{"Identifier":"2021MNRAS.501.4035R__Vigren_&_Galand_2013_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":["Vigren & Galand 2013"],"Functions Text":["Several modelling works on this comet have shown that ion composition in the coma varies based on the sublimation rate of the nucleus"],"Functions Label":["Background"],"Citation Start End":[[1242,1262]],"Functions Start End":[[1107,1240]]}
{"Identifier":"2022MNRAS.512.2222V__Granato_et_al._2004_Instance_1","Paragraph":"As for the ETGs, which are spheroid-dominated, most of these that are more massive than \u223c2 \u00d7 1010 M\u2299 are old, implying that the SF has strongly declined long time ago and no recent bursts of SF have occurred. Early formation and rapid quenching mechanisms are expected for those galaxies, where the quenching mechanisms are likely to be associated with their morphological transformation through wet major mergers and disc instabilities. These processes lead to compaction and strong bursts of SF that consume the gas (e.g. Hopkins et al. 2008; Barro et al. 2013; Dekel & Burkert 2014), as well as strong AGN\/QSO and Supernova feedback that will heat and expel the gas (e.g. Granato et al. 2004; Sijacki et al. 2007; Somerville et al. 2008; Vogelsberger et al. 2014). At intermediate and low masses, the environment-driven quenching mechanisms (e.g. ram-pressure and tidal striping, strangulation, etc.) can also be relevant (e.g. Kauffmann et al. 2004; Peng et al. 2010; Schawinski et al. 2014). Note that these mechanisms are expected to lead also to morphological transformation (e.g. Gunn & Gott 1972; Moore et al. 1996; Abadi, Moore & Bower 1999; Bekki, Couch & Shioya 2002; Arag\u00f3n-Salamanca, Bedregal & Merrifield 2006). However, as M* is lower there is an increasing fraction of ETGs with intermediate values of Agelw, or even low values (Table 3), though not as low as in the case of low-mass LTGs (see Fig. 16). This population of ETGs is partially associated to the Blue Star-forming and Recently Quenched Early-type galaxies, for which rejuvenation processes (Thomas et al. 2010) by gas infall or late gas-rich mergers were proposed in Lacerna et al. (2016, 2020). In fact, the $Age_{\\rm mw}\/{Age_{\\rm lw}}$ ratio of these galaxies is large, which implies that they did not form so late but they contain some (small) fractions of very young stellar populations. Finally, for the ITGs, the Agelw distribution in the \u223c0.5\u20135 \u00d7 1010 M\u2299 mass range is roughly bimodal (Fig. 16), with a large fraction of them in the intermediate age region (or green valley in the colour\u2013M* diagram), showing that the quenching time-scales for these galaxies with a significant bulge are slow (quenching mechanisms like those driven by morphology, halo mass, and environment; see above). For the most massive S0\u2013Sa galaxies, ${Age_{\\rm lw}}\\gt 4$ Gyr, that is, they are mostly quenched, while for the lest massive, ${Age_{\\rm lw}}\\lesssim 2$ Gyr, which imply some rejuvenation processes.","Citation Text":["Granato et al. 2004"],"Functions Text":["as well as strong AGN\/QSO and Supernova feedback that will heat and expel the gas (e.g."],"Functions Label":["Background"],"Citation Start End":[[675,694]],"Functions Start End":[[587,674]]}
{"Identifier":"2016AandA...595A..72M__Vergani_et_al._2015_Instance_1","Paragraph":"On the other hand, the Australia Telescope Compact Array (ATCA) 21\u2009cm line survey of GRB host galaxies revealed high levels of atomic hydrogen (H\u2009i), suggesting that the connection between atomic gas and star formation is stronger than previously thought (Micha\u0142owski et al. 2015). Star formation may be directly fuelled by atomic gas, as has been theoretically shown to be possible (Glover & Clark 2012; Krumholz 2012; Hu et al. 2016), and this is supported by the existence of H\u2009i-dominated, star-forming regions in other galaxies (Bigiel et al. 2008, 2010; Fumagalli & Gavazzi 2008; Elmegreen et al. 2016). This can happen in a low metallicity gas that is recently acquired, even if the metallicity in other parts of a galaxy is higher, near the onset of star formation because cooling of gas (necessary for star formation) is faster than the H\u2009i-to-H2 conversion (Krumholz 2012). Indeed, large atomic gas reservoirs, together with low molecular gas masses (Hatsukade et al. 2014; Stanway et al. 2015b) and stellar masses (Perley et al. 2013, 2015; Vergani et al. 2015), indicate that GRB hosts are preferentially galaxies that have very recently started a star formation episode. This provides a natural route for forming GRBs in low-metallicity environments, as found for most GRB hosts (Fruchter et al. 2006; Modjaz et al. 2008; Levesque et al. 2010a; Han et al. 2010; Boissier et al. 2013; Schulze et al. 2015; Vergani et al. 2015; Japelj et al. 2016; Perley et al. 2016), except of a few examples of hosts with solar or super-solar metallicities (Prochaska et al. 2009; Levesque et al. 2010b; Kr\u00fchler et al. 2012; Savaglio et al. 2012; Elliott et al. 2013; Schulze et al. 2014; Hashimoto et al. 2015; Schady et al. 2015; Stanway et al. 2015a). Indeed, the GRB collapsar model requires that most of the GRB progenitors have low metallicity (below solar) in order to reduce the loss of mass and angular momentum that is required for launching the jet (Yoon & Langer 2005; Yoon et al. 2006; Woosley & Heger 2006). We note however that other models, while still predicting the metallicity preference (e.g. Izzard et al. 2004; Podsiadlowski et al. 2004; Detmers et al. 2008), allow higher metallicities owing to differential rotation (Georgy et al. 2012), binary evolution (Podsiadlowski et al. 2010; van den Heuvel & Portegies Zwart 2013), or weaker magnetic fields (Petrovic et al. 2005). ","Citation Text":["Vergani et al. 2015"],"Functions Text":["Indeed, large atomic gas reservoirs, together with low molecular gas masses","and stellar masses","indicate that GRB hosts are preferentially galaxies that have very recently started a star formation episode."],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1052,1071]],"Functions Start End":[[884,959],[1006,1024],[1074,1183]]}
{"Identifier":"2015AandA...582A..22L__Todorov_et_al._2014_Instance_2","Paragraph":"Dust distribution:We employed the standard flared disk model with well-mixed gas and dust, which has been successfully used to explain the observed SEDs of a large sample of young stellar objects and BDs (e.g., Wolf et al. 2003; Sauter et al. 2009; Harvey et al. 2012a; Joergens et al. 2013; Liu et al. 2015). The structure of the dust density is assumed with a Gaussian vertical profile (1)\\begin{equation} \\rho_{\\rm{dust}}=\\rho_{0}\\left(\\frac{R_{*}}{\\varpi}\\right)^{\\alpha}\\exp\\left[-\\frac{1}{2}\\left(\\frac{z}{h(\\varpi)}\\right)^2\\right], \\label{dust_density} \\end{equation}\u03c1dust=\u03c10R\u2217\u03d6\u03b1exp\u221212zh(\u03d6)2,and the surface density is described as a power-law function (2)\\begin{equation} \\Sigma(\\varpi)=\\Sigma_{0}\\left(\\frac{R_{*}}{\\varpi}\\right)^p, \\end{equation}\u03a3(\u03d6)=\u03a30R\u2217\u03d6p,where \u03d6 is the radial distance from the central star measured in the disk midplane, and h(\u03d6) is the scale height of the disk. The disk extends from an inner radius Rin to an outer radius Rout. To the best of our knowledge, among our sample, there are five objects that have been identified as binary systems so far. They are 2M1207 (a~55 AU, Chauvin et al. 2004), J04221332+1934392 (a~7 AU, Todorov et al. 2014), J04414489+2301513 (a~15 AU, Todorov et al. 2014), USD161833 (a~134 AU, Bouy et al. 2006), and USD161939 (a ~ 26 AU, Bouy et al. 2006), where a refers to the separation within the system. The disks around individual components in binary systems are expected to have truncation radii of the order of (0.3 \u2212 0.5)a (Papaloizou & Pringle 1977). We adopted 0.5 a as the disk outer radii for 2M1207, USD161833, and USD161939. For the close pairs (a \u2272 15 AU, J04221332+1934392 and J04414489+2301513), dynamical simulations of star-disk interactions suggest that individual disks are unlikely to survive (e.g., Artymowicz & Lubow 1994). Disk modeling is complicated in those close multiple systems. For simplicity, we assume that the emission is associated with circumbinary disks of 100 AU in size. For other objects, we fix Rout = 100 AU in the modeling because the choice of this parameter value makes essentially no difference to the synthetic SEDs in the simulated wavelength range (Harvey et al. 2012a). The scale height follows the power-law distribution(3)\\begin{equation} h(\\varpi) = H_{100}\\left(\\frac{\\varpi}{100\\,\\rm{AU}}\\right)^\\beta,\\\\ \\end{equation}h(\u03d6)=H100\u03d6100\u2009AU\u03b2,with the exponent \u03b2 characterizing the degree of flaring and H100 representing the scale height at a distance of 100 AU from the central star. The indices \u03b1, p, and \u03b2 are codependent through p = \u03b1 \u2212 \u03b2. We fix p = 1, which is the typical value found for T Tauri disks in the sub-millimeter (e.g., Isella et al. 2009; Guilloteau et al. 2011), since only spatially resolved data can place constraints on this parameter (e.g., Ricci et al. 2013, 2014).               Dust properties:We assume the dust grains to be a homogeneous mixture of 75% amorphous silicate and 25% carbon with a mean density of \u03c1grain = 2.5 g cm-3 and the complex refractive indices given by J\u00e4ger et al. (1994, 1998), and Dorschner et al. (1995). Porous grains are not considered because the fluxes at wavelengths beyond ~ 2 \u03bcm are almost independent of the degree of grain porosity in low-mass disks, as shown by Kirchschlager & Wolf (2014). The grain size distribution is given by the standard power law dn(a) \u221d a-3.5da with minimum and maximum grain sizes amin = 0.1 \u03bcm and amax = 100 \u03bcm, respectively. The choice of the minimum value for the grain size, amin, ensures that its exact value has a negligible impact on the synthetic SEDs. Since there is no information about the maximum grain sizes of our target disks, as provided, e.g., by the (sub)millimeter spectral index, we adopt amax = 100 \u03bcm. The Herschel\/PACS far-IR observations are sensitive to dust grains with this assumed sizes. Strong grain growth up to millimeter sizes as detected in some BD disks (e.g., Ricci et al. 2012, 2013, 2014; Broekhoven-Fiene et al. 2014) would remain undetected in our data and could affect the disk mass. Our prescription for the dust properties is identical to those used in Liu et al. (2015). ","Citation Text":["Todorov et al. 2014"],"Functions Text":["J04414489+2301513 (a~15 AU,"],"Functions Label":["Uses"],"Citation Start End":[[1210,1229]],"Functions Start End":[[1182,1209]]}
{"Identifier":"2019ApJ...882..144K__Berk_et_al._2001_Instance_2","Paragraph":"The FOCAS and NIRSPEC spectra of PSO J006+39 were obtained at two different epochs separated by 1 yr and 9 months (by slightly less than 3 months in the quasar rest frame). Previously, we found that the PS1 y-band light curve of PSO J006+39 shows brightness variations with a peak-to-peak amplitude of \u223c0.7 mag over \u223c4 yr (Koptelova et al. 2017), which might be due to the flux variations of both continuum and Ly\u03b1 line of PSO J006+39. To infer the brightness state of PSO J006+39 at the epochs of its FOCAS and NIRSPEC observations, we first calculated the spectral slope of the quasar continuum from the NIRSPEC spectrum with a wider wavelength coverage than that of the FOCAS spectrum. Using wavelength intervals of 11100\u201311300, 11400\u201311600, 13085\u201313400, and 14700\u201315200 \u212b we measured a spectral slope of \u03b1\u03bb = \u22121.35 \u00b1 0.26, where the quoted uncertainty is the statistical error of the fit. The fitted power law is shown in Figure 2 with a solid line. The estimated continuum slope is consistent but somewhat flatter than the typical slope of luminous quasars (Zheng & Malkan 1993; Vanden Berk et al. 2001; Selsing et al. 2016). We then fitted the FOCAS data using the power law with a fixed spectral slope of \u03b1\u03bb = \u22121.35 and spectral windows of 9700\u20139850 and 10050\u201310100 \u212b. The spectral windows adopted for the analysis of the FOCAS and NIRSPEC spectra were taken to be similar to the rest-frame wavelength intervals commonly used to fit the continua of quasars (Vanden Berk et al. 2001; Decarli et al. 2010; Lusso et al. 2015) and less affected by the contribution from emission lines on the red side of the Big Blue Bump (BBB; e.g., Malkan 1983). The estimated continuum flux of PSO J006+39 at the epoch of its FOCAS observations is shown in Figure 2 with a dashed line. By comparing the continuum flux at the epochs of the FOCAS and NIRSPEC observations, we find that the brightness state of PSO J006+39 was different at these two epochs. PSO J006+39 was brighter by about 0.8 mag during the FOCAS observations than during the NIRSPEC observations. Thus, the continuum flux of PSO J006+39 might be different at different epochs depending on the brightness state of the quasar. Figure 2 also shows the fluxes of PSO J006+39 in the FOCAS Y, and NIRSPEC N2, N4, and N6 bands at the epochs of the FOCAS and NIRSPEC observations (see also Table 1).","Citation Text":["Vanden Berk et al. 2001"],"Functions Text":["The spectral windows adopted for the analysis of the FOCAS and NIRSPEC spectra were taken to be similar to the rest-frame wavelength intervals commonly used to fit the continua of quasars"],"Functions Label":["Uses"],"Citation Start End":[[1465,1488]],"Functions Start End":[[1276,1463]]}
{"Identifier":"2017ApJ...835..246Y__Yoon_&_Seough_2014_Instance_1","Paragraph":"The present analysis builds upon the macroscopic-kinetic model of the solar wind, originally formulated by Yoon & Seough (2014) for the proton temperatures. The same model was recently generalized to include collisional dissipation (Yoon 2016a, 2016b). The basic methodology is similar to that of Marsch & Tu (2001) and Jasperse et al. (2006), especially in regards to treating the particle aspect, but unlike earlier works (Marsch & Tu 2001; Jasperse et al. 2006), which do not treat the waves self-consistently, Yoon & Seough (2014) and Yoon (2016a, 2016b) discuss the wave generation in a self-consistent manner by solving the adiabatic dispersion relation and wave kinetic equation for each spatial location. We now extend the original formalism (Yoon & Seough 2014) in another direction. We include dynamic electrons, but unlike the two later works (Yoon 2016a, 2016b), we ignore collisional dissipation. For an inhomogeneous plasma immersed in a diverging or converging magnetic field, the kinetic equation for the particles subject to perturbations propagating in parallel direction, s, is given by\n1\n\n\n\n\n\nwhere, for cylindrical velocity coordinate system, the velocity diffusion coefficient tensor is given by\n2\n\n\n\n\n\nand where \n\n\n\n\n\n is the complex angular frequency, which must be determined by the local dispersion relation,\n3\n\n\n\n\n\nIn the above relation, ea and ma are unit electric charge and mass for particles species a (a = p for protons and a = e for electrons, \n\n\n\n\n\n for protons and \n\n\n\n\n\n for electrons); \n\n\n\n\n\n stands for cyclotron frequency for species a, B0 and c being the ambient magnetic field intensity and the speed of light in vacuo; \n\n\n\n\n\n is the square of the plasma frequency defined for species a, n0 being the ambient plasma density, and \n\n\n\n\n\n being the perturbed electric field associated with the unstable transverse mode propagating parallel (and anti-parallel) to the ambient magnetic field, \u00b1 denoting the right\/left-hand circular polarization. The spectral electric field wave energy density \n\n\n\n\n\n must be determined by solving the wave kinetic equation, the simplest form of which is given by the quasilinear theory,\n4\n\n\n\n\n\n\n","Citation Text":["Yoon & Seough (2014)","Yoon & Seough 2014"],"Functions Text":["and Yoon (2016a, 2016b) discuss the wave generation in a self-consistent manner by solving the adiabatic dispersion relation and wave kinetic equation for each spatial location.","We now extend the original formalism","in another direction."],"Functions Label":["Motivation","Extends","Extends"],"Citation Start End":[[514,534],[751,769]],"Functions Start End":[[535,712],[713,749],[771,792]]}
{"Identifier":"2021MNRAS.503.1319G__Chae_&_Mao_2003_Instance_2","Paragraph":"In the first scenario, we assume that neither the characteristic velocity dispersion (\u03c3*) nor the number density (n*) of galaxies evolves with redshifts (\u03bdn = \u03bdv = 0). Given the redshift coverage of the lensing galaxies in the lens sample (0.06  zl  1.0), if we constrain a non-evolving VDF using the lens data, then, assuming the VDF evolution with redshift is smooth, the fits on the VDF parameters may represent the properties of ETGs at an effective epoch of z \u223c 0.5. Such non-evolving VDF has been extensively applied in the previous studies on lensing statistics (Chae & Mao 2003; Ofek et al. 2003; Capelo & Natarajan 2007; Cao et al. 2012a). By applying the above-mentioned \u03c72 \u2013 minimization procedure to Sample A \u2013 we obtain the best-fitting values and corresponding 1\u03c3 uncertainties (68.3 per cent confidence level): $\\alpha =0.66^{+2.13}_{-0.66}$, $\\beta =2.28^{+0.24}_{-0.18}$. It is obvious that the full sample analysis has yielded improved constraints on the high-velocity exponential cut-off index \u03b2, compared with the previous analysis of using the distribution of image separations observed in CLASS and PANELS to constrain a model VDF of ETGs (Chae 2005). Suffering from the limited size of lens sample, such analysis (Chae 2005) found that neither of the two VDF parameters (\u03b1, \u03b2) can be tightly constrained, due to the broad regions in the \u03b1 \u2212 \u03b2 plane. Consequently, the image separation distribution is consistent with the SDSS measured stellar VDF (Sheth et al. 2003) and the Second Southern Sky Redshift Survey (SSRS2) inferred stellar VDF (Chae & Mao 2003), although the two stellar VDFs are significantly different from each other concerning their corresponding parameter values. We also consider constraints obtained for the Sample B (defined in previous section), with the likelihood is maximized at $\\alpha =1.00^{+2.38}_{-1.00}$ and $\\beta =2.34^{+0.26}_{-0.24}$, from which one could clearly see the marginal consistency between our fits and recent measurements of three stellar VDFs (especially the SDSS DR5 VDF of ETGs).","Citation Text":["Chae & Mao 2003"],"Functions Text":["Consequently, the image separation distribution is consistent with the SDSS measured stellar VDF","and the Second Southern Sky Redshift Survey (SSRS2) inferred stellar VDF","although the two stellar VDFs are significantly different from each other concerning their corresponding parameter values."],"Functions Label":["Similarities","Similarities","Compare\/Contrast"],"Citation Start End":[[1564,1579]],"Functions Start End":[[1373,1469],[1490,1562],[1582,1704]]}
{"Identifier":"2020AandA...639A.107S__Helling_et_al._2019a_Instance_1","Paragraph":"Micro-porosity is the porosity arising from the organisation of the condensate monomers (e.g. Mg2 SiO4 in Mg2 SiO4[s]) within a cloud particle during growth. This is different from the porosity that can be used to characterise aggregates that originate from particle-particle collision processes (coagulation, e.g. Dominik & Tielens 1997; Blum & Wurm 2000), which we do not considerhere. On Earth, the material density of water ice is dependent on the ambient temperature at formation. Snowflakes are known to form many types of crystal structures that can be up to 84% porous for millimetre-sized cloud particles when compared toice material density (Hales 2005), leading to the possibility of altitude-dependent porosity in terrestrial snow clouds. Earth-like exoplanets, mini-Neptunes, and T-type brown dwarfs may form water clouds, composed of liquid or solid particles, but warmer planets and brown dwarfs of L-type and later have been shown to form cloud particles made of a mix of materials that is dominated by Mg, Si, Fe, and O and to a lesser extent by Ti, Al, K and other elements (e.g. Witte et al. 2009; Lee et al. 2015; Helling et al. 2019a). There are many ways in which this micro-porosity might be incorporated into mineral cloud particles, for example lattice faults at the interfaces between two different condensation species owing to the different lattice structures. Even for homogeneous growth, single species often have multiple crystal structures (Sood & Gouma 2013), which can also generate lattice faults at their interfaces. For example, the TiO2 [s] rutile and anatase forms are both stable at atmospheric pressures for temperatures greater than 1100 K (Jung & Imaishi 2001; Hanaor & Sorrell 2010). Additionally, within crystal structures, there are many known types of defect that might further decrease material density (e.g. Schottky defects in TiO2 [s] and MgO[s] crystals M\u00e9n\u00e9trey et al. 2004). Furthermore, these cloud particles not only change their material composition when falling through the atmosphere, but their particle sizes will also change such that the largest cloud particles are forming the innermost part of the cloud, which often sits deep inside the optically thick part of the atmosphere. Because cloud particles made of a mix of many thermally stable materials fall into warmer atmospheric regions, the low-temperature materials (such as SiO[s], MgSiO3[s]) become thermally unstable, they evaporate and leave behind a skeleton made of high-temperature materials (such as Fe[s], TiO2[s], Al2O3[s]). Whilst this may be a source of micro-porosity of cloud particles, Juncher et al. (2017) noted that this may also lead to a reduction in micro-porosity because the structural integrity of the particle is weakened and dangling structures break off. These micro-porous mineral cloud particles we call \u201cmineral snowflakes\u201d.","Citation Text":["Helling et al. 2019a"],"Functions Text":["Earth-like exoplanets, mini-Neptunes, and T-type brown dwarfs may form water clouds, composed of liquid or solid particles, but warmer planets and brown dwarfs of L-type and later have been shown to form cloud particles made of a mix of materials that is dominated by Mg, Si, Fe, and O and to a lesser extent by Ti, Al, K and other elements"],"Functions Label":["Background"],"Citation Start End":[[1134,1154]],"Functions Start End":[[751,1091]]}
{"Identifier":"2018ApJ...852...45W__Ghisellini_et_al._2013_Instance_1","Paragraph":"Similarly to former works (e.g., Zhang et al. 2010, 2012; Kang et al. 2014, 2016; Yan et al. 2016), we neglect the low-frequency radio data and consider the data with \n\n\n\n\n\n (or \n\n\n\n\n\n) in our SED fitting with the one-zone model due to the fact that radio emission should come from the large-scale jet and cannot be accounted for with a one-zone model. The variability correlation between millimeter, optical, X-ray, and \u03b3-ray emission support the fact that they come from more or less the similar region (e.g., Sikora et al. 2008; Le\u00f3n-Tavares et al. 2012; Wehrle et al. 2012; D\u2019Ammando et al. 2013; Orienti et al. 2013). In four LSP blazars (0333 + 321, 0430 + 052, 2145 + 067, 2230 + 114), the putative UV excesses are not included in our SED fitting, as they should come from the cold accretion disk (Shakura & Sunyaev 1973; Ghisellini et al. 2013; Ajello et al. 2016). On average, there are 17 data points in our fitting. As an example, we show the multi-wavelength SED and its fitting for 2200 + 420 (BL Lac) in Figure 1 (left panel), where only the SSC process is considered due to the EC not being important in this source. In the right panel of Figure 1, we show the probability distributions of the model parameters, where \n\n\n\n\n\n G, \n\n\n\n\n\n, \n\n\n\n\n\n, \n\n\n\n\n\n, \n\n\n\n\n\n, \n\n\n\n\n\n, and \n\n\n\n\n\n (upper and lower limits represent 1\u03c3 errors, see also Table 1). The source of 1219+285 is also a BL Lac object, for which the EC component is negligible. For the remaining 23 LSP sources, both the SSC and EC component are considered, where the SEDs and the fitting are shown in Figures 2\u20136. For each source, the SED fitting with seed photons from the torus and BLR are presented in left and right panels, respectively. In our SED fitting, we find that most LSP blazars have \n\n\n\n\n\n (20 of 23 sources, \n\n\n\n\n\n for the two other sources and only one has the ratio \u223c0.6, see Table 1), where the distribution of the ratio \n\n\n\n\n\n is shown in Figure 7. Our results suggest that the SED fitting with the seed photons from torus are better than with those from the BLR, which supports the notion that the location of the \u03b3-ray emitting region should stay outside of the BLR. In the following work, therefore, we consider the model parameters from the SED fittings with the torus seed photons.","Citation Text":["Ghisellini et al. 2013"],"Functions Text":["In four LSP blazars (0333 + 321, 0430 + 052, 2145 + 067, 2230 + 114), the putative UV excesses are not included in our SED fitting, as they should come from the cold accretion disk"],"Functions Label":["Uses"],"Citation Start End":[[829,851]],"Functions Start End":[[623,803]]}
{"Identifier":"2015AandA...575A.111D__Drimmel_&_Spergel_2001_Instance_1","Paragraph":"Using our stellar parameters, we derived an estimate of the spectroscopic distances of XO-2N and XO-2S by means of the following procedure. We generated Monte Carlo (MC) normal distributions for each spectroscopic parameter Teff, [Fe\/H], and log\u2009 g, composed of 10\u2009000 random values and centred on the best estimates (Table 2). By keeping the stellar radii fixed to the values listed in Table 2 (for XO-2N we used the most accurate estimate), for each MC simulation we first determined the stellar bolometric luminosity L\u2217 (in solar units) from the Stefan-Boltzmann law, and then we derived the absolute bolometric magnitude Mbol from the relation Mbol = 4.75 \u2212 2.5\u00b7log\u2009(L\u2217). By estimating the appropriate bolometric correction (BC), a value for the absolute magnitude in V-band MV was then obtained. The BC term was evaluated using the code provided by Casagrande & VandenBerg (2014). An additional input is the colour excess E(B \u2212 V) of the star, which we derived through the relation E(B \u2212 V) = AV(s)\/3.1, where AV(s) is the interstellar dust extinction in V-band integrated at the distance s of the star (in pc) and measured along the line of sight. We derived AV(s) by adopting a simplified model of the local distribution of the interstellar dust density (Drimmel & Spergel 2001), expressed by the relation \u03c1 = \u03c10\u00b7sech2(z\/hs), where z is the height of the star above the Galactic plane and hs is the scale-height of the dust, for which we adopted the value of 190 pc. The term z is related to the distance s and the Galactic latitude of the star b by the formula z = s\u00b7sinb. From this model we obtained the relation AV(s) = AV(tot)\u00b7sinh(s\u00b7sin b\/hs)\/cosh (s\u00b7sin b\/hs), where AV(tot) is the interstellar extinction in V band along the line of sight integrated through the Galaxy, and can be estimated from 2D Galactic maps. For this purpose we used the value AV (tot) = 0.16 mag derived from the maps of Schlafly & Finkbeiner (2011)6. By assuming s = 150 pc as a prior distance of the stars (Burke et al. 2007), we obtained AV (150 pc) ~ 0.06 mag, corresponding to E(B \u2212 V) = 0.019 mag. This is the value used as input to the code of Casagrande & VandenBerg (2014) to obtain a first guess of the BC in V-band. This in turn was used in the distance modulus formula V \u2212 (Mbol \u2212 BCV) = 5\u00b7log\u2009(s) \u22125 \u2212 AV(s), to obtain a new value for the stellar distance s. The new distance was used to repeat the procedure iteratively, by determining at each step a new value of AV(s) and BCV(s), and finally another estimate of s. When the absolute difference between the last and previous calculated values of s was below 0.1 pc, the iterative process was interrupted and the last derived value for s was assumed as the distance of the star for the Nth Monte Carlo simulation. The adopted estimates for the distance of the XO-2 components are the median of the distributions of the 10\u2009000 MC values, and the asymmetric error bars defined as the 15.85th and 68.3th percentile (see Table 1). Because they are model-dependent, we do not argue here whether the difference of ~2.5 pc between the XO-2S and XO-2S distances is real. We only note that the two values are compatible within the uncertainties and that our best estimate for the distance of XO-2N locates the star a few parsec closer than reported by Burke et al. (2007). ","Citation Text":["Drimmel & Spergel 2001"],"Functions Text":["We derived AV(s) by adopting a simplified model of the local distribution of the interstellar dust density","expressed by the relation \u03c1 = \u03c10\u00b7sech2(z\/hs), where z is the height of the star above the Galactic plane and hs is the scale-height of the dust, for which we adopted the value of 190 pc."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1262,1284]],"Functions Start End":[[1154,1260],[1287,1473]]}
{"Identifier":"2019ApJ...875...90L__Pontieu_et_al._2011_Instance_1","Paragraph":"When energy flows from the interior of the Sun outward into the solar atmosphere, why is the Sun\u2019s outer atmosphere, the corona, much hotter than the inner atmosphere, the underlying chromosphere and photosphere? This is the long-standing problem of the coronal heating, which is one of the eight key mysteries in modern astronomy (Kerr 2012). For about 80 yr since the discovery of the extremely hot corona around the late 1930s (Grotian 1939; Edlen 1945), people have worked hard on addressing this issue, and great advances have been made in observation and theoretical studies (Parnell & De Moortel 2012; Amari et al. 2015; Arregui 2015; Cargill et al. 2015; De Moortel & Browning 2015; Jess et al. 2015; Klimchuk 2015; Longcope & Tarr 2015; Peter 2015; Schmelz & Winebarger 2015; Velli et al. 2015; Wilmot-Smith 2015). Especially during recent decades, high-resolution observations of solar super-fine structures indicate that small spicules, minor hot jets along small-scale magnetic channels from the low atmosphere upwards to the corona, petty tornados and cyclones, and small explosive phenomena such as mini-filament eruptions and micro- and nano-flares\u2014all of these small-scale magnetic activities contribute greatly to coronal heating (De Pontieu et al. 2011; 2018; Zhang & Liu 2011; Parnell & De Moortel 2012; Klimchuk 2015; Peter 2015; Schmelz & Winebarger 2015; Henriques et al. 2016; Li et al. 2018a). Additionally, contributions of MHD waves to heating the corona have been observationally illustrated (van Ballegooijen et al. 2011; Jess et al. 2015; Kubo et al. 2016; Morton et al. 2016; Soler et al. 2017; Morgan & Hutton 2018). Meanwhile, with the progress of observational studies, two groups of theoretical models, magnetic reconnection models and magnetohydrodynamic wave models, have traditionally attempted to explain coronal heating, but so far no models can address the key mystery perfectly (van Ballegooijen et al. 2011; Arregui 2015; Cargill et al. 2015; Peter 2015; Velli et al. 2015; Wilmot-Smith 2015). Maybe we do not need to intentionally take to heart such the classical dichotomy, because waves and reconnections may interact with each other (De Moortel & Browning 2015; Velli et al. 2015). Additionally, statistical studies may look at coronal heating from a comprehensive perspective. Li et al. (2018b) found that the long-term variation of the heated corona, which is represented by coronal spectral irradiances, and that of small-scale magnetic activity are in lockstep, indicating that the corona should statistically be effectively heated by small-scale magnetic activity. Observational and theoretical model studies through heating channels and modes, and statistical studies by means of heating effect (performance of the heated corona), both suggest that coronal heating originates from small-scale magnetic activity.","Citation Text":["De Pontieu et al. 2011"],"Functions Text":["Especially during recent decades, high-resolution observations of solar super-fine structures indicate that small spicules, minor hot jets along small-scale magnetic channels from the low atmosphere upwards to the corona, petty tornados and cyclones, and small explosive phenomena such as mini-filament eruptions and micro- and nano-flares\u2014all of these small-scale magnetic activities contribute greatly to coronal heating"],"Functions Label":["Background"],"Citation Start End":[[1248,1270]],"Functions Start End":[[824,1246]]}
