). These can trigger the formation of jets (Shen et al. 2017; Yang et al. 2017) and other phenomena, which are then observed as spicules (Mart\u00ednez-Sykora et al. 2017; Bose et al. 2021) or fibrils (Hansteen et al. 2006; Leenaarts et al. 2015).","Citation Text":["Rempel 2014"],"Citation Start End":[[845,856]]}
{"Identifier":"2015ApJ...804..130C__Bertschinger_1985_Instance_2","Paragraph":"We have rigorously developed the embedded gravitational lensing theory for point mass lenses in a series of recent papers (Chen et al. 2010, 2011, 2015; Kantowski et al. 2010, 2012, 2013) including the embedded lens equation, time delays, lensing magnifications, shears, etc. We successfully extended the lowest-order embedded point mass lens theory to arbitrary spherically symmetric distributed lenses in Kantowski et al. (2013). The gravitational correctness of the theory follows from its origin in Einstein\u2019s gravity. The embedded lens theory is based on the Swiss cheese cosmologies (Einstein & Straus 1945; Sch\u00fccking 1954; Kantowski 1969). The idea of embedding (or Swiss cheese) is to remove a co-moving sphere of homogeneous dust from the background Friedmann\u2013Lema\u00eetre\u2013Robertson\u2013Walker (FLRW) cosmology and replace it with the gravity field of a spherical inhomogeneity, maintaining the Einstein equations. In a Swiss cheese cosmology the total mass of the inhomogeneity (up to a small curvature factor) is the same as that of the removed homogeneous dust sphere. For a galaxy cluster, embedding requires the overdense cluster to be surrounded by large underdense regions often modeled as vacuum. For a cosmic void, embedding requires the underdense interior to be \u201ccompensated\u201d by an overdense bounding ridge, i.e., a compensated void (Sato & Maeda 1983; Bertschinger 1985; Thompson & Vishniac 1987; Mart\u00ednez-Gonz\u00e1lez et al. 1990; Amendola et al. 1999; Lavaux & Wandelt 2012). A low-density region without a compensating overdense boundary, or with an overdense boundary not containing enough mass to compensate the interior mass deficit, has a negative net mass (with respect to the homogeneous background) and is known as an \u201cuncompensated\u201d or \u201cundercompensated\u201d void (Fillmore & Goldreich 1984; Bertschinger 1985; Sheth & van de Weygaert 2004; Das & Spergel 2009).3\n\n3\nThis dichotomy between compensated and uncompensated voids is slightly different from one based on the classification of the small initial perturbations from which voids are thought to be formed. The initial perturbation can be compensated or uncompensated, which leads to different void growth scenarios (Bertschinger 1985), but if the evolved void formed from either perturbation is surrounded by an overdense shell that \u201clargely\u201d compensates the underdense region (i.e., the majority of the void mass is swept into the boundary shell in the snowplowing fashion when the void is growing), we still call it compensated because the small mass deficit originating in the initial perturbation is unimportant for gravitational lensing.\n Similarly, an overcompensated void has positive net mass with respect to the homogeneous FLRW background. Numerical or theoretical models of over-or undercompensated voids do commonly exist (e.g., Sheth & van de Weygaert 2004; Cai et al. 2010, 2014; Ceccarelli et al. 2013; Hamaus et al. 2014). We focus on compensated void models in this paper, given that uncompensated void models do not satisfy Einstein\u2019s equations. The critical difference between an embedded lens and a traditional lens lies in the fact that embedding effectively reduces the gravitational potential\u2019s range, i.e., partially shields the lensing potential because the lens mass is made a contributor to the mean mass density of the universe and not simply superimposed upon it. At lowest order, this implies that the repulsive bending caused by the removed homogeneous dust sphere must be accounted for when computing the bending angle caused by the lens mass inhomogeneity and legitimizes the prior practice of treating negative density perturbations as repulsive and positive perturbations as attractive. In this paper we investigate the gravitational lensing of cosmic voids using the lowest-order embedded lens theory (Kantowski et al. 2013). We introduce the embedded lens theory in Section 2, build the simplest possible lens model for a void in Section 3, and study the lensing of the CMB by individual cosmic voids in Section 4. Steps we outline can be followed for many void models of current interest.","Citation Text":["Bertschinger 1985"],"Citation Start End":[[1808,1825]]}
{"Identifier":"2022AandA...660A.108Z__Bouvier_et_al._1986_Instance_1","Paragraph":"The TESS light curves present the best cadence of the datasets available for CR Cha. They cover 56 days. The 2019 TESS light curve (Fig. 1, middle panel), taken with a 30 min cadence, reveals both periodic and stochastic variations with a peak-to-peak amplitude of \u223c0.04 mag. In order to find periodic signals in these photometric data, we computed the Lomb-Scargle periodogram (Lomb 1976; Scargle 1982), which is shown in the bottom panel of Fig. 1. The periodogram resulted in significant peaks at 2.314 \u00b1 0.033 days and 1.156 \u00b1 0.009 days. An analysis of the more recent TESS light curve taken with a 10 min cadence between 2021 April and June (Fig. 2), obtained with the same analysis technique as for the 2019 data (see Sect. 2.3), confirms the presence of these two peaks. The peak at 2.314 days is consistent with the stellar rotational period found in previous studies (P\u2004=\u20042.3 days; Bouvier et al. 1986), suggesting the presence of starspots on the stellar surface. The 1.156-day period is half the stellar rotational period. The TESS light curve also indicates a long-term oscillation with a timescale of \u223c25 days, but the time coverage of our dataset is not long enough to probe these longer timescales. Robinson et al. (2021) modeled the effect of inclination on how periodic a light curve appears. They found that higher inclinations lead to more burst-dominated light curves. The moderate 31\u00b0 inclination of CR Cha indicates that the light curve is expected to be not purely periodic. We were able to detect a rotational period in our analysis of the light curve, but at the same time, additional effects, such as accretion variation and flares also significantly contribute to the observed light curve. The TESS light curve shows a few flare-like events as well, indicated by the blue points in the middle panel of Fig. 1. During the 56-day-long TESS observing period in 2019, we found five flares with durations ranging from 3.8\u22127.2 h, which gives a flare rate of 0.09 days\u22121. However, these were events showing \u223c0.02 mag brightening.","Citation Text":["Bouvier et al. 1986"],"Citation Start End":[[892,911]]}
{"Identifier":"2022AandA...662A..94P__Molina_et_al._(2020)_Instance_1","Paragraph":"In the \u03c30\u2013SFR plane, local (U)LIRGs and high-z galaxies tend to occupy the same region, regardless of their intrinsic difference in terms of morphology, gas fraction and starburstiness. In order to distinguish between normal MS and SB galaxies, we show in Fig. 7, bottom, the velocity dispersion \u03c30 normalised to \u03c30,\u2006MS (solid line in the top panel; \u00dcbler et al. 2019) as a function of the starburstiness \u03b4MS\u2004=\u2004sSFR\/sSFR|MS for all targets already reported in the previous panels. The sSFR|MS is derived from the Speagle et al. (2014) relation, starting from the available stellar mass measurements from the literature for Johnson et al. (2016), F\u00f6rster Schreiber et al. (2018), Molina et al. (2020), Cochrane et al. (2021), and KMOS3D individual targets; stellar masses of LIRGs and ULIRGs from this study, from Bellocchi et al. (2013), Pereira-Santaella et al. (2019) and Crespo et al. (2021), are instead derived from the dynamical mass estimates, assuming M*\u2004=\u2004(1\u2005\u2212\u2005fgas)\u00d7(1\u2005\u2212\u2005fDM)\u00d7Mdyn, where fgas and fDM are the gas and dark matter fractions4, respectively, and Mdyn is the dynamical mass within 2Re (see next section). For the gas fraction, we considered a conservative fgas\u2004=\u20040.1 (Isbell et al. 2018; higher fgas values would further increase their \u03b4MS). This gas fraction is consistent with the estimate we obtain considering the molecular gas mass inferred by ALMA data (Paper II; Lamperti et al., in prep.) and the M* measurements available for a few PUMA targets (see Table 2), \n\n\n\n\n\n\nf\n\n\n\u00af\n\n\ngas\n\n=\n0.11\n\u00b1\n0.05\n\n\n$ {\\bar{f}}_{\\mathrm{gas}} = 0.11 \\pm 0.05 $\n\n\n. For the dark matter fraction (within 2Re) we assumed fDM\u2004=\u20040.26, defined as 1\u2005\u2212\u2005Mbar\/Mdyn, with Mbar\u2004=\u2004M*\u2005+\u2005Mgas\u2004=\u2004M*\/(1\u2005\u2212\u2005fgas). The fDM estimate was derived for the PUMA and Johnson et al. (2016) SB galaxies for which stellar masses are available, and considering the dynamical masses within 2Re (see Sect. 3.7). Because of these assumptions, we considered a factor of 3 uncertainties for the stellar mass measurements of (U)LIRGs. These uncertainties, however, play a minor role in the derived \u03b4MS: at low z, the MS has a soft slope, and normal and massive MS galaxies have similar SFRs (e.g. at z\u2004\u223c\u20040.1, galaxies of 1010 and 1011\u2006M\u2299 have SFR|MS\u2004\u223c\u20041 and \u223c3.5\u2006M\u2299 yr\u22121, respectively); on the other hand, local (U)LIRGs have much higher SFRs, from 10 s to 100 s M\u2299 yr\u22121, and therefore \u03b4MS of the order of 10\u2005\u2212\u2005100. Finally, for the Alaghband-Zadeh et al. (2012) and Harrison et al. (2012) high-z galaxies we assumed that M*\u2004=\u20041011\u2006M\u2299, following Harrison et al. (2012).","Citation Text":["Molina et al. (2020)"],"Citation Start End":[[679,699]]}
{"Identifier":"2019ApJ...872..151M__Jiang_et_al._2016_Instance_1","Paragraph":"Regardless of the process or combination of processes responsible for generating the emission, the kinetic energy of the returning debris must eventually be dissipated in order to be observed. Even if some energy is deposited by circularization at large distances (Piran et al. 2015), the energy will be primarily dissipated by processes that operate closest to the black hole simply because the velocities there are the greatest. However, this would imply that most of the radiation would be emitted at very high energies (X-rays), and instead we observe many TDEs with significant (and sometimes dominant) optical\/UV flux. A reprocessing layer, either static or outflowing (Miller 2015; Metzger & Stone 2016), can help explain the observed emission by reprocessing the luminosity generated by the various dissipation processes at play (Loeb & Ulmer 1997; Ulmer et al. 1998; Bogdanovi\u0107 et al. 2004; Guillochon et al. 2014; Jiang et al. 2016; Coughlin & Begelman 2014; Strubbe & Quataert 2009). The reprocessing of the radiation has also been used to successfully explain the line ratios observed in PS1-10jh (Gaskell & Rojas Lobos 2014; Roth et al. 2016). In this work we assume a simple blackbody photosphere for the reprocessing layer, so that the observed flux becomes\n8\n\n\n\n\n\nwith an effective blackbody temperature\n9\n\n\n\n\n\nIn the above equations, F\u03bd is the specific flux, Rphot is the photospheric radius, D is the distance from the source, L is the bolometric luminosity from our fit, and Teff is the temperature of the photosphere. Most observations of TDEs have thermal temperatures that do not exhibit significant variation. For blackbody emission, the radius must increase as the luminosity (and \n\n\n\n\n\n) increases and decrease as the luminosity decreases, in order for the temperature not to change significantly as the luminosity evolves. This simple behavior also explains the rise in temperatures at late times as the photospheric radius decreases and the bulk of the observed radiation shifts to higher energies. To model this dependence, we assume that the radius of the photosphere has a power-law dependence on the luminosity and fit for both the power-law exponent l and radius normalization Rph0,\n10\n\n\n\n\n\nHere ap is the semimajor axis of the accreting mass at peak \n\n\n\n\n\n. This provides a reasonable typical scaling for the radius of the photosphere, with a minimum photosphere size set by Risco and a maximum photosphere size set by the semimajor axis of the accreting mass.","Citation Text":["Jiang et al. 2016"],"Citation Start End":[[924,941]]}
{"Identifier":"2019AandA...627A.114K__Gorman_et_al._(2015)_Instance_2","Paragraph":"Feature VY may not be a dusty clump in the same sense as the other features since it has a significant flux contribution from the star. In our model, VY is modeled as dusty medium of a Gaussian density distribution centered on VY CMa and uniformly surrounding the star, as in an idealized spherical outflow. Alternatively, however, it could be a clump located in front of or behind the star and seen along the same line of sight. In our implementation, we clear the central space of dust within a radius of 90 AU (equivalent to a diameter of \n\n\n\n0\n\n.\n\u2033\n\n15\n\n\n$ 0{{\\overset{\\prime\\prime}{.}}}15 $\n\n\n) where solids would be too warm to exist heated by the cool but very luminous star. The peak surface brightness simulated with a beam of 152 mas are 34 and 140 mJy beam\u22121 in Bands 6 and 7, respectively. Of that, 13.0 and 30.1 mJy beam\u22121 comes from the blackbody radiation of the model star. O\u2019Gorman et al. (2015) noted that stellar fluxes are underestimated in a blackbody model at millimeter and submillimeter wavelengths owing to the unaccounted presence of an extended radiophotosphere. We extrapolated fluxes expected from their radiophotosphere model getting 12 mJy in Band 6 and 44 mJy in Band 7. Given the small difference with our assumed blackbody fluxes and uncertainties in the radiophotosphere model of O\u2019Gorman et al. (2015), we did not correct our model for this effect. We required 0.014 M\u2299 of dust to explain the observed fluxes which is rather large compared to most other clumps (except C). The stellar contribution partially explains the very low spectral index of 1.7 toward VY because the stellar spectrum is expected to have \u03b1\u2004\u2248\u20042 (Lipscy et al. 2005). At the modeled parameters of VY, the medium must be very thick at visual light and obscure the stellar photosphere from a direct view for a terrestrial observer, as visual observations indeed seem to indicate. The stellar photosphere and its light variations are seen mainly through light scattered in the surrounding medium with a high degree of inhomogeneity (Humphreys et al. 2005).","Citation Text":["O\u2019Gorman et al. (2015)"],"Citation Start End":[[1315,1337]]}
{"Identifier":"2017AandA...601A..72I__Kobayashi_&_Tanaka_2010_Instance_1","Paragraph":"Small grains, which contribute most to infrared emission, are removed by collisional fragmentation and blown out by radiation pressure. The removal timescale is much shorter than the ages of host stars. Disruptive collisions among underlying large bodies, which are called planetesimals, produce smaller bodies and collisional fragmentation among them results in even smaller bodies. This collisional cascade continues to supply small grains. The evolution of debris disks has been explained by the steady-state collisional cascade model (e.g., Wyatt 2008; Kobayashi & Tanaka 2010): the total mass of bodies decreases inversely proportional to time t. Therefore, the excess ratio (Fdisk\/F\u2217) is given by (2)\\begin{equation} \\frac{F_{\\rm disk}}{F_{*}} = \\frac{t_0}{t},\\label{cc} \\end{equation}FdiskF\u2217=t0t,where t0 is the dissipation timescale that is determined by the collisional cascade. Under the assumption of the steady state of collisional cascade, the power-law size distribution of bodies is analytically obtained and the power-law index depends on the size dependence of the collisional strength of bodies (see Eq. (32) of  Kobayashi & Tanaka 2010). In the obtained size distribution, erosive collisions are more important than catastrophic collisions (see Fig. 10 of Kobayashi & Tanaka 2010). Taking into account the size distribution and erosive collisions, we derive t0 according to the collisional cascade (see Appendix E for derivation), (3)\\begin{eqnarray} t_0&amp;\\sim&amp; 1.3 \\left( \\frac{s_{\\rm p}}{\\rm 3000\\,km} \\right)^{0.96} \\left( \\frac{R}{\\rm 2.5\\,au} \\right)^{4.18}\\nonumber\\\\ &amp;&amp;\\quad\\times \\left(\\frac{\\Delta R}{0.4 R}\\right) \\left( \\frac{e}{\\rm 0.1} \\right)^{-1.4} {\\rm Gyr},\\label{eq:t0} \\end{eqnarray}t0~1.3sp3000\u2009km0.96R2.5\u2009au4.18where sp is the size of planetesimals, R is the radius of the planetesimal belt, and e is the eccentricity of planetesimals. Interestingly, t0 is independent of the initial number density of planetesimals (Wyatt et al.  2007). Note that the perturbation from Moon-sized or larger bodies is needed to induce the collisional fragmentation of planetesimals (Kobayasi & L\u00f6hne 2014), which is implicitly assumed in this model. ","Citation Text":["Kobayashi & Tanaka 2010"],"Citation Start End":[[557,580]]}
{"Identifier":"2016AandA...593A..42T__Gonz\u00e1lez_et_al._1998_Instance_1","Paragraph":"Molecular clouds can be formed either directly from neutral medium or by assembling pre-existing, cold molecular clumps. The first scenario is commonly accepted (e.g., Hollenbach & McKee 1979). The second scenario is outlined by Pringle et al. (2001), who proposed that clouds formed out of pre-existing, CO-dark, molecular gas. Compared to the first scenario, the second scenario allows for fast cloud formation in a few Myr, which was suggested by observations (Beichman et al. 1986; Lee et al. 1999; Hartmann et al. 2001). The key problem for the second evolutionary scenario is how molecular gas can exist before it is collected together. Pringle et al. (2001) argued that the pre-existing molecular gas should be cold (10 K) and was shielded from photodissociation by patch dust with AV of ~0.5 mag (White et al. 1996; Berlind et al. 1997; Gonz\u00e1lez et al. 1998; Keel & White 2001). Under AV of ~0.5 mag, H2 should exist substantially while CO is in very low abundance. This is because the self-shielding threshold of AV = 0.02 and 0.5 mag (Wolfire et al. 2010) are widely considered as a condition for maintaining a stable population of abundant H2 and CO gas, respectively. Listz & Pety (2012) detected strong CO(1\u22120) emission (4\u22125 K) in regions with equivalent visual extinction less than 0.5 mag. Two obvious possibilities are 1) the CO gas is transient. Such gas may also have been seen in Goldsmith et al. (2008), who detected a 40% of the total CO mass in Taurus in regions with low to intermediate AV (mask 0 and 1 in their terminology). 2) Such CO gas lies in a highly clumpy medium with lower apparent averaged extinction when photons travel through lower density interclump medium. When small agglomerations of molecular gas are compressed and heated by shock, e.g., in a spiral arm, they become detectable. This scenario is supported by observations of GMC formation in spiral arms (Dobbs 2008) and simulations of molecular clouds formation (e.g., Clark et al. 2012). ","Citation Text":["Gonz\u00e1lez et al. 1998"],"Citation Start End":[[845,865]]}