{"Identifier":"2016ApJ...825..150C__Berm\u00fadez_et_al._2013_Instance_1","Paragraph":"The rotational spectrum of NaCl has been obtained using two different FTMW spectrometers constructed at the University of Valladolid. A solid rod was prepared by pressing the NaCl fine powder mixed with a small amount of commercial binder and was placed in the ablation nozzle (Alonso et al. 2009; Mata et al. 2012). A picosecond Nd:YAG laser (12 mJ per pulse, 20 ps pulse width) was used as a vaporization tool. Then, NaCl neutral molecules were supersonically expanded using the flow of a carrier gas (Ne at backing pressure of 15 bars) into the spectrometer chamber. NaCl was first investigated using a chirped-pulse Fourier transform microwave (CP-FTMW) spectrometer with a laser ablation source (Mata et al. 2012; Berm\u00fadez et al. 2013) operating between 6.0 and 12.0 GHz to sample swiftly the rotational spectra of the different species present in the supersonic expansion. Chirped pulses of 4 \u03bcs directly generated by the 24 Gs s\u22121 AWG were amplified to about 300 W peak power using a traveling wave tube amplifier. The amplified pulse is broadcasted into the vacuum chamber through two microwave horns, interacting with the vaporized molecules in the pulsed jet. Finally, a total of 40,000 free induction decays (4 FID emissions per gas pulse), of 10 \u03bcs length duration, were averaged and digitized using a 50 Gs s\u22121 digital oscilloscope. Line widths of the order of 100 kHz FWHM were achieved. The sub-Doppler resolution LA-MB-FTMW spectrometer, described elsewhere (Alonso et al. 2009), operating from 4 to 26 GHz, was used to record the NaCl spectra with the resolution necessary to analyze the hyperfine structure due to the presence of two nuclei with I = 3\/2 in the molecule. Microwave pulses of 0.3 \u03bcs duration with powers of 140 mW were applied to polarize the molecules in the jet. The free induction decay (FID) was recorded for 100 \u03bcs in the time domain at 40\u2013100 ns sample intervals and then converted to the frequency domain by Fourier transformation. All the transitions appeared as Doppler doublets due to the parallel configuration of the molecular beam and the microwave radiation. The resonance frequency was determined as the arithmetic mean of the two Doppler components. The estimated accuracy of the frequency measurements is greater than 3 kHz. From 10 (in the case of the ground state and the lower vibrational states) to 2500 (for the higher vibrational states) averages were phase-coherently co-added to achieve reasonable signal-to-noise ratios.","Citation Text":["Berm\u00fadez et al. 2013"],"Functions Text":["NaCl was first investigated using a chirped-pulse Fourier transform microwave (CP-FTMW) spectrometer with a laser ablation source","operating between 6.0 and 12.0 GHz to sample swiftly the rotational spectra of the different species present in the supersonic expansion."],"Functions Label":["Background","Background"],"Citation Start End":[[719,739]],"Functions Start End":[[570,699],[741,878]]}
{"Identifier":"2019MNRAS.489..855C__Husemann_et_al._2013_Instance_1","Paragraph":"The size of ENLRs have been defined in different ways in the literature. Bennert et al. (2002) and Schmitt et al. (2003b) used the Hubble Space Telescope (HST) to obtain narrow band images of $\\rm [O\\, III]$, and adopted the maximum 3\u03c3 detected radius as the radius of the ENLR. This method is subject to the instrumental sensitivity limit that could be very different in different observations. Studies with long-slit spectroscopic observations define the radius based on isophote (Greene et al. 2011; Hainline et al. 2013, 2014), or the distance at which the ionization state changes from AGN to star-forming activities (Bennert et al. 2006a,b). The long-slit based observations also have drawbacks: the morphology of ENLR is sometimes irregular so that the derived size depends on the orientation of slits (Greene et al. 2011; Husemann et al. 2013). We have compared the measured size based on the IFU and the mock long-slit observation in Fig. 4 following the method discussed below. In most cases, long-slit observations tend to underestimate the true size of ENLR. IFU spectroscopic data allow us to use 2D maps to define the sizes of ENLRs. Common definitions include the radius of a specified $\\rm [O\\, III]$ surface brightness isophote (Liu et al. 2013, 2014), or the $\\rm [O\\, III]$ flux weighted radius (Husemann et al. 2013, 2014; Bae et al. 2017). We followed the same method as Liu et al. (2013) but chose a different threshold. The isophote threshold of 10\u221215$\\rm erg\\, s^{-1}cm^{-2}arcsec^{-2}$ was used for quasars related studies. This is suitable for such bright objects but are not as useful for fainter Syferts in our sample as it will leave a large number of AGN undetected. The typical 3\u03c3 depth of the MaNGA observation in $\\rm [O\\, III]$ surface brightness can reach 10\u221217$\\rm erg\\, s^{-1}cm^{-2}arcsec^{-2}$. For our AGN sample, the majority of AGN spaxels have surface brightnesses above 10\u221216$\\rm erg\\, s^{-1}cm^{-2}arcsec^{-2}$ which is thus adopted in this work as the threshold to define the sizes of the ENLRs (hereafter R16). If all spaxels are above this threshold, we extrapolated the fitted $\\rm [O\\, III]$ surface brightness profile to determine R16 (see Section 3.4 for more detail). It should be noted that the surface brightness can be affected by cosmological dimming, which has a scale factor of (1 + z)4 (Liu et al. 2013; Hainline et al. 2014). That is important for works trying to compare sample with different redshift, especially for high redshift quasars.","Citation Text":["Husemann et al. 2013"],"Functions Text":["The long-slit based observations also have drawbacks: the morphology of ENLR is sometimes irregular so that the derived size depends on the orientation of slits"],"Functions Label":["Motivation"],"Citation Start End":[[830,850]],"Functions Start End":[[648,808]]}
{"Identifier":"2021AandA...656A.122D__Triana_et_al._2015_Instance_1","Paragraph":"Understanding how angular momentum and chemicals are transported in the interiors of stars (and planets) along their evolution is one of the key challenges of modern stellar (and planetary) astrophysics. Indeed, rotation modifies their structure, their chemical stratification, their internal flows and magnetism, and their mass losses and winds (e.g. Maeder 2009; Mathis et al. 2013; Aerts et al. 2019, and references therein). In this quest, asteroseismology has bought a fundamental breakthrough by demonstrating that all stars are the seat of a strong extraction of angular momentum during their evolution in comparison with the predictions by stellar models taking the rotation into account following the standard rotational transport and mixing theory (Eggenberger et al. 2012; Marques et al. 2013; Ceillier et al. 2013; Cantiello et al. 2014; Ouazzani et al. 2019). This was first obtained thanks to mixed pulsation modes splitted by rotation propagating in evolved low- and intermediate-mass stars (Beck et al. 2012, 2014, 2018; Mosser et al. 2012; Deheuvels et al. 2012, 2014, 2015; Deheuvels et al. 2020; Di Mauro et al. 2016; Triana et al. 2017; Gehan et al. 2018; Tayar et al. 2019). Then, observations of oscillation modes in F- and A-type stars (Kurtz et al. 2014; Saio et al. 2015; Bedding et al. 2015; Keen et al. 2015; Van Reeth et al. 2015, 2016, 2018; Schmid & Aerts 2016; Murphy et al. 2016; Sowicka et al. 2017; Guo et al. 2017; Saio et al. 2018, 2021; Mombarg et al. 2019; Li et al. 2019, 2020; Ouazzani et al. 2020) and in B-type stars (P\u00e1pics et al. 2015, 2017; Triana et al. 2015; Moravveji et al. 2016; Kallinger et al. 2017; Buysschaert et al. 2018; Szewczuk & Daszy\u0144ska-Daszkiewicz 2018; Pedersen et al. 2021; Szewczuk et al. 2021) provided us new Rosetta stones to constrain the transport of angular momentum in the whole Hertzsprung-Russell diagram. More particularly, this pushes gravity- and gravito-inertial mode pulsators such as \u03b3-Doradus and SPB stars at the forefront of this research. For instance, recent theoretical developments have demonstrated how it is possible to probe stellar internal rotation in \u03b3-Doradus stars from their surface to their convective core (Ouazzani et al. 2020; Saio et al. 2021). These stars are rapid rotators for a large proportion of them. Therefore, it is necessary to study gravito-inertial modes. These modes are gravity modes, which propagate only in stably stratified stellar radiation zones under the action of the restoring buoyancy force in the absence of rotation, which are modified by rotation. If their frequency is super-inertial (i.e. above the inertial frequency 2\u03a9, \u03a9 being the stellar angular velocity), they are propagating in stellar radiation zones and evanescent in convective regions. If their frequency is sub-inertial (below 2\u03a9) they propagate in an equatorial belt in radiation zones and they become propagative inertial waves in convective regions (e.g. Dintrans & Rieutord 2000; Mathis et al. 2014). The challenge of studying these waves is that the equation describing their dynamics are intrinsically bi-dimensional and non-separable (Dintrans et al. 1999; Prat et al. 2016, 2018; Mirouh et al. 2016). This makes the development of seismic diagnosis difficult analytically (Prat et al. 2017) or expansive in computation time when using 2D oscillation and stellar structure codes (e.g. Ouazzani et al. 2017; Reese et al. 2021) in the general case.","Citation Text":["Triana et al. 2015"],"Functions Text":["Then, observations of oscillation modes in","and in B-type stars","provided us new Rosetta stones to constrain the transport of angular momentum in the whole Hertzsprung-Russell diagram."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[1586,1604]],"Functions Start End":[[1196,1238],[1539,1558],[1760,1879]]}
{"Identifier":"2018AandA...613A..15S__Becker_et_al._2016_Instance_1","Paragraph":"In this study, we have outlined and successfully tested a refined technique to measure in contemporary lensing surveys the scale-dependent galaxy bias down to non-linear scales of k ~ 10 h\u22121 Mpc for lens galaxies at z \u2272 0.6. To test our reconstruction technique, we employ a fiducial survey with a sky coverage of ~ 1000 deg2, and a photometry and a survey depth as in CFHTLenS. To construct realistic samples of lenses and sources, we have prepared mock catalogues that are consistent with those used in SES13 and Saghiha et al. (2017). Despite some variations in survey depth and area, these survey parameters are similar to the ongoing Kilo-Degree Survey (KiDS), Dark Energy Survey (DES), or the survey with the Hyper Suprime-Cam (Kuijken et al. 2015; Becker et al. 2016; Aihara et al. 2018). If the galaxy-bias normalisation is perfect, our technique applied to these data can achieve a statistical precision within the range of 5\u201310% (68% CL), if similar lens and source samples are targeted, and a slightly better accuracy of 3\u22127% (68% CL; see Table 3). For the high-z samples, the accuracy will be somewhat higher with 3\u22125%. On the other hand, it is clear from our overview Table 4 that the accuracy of the galaxy-bias normalisation is in fact limited, mainly by our knowledge of the intrinsic alignment of sources, cosmological parameters, and the galaxy redshift distributions. With a broad knowledge of |Aia|\u2272 2 and the specifications for the normalisation errors in Table 4, we conclude that systematic errors would potentially degrade the overall accuracy to approximately 15% for b(k) and 10% for r(k). For fully controlled intrinsic alignment of sources, these errors could be reduced by 5%. An additional reduction by 3% may be possible by controlling the redshift distributions (their mean and variance) in the normalisation to 1% accuracy. For the fiducial cosmology, the knowledge of \u03a9m is of most importance while the normalisation of the ratio statistics is less affected by \u03c38.","Citation Text":["Becker et al. 2016"],"Functions Text":["Despite some variations in survey depth and area, these survey parameters are similar to the ongoing Kilo-Degree Survey (KiDS), Dark Energy Survey (DES), or the survey with the Hyper Suprime-Cam"],"Functions Label":["Similarities"],"Citation Start End":[[755,773]],"Functions Start End":[[538,732]]}
{"Identifier":"2016ApJ...817..156W__Yan_et_al._2014a_Instance_1","Paragraph":"Recently, one of the hot topics in solar physics is the understanding of solar filaments in the corona, including their distribution, formation, eruption, and stability (Yang et al. 2008; Kong et al. 2015; Su et al. 2015; Yan et al. 2015). Martin (1998) and Gaizauskas (2002) have shown that convergence and cancellation of flux play an important role in the formation of filament channels and filaments. Flux ropes and magnetic dips represent the magnetic structures of filaments, which were reported by many authors (van Ballegooijen & Martens 1989; Mackay et al. 1999; Litvinenko & Wheatland 2005; Aulanier et al. 2006; Canou & Amari 2010). Many reports on the eruption of filaments are concerned with torus instability or\/and kink instability (T\u00f6r\u00f6k & Kliem 2003; Kliem & T\u00f6r\u00f6k 2006; T\u00f6r\u00f6k et al. 2010; Yan et al. 2014a, 2014b), and magnetic flux emergence and cancellation are also known to play a key role in these eruptions (Magara & Longcope 2003; Archontis & T\u00f6r\u00f6k 2008; Yan et al. 2011). The transverse component of photospheric magnetic fields near the PIL increase after a filament\u2019s eruption or flares (Liu et al. 2012; Sun et al. 2012; Wang et al. 2013). Sun et al. (2012) indicated that the substantial electric current increases with the emergence of flux during the formation of the filament. A downward collapse of coronal current after the eruption of the filament was also reported by Liu et al. (2012). Nonlinear force-free field (NLFFF) model extrapolation is the most powerful tool to reconstruct the magnetic field above the photosphere from photospheric vector magnetograms (VMs) thus far (Sakurai 1981; Wheatland et al. 2000; Amari et al. 2006; Canou & Amari 2010; Jiang et al. 2014), since the chromospheric and coronal magnetic fields are hard to measure exactly. Even so, it is still a long way to completely understand filaments. Regardless of its formative or eruptive process, and the variation of parameters including the electric current, magnetic field and plasma motion in the evolution of filament are not yet really clear. Investigating the electric current associated with the filament is helpful for understanding the characteristic of solar filaments.","Citation Text":["Yan et al. 2014a"],"Functions Text":["Many reports on the eruption of filaments are concerned with torus instability or\/and kink instability"],"Functions Label":["Background"],"Citation Start End":[[807,823]],"Functions Start End":[[644,746]]}
{"Identifier":"2021MNRAS.507..175S___2008_Instance_1","Paragraph":"Momentum and kinetic energy can be directly transferred to the gas, suppressing inflows. The fast-moving jets can also shock heat the surrounding gas. Many models have invoked kinetic jets to suppress cooling flows and SFRs in massive haloes (e.g. Dubois et al. 2010; Gaspari et al. 2012a; Li & Bryan 2014a; Prasad et al. 2015; Yang & Reynolds 2016a). Many models in the literature also invoke the idea that AGN can effectively drive strong pressure-driven outflows and offset cooling if a large fraction of the accretion energy is thermalized (Begelman 2004; Springel, Di Matteo & Hernquist 2005; Di Matteo, Springel & Hernquist 2005; Hopkins et al. 2006a, b, 2007, 2008; Johansson, Naab & Burkert 2009; Hopkins & Elvis 2010; Ostriker et al. 2010; Faucher-Gigu\u00e8re & Quataert 2012; Dubois et al. 2013; Barai et al. 2014; Weinberger et al. 2017a; Pillepich et al. 2018; Richings & Faucher-Gigu\u00e8re 2018a, b; Torrey et al. 2020). Physically, as the jet propagates, part of the kinetic energy can thermalize through shocks. Some studies have argued that the heat from those weak shocks can suppress cooling flows and SFRs in massive haloes (Yang & Reynolds 2016b; Li, Ruszkowski & Bryan 2017; Martizzi et al. 2019). The magnetic fields carried by the jet at its launch might also help suppress cooling flows by providing additional pressure support (Soker & Sarazin 1990; Beck et al. 1996, 2012), although our studies find that they have limited effects on global star formation properties of sub-L* galaxies (Su et al. 2017).1 Finally, CRs arise generically from processes that occur in fast shocks, so they could come from shocked winds or outflows. But they are particularly associated with relativistic jets from AGN (where they can make up the bulk of the jet energy; Berezinsky, Gazizov & Grigorieva 2006; Ruszkowski, Yang & Reynolds 2017b) and hot, relativistic plasma-filled \u2018bubbles\u2019 or \u2018cavities\u2019 (perhaps inflated by jets in the first place) around AGN. Different authors have argued that they could help suppress cooling flows by providing additional pressure support to the gas, driving pressurized outflows in the galaxy or CGM, or via heating the CGM\/ICM directly via collisional (hadronic & Coulomb) and streaming-instability losses (Guo & Oh 2008; Sharma, Parrish & Quataert 2010; En\u00dflin et al. 2011; Fujita & Ohira 2011; Fujita, Kimura & Ohira 2013; Pfrommer 2013; Wiener, Oh & Guo 2013; Jacob & Pfrommer 2017a, b; Pfrommer et al. 2017; Ruszkowski et al. 2017a, b; Jacob et al. 2018).","Citation Text":["Hopkins et al.","2008"],"Functions Text":["Many models in the literature also invoke the idea that AGN can effectively drive strong pressure-driven outflows and offset cooling if a large fraction of the accretion energy is thermalized"],"Functions Label":["Background"],"Citation Start End":[[636,650],[667,671]],"Functions Start End":[[352,543]]}
{"Identifier":"2022AandA...663A..70F___2017_Instance_1","Paragraph":"Out of these sites (a) might provide the conditions for a very weak r-process and \u03bdp-process, whether only up to Sr, Y, Zr or up to (but not beyond) the A\u2004=\u2004130 peak is still debated (Wanajo et al. 2018; Curtis et al. 2019; Fischer et al. 2020a; Ghosh et al. 2022). (b) is a class of supernovae whose existence is put into question after recent re-determinations of the electron capture rate of 20Ne (Kirsebom et al. 2019a,b), but is not firmly excluded, however, leading to a too strong decline in abundances as a function of A for realistic Ye-conditions (Wanajo et al. 2011). (c) could plausibly lead to magnetars, neutron stars with surface magnetic fields of the order 1014 G, which form in \u223c1 out of 10 of core collapse supernovae (e.g., Beniamini et al. 2019). Dependent on the initial fields, varying weak (probably dominating) to strong r-process conditions can be obtained, the latter, however, only for precollapse fields beyond 1012 G (Winteler et al. 2012; M\u00f6sta et al. 2014, 2015, 2018; Halevi & M\u00f6sta 2018; Nishimura et al. 2015, 2017; Bugli et al. 2020; Reichert et al. 2021). Case (d) has been proposed for a while. Dependent on the nuclear equation of state for massive core-collapse events, the collapse of the proto-neutron star to a black hole can be avoided (in a narrow stellar mass range) due to a quark-hadron phase transition with the right properties. The ejecta would experience a weak r-process, but populating even the actinides, however, with negligible abundances (Fischer et al. 2020b). Case (e) has been extensively discussed in the context of long-duration gamma-ray bursts (Woosley 1993; MacFadyen & Woosley 1999; MacFadyen et al. 2001). They involve the collapse of massive stars that rotate rapidly enough so that an accretion torus can form outside of the last stable orbit of a forming black hole, and they go along with relativistic polar and nonrelativistic torus outflows. This scenario has been proposed by Cameron (2003) as an r-process site and recently been examined in more detail by Siegel et al. (2019) and Siegel (2019). The remaining site, (f), is related to compact binary mergers (see Thielemann et al. 2017; Rosswog et al. 2017; Cowan et al. 2021, for overviews).","Citation Text":["Nishimura et al.","2017"],"Functions Text":["Dependent on the initial fields, varying weak (probably dominating) to strong r-process conditions can be obtained, the latter, however, only for precollapse fields beyond 1012 G"],"Functions Label":["Uses"],"Citation Start End":[[1022,1038],[1045,1049]],"Functions Start End":[[768,946]]}
{"Identifier":"2021ApJ...921...18K__Kushwaha_et_al._2018a_Instance_1","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"],"Functions Text":["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","and the activity of OJ 287 from the end of 2015 to the middle of 2017"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[919,940]],"Functions Start End":[[677,809],[848,917]]}
{"Identifier":"2017MNRAS.472..772M__Cassata_et_al._2015_Instance_1","Paragraph":"Fig. 7 also shows that there is little evidence for a relation between the Ly\u2009\u03b1 luminosity and M1500 for Ly\u2009\u03b1-selected sources at z = 5.7 in our UV and Ly\u2009\u03b1 luminosity range. As both M1500 and LLy\u2009\u03b1 are, to first order, related to the SFR, we would have expected a correlation. To illustrate this, we show lines at constant Ly\u2009\u03b1 escape fractions (based on the assumption that SFRUV = SFRH\u2009\u03b1, case B recombination with T = 10\u2009000 K and ne = 100\u2009cm\u22123 and no attenuation due to dust). This result resembles the well-known Ando et al. (2006) diagram, which reveals a deficiency of luminous LAEs with bright UV magnitudes between z \u2248 5 and 6. More recently, other surveys also revealed that the fraction of high-EW Ly\u2009\u03b1 emitters increases towards fainter UV magnitudes (e.g. Schaerer, de Barros & Stark 2011; Stark, Ellis & Ouchi 2011; Cassata et al. 2015). The lack of a strong correlation between M1500 and LLy\u2009\u03b1 might indicate that the SFRs are bursty (because emission-line luminosities trace SFR over a shorter time-scale than UV luminosity), or that the Ly\u2009\u03b1 escape fraction is anti-correlated with M1500 (such that Ly\u2009\u03b1 photons can more easily escape from galaxies that are fainter in the UV). A possible explanation for the latter scenario is that slightly more evolved galaxies (which are brighter in the UV) have a slightly higher dust content (e.g. Bouwens et al. 2012), affecting their Ly\u2009\u03b1 luminosity more than the UV luminosity. It is interesting to note that several galaxies lie above the 100 \u2009per\u2009cent Ly\u2009\u03b1 escape fraction line. This implies bursty or stochastic star formation (which is more likely in lower mass galaxies with faint UV luminosities, e.g. Mas-Ribas, Dijkstra & Forero-Romero 2016), alternative Ly\u2009\u03b1 production mechanisms to star formation (such as cooling), a higher ionizing production efficiency (for example due to a top-heavy IMF or binary stars, e.g. Gotberg, de Mink & Groh 2017), or dust attenuating Ly\u2009\u03b1 in a different way than the UV continuum (e.g. Neufeld 1991; Finkelstein et al. 2008; Gronke et al. 2016).","Citation Text":["Cassata et al. 2015"],"Functions Text":["More recently, other surveys also revealed that the fraction of high-EW Ly\u2009\u03b1 emitters increases towards fainter UV magnitudes"],"Functions Label":["Background"],"Citation Start End":[[831,850]],"Functions Start End":[[638,763]]}
{"Identifier":"2022ApJ...934..126K__Lotekar_et_al._2016_Instance_1","Paragraph":"We considered a homogeneous, collisionless two species plasma consisting of electrons and ions (H+ ions) in the simulation model. The ambient plasma parameters for ions and electrons are given in Table 2. The ions and electrons are considered to be fluid and their dynamic is incorporated into the simulation model using the following model equations viz., the continuity, momentum, and pressure equations of each species, and the Poisson equation (Kakad et al. 2014) given by\n1\n\n\n\n\u2202nj\u2202t+\u2202(njvj)\u2202x=0,\n\n\n\n2\n\n\n\n\u2202vj\u2202t+vj\u2202vj\u2202x+1mjnj\u2202Pj\u2202x\u2212qjmjE=0,\n\n\n\n3\n\n\n\n\u2202Pj\u2202t+vj\u2202Pj\u2202x+\u03b3jPj\u2202vj\u2202x=0,\n\n\n\n4\n\n\n\n\u2202E\u2202x=\u2211jqjnj\/\u03f50.\n\nHere, the electric field E = \u2212\u2202\u03d5\/\u2202x and the variables n\n\nj\n, P\n\nj\n, and v\n\nj\n are the plasma density, thermal pressure, and velocity of species j, respectively. The subscripts j = e and j = i are, respectively, used for electrons and ions. m\n\nj\n and q\n\nj\n represent the mass and the charge of species j, respectively. For electrons q\n\ne\n = \u2212e and ions q\n\ni\n = e., \u03f5\n0 is the electric permittivity. In Equation (3), electrons and ions are treated as adiabatic with the same adiabatic index \u03b3\n\ne\n = \u03b3\n\ni\n = 3. The above set of equations is solved numerically. The spatial derivatives in these equations are solved using the fourth order central finite difference method and time derivatives are integrated using the leap-frog method to achieve second order accuracy. The details of the development of the simulation model are given in Kakad et al. (2013). We used a compensated filter to eliminate the small wavelength modes linked with such numerical noise (Lotekar et al. 2016; Kakad et al. 2016b). These numerical schemes are highly stable, and in the past, several electrostatic solitary wave structures have been modeled using such fluid simulations in multispecies plasmas (Kakad et al. 2014, 2016a). \u0394x and \u0394t are respectively considered as the grid size in spatial and time domain, and their values are taken in such a way that it fulfills the Courant\u2013Friedrichs\u2013Lewy condition, i.e., \n\n\n\nc\u0394t\u0394x\u22641\n\n, which is necessary for the convergence of the explicit finite difference method. Here, c is the speed of light. We performed two simulation runs in a one-dimensional system with the periodic boundary conditions by considering the observed ambient parameters as (i) T\n\ne\n = 60 eV, T\n\ni\n = 300 eV, n\n\ni0 = 15\/cc, n\n\ne0 = 15\/cc, V\ndi = V\nde = \u221250 km s\u22121 (2) T\n\ne\n = 30 eV, T\n\ni\n = 300 eV, n\n\ni0 = 15\/cc, n\n\ne0 = 15\/cc, V\ndi = V\nde = \u221250 km s\u22121 (see Table 2 for more details). In the simulations, we considered the real mass ratio, i.e., m\n\ni\n\/m\n\ne\n = 1836. The background electron and ion densities satisfy the quasi-neutrality, i.e., n\n\ne0 = n\n\ni0 = n\n0. For electron temperature, two values are considered based on the observations, i.e., 60 and 30  eV (see Figure 6(d)). The values of \u03c9\npi, \u03c9\npe, \u03bb\n\ni\n, and \u03bb\n\ne\n for the considered parameters are 5.1 \u00d7 103 rad s\u22121, 2.18 \u00d7 105 rad s\u22121, 33.3 m, and 14.9 m, respectively. To initiate the simulations, we used a localized Gaussian-type initial density perturbation in the equilibrium electron and ion densities given by\n5\n\n\n\n\u03b4n=\u0394nexp\u2212x\u2212xcl02.\n\nHere, \u0394n and l\n0 are the amplitude and width of the initial density perturbations, respectively. Thus, the perturbed density at t = 0 is n\n\nj\n(x) = n\n\nj0 + \u03b4n. We consider the simulation system length as L\n\nx\n, and x\n\nc\n = L\n\nx\n\/2 is the center of the simulation system. We performed two simulation runs for the parameters given in Table 2. For these simulation runs, we consider the time interval dt = 1 \u00d7 10\u22123, the grid spacing dx = 0.2, system length L\n\nx\n = 7000, l\n0 = 1, and \u0394n = 0.1. Here, time is expressed in units of \u03c9\npi\n\u22121, space is in \u03bb\ndi, and density is in units of n\n\ni0.","Citation Text":["Lotekar et al. 2016"],"Functions Text":["We used a compensated filter to eliminate the small wavelength modes linked with such numerical noise"],"Functions Label":["Uses"],"Citation Start End":[[1560,1579]],"Functions Start End":[[1457,1558]]}
{"Identifier":"2017AandA...607A..71G__Hansen_&_Oh_(2006)_Instance_2","Paragraph":"An implication of the respective escape fractions of the two regimes is visible in Fig. 12. Here we show several values of NHI,cl for the static setup using \u03c4d,cl = 10-4 (empty symbols) and \u03c4d,cl = 1 (filled symbols), which correspond to metallicities of \\hbox{$Z\/Z_\\odot = 0.63\\left(\\tau_{\\rm d}\/10^{-4}\\right)\\left(10^{17}\\cm^{-2}\/N_{\\HI,\\cl}\\right)$}Z\/Z\u2299=0.63\u03c4d(\/10-4)(1017\u2009cm-2\/NHI,cl) (Pei 1992; Laursen et al. 2009); this reaches clearly unrealistic values. However, as in this paper we are interested in the fundamental impact of the individual parameters, we also study these extreme values. Also shown in Fig. 12 (with a black [gray] solid line for the low [high] dust content) is the proposed analytic solution for fesc by Hansen & Oh (2006)(18)\\begin{equation} f_{\\rm esc}^{\\rm HO06} = 1\/{\\rm cosh}(\\!\\sqrt{2 N_{\\cl}\\epsilon}), \\label{eq:fescHO06} \\end{equation}fescHO06=1\/cosh(2Ncl\u03f5),where for Ncl we used Eq. (8)(with (a1,b1) = (3\/2, 2) as found in Sect. 4.1) and for the clump albedo (i.e., the fraction of incoming photons that are reflected) \u03f5, we adopted a value of c1(1\u2212e\u2212 \u03c4d,cl) with c1 = 1.6 [c1 = 0.06] to match the NHI,cl = 1022 cm-2 data points for \u03c4d,cl = 10-4 [\u03c4d,cl = 1]. The behavior for the low- and high-dust contents is quite different. On the one hand, the escape fractions versus NHI,cl scales for \u03c4d,cl = 1 (filled symbols in Fig. 12) as predicted by Hansen & Oh (2006) in their \u201csurface scatter\u201d approximation, that is, a larger clump hydrogen column density \u201cshields\u201d the dust better from the Ly\u03b1 photons and thus prevents their destruction more efficiently. On the other hand, however, this is not the case for the low-dust scenario presented in Fig. 12 (with unfilled symbols) where a larger value of NHI,cl implies a lower fesc. This is because here the dust optical depth through all the clumps (shown in the black dotted line in Fig. 12) is lower than the accumulated dust optical depth through the subsequent random-walk clump encounters (black solid line), i.e., \\hbox{$\\exp(-4\/3 \\fc \\tau_{\\rm d, cl}) \\lesssim f_{\\rm esc}^{\\rm HO06}$}exp(\u22124\/3fc\u03c4d,cl)\u2272fescHO06. Consequently, configurations in the \u201cfree-streaming\u201d regime can possess enhanced Ly\u03b1 escape fractions compared to the \u201crandom walk\u201d regime (see Sect. 5.2 for a more detailed discussion). Still, both cases possess (much) larger escape fractions than a homogeneous slab, which is shown in Fig. 12 with a black dashed line. Here, we use the derived escape fraction by Neufeld (1990) with NHI = 4\/3 \u00d7 fc1022 cm-2 and \u03c4d = 4\/3fc\u03c4d,cl, i.e., with equal column densities as in the NHI,cl = 1022 cm-2 case. ","Citation Text":["Hansen & Oh (2006)"],"Functions Text":["The behavior for the low- and high-dust contents is quite different.","On the one hand, the escape fractions versus NHI,cl scales for \u03c4d,cl = 1 (filled symbols in Fig. 12) as predicted by","in their \u201csurface scatter\u201d approximation","On the other hand, however, this is not the case for the low-dust scenario presented in Fig. 12 (with unfilled symbols) where a larger value of NHI,cl implies a lower fesc."],"Functions Label":["Differences","Similarities","Similarities","Differences"],"Citation Start End":[[1384,1402]],"Functions Start End":[[1198,1266],[1267,1383],[1403,1443],[1594,1766]]}