{"Identifier":"2015MNRAS.451..353R___2010_Instance_1","Paragraph":"UGC 4284 is classified as Scd(s) in the RC3 (de Vaucouleurs et al. 1991), and as SABc by Paturel et al. (2003). Kandalian & Kalloghlian (1998), Chapelon, Contini & Davoust (1999), and Cabrera-Lavers & Garz\u00f3n (2004) found evidence of a central isophotal distortion that they explain postulating the existence of a bar with diameter of 38 arcsec. Other authors did not confirm the presence of a bar in UGC 4284, nevertheless they encountered non-circular and off-plane gas motion in the inner region of UGC 4284 that they discuss as originated by a warped gaseous disc (Coccato et al. 2004; J\u00f3zsa 2007). Bearing in mind the purpose of the present study, it is important to determine which kind of baryonic components do actually exist for UGC 4284, and given the incongruity of the above-reported results we required to perform 2D isophotal analysis of the surface brightness of this galaxy, using galfit (Peng et al. 2002, 2010) and considering at least three principal components: disc, bulge, and bar. UGC 4284 has been observed with Hubble Space Telescope (HST) in the filters F555W (WFPC2 V), F814W (WFPC2 wide I), and F606W (WFPC2 wide V), none the less the filters F555W and F814W show only a small part of the periphery of the galaxy. In contrast the filter F606W exhibits a larger portion of the disc of UGC 4284, so we use that filter to decompose the disc of UGC 4284. The results of the 2D exponential disc fit are presented in Table 1. In the case of UGC 4284, we also examined the possibility of a bulge and a bar components given the inconsistency in the morphological classification of this galaxy. We use the Spitzer IRAC 4.5 \u03bcm image because UGC 4284 is indiscernible in the 2MASS J, H, K bands, and in all the SDSS filters are too faint to allow an obvious separation from the background. Neither the S\u00e9rsic component nor the Ferrers (Ferrers 1877) component produces acceptable fit parameters to reveal the presence of a bulge and a bar, and for this reason we conclude that UGC 4284 has only an exponential disc component in total agreement with the RC3 classification.","Citation Text":["Peng et al.","2010"],"Citation Start End":[[904,915],[922,926]]}
{"Identifier":"2016ApJ...824...92H__DeMastus_et_al._1973_Instance_1","Paragraph":"The idea of ionized material arising from solar flares and causing geomagnetic storms appears to date from around the 1930s (Hale 1931; Dellinger 1937; Newton 1939, 1943; Chapman 1950) and the idea persisted into the space age, where interplanetary shocks were attributed to being caused by flares (e.g., Sonnet et al. 1964; Gosling et al. 1968). By the 1970s, however, some workers had begun to express doubts that solar flares could be the source for interplanetary shocks, geomagnetic storms, and, after their discovery (Tousey 1973), CMEs observed by coronagraphs. Evidence against the flare being the source included: the absence of a 1:1 relationship between geomagnetic storms and flares (e.g., DeMastus et al. 1973; Gosling et al. 1974; MacQueen 1980; Joselyn & McIntosh 1981); the energy required to launch the mass ejection was much larger than that of the flare (e.g., MacQueen 1980; Webb et al. 1980); the location of the flare being at only one footpoint of the larger expanding structure (Harrison & Simnett 1984; Simnett & Harrison 1984, 1985); the lack of correlation between the time of the projected onset of the CME and that of the flare (Harrison et al. 1985, 1990; Harrison & Sime 1989); and a lack of lateral movement of the expanding bright front as its leading edge expanded outwards (Sime et al. 1984). These and other measurements were presented in a review paper by Kahler (1992) and another by Gosling (1993a), the latter now well known as the \u201csolar flare myth\u201d paper. In that paper, flares are described as being either connected to CMEs via a common primary physical process (labeled in his Figure 16 as \u201cevolving solar magnetic fields\u201d), or as being caused by the CME launch. By around 1996 Gosling\u2019s \u201cmodern paradigm\u201d was widely accepted as the correct general physical model for CMEs. This picture has changed little in the last two decades, apart from perhaps the general acceptance that a magnetic flux rope is the driving core of a CME, although this too is an idea dating from before the space age (e.g., Morrison 1954; Cocconi et al. 1958; Gold 1959).","Citation Text":["DeMastus et al. 1973"],"Citation Start End":[[702,722]]}
{"Identifier":"2015MNRAS.447.2198E__Buote_&_Humphrey_2012_Instance_1","Paragraph":"Both for the hydrodynamical and the synthetic simulations, a blind analysis was performed on the simulated images. First, surface-brightness profiles were extracted by computing the azimuthal median in concentric annuli. The profiles were then deprojected using a non-parametric onion-peeling technique (Kriss, Cioffi & Canizares 1983; McLaughlin 1999) and converted into radial density profiles. For the synthetic simulations, we found a very good agreement between the true and the reconstructed density profiles, in all of the situations considered here. Even when adding as many as 103 randomly distributed clumps in the simulated image, our method is able to recover the underlying density profile with an accuracy of a few\u2009per\u2009cent at all radii. The addition of significant triaxiality to the halo did not affect our results (in agreement with the findings of Buote & Humphrey 2012), and similar conclusions were reached for the inclusion of a filamentary structure. In the case of the hydrodynamical simulations, we show in Fig. 1 the average surface-brightness profiles (in arbitrary units) obtained using our azimuthal median method and with the standard azimuthal mean. We can see that the azimuthal median results in a much more regular profile, showing the expected smooth radial decline. In Fig. 2 we show the comparison between the average true 3D density profile in the sample of Vazza et al. (2011b) and the average profile obtained with our method through deprojection. On the same figure, we also show the same comparison in the case of one particularly perturbed cluster, E2 (see Vazza et al. 2011b). As shown in the bottom panel, the deviations are typically at the level of 10 per cent or less, even in the case of the very perturbed system E2, and do not show any particular trend with radius. A similar result is found for the integrated gas mass. This demonstrates the ability of our method to recover the true density profiles from X-ray observations in the presence of inhomogeneities.","Citation Text":["Buote & Humphrey 2012"],"Citation Start End":[[866,887]]}
{"Identifier":"2019AandA...629A..54U__Marinucci_et_al._2015_Instance_2","Paragraph":"NGC 2110. NGC 2110 is another nearby (z\u2004=\u20040.00779, Gallimore et al. 1999), X-ray bright Seyfert galaxy. Diniz et al. (2015) report a black hole mass of \n\n\n\n2\n.\n\n7\n\n\u2212\n2.1\n\n\n+\n3.5\n\n\n\u00d7\n\n10\n8\n\n\n\nM\n\u2299\n\n\n\n$ 2.7^{+ 3.5}_{- 2.1} \\times 10^{8}\\,{{M}_{\\odot}} $\n\n\n, from the relation with the stellar velocity dispersion. From BeppoSAX data, Malaguti et al. (1999) found the X-ray spectrum to be affected by complex absorption. This has been later confirmed by Evans et al. (2007), who find the Chandra+XMM\u2013Newton data to be well fitted with a neutral, three-zone, partial-covering absorber. Rivers et al. (2014) find the Suzaku data to be well fitted with a stable full-covering absorber plus a variable partial-covering absorber. A soft excess below 1.5 keV is also present (Evans et al. 2007), and possibly due to extended circumnuclear emission seen with Chandra (Evans et al. 2006). No Compton reflection hump has been detected with Suzaku (Rivers et al. 2014) or NuSTAR (Marinucci et al. 2015), despite the presence of a complex Fe K\u03b1 line. According to the multi-epoch analysis of Marinucci et al. (2015), the Fe K\u03b1 line is likely the sum of a constant component (from distant, Compton-thick material) and a variable one (from Compton-thin material). Concerning the high-energy cut-off, ambiguous results have been reported in literature (see Table 1). Ricci et al. (2017) report a value of \n\n\n\n448\n\n\u2212\n55\n\n\n+\n63\n\n\n\n$ 448^{+63}_{-55} $\n\n\n keV, while Lubi\u0144ski et al. (2016) report a coronal temperature of \n\n\n\n230\n\n\u2212\n57\n\n\n+\n51\n\n\n\n$ 230^{+51}_{-57} $\n\n\n keV and an optical depth of \n\n\n\n0\n.\n\n52\n\n\u2212\n0.13\n\n\n+\n0.14\n\n\n\n\n$ 0.52^{+ 0.14}_{- 0.13} $\n\n\n. From 2008\u20132009 INTEGRAL data, Beckmann & Do Cao (2010) report a cut-off of \u223c80 keV with a hard photon index, but these results are not confirmed by NuSTAR (Marinucci et al. 2015). Indeed, only lower limits to the high-energy cut-off have been found with NuSTAR (210 keV: Marinucci et al. 2015), Suzaku (250 keV: Rivers et al. 2014) and BeppoSAX (143 keV: Risaliti 2002). No hard X-ray spectral variability has been detected by Caballero-Garcia et al. (2012) and Soldi et al. (2014) from BAT data, despite the significant flux variability.","Citation Text":["Marinucci et al. (2015)"],"Citation Start End":[[1077,1100]]}
{"Identifier":"2022ApJ...930...30Z__McIntosh_et_al._2014_Instance_1","Paragraph":"Figure 4(a) shows a time\u2013latitude diagram of the subsurface toroidal flux density calculated by integrating the toroidal field over the range of 0.7\u20131.0R\n\u2299. The toroidal field is antisymmetric because the dipolar parity always prevails in our model. The latitudinal migration pattern of the toroidal field is different from that of FTD models, which usually show concentrated flux at high latitudes due to the radial shear in the polar portion of the tachocline. For each cycle, a weak toroidal field starts from high latitudes along with another overlapped branch of the previous cycle near the equator. With the equatorward migration of the toroidal field, its flux density increases as the cycle evolutes. This scenario is consistent with that of the solar cycle and the extended solar cycle (Howard & Labonte 1980; Wilson et al. 1988; McIntosh et al. 2014). Below 40\u00b0 latitude, the toroidal flux density is larger than 2.5 \u00d7 1023 Mx Radian\u22121. Considering the spatial scale of a typical active region to be a few degrees, these fluxes are comparable with that in active regions. According to Cameron & Jiang (2019), the toroidal field at lower latitudes corresponds to the expected amount of flux residing in active regions. The weaker toroidal field at higher latitudes corresponds to ephemeral regions, which are small bipolar magnetic regions that contain a maximum total flux of the order of 1020 Mx (Harvey & Martin 1973). In Figure 4, the vertical dashed line represents cycle minimum (t\nmin) when the toroidal flux of the last cycle disappears, and vertical dotted line represents cycle maximum (t\nmax) when the toroidal flux of the new cycle reaches a maximum. The interval between t\nmin and t\nmax is 4 yr, which is close to the average period of the rising phase for solar cycles (Hathaway 2015; Jiang et al. 2018). The solid and dashed curves of Figure 5 show the time evolution of unsigned toroidal flux, \n\n\n\n\u222b5090\u00b0\u222b0.7R\u2299R\u2299r\u2223B\u03d5\u2223drd\u03b8\n\n, of each hemisphere for odd and even cycles, respectively. The cycle period is 10.96 yr based on the interval between successive maxima of the toroidal flux. This period is consistent with the average period of observed solar activity (Hathaway 2015). If the cycle period is calculated as the interval between the start and end of the toroidal flux of a cycle, then it is 18.1 yr. This result is consistent with approximately the 18\u201322 yr period of the extended cycle (Wilson et al. 1988). The maximum toroidal flux produced in one hemisphere is about 2.6 \u00d7 1023 Mx, which matches the observed magnetic flux at the solar surface (Cameron & Sch\u00fcssler 2015).","Citation Text":["McIntosh et al. 2014"],"Citation Start End":[[839,859]]}
{"Identifier":"2021ApJ...912..125G__Lieshout_et_al._2016_Instance_1","Paragraph":"ZTF J092311.41+423634.16 (ZTF J0923+4236), the fifth-most variable white dwarf within 200 pc according to our Gaia-ZTF ranking parameter, demonstrates multiple transit-like phenomena of varying durations and depths. The highest-resolved of these features at MJD \u2248 58450 closely resembles a mirror image of the dips observed for ZTF J0139+5245 (Vanderbosch et al. 2020): this object displays a gradual ingress and sharp egress, also contrasting the transit shapes of WD 1145+017 (Figure 6) and the K dwarf KIC 12557548 (Rappaport et al. 2012; van Lieshout et al. 2016). Unlike those in ZTF J0139+5245, the observed transits in ZTF J0923+4236 appear more incongruous and complex, repeating on irregular intervals with varying shapes. The dip at MJD \u2248  58450 appears to have a secondary feature before the onset of the larger dip. This more complicated structure, which appears to be transiting on the order of days, implies ZTF J0923+4236 may be at an evolutionary stage in between WD 1145+017 (roughly 4.5 hr orbital period) and ZTF J0139+5245 (roughly 107 day orbital period). We followed up ZTF J0923+4236 with high-speed photometry on MJD  =  59138 and 59139, using the McDonald 2.1 m telescope to assess short-term variability, observing small amplitude variations on timescales of \u22481 hr (Appendix C). These observations, however, were constrained to less than two hours in length, due to the observing conditions. McDonald 2.1 m observations on MJD  =  59168 spanning 3 hr in length do not show this same possible short-term variability to within the 4\u2329A\u232a significance threshold (Appendix C). We observe broad Balmer lines in a relatively low-resolution follow-up spectrum from HET (see Figure 7), and classify ZTF J0923+4236 as a DA white dwarf. The lack of emission lines disfavors cataclysmic variable activity as the cause of photometric variability, leaving transiting planetary debris as a likely explanation. Future monitoring is suggested to place stronger constraints on the source of these transit-like features and any metal pollution or circumstellar gas around this white dwarf.","Citation Text":["van Lieshout et al. 2016"],"Citation Start End":[[542,566]]}
{"Identifier":"2015ApJ...803L...3P__Kafle_et_al._2014_Instance_1","Paragraph":"Mapping out the distribution of dark matter in our Galaxy rests as a paramount task with potentially far reaching implications for astroparticle physics and cosmology. This is important to understand galaxy formation and to feed searches aimed at unveiling the very nature of dark matter. In particular, direct and indirect dark matter searches rely heavily on the findings of numerical simulations. It is therefore essential to extract the Galactic dark matter distribution directly from observations. In the outer Milky Way (at Galactocentric radii greater than \u223c20 kpc), where baryons contribute little to the total mass budget, the gravitational potential traces closely the dark matter component and the total mass enclosed can be constrained using convenient tracers (e.g., Sakamoto et al. 2003; Dehnen et al. 2006; Xue et al. 2008; Bhattacharjee et al. 2014; Kafle et al. 2014), although with important degeneracies in the tracer population modeling. By contrast, in the inner Galaxy (i.e., in the inner \u223c20 kpc) the baryonic contribution is very significant and its morphology rather uncertain, which makes the evidence for dark matter difficult to establish and the extraction of its distribution a delicate undertaking (Iocco et al. 2015). This has been addressed by many authors with different methods (e.g., Dehnen & Binney 1998; Sofue et al. 2009; Catena & Ullio 2010; Weber & de Boer 2010; Iocco et al. 2011; Bovy & Rix 2013; Nesti & Salucci 2013; Loebman et al. 2014), all of which do, however, make explicit assumptions about the underlying dark matter profile: typically, a multi-parameter profile is fitted to the observations together with a given baryonic component. The class of \u201clocal\u201d methods to measure the dark matter density in the solar neighborhood (e.g., Salucci et al. 2010; Garbari et al. 2012; Smith et al. 2012; Zhang et al. 2013; Read 2014) avoids this bias, yet such methods are not easily applicable elsewhere in the Galaxy. An approach free of profile assumptions has been developed and successfully tested in external galaxies (Persic et al. 1996; Salucci et al. 2007), but never applied to our own Galaxy given the sizeable uncertainties of both kinematic data and baryonic modeling.","Citation Text":["Kafle et al. 2014"],"Citation Start End":[[866,883]]}
{"Identifier":"2022MNRAS.512.4573F__Sana_et_al._2012_Instance_1","Paragraph":"We use the state-of-the-art open-source stellar-evolution code Modules for Experiments in Stellar Astrophysics (mesa, Paxton et al. 2011, 2013, 2015, 2018, 2019) in order to investigate a broader range of rotational velocities and metallicities. Moreover, since very massive stars have been observed in the Tarantula Nebula (e.g. Doran et al. 2013) in the Large Magellanic Cloud (LMC), we consider stellar models with initial masses up to nearly 160 M\u2299 (a factor of 1.3 larger than in Starburst99). Finally, we account for the fact that the majority of stars are born in binary systems (Sana et al. 2012, 2013). The observational evidence for such large binary fraction led to a paradigm change in our understanding of stellar evolution. As one (or both) of the stars in a binary system fill their Roche lobes, a phase of mass transfer takes place. Material is exchanged between the companions through the first Lagrangian point or lost by the system. The outermost layers of the donor are stripped off and eventually accreted on to the secondary star. This significantly alters the masses and spectroscopic appearances of the stars and generates evolutionary sequences otherwise unattainable in a single-star scenario (e.g. De Marco & Izzard 2017), influencing the mass-loss and the rotation rates of stars. In some cases, the interaction and mass exchange can be unstable, leading to merger events. Although numerous uncertainties still exist regarding the modelling of binary systems, it is becoming increasingly clear that more realistic estimates of stellar feedback cannot ignore the impact of stellar multiplicity. This applies also to radiative feedback as interacting binaries enhance the production of hydrogen- and helium-ionizing photons (e.g. Eldridge et al. 2017; G\u00f6tberg, de Mink & Groh 2017) and harden the spectra of a stellar population (G\u00f6tberg et al. 2019). Of particular interest are stripped helium stars, i.e. massive helium stars produced by binary interaction which emit the majority of their light at wavelengths shorter than the Lyman limit (Stanway, Eldridge & Becker 2016; G\u00f6tberg et al. 2019) on a time-scale beyond few Myr.","Citation Text":["Sana et al. 2012"],"Citation Start End":[[587,603]]}