{"Identifier":"2016MNRAS.457.2433P__Nolan_et_al._2012_Instance_1","Paragraph":"From the result of the \u03c72 minimization, we found that the minimized \u03c72 values agree with the expected values, i.e. the computed \u03c72 are typically in the range of ($\\mathrm{d.o.f.}-\\sqrt{2 \\mathrm{d.o.f.}}$, $\\mathrm{d.o.f.}+\\sqrt{2 \\mathrm{d.o.f.}}$), where d.o.f. is the number of degrees of freedom. This means that the fits describe the observed data rather well. The only exception is with \u03c72 \u2248 20, which occurs for nearby AGN, z  0.2, and for the highest energy band, E > 10 GeV. Note that there is a strong contribution of the source, Mrk 421, in the first redshift interval at high energies, E > 10 GeV for quiescent states. Mrk 421 is a very hard spectrum \u03b3-ray source with a photon index of \u22481.77 and its semiminor and semimajor axes at 68 per cent confidence are of 0$_{.}^{\\circ}$0067 as derived in Nolan et al. (2012). Semiminor and semimajor axes of many 2FGL sources are derived with one order of magnitude higher uncertainties than those for Mrk 421 in the 2FGL catalogue (Nolan et al. 2012). We noted that the discrepancy between the observation and model is particularly strong in the annular bin, r  0$_{.}^{\\circ}$05, for the redshift interval z  0.2 and for the highest energy band, E > 10 GeV. If we exclude photons from Mrk 421, then the minimized \u03c72 value is 7.5 and is consistent with the expected one. In the limit of a large number of counts in each bin, the likelihood is given by $\\mathcal {L}=\\text{exp}(-\\chi ^{2}\/2)$, so that minimizing \u03c72 is equivalent to maximizing the likelihood, $\\mathcal {L}$. We found that the inclusion of a pair halo component in the model does not improve the likelihood value sufficiently to establish the presence of this pair halo component in the data. Therefore, we derived the one-sided 95 per cent upper limit on the fraction of photons attributable to a pair halo component by fitting the normalization of this component, for which we increase its normalization until the maximum likelihood decreases by 2.71\/2 in logarithm. The computed upper limits are between 2 and 6 per cent depending on energy band and redshift interval. These upper limits are stronger than those obtained before. Note that the model for a point-like source used in the likelihood analysis is considered to be precisely established, however, the number of photons recorded during flaring states is close to those numbers of photons recorded during quiescent states for each of these redshift intervals. The expression, such as equation (1), leads to more conservative upper limits on the fraction of photons attributable to a pair halo component, since it takes the error bars assigned to the model into account. If the point-like source model is considered well established, then the error bars shown in Table 3 would decrease by a factor of \u22481.5.","Citation Text":["Nolan et al. (2012)"],"Functions Text":["Mrk 421 is a very hard spectrum \u03b3-ray source with a photon index of \u22481.77 and its semiminor and semimajor axes at 68 per cent confidence are of 0$_{.}^{\\circ}$0067 as derived in"],"Functions Label":["Uses"],"Citation Start End":[[809,828]],"Functions Start End":[[631,808]]}
{"Identifier":"2016AandA...586A..89C__Lo_2005_Instance_1","Paragraph":"The knowledge of physical properties, the structure, and the kinematics of the matter in the vicinity of supermassive black holes (SMBH) is essential to build detailed models of the clumpy outflow and to test the disc-wind scenario. While X-ray variability studies can provide accurate information on the atomic and ionized matter on scales of the BLR, the radio emission from luminous H2O masers (the so-called \u201cmegamasers\u201d) constitutes a fundamental instrument to study the geometry and kinematics of the molecular gas at sub-parsec distance from SMBH. H2O masers may trace distinct regions in the AGN environment, from nearly edge-on accretion discs to nuclear outflows in the form of jets or winds (for recent reviews see Henkel et al. 2005; Lo 2005; Greenhill 2007; Tarchi 2012). Very Long Baseline Interferometry (VLBI) and single-dish monitoring studies of disc-masers allow us to map accretion discs and to determine the enclosed dynamical masses (e.g. Kuo et al. 2011). Jet-masers observations, instead, can provide estimates of the velocity and density of jet material (Peck et al. 2003). H2O maser emission have been also found to be associated with nuclear winds at 1 pc from the nuclear engine. In particular, water maser observations in Circinus (Greenhill et al. 2003) and NGC 3079 (Kondratko et al. 2005) seem to have resolved individual outflowing torus clouds. In fact, Greenhill et al. (2003) discovered that the H2O masers in Circinus trace both a Keplerian disc and a wide-angle outflow which appears to be collimated by the warps of the disc. In NGC 3079, VLBI observations of the maser emission revealed a clumpy thick disc. In addition to that, four maser features were found to be located at high latitude above the disc (at ~0.5 pc from the disc plane) and interpreted as part of a nuclear wind by Kondratko et al. (2005). Proper motion measurements and comparison of these outflow-masers with their disc counterpart provide the most promising method for probing the structure and kinematics of the torus molecular clouds (Nenkova et al. 2008). ","Citation Text":["Lo 2005"],"Functions Text":["While X-ray variability studies can provide accurate information on the atomic and ionized matter on scales of the BLR, the radio emission from luminous H2O masers (the so-called \u201cmegamasers\u201d) constitutes a fundamental instrument to study the geometry and kinematics of the molecular gas at sub-parsec distance from SMBH. H2O masers may trace distinct regions in the AGN environment, from nearly edge-on accretion discs to nuclear outflows in the form of jets or winds (for recent reviews see"],"Functions Label":["Motivation"],"Citation Start End":[[746,753]],"Functions Start End":[[233,725]]}
{"Identifier":"2021MNRAS.501.2522J__Mukherjee_&_Paul_2004_Instance_1","Paragraph":"GX 301-2 is an HMXB consisting of a highly magnetized (B \u223c 4 \u00d7 1012\u2009G, or even larger Doroshenko et al. 2010) pulsar and a B-type hyper-giant star Wray 977 (Vidal 1973; Kaper et al. 1995; Staubert et al. 2019). According to modelling of high-resolution optical spectra, Wray 977 has a mass of 43 \u00b1 10 $\\, \\mathrm{M}_{\\odot }$, a radius of 62 R\u2299, and looses mass through powerful stellar winds at a rate of $\\sim \\! 10^{-5}\\, \\mathrm{M}_\\odot \\, {\\rm yr}^{-1}$ with terminal velocity of 300\u2009$\\rm km\\, s^{-1}$ (Kaper, van der Meer & Najarro 2006). The system is highly eccentric (e \u223c 0.46), with an orbital period of \u223c41.5\u2009d, and exhibits strong variation of the X-ray flux with orbital phase (Koh et al. 1997; Doroshenko et al. 2010). In particular, periodic outbursts at the orbital phase \u223c1.4\u2009d before the periastron passage (Sato et al. 1986), and a fainter one near the apastron passage are observed (Pravdo et al. 1995). The broad-band X-ray spectrum is orbital phase-dependent and can be approximately described as a power law with a high-energy cutoff and a cyclotron resonant scattering feature around 40\u2009keV (Kreykenbohm et al. 2004; Mukherjee & Paul 2004; La Barbera et al. 2005; Doroshenko et al. 2010; Suchy et al. 2012; Islam & Paul 2014; F\u00fcrst et al. 2018; Nabizadeh et al. 2019). During the periastron flares, the source exhibits strong variability with an amplitude of up to a factor of 25, reaching a few hundreds mCrab in the energy band of 2\u201310\u2009keV (e.g. Rothschild & Soong 1987; Pravdo et al. 1995). The flares are accompanied by the variability of the equivalent hydrogen column density ($\\rm \\mathit{ N}_{\\rm H}$) and of the fluorescent iron lines, which is believed to be associated with clumpiness of the stellar wind, launched from the donor star (Mukherjee & Paul 2004). We note the clumpiness in this paper refers to any inhomogeneities in the stellar wind\/stream, which are higher density regions, regardless of its specific formation mechanisms. On the other hand, F\u00fcrst et al. (2011) reported a long XMM\u2013Newton observation in GX 301-2 around its periastron, which also exhibits systematic variations of the flux and $\\rm \\mathit{ N}_{H}$ at a time-scale of a few kiloseconds. Several wind accretion models, consisting of stellar winds and a gas stream, were proposed to explain the observed flares (e.g. Haberl 1991; Leahy 1991; Leahy & Kostka 2008; M\u00f6nkk\u00f6nen et al. 2020).","Citation Text":["Mukherjee & Paul 2004"],"Functions Text":["The broad-band X-ray spectrum is orbital phase-dependent and can be approximately described as a power law with a high-energy cutoff and a cyclotron resonant scattering feature around 40\u2009keV"],"Functions Label":["Background"],"Citation Start End":[[1142,1163]],"Functions Start End":[[925,1115]]}
{"Identifier":"2019MNRAS.486.3741H__Susa_et_al._2015_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":["Susa et al. 2015"],"Functions Text":["The temperature Tcl of each cloud is determined as the result of a one-zone calculation (for details, see"],"Functions Label":["Uses"],"Citation Start End":[[672,688]],"Functions Start End":[[566,671]]}
{"Identifier":"2019AandA...629A..54U__Marinucci_et_al._2015_Instance_4","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"],"Functions Text":["Indeed, only lower limits to the high-energy cut-off have been found with NuSTAR (210 keV"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1910,1931]],"Functions Start End":[[1819,1908]]}
{"Identifier":"2018ApJ...864..158L__Pino_&_Lazarian_2005_Instance_1","Paragraph":"Inspection of Equation (25) reveals that for pitch angles satisfying \n\n\n\n\n\n, \n\n\n\n\n\n so that curvature drift energization is more efficient than generalized betatron energy loss during incompressible contraction\/merging of small-scale flux ropes. When \u03bc2  1\/3, betatron energy loss dominates curvature drift energization. That is why the net acceleration obtained from the two competing acceleration mechanisms depends sensitively on the anisotropy characteristics of the energetic particle pitch-angle distribution as discussed above. Consider the following three possibilities: (1) If energetic particles maintain a highly beamed pitch-angle distribution (which requires negligible pitch-angle scattering), curvature drift energy gain strongly dominates generalized betatron energy loss, and for all practical purposes we have a first-order Fermi acceleration mechanism as a consequence of incompressible contraction or merging of curved flux-rope magnetic fields (de Gouveia dal Pino & Lazarian 2005; Drake et al. 2006, 2010). (2) If the energetic particle distribution stays purely isotropic (extremely strong pitch-angle particle scattering), we can average the last terms in Equation (25) over all \u03bc values to find \n\n\n\n\n\n, indicating that the probability for curvature drift energy gain equals the probability for betatron energy loss (Drake et al. 2010). This supports the conclusion made above that net acceleration requires and depends only on the anisotropic part of the distribution. (3) The energetic particle distribution maintains a particle distribution with a small pitch-angle anisotropy (efficient pitch-angle scattering consistent with the diffusion approximation). In this case particle energization by incompressible contraction or merging of curved flux-rope magnetic fields becomes a second-order Fermi acceleration process (Drake et al. 2013; Zank et al. 2014; le Roux et al. 2015a). The small anisotropy option is supported by self-consistent particle simulations of turbulent magnetic reconnection and island formation at stacked primary current sheets in the absence of a guide field (Schoeffler et al. 2011; Drake et al. 2013), because energetic particles are scattered by fluctuations generated by plasma instabilities such as the firehose and magnetic mirror instabilities, resulting in energetic charged particle distributions with small anisotropies. However, in the presence of a strong guide field, particle simulations suggest larger anisotropies owing to weaker instabilities (Dahlin et al. 2017; Li et al. 2018).","Citation Text":["de Gouveia dal Pino & Lazarian 2005"],"Functions Text":["Consider the following three possibilities: (1) If energetic particles maintain a highly beamed pitch-angle distribution (which requires negligible pitch-angle scattering), curvature drift energy gain strongly dominates generalized betatron energy loss, and for all practical purposes we have a first-order Fermi acceleration mechanism as a consequence of incompressible contraction or merging of curved flux-rope magnetic fields"],"Functions Label":["Uses"],"Citation Start End":[[966,1001]],"Functions Start End":[[535,964]]}
{"Identifier":"2021MNRAS.506..813D__Fabricius_et_al._2014_Instance_1","Paragraph":"Traditionally, GCs have been considered as relatively simple spherical, non-rotating, and almost completely relaxed systems. However, observational results obtained in the past few years are demonstrating that they are much more complex than previously thought. In particular, the classical simplified approach of neglecting rotation in GCs has become untenable from the observational point of view. In fact, there is an increasing wealth of observational results suggesting that, when properly surveyed, the majority of GCs rotate at some level. As of today, more than $50{{\\ \\rm per\\ cent}}$ of the sampled GCs show clear signatures of internal rotation (e.g. Anderson & King 2003; Bellazzini et al. 2012; Fabricius et al. 2014; Bianchini et al. 2018; Ferraro et al. 2018; Kamann et al. 2018a; Lanzoni et al. 2018a,b; Sollima, Baumgardt & Hilker 2019). Moreover, evidence of rotation has also been reported for intermediate-age clusters (Mackey et al. 2013; Kamann et al. 2018b), young massive clusters (H\u00e9nault-Brunet et al. 2012; Dalessandro et al. 2021), and nuclear star clusters (Nguyen et al. 2018; Neumayer, Seth & B\u00f6ker 2020) indicating that internal rotation is a common ingredient across dense stellar systems of different sizes and ages. On the theoretical side, the presence of internal rotation has strong implications on our understanding of the formation and dynamics of GCs and affects, for example, their long-term evolution (Einsel & Spurzem 1999; Ernst et al. 2007; Breen, Varri & Heggie 2017) and their present-day morphology (e.g. Hong et al. 2013; van den Bergh 2008). Moreover, signatures of internal rotation could be crucial in revealing the formation mechanisms of the so-called multiple stellar populations (MPs) in GCs (Bekki 2010; Mastrobuono-Battisti & Perets 2013; H\u00e9nault-Brunet et al. 2015) differing in terms of their light-element (such as He, Na, O, C, N) abundances (see Bastian & Lardo 2018; Gratton et al. 2019 for recent reviews on the subject), and which are observed in almost all GCs now. Differences in the rotation amplitudes of MPs have been observed in two cases so far, namely M 13 and M 80 (Cordero et al. 2017; Kamann et al. 2020)1 and in both clusters the Na-rich population (also known as second population or generation \u2013 SP) is found to rotate with a larger amplitude than the first population FP (Na-poor).","Citation Text":["Fabricius et al. 2014"],"Functions Text":["In particular, the classical simplified approach of neglecting rotation in GCs has become untenable from the observational point of view. In fact, there is an increasing wealth of observational results suggesting that, when properly surveyed, the majority of GCs rotate at some level. As of today, more than $50{{\\ \\rm per\\ cent}}$ of the sampled GCs show clear signatures of internal rotation (e.g."],"Functions Label":["Background"],"Citation Start End":[[708,729]],"Functions Start End":[[262,661]]}
{"Identifier":"2017MNRAS.469..521K__Redfield_2007_Instance_1","Paragraph":"Molecular CO gas is observed in the sub-mm with both single-dish telescopes (JCMT, APEX) and interferometers such as ALMA, the SMA or NOEMA. For the brightest targets, ALMA's high-resolution and unprecedented sensitivity allow us to obtain CO maps for different lines and isotopes showing the location of CO belts and giving an estimate of their mass (see the CO gas disc around \u03b2 Pic, Dent et al. 2014; Matr\u00e0 et al. 2017). Atomic species are also detected around a few debris disc stars. In particular, Herschel was able to detect the O\u2009i and C\u2009ii fine structure lines in two and four systems, respectively (e.g. Riviere-Marichalar et al. 2012, 2014; Roberge et al. 2013; Cataldi et al. 2014; Brandeker et al. 2016). Also, metals have been detected, using UV\/optical absorption lines, around \u03b2 Pictoris (Na, Mg, Al, Si and others, Roberge et al. 2006), 49 Ceti (Ca\u2009ii, Montgomery & Welsh 2012) and HD 32297 (Na\u2009i, Redfield 2007). Some of these metals are on Keplerian orbits but should be blown out by the ambient radiation pressure (Olofsson, Liseau & Brandeker 2001). It is proposed that the overabundant ionized carbon observed around \u03b2 Pic, which is not pushed by radiation pressure could brake other ionized species due to Coulomb collisions with them (Fern\u00e1ndez, Brandeker & Wu 2006). A stable disc of hydrogen has not yet been observed in these systems (Freudling et al. 1995; Lecavelier des Etangs et al. 2001) but some high velocity H\u2009i component (presumably falling on to the star) was detected recently with the HST\/COS around \u03b2 Pic (Wilson et al. 2017). All these observations need to be understood within the framework of a self-consistent model. Models of the emission of the gas around main sequence stars have been developed, but gas radial profiles were not derived self-consistently and often assumed to be Gaussian (e.g. Zagorovsky, Brandeker & Wu 2010) or not to be depleted in hydrogen compared to solar (as expected in debris discs, e.g. Gorti & Hollenbach 2004) or both (e.g. Kamp & Bertoldi 2000).","Citation Text":["Redfield 2007"],"Functions Text":["Also, metals have been detected, using UV\/optical absorption lines, around","and HD 32297 (Na\u2009i,"],"Functions Label":["Background","Background"],"Citation Start End":[[915,928]],"Functions Start End":[[718,792],[895,914]]}
{"Identifier":"2018MNRAS.475.1104B__Leonard_et_al._2001_Instance_1","Paragraph":"Observational evidence suggests that SNe IIn are aspherical and may have high polarization signals. An \u223c 20\u2009per\u2009cent level of SN asphericity may result in a detectable 1\u2009per\u2009cent linear polarization signal (H\u00f6flich 1991; Leonard & Filippenko 2005). While a number of efforts have been made to explain core-collapse SNe in terms of axisymmetric jets (Khokhlov et al. 1999; Wheeler, Meier & Wilson 2002; Wang et al. 2002), observational evidence in the form of loops in the Q\/U plane suggests that even these axisymmetric models may not be sufficient for all types of core-collapse SNe (Hoffman et al. 2008; Wang & Wheeler 2008; Maund et al. 2009). In contrast to SNe IIn, SNe II-P generally have shown very low levels of polarization at early times (Leonard & Filippenko 2001; Leonard et al. 2002a; however, see Leonard et al. 2016; Mauerhan et al. 2017). The initially low polarization levels often rise during the plateau phase (e.g. Leonard et al. 2016), with a polarization angle that typically remains nearly fixed throughout (e.g. Leonard et al. 2001; Mauerhan et al. 2017). Occasionally, a sharp rise in the polarization signal is seen during the transition to the nebular phase (Leonard et al. 2006; Chornock et al. 2010), perhaps suggesting that the core of the SN is more aspherical than the early-time photosphere (Chugai et al. 2005; Chugai 2006). However, as demonstrated by the modelling of Dessart & Hillier (2011), it is also possible that even large asymmetries during the plateau phase will produce very little polarization, owing to the high optical depth to electron scattering and the fact that geometric information is lost due to multiple scatters. The \u2018spike\u2019 that is sometimes seen during the drop off of the plateau may, therefore, be more of an optical-depth effect (i.e. the \u2018spike\u2019 occurs when $\\tau _{e^-}$ has decreased to unity) than a demonstration of increasing asphericity with depth in the atmosphere (Leonard et al. 2012). The picture for SNe IIn, on the other hand, is not as well understood. The primary reason for this is that an effective model for a SN IIn must not only account for the geometry of the SN ejecta, but also the geometry of the CSM interaction region (Chugai 2001). In such cases, the temporal evolution that multi-epoch spectropolarimetry provides becomes particularly important in establishing a physical model.","Citation Text":["Leonard et al. 2001"],"Functions Text":["The initially low polarization levels often rise during the plateau phase","with a polarization angle that typically remains nearly fixed throughout (e.g."],"Functions Label":["Background","Background"],"Citation Start End":[[1036,1055]],"Functions Start End":[[855,928],[957,1035]]}
{"Identifier":"2020MNRAS.499.5230F__Tripp,_Savage_&_Jenkins_2000_Instance_1","Paragraph":"Different methods have been proposed to detect the hot, highly ionized WHIM gas: detection in galaxy groups with Sunyaev\u2013Zeldovich effect (Hill et al. 2016; de Graaff et al. 2019; Lim et al. 2020; Tanimura et al. 2020) using autocorrelation function measurements (Galeazzi et al. 2010), with absorption lines in quasar sightline (Kov\u00e1cs et al. 2019) and using CMB as a backlight (Ho, Dedeo & Spergel 2009). Given the challenges of X-ray data, observations at longer wavelengths (UV and optical) benefit from higher instrumental throughout, enhanced spectral resolution. By reverting to ground-based facilities, longer exposure times and larger number of targets become possible. Nevertheless, the UV lines have so far mostly been used to detect absorbing gas with temperature range $10^5\\, \\mathrm{K} \\lt T \\lt 10^6\\, \\mathrm{K}$ from either O\u2009vi (Tripp, Savage & Jenkins 2000; Danforth & Shull 2005; Danforth & Shull 2008; Tripp et al. 2008; Savage et al. 2014; Werk et al. 2014; Danforth et al. 2016; Sanchez, Morisset & Delgado-Inglada et al. 2016) or BLAs (Lehner et al. 2007; Danforth, Stocke & Shull 2010). Recently, Zastrocky et al. (2018) have constrained the Milky Way\u2019s hot (T = 2 \u00d7 106 K) coronal gas using the forbidden 5302\u2009\u00c5 transition of Fe\u2009xiv. Only recently, some highly ionized iron UV lines detected in emission have been used as diagnostics of gas at temperatures of T = 107 K. Out of several forbidden lines in the UV that could trace this gas temperature range, and from various species of highly ionized iron, the emission of [Fe\u2009xxi] is the brightest (Anderson & Sunyaev 2016). Anderson & Sunyaev (2018) report the discovery of [Fe\u2009xxi] in emission in a filament projected 1.9 kpc from the nucleus of M87. Theoretically, the highly ionized iron UV lines can be observed in absorption as well. The forbidden line of [Fe\u2009xxi], in particular, has the largest effective cross-section for absorption and a rest wavelength \u03bb1354\u2009\u00c5, conveniently close to Ly \u03b1 \u03bb1215\u2009\u00c5.","Citation Text":["Tripp, Savage & Jenkins 2000"],"Functions Text":["Nevertheless, the UV lines have so far mostly been used to detect absorbing gas with temperature range $10^5\\, \\mathrm{K} \\lt T \\lt 10^6\\, \\mathrm{K}$ from either O\u2009vi"],"Functions Label":["Background"],"Citation Start End":[[848,876]],"Functions Start End":[[679,846]]}
{"Identifier":"2017AandA...602A..29B__Shepherd_1997_Instance_1","Paragraph":"The MOJAVE survey provides access to excellent Very Long Baseline Array (VLBA) data taken at 15 GHz. This is of great value for investigating AGN properties on a statistical basis (e.g., Lister et al. 2016; Homan et al. 2015). Whereas the MOJAVE team is providing a statistical analysis of the complete sample, our approach is to select and focus on specific sources which appear to be special due to unique or rare properties. Although a statistical analysis of large numbers of sources is certainly of utmost importance and great value, a detailed analysis concentrating on peculiar properties \u2013 that are not necessarily common to the AGN population \u2013 can produce complementary results. We re-modeled fifty VLBA observations of 1308+326 obtained at 15 GHz (taken from the online MOJAVE archive webpage) between 1995.05 and 2014.07 with Gaussian components within the difmap-modelfit program (Shepherd 1997). The modelfit program fits image-plane model components to the visibilities in the uv plane. We did not apply self-calibration but used the calibration as provided by the MOJAVE team. Only circular components were used. The use of exclusively circular Gaussian components proved to be the most efficient way to parameterize the component properties. It reduces the number of free parameters, compared to the use of elliptical Gaussians, and allows a more secure identification of components across the epochs. Although it might have advantages to model the data of some of the epochs with elliptical components, our experience is that circular component modeling allows a more homogeneous analysis of all the epochs and more trustworthy identification of individual components in their long-term motion and evolution. Every epoch was modeled independently starting from a point source model. The errors were estimated from deviations in all parameters derived by calculating fits to models with \u00b11 component. All the images with model-fits superimposed are displayed in Figs. A.1\u2013A.13. The parameters and corresponding uncertainties of the model-fits are listed in Table A.1. Components labeled with an \u201cx\u201d denote the presence of features that could not be reliably traced across the epochs. In addition we analyzed single-dish total intensity and polarization radio monitoring data at three radio frequencies, which was observed by the UMRAO monitoring program. ","Citation Text":["Shepherd 1997"],"Functions Text":["We re-modeled fifty VLBA observations of 1308+326 obtained at 15 GHz (taken from the online MOJAVE archive webpage) between 1995.05 and 2014.07 with Gaussian components within the difmap-modelfit program","The modelfit program fits image-plane model components to the visibilities in the uv plane.","We did not apply self-calibration but used the calibration as provided by the MOJAVE team. Only circular components were used. The use of exclusively circular Gaussian components proved to be the most efficient way to parameterize the component properties. It reduces the number of free parameters, compared to the use of elliptical Gaussians, and allows a more secure identification of components across the epochs. Although it might have advantages to model the data of some of the epochs with elliptical components, our experience is that circular component modeling allows a more homogeneous analysis of all the epochs and more trustworthy identification of individual components in their long-term motion and evolution."],"Functions Label":["Uses","Background","Compare\/Contrast"],"Citation Start End":[[894,907]],"Functions Start End":[[689,892],[910,1001],[1002,1726]]}
{"Identifier":"2016AandA...586A.156K__Osorio_et_al._(2011)_Instance_1","Paragraph":"In this study, we use a model atom of Li\u2009i which was originally developed and tested by Cayrel et al. (2007) and Sbordone et al. (2010). For the purposes of the current work, the model atom was updated and now consists of 26 levels and 123 (96 of which are radiative) bound-bound transitions of Li\u2009i and the ground level of Li\u2009ii, with each level of Li\u2009i coupled to the continuum via bound-free transitions. (The ground state of Li\u2009ii in the current model atom is always in LTE, since lithium is mostly fully ionized throughout the model atmospheres studied in this work.) This renders the model atom complete up to the principal quantum number n = 6 and spectroscopic term \\hbox{$^2{\\rm F}^{\\rm o}$}Fo2, with additional energy levels up to n = 9 and term 2D (Fig. 1). Data concerning atomic energy levels and transitions (level energies and statistical weights; wavelengths and Einstein coefficients of the bound-bound transitions) were taken from the NIST database. We used electron collisional excitation and ionization rates from the quantum mechanical computations of Osorio et al. (2011) for the energy levels of up to 5s (2S). Elsewhere, collisional excitation by electrons for radiatively permitted transitions was accounted for by using the classical formula of van Regemorter (1962), while the formula of Seaton (1962) was used to compute collisional electron ionization rates. To account for the collisional excitation by hydrogen, we used collisional excitation rates computed by Barklem et al. (2003), while the classical formula of Drawin (in the formulation of Lambert 1993) was used for radiatively permitted transitions when no quantum mechanical data were available. Hydrogen H\u2013Li charge transfer rates were taken from Barklem et al. (2003) for the atomic levels up to 4p inclusive. Bound-free transitions resulting from collisions with hydrogen were expected to be inefficient and thus were ignored. Photoionization cross sections were taken from TOPBASE (Cunto et al. 1993). No scaling of collisional rates was applied in the calculations of bound-free and bound-bound transitions. Information about the energy levels and bound-bound radiative transitions, included in the present version of the Li\u2009i model atom, are provided in Tables A.1 and A.2, respectively. Twenty-seven transitions in the model atom are purely collisional. Collisional radiatively-forbidden transitions involving Li\u2009i levels beyond 5s were not accounted for since reliable quantum-mechanical data for these transitions are not available. We note that the role of the omitted transitions between the higher levels is minor: when they are taken into account using the formula of Allen (1973), collision strength \u03a9 = 1, the change in the estimated abundance (which directly applies to abundance corrections, too) is always less than 0.05 dex, with typical values being significantly smaller. ","Citation Text":["Osorio et al. (2011)"],"Functions Text":["We used electron collisional excitation and ionization rates from the quantum mechanical computations of","for the energy levels of up to 5s (2S)."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1073,1093]],"Functions Start End":[[968,1072],[1094,1133]]}
{"Identifier":"2016MNRAS.462.2275K__Porubcan,_Kornos_&_Williams_2004_Instance_1","Paragraph":"Currently, the concept of the existence of genetic relations between comets and meteoroid streams as well as near-Earth asteroids (NEAs) is generally accepted. As a consequence, a new definition, the \u2018comet\u2013asteroid\u2013meteoroid complex\u2019, was introduced for the indication of the families of objects with a common origin. The certainty of the existence of such associations has been proved by numerous studies of the dynamical properties of some small bodies. For instance, the association of comet 2P\/Encke with the Taurid meteoroid stream was confirmed by a number of investigations and does not raise any doubts. Furthermore, it was shown that more than 40 asteroids belong to the Taurid complex, i.e. they move on orbits close to those of the comet Encke and the Taurid stream. A cometary nature of these NEAs was suggested and a conclusion was made that they are either extinct fragments of comet Encke or represent (together with comet Encke) the remnants of a larger comet-progenitor of the stream (Asher, Clube & Steel 1993a; Babadzhanov 2001; Porubcan, Kornos & Williams 2004, 2006; Babadzhanov, Williams & Kokhirova 2008a; Rudawska, Vaubaillon & Jenniskens 2012a,b; Madiedo et al. 2013). This family of near-Earth objects with a common origin is named the Taurid comet\u2013asteroid\u2013meteoroid complex. The Quadrantid comet\u2013asteroid\u2013meteoroid complex is another well-known set of related NEOs that very probably includes comet 96P\/Machholz 1 and certainly NEA 2003EH1 and the Quadrantid meteoroid stream, producing eight meteor showers observable on the Earth. It was shown that asteroid 2003EH1 is in fact the extinct fragment of a parent comet of the Quadrantid stream (Jenniskens 2004; Williams et al. 2004; Babadzhanov, Williams & Kokhirova 2008b; Neslusan, Hajdukova & Jakubik 2013). It turned out that the Piscid meteoroid stream contains three related NEAs moving within the stream and representing the extinct debris of a larger comet-progenitor of this complex (Babadzhanov & Williams 2007; Babadzhanov, Williams & Kokhirova 2008c). Based on investigation of the dynamical properties, it was found that three NEAs with similar comet-like orbits are associated with the \u03b9 Aquariid meteoroid stream and this asteroid\u2013meteoroid complex is the result of the break-up of a parent comet (Babadzhanov, Williams & Kokhirova 2009).","Citation Text":["Porubcan, Kornos & Williams 2004"],"Functions Text":["A cometary nature of these NEAs was suggested and a conclusion was made that they are either extinct fragments of comet Encke or represent (together with comet Encke) the remnants of a larger comet-progenitor of the stream"],"Functions Label":["Background"],"Citation Start End":[[1049,1081]],"Functions Start End":[[779,1001]]}
{"Identifier":"2016ApJ...822...22O__Orlando_et_al._2015_Instance_1","Paragraph":"We used the FLASH code (Fryxell et al. 2000) to perform the calculations. In particular we solved the equations for compressible gas dynamics with the FLASH implementation of the piecewise-parabolic method (Colella & Woodward 1984). The radiative losses \u039b in Equation (2) are calculated through a table lookup\/interpolation method. Also we extended the code by additional computational modules to calculate the deviations from electron-proton temperature-equilibration and the deviations from equilibrium of ionization of the most abundant ions. For the former, we calculated the ion and electron temperatures in each cell of the post-shock medium, taking into account the effects of Coulomb collisions (see Orlando et al. 2015 for the details of the implementation). According to Ghavamian et al. (2007), first the electrons are assumed to be heated at the shock front almost istantaneously up to kT \u223c 0.3 keV by lower hybrid waves. This istantaneous heating does not depend on the shock Mach number and is expected for fast shocks (i.e., >103 km s\u22121) as those simulated here. Then we considered the effects of the Coulomb collisions to calculate the evolution of ion and electron temperatures in each cell of the post-shock medium in the time \u0394tj = t \u2212 tshj, where tshj is the time when the plasma in the jth domain cell was shocked and t is the current time. The time tshj is stored in an additional passive tracer added to the model equations. To estimate the deviations from equilibrium of ionization of the most abundant ions, we adopted the approach suggested by Dwarkadas et al. (2010). In fact, this approach ensures high efficiency in the calculation (expecially in the case of 3D simulations as in our case) as well as a reasonable accuracy in the evaluation of the non-equilibrium of ionization effects. The approach consists of the computation of the maximum ionization age in each cell of the spatial domain \u03c4j = nej\u0394tj, where nej is the electron density in the jth cell and \u0394tj is the time since when the plasma in the cell was shocked (see above).","Citation Text":["Orlando et al. 2015"],"Functions Text":["For the former, we calculated the ion and electron temperatures in each cell of the post-shock medium, taking into account the effects of Coulomb collisions (see","for the details of the implementation)."],"Functions Label":["Uses","Uses"],"Citation Start End":[[708,727]],"Functions Start End":[[546,707],[728,767]]}
{"Identifier":"2022AandA...660A.135V__Spina_et_al._2016_Instance_1","Paragraph":"The first study, to our knowledge, to notice the net increase in the abundance of slow (s) neutron capture elements in young stellar populations is D\u2019Orazi et al. (2009), in which the abundance of barium in young star clusters was seen to be higher than in the older ones. Maiorca et al. (2011, 2012) added a few more elements with important s-process contributions (yttrium, zirconium, lanthanum, and cerium), confirming the increasing trend towards younger ages. Subsequently, a number of works have attempted to both clarify the origin of this increase (see, e.g., Bisterzo et al. 2014; Mishenina et al. 2015; Trippella et al. 2016; Magrini et al. 2018; Spina et al. 2018; Busso et al. 2021) and to use their abundances to estimate the ages of stars, often using neutron capture s-process elements in combination with other elements with opposite behaviours, such as \u03b1 elements \u2013 that we indicate as chemical clocks \u2013 and thus maximising the dependence of the relationship with age (see, e.g., Tucci Maia et al. 2016; Nissen 2016; Feltzing et al. 2017; Fuhrmann et al. 2017; Slumstrup et al. 2017; Titarenko et al. 2019). Once the existence of a relationship between age and chemical clocks was established (see, e.g., Spina et al. 2016; Delgado Mena et al. 2019; Jofr\u00e9 et al. 2020), the next steps were the following: (i) to clarify the applicability of these relationships with luminosity class (dwarf or giant) (see, e.g., Tucci Maia et al. 2016; Slumstrup et al. 2017; Casamiquela et al. 2021), metallicity (see, e.g., Feltzing et al. 2017; Delgado Mena et al. 2019; Casali et al. 2020), and population type (thin disc, thick disc, halo) (see, e.g., Titarenko et al. 2019; Nissen et al. 2020; Tautvai\u0161ien\u0117 et al. 2021), or even in dwarf galaxies (Sk\u00falad\u00f3ttir et al. 2019; ii) to calibrate them with a sample of stars with reliable age determination, which are usually open star clusters (OCs), solar twins, or targets with asteroseismic observations. Finally, it is essential to understand whether these relationships are valid throughout the Galactic disc, or whether they are necessarily position-dependent. For the first time, Casali et al. (2020) applied the relations derived from a large sample of solar-like stars located in the solar neighbourhood and noted that they fail to reproduce the ages of star clusters in the inner disc. They concluded that the relationship between age and chemical clocks is not universal and that it varies with galactocentric position. Later, Magrini et al. (2021b) suggested that the differences in the relationships between age and chemical clocks in different parts of the Galactic disc are due to the strong dependence on the metallicity of the yields of low-mass stars, which produce s-process elements during the final stages of their evolution. Casamiquela et al. (2021) used red clump stars in open clusters to investigate the age dependence of several abundance ratios, including those that contain s-process and \u03b1 elements. They found that the relationship between [Y\/Mg] and ages outlined by open clusters is similar to the one found using solar twins in the solar neighbourhood. They also found that the abundance ratios involving Ba are those with the highest correlation with age. However, they also note that as one moves away from the solar neighbourhood, the dispersion increases and is in agreement with the findings of Casali et al. (2020), which attributed this to the spatial variation of the star formation history along the galactocentric radius.","Citation Text":["Spina et al. 2016"],"Functions Text":["Once the existence of a relationship between age and chemical clocks was established (see, e.g.,"],"Functions Label":["Background"],"Citation Start End":[[1222,1239]],"Functions Start End":[[1125,1221]]}
{"Identifier":"2021MNRAS.505.5833F__Lesgourgues_2011_Instance_1","Paragraph":"Besides the Patchy and the LN mocks, we also model the multipoles of the BOSS CMASS two-point correlation function using an analytic approach, which is required to run the Monte Carlo analysis (see Section 5). The 2PCF can be obtained from the Fourier transform of the matter power spectrum, P(k), for which we assume the template from Padmanabhan & White (2008):\n(10)$$\\begin{eqnarray*}\r\nP(k)=\\left[P_{\\rm {lin}}(k)-P_{\\rm {dw}}(k)\\right]e^{-k^2\\Sigma _{\\rm {nl}}^2\/2}+P_{\\rm {dw}}(k) .\r\n\\end{eqnarray*}$$In the equation above, Plin(k) is the linear matter power spectrum computed using the Boltzmann code CLASS (Lesgourgues 2011), assuming the Planck 2015 (Ade et al. 2016) fiducial cosmology. The Pdw(k) term is the de-wiggled power spectrum (Eisenstein & Hu 1998), while the \u03a3nl parameter encodes the smoothing of the BAO peak due to non-linear effects (Crocce & Scoccimarro 2006). The multipoles of the analytic 2PCF are defined as\n(11)$$\\begin{eqnarray*}\r\n\\xi _l(s) = \\frac{i^l}{2\\pi ^2}\\int _0^{\\infty } P_l(k)j_l(ks)k^2{\\rm d}k ,\r\n\\end{eqnarray*}$$from which we recover the monopole (l = 0) and the quadrupole (l = 2). In equation (11), jl(x) represents the spherical Bessel function of first kind and order l, while Pl(k) are the multipoles of the power spectrum defined as\n(12)$$\\begin{eqnarray*}\r\nP_l(k)=\\frac{2l+1}{2}\\int ^1_{-1}\\left(1+f\\mu ^2\\right)^2P(k)L_l(\\mu){\\rm d}\\mu ,\r\n\\end{eqnarray*}$$where Ll(x) is the Legendre polynomial of order l and P(k) is the template given in equation (10). By replacing equation (12) in equation (11), the analytic expressions for monopole (l = 0) and quadrupole (l = 2) are respectively (Xu et al. 2012):\n(13)$$\\begin{eqnarray*}\r\n\\xi _{\\rm {model}}^{(0)}(s) = B_0\\xi _0(\\alpha s)+a_0^{(0)}+\\frac{a_1^{(0)}}{s}+\\frac{a_2^{(0)}}{s^2} ,\r\n\\end{eqnarray*}$$(14)$$\\begin{eqnarray*}\r\n\\xi _{\\rm {model}}^{(2)}(s) = B_2\\xi _2(\\alpha s)+a_0^{(2)}+\\frac{a_1^{(2)}}{s}+\\frac{a_2^{(2)}}{s^2} ,\r\n\\end{eqnarray*}$$where \u03b1 is the shift parameter, while $(a_1^{(i)},a_2^{(i)},a_3^{(i)})$ are linear nuisance parameters.","Citation Text":["Lesgourgues 2011"],"Functions Text":["In the equation above, Plin(k) is the linear matter power spectrum computed using the Boltzmann code CLASS","assuming the Planck 2015 (Ade et al. 2016) fiducial cosmology."],"Functions Label":["Uses","Uses"],"Citation Start End":[[614,630]],"Functions Start End":[[506,612],[633,695]]}
{"Identifier":"2019ApJ...875...63M__Stern_et_al._2019_Instance_1","Paragraph":"The technique of weak lensing offers direct measurement of the total matter distribution of a galaxy cluster (baryonic and dark matter), and can thus provide an unbiased mass calibration. Weak lensing manifests itself as small but coherent distortions of distant galaxies that result from the gravitational deflection of light due to foreground structures (e.g., Kaiser 1992). Cluster weak lensing appears as a tangential shear of background galaxy shapes around a cluster. Numerous attempts to calibrate SZ masses have been made in the literature using ACT clusters (Miyatake et al. 2013; Jee et al. 2014; Battaglia et al. 2016), SPT clusters (McInnes et al. 2009; High et al. 2012; Schrabback et al. 2018; Stern et al. 2019; Dietrich et al. 2019), Planck clusters (von der Linden et al. 2014b; Hoekstra et al. 2015; Penna-Lima et al. 2017; Sereno et al. 2017; Medezinski et al. 2018a), Planck and SPT clusters (Gruen et al. 2014), and other massive cluster samples (Marrone et al. 2009, 2012; Hoekstra et al. 2012; Smith et al. 2016). The mass calibration is often parameterized as\n1\n\n\n\n\n\nwhere MSZ is the SZ mass and Mtrue is the true cluster mass, which for this paper we take to be the weak-lensing mass MWL. This ratio can be taken for individual clusters or for an ensemble average and these values will be consistent as long as the appropriate weights are used (Medezinski et al. 2018a). Recently, Planck Collaboration et al. (2016d) reported a disagreement between 1\u2212b obtained by weak-lensing calibrations of Planck SZ cluster masses (e.g., von der Linden et al. 2014b; Hoekstra et al. 2015) and that inferred from reconciling the Planck primary CMB parameters with the Planck SZ cluster counts. This disagreement is not statistically significant (\u223c2\u03c3) and will decrease after accounting for additional bias corrections, like Eddington bias (Battaglia et al. 2016) and new optical depth measurements (Planck Collaboration et al. 2016e). However, if such a disagreement persists as the precision of cluster measurements improves, then it could reveal the need for extensions to the standard cosmological model (Planck Collaboration et al. 2016c), like a non-minimal sum of neutrino masses (e.g., Wang et al. 2005; Shimon et al. 2011; Carbone et al. 2012; Mak & Pierpaoli 2013; Louis & Alonso 2017; Madhavacheril et al. 2017), or illuminate additional systematic effects in cluster abundance measurements.","Citation Text":["Stern et al. 2019"],"Functions Text":["Numerous attempts to calibrate SZ masses have been made in the literature using","SPT clusters"],"Functions Label":["Background","Background"],"Citation Start End":[[708,725]],"Functions Start End":[[474,553],[631,643]]}
{"Identifier":"2019AandA...625A.148D__Li_et_al._2011_Instance_1","Paragraph":"With a stellar mass of 5\u2005\u00d7\u20051010\u2006M\u2299 (Viaene et al. 2014), Andromeda belongs to the transition regime between the active blue-sequence galaxies and passive red-sequence galaxies (e.g. Bower et al. 2017; Baldry et al. 2006), which happens around the stellar mass of 3\u2005\u00d7\u20051010\u2006M\u2299 (e.g. Kauffmann et al. 2003). This galaxy is a prototype galaxy from the Local Group where the star formation has been quenched in the central part. Andromeda hosts both very little gas and very little star formation, while the black hole is basically quiet and has some murmurs (Li et al. 2011). In a previous study about M 31 nucleus, Melchior & Combes (2017) have shown that there is no gas within the sphere of influence of the black hole. Indeed, the gas has been exhausted. Most scenarios of the past of evolution of Andromeda reproduce the large scale distribution and show evidence of a past activity rich in collisions (Ibata et al. 2001, 2014; Thilker et al. 2004; Gordon et al. 2006; McConnachie et al. 2009; Miki et al. 2016; Hammer et al. 2018). However, the exact mechanism quenching the activity in the central kiloparsec is still unknown (Tenjes et al. 2017). Block et al. (2006) proposed a frontal collision with M 32, which could account for the two ring structures observed in the dust distribution. Melchior & Combes (2011, 2016) showed the presence of gas along the minor axis and support the scenario of the superimposition of an inner 1 kpc ring with an inner disc. Melchior & Combes (2013) estimated a minimum total mass of 4.2\u2005\u00d7\u2005104\u2006M\u2299 of molecular gas within a (projected) distance to the black hole of 100 pc. This is several orders of magnitude smaller than the molecular gas present in the central molecular zone of the Milky Way (Pierce-Price et al. 2000; Molinari et al. 2011). In the Galaxy, while large amounts of dense gas are present in the central region, Kruijssen et al. (2014) discussed the different processes that combine to inhibit star formation, which was observed to be a factor ten times weaker than expected (e.g. Leroy et al. 2008).","Citation Text":["Li et al. 2011"],"Functions Text":["Andromeda hosts both very little gas and very little star formation, while the black hole is basically quiet and has some murmurs"],"Functions Label":["Background"],"Citation Start End":[[555,569]],"Functions Start End":[[424,553]]}
{"Identifier":"2021MNRAS.502..915C__Cisneros-Parra_1970_Instance_1","Paragraph":"Under the Applegate model, the change in orbital period is directly related to the change in the companion star\u2019s gravitational quadrupole moment Q (Applegate & Patterson 1987),\n(8)$$\\begin{eqnarray*}\r\n\\frac{\\Delta P_{\\rm orb}}{P_{\\rm orb}} = -9\\frac{\\Delta Q}{M_{\\rm c} A^2},\r\n\\end{eqnarray*}$$where A = x(1 + q)\/sin\u2009i is the orbital separation. For comparison, the total quadrupole moment induced by the spin of the companion star and the tidal distortion in the pulsar\u2019s gravitational field is (Voisin, Breton & Summers 2020a)\n(9)$$\\begin{eqnarray*}\r\n\\frac{Q}{M_{\\rm c} A^2} = -\\frac{2}{9} k_2 \\left(\\frac{R_{\\rm c}}{A}\\right)^5 \\left(4 q + 1\\right),\r\n\\end{eqnarray*}$$where Rc is the radius of the companion star and k2 is the apsidal motion constant, a parameter describing the deformability of the companion star (Sterne 1939). For solar-type stars k2 \u223c 0.035 (Ogilvie 2014), while if we assume that redback companions are akin to the companions in CV systems whose outer envelopes have also been stripped through accretion, then we may expect a smaller value k2 \u223c 10\u22123 (Cisneros-Parra 1970). For J2039, the hyperparameter $h = 3.9^{+2.2}_{-1.2}$\u2009s corresponds to the typical fractional amplitude for the variations in orbital phase. Taking the simpler squared exponential covariance function of equation (4) corresponding to n \u2192 \u221e, then the deviations in orbital period have covariance function,\n(10)$$\\begin{eqnarray*}\r\nK_{\\Delta P_{\\rm orb}\/P_{\\rm orb}}(t_1,t_2) &=& \\frac{\\partial ^2 K}{\\partial t_1 \\partial t_2} \\nonumber\\\\\r\n&=& \\frac{h^2}{l^2} \\exp \\left(\\!-\\frac{(t_1 - t_2)^2}{2\\ell ^2}\\!\\right) \\left(\\!1 - \\frac{(t_1 - t_2)^2}{\\ell ^4}\\!\\right).\r\n\\end{eqnarray*}$$The typical (fractional) amplitude of the orbital period variations is therefore \u0394Porb\/Porb \u223c h\/\u2113 = (3 \u00b1 1) \u00d7 10\u22127, corresponding to $\\Delta Q \/ Q \\sim 3\\times 10^{-5} k_2^{-1}$. The time-varying component to the gravitational quadrupole moment is therefore required to be of order a few per\u2009cent of the total expected quadrupole moment at most to explain the observed orbital period variations. From this, it seems plausible that the observed period variations can be powered by quadrupole moment changes, without requiring that a large fraction of the star be involved in the process. The required fractional quadrupole moment changes are very similar to those recently calculated for the companion to the black widow PSR J2051\u20130827 by Voisin et al. (2020b), despite the large difference in their masses.","Citation Text":["Cisneros-Parra 1970"],"Functions Text":["while if we assume that redback companions are akin to the companions in CV systems whose outer envelopes have also been stripped through accretion, then we may expect a smaller value k2 \u223c 10\u22123"],"Functions Label":["Uses"],"Citation Start End":[[1077,1096]],"Functions Start End":[[882,1075]]}
{"Identifier":"2018MNRAS.478...69A__Borucki_2016_Instance_1","Paragraph":"The complexity of non-adiabatic pulsations and their coupling to the convection has posed many problems since the field\u2019s inception and still does. The main problem lies in our, so far, limited understanding of the interaction between convection and pulsations. However, several important steps forward have already been taken, and several recent reviews on the topic exist (see for example Houdek & Dupret 2015; Samadi, Belkacem & Sonoi 2015). The case of solar pulsational stability has been studied in detail both theoretically (Balmforth 1992) and observationally (Chaplin et al. 1997; Komm, Howe & Hill 2000), while the space missions CoRoT (Baglin et al. 2006) and Kepler (Borucki et al. 2010; Borucki 2016) have provided high-quality seismic data for stars of different flavours against which we can test models and further our understanding of stellar pulsations. Appourchaux et al. (2014) analysed oscillation mode linewidths for a number of Kepler main-sequence solar-like stars and found interesting relationships between linewidths, frequencies, and effective temperatures. Using CoRoT observations Samadi et al. (2012) showed that non-adiabatic effects are present and non-negligible in red giant stars. Houdek & Gough (2002) modelled damping rates and velocity amplitudes of the red giant \u03be Hydrae, while Dupret et al. (2009) computed theoretical amplitudes, lifetimes, and heights in the frequency power spectrum of oscillation modes at different stages of red giant evolution, including the phase of core helium burning, and Grosjean et al. (2014) computed synthetic power spectra for mixed modes in red giants. Belkacem et al. (2012) were able to reproduce observed \u0393 versus Teff across the HR-diagram including both main-sequence and red giant stars. However, calculations of frequency-dependent damping rates for red giant stars have so far not been able to survive comparisons with observations. Handberg et al. (2017), hereafter referred to as H17, obtained precise frequencies and linewidths for a sample of red giants in NGC 6819 by means of extensive, careful peakbagging. Here, we compute frequency-dependent damping rates for a selection of red giant branch (RGB) stars in the H17 sample. This is done via a non-adiabatic stability calculation (Houdek et al. 1999) for which we obtain the convective fluxes from a non-local, time-dependent convection model (Gough 1977a,b) partly calibrated through 3D convection simulations (Trampedach et al. 2013).","Citation Text":["Borucki 2016"],"Functions Text":["The case of solar pulsational stability has been studied in detail","while the space missions","and Kepler","have provided high-quality seismic data for stars of different flavours against which we can test models and further our understanding of stellar pulsations."],"Functions Label":["Background","Background","Background","Background"],"Citation Start End":[[700,712]],"Functions Start End":[[445,511],[615,639],[667,677],[714,871]]}
{"Identifier":"2020AandA...635A.121M__Stolker_et_al._2016_Instance_1","Paragraph":"As scattered light imaging is sensitive to the stellar irradiation, it allows one to search for misalignments between various disk regions. While studying the morphology of the innermost disk region is challenging due to its very small radial extent, often marginally resolvable by optical interferometry (Lazareff et al. 2017), scattered light imaging of the outer disk can indirectly reveal the presence of a misaligned inner disk. In this scenario, depending on the misalignment angle, the outer disk image will show narrow shadow lanes (e.g., Pinilla et al. 2015; Stolker et al. 2016; Benisty et al. 2017; Casassus et al. 2018), broad extended shadows (Benisty et al. 2018) or low-amplitude azimuthal variations (Debes et al. 2017; Poteet et al. 2018). In some cases, studies of the CO line kinematics support a misalignment between inner and outer disk regions (Loomis et al. 2017; P\u00e9rez et al. 2018). The exact origin of such a misalignment is still unclear. In the case of T Tauri stars, if the stellar magnetic field is inclined, it can warp the innermost edge of the disk, which would then rotate at the stellar period (AA Tau; Bouvier et al. 2007). Alternatively, inner and outer disk regions can have different orientations if the primordial envelope had a different angular momentum vector orientation at the time of the inner\/outer disk formation (Bate 2018). Other scenarios involve the presence of a massive companion\/planet that is inclined with respect to the disk. If the companion is massive enough, the disk can break into two separate inner and outer disk regions, that can then precess differently and result in a significant misalignment between each other (e.g., Nixon et al. 2012; Facchini et al. 2013; Nealon et al. 2018; Zhu 2019). A clear example of such a scenario is the disk around HD 142527, in which an M-star companion was detected (Biller et al. 2012), likely on an inclined and eccentric orbit (Lacour et al. 2016; Claudi et al. 2019). Dedicated hydrodynamical simulations successfully reproduce most of the observed features in this disk (eccentric cavity, spiral arms, misaligned inner disk and shadows; Price et al. 2018).","Citation Text":["Stolker et al. 2016"],"Functions Text":["In this scenario, depending on the misalignment angle, the outer disk image will show narrow shadow lanes (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[568,587]],"Functions Start End":[[434,546]]}
{"Identifier":"2020ApJ...897...73M__Dugair_et_al._2013_Instance_1","Paragraph":"Many Be-binary systems were observed during their outbursts, which offered interesting results (Bildsten et al. 1997; Reig 2011). Bright X-ray outbursts are observed in Be binaries, most likely during the periastron passage of its neutron star through the circumstellar disk of its companion. Depending on the amount of matter released from its companion and the geometry of the binary system, rare as well as regular outbursts are observed during its binary orbit (Okazaki & Negueruela 2001; Okazaki et al. 2002; Okazaki 2016). The pulse characteristics of some of these, such as EXO 2030+375 (Parmar et al. 1989), Cepheus X-4 (Mukerjee et al. 2000), and XTE J1946+274 (Paul et al. 2001), were studied in detail during their outburst activities. Studies on pulse characteristics offer information on the pulsar geometry and the mechanism underlying its emitted pulse profile. The shape of the emitted pulse depends on modes of accretion inflows, source luminosity, geometry of accretion column, and the configuration of its magnetic field with respect to an observer\u2019s line of sight (Parmar et al. 1989). Therefore, such studies offer understanding of pulsars in binary system and disk\u2013magnetosphere interaction during the process of mass accretion, which affects its emitted radiation. Quasi-periodic oscillations (QPOs) have been detected from many Be binaries, such as A0535+262 (Finger et al. 1996; Finger 1998), EXO 2030+375 (Angelini et al. 1989), 4U 0115+63 (Soong & Swank 1989; Heindl et al. 1999; Dugair et al. 2013), and V0032+53 (Qu et al. 2005). Studies of QPOs offer rich information about the accretion torque onto the neutron star, thermodynamic properties of the inner accretion disk, and electrodynamics of the disk\u2013magnetosphere interaction of the neutron star. Details on sources with observed QPOs, their frequencies, and other features along with pulsar spin frequencies, etc. have been given in tabular form by Devasia et al. (2011), Ghosh (1998), and Mukerjee et al. (2001). Some of these transient Be-binary pulsars such as A0535+26 (50 mHz, 9.7 mHz), 1A 1118\u201361 (92 mHz, 2.5 mHz), XTE J1858+034 (110 mHz, 4.53 mHz), EXO 2030+375 (200 mHz, 24 mHz), SWIFT J1626.6\u20135156 (1000 mHz, 65 mHz), XTE J0111.2\u20137317 (1270 mHz, 32 mHz), as well as a persistent Be binary, X Per (54 mHz, 1.2 mHz), and an OB-type binary, 4U 1907+09 (69 mHz, 2.27 mHz), showed higher QPO frequency compared to their respective spin frequency as mentioned in order inside parentheses (Devasia et al. 2011). These cover an interestingly wide range of QPO frequencies between 50 and 1270 mHz for these pulsars. Studies of cyclotron absorption features, if present in the spectrum, enable us to determine the strength of the surface magnetic field of the neutron star and offer insight into the line-producing region and the structure of the accretion column and its geometry (Staubert et al. 2019). Cyclotron absorption features thus have provided an important diagnostic probe for detailed studies of neutron star binaries since their discovery in the spectrum of Her X-1 (Truemper et al. 1978). There are several reports on the detection of cyclotron absorption features in the spectrum of many Be binaries, starting at a lower energy, from \u223c10 keV (Jun et al. 2012; DeCesar et al. 2013), to a higher energy, at \u223c100 keV (La Barbera et al. 2001). A detailed compilation of such sources and studies is given in Staubert et al. (2019) and Maitra (2016). It has been observed in detailed studies that some sources show a wide variation in their cyclotron line energy with respect to their pulse phase and source luminosity, and time, such as Vela X-1, Cen X-3, and Her X-1 (Staubert et al. 2019). These interesting properties help in understanding the nature of these sources and also offer insight into their underlying physical properties governing such changes.","Citation Text":["Dugair et al. 2013"],"Functions Text":["Quasi-periodic oscillations (QPOs) have been detected from many Be binaries, such as","4U 0115+63"],"Functions Label":["Background","Background"],"Citation Start End":[[1507,1525]],"Functions Start End":[[1288,1372],[1455,1465]]}
{"Identifier":"2021ApJ...921...25C__Way_et_al._2017_Instance_1","Paragraph":"Only a few GCM studies have previously considered isolated examples of higher-order spin\u2013orbit resonance effects on climate (Wordsworth et al. 2010; Yang et al. 2013, 2020; Turbet et al. 2016; Boutle et al. 2017; Del Genio et al. 2019b). To our knowledge, no previous work has incorporated geothermal heating into a 3D GCM in the context of evaluating IHZ limits. Haqq-Misra & Kopparapu (2014) did report the impact of a 2 W m\u22122 surface heating in a highly idealized GCM for a synchronous rotation planet, while Haqq-Misra & Heller (2018) conducted idealized GCM simulations of tidally heated exomoons in synchronous rotation with the host planet. Yang et al. (2013) showed that at 2:1 and 6:1 resonances with a static\/slab ocean (see Section 2.2.2 of Way et al. 2017), Bond albedo is lower than it is for synchronous rotation and decreases rather than increases with incident stellar flux, thus destabilizing the climate as the planet approaches the IHZ. Turbet et al. (2016) considered a 3:2 resonance state and static ocean for Proxima Centauri b, assuming zero eccentricity. Boutle et al. (2017) simulated the same planet in 3:2 resonance and 0.3 eccentricity; that study uses a thin static ocean surface, which produced a longitudinal double-eyeball pattern of surface liquid water roughly coincident with the maxima in stellar heating. Wang et al. (2014) found zonally symmetric temperatures for 3:2 and 5:2 resonances with a static ocean, but Dobrovolskis (2015) showed that this was the result of an incorrect spatial pattern of instellation. Del Genio et al. (2019b) performed the first dynamic ocean simulation of a planet in a higher-order resonance, showing that despite the double maximum in instellation at 3:2 resonance, the resulting climate has a tropical liquid waterbelt spanning the planet because of the ocean thermal inertia and heat transport. Yang et al. (2020) used a dynamic ocean and focused on the outer edge of the habitable zone by considering the effect of sea ice drift on snowball transitions for nine exoplanets, including a sampling of the 3:2 resonance, also with zero eccentricity.","Citation Text":["Way et al. 2017"],"Functions Text":["Yang et al. (2013) showed that at 2:1 and 6:1 resonances with a static\/slab ocean (see Section 2.2.2 of","), Bond albedo is lower than it is for synchronous rotation and decreases rather than increases with incident stellar flux, thus destabilizing the climate as the planet approaches the IHZ."],"Functions Label":["Background","Background"],"Citation Start End":[[752,767]],"Functions Start End":[[648,751],[767,955]]}