{"Identifier":"2017ApJ...850...97B__Dutton_et_al._2016_Instance_1","Paragraph":"However, recent work by Brook & Di Cintio (2015a, 2015b) used results from simulations in which stellar feedback processes alter the DM content of dwarf galaxies to show that they can simultaneously match the densities and velocities of observed dwarfs. In this scenario, feedback from stars and supernovae creates bursty star formation histories in dwarf galaxies that fluctuate the gravitational potential well at the center of the dwarf (Pontzen & Governato 2012, 2014; Teyssier et al. 2013; Chan et al. 2015; O\u00f1orbe et al. 2015; Dutton et al. 2016). DM core creation leads to a better match between theory and observed rotation curves (Katz et al. 2017; Santos-Santos et al. 2017). Feedback is particularly effective in dwarf galaxies with halo masses of a few 1010 \n\n\n\n\n\n (Governato et al. 2012; Di Cintio et al. 2014b), where it can transform an initially steep inner DM density profile into a flatter \u201ccored\u201d profile. At lower halo masses, there is less star formation, leading to less energy injection and lower core formation efficiencies (Pe\u00f1arrubia et al. 2012; Maxwell et al. 2015). At higher masses, the deeper potential wells of galaxies make core formation increasingly difficult (Di Cintio et al. 2014b; Pontzen & Governato 2014), at least if an additional source of feedback is neglected, such as active galactic nuceli (AGN) (Martizzi et al. 2013). In this model with baryonic feedback, it is possible to assign dwarf galaxies to relatively massive halos, despite the low rotational velocities measured from their spatially resolved stellar kinematics. This is because baryonic feedback can push DM out of the central regions, lowering the enclosed mass at the radii that stellar kinematics probe (but without affecting the total halo mass8\n\n8\nModulo a slight reduction in halo mass caused by the loss of baryons or preventive feedback (Munshi et al. 2013; Sawala et al. 2013).\n). Hence, the densities are lowered, and the apparent velocities of the galaxies, reconciling the observations with theory.","Citation Text":["Dutton et al. 2016"],"Citation Start End":[[533,551]]}
{"Identifier":"2015ApJ...811L..13W__Qiu_et_al._2012_Instance_1","Paragraph":"In summary, we have found quantitative evidence of thermal conduction suppression in a hot flare loop by coronal seismology of the slow-mode waves. This result suggests that the flare loop should cool much slower than expected from the classical Spitzer conductive cooling. Our studied event is indeed of such long-duration events (LDEs) which are the flares with a slower-than-expected decay rate in soft X-ray and EUV radiation (e.g., Forbes & Acton 1996; Takahashi & Watanabe 2000; Qiu et al. 2012). To explain the LDEs and flare loop-top sources, some previous studies have suggested the mechanism of continuous heating (Warren 2006; Liu et al. 2013; Sun et al. 2013) or conduction suppression (McTiernan et al. 1993; Jiang et al. 2006; Li et al. 2012). Our study confirms the effect of the latter mechanism in a more direct way. Laboratory experiments and numerical studies showed that the actual conductivity is smaller (by at least a factor of two) than that given by Spitzer when l \u2272 30\u03bb where l is the temperature gradient scale length and \u03bb the mean free path of thermal electrons (Bell et al. 1981; Luciani et al. 1983). We estimate l\/\u03bb \u223c 30 in this case, suggesting that the nonlocal conduction may account for the observed conduction suppression (Matte & Virmont 1982; Rosner et al. 1986). By studying the evolution of flare loop-top sources, Jiang et al. (2006) suggested that plasma waves or turbulence may play an important role in suppressing the conduction during the decay phase of flares. The mechanism is similar to that used for interpreting the significant reduction of thermal conductivity in galaxy-cluster cooling flows by a tangled magnetic field (Chandran & Cowley 1998). Finally, we conclude that the result of conduction suppression may also shed light on the coronal heating problem (see the review by Klimchuk 2015), because weak thermal conductivity implies smaller conductive losses, and an extended lifetime of individual nanoflares, increasing the average coronal temperature for the same heating rate.","Citation Text":["Qiu et al. 2012"],"Citation Start End":[[485,500]]}
{"Identifier":"2020MNRAS.496..708L__Lusso_et_al._2010_Instance_1","Paragraph":"Following the approach described in Risaliti & Lusso (2015), there exists a relation between the luminosities in the X-rays (LX) and UV band (LUV)\n(8)$$\\begin{eqnarray*}\r\n\\log (L_\\mathrm{ X})=\\hat{\\gamma }\\log (L_{\\mathrm{ UV}})+\\beta ^{\\prime },\r\n\\end{eqnarray*}$$where $\\hat{\\gamma }$ and \u03b2\u2032 denote the slope parameter and the intercept. Combining equation (8) with the well-known expression of $L=F\\times 4\\pi D_L^2$, the luminosity distance can be rewritten as\n(9)$$\\begin{eqnarray*}\r\n\\log (D_L)=\\frac{1}{2-2\\hat{\\gamma }}\\times [\\hat{\\gamma }\\log (F_{\\mathrm{ UV}}) - \\log (F_\\mathrm{ X}) + \\hat{\\beta }],\r\n\\end{eqnarray*}$$a function of the respective fluxes (F), the slope parameter ($\\hat{\\gamma }$), and the normalization constant ($\\hat{\\beta }=\\beta ^{\\prime }+(\\hat{\\gamma }-1)\\log _{10} 4\\pi$). Therefore, from theoretical point of view, the luminosity distance can be directly determined from the measurements of the fluxes of FX and FUV, with a reliable knowledge of the dispersion \u03b4 in this relation and the value for the two parameters ($\\hat{\\gamma }$, $\\hat{\\beta }$) characterizing the LX\u2013LUV relation. However, it has been established that the LX\u2013LUV relation was characterized by a high dispersion. Through the analysis of different quasar samples with multiple observations available, previous works derived a consistent value for the slope parameter ($\\hat{\\gamma }=0.599\\pm 0.027$) and the intrinsic dispersion of the relation (0.35\u20130.40 dex) (Lusso et al. 2010; Young, Risaliti & Elvis 2010). It was found in subsequent analysis quantifying the observational effects (Lusso & Risaliti 2016) that the magnitude of the intrinsic dispersion can be eventually decreased to the level of 0.15 dex. They identified a subsample of quasars without the major contributions from uncertainties in the measurement of the (2\u2009keV) X-ray flux, absorption in the spectrum in the UV and in the X-ray wavelength ranges, variability of the source and non-simultaneity of the observation in the UV and X-ray bands, inclination effects affecting the intrinsic emission of the accretion disc, and the selection effects due to the Eddington bias (Risaliti & Lusso 2019). Besides the lower dispersion in the relation, the reliability and effectiveness of the method strongly depend on the lack of evolution of the relation with redshift (Bisogni et al. 2018). Finally, Risaliti & Lusso (2019) produced a final, high-quality catalogue of 1598 quasars, by applying several filters (X-ray absorption, dust-reddening effects, observational contaminants in the UV, and Eddington bias) to the parent sample from the Sloan Digital Sky Survey (SDSS) quasar catalogues (Shen et al. 2011; Paris et al. 2017) and the XMM\u2013Newton Serendipitous Source Catalogue (Rosen et al. 2016). The final results indicated that such refined selection of the sources could effectively mitigate the large dispersion in the LX\u2013LUV relation, with a tractable amount of scatter avoiding possible contaminants and unknown systematics (see Risaliti & Lusso 2019 for more details). More importantly, the similar analysis has supported the non-evolution of LX\u2013LUV relation with the redshift, which is supported by the subsequent study involving the intercept parameter $\\hat{\\beta } = 8.24\\pm 0.01$, the slope parameter $\\hat{\\gamma }=0.633\\pm 0.002$, and smaller dispersion $\\hat{\\delta }=0.24$ in a new, larger quasar sample (Risaliti & Lusso 2019).","Citation Text":["Lusso et al. 2010"],"Citation Start End":[[1469,1486]]}