{"Identifier":"2021MNRAS.503..354G__Cantat-Gaudin_et_al._2020_Instance_1","Paragraph":"The spatial distribution of OB stars and associations, young long-period Cepheids and open clusters, star-forming regions, H\u2009ii regions, interstellar dust, and giant molecular and neutral gas clouds in the solar vicinity that have been in existence generally \u03c4 \u2272 108 yr is known to correlate with the location of the inner Sagittarius, the closest Orion, and outer Perseus spiral arm segments. (The distances for the vast majority of these spiral tracers have been determined in the literature with trigonometric or photometric methods.) The Sun is situated at the inner edge of the Orion arm (Levine et al. 2006; Hou & Han 2014; Nakanishi & Sofue 2016; Xu et al. 2018, 2021; Lallement et al. 2019; Reid et al. 2019; Skowron et al. 2019; Cantat-Gaudin et al. 2020; Fig. 2 above).3 These three spatial features nearby to the Sun appear to form part of the global spiral structure in the Galaxy. Contrary, the objects of older population with larger random velocities, for instance, main-sequence A\u2013K stars or the oldest Cepheids and open clusters, do not currently follow the exact location of those arms (e.g. Cantat-Gaudin et al. 2020, fig. 8 therein; Griv et al. 2020, fig. 7 therein). The latter can be explained by the difference in rotation velocity between the spiral density waves and the objects. Investigating the velocity field of Xu et al.\u2019s (2018) O and early B-type stars in the framework of the Lin\u2013Shu density-wave proposal, we also found that the Sun lies within the Orion arm, at the inner edge of this spiral. The radial distance from the Sun to the centre of the Orion arm is \u22480.2 kpc in the direction of the Galactic anticentre, the centre of the Sagittarius arm is \u22481.8 kpc from the Sun in the direction of the GC, and the width of the arms is \u22480.5 kpc. The radial distance between the centres of the Orion and Sagittarius arms near the Sun is \u03bbrad \u2248 2 kpc (cf. Hou & Han 2014; Wu et al. 2014; Bovy et al. 2015). As for us, the nearest Orion spiral arm forms part of the dominant density-wave structure of the system.","Citation Text":["Cantat-Gaudin et al. 2020"],"Functions Text":["The Sun is situated at the inner edge of the Orion arm"],"Functions Label":["Background"],"Citation Start End":[[738,763]],"Functions Start End":[[538,592]]}
{"Identifier":"2018MNRAS.476.4510P__Hobbs,_Edwards_&_Manchester_2006_Instance_1","Paragraph":"The modulation of an extra-solar signal can, if working in terms of signal phase, be expressed as a time modulation, e.g. for a phase evolution given by\n(1)\r\n\\begin{eqnarray}\r\n\\phi (t) = \\phi _0 + 2\\pi f_0\\left( t - t_0 + \\Delta \\tau (t) \\right),\r\n\\end{eqnarray}\r\nwhere t is the time of arrival of the signal at the observer, and \u03d50 and f0 are an initial phase and frequency at the epoch t0 in a reference frame at rest with respect to the source, the time modulation term is \u0394\u03c4(t). Assuming, for now, that the source is at rest with respect to the SSB, the time modulation can be expressed as a combination of terms\n(2)\r\n\\begin{eqnarray}\r\n\\Delta \\tau = \\Delta _{\\rm R} + \\Delta _{\\rm E} - \\Delta _{\\rm S},\r\n\\end{eqnarray}\r\nwhere \u0394R (the Roemer delay) is a geometric retardation term, \u0394E (the Einstein delay) is a relativistic frame transformation term taking into account relativistic time dilation, and \u0394S (the Shapiro delay) is the delay due to passing through curved space\u2013time. These terms are discussed in, for example, chapter 5 of Lyne & Graham-Smith (1998), while Edwards, Hobbs & Manchester (2006) provide more detailed discussion of time delays accounting for more effects with particular relevance to pulsar observations. Here, for each of the terms we use the sign conventions given in the source code for the pulsar timing software tempo21 (Hobbs, Edwards & Manchester 2006) and in the LALSuite gravitational wave software library functions (LIGO Scientific Collaboration 2017), rather than those used in the equation of Edwards et al. (2006).2 The Roemer delay is given by\n(3)\r\n\\begin{eqnarray}\r\n\\Delta _{\\rm R} = \\frac{\\boldsymbol {r}\\cdot \\hat{\\boldsymbol {s}}}{c},\r\n\\end{eqnarray}\r\nwhere $\\boldsymbol {r}$ is a vector giving the position of the observer with respect to the SSB, and $\\hat{\\boldsymbol {s}}$ is a unit vector pointing from the observer to the source. The Einstein delay (see e.g. equations 9 and 10 of Edwards et al. 2006) converts to a new time coordinate frame, and depends on the choice of frame you want, i.e. Barycentric Coordinate Time (TCB), in which the effect of the presence of the Sun's gravitational potential is removed. The Shapiro delay (for which we will only consider the contribution from the Sun) is to first order given by\n(4)\r\n\\begin{eqnarray}\r\n\\Delta _{\\rm S} \\equiv \\Delta _{\\rm S_{\\odot }} = -\\frac{2G {\\rm M}_{\\odot }}{c^3}\\ln {\\left({\\boldsymbol {r}}_{\\rm se}\\cdot \\hat{\\boldsymbol {s}} + |\\boldsymbol {r}_{\\rm se}| \\right)}\r\n\\end{eqnarray}\r\nfor waves passing around the Sun, where $\\boldsymbol {r}_{\\rm se} = \\boldsymbol {r}_{{\\oplus}} - \\boldsymbol {r}_{\\odot }$ is the vector from the centre of the Sun to the geocentre.3 Unlike electromagnetic waves, gravitational waves will pass through matter, and therefore a different term is required for a wave passing through the Sun, i.e. when $\\left|\\boldsymbol {r}_{\\rm se}\\right|^2 - \\left(\\boldsymbol {r}_{\\rm se}\\cdot \\hat{\\boldsymbol {s}}\\right)^2 < R_{\\odot }^2$ and $\\boldsymbol {r}_{\\rm se}\\cdot \\hat{\\boldsymbol {s}} < 0$, giving\n(5)\r\n\\begin{eqnarray}\r\n\\Delta _{\\rm S_{\\odot }} &amp;=&amp; -\\frac{2G {\\rm M}_{\\odot }}{c^3}\\Bigg [\\ln {\\left(\\boldsymbol {r}_{\\rm se}\\cdot \\hat{\\boldsymbol {s}} + \\sqrt{R_{\\odot }^2 + \\left(\\boldsymbol {r}_{\\rm se}\\cdot \\hat{\\boldsymbol {s}}\\right)^2} \\right)} \\nonumber \\\\\r\n&amp;&amp;- \\,2\\left(1 - \\frac{\\sqrt{\\left|\\boldsymbol {r}_{\\rm se}\\right|^2 - \\left(\\boldsymbol {r}_{\\rm se}\\cdot \\hat{\\boldsymbol {s}}\\right)^2}}{R_{\\odot }}\\right)\\Bigg]\\!,\r\n\\end{eqnarray}\r\nwhere R\u2299 is the radius of the Sun.","Citation Text":["Hobbs, Edwards & Manchester 2006"],"Functions Text":["Here, for each of the terms we use the sign conventions given in the source code for the pulsar timing software tempo21"],"Functions Label":["Uses"],"Citation Start End":[[1355,1387]],"Functions Start End":[[1234,1353]]}
{"Identifier":"2022ApJ...937...76W__Verdini_et_al._2015_Instance_1","Paragraph":"Direct measures of cascade rates in turbulent systems often employ theoretical formulations related to Kolmogorov\u2019s \u201c4\/5\u201d law (Kolmogorov 1941b; Frisch 1995) and its variants, in which the inertial range cascade rate is related to a signed third-order structure function. This so-called exact law is derived from the fluid equations without appeal to dimensional analysis, assumptions about scaling behavior, or any ansatz concerning timescales; however, this law does require time stationarity, spatial homogeneity, the existence of an inertial range, and a finite dissipation rate. The original formulation for isotropic incompressible hydrodynamics has been extended to magnetohydrodynamics (MHD; Politano & Pouquet 1998a,1998b) and related models. The MHD version is frequently applied to in situ observations of plasma turbulence in the solar wind (Sorriso-Valvo et al. 2007; MacBride et al. 2008; Bandyopadhyay et al. 2020) to obtain cascade rates that inform theories of heating and acceleration of the solar wind (Osman et al. 2011), providing ground truth for related approximations in space physics (Vasquez et al. 2007). Frequently a major issue in these applications is the use of formulations derived assuming isotropy in turbulence that is actually anisotropic (Verdini et al. 2015), this being the typical case for solar wind and magnetosheath turbulence. Usually this potential inconsistency is disregarded in favor of extensive averaging, whenever possible. Another more practical limitation is the challenging requirement of a sufficient volume of data (Podesta et al. 2009), a kinematic and statistical issue further complicated by potential sensitivity to the tails of the probability distribution of the fluctuations (Dudok de Wit 2004). Taking these challenges into account, we note that the ability to extract cascade rates from observational data is of increasing importance due to the centrality of fundamental questions relating to heating and dissipation in space and astrophysical plasmas (e.g., Kiyani et al. 2015). Therefore, in the present study we revisit several related issues that are pertinent to the evaluation of third-order laws using single-point or multi-point measurements. We reexamine the issue of averaging by focusing on conditions for obtaining accurate results in both isotropic and anisotropic turbulence. The strategies we examine are implemented using data from three-dimensional (3D) MHD turbulence simulations. A motivation for this approach is that for such cases we have an unambiguous determination of the underlying turbulence symmetry as well as a straightforward method to quantify the absolute dissipation rate.","Citation Text":["Verdini et al. 2015"],"Functions Text":["Frequently a major issue in these applications is the use of formulations derived assuming isotropy in turbulence that is actually anisotropic",", this being the typical case for solar wind and magnetosheath turbulence."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1276,1295]],"Functions Start End":[[1132,1274],[1296,1370]]}
{"Identifier":"2017MNRAS.470.3882T__Breen_et_al._2015_Instance_1","Paragraph":"A sample of 17 starless clump candidates has been selected from the Traficante et al. (2015a) as objects with \u03a3 \u2265 0.05 g cm\u22122, mass ${{{M}}}\\ge 300$ M\u2299, bolometric luminosity over envelope mass ratio L\/M \u2264 0.3 and very low dust temperature (T\u2009\u200915 K), indicative of very young stage of evolution (see Section 4), no (or faint) emission at 70 \u03bcm after visual inspection of each source and no counterparts in the MSX and WISE catalogues in correspondence of the Herschel dust column density peak. In addition, we checked for different masers emission associated with these clumps, as they are an indication of on-going star formation activity. We searched in the methanol multibeam survey (MMB; Green et al. 2009) and found no Class II CH3OH masers in the sources of our sample (Breen et al. 2015); from the MMB survey we also searched for hydroxyl (OH) masers at 6035 MHz (Avison et al. 2016), a transition often associated with high-mass star-forming regions, and also found no associations. We searched for CH3OH and OH masers using also the Arecibo surveys of Olmi et al. (2014), which is more sensitive than the MMB survey and it is targeted to identify weak masers associated with Hi-GAL high-mass objects. We found no CH3OH masers at distances less than 100 arcsec from the source centroids. We found one source, 34.131 + 0.075, with a weak OH maser (peak emission of 20 mJy, \u22433\u03c3 above the rms of the observations made with the Arecibo telescope, Olmi et al. 2014). Finally, we checked several surveys of water masers in the first quadrant (Merello et al., in preparation, and references therein) and found that only one source, 23.271 \u2212 0.263, has a H2O maser association (at \u22433 arcsec from the source centroid), identified in the survey of Svoboda et al. (2016). Note that the source 18.787\u22120.286 is classified as starless in Traficante et al. (2015a) catalogue and has no maser associations, although it shows a 70 \u03bcm counterpart. The 70 \u03bcm source however is faint, with a peak emission of \u224360 mJy pixel\u22121 compared to a background of \u2243130 mJy pixel\u22121. The clump follows all the other selection criteria and has very low dust temperature (T = 10.6 K, see Section 4), so we include it in the analysis. The clump embedded in the cloud SDC19.281\u22120.387, which follows the same criteria but it is not in the Traficante et al. (2015a) catalogue, was also included in our selection. The final sample of 18 clumps presented here contains some of the most massive 70 \u03bcm quiet clumps observed in the Galaxy.","Citation Text":["Breen et al. 2015"],"Functions Text":["and found no Class II CH3OH masers in the sources of our sample"],"Functions Label":["Uses"],"Citation Start End":[[776,793]],"Functions Start End":[[711,774]]}
{"Identifier":"2022MNRAS.515.2633P__Blandford_&_K\u00f6nigl_1979_Instance_1","Paragraph":"Active galactic nuclei (AGN) have been discovered a century ago and still remain a hot topic of research. The theoretical predictions and the observational results suggest that the AGN have three main components namely, a central supermassive black hole (SMBH), an accretion disc around the SMBH, bipolar high relativistic jets of particles. The launching of the jets perpendicular to the accretion disc plane remains a long-standing puzzle to the astronomers. The GRMHD simulations have made some progress in understanding the launching of the jets and it is suggested to be the interplay between the accretion flow, magnetic field lines, and the SMBH (Mo\u015bcibrodzka & Falcke 2013; Mo\u015bcibrodzka et al. 2014). Simulations and theory together suggest that the magnetic field plays an important role in collimating the particles and giving them a jet like shape (Blandford & K\u00f6nigl 1979; Blandford & Payne 1982; Koide, Shibata & Kudoh 1998; Mo\u015bcibrodzka, Falcke & Shiokawa 2016). The jets are believed to be highly relativistic and the process of particle acceleration inside the jets are still unclear. The AGNs are classified in various types under the AGN unification scheme developed by Urry & Padovani (1995). According to this scheme, the observer\u2019s viewing angle with respect to jet axis divides the sources in various class such as quasars, seyfert galaxies, blazars, etc. In the case of blazars, the jet axis is within few degrees to the observer\u2019s line of sight. Due to the direct view of the jet and because of the superluminal motion, blazars are among the brightest sources in the Universe. Blazars are classified in two main types namely flat-spectrum radio quasars (FSRQs) and the BL Lacertae (BL Lac) objects depending upon the presence or absence of optical emission lines in their spectra (Stickel et al. 1991; Weymann et al. 1991). Based on the location of the synchrotron peak, the BL Lac objects are further divided into three parts known as low BL Lac (LBL), intermediate BL Lac (IBL), and high BL Lac (HBL) by Padovani & Giommi (1995).","Citation Text":["Blandford & K\u00f6nigl 1979"],"Functions Text":["Simulations and theory together suggest that the magnetic field plays an important role in collimating the particles and giving them a jet like shape"],"Functions Label":["Background"],"Citation Start End":[[860,883]],"Functions Start End":[[709,858]]}
{"Identifier":"2022MNRAS.516.5874W__Haasteren_2017_Instance_1","Paragraph":"An outlier is defined to be an anomalous event or observation that arises from a process that differs from the majority of the data generation. Outlier detection has always been an indispensable part of data analysis; contaminated data sets without proper outlier treatment can increase modelling uncertainties and produce misleading results. A common mitigation strategy is to model the data as a Gaussian mixture of inlier and outlier distributions with differing variance (Hogg, Bovy & Lang 2010; Vallisneri & van Haasteren 2017). For example, assume a model with an inlier data model, yi = \u03bci + \u03f5i, where yi is the i-th observation, \u03bci is the mean, and \u03f5i is Gaussian random noise with variance $\\sigma _i^2$. We can then account for possible outlier contamination in the data collection by modelling such fluctuations with $\\epsilon _\\mathrm{out}\\sim \\mathcal {N}(0, \\sigma _\\mathrm{out}^{2})$, where our full data model now includes a latent outlier indicator zi for each observation, where zi = 1 for an outlier and 0 otherwise. This indicator is modelled with a certain prior outlier probability, \u03b8, e.g. zi \u223c Bernoulli(\u03b8). This data model is a mixture of two Gaussians with different mean and different variances, which can be expressed in the form\n(1)$$\\begin{eqnarray}\r\ny_i = (1-z_i) \\mu _i +(1-z_i) \\epsilon _i + z_i \\epsilon _\\mathrm{out}.\r\n\\end{eqnarray}$$Incorporating the outlier indication parameter zi allows us to assess the possibility of an individual measurement being an outlier and to formulate the likelihood as a mixture of two Gaussians that respectively model the inlier and outlier distributions:\n(2)$$\\begin{eqnarray}\r\np \\left(y_{i} \\mid \\mu _{i}, z_{i}, \\sigma _{i}, \\sigma _{\\mathrm{out }}\\right)=\\left\\lbrace \\begin{array}{l{@}{\\quad}l}\\mathrm{e}^{-\\left(y_{i}- \\mu _{i}\\right)^{2} \/\\big(2 \\sigma _{i}^{2}\\big)} \/ \\sqrt{2 \\mathrm{\\pi} \\sigma _{i}^{2}}, &amp; \\text{for } z_{i}=0 \\\\\r\n\\mathrm{e}^{-y_{i}^{2} \/\\big(2 \\sigma _{\\mathrm{out }}^{2}\\big)} \/ \\sqrt{2 \\mathrm{\\pi} \\sigma _{\\mathrm{out }}^{2}}, &amp; \\text{for } z_{i}=1 .\\end{array}\\right. \\\\\r\n\\end{eqnarray}$$We use this model for the toy sine wave example in Section 2.3.","Citation Text":["Vallisneri & van Haasteren 2017"],"Functions Text":["A common mitigation strategy is to model the data as a Gaussian mixture of inlier and outlier distributions with differing variance"],"Functions Label":["Uses"],"Citation Start End":[[500,531]],"Functions Start End":[[343,474]]}
{"Identifier":"2020AandA...639A..88C__Chatzistergos_et_al._2019b_Instance_2","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"],"Functions Text":["In",", 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."],"Functions Label":["Background","Background"],"Citation Start End":[[960,968]],"Functions Start End":[[957,959],[968,1476]]}
{"Identifier":"2018AandA...616A..96R__Haardt_&_Madau_(2012)_Instance_1","Paragraph":"As expected, the final luminosity (LV) of our model dwarfs strongly correlates with the shape of their formation histories. We divide our models into three categories dependent on their LV range. In the following, we will refer to them as sustained, extended and quenched. A few representative cases of each of these three categories are shown in Fig. 9. The strength of the UV-background heating is indicated by the dotted black curve. It represents the hydrogen photo-heating rate due to the UV-background photons following the model of Haardt & Madau (2012).\n\n(a)LV > 108 L\u2299, sustained: the star formation rate of those massive and luminous dwarfs increases over 1 to 2 Gyr. This period is followed by a rather constant SFR plateau. These systems are massive enough to resist the UV-background heating and, once formed, to form stars continuously. This sustained star formation activity that lasts up to z = 0 results from the self-regulation between stellar feedback and gas cooling (Revaz et al. 2009; Revaz & Jablonka 2012).\n(b)106 L\u2299 LV  108 L\u2299, extended: in this luminosity range, the star formation is clearly affected by the UV-background. After a rapid increase, the star formation activity is damped owing to the increase of the strength of the UV-heating. However, at the exception of the h070 halo which is definitively quenched after 6.5 Gyr, the potential well of those dwarfs is sufficiently deep to avoid a complete quenching. The star formation activity extends to z = 0, however, at a much lower rate than the original one.\n(c)LV  106 L\u2299, quenched: the potential well of those galaxies is so shallow that the gas heated by the UV - photons escape the systems. Star formation is generally rapidly quenched after 2 or 3 Gyr. Only halo h064 shows signs of activity up to 4 Gyr. Those galaxies may be considered as true fossils of the re-ionization in the nomenclature of Ricotti & Gnedin (2005). They are all faint objects with only old stellar populations.\n","Citation Text":["Haardt & Madau (2012)"],"Functions Text":["The strength of the UV-background heating is indicated by the dotted black curve. It represents the hydrogen photo-heating rate due to the UV-background photons following the model of"],"Functions Label":["Uses"],"Citation Start End":[[539,560]],"Functions Start End":[[355,538]]}
{"Identifier":"2020AandA...638A..16T__Barnes_(2017)_Instance_2","Paragraph":"Figure 12 shows the results of our tidal evolution calculations. The left panel of Fig. 12 shows the planetary rotational evolution of GJ 1148 b due to star\u2013planet tides. After ~850 Myr, GJ 1148 b reaches a rotation period that is 2\u22153 of the orbital period, and remains there with Prot = 27.5 d. During the integration the planetary semi-major axis and eccentricity are mostly unaffected. An asymptotic rotation period that is shorter than synchronous and 2\/3 of the orbital period is expected for eb \u22730.24 in the constant Q tidal model (Goldreich & Peale 1966; Cheng et al. 2014). The time for GJ 1148 b to reach asymptotic rotation is inversely proportional to the initial Prot, as long as the initial Prot is much less than 27.5 d, and it depends on the other parameters of GJ 1148 b according to Eq. (3) of Barnes (2017) and Eq. (15) of Cheng et al. (2014). The rotational period of GJ 1148 b is thus very likely much longer than the orbital periods of the hypothetical exomoons, which could be dynamically stable only with orbital periods between 0.7 and 2 d. The right panel of Fig. 12 shows that the longer rotational period of GJ 1148 b (Prot = 27.5 d) leadsto strong orbital decay of the stable exomoon orbits due to tidal interactions with the planet. An exomoon eventually reaches the Roche limit where it is tidally disrupted by the gas giant. Not even one hypothetical \u201cstable\u201d exomoon in the context of Sect. 5.3.1 had survived this test. The maximum time a Mars-like exomoon could survive is ~55 M yr, while for Titan-like moons the maximum survival time is longer, ~255 M yr. The latter is longer by roughly the mass ratio of Mars to Titan, which can be understood from Eq. (2) of Barnes (2017) and Eq. (16) of Cheng et al. (2014). These timescales are optimistic since the orbital decay would start before the planet reaches the asymptotic spin state. In both cases the survival times are much shorter than the age of the system. Therefore, given the relatively fast orbital decay in the small stable region around the planet, we conclude that exomoons around GJ 1148 b are unlikely to exist.","Citation Text":["Barnes (2017)"],"Functions Text":["The latter is longer by roughly the mass ratio of Mars to Titan, which can be understood from Eq. (2) of"],"Functions Label":["Uses"],"Citation Start End":[[1697,1710]],"Functions Start End":[[1592,1696]]}
{"Identifier":"2015MNRAS.454.2691M__Leitherer_et_al._1999_Instance_1","Paragraph":"In order for simulations to play a role in improving our understanding of the formation and dynamics of the CGM, particularly given the complex, multiphase picture emerging from the latest observations (Tumlinson et al. 2011; Werk et al. 2014), the level of detail and sophistication in stellar feedback models must improve. In this work, we analyse the outflowing (and infalling) gas seen in the galaxies and CGM of the Feedback in Realistic Environments (FIRE) simulations,1 first presented in Hopkins et al. (2014). Unlike the subgrid recipes which involve kinetically prescribed decoupled winds and cooling-suppressed blastwaves, the FIRE simulations solve the \u2018overcooling\u2019 problem by explicitly modelling the radiation pressure, stellar winds, and ionizing feedback from young stars as taken directly from the population synthesis code starburst99 (Leitherer et al. 1999). These \u2018early feedback\u2019 mechanisms act before SNe, heating and stirring the surrounding ISM which is necessary to match conditions in star-forming regions such as Carina (Harper-Clark & Murray 2009) and 30 Dor (Lopez et al. 2011; Pellegrini, Baldwin & Ferland 2011). SNe are implemented by taking into account their energy and momentum input. When the cooling radius of SNe is resolved, SN energy injected is free to expand and generate momentum in the ISM before too much energy is radiated away. When this scale is poorly resolved, momentum accumulation from SN remnant evolution below the resolution scale is added to the surrounding gas. This model is physically realistic when it is applied on the scale of giant molecular clouds, meaning that a resolution of several to tens of parsecs is required. The physical feedback implementation in FIRE successfully regulates mass accumulation in galaxies and provides a physical explanation for the inefficiency of star formation in galactic discs (Kennicutt 1983, 1998; Genzel et al. 2010). We stress that we allow hydrodynamical interactions and cooling of all gas at all times, unlike in typical subgrid models. This is critical to make meaningful predictions for the phase structure of circumgalactic gas.","Citation Text":["Leitherer et al. 1999"],"Functions Text":["Unlike the subgrid recipes which involve kinetically prescribed decoupled winds and cooling-suppressed blastwaves, the FIRE simulations solve the \u2018overcooling\u2019 problem by explicitly modelling the radiation pressure, stellar winds, and ionizing feedback from young stars as taken directly from the population synthesis code starburst99"],"Functions Label":["Uses"],"Citation Start End":[[855,876]],"Functions Start End":[[519,853]]}
{"Identifier":"2021AandA...653A..36M__Goulding_&_Alexander_(2009)_Instance_4","Paragraph":"The SFG sample was constructed using the Great Observatories All-Sky LIRG Survey (GOALS sample, Armus et al. 2009), from which we extracted 158 galaxies, with data from Inami et al. (2013), who report the fine-structure lines at high resolution in the 10\u2005\u2212\u200536\u2006\u03bcm interval, and Stierwalt et al. (2014), who include the detections of the H2 molecular lines and the PAH features at low spectral resolution. For those galaxies in the GOALS sample that have a single IRAS counterpart, but more than one source detected in the emission lines, we have added together the line or feature fluxes of all components, to consistently associate the correct line or feature emission to the total IR luminosity computed from the IRAS fluxes. To also cover lower luminosity galaxies, as the GOALS sample only includes luminous IR galaxies (LIRGs) and ultra-luminous IR galaxies (ULIRGs), we included 38 galaxies from Bernard-Salas et al. (2009) and Goulding & Alexander (2009), to reach the total sample of 196 galaxies with IR line fluxes in the 5.5\u2005\u2212\u200535\u2006\u03bcm interval in which an AGN component is not detected. For the Bernard-Salas et al. (2009), Goulding & Alexander (2009), and the GOALS samples, we excluded all the composite starburst-AGN objects identified as those with a detection of [NeV] either at 14.3 or 24.3 \u03bcm. It is worth noting that the original samples from Goulding & Alexander (2009) and Bernard-Salas et al. (2009) have spectra solely covering the central region of the galaxies. To estimate the global SFR, we corrected the published line fluxes of the Spitzer spectra by multiplying them by the ratio of the continuum reported in the IRAS point source catalogue to the continuum measured on the Spitzer spectra extracted from the CASSIS database (Lebouteiller et al. 2015). We assumed here that the line emission scales (at first order) with the IR brightness distribution. In particular, we considered the continuum at 12\u2006\u03bcm for the [NeII]12.8\u2006\u03bcm and [NeIII]15.6\u2006\u03bcm lines, and the continuum at 25 \u03bcm for the [OIV]25.9\u2006\u03bcm, [FeII]26\u2006\u03bcm, [SIII]33.5\u2006\u03bcm, and [SiII]34.8\u2006\u03bcm lines. This correction was not needed for the AGN sample and the GOALS sample because of the greater average redshift of the galaxies in the 12MGS and GOALS samples. In particular, the 12MGS active galaxy sample has a mean redshift of 0.028 (Rush et al. 1993), while the GOALS sample has a mean redshift of 0.026. The galaxies presented by Bernard-Salas et al. (2009) have instead an average redshift of 0.0074, while the sample by Goulding & Alexander (2009) has an average redshift of 0.0044. For the other lines in the 10\u2005\u2212\u200536\u2006\u03bcm interval, Goulding & Alexander (2009) did not report a detection, and we used the data presented in Bernard-Salas et al. (2009) for a total of 15 objects. Both Bernard-Salas et al. (2009) and Goulding & Alexander (2009) reported data from the high-resolution Spitzer-IRS spectra. Data in the 50\u2005\u2212\u2005205\u2006\u03bcm interval were taken from D\u00edaz-Santos et al. (2017). For the GOALS sample, 20 starburst galaxies were taken from Fern\u00e1ndez-Ontiveros et al. (2016), and 23 objects were taken from the ISO-LWS observations of Negishi et al. (2001). As a result, we obtained a total sample of 193 objects. Lastly, the PAH features\u2019 fluxes were measured from the low-resolution Spitzer-IRS spectra by Brandl et al. (2006), including 12 objects from the sample of Bernard-Salas et al. (2009) and 179 objects from Stierwalt et al. (2014).","Citation Text":["Goulding & Alexander (2009)"],"Functions Text":["The galaxies presented by Bernard-Salas et al. (2009) have instead an average redshift of 0.0074, while the sample by","has an average redshift of 0.0044."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[2507,2534]],"Functions Start End":[[2389,2506],[2535,2569]]}
{"Identifier":"2019MNRAS.486.1781R__Bonning_et_al._2012_Instance_1","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"],"Functions Text":["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"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[922,941]],"Functions Start End":[[734,904]]}
{"Identifier":"2020ApJ...898L..25K__Townsley_&_Bildsten_2003_Instance_1","Paragraph":"Once He core burning finishes the core contracts and hydrogen burning starts, the radius of the hot subdwarf expands beyond its Roche radius, and mass transfer starts at an orbital period close to the observed orbital period. Mass transfer will continue for \u22481 Myr at a rate exceeding 10\u22129 M\u2299 yr\u22121 until hydrogen shell burning is finished and the star contracts to become a carbon\u2013oxygen WD with a thick helium layer and a small residual layer of hydrogen. The high accretion rate will heat the accreting WD significantly (Townsley & Bildsten 2003). Models predict a Teff \u2248 50,000 K for accretion rates of 10\u22129 M\u2299 yr\u22121 (Burdge et al. 2019) consistent with the high blackbody temperature of the accretor observed in ZTF J2055+4651. Accretion onto the WD companion at this rate will cause unstable hydrogen ignition after \u224810\u22124 M\u2299 accumulates, leading to classical novae eruptions (Nomoto 1982; Nomoto et al. 2007; Wolf et al. 2013). This accretion rate predicts a recurrence time of order 105 yr for a total of approximately 10 novae. Our binary model suggests that this system is within the first \u224810% of the 1 Myr accretion phase. At the current state the orbit will shrink with \n\n\n\n\n\n s s\u22121, which will be detectable after a few years of monitoring. The right panel of Figure 5 shows the evolution of the donor through this mass transfer phase. After accretion has ceased, the orbit of the system will continue to shrink due to the radiation of gravitational waves and the system will merge in \u224830 Myr. Our models predict that there is a substantial He layer of \u22480.05 M\u2299 left in the former hot subdwarf and the total mass of the system is relatively high (Mtotal \u2248 1.1 M\u2299). Recent models predict that such a system explodes as a subluminous thermonuclear supernova (Perets et al. 2019; Zenati et al. 2019). If the system avoids a thermonuclear supernova it will merge and could evolve into a rapidly rotating single high-mass carbon\u2013oxygen WD (Saio 2008; Clayton 2012; Schwab 2019).","Citation Text":["Townsley & Bildsten 2003"],"Functions Text":["Mass transfer will continue for \u22481 Myr at a rate exceeding 10\u22129 M\u2299 yr\u22121 until hydrogen shell burning is finished and the star contracts to become a carbon\u2013oxygen WD with a thick helium layer and a small residual layer of hydrogen. The high accretion rate will heat the accreting WD significantly"],"Functions Label":["Background"],"Citation Start End":[[523,547]],"Functions Start End":[[226,521]]}
{"Identifier":"2018ApJ...866L...1S__Pecharrom\u00e1n_et_al._1999_Instance_1","Paragraph":"It was found that the complex dielectric function from Pecharrom\u00e1n et al. (1999) for the sample obtained by heating bayerite at 1273 K, assuming a spheroid with depolarization parameters of (0.35, 0.003), produced an opacity with 11, 20, 28, and 32 \u03bcm features, so this component was included in the models. However, with only this component, the observed 20 \u03bcm features in the residual spectra were found to be wider than those in the models. By adding the opacity of the sample obtained by heating boehmite at 1173 K, the width of the 20 \u03bcm feature could be matched. This was done using the complex dielectric function for the sample obtained by heating boehmite at 1173 K from Pecharrom\u00e1n et al. (1999), assuming a spheroid with depolarization parameters of (0.35, 0.035). The complex dielectric functions of the samples obtained by heating bayerite and boehmite to various temperatures (Pecharrom\u00e1n et al. 1999) were derived by modeling the reflectance spectra of pellets obtained by pressing powders of these materials under great pressure. This method required Pecharrom\u00e1n et al. (1999) to assume an effective medium theory, such that a pellet is a mixture of one of their samples with a matrix of air. Pecharrom\u00e1n et al. (1999) noted that heating bayerite at 500\u00b0C eliminates the XRD pattern of bayerite, and they note that at 700\u00b0C, the infrared reflectance spectrum of the boehmite sample no longer shows OH\u2212 stretching bands. This must mean that the samples obtained from heating bayerite at 1273 K and from heating boehmite at 1173 K are no longer bayerite or boehmite, respectively. XRD performed by Pecharrom\u00e1n et al. (1999) of the sample of bayerite prepared at 1273 K suggests only \u03b8-alumina was present, and their infrared and NMR spectroscopy confirms this. XRD of their sample obtained from heating boehmite to 1173 K (Pecharrom\u00e1n et al. 1999) suggests \u03b4-alumina to be present, though some amounts of \u03b8-alumina and \u03b1-alumina are present, as they deduce from XRD and infrared and NMR spectroscopy.","Citation Text":["Pecharrom\u00e1n et al. (1999)"],"Functions Text":["It was found that the complex dielectric function from","for the sample obtained by heating bayerite at 1273 K, assuming a spheroid with depolarization parameters of (0.35, 0.003), produced an opacity with 11, 20, 28, and 32 \u03bcm features, so this component was included in the models."],"Functions Label":["Uses","Uses"],"Citation Start End":[[55,80]],"Functions Start End":[[0,54],[81,307]]}
{"Identifier":"2018MNRAS.479.3254V__Schneider_et_al._2018_Instance_1","Paragraph":"The lifetime of molecular clouds (MCs) remains an active research topic in the study of the interstellar medium and star formation, and most recent studies, both observational and theoretical, place this lifetime at a few times 107 yr for clouds in the 105\u2013106M\u2299 mass range (e.g. Blitz & Shu 1980; Kawamura et al. 2009; Zamora-Avil\u00e9s, V\u00e1zquez-Semadeni & Col\u00edn 2012; Zamora-Avil\u00e9s & V\u00e1zquez-Semadeni 2014; Lee, Miville-Desch\u00eanes & Murray 2016). In addition, several observational studies have suggested that the star formation rate (SFR) of the clouds appears to increase over their lifetimes. For example, studies of young clusters embedded in moderate-mass MCs (\u223c104M\u2299) (e.g. Palla & Stahler 1999, 2000; Da Rio et al. 2010) have shown that their age histograms contain a large majority of young (1\u20132 Myr) objects, but also a tail of older (up to several Myr) ones suggesting an accelerating star formation activity, sometimes followed by a subsequent decline (see also Povich et al. 2016; Schneider et al. 2018). In addition, Kawamura et al. (2009) reported a clear evolutionary process over the lifetime of giant molecular clouds (GMCs; of masses \u223c105\u2013106M\u2299) in the Large Magellanic Cloud, evidenced by the increasing number of massive stars across the sequence of GMC \u2018classes\u2019 proposed by those authors. Finally, on the basis of the large scatter in the observed star formation efficiency (SFE) in Milky Way GMCs, Lee et al. (2016) have concluded that the SFR in those clouds must also be time variable. Numerical simulations of MC formation and evolution also exhibit time varying, increasing SFRs during their early stages (e.g. V\u00e1zquez-Semadeni et al. 2007; Hartmann, Ballesteros-Paredes & Heitsch 2012). Also, in the presence of stellar feedback, at late times the SFRs reach a maximum and begin to decrease again (e.g. V\u00e1zquez-Semadeni et al. 2010; Col\u00edn, V\u00e1zquez-Semadeni & G\u00f3mez 2013). V\u00e1zquez-Semadeni, Gonz\u00e1lez-Samaniego & Col\u00edn (2017) have recently shown that the simulations of Col\u00edn et al. (2013) in fact produce stellar age histograms highly resemblant of the observed ones (Palla & Stahler 1999, 2000; Da Rio et al. 2010), and reproduce observed radial age gradients in clusters (Getman et al. 2014) as well as bottom-heavy stellar initial mass functions (IMFs) in scattered regions of massive star formation (Povich et al. 2016).","Citation Text":["Schneider et al. 2018"],"Functions Text":["For example, studies of young clusters embedded in moderate-mass MCs (\u223c104M\u2299)","have shown that their age histograms contain a large majority of young (1\u20132 Myr) objects, but also a tail of older (up to several Myr) ones suggesting an accelerating star formation activity, sometimes followed by a subsequent decline (see also"],"Functions Label":["Background","Background"],"Citation Start End":[[990,1011]],"Functions Start End":[[593,670],[725,969]]}
{"Identifier":"2019MNRAS.483..971V__al_2002_Instance_1","Paragraph":"We have also considered early radio emission from gamma-ray bursts (GRBs), which can have higher brightness temperatures at early times than blazars, owing to their ultrarelativistic velocities. They can therefore be brighter and easier to measure while still at small angular sizes, and are consequently observed to show interstellar scintillation in their first days at ${\\sim } 5\\, \\mbox{GHz}$ (Granot & van der Horst 2014). Before deceleration to Lorentz factor \u0393  1\/\u03b8j (before the \u2018jet break\u2019 for a jet of opening half-angle \u03b8j), the projected source angular size \u03b8 at (earth) time T after explosion of a GRB at redshift $z$ is \u03b8 \u223c 2cT\u2009\u0393\/DM($z$), where DM($z$) = DA($z$)(1 + $z$) is the proper motion distance, and DA($z$) the angular diameter distance. The Blandford\u2013McKee blast wave of the ultrarelativistic shock moving into a medium of uniform external density \u03c10 has radius R \u2243 2cT\u2009\u03932\/(1 + $z$) and explosion energy per unit solid angle E\/\u03a9 \u2243 \u03c10R3\u03932c2, which gives \u0393 \u2243 9(Eiso, 53\/n0)1\/8(T\/[(1 + $z$)day])\u22123\/8, where $E=10^{53}\\, \\mbox{erg}(\\Omega \/4\\pi)E_{iso,53}$ and n0 is the external density in cm\u22123 (Granot et al 2002, cf.). At DM($z$ = 1) = 3.3 Gpc, $\\theta = [0.2,\\, 1,\\, 4]\\, \\mu \\mbox{as}$ at $T=[0.1,\\, 1,\\, 10]\\, \\mbox{d}$. Thus at $\\lambda \\lt 4\\, \\mbox{cm}$ (the transition wavelength below which Milky Way scintillation becomes unimportant), the GRB will be smaller than our fiducial scattering angle $\\theta =20\\,\\mu \\mbox{as}(\\lambda \/30\\, \\mbox{cm})^{11\/5}\\lt 0.25\\,\\mu \\mbox{as}$ for less than 0.1 d. During this time, the scintillation time-scale will be set by the rapidly expanding source, expanding across the refractive screen at a projected speed of \u223c\u0393cDl\/Ds. This is many times c for our cosmological lenses with Dl \u223c 0.5Ds (but less than $1\\, \\mbox{km s}^{-1}$ for Milky Way interstellar plasma at $D_l\\sim 100\\, \\mbox{pc}$, so Milky Way scintillation time-scales are dominated by gas motions in the Milky Way, not the apparent source expansion). The refractive scintillation time-scale is thus the same as the time-scale for the source to expand to a size larger than the refractive scale \u2013 i.e. the source will have only about 1 speckle before becoming too large to display refractive scintillation. This would be difficult to convincingly detect in a GRB.","Citation Text":["Granot et al 2002"],"Functions Text":["The Blandford\u2013McKee blast wave of the ultrarelativistic shock moving into a medium of uniform external density \u03c10 has radius R \u2243 2cT\u2009\u03932\/(1 + $z$) and explosion energy per unit solid angle E\/\u03a9 \u2243 \u03c10R3\u03932c2, which gives \u0393 \u2243 9(Eiso, 53\/n0)1\/8(T\/[(1 + $z$)day])\u22123\/8, where $E=10^{53}\\, \\mbox{erg}(\\Omega \/4\\pi)E_{iso,53}$ and n0 is the external density in cm\u22123"],"Functions Label":["Uses"],"Citation Start End":[[1115,1132]],"Functions Start End":[[759,1113]]}
{"Identifier":"2016MNRAS.463L..26B__Ackermann_et_al._2014_Instance_1","Paragraph":"We calculate the corresponding \u03b3-ray spectrum constructed from the photons arriving at the observer within the observation time of MAGIC and plot it in the bottom panel of Fig. 5. Although it is possible to explain the very fast variability of the emission even with a moderate Lorentz factor of the blob, strong constraints are put by the level of the observed flux. If the emission region has a radius of Rb = 3 \u00d7 1014 cm, and is moving with a Lorentz factor of \u03b3b = 100, we require an energy density of \u03c1E = 20\u2009erg\u2009cm\u22123 (measured in the blob's frame of reference) to reproduce the flux observed by MAGIC and Fermi-LAT (see the bottom panel of Fig. 5). Note that such large values of the Lorentz factor of the emission region in the jet have been already postulated in terms of other models in order to explain extremely short flares observed in this source (Ackermann et al. 2014) or in the other sources, e.g. PKS 2155\u2212304 (Aharonian et al. 2007). Such large Lorentz factors of the blob find also some observational support from the observations of the superluminal motion in PKS 1510\u2212089 which represents similar type of blazar (Jorstad et al. 2005). We can estimate the power of the blob in the observer's reference frame on $L_{\\rm blob} = \\pi R_{\\rm b}^2c\\rho _E \\gamma _{\\rm b}^2\\approx 1.7\\times 10^{45}$ erg s\u22121. On the other hand, the Eddington luminosity of the black hole in PKS 1222+21, with the mass 6\u20138 \u00d7 108 M\u2299, is LEdd = 1.3 \u00d7 1047M9 \u2248 8\u201310 \u00d7 1046 erg s\u22121. Therefore, the blob has to contain about \u223c2\u2009per\u2009cent of the Eddington power. This is quite demanding but seems not to be excluded, especially if Rb \u2248 R\u22a5. Lower values of \u03b3b require a much higher energy density in the blob (e.g. \u03c1E = 340\u2009erg\u2009cm\u22123 for \u03b3b = 50 and Rb = 3 \u00d7 1014 cm). The strong dependence on the \u03b3b is a combined result of the transformation of the energy density to the reference frame of the observer and the beaming of the emission in a narrower cone. Note that a larger radius of the blob will lower the energy density constraint, e.g. for Rb = 1015 cm we obtain \u03c1E = 4.9\u2009erg\u2009cm\u22123 for \u03b3b = 100 and \u03c1E = 80\u2009erg\u2009cm\u22123 for \u03b3b = 50, at the assumption that there is no competing energy loss process of the electrons at such a large distance from the star.","Citation Text":["Ackermann et al. 2014"],"Functions Text":["Note that such large values of the Lorentz factor of the emission region in the jet have been already postulated in terms of other models in order to explain extremely short flares observed in this source"],"Functions Label":["Similarities"],"Citation Start End":[[861,882]],"Functions Start End":[[655,859]]}
{"Identifier":"2015ApJ...799...55G__Klassen_et_al._2000_Instance_1","Paragraph":"While the angular extent of IP shocks can be directly investigated using multi-point in situ measurements, the size of coronal shocks can only be indirectly inferred via remote-sensing observations of the electromagnetic emissions associated with them. According to Nelson & Robinson (1975), the average angle subtended at the solar surface by fundamental metric type II radio emission sources is 43\u00c2\u00b0. Aurass et al. (1994) found particular cases with larger, double type II source structures covering a separation angle beyond 90\u00c2\u00b0. Type II radio sources often show non-radial propagation trajectories (see Mann et al. 2003, and references therein). Wave-like large-scale disturbances propagating over the solar disk in extreme ultraviolet observations (usually referred to as \u00e2\u0080\u009cEIT waves\u00e2\u0080\u009d or \u00e2\u0080\u009cEUV waves\u00e2\u0080\u009d) are in close empirical correlation with type II radio bursts (Klassen et al. 2000). Most EIT waves are accompanied by CMEs, and observations and MHD modeling suggest that they are driven by the lateral expansion of CMEs, while the ultimate nature of the phenomenon remains under debate and could consist of true waves, pseudo waves (e.g., reconnection fronts), or a combination of both (Patsourakos & Vourlidas 2012; Nitta et al. 2013b, and references therein). According to Patsourakos & Vourlidas (2012), EIT waves can reach distances up to 1.3 R (850 Mm) from the source. Single-case studies reported some EIT waves covering a whole solar hemisphere (Klassen et al. 2000; Kienreich et al. 2009; Thompson & Myers 2009). Connections between EIT waves and SEP events have been often suggested (e.g., Bothmer et al. 1997; Krucker et al. 1999), and recently Rouillard et al. (2012) hypothesized that the EIT wave can be used to track the expansion of a coronal shock responsible for particle acceleration. Other authors question the EIT wave acceleration scenario for SEPs, with many EIT waves being observed at well-connected positions having no associated SEP increase (Miteva et al. 2014).","Citation Text":["Klassen et al. 2000"],"Functions Text":["Type II radio sources often show non-radial propagation trajectories (see Mann et al. 2003, and references therein). Wave-like large-scale disturbances propagating over the solar disk in extreme ultraviolet observations (usually referred to as \u00e2\u0080\u009cEIT waves\u00e2\u0080\u009d or \u00e2\u0080\u009cEUV waves\u00e2\u0080\u009d) are in close empirical correlation with type II radio bursts"],"Functions Label":["Background"],"Citation Start End":[[876,895]],"Functions Start End":[[534,874]]}
{"Identifier":"2015AandA...582A.104R__Stix_2002_Instance_1","Paragraph":"Thermal motions of atoms produce a Doppler broadening of spectral lines with a Gaussian profile. Other unresolved velocities of a random nature are usually described as a turbulence broadening with a Gaussian or Lorentzian profile (see Rutten 2003, for a detailed description). The instrumental broadening encompasses the broadening caused by the finite spectral resolution of the instrument and is commonly approximated by a Gaussian. When other line broadening mechanisms are negligible, the total line broadening is the convolution of the line broadening profiles for the thermal and turbulence motions as well as the instrumental profile (Sect. 10.5 in  B\u00f6hm-Vitense 1989). In other words, the velocity equivalent of the observed line width, Wobs = c \u00d7 \u0394\u03bb\/\u03bb, at 1\/e of the peak intensity results from instrumental, Doppler, and turbulence (non-thermal) broadenings by (1)\\begin{equation} W_{\\rm{obs}}^2 = \\left(c\\times \\frac{\\Delta\\lambda}{\\lambda}\\right)^2 = W_{\\rm{instrumental}}^2 + W_{\\rm{Doppler}}^2 + W_{\\rm{turbulence}}^2, \\end{equation}Wobs2=c\u00d7\u0394\u03bb\u03bb2=Winstrumental2+WDoppler2+Wturbulence2,assuming that all three line broadening components have a Gaussian profile (Eq. (4.17) of  Stix 2002). The POLIS instrumental width is less than 2 km\u2009s-1 (Beck 2006; Rezaei 2008; Beck et al. 2013a), while that of the Echelle spectrograph is of the same order. Assuming a generic chromospheric temperature of 104 K, we estimate a turbulence velocity using Eq. (2), (2)\\begin{equation} \\label{eq:one}  \\Delta\\lambda=\\frac{\\lambda}{c}\\sqrt{2\\,k_{\\rm B}\\,T\/m+ W_{\\rm{turbulence}}^2\\,\\,+\\,\\,W_{\\rm{instrumental}}^2 }, \\end{equation}\u0394\u03bb=\u03bbc2\u2009kB\u2009T\/m+Wturbulence2\u2009\u2009+\u2009\u2009Winstrumental2,where \u0394\u03bb is the observed line width, m is mass of the calcium atom, c the speed of light, and kB is the Boltzmann constant (Tandberg-Hanssen 1960). Using the measured width of Ca\u2009ii\u2009H & IR lines of 1.0 and 0.7 \u00c5, respectively (Sect. 4), we estimate a turbulence velocity of about 45 km\u2009s-1 for Ca\u2009ii\u2009H and 24 km\u2009s-1 for Ca\u2009ii\u2009IR lines (we also note that the H1 minima of the EB profile is about 1 \u00c5\u2009\u2009 wider than in the quiet Sun profile). The width of the emission peaks on either side of the H\u03b1 line (1 \u00c5) amounts to a turbulence velocity of 15 km\u2009s-1. Attributing the line width to the temperature and the instrumental profiles (omitting the turbulence broadening), we get a temperature of about 5 \u00d7  105 K, which is too hot for chromospheric heights. An increased turbulence velocity as a function of height in the chromosphere is part of standard models, either in the quiet Sun or sunspots (Kneer & Mattig 1978; Vernazza et al. 1981; Lites & Skumanich 1982). Our measured values, however, are larger than the turbulence velocity in an average atmosphere. The estimated turbulence velocity changes from one profile to another, but the general result remains the same: the observed width of the emission peaks is far in excess of any instrumental or thermal Doppler profile and to the first order has to have a turbulent nature. The Doppler width of a calcium line at 1\u20132 \u00d7 104 K is only about 2\u20133 km\u2009s-1, far from the observed widths of >20\u2009 km\u2009s-1. ","Citation Text":["Stix 2002"],"Functions Text":["In other words, the velocity equivalent of the observed line width, Wobs = c \u00d7 \u0394\u03bb\/\u03bb, at 1\/e of the peak intensity results from instrumental, Doppler, and turbulence (non-thermal) broadenings by (1)\\begin{equation} W_{\\rm{obs}}^2 = \\left(c\\times \\frac{\\Delta\\lambda}{\\lambda}\\right)^2 = W_{\\rm{instrumental}}^2 + W_{\\rm{Doppler}}^2 + W_{\\rm{turbulence}}^2, \\end{equation}Wobs2=c\u00d7\u0394\u03bb\u03bb2=Winstrumental2+WDoppler2+Wturbulence2,assuming that all three line broadening components have a Gaussian profile (Eq. (4.17) of"],"Functions Label":["Uses"],"Citation Start End":[[1190,1199]],"Functions Start End":[[678,1188]]}
{"Identifier":"2017AandA...605A..20C__Lattanzi_et_al._(2015)_Instance_1","Paragraph":"As mentioned above, molecular oxygen was used to calibrate the magnetic field applied. A total of 155 Zeeman components (for the three transitions considered) was measured, with the magnetic field varied from B = 2.3 G (Itot = 0.2\u2009Amp) to B = 113.5\u2009G (Itot = 10 Amp). Figure 5 shows the Zeeman spectrum for the N,J = 3, 2 \u2190 1, 2 transition when a magnetic field of 5.7 G is applied. The fit was carried out with the program described in the previous section, with the spectroscopic parameters and g factors fixed at the values of Yu et al. (2012) and Christensen & Veseth (1978), Evenson & Mizushima (1972), respectively; the only free parameter was the correction factor to be applied to the theoretical magnetic field (see Eq. (1)). The fit reproduces in a satisfactory manner the measured Zeeman components, with a standard deviation of 73 kHz (the uncertainty for the measured frequencies was in most cases set to 70 kHz). Moving to SO, for the seven transitions considered, a total of 353 Zeeman components were measured, with the magnetic field varied from B = 5.7 Gauss (Itot = 0.5 Amp) to B = 124.8 Gauss (Itot = 11 Amp), and fitted as described above. In the fitting procedure the spectroscopic parameters (i.e., the rotational, centrifugal distortion, and fine structure constants) were kept fixed at the values derived by Lattanzi et al. (2015). The g factors resulting from the fit are given in Table 4 together with the best-estimated values discussed above, while the complete set of the measured Zeeman components is available in the Supplementary Material. Alternative fits were carried out: In the first fit, the three g factors, gs, gl, and gr, were fitted. In a second step, three different fits were carried out by fixing one of the three g factors and fitting the other two. In the last fit, only gs was determined. We note that in the first fit, based on the comparison with theory, the gs value is overestimated and the gl value is underestimated; the two terms therefore seem to be correlated. For this reason, we performed the additional fits described above. We note that in all cases, gr is determined with a limited accuracy, that is, with a relative uncertainty of ~20%. We also note that, if we fix gs at the best computed value, a gl value in good agreement with theory is obtained. However, by fixing gl at the best estimate, the resulting gs is still slightly overestimated. Simulations using values of gs in the 2.0020\u22122.0030 range show that the Zeeman splittings only change by a few tens of kHz, that is, in most cases within the typical uncertainty affecting the frequency measurements. The last comment concerns the standard deviation of the fits that, in all cases, is about 50 kHz. ","Citation Text":["Lattanzi et al. (2015)"],"Functions Text":["In the fitting procedure the spectroscopic parameters (i.e., the rotational, centrifugal distortion, and fine structure constants) were kept fixed at the values derived by"],"Functions Label":["Uses"],"Citation Start End":[[1333,1355]],"Functions Start End":[[1161,1332]]}
{"Identifier":"2021ApJ...915...93L__Magdziarz_&_Zdziarski_1995_Instance_1","Paragraph":"Table 3 shows the best-fit model parameters for the baseline fit for the three observations. We note that we did not find any statistically significant neutral or ionized absorption intrinsic to the source. On the addition of an ionized absorption model generated using the latest atomic data with the photoionization modeling code CLOUDY (Ferland et al. 2013; Laha et al. 2016c, 2016b, 2017, 2019b), we did not detect any improvement in the fit, as also found in previous works on this source (Rivers et al. 2011). Two blackbody components were necessary to describe the soft excess. We detected a narrow FeK\u03b1 emission line in the Suzaku observation at an energy E = 6.41 \u00b1 0.01  keV, while a marginally broad line (\u03c3 = 0.35 \u00b1 0.08  keV) was required in the XMM-Newton and NuSTAR observations. We detected a neutral reflection hump at energy E > 10  keV with Suzaku that was modeled with pexrav, with an improvement  \u0394\u03c72 = 32 for one additional degree of freedom. The pexrav model assumes an exponentially cutoff power-law spectrum reflected from neutral material (Magdziarz & Zdziarski 1995). The output spectrum is the sum of the cutoff power law and the reflection component. However, the reflection component alone can be obtained by setting the relative reflection value negative. The power-law cutoff energy for pexrav is assumed to be 300  keV, the abundance of the reprocessor set to solar value, and we allowed the inclination angle of this model to vary between 0\u00b0 and 45\u00b0 (being a Seyfert 1 galaxy). We have tied the power-law photon index normalization of pexrav to that of the primary power-law component. The pexrav reflection fraction is 0.17 \u00b1 0.03. Note that throughout the text, we refer to the positive value of the reflection fraction. We did not detect any significant neutral reflection component in the XMM-Newton+NuSTAR spectra. Table 2 shows the fluxes of the soft excess, power law, FeK\u03b1 emission line, and hard X-ray excess. In XSPEC notation, the best-fit baseline model is written as constant \u00d7 tbabs \u00d7 (powerlaw+bbody+pexrav+Gaussian(s)). We note that we required an emission line at 0.56  keV for XMM-Newton observation, which corresponds to O vii emission and is found in several bright AGNs, and we did not need that model for Suzaku and NuSTAR, as they do not cover that energy range.","Citation Text":["Magdziarz & Zdziarski 1995"],"Functions Text":["The pexrav model assumes an exponentially cutoff power-law spectrum reflected from neutral material"],"Functions Label":["Uses"],"Citation Start End":[[1066,1092]],"Functions Start End":[[965,1064]]}
{"Identifier":"2015MNRAS.448..629G__Cenko_et_al._2011_Instance_1","Paragraph":"The remaining six parameters are less constrained. They are \u03f5e, \u03f5B, n0, L (where L = \u03b7L0), \u03ba and Ek. We apply constraints to the range of allowed values for these parameters. \u03f5B has been found to be as low as 10\u22128 (Barniol Duran 2014; Santana, Barniol Duran & Kumar 2014) and as a fraction can be as high as 1. In practice, \u03f5e tends towards higher values than \u03f5B. We set an upper limit of 1, noting that \u03f5e actually refers to the electron population that is emitting synchrotron radiation, rather than the electron population as a whole, and set a lower limit of 10\u22123 (Kumar 2000). n0 is limited between 10\u22125 and 100\u2009cm\u22123, in line with what has been found in these sources (Cenko et al. 2011). The upper limit of L is set by the argument in equation (5), and values of this parameter below \u223c1047\u2009erg\u2009s\u22121 are never energetic enough to match the data, so we set the lower limit as 1047\u2009erg\u2009s\u22121. Within these limits for L, we find that EE ceases to have any influence on the light curve if \u03ba \u2272 10\u22122. If EE is isotropic, and the observed luminosity is only 1 per cent of the true energy (i.e. the conversion efficiency of kinetic to potential energy in the internal shocks is 1 per cent), then the energy delivered to the synchrotron shock front could be up to 100 times higher than observed in the light curve. In practice, however, the emission is (a) unlikely to be fully isotropic, (b) likely to shock more efficiently than 1 per cent, and (c) certain to be less than 100 per cent efficient at delivering its energy to the synchrotron shock front. For these reasons, we set the upper limit of \u03ba at a still fairly generous factor of 10. Finally, we limit the energy in the shock from prompt emission to 1048\u2009erg  Ek  1052\u2009erg. The arguments for these limits are identical to those used for \u03ba, except that the prompt emission is known to be beamed (Sari, Piran & Halpern 1999; Frail et al. 2001) so the upper limit is lower, and because the injected energy at early times is negligible, Ek dominates the early light curve so the lower limit can be much less energetic before its influence vanishes. These limits are summarized in Table 3.","Citation Text":["Cenko et al. 2011"],"Functions Text":["n0 is limited between 10\u22125 and 100\u2009cm\u22123, in line with what has been found in these sources"],"Functions Label":["Uses"],"Citation Start End":[[674,691]],"Functions Start End":[[582,672]]}
{"Identifier":"2016ApJ...827..151Y__Couvidat_et_al._2015_Instance_1","Paragraph":"The data produced by the HMI instrument are of high quality; however, there are known uncertainties, limitations and systematic errors present that affect the measurement of magnetic flux. These include a sinusoidal variation in the total magnetic flux with a periodicity of 12 and 24 hr, due to a Doppler shift present in the Fe spectral line (Hoeksema et al. 2014). The main contribution to this shift is the geosynchronous orbit of SDO. The daily variation of \u00b13 km s\u22121 in spacecraft orbital velocity causes a sinusoidal variation in the total flux measured. This affects weak and strong magnetic fields in the LoS magnetograms differently, with the daily variation remaining less than 30 G for field strengths below 1000 G and less than 75 G for field strengths below 2250 G. On average during a day this is roughly \u00b135 G (Couvidat et al. 2015). Nevertheless, the strength of these instrumental and observational effects does not account for the strong flux cancellation we observe in AR 11226 prior to eruption. The leading edge of the positive polarity of AR 11226 reaches \u223c60\u00b0 from central meridian on June 6. Due to the spatially dependent sensitivity of HMI the noise level increases as a function of the center-to-limb angle and the spacecraft\u2019s orbital velocity. This increases the value of pixels in low and moderate fields (between 250 and 750 G) by a few tens of percent and manifests itself as broad peaks that are centered at \u223c\u00b160\u00b0 in the magnetic flux. This could be responsible for the increase in flux seen in Figure 5 at this time. However, the increase also coincides with the emergence of the two anti-Hale bipoles and so it is difficult to disentangle these effects. The proximity of AR 11226 to the limb on June 6 means that the magnetic flux was not measured right up until the time of the eruption but only until approximately a day beforehand. Therefore, these results provide a lower limit to the total flux canceled in the lead-up to the eruption. Furthermore, due to the fact that cancellation persists for several days, we can extrapolate that, if the cancellation process continued at the same rate (2.0 \u00d7 1019  Mx hr\u22121) as we observed in the period from June 1 09:00 UT to June 5 00:00 UT up until eruption, then an extra 1.1 \u00d7 1021 Mx could have been canceled. This would therefore result in a total amount of 2.8 \u00d7 1021 Mx flux canceled between two consecutive CMEs.","Citation Text":["Couvidat et al. 2015"],"Functions Text":["This affects weak and strong magnetic fields in the LoS magnetograms differently, with the daily variation remaining less than 30 G for field strengths below 1000 G and less than 75 G for field strengths below 2250 G. On average during a day this is roughly \u00b135 G","Nevertheless, the strength of these instrumental and observational effects does not account for the strong flux cancellation we observe in AR 11226 prior to eruption."],"Functions Label":["Uses","Compare\/Contrast"],"Citation Start End":[[827,847]],"Functions Start End":[[562,825],[850,1016]]}