{"Identifier":"2017MNRAS.470.2959K___2013_Instance_2","Paragraph":"At large distances, the proper motions of the halo stars are either unreliable or generally unavailable, which hinders a direct measurement of their velocity dispersions. However, our off-centric location in the Galaxy means that the galactocentric radial (r) and heliocentric radial (s) directions are not the same. This difference is more significant in the inner halo, at a distance of r \u2272 a couple of times of R0, where R0 is the distance of the Sun from the Galactic Centre. Hence, in the inner halo the observed line-of-sight velocities of the stars can be expressed in terms of all three orthogonal galactocentric velocities (vr, v\u03b8, v\u03d5), or in other words the line-of-sight velocities have some contribution from the tangential galactocentric velocities. Provided we have a model that well represents the distribution of the halo stars, we can fit a model marginalized over the unknown tangential motions to the available four-dimensional data (position vector and a line-of-sight velocity), and thus estimate the velocity moments of the system. In the absence of proper motion, the approach of estimating moments of the velocities has been extensively used to predict the kinematics of the MW halo. For example, Sirko et al. (2004), Kafle et al. (2012), Kafle et al. (2013), Kafle et al. (2014) and King et al. (2015) fit an ellipsoidal distribution of velocities and similarly, Deason, Belokurov & Evans (2011a) apply an alternative power-law model to derive the halo kinematics. Using the marginalisation scheme, Kafle et al. (2012, 2013) studied halo Blue Horizontal Branch stars (BHBs), Kafle et al. (2014) studied both BHBs and K-Giant stars (KGs), while King et al. (2015) analysed a mixed bag of BHB and F-type stars to cumulatively construct the velocity dispersion and anisotropy of the outer halo. Interestingly, Kafle et al. (2012) found that the velocity anisotropy parameter of the Galactic halo is non-monotonic and has a prominent dip at a galactocentric radius of r \u2243 18\u2009kpc. In their studies, King et al. (2015) find that the value of \u03b2 at 15 \u2272 R\/kpc \u2272 25 is more tangentially biased, which they attribute to the difference in the spatial resolutions of the data sets and adopt a marginalization technique. A varying level of undulations in the anisotropy parameter has also been observed in simulated haloes (Rashkov et al. 2013; Loebman et al. 2017). There are a number of proposed scenarios that could explain such a feature, e.g. a transition from inner to outer halo or a local shell-like structure at the given radius. Moreover, it can also be due to the unrelaxed stars dispersed from the kinematically coherent satellite galaxies that are aligned with kinematically coherent planar structures; assuming that such planar structures have strong rotation as suggested by Ibata et al. (2013), Pawlowski, McGaugh & Jerjen (2015), Libeskind et al. (2015) and Ibata et al. (2015), etc. Recently, Loebman et al. (2017) suggest that a major merger as early as redshift z \u223c 1 can also result in a tangential dip in the value of \u03b2 over a wide range of radii. While Bird & Flynn (2015) suggest that such a feature in the velocity anisotropy run of the halo is a transitory phase, Loebman et al. (2017) conclude that such dips are long-lived in the in situ stellar halo. In any case, there is currently no consensus view as to what causes such velocity anisotropy changes. Finally, in the outer halo there have been recent attempts to utilize multi-epoch Hubble Space Telescope data to estimate the halo velocity dispersion. In particular, recently Cunningham et al. (2016) used the Galactic foreground stars along the M31 galaxy and found that the halo is isotropic at r\/kpc \u223c 25. In Fig. 10, we summarize the recent (this paper inclusive) measurements of the halo velocity anisotropy.","Citation Text":["Kafle et al.","2013"],"Citation Start End":[[1524,1536],[1544,1548]]}
{"Identifier":"2020AandA...639A..46B__\u0160tver\u00e1k_et_al._(2009)_Instance_3","Paragraph":"The linear relationship that we observe between breakpoint energy and core temperature is in line with previous measurements (e.g. McComas et al. 1992; \u0160tver\u00e1k et al. 2009), for both the halo and strahl. According to Scudder & Olbert (1979), a linear trend in the halo relation also follows under the assumption that binary Coulomb collisions dominate electron dynamics in the solar wind. However, in order to align with available experimental data, Scudder & Olbert (1979) set a scaling factor of Ebp\/kBTc = 7, which differs from our scaling factor of Ebp\/kBTc = 5.5 \u00b1 0.1. With a scaling factor of Ebp\/kBTc = 7, Scudder & Olbert (1979) predict that a transformation of thermal electrons into the suprathermal population occurs as the solar wind flows out from the Sun. Findings by \u0160tver\u00e1k et al. (2009), on the other hand, show that the (nh\u2005+\u2005ns)\/nc ratio remains roughly constant with heliocentric distance in the slow wind, suggesting a lack of interchange between the thermal and suprathermal populations. However \u0160tver\u00e1k et al. (2009) observes some variability in the (nh\u2005+\u2005ns)\/nc ratio in the fast wind, which they attribute to either statistical effects due to a lack of samples or a possible \u201cinterplay\u201d between thermal and suprathermal electrons. Scudder & Olbert (1979) also predict that the halo Ebp\/kBTc ratio remains constant with heliocentric distance, whereas \u0160tver\u00e1k et al. (2009) find that the halo Ebp\/kBTc ratio decreases with heliocentric distance. These findings by \u0160tver\u00e1k et al. (2009), along with the discrepancy between our calculated ratio of Ebp\/kBTc = 5.5 \u00b1 0.1 and the prediction of Ebp\/kBTc = 7, suggest that the model of Scudder & Olbert (1979) requires a minor update to either the theory or to the input parameters. The discrepancy, however, may also be indicative of other processes, such as wave-particle scattering (e.g. Gary et al. 1994), that possibly modifies the ratio between breakpoint energy and core temperature while preserving its linear relationship.","Citation Text":["\u0160tver\u00e1k et al. (2009)"],"Citation Start End":[[1019,1040]]}
{"Identifier":"2022AandA...666A.153D___1999_Instance_1","Paragraph":"The inventory of solid-state material sometimes combines ice feature investigation with radiative transfer codes to simulate the observed spectral energy distribution and\/or chemical models (e.g. Pontoppidan et al. 2005; Ballering et al. 2021). The ice band spectroscopic profiles observed at medium to high spectral resolution intrinsically contain information that can be used to constrain the extent of grain growth since they are affected by this growth. Carbon dioxide displays several characteristics that are particularly interesting for probing grain growth. It is one of the main and ubiquitous ice mantle constituents, along with water and, depending on the line of sight, carbon monoxide. The carbon dioxide stretching mode around 4.27 \u00b5m possesses a fairly narrow absorption band with a typical full width at half maximum (FWHM) of tens of cm\u22121, depending on the exact ice mixture environment (e.g. Ehrenfreund et al. 1996, 1999), whereas the water ice FWHM is several hundred cm\u22121. In addition to the relatively high contrast expected in the CO2 ice profile due its narrowness, carbon dioxide absorbs in a relatively clean region of the infrared spectrum. For the absorption band of water ice centred at 3.1 the red wing of the profile is modified not only by grain growth but also by additional absorption from, for example, methanol and the 3.47 \u00b5m band assigned to the presence of ammonia in the water mantle. The carbon monoxide stretching mode, lying at slightly higher wavelengths than that of carbon dioxide, is also in a relatively clean region. Some sources show a significant absorption at 4.62 \u00b5m, attributed to the presence of OCN\u2212, that can mainly affect the blue side of the CO absorption profile. It has been investigated and discussed in Dartois (2006), where it was shown that grain growth to micron sizes can still produce an observed large red component in its absorption profiles towards some lines of sight. Some young stellar object (YSO) spectra can also harbour hydrogen lines in emission, such as Pfund \u03b2 (4.654 \u00b5m) and Brackett \u03b1 (4.051 \u00b5m), which has to be taken into account. Their contribution can either be estimated from the set of observed hydrogen lines and\/or taken out of the profile analysis if significant for spectra with high enough spectral resolution given their narrow profile widths. The integrated absorption cross-section of the carbon dioxide band is relatively high, higher than carbon monoxide, an additional factor that makes it a good target for investigating how grain growth affects spectroscopic band profiles.","Citation Text":["Ehrenfreund et al.","1999"],"Citation Start End":[[911,929],[936,940]]}
{"Identifier":"2022MNRAS.510.3039K__Sakai_et_al._2018_Instance_2","Paragraph":"Finally, our models assume the absence of planetary magnetic fields. The early paradigm considering the evolution of terrestrial planets has implied that the planetary magnetic field is necessary to protect planetary atmospheres and reduce the atmospheric mass loss (see e.g. Dehant et al. 2007, and references therein). The later studies, however, show that this point of view is ambiguous. Thus, the effect of the magnetic field on the atmospheric escape can be considered as a result of the two concurring processes: reducing the escape by capturing the ionised atmospheric species within the closed magnetic field lines, and enhancing the escape of the atmospheric ions through the regions of the open magnetic lines (polar cusps, in the case of a dipole field) and the reconnection on the night-side (see, e.g. Khodachenko et al. 2015; Sakai et al. 2018; Carolan et al. 2021). Thus, for planets in the Solar System, it was shown both in the observations (Gunell et al. 2018; Ramstad & Barabash 2021) and by modelling (Sakai et al. 2018; Egan et al. 2019) that the presence of a weak magnetic field can intensify atmospheric escape. These results, however, should be taken with caution for young planets, and in particular those in the sub-Neptune range, because of the different atmospheric structures and the non-thermal mechanisms dominating the atmospheric mass loss in the Solar System, which are contrary to the planets considered in this study (see, e.g. Scherf & Lammer 2021, for the discussion). For hot Jupiters, Khodachenko et al. (2015) predict a significant suppression of escape for intrinsic magnetic fields larger than 0.3 G. The model with the closest setup to this study by Carolan et al. (2021) predicts, however, for the 0.7Mjup planet experiencing XUV (thermally) driven atmospheric escape, a small increase in the atmospheric mass-loss rate with increasing dipole field strength (about twice between 0 and 5\u2009G). We therefore expect that the possible effect from the planetary intrinsic magnetic field depends largely on the strength and configuration of the planetary and stellar magnetic fields, but, according to the numbers reported in the literature, might not affect our results dramatically. The lack of studies for close-in sub-Neptune-like planets, however, holds us from making final conclusions.","Citation Text":["Sakai et al. 2018"],"Citation Start End":[[1023,1040]]}