{"Identifier":"2020AandA...644A..59K__Heays_et_al._2017_Instance_3","Paragraph":"Analyzing optical emission lines, emanating from within the northern lobe, Tylenda et al. (2019) found a reddening with EB\u2005\u2212\u2005V\u2004\u2248\u20040.9 mag or AV\u2004\u2248\u20042.8 mag, which we assume is mainly circumstellar in origin. Hajduk et al. (2013) observed two stars shining through the southern lobe and found AV\u2004=\u20043.3\u2005\u2212\u20054.4 mag with unknown contribution from the interstellar component. We assume here that those observations quantify the amount of circumstellar dust that is the main actor in shielding molecules from the central source. We recalculated the lifetimes of molecules assuming AV\u2004=\u20043 mag, and with (1) standard dust properties (i.e. with composition and size distribution as of interstellar dust) or (2) with larger and less opaque grains, at the gas-to-dust mass ratio of 124 (see Heays et al. 2017, for more details on the assumed dust properties). We used shielding functions from Heays et al., which include effects in lines. Results are shown in Cols. (3) and (4) of Table 3. The presence of ISM grains makes it possible for the observed molecules to survive for a very long time, longer than 350 yr. The lifetimes in the presence of the large grains considered by Heays et al. are typically a few times shorter than the age of the remnant. It is uncertain what kind of grains populate the dusty remnant of CK Vul, but given its anomalous elemental and molecular compositions and eruptive history, dust may have a peculiar chemical composition and size distribution. In such a case, the total to selective extinction law would also be different and the assumed AV may not be adequate. Nevertheless, if the molecules formed 350 yr ago and are shielded by big grains, with the calculated lifetimes a considerable fraction of molecular species would survive, except perhaps for a few most fragile ones which indeed are almost absent in the lobes. We conclude that the lifetimes in Table 3 that were calculated with an attenuated ISRF are consistent with the molecule formation 350 yr ago or more recently.","Citation Text":["Heays et al."],"Functions Text":["The presence of ISM grains makes it possible for the observed molecules to survive for a very long time, longer than 350 yr. The lifetimes in the presence of the large grains considered by","are typically a few times shorter than the age of the remnant."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1164,1176]],"Functions Start End":[[975,1163],[1177,1239]]}
{"Identifier":"2016MNRAS.457..212S__Ijjas,_Steinhardt_&_Loeb_2013_Instance_1","Paragraph":"One of the first and simplest proposed Friedmann\u2013Robertson\u2013Walker (FRW) cosmological model is the \u039b cold dark matter (\u039bCDM) universe, which involves Einstein's cosmological constant \u039b. This standard model of cosmology, which is also referred to as the concordance model, assumes that the total energy density \u03c1 of the universe is made up of three components, namely matter \u03c1m (baryonic and dark matter), radiation \u03c1r, and dark energy or vacuum energy \u03c1\u039b, which produces the necessary gravitational repulsion. In this model, dark energy which has an equation of state (EOS) \u03c9\u039b = p\u039b\/\u03c1\u039b = \u22121, is a property of the space itself and its density \u03c1\u039b = \u2212p\u039b = \u039bc4\/8\u03c0G is constant, such that as the universe expands the constant vacuum energy density will eventually exceed the matter density of the universe which is ever decreasing. The spatially flat \u039bCDM model dominated by vacuum energy with \u03a9\u039b \u223c 0.70, with the rest of the energy density being in the form of non-relativistic cold dark matter with \u03a9m \u223c 0.25 and non-relativistic baryonic matter with \u03a9b \u223c 0.05, fits observational data reasonably well (Riess et al. 1998; Permutter et al. 1999; Knop et al. 2003; Riess et al. 2004). However, the main problem in this model is the huge difference of about 10120 orders of magnitude between the observed value of the cosmological constant and the one predicted from quantum field theory; known as the cosmological constant problem (Weinberg 1989). Another issue is the so called coincidence problem which expresses the fact that although in this model the matter and dark energy components scale differently with redshift during the evolution of the universe, both components today have comparable energy densities, and it is unclear why we happen to live in this narrow window of time. Besides these main issues, there are other inherent problems faced by the \u039bCDM, some of which arose as a result of recent observations that are in disagreement with the model's predictions. For example, in order to account for the general isotropy of the cosmic microwave background (CMB), the standard model invokes an early period of inflationary expansion (Kazanas 1980; Guth 1981; Linde 1982). However, the latest observations by Planck (Planck Collaboration XXIII 2003) indicate that there may be some problems with such an inflationary scenario (Ijjas, Steinhardt & Loeb 2013; Guth, Kaiser & Nomura 2014). It was partly due to these issues of the standard \u039bCDM, that during the last decade several alternative dark energy models have been proposed and tested with observations. In these models the dark energy density component \u03c1de is not constant and in most cases \u03c9de = pde\/\u03c1de depends on time, redshift, or scale factor. For example in some of these so called dynamical dark energy models, late time inflation is achieved using a variable cosmological term \u039b(t) (Ray et al. 2011; Basilakos 2015) sometimes taken in conjunction with a time dependent gravitational constant G(t) (Ray, Mukhopadhyay & Dutta Choudhury 2007; Ibotombi Singh, Bembem Devi & Surendra Singh 2013). Other sources of dark energy include scalar fields such as quintessence (Peebles & Ratra 2003), K-essence (Armendariz-Picon et al. 2001) and phantom fields (Singh, Sami & Dadhich 2003). An alternative approach to the dark energy problem relies on the modification of Einstein's theory itself such that in these alternative theories of gravity, cosmic acceleration is not provided solely by the matter side T\u03bc\u03bd of the field equations, but also by the geometry of spacetime. These theories include the scalar-tensor theory with non-minimally coupled scalar fields (Barrow & Parsons 1997; Bertolami & Martins 2000), f(R) theory (Tsujikawa 2008), conformal Weyl gravity (Mannheim 2000) and higher dimensional theories such as the Randall\u2013Sundrum (RS) braneworld model (Randall & Sundrum 1999), and the braneworld model of Dvali\u2013Gabadadze\u2013Porrati (DGP) (Dvali, Gabadadze & Porrati 2000). Over the last few years considerable interest has been shown in the simple FRW linearly expanding (coasting) model in Einstein's theory with a(t)\u2009\u221d\u2009t, H(z) = H0(1 + z). Like the \u039bCDM the total energy density and pressure in this model are expressed in terms of matter, radiation and dark energy components, such that p = \u03c9\u03c1 with \u03c1 = \u03c1m + \u03c1r + \u03c1de and p = pr + pde (since pm \u2248 0), but it includes the added assumption \u03c9 = \u22121\/3, i.e. the cosmic fluid acting as the source has zero active gravitational mass. So this would definitely exclude a cosmological constant as the source of the dark energy component in this case. The model was first discussed by Kolb (1989) who referred to this zero active mass cosmic fluid as \u2018K-matter\u2019. Interest in this model has been revived recently after it was noted (Melia 2003) that in the standard model the radius of the gravitational horizon Rh(t0) (also known as the Hubble radius) is equal to the distance ct0 that light has travelled since the big bang, with t0 being the current age of the universe. In the \u039bCDM this equality is a peculiar coincidence because it just happens at the present time t0. It was then proposed (Melia 2007, 2009; Melia & Shevchuk 2012) that this equality may not be a coincidence at all, and should be satisfied at all cosmic time t. This was done by an application of Birkhoff's theorem and its corollary, which for a flat universe allows the identification of the Hubble radius Rh with the gravitational radius Rh = 2GM\/c2, given in terms of the Misner\u2013Sharp mass $M = (4\\pi \/3)R_{{\\rm h}}^{3}(\\rho \/c^2)$ (Misner & Sharp 1964). The added assumption of a zero active gravitational mass \u03c1 + 3p = 0 implies (Melia & Shevchuk 2012) that Rh = ct or H = 1\/t for any cosmic time t. This linear model became known as the Rh = ct universe. Unlike the \u039bCDM\/\u03c9CDM which contains at least the three parameters H0, \u03a9m and \u03c9de, the Rh = ct model depends only on the sole parameter H0, so that for example the luminosity distance used to fit Type Ia supernova data (Melia 2009) is given by the simple expression dL = (1 + z)Rh(t0)ln\u2009(1 + z). Also while the \u039bCDM would need inflation to circumvent the well-known horizon problem, the Rh = ct universe does not require inflation. One should also point out that the condition Rh = ct is also satisfied by other linear models such as the Milne universe (Milne 1933), which however has been refuted by observations. Unlike the Rh = ct model discussed here, the Milne universe is empty (\u03c1 = 0) and with a negative spatial curvature (k = \u22121). As a result of these properties its luminosity distance is given by $d_{L}^{\\rm {Milne}} = R_{{\\rm h}}(t_0)(1 + z)\\sinh [\\ln (1+z)]$, and it was shown that this is not consistent with observational data (Melia & Shevchuk 2012). In the last few years the Rh = ct universe received a lot of attention when it was shown (Melia & Maier 2013; Wei, Wu & Melia 2013, 2014a, 2014b, 2015; Melia, Wei & Wu 2015) that it is actually favoured over the standard \u039bCDM (and its variant \u03c9CDM with \u03c9 \u2260 \u22121) by most observational data. This claim has been contested by Bilicki & Seikel (2012) and Shafer (2015) who argued that measurement of H(z) as a function of redshift and the analysis of Type Ia supernovae favoured the \u039bCDM over the Rh = ct universe. However, this was later contested by Melia & McClintock (2015) who showed that the Rh = ct was still favoured when using model-independent measurements that are not biased towards a specific model. Others (see for example van Oirschot, Kwan & Lewis 2010; Lewis & van Oirschot 2012; Mitra 2014) have also criticized the model itself, particularly the validity of the EOS \u03c9 = \u22121\/3 (Lewis 2013). These and other criticisms have been addressed by Bikwa, Melia & Shevchuk (2012); Melia (2012) (see also Melia 2015 and references therein.) As pointed out above the Rh = ct model would still require a dark energy component \u03c1de, albeit not in the form of a cosmological constant. So the obvious question at this point would be: what are the possible sources for this component that together with the matter and radiation components will give the required total EOS, \u03c9 = \u22121\/3? The purpose of this paper is to answer this question by discussing the various possible sources of dark energy that are consistent with this EOS. Since the radiation component \u03c1r at the present time t0 is insignificant (at least for the \u039bCDM with which this model has been compared) we assume that the total energy density \u03c1 = \u03c1de + \u03c1m and the total pressure p = pde (pm \u2248 0), as is normally done in the other alternative dynamical dark energy models found in the literature. So in the next three sections we examine three possibilities for the source of dark energy in the Rh = ct model, namely a variable cosmological term \u039b(t), a non-minimally coupled scalar field in Brans\u2013Dicke theory which is equivalent to a variable gravitational constant G(t), and finally quintessence represented by a minimally coupled scalar field \u03d5. We show that although the first two sources are consistent with the model, they are both unphysical, which leaves the third source of quintessence as the viable source of dark energy in the Rh = ct universe. Results are then discussed in the Conclusion. Unless otherwise noted we use units such that G = c = 1.","Citation Text":["Ijjas, Steinhardt & Loeb 2013"],"Functions Text":["However, the latest observations by Planck","indicate that there may be some problems with such an inflationary scenario"],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[2332,2361]],"Functions Start End":[[2178,2220],[2255,2330]]}
{"Identifier":"2022ApJ...936..102A__Williams_et_al._2006_Instance_3","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)"],"Functions Text":["Their simulation results also reveal that the bipolar pulses reported in the Enceladus plume by","and Pickett et al. (2015) could most probably be ion holes."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[3115,3137]],"Functions Start End":[[3019,3114],[3138,3197]]}
{"Identifier":"2022MNRAS.509.3488I__Terzi\u0107_&_Graham_2005_Instance_1","Paragraph":"The values rinf, \u03c1inf, and \u03c3inf of equations (9) and (10) were computed assuming a bulge mass profile. Unlike other works that use isothermal sphere or Dehnen profiles (see e.g Volonteri, Haardt & Madau 2003; Sesana 2010; Sesana & Khan 2015; Bonetti et al. 2018b; Volonteri et al. 2020), here we decided to use a S\u00e9rsic model. This choice is motivated by observational studies that found it to be a good approximation for fitting the bulge light distribution of different galaxies (Drory & Fisher 2007; Drory & Alvarez 2008; Gadotti 2009). The analytical expressions for the S\u00e9rsic model are taken from Prugniel & Simien (1997) (see also Terzi\u0107 & Graham 2005):\n(11)$$\\begin{eqnarray*}\r\n\\rm \\rho (\\mathit {r}) = \\rho _0 \\left(\\frac{\\mathit {r}}{{\\it R}_e}\\right)^{-{\\it p}} \\mathit {\\rm e}^{-{\\it b}\\left({\\mathit {r}}\/{{\\it R}_e}\\right)^{1\/{\\it n}}},\r\n\\end{eqnarray*}$$(12)$$\\begin{eqnarray*}\r\n\\rm \\sigma ^2(\\mathit {r}) &amp;=&amp; \\frac{4 \\pi G \\rho _0^2 \\mathrm{{\\it R}_{\\rm e}}^2 n^2 b^{2\\mathit {n}(p-1)}}{\\rho (\\mathit {r})} \\nonumber\\\\&amp;&amp;\\times \\, \\int _Z^{\\infty } \\mathcal {Z}^{-\\mathit {n}(p+1)-1} {\\rm e}^{-\\mathcal {Z}} \\gamma (\\mathit {n}(3-p),\\mathcal {Z}) {\\rm d}\\mathcal {Z},\r\n\\end{eqnarray*}$$where $\\rm {\\it R}_e$ is the bulge effective radius,8 \u03c10 is the central bulge density, and n is its S\u00e9rsic index. This index correlates with the central concentration of the bulge, being the bulges with smaller n, the ones less centrally concentrated. Finally, the variable \u03b3 represents the incomplete gamma function, whereas Z, p, and b are three different quantities that depend on the bulge properties: $\\rm {\\it Z} = {\\it b}(\\mathit {r}\/{\\it R}_e)^{1\/{\\it n}}$, $p = 1 - 0.6097 n^{-1} + 0.05563 n^{-2}$, and $\\mathrm{\\it b} = 2n - 0.33 + 0.009\\,876 n^{-1}$. This S\u00e9rsic model causes that smaller MBHs spend more time in the hardening phase than the most massive ones. To guide the reader, for an MBHB system with total mass $\\rm {\\it M}_{bin} = 10^9\\, M_{\\odot }$, $q = 1$, and $e_{\\rm BH} = 0.3$, the hardening time-scale is $\\rm {\\sim }0.2\\, Gyr$. For the same system but with $\\rm {\\it M}_{bin} = 10^6\\, M_{\\odot }$, the time increases up to $\\rm 10\\, Gyr$. For further details, we refer to Biava et al. (2019), where a detailed study of hardening time-scales in different bulge profiles was performed.","Citation Text":["Terzi\u0107 & Graham 2005"],"Functions Text":["The analytical expressions for the S\u00e9rsic model are taken from Prugniel & Simien (1997) (see also"],"Functions Label":["Uses"],"Citation Start End":[[638,658]],"Functions Start End":[[540,637]]}
{"Identifier":"2022MNRAS.515...71S___2017_Instance_1","Paragraph":"\nSteady and smooth decline: Since falling from the peak of the 2012b event, SN 2009ip has only continued to fade and is now the faintest it has ever been in the optical. Immediately after the 2012 event (around day 200), it was declining somewhat slower than the rate of 56Co decay for a 56Ni mass of 0.04 $\\, {\\rm M}_{\\odot }$ (shown by the dashed grey line in Fig. 2). During that time, its light curve and spectral properties were consistent with those of late-time SNe IIn, and it had an underlying broad-line spectrum similar to that of SN 1987A (Graham et al. 2014; Smith et al. 2014). Up to about day 1000 it continued to fade smoothly at a slower rate around 0.003 mag d\u22121, and spectroscopically it continued to resemble late-time interaction in SNe IIn (Fox et al. 2015; Smith et al. 2016b; Graham et al. 2017). Our new HST\u2009photometry shows that from about day 1000 to day 3000, SN 2009ip has continued to fade smoothly and steadily in the optical, at an even slower rate. While the HST cadence is obviously sparse at these late times, the filters with more than two observations (F555W and F814W) show no significant deviation from a steady decline; there is no evidence of any rebrightening or irregularity in the fading rate. From 2015 to 2021, the decline rates in the various optical filters are 0.00051 \u00b1 0.00001 mag d\u22121 in F555W, 0.00068 \u00b1 0.00002 mag d\u22121 in F606W, 0.00092 \u00b1 0.00001 mag d\u22121 in F657N, and 0.00050 \u00b1 0.00001 mag d\u22121 in F814W. (The UV is an exception, as discussed below.) While this decline is much slower than radioactive decay, such slow decline rates are not at all unusual for SNe IIn with late-time CSM interaction. SNe IIn span a wide diversity of late-time decay rates, ranging from some SNe IIn that have essentially flat light curves for many years, like SN 2005ip (Smith et al. 2009, 2017; Stritzinger et al. 2012; Fox et al. 2020), down to those that have only weak CSM interaction and faster decline rates that are difficult to distinguish from radioactive decay or light echoes. Some SNe IIn even rebrighten at late times, like SN 2006jd (Stritzinger et al. 2012). In any case, it seems as if SN 2009ip is levelling off and approaching a constant V absolute magnitude of \u22127.5 or so. In all three broad continuum filters, the object is now fainter than the progenitor, and the light curve has not exhibited any additional eruptive variability since the 2012b event.","Citation Text":["Smith et al.","2017"],"Functions Text":["SNe IIn span a wide diversity of late-time decay rates, ranging from some SNe IIn that have essentially flat light curves for many years, like SN 2005ip","down to those that have only weak CSM interaction and faster decline rates that are difficult to distinguish from radioactive decay or light echoes."],"Functions Label":["Background","Background"],"Citation Start End":[[1807,1819],[1826,1830]],"Functions Start End":[[1653,1805],[1875,2023]]}
{"Identifier":"2019AandA...627A.130D__Broadhurst_et_al._2019_Instance_2","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"],"Functions Text":["Recent works have studied lensing effects in the existing LIGO\/Virgo O1 and O2 events"],"Functions Label":["Background"],"Citation Start End":[[741,763]],"Functions Start End":[[630,715]]}
{"Identifier":"2019MNRAS.482..194B__Mallmann_et_al._2018_Instance_1","Paragraph":"Observationally, it is a big challenge to conclude whether AGN\u2019s feedback could regulate star formation or not, and how it regulates star formation. On one hand, if strong outflows emerge, they could clear out the star-forming gas to suppress star formation (Alexander & Hickox 2012; Garc\u00eda-Burillo et al. 2014; Alatalo et al. 2015; Hopkins et al. 2016; Wylezalek & Zakamska 2016). The heating by jets propagating through the galaxies could prevent gas from cooling and cut-off the gas supply for further star formation (Karouzos et al. 2014; Choi et al. 2015). On the other hand, the outflows and jets interact with the gas in host galaxies and compress it to trigger new star formation (Silk 2013; Zubovas et al. 2013; Zubovas & Bourne 2017). In fact observations show either no or positive relationships between star formation rates and SMBH accretion rates but no negative trends are seen (Shi et al. 2007, 2009; Baum et al. 2010; Xu et al. 2015; Zhang et al. 2016; Mallmann et al. 2018). Whether AGN\u2019s feedback plays the role may also depends on the spatial scale that observations could resolve and time-scale that the observed tracers could probe (Harrison et al. 2012; Cresci et al. 2015; Feruglio et al. 2015). For example, radiation from AGNs nearly instantaneously impact the surrounding ISM while the attenuation from ISM probably limits their impact to the nuclear regions (Roos et al. 2015). Outflows or jets travel slowly and may be decelerated after interactions with ambient gas, which delays their effects on star formation at large distances from the nuclei (Harrison 2017; Harrison et al. 2018). Feedback by jets or outflows on ISM also depends on their orientation relative to the dusty torus, making their effects on star formation to be anisotropic. The short duty cycle of AGNs could also make feedback by radiation from AGNs temporally variable in strength. Case studies of individual AGNs find evidence of coexistence of positive and negative feedback on star formation (Zinn et al. 2013), suggesting the complicated nature of AGN feedback.","Citation Text":["Mallmann et al. 2018"],"Functions Text":["In fact observations show either no or positive relationships between star formation rates and SMBH accretion rates but no negative trends are seen"],"Functions Label":["Motivation"],"Citation Start End":[[970,990]],"Functions Start End":[[745,892]]}
{"Identifier":"2019MNRAS.487..845B__Acero_et_al._2015_Instance_1","Paragraph":"\nFermi\u2013LAT (Atwood et al. 2009) is a pair conversion telescope, with a field of view (FoV) of above 2 sr, operating in the energy range from 20 MeV to 300 GeV. It is the most sensitive instrument available in this energy range (Ackermann et al. 2012). A few months after the launch in 2008 June, Fermi\u2013LAT started to operate in all sky survey mode. The telescope scans the whole sky in 3 h (Atwood et al. 2009). For this paper we have analysed the Pass81 data from 2010 February 13 to 17. We analysed the data for Mrk 421 using the latest Fermi Science Tool2 software package version v10r0p5. In order to determine the flux and the spectrum of the source, maximum likelihood optimization has been used (Abdo et al. 2009). The data selection and quality checks have been made using the gtselect tool.3 Since the telescope is sensitive to \u03b3-rays from the interactions of cosmic rays with the ambient matter, we set our maximum zenith angle at 105 deg to remove the background \u03b3-ray events from Earth\u2019s limb. The analysis included all photons from a circular region of 10 deg around Mrk 421, which we call the region of interest (ROI). Only photons of energy above 100 MeV were considered for further analysis. The latest LAT instrument response function \u2018P8R2_SOURCE_V6\u2019 has been used. The third Fermi\u2013LAT catalogue (3FGL catalogue: Acero et al. 2015) has been used to include the contributions of sources inside the ROI. The spectral model of the source has been considered as a simple power law of the form $\\mathrm{ d}N(E)\/\\mathrm{ d}E = N_0 (E\/E_0)^{-\\Gamma _{\\mathrm{ ph}}}$, where N0 is called the prefactor, \u0393ph is the index, and E0 is the scale in energy. The spectral parameters of the sources including Mrk 421 inside the ROI are kept free, whereas the spectral parameters of the sources beyond 10 deg from Mrk 421 are kept fixed to the values according to the 3FGL catalogue. The unbinned likelihood4 method has been used in order to estimate the detection significance of the sources. The test statistics parameter determined from the aforementioned method is given by TS = 2\u0394log(L), where L denotes the likelihood function between the model with the source and without the source. According to the definition TS = 9 corresponds to a detection significance of \u223c3\u03c3 (Mattox et al. 1996). All the sources with TS  9 are excluded from the likelihood analysis. In order to model the spectrum of sources, we used the latest Galactic diffuse emission model \u2018gll_iem_v06\u2019 and the isotropic background model \u2018iso_P8R2_SOURCE_V6_v06\u2019. We estimated the corresponding butterfly plots for the source using the methods mentioned in the Fermi\u2013LAT webpage.5 VERITAS observed the high-flux state of Mrk 421 for \u223c6 h. In order to obtain quasi-simultaneous data with VERITAS, we also analysed the Fermi\u2013LAT data for this state by taking 12 h around the VERITAS observation period. In Fig. 2(d), we show the butterfly plots for both 24 and 12 h by closed butterfly (grey) and open butterfly with cross at the edges (blue), respectively, for the high-flux state.","Citation Text":["Acero et al. 2015"],"Functions Text":["The third Fermi\u2013LAT catalogue (3FGL catalogue:","has been used to include the contributions of sources inside the ROI."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1331,1348]],"Functions Start End":[[1284,1330],[1350,1419]]}
{"Identifier":"2020ApJ...903L..22T__Vuitton_et_al._2007_Instance_2","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"],"Functions Text":["Aside from electron dissociative recombination of C4H3NH+","neutral production of CH3C3N can occur in a few ways"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[568,587]],"Functions Start End":[[509,566],[590,642]]}
{"Identifier":"2020MNRAS.493L..98S__Sourie,_Oertel_&_Novak_2016_Instance_1","Paragraph":"To solve equations (9)\u2212(11), the pulsar rotation rate \u03a90, the long-term spin-down rate $\\dot{\\Omega }_\\infty$, the initial lag \u03b4\u03a90, the mutual-friction coefficients $\\mathcal {B}_{\\operatorname{\\text{f}}}$ and $\\mathcal {B}_{\\operatorname{\\text{pin}}}$, and the ratios $I_{\\operatorname{\\text{n}}}^{\\operatorname{\\text{f}}}\/I$ and $I_{\\operatorname{\\text{n}}}^{\\operatorname{\\text{pin}}}\/I$ need to be specified. In what follows, \u03a90 and $\\dot{\\Omega }_\\infty$ are directly taken from pulsar timing. The coefficient $\\mathcal {B}_{\\operatorname{\\text{f}}}$ in the non-pinned region is given by $\\mathcal {B}_0$, and the corresponding drag-to-lift ratio by equation (8). In the pinned region, the coefficient $\\mathcal {B}_{\\operatorname{\\text{pin}}}$ is given by equation (6), with the prescription (7) and suitable parameters. Typical values for the underlying parameters are: $\\varepsilon _{\\operatorname{\\text{p}}}^{\\operatorname{\\text{pin}}}\\simeq 0.05-0.2$, $\\varepsilon _{\\operatorname{\\text{p}}}^{\\operatorname{\\text{f}}}\\simeq 0.1-0.5$ (see e.g. Chamel & Haensel 2006; Sourie, Oertel & Novak 2016), $x_{\\operatorname{\\text{p}}}^{\\operatorname{\\text{pin}}} \\simeq 0.05 - 0.1$, $x_{\\operatorname{\\text{p}}}^{\\operatorname{\\text{f}}} \\simeq 0.05 - 0.4$, $\\rho _{\\operatorname{\\text{pin}}} \\simeq (0.5-2) \\rho _0$, and $\\rho _{\\operatorname{\\text{f}}} \\simeq (2-6) \\rho _0$, \u03c10 \u2243 2.7 \u00d7 1014 g cm\u22123 being the nuclear saturation density (see e.g. Pearson et al. 2018). The ratios $I_{\\operatorname{\\text{n}}}^{\\operatorname{\\text{pin}}}\/I$ and $I_{\\operatorname{\\text{n}}}^{\\operatorname{\\text{f}}}\/I$ are computed using the relations $I_{\\operatorname{\\text{n}}}^{_X} \/I^{_X} = 1-x_{\\operatorname{\\text{p}}}^{_X}$ (assuming uniform densities in each region), where $I^{_X}$ is the total moment of inertia of region X, and $I^{\\operatorname{\\text{f}}} = I-I^{\\operatorname{\\, \\text{cr}}}-I^{\\operatorname{\\text{pin}}}$, $I^{\\operatorname{\\, \\text{cr}}}$ denoting the crustal moment of inertia of the star. Typical values are: $I^{\\operatorname{\\, \\text{cr}}}\/I\\simeq 0.01-0.05$ (Delsate et al. 2016) and $I^{\\operatorname{\\text{pin}}}\/I\\sim 0.05$ (G\u00fcgercino\u011flu & Alpar 2014). Unlike the previous quantities, both the initial lag \u03b4\u03a90 and number $N_{\\operatorname{\\text{p}}}$ of pinned fluxoids are essentially unknown. As shown in the next section, the large range of possible values for $N_{\\operatorname{\\text{p}}}$ could account for the very different spin-up behaviours in the Crab and Vela pulsars.","Citation Text":["Sourie, Oertel & Novak 2016"],"Functions Text":["Typical values for the underlying parameters are: $\\varepsilon _{\\operatorname{\\text{p}}}^{\\operatorname{\\text{pin}}}\\simeq 0.05-0.2$, $\\varepsilon _{\\operatorname{\\text{p}}}^{\\operatorname{\\text{f}}}\\simeq 0.1-0.5$ (see e.g."],"Functions Label":["Uses"],"Citation Start End":[[1076,1103]],"Functions Start End":[[827,1052]]}
{"Identifier":"2021MNRAS.505.2111L__Xu_et_al._2018_Instance_1","Paragraph":"Recently, quasars observed with multiple measurements, another potential cosmological probe with a higher redshift range that reaches to z \u223c 5, is becoming popular to constrain cosmological models in the largely unexplored portion of redshift range from z \u223c 2 to z \u223c 5. A sample that contains 120 angular size measurements in intermediate-luminosity quasars from the very long baseline interferometry (VLBI) observations (Cao et al. 2017a,b) has become an effective standard ruler, which have been extensively applied to test cosmological models (Li et al. 2017; Melia et al. 2017; Qi et al. 2017; Zheng et al. 2017; Xu et al. 2018; Ryan, Chen & Ratra 2019), measuring the speed of light (Cao et al. 2017a, 2020a), exploring cosmic curvature at different redshifts (Cao et al. 2019; Qi et al. 2019), and the validity of cosmic distance duality relation (Zheng et al. 2020). Then, Risaliti & Lusso (2019) put forward a new compilation of quasars containing 1598 quasi-stellar object (QSO) X-ray and ultraviolet (UV) flux measurements in the redshift range of 0.036 \u2264 z \u2264 5.1003, which have been used to constrain cosmological models (Khadka & Ratra 2020b) and cosmic curvature at high redshifts (Liu et al. 2020a,c), as well as test the cosmic opacity (Geng et al. 2020; Liu et al. 2020b). Making use of this data to explore cosmological researches mainly depends on the empirical relationship between the X-ray and UV luminosity of these high-redshift quasars proposed by Avni & Tananbaum (1986), which leads to the Hubble diagram constructed by quasars (Risaliti & Lusso 2015, 2017; Lusso & Risaliti 2016; Bisogni, Risaliti & Lusso 2017). In general, the advantage of these two QSO measurements over other traditional cosmological probes is that QSO has a larger redshift range, which may be rewarding in exploring the behaviour of the non-standard cosmological models at high redshifts, providing an important supplement to other astrophysical observations and also demonstrating the ability of QSO as an additional cosmological probe (Zheng et al. 2021).","Citation Text":["Xu et al. 2018"],"Functions Text":["Recently, quasars observed with multiple measurements, another potential cosmological probe with a higher redshift range that reaches to z \u223c 5, is becoming popular to constrain cosmological models in the largely unexplored portion of redshift range from z \u223c 2 to z \u223c 5. A sample that contains 120 angular size measurements in intermediate-luminosity quasars from the very long baseline interferometry (VLBI) observations","has become an effective standard ruler, which have been extensively applied to test cosmological models"],"Functions Label":["Background","Background"],"Citation Start End":[[617,631]],"Functions Start End":[[0,420],[442,545]]}