{"Identifier":"2022MNRAS.511.2105K__Husemann_et_al._2019_Instance_1","Paragraph":"The calculation of mass outflow rates, especially in the ionized gas phase, have often come from measurements using integrated fibre or long-slit spectra, where several assumptions are invoked in the outflow modelling. These assumptions, briefly described here, result in \u2018time-averaged global mass outflow rate\u2019 with large systematic uncertainties. First, due to the limitations of the current instruments even on large telescopes, an accurate modelling of the outflow geometry is not possible. This is especially true for high redshift galaxies (z\u223c2) where, with currently available adaptive optics (AO) technology, one can at best achieve a spatial resolution of \u223c2 kpc where the bulk of the outflow might reside (e.g. Brusa et al. 2016; Davies et al. 2020b). Therefore, the outflow geometry is either assumed to be a uniformly filled conical, bi-conical, or spherical thin shells (e.g. Veilleux, Shopbell & Miller 2001; Fischer et al. 2013; Riffel, Storchi-Bergmann & Winge 2013; Ishibashi & Fabian 2015; Thompson et al. 2015; Bae & Woo 2016; Husemann et al. 2019; Mingozzi et al. 2019). Secondly, if the data are obtained from fibre and single-slit spectroscopy, the size of the outflow is largely unconstrained. For long-slit observations, as an example, the outflow size depends on whether the slit is oriented along the outflow direction. This can be mitigated by using integral field spectroscopy (IFS) which is being increasingly used for extragalactic studies (e.g. Liu et al. 2013; Rupke & Veilleux 2013; Harrison et al. 2014; Maiolino et al. 2017; Husemann et al. 2019; Sch\u00f6nell et al. 2019; Rupke, Thomas & Dopita 2021), although there could still be projection effects with the IFS data. Thirdly, accurate determination of electron density and electron temperature is required for the ionized mass outflow rate calculations. Electron density is usually derived from emission lines that arise out of two closely spaced \u2018metastable\u2019 energy levels such as [S\u2009ii]\u03bb\u03bb6716,6731 ([S\u2009ii] doublet hereafter). Density measured from the [S\u2009ii] doublet is sensitive to values between \u223c10 and 5000 cm\u22123, typical in the Narrow Line Region (NLR) of AGN host galaxies (e.g. Osterbrock & Ferland 2006; Perna et al. 2017; Baron & Netzer 2019; Davies et al. 2020a). The [S\u2009ii] doublet is significantly weaker than the lines used to trace ionized outflows such as the [O\u2009iii]\u03bb5007 and H\u2009\u03b1. In high redshift galaxies, it is extremely challenging to detect these doublet lines, despite hours of observations on a single target. Therefore, nominal density values are often assumed in mass outflow rate calculations, resulting in systematic uncertainties of up to 2\u20133 orders of magnitude. Furthermore, the density structure within the outflowing medium is often non-uniform, when resolved in low redshift galaxies (e.g. Kakkad et al. 2018). Therefore assuming a constant density within the outflowing medium often leads to inaccurate outflow rate and kinetic energy values. Collectively, these assumptions result in a systematic uncertainty of approximately 3\u20134 orders of magnitude (e.g. Harrison et al. 2018). This implies that the quoted values of coupling efficiency in the literature have a wide range, with the actual efficiency still an unknown in most of the studies.","Citation Text":["Husemann et al. 2019"],"Citation Start End":[[1047,1067]]}
{"Identifier":"2016ApJ...824...10F__Martin_et_al._2015_Instance_1","Paragraph":"The number of known Local Group galaxies continues to rise in the era of deep, panchromatic, wide-field surveys such as the SDSS (York et al. 2000), Pan-STARRS (Kaiser et al. 2010), the Dark Energy Survey (DES; Diehl et al. 2014), and PAndAS (McConnachie et al. 2009). A complete catalog of Local Group galaxies discovered up to 2012 was compiled by McConnachie (2012);26\n\n26\nAn updated list is maintained online at http:\/\/www.astro.uvic.ca\/alan\/Nearby_Dwarf_Database.html\n.\n\n since then, four new M31 satellites have been discovered: Andromeda XXXI, Andromeda XXXII, and Andromeda XXXIII from Pan-STARRS data (Martin et al. 2013b, 2013a), and Andromeda XXX from PAndAS data (Conn et al. 2012; M. J. Irwin et al. 2016, in preparation). More dramatically, 23 new \u201cultrafaint\u201d satellites have since been discovered in the immediate vicinity of the Milky Way. Four of these detections were made using Pan-STARRS data (Laevens et al. 2015a, 2015b; Martin et al. 2015) while the remaining objects were discovered in the first and second years of the DES (Bechtol et al. 2015; Drlica-Wagner et al. 2015; Kim & Jerjen 2015; Kim et al. 2015a, 2015b; Koposov et al. 2015). For our purposes\u2014where we seek a homogeneous comparison of the Local Group and Virgo luminosity functions, based on NGVS data\u2014we may limit the analysis to Local Group satellites brighter than Mg = \u22129.13 mag, the 50% detection limit that applies to the Virgo sample. This reduces the Local Group sample to the galaxies listed in Table 5, divided according to whether they are M31 satellites (20), Galactic satellites (8), or members of the Local Group but not associated with either galaxy specifically (10), following the prescription of McConnachie (2012).27\n\n27\nA virial radius of 300 kpc is assumed for both the Milky Way and M31. The zero-velocity radius of the Local Group is taken to be 1060 kpc.\n For comparison with Virgo, we assume a color term g \u2212 V = 0.14 mag to convert the V-band magnitudes in Table 5 to g-band magnitudes, although, as we will show, the slope is not sensitive to the exact value used for the conversion (or, equivalently, to the value of \n\n\n\n\n\n).","Citation Text":["Martin et al. 2015"],"Citation Start End":[[944,962]]}
{"Identifier":"2019MNRAS.483.2362R__Wang,_Brinkmann_&_Bergeron_1996_Instance_1","Paragraph":"Various types of AGNs are known and one among them are the narrow-line Seyfert 1 (NLSy1) galaxies, which are classified based on the presence of narrow H\u03b2 emission line with full width at half-maximum (FWHM) 2000\u2009km s\u22121 and weak [O\u2009iii] emission line, with F([O\u2009iii])\/F(H\u03b2) 3 (Osterbrock & Pogge 1985; Goodrich 1989). They are believed to be powered by low-mass black holes (${\\sim } 10^7 \\, \\mathrm{M}_{\\odot }$) having higher accretion rate and generally showing strong Fe\u2009ii emission compared to their broad line counterparts namely the broad-line Seyfert 1 (BLSy1) galaxies (Xu et al. 2012; Rakshit et al. 2017a). However, from spectro-polarimetric observations of a \u03b3-ray emitting NLSy1 galaxy, PKS 2004 \u2212 447 (Baldi et al. 2016) and accretion disc modelling of a sample of 23 radio-detected NLSy1 galaxies (Calderone et al. 2013) indicate that they have masses similar to the blazar class of AGNs. Other characteristics that make NLSy1 galaxies different from the BLSy1 galaxies are their rapid soft X-ray variability (Pounds, Done & Osborne 1995; Leighly 1999a), steep soft X-ray spectra (Boller, Brandt & Fink 1996; Wang, Brinkmann & Bergeron 1996; Leighly 1999b) and low amplitude optical variability (Grupe 2004; Rakshit & Stalin 2017). Also, the fraction of NLSy1 galaxies detected in radio is much lower (${\\sim } 7{{\\ \\rm per\\ cent}}$) compared to the fraction of radio detected BLSy1 galaxies (Komossa et al. 2006; Rakshit et al. 2017a). Among radio-loud NLSy1 galaxies, about a dozen (\u223c2 per\u2009cent) have been detected in \u03b3-ray by the Fermi-Large Area Telescope (e.g. Abdo et al. 2009; D\u2019Ammando et al. 2015; Paliya et al. 2018) suggesting the unambiguous presence of relativistic jets in them. Multiband broad-band SED modelling of these \u03b3-ray detected NLSy1 galaxies indicate that these sources have many properties similar to the blazar class of AGNs (Paliya et al. 2013b) and specifically resembling the flat spectrum radio quasar (FSRQ) category (Paliya et al. 2018). In the radio, these \u03b3-ray emitting NLSy1 galaxies have a compact core jet morphology, high brightness temperature, show superluminal motion and significant radio variability (Doi et al. 2006; Komossa et al. 2006). Detailed investigations of the population of NLSy1 galaxies need to be undertaken to understand more about their peculiar characteristics.","Citation Text":["Wang, Brinkmann & Bergeron 1996"],"Citation Start End":[[1124,1155]]}
{"Identifier":"2022ApJ...937L..35M__Boylan-Kolchin_2017_Instance_1","Paragraph":"The most interesting interpretation of the clumps in and around the Sparkler is that the bulk of them are evolved (maximally old, given the 4.6 Gyr age of universe at the epoch of observation) globular clusters. This is also consistent with Sameie et al. (2022), who find that globular cluster formation in lower-mass galaxies occurs at redshifts of \u223c5\u201310. If this interpretation is correct, JWST observations of quiescent, evolved globular clusters around z \u223c 1.5 galaxies can be used to explore the formation history of globular clusters in a manner that is complementary to searching directly for the earliest stages of globular cluster formation (e.g., by examining young massive star formation complexes at z \u223c 6 and higher; Boylan-Kolchin 2017, 2018). Young star formation complexes may, or may not, evolve eventually into globular clusters, but there can be little doubt about the identity of an isolated and quiescent compact system if its mass is around 106\nM\n\u2299 and its scale length is a few parsecs. JWST observations of evolved globular clusters at z \u223c 1.5 are also complementary to exploring the ages of local globular clusters, as models fit to local globular clusters cannot distinguish between old and very old systems. For example, distinguishing between an \u223c11.5 Gyr old stellar population that formed at z = 3 and a 13.2 Gyr old stellar population that formed at z = 9 is not possible with current models and data, because they are degenerate with respect to a number of physical parameters (Ocvirk et al. 2006; Conroy et al. 2009, 2010). JWST observations of evolved globular clusters, seen when the universe was about one third of its present age, provide an opportunity for progress by \u201cmeeting in the middle\u201d, because population synthesis models of integrated starlight from SSPs can distinguish rather easily between the ages of young-intermediate stellar populations. This is because intermediate-mass stars with very distinctive photospheric properties are present at these ages. At z = 1.378, the lookback time to the Sparkler is 9.1 Gyr, and the age of the universe at that epoch is 4.6 Gyr. Distinguishing between z = 3 and z = 9 formation epochs for the globular cluster system corresponds to distinguishing between 2.4 Gyr and 4.1 Gyr old populations, which is relatively straightforward for population synthesis models in the JWST bands. In the case of the Sparkler, the striking conclusion is that at least foir of its globular clusters have likely formed at z > 9.","Citation Text":["Boylan-Kolchin 2017"],"Citation Start End":[[730,749]]}
{"Identifier":"2021AandA...651A..52B__Fritz_et_al._2006_Instance_1","Paragraph":"In this section we describe the method used to retrieve the various physical parameters associated with each LIR\u2005\u2212\u2005z bin of the master catalogue. For all templates, the total LIR is derived directly integrating the templates between 8 and 1000 \u03bcm. In order to derive the other main physical parameters such as stellar mass and colour excess for all simulated galaxies, we applied the Multi-wavelength Analysis of Galaxy Physical Properties (MAGPHYS, da Cunha et al. 2008) package, or the software SED3FIT (Berta et al. 2013) based on MAGPHYS when an AGN contribution is present to the set of templates associated with each galaxy population. This SED fitting procedure serves to infer the physical properties associated with each SED template because this information is not available a priori because of the empirical nature of the templates described in Sect. 2.1.1. To perform the fit, we created a large set of simulated observations of custom filters spanning from FUV to far-IR wavelengths and assigning S\/N of 10 to each of them. For the AGN component, we considered a library with smooth (Fritz et al. 2006; Feltre et al. 2012) and clumpy tori (Nenkova et al. 2008a,b), as explained in more detail in the next section. The analysis of the AGN component is necessary to disentangle the contribution of the AGN and star formation to the total IR-luminosity and to take into account the contamination of the AGN when deriving the stellar mass. In the fits by both MAGPHYS and SED3FIT we considered Bruzual & Charlot (2003) stellar templates with Chabrier (2003) IMF. For all considered empirical templates we retrieved the best fit value of the different physical parameters estimated by SED3FIT or MAGPHYS, as well as the 2.5%, 16%, 84%,\u2006 and 97.5% percentile when available3. When assigning a physical property to a simulated galaxy, we randomised each parameter value following the corresponding probability distribution to take into account the uncertainties on the SED fitting procedure. When a probability distribution is not available, we assumed a Gaussian distribution centred to the best value and an arbitrary \u03c3\u2004=\u20040.3 dex. In the master catalogue we considered only the median value of each parameter, as derived for each galaxy population and LIR\u2005\u2212\u2005z bin.","Citation Text":["Fritz et al. 2006"],"Citation Start End":[[1097,1114]]}
{"Identifier":"2021MNRAS.500.5009M__Bono_et_al._1997_Instance_1","Paragraph":"An extensive set of non-linear convective pulsation models for RR Lyrae at different metal and helium abundances has been taken into account. The transformation of bolometric magnitude variations into the Gaia filters allowed us to derive the first theoretical light curves directly comparable with Gaia time-series data. In particular, we built the first theoretical Bailey diagrams and PW relations in the GBP, G, and GRP filters, varying both the metallicity and the helium content. As for the Bailey diagram, we conclude that an increase in the metal abundance and\/or in the helium abundance produces a decrease in the pulsation amplitudes, whereas an increase in the luminosity level produces a period shift towards longer values. In particular, in the case of FO RR Lyrae, the location of the described bell-shape in the Bailey diagram can be used to constrain the luminosity level (see e.g. Bono et al. 1997). The theoretical PW relations in the Gaia bands show a dependence of the zero-point on metal abundance, in the sense that brighter Wesenheit functions are predicted for more metal poor chemical composition and a lower effect due to variations in the helium content, with helium-enriched models characterized by longer periods and brighter Wesenheit functions. The theoretical PWZ relations are applied to a subset of Gaia DR2 RR Lyrae (293 F and 50 FO pulsators) with complementary metallicity information to infer individual theoretical parallaxes, which are in very good agreement with Gaia results. In particular, the inferred zero-point parallax offset is consistent with zero both in the case of F and FO pulsators. Even if more stringent conclusions could be drawn in the future from the next Gaia data releases, the obtained results seem on one side to support the predictive capabilities of current pulsation models and on the other to suggest that a smaller parallax offset could be required for the bluer older and lower mass RR Lyrae stars than for classical Cepheids.","Citation Text":["Bono et al. 1997"],"Citation Start End":[[898,914]]}
{"Identifier":"2019AandA...629A..54U__Evans_et_al._2007_Instance_1","Paragraph":"NGC 2110. NGC 2110 is another nearby (z\u2004=\u20040.00779, Gallimore et al. 1999), X-ray bright Seyfert galaxy. Diniz et al. (2015) report a black hole mass of \n\n\n\n2\n.\n\n7\n\n\u2212\n2.1\n\n\n+\n3.5\n\n\n\u00d7\n\n10\n8\n\n\n\nM\n\u2299\n\n\n\n$ 2.7^{+ 3.5}_{- 2.1} \\times 10^{8}\\,{{M}_{\\odot}} $\n\n\n, from the relation with the stellar velocity dispersion. From BeppoSAX data, Malaguti et al. (1999) found the X-ray spectrum to be affected by complex absorption. This has been later confirmed by Evans et al. (2007), who find the Chandra+XMM\u2013Newton data to be well fitted with a neutral, three-zone, partial-covering absorber. Rivers et al. (2014) find the Suzaku data to be well fitted with a stable full-covering absorber plus a variable partial-covering absorber. A soft excess below 1.5 keV is also present (Evans et al. 2007), and possibly due to extended circumnuclear emission seen with Chandra (Evans et al. 2006). No Compton reflection hump has been detected with Suzaku (Rivers et al. 2014) or NuSTAR (Marinucci et al. 2015), despite the presence of a complex Fe K\u03b1 line. According to the multi-epoch analysis of Marinucci et al. (2015), the Fe K\u03b1 line is likely the sum of a constant component (from distant, Compton-thick material) and a variable one (from Compton-thin material). Concerning the high-energy cut-off, ambiguous results have been reported in literature (see Table 1). Ricci et al. (2017) report a value of \n\n\n\n448\n\n\u2212\n55\n\n\n+\n63\n\n\n\n$ 448^{+63}_{-55} $\n\n\n keV, while Lubi\u0144ski et al. (2016) report a coronal temperature of \n\n\n\n230\n\n\u2212\n57\n\n\n+\n51\n\n\n\n$ 230^{+51}_{-57} $\n\n\n keV and an optical depth of \n\n\n\n0\n.\n\n52\n\n\u2212\n0.13\n\n\n+\n0.14\n\n\n\n\n$ 0.52^{+ 0.14}_{- 0.13} $\n\n\n. From 2008\u20132009 INTEGRAL data, Beckmann & Do Cao (2010) report a cut-off of \u223c80 keV with a hard photon index, but these results are not confirmed by NuSTAR (Marinucci et al. 2015). Indeed, only lower limits to the high-energy cut-off have been found with NuSTAR (210 keV: Marinucci et al. 2015), Suzaku (250 keV: Rivers et al. 2014) and BeppoSAX (143 keV: Risaliti 2002). No hard X-ray spectral variability has been detected by Caballero-Garcia et al. (2012) and Soldi et al. (2014) from BAT data, despite the significant flux variability.","Citation Text":["Evans et al. (2007)"],"Citation Start End":[[450,469]]}