{"Identifier":"2022AandA...666A..67P__Gall_et_al._2017_Instance_1","Paragraph":"The KN coincident with GW170817 showed a rapidly fading EM transient in the optical and infrared bands. The term kilonova was historically coined from the perception of its brightness being a thousand times larger than a nova (Li & Paczy\u0144ski 1998; Rosswog 2005; Metzger et al. 2010). We discuss in this work that, even if far from such a standardized luminosity, clear dependencies on physical BNS quantities can be extracted. For AT 2017gfo, the observed spectral evolution is suggested to arise from a mixed composition of r-process elements in the ejected material (Gillanders et al. 2022); while early-time spectra are consistent with light r-process elements (nuclear masses A\u2004\u223c\u200490\u2005\u2212\u2005140), later spectra require intermediate composition, producing even heavier elements such as lanthanides. An alternative scenario, such as dust formation in the KN that could also explain such a blue-red evolution, has been ruled out (Gall et al. 2017). In addition, a broad feature observed at \u223c0.7\u2005\u2212\u20051\u2006\u03bcm has been interpreted as due to Sr\u202fII (Watson et al. 2019; Domoto et al. 2021; Gillanders et al. 2022). Since then, no other KN has been firmly detected (see e.g., Coughlin 2020 and references therein for a summary of the follow-up efforts during LIGO and Virgo\u2019s third observing campaign). Another merger candidate is linked to GW190814, located farther away at \u223c240 Mpc, whose origin remains uncertain. However, it lacks of any EM counterpart and could likely be caused by a binary black hole merger (Essick & Landry 2020; Tews et al. 2021). GW190425 was the first BNS candidate in O3. The LALInference localization pipeline (Veitch et al. 2015) provided a 90% credible region of 7461 deg2, estimated luminosity distance of 156 \u00b1 41 Mpc, much larger than that for GW170817, which complicated the search campaign resulting in no EM counterpart identified. The inferred merger parameters from the GW analysis published in Abbott et al. (2020a) for chirp mass \n\n\n\n1\n.\n\n44\n\n\u2212\n0.02\n\n\n+\n0.02\n\n\n\n\nM\n\u2299\n\n\n\n$ 1.44^{+0.02}_{-0.02}\\,M_{\\odot} $\n\n\n are the following. Total mass \n\n\n\n\nM\n1\n\n+\n\nM\n2\n\n\u2243\n3\n.\n\n4\n\n\u2212\n0.1\n\n\n+\n0.3\n\n\n\n\nM\n\u2299\n\n\n\n$ M_{1}+M_{2} \\simeq 3.4_{-0.1}^{+0.3}\\,M_{\\odot} $\n\n\n, with M1\u2004\u2208\u2004(1.62,\u20061.88)\u2006M\u2299, M2\u2004\u2208\u2004(1.45,\u20061.69)\u2006M\u2299, and \n\n\n\n\n\u039b\n\u223c\n\n\u2264\n600\n\n\n$ \\tilde{\\Lambda} \\leq 600 $\n\n\n for low spin prior, or M1\u2004\u2208\u2004(1.61,\u20062.52)\u2006M\u2299, M2\u2004\u2208\u2004(1.12,\u20061.68)\u2006M\u2299, and \n\n\n\n\n\u039b\n\u223c\n\n\u2264\n1100\n\n\n$ \\tilde{\\Lambda} \\leq 1100 $\n\n\n for high-spin prior. Note that for the other confirmed BNS event, GW170817, the inferred total mass is \n\n\n\n\nM\n1\n\n+\n\nM\n2\n\n\u2243\n2\n.\n\n73\n\n\u2212\n0.01\n\n\n+\n0.04\n\n\n\n\nM\n\u2299\n\n\n\n$ M_{1}+M_{2} \\simeq 2.73_{-0.01}^{+0.04}\\,M_{\\odot} $\n\n\n with M1\u2004\u2208\u2004(1.36,\u20061.60)\u2006M\u2299, M2\u2004\u2208\u2004(1.16,\u20061.36)\u2006M\u2299, and \n\n\n\n\n\u039b\n\u223c\n\n=\n\n300\n\n\u2212\n230\n\n\n+\n420\n\n\n\n\n$ \\tilde{\\Lambda} = 300_{-230}^{+420} $\n\n\n for low-spin prior (Abbott et al. 2019).","Citation Text":["Gall et al. 2017"],"Functions Text":["An alternative scenario, such as dust formation in the KN that could also explain such a blue-red evolution, has been ruled out"],"Functions Label":["Background"],"Citation Start End":[[925,941]],"Functions Start End":[[796,923]]}
{"Identifier":"2019ApJ...881...42J__Kelson_et_al._2000_Instance_1","Paragraph":"Stellar population evolution studies beyond z \u2248 1 have primarily focused on ages through studies of luminosity changes. Beifiori et al. (2017) used new data for 19 galaxies in z = 1.3\u20131.6 clusters obtained with the Very Large Telescope\/KMOS to extend the redshift coverage of the results regarding the evolution of the mass-to-light (M\/L) ratios of bulge-dominated passive galaxies. The authors used their new results together with the available literature results covering up to z = 1.3 (van Dokkum & Franx 1996; J\u00f8rgensen et al. 1999, 2006, 2014; Kelson et al. 2000; Wuyts et al. 2004; Holden et al. 2005, 2010; Barr et al. 2006; van Dokkum & van der Marel 2007; Saglia et al. 2010; J\u00f8rgensen & Chiboucas 2013) and low-redshift reference data for the Coma cluster (J\u00f8rgensen 1999; J\u00f8rgensen et al. 2006) to further solidify the evidence supporting passive evolution and a formation redshift zform \u2248 2. The formation redshift should be understood as the epoch of the last major star formation episode. At z \u2248 1 the massive (Mass > 1011 M) bulge-dominated galaxies in clusters appear to be in place and mostly passively evolving. Lower mass galaxies may still be added to the red sequence and from then on passively evolve (e.g., S\u00e1nchez-Bl\u00e1zquez et al. 2009; Choi et al. 2014), but see also Cerulo et al. (2016) for results supporting that the red sequence well below L\u22c6 is fully populated in rich clusters already at \n\n\n\n\n\n. Ultimately, the properties of galaxies mapped over a large fraction of the age of the universe, may constrain the models for building the galaxies. It is difficult to understand within the prevailing hierarchical model favored by the \u039bCDM (cold dark matter) cosmology, the existence of such massive passive galaxies with relatively old stellar populations at z \u2248 1, while less massive galaxies appear to harbor younger stellar populations, e.g., J\u00f8rgensen et al. (2017, and references therein), see Kauffmann et al. (2003) for a discussion of this tension between the observational results and the hierarchical models of galaxy formation. However, more recent cosmological simulations like Illustris (Genel et al. 2014; Vogelsberger et al. 2014; Wellons et al. 2015) and UniverseMachine (Behroozi et al. 2019) find that massive quiescent galaxies can be in place by z \u2273 2.","Citation Text":["Kelson et al. 2000"],"Functions Text":["The authors used their new results together with the available literature results covering up to z = 1.3","to further solidify the evidence supporting passive evolution and a formation redshift zform \u2248 2."],"Functions Label":["Uses","Compare\/Contrast"],"Citation Start End":[[549,567]],"Functions Start End":[[383,487],[806,903]]}
{"Identifier":"2022ApJ...928...38T__Tiwari_&_Nusser_2016_Instance_1","Paragraph":"The LoTSS survey, homogeneously covering the whole northern sky complete down to the sub mJy limit will overcome statistical limitations due to shot noise. The large galaxy number density and large sky coverage will substantially reduce cosmic variance in cosmological analysis. The radio galaxies, tracing the background dark matter, will constrain the shape of power spectrum, i.e., the early universe physics, dark matter, baryon, density of neutrinos, the inflation power spectrum, and the degree of non-Gaussianity in density fluctuations. The upcoming LoTSS catalogs, covering a large sky area, will help us to explore further regarding large-scale anomalies (de Oliveira-Costa et al. 2004; Ralston & Jain 2004; Schwarz et al. 2004; Tiwari & Aluri 2019) and the current puzzling dipole signal observed with radio catalogs (Blake & Wall 2002; Singal 2011; Gibelyou & Huterer 2012; Rubart & Schwarz 2013; Tiwari & Jain 2015; Tiwari et al. 2015; Tiwari & Nusser 2016; Colin et al. 2017; Siewert et al. 2021). Furthermore, LoTSS will significantly improve on present low-frequency radio catalogs, e.g., TIFR GMRT Sky Survey (TGSS; Intema et al. 2017) and GaLactic and Extragalactic All-sky MWA (GLEAM; Hurley-Walker et al. 2017), and analyses based on these surveys (Dolfi et al. 2019; Rana & Bagla 2019; Tiwari 2019; Tiwari et al. 2019; Choudhuri et al. 2020). Unfortunately, the link between the galaxy and total matter power spectra depends on some unknowns from astrophysics, such as the galaxy bias factor, which depends on galaxy type and is quite different for radio AGNs and star-forming galaxies. The LoTSS population is a mixture of AGNs and star-forming galaxies, and therefore understanding galaxy bias, relative number densities, and luminosity evolution is nontrivial. The purpose of this work is to present a detailed cosmological analysis of LoTSS galaxies and study the effect of survey footprint, shot noise, and other systematics. We have produced galaxy mocks for the survey and have customized and calibrated the data pipeline for galaxy clustering statistics recovery.","Citation Text":["Tiwari & Nusser 2016"],"Functions Text":["The upcoming LoTSS catalogs, covering a large sky area, will help us to explore further regarding","and the current puzzling dipole signal observed with radio catalogs"],"Functions Label":["Future Work","Future Work"],"Citation Start End":[[949,969]],"Functions Start End":[[545,642],[760,827]]}
{"Identifier":"2022MNRAS.510.4723P__Henri_&_Petrucci_1997_Instance_1","Paragraph":"A new result brought by our work is the role of incident polarization on the reflected spectral outcome. Although some X-ray polarization of the primary radiation is expected to be due to Comptonization of thermal radiation inside the corona, general estimates for the polarization degree of the primary source are rather low ($\\lesssim 10\\,{\\rm {per \\, cent}}$) according to the recent analysis in Beheshtipour et al. (2017), Tamborra et al. (2018), and Beheshtipour (2018). If some polarization was present, then, for symmetric reasons, vertically or horizontally polarized light should be dominant in case of extended coronal models (Dabrowski & Lasenby 2001; Niedzwiecki & Zycki 2008; Schnittman & Krolik 2010) at low heights. In case of lamp-post models (Matt, Perola & Piro 1991; Martocchia & Matt 1996; Henri & Petrucci 1997; Martocchia, Karas & Matt 2000; Dov\u010diak, Karas & Yaqoob 2004a; Miniutti & Fabian 2004) general-relativistic effects will rotate the polarization position angle along null geodesics from the corona towards the disc (Connors et al. 1980) and the situation becomes more general for incident disc irradiation, i.e. any incident state of polarization is possible. Our computations, apart from the unpolarized case, assumed two extreme cases of initial $p = 100\\,{\\rm {per \\, cent}}$ polarization in order to estimate its possible effects in comparison with the completely unpolarized light and in order to test that the code adheres to the stated orientation conventions. Having appended these two initially polarized cases to our tables, it also allows for interpolation of reflection results for any input polarization state from this basis of the three computed independent polarization states, which will be necessary in future construction of global spectropolarimetric models. We aim to address any impact of light bending and other relativistic effects for a distant observer in our future works that will introduce the stokes local tables integrated over the accretion disc with some adopted global geometry.","Citation Text":["Henri & Petrucci 1997"],"Functions Text":["In case of lamp-post models","general-relativistic effects will rotate the polarization position angle along null geodesics from the corona towards the disc","and the situation becomes more general for incident disc irradiation, i.e. any incident state of polarization is possible. Our computations, apart from the unpolarized case, assumed two extreme cases of initial $p = 100\\,{\\rm {per \\, cent}}$ polarization in order to estimate its possible effects in comparison with the completely unpolarized light and in order to test that the code adheres to the stated orientation conventions."],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[810,831]],"Functions Start End":[[731,758],[919,1045],[1068,1498]]}
{"Identifier":"2021MNRAS.501.3781R__Nisini_et_al._2005_Instance_1","Paragraph":"While spatially extended optical jets and bipolar CO molecular outflows have been observed in numerous Class 0\/I protostars (e.g. Reipurth & Bally 2001; Bally 2016, and references therein), near-infrared high-resolution spectroscopy and spectroimaging observations in the past two decades have made it possible to study the kinematics of the outflowing gas and physical properties at the base of the jet within a few hundred au of the driving source in Class 0\/I protostars (e.g. Davis et al. 2001, 2003, 2011; Nisini et al. 2005, 2016; Caratti o Garatti et al. 2006; Takami et al. 2006; Antoniucci et al. 2008, 2011, 2017; Garcia Lopez et al. 2008, 2013). These microjets are bright in [Fe\u2009ii] forbidden and H2 rovibrational emission lines, hence showing the presence of forbidden emission-line (FEL) regions and molecular hydrogen emission-line (MHEL) regions in low-mass Class 0\/I protostars. While multiple low- and high-velocity components are observed in both MHELs and FELs, the higher velocity gas is slightly further offset from the driving source than the slower gas, and the kinematics of the H2 emission differs from [Fe\u2009ii] emission, revealing complicated kinematic structures. Evidence of H2 emission from cavity walls is also seen in some protostars, suggesting the presence of a wide-angled wind. Strong emission in the well-known accretion diagnostics of Paschen and Brackett hydrogen recombination lines is observed in protostars, with the ratio of the accretion luminosity to bolometric luminosity spanning from \u223c0.1 to \u223c1. The mass accretion and loss rates for Class 0\/I low-mass protostars span the range of 10\u22126\u201310\u22128 M\u2299 yr\u22121, and the derived jet efficiencies (ratio between mass ejection and accretion rates) range between \u223c1 per\u2009cent and 10 per\u2009cent (e.g. Davis et al. 2001, 2003, 2011; Nisini et al. 2005, 2016; Caratti o Garatti et al. 2006; Takami et al. 2006; Antoniucci et al. 2008, 2011, 2017; Garcia Lopez et al. 2008, 2013). These measurements are within the range predicted by the magnetohydrodynamic jet launching models (e.g. Frank et al. 2014).","Citation Text":["Nisini et al. 2005"],"Functions Text":["While spatially extended optical jets and bipolar CO molecular outflows have been observed in numerous Class 0\/I protostars","near-infrared high-resolution spectroscopy and spectroimaging observations in the past two decades have made it possible to study the kinematics of the outflowing gas and physical properties at the base of the jet within a few hundred au of the driving source in Class 0\/I protostars (e.g."],"Functions Label":["Background","Background"],"Citation Start End":[[511,529]],"Functions Start End":[[0,123],[190,479]]}
{"Identifier":"2018ApJ...866L...1S__Pecharrom\u00e1n_et_al._1999_Instance_7","Paragraph":"It was found that the complex dielectric function from Pecharrom\u00e1n et al. (1999) for the sample obtained by heating bayerite at 1273 K, assuming a spheroid with depolarization parameters of (0.35, 0.003), produced an opacity with 11, 20, 28, and 32 \u03bcm features, so this component was included in the models. However, with only this component, the observed 20 \u03bcm features in the residual spectra were found to be wider than those in the models. By adding the opacity of the sample obtained by heating boehmite at 1173 K, the width of the 20 \u03bcm feature could be matched. This was done using the complex dielectric function for the sample obtained by heating boehmite at 1173 K from Pecharrom\u00e1n et al. (1999), assuming a spheroid with depolarization parameters of (0.35, 0.035). The complex dielectric functions of the samples obtained by heating bayerite and boehmite to various temperatures (Pecharrom\u00e1n et al. 1999) were derived by modeling the reflectance spectra of pellets obtained by pressing powders of these materials under great pressure. This method required Pecharrom\u00e1n et al. (1999) to assume an effective medium theory, such that a pellet is a mixture of one of their samples with a matrix of air. Pecharrom\u00e1n et al. (1999) noted that heating bayerite at 500\u00b0C eliminates the XRD pattern of bayerite, and they note that at 700\u00b0C, the infrared reflectance spectrum of the boehmite sample no longer shows OH\u2212 stretching bands. This must mean that the samples obtained from heating bayerite at 1273 K and from heating boehmite at 1173 K are no longer bayerite or boehmite, respectively. XRD performed by Pecharrom\u00e1n et al. (1999) of the sample of bayerite prepared at 1273 K suggests only \u03b8-alumina was present, and their infrared and NMR spectroscopy confirms this. XRD of their sample obtained from heating boehmite to 1173 K (Pecharrom\u00e1n et al. 1999) suggests \u03b4-alumina to be present, though some amounts of \u03b8-alumina and \u03b1-alumina are present, as they deduce from XRD and infrared and NMR spectroscopy.","Citation Text":["Pecharrom\u00e1n et al. 1999"],"Functions Text":["XRD of their sample obtained from heating boehmite to 1173 K","suggests \u03b4-alumina to be present, though some amounts of \u03b8-alumina and \u03b1-alumina are present, as they deduce from XRD and infrared and NMR spectroscopy."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1837,1860]],"Functions Start End":[[1775,1835],[1862,2014]]}
{"Identifier":"2019MNRAS.487.1210T__McNamara_&_Nulsen_2007_Instance_2","Paragraph":"On larger scales, the clusters in which BCGs reside can generally be divided into two categories: cool core clusters, which exhibit very peaked surface brightness distributions at X-ray wavelengths, and non cool core clusters, with similar overall X-ray luminosities but with smoother, less peaked X-ray surface brightness distributions. Some authors (e.g. Hudson et al. 2010; Santos et al. 2010) define an intermediate category called moderate or weak cool core clusters. Since cool core clusters have short radiative cooling time-scales on the order of 108 yr in their centres (e.g. Voigt & Fabian 2004; McNamara & Nulsen 2007, 2012; Hlavacek-Larrondo et al. 2012), starbursts are expected to be common at the centre of such clusters. Indeed, the central cool gas in these clusters should condense onto the BCG, forming stars at rates of hundreds of solar masses per year (e.g. Fabian 1994). However, most BCGs are relatively quiescent and those that do show evidence of star formation generally tend to have star formation rates 1 order of magnitude smaller, on the order of $1-150 \\, \\mathrm{M_{\\odot }\\, {yr}^{-1}}$ (e.g. Donahue et al. 2007; Bildfell et al. 2008; O\u2019Dea et al. 2008, 2010; Rawle et al. 2012). This mismatch between expected and observed star-forming rates, known as the cooling flow problem, is thought to be caused by active galactic nuclei (AGNs) feedback processes from the BCG. AGNs can release copious amounts of energy into the intracluster medium (ICM) through many ways, including: jetted outflows that inflate cavities, weak shocks, sound waves, or turbulence in the ICM (e.g. Markevitch & Vikhlinin 2007; McNamara & Nulsen 2007, 2012; Zhuravleva et al. 2014; Fabian et al. 2017). Alone, the energy released by jetted outflows appears to be on the same order as the energy needed to offset cooling (e.g. Rafferty et al. 2006; McNamara & Nulsen 2007; Hlavacek-Larrondo et al. 2012), therefore suggesting that AGN feedback is a good candidate for solving the cooling flow problem.","Citation Text":["McNamara & Nulsen 2007"],"Functions Text":["AGNs can release copious amounts of energy into the intracluster medium (ICM) through many ways, including: jetted outflows that inflate cavities, weak shocks, sound waves, or turbulence in the ICM (e.g."],"Functions Label":["Background"],"Citation Start End":[[1637,1659]],"Functions Start End":[[1404,1607]]}
{"Identifier":"2021MNRAS.506.5935R__Ellison_et_al._2011_Instance_1","Paragraph":"The discovery of a correlation between the mass of supermassive black holes (SMBHs) and several properties of their host galaxies (e.g. Magorrian et al. 1998; Ferrarese & Merritt 2000; Gebhardt et al. 2000; Kormendy & Ho 2013) has suggested that the growth of SMBHs and their host galaxies are tightly connected. Mergers of galaxies are thought to be one of the most important mechanisms with which galaxies build up their stellar masses (White & Rees 1978). Both observational (e.g. Lonsdale, Persson & Matthews 1984; Joseph & Wright 1985; Armus, Heckman & Miley 1987; Clements et al. 1996; Alonso-Herrero et al. 2000; Ellison et al. 2008) and theoretical (e.g. Mihos & Hernquist 1996; Di Matteo et al. 2007) studies have shown that galaxy mergers enhance star formation (SF). Simulations have also shown that the interaction between two or more galaxies can reduce the angular momentum of the circumnuclear material (e.g. Barnes & Hernquist 1991; Blumenthal & Barnes 2018), thus providing an effective mechanism to trigger accretion on to SMBHs (e.g. Di Matteo, Springel & Hernquist 2005). Observationally, several works have confirmed this scenario. Koss et al. (2010) and Silverman et al. (2011) found a higher AGN fraction in pairs than in isolated galaxies with similar stellar masses. It has been shown that the fraction of AGN in mergers tends to increase as the separation between the two galaxies decreases (Ellison et al. 2011), and peaks after coalescence (Ellison et al. 2013). Koss et al. (2012) have shown that the average luminosity of dual AGN also increases with decreasing separation (see also Hou, Li & Liu 2020), and it is higher for the primary (i.e. more massive) component of the system (see also De Rosa et al. 2019, for a recent review). While AGNs with moderate X-ray luminosities are typically found in non-interacting disc galaxies (e.g. Koss et al. 2011; Kocevski et al. 2012; Schawinski et al. 2012), more luminous objects are commonly found in merging systems (e.g. Treister et al. 2012; Glikman et al. 2015; Hong et al. 2015). Treister et al. (2012) showed that, while for 2\u201310\u2009keV AGN luminosities of $L_{2-10}\\sim 10^{41}\\rm \\, erg\\, s^{-1}$ only a small fraction (${\\lt}1{{\\ \\rm per\\ cent}}$) of AGNs are in mergers, at $L_{2-10}\\sim 10^{46}\\rm \\, erg\\, s^{-1}$ \u223c 70\u201380 per cent of the sources are found in interacting systems (see also Glikman et al. 2015). Recent evidence has suggested that hot dust obscured galaxies (Hot DOGs; Wu et al. 2012; Assef et al. 2015), which are some of the most luminous galaxies observed so far ($L_{\\rm \\, IR} \\gt 10^{13} \\mathrm{ L}_{\\odot }$), are also found in mergers (e.g. Fan et al. 2016). These observations suggest that, while at low luminosities SMBH accretion is triggered by secular processes, at high luminosities mergers can play a dominant role. This is in agreement with the evolutionary scenario proposed by Sanders et al. (1988) for ultra-luminous [$L_{\\rm \\, IR}(8\\!-\\!1000\\, \\mu \\rm m)\\ge 10^{12}\\,L_{\\odot }$] infrared galaxies (ULIRGs; e.g. Sanders & Mirabel 1996; P\u00e9rez-Torres et al. 2021). In this scheme, two gas-rich disc galaxies collide, triggering SF and accretion on to the SMBH. The strong accretion on to the SMBH would lead the source to evolve first in a luminous red quasar (e.g. Urrutia, Lacy & Becker 2008; Glikman et al. 2015; LaMassa et al. 2016) and then in an unobscured blue quasar.","Citation Text":["Ellison et al. 2011"],"Functions Text":["It has been shown that the fraction of AGN in mergers tends to increase as the separation between the two galaxies decreases"],"Functions Label":["Background"],"Citation Start End":[[1418,1437]],"Functions Start End":[[1292,1416]]}
{"Identifier":"2021ApJ...911...79P___2007_Instance_1","Paragraph":"Now we turn to the question of whether there are any general trends or strong correlations between the degree of HSP compliance, the average Fidx, and the average \u03a6 derived from the AR samples contained in the defined HRs. As shown in the scatter plot of the average Fidx versus the HSP (Figure 5(a)), we find a weak tendency that HRs with lower degrees of HSP compliance show larger values of the average Fidx, with the linear Pearson correlation coefficient (PCC) of \u22120.55. In the scatter plot, the anti-HSP region with the largest average Fidx can be considered as an extreme case, which lies far away from both the vertical and horizontal dashes lines used to separate data points of the five lowest HSP and the five largest average Fidx, respectively. On the other hand, a positive correlation exists between the average \u03a6 and the average Fidx with the linear PCC = 0.65 (refer to Figure 5(b)). Such correlation of \u03a6 with flaring activity in individual ARs has been reported in many previous studies (e.g., Leka & Barnes 2003, 2007; Park et al. 2010; Liu et al. 2017; Lee et al. 2018), but it is reported here for the first time on this larger spatial scale of the HRs over a much longer period of solar cycle 24. As shown in Figure 5(c), a negative, although weak, correlation appears between the HSP and the average \u03a6 with the linear PCC = \u22120.36. In contrast, a weak trend was reported in Paper I that ARs with larger values of \u03a6 show a higher HSP. These two contrasting trends of the HSP with the average \u03a6 are mainly due to the different ways of dividing the same AR samples into smaller subsets. The AR subsets in this study were selected based on heliographic locations of the AR samples in the defined Carrington longitude\u2013latitude plane, while those in Paper I as a function of \u03a6 values. The two highest flare-productive HRs with the average Fidx \u226515 have the average \u03a6 values ranked in the top two as well as lower degrees of HSP compliance (i.e., one with the lowest HSP and the other with the HSP in the bottom 20%). Meanwhile, the two highly flare-productive regions can be considered as obvious outliers, compared to the positive trend of the HSP with respect to the average \u03a6 in Paper I. This may indicate that the HSP for those regions is obscured by vigorous turbulent convective flows interacting with rising flux tubes therein. As mentioned earlier, the two X-class flaring ARs, NOAA 12673 and NOAA 12192, are located at the two HRs, respectively. Even excluding these two influential ARs, however, all of the trends as described in Figure 5 remain the same, although the correlations become less strong (i.e., the linear PCCs of \u22120.25, 0.51, and \u22120.27 for the cases in Figures 5(a), (b), and (c)).","Citation Text":["Leka & Barnes","2007"],"Functions Text":["On the other hand, a positive correlation exists between the average \u03a6 and the average Fidx with the linear PCC = 0.65 (refer to Figure 5(b)). Such correlation of \u03a6 with flaring activity in individual ARs has been reported in many previous studies (e.g.,"],"Functions Label":["Similarities"],"Citation Start End":[[1012,1025],[1032,1036]],"Functions Start End":[[757,1011]]}
{"Identifier":"2016MNRAS.463.3783B__Reid_&_White_2011_Instance_2","Paragraph":"It is well known that a per cent level understanding of the anisotropy of the redshift-space galaxy clustering is needed to accurately recover cosmological information from the RSD signal in order to shed light on the issue of dark energy versus modified gravity. From a statistical point of view, the source of the anisotropy is the galaxy line-of-sight pairwise velocity distribution. It is therefore important to adopt a realistic functional form for this velocity PDF when fitting models to the data. To this purpose, in Paper I we introduced the GG prescription for the velocity PDF. In this work, we have continued the development of this model by making explicit the dependence of the GG distribution on quantities predictable by theory, namely its first three moments, and extending it to the more general concept of GQG. To keep the model as simple as possible, we have proposed an ansatz with two free dimensionless parameters that describe how infall velocity and velocity dispersion vary when moving from one place to another in our Universe. Since their interpretation is clear, these parameters can be theoretically predicted or, assuming a more pragmatic approach, tuned to simulations or used as nuisance parameters. State-of-the-art PT has proven successful in predicting the large-scale behaviour of the velocity PDF and the correspondent monopole and quadrupole of the redshift-space correlation function (e.g. Reid & White 2011; Wang, Reid & White 2014), at least for massive haloes, M \u223c 1013\u2009M\u2299. Unfortunately, by definition, any PT breaks down for small separations. Consequently, alternative approaches have been suggested in the literature, spanning from purely theoretical (e.g. Sheth 1996) to hybrid techniques in which N-body simulations plus an HOD are employed to deal with the issue of non-linearities (e.g. Tinker 2007; Reid et al. 2014). One of the main results from our work is to provide a framework in which perturbation and small-scale theories are smoothly joined, so that all available RSD information can be coherently extracted from redshift surveys. A fundamental requirement for a redshift-space model is that it must be precise on all scales interest, and it should inform the user of the scales on which the model can be trusted. We have compared to N-body simulations the well-known GSM (Reid & White 2011), the more recent ESM (Uhlemann et al. 2015) and the GQG prescription over a broad range of separations, from 0 to 80\u2009h\u22121 Mpc. Different redshifts, from z = 0 to 1, and different tracers, namely DM particles and two mass-selected catalogues of DM haloes, have been considered. We have concluded that, among the three, QGQ is the only model capable of providing a precise redshift-space correlation function on scales down to \u223c5\u2009h\u22121 Mpc over the range of redshifts covered by future surveys. Keeping in mind that the range of validity of the models depends on tracer, redshift and order of the Legendre multipoles we are interested in, for finiteness, we can say that all the models converge to the expected amplitude on scales \u227330\u2009h\u22121 Mpc, at least for multipole and quadrupole. Since these scales roughly coincide with the range of validity of state-of-the-art PTs, if we rely only on PT and if we are not interested in higher order multipoles, the most natural choice is the simplest model among the three, i.e. the GSM. As for the ESM, we have found it to be unbiased down to smaller scales and for higher order multipoles than the GSM, thus confirming the results by Uhlemann et al. (2015), but, on the other hand, it seems to behave even worse than the GSM on the smallest scales. We can therefore think of it as a natural extension of the GSM in the perspective of further PT developments. In particular, a better prediction of the third moment of the velocity PDF is required before the ESM can be applied to data on smaller scales. Formally, the same argument holds for the GQG model, none the less, since this latter is meant to include non-linear scales, it could be possible to obtain a prediction for the third moment by interpolating between (very) small and (very) large scales. More precisely, as shown in the lower-right panel of Fig. A1, the functions $c^{(3)}_t$ and $c^{(3)}_r$, which fully characterize the third moment, are peaked at r \u2272 10\u2009h\u22121 Mpc. By adopting a model for the small-scale limit that includes those separation, most likely using simulations in a similar way to that proposed in Reid et al. (2014), we would then be able to interpolate between these peaks and their large-scale limit, which is trivially 0.","Citation Text":["Reid & White 2011"],"Functions Text":["We have compared to N-body simulations the well-known GSM","We have concluded that, among the three, QGQ is the only model capable of providing a precise redshift-space correlation function on scales down to \u223c5\u2009h\u22121 Mpc over the range of redshifts covered by future surveys."],"Functions Label":["Compare\/Contrast","Differences"],"Citation Start End":[[2333,2350]],"Functions Start End":[[2274,2331],[2628,2841]]}
{"Identifier":"2018ApJ...856..140H__Leary_et_al._2009_Instance_1","Paragraph":"The number of BHs, their mass distribution, and the number density are poorly known in NSCs. Theoretically, a single-mass distribution of objects forms a power-law density cusp around a massive object with n(r) \u221d r\u22121.75 (Bahcall & Wolf 1976), where n(r) is the number density, and r is the distance from the MBH. For multi-mass distributions, lighter and heavier objects develop shallower (\u221dr\u22121.5) and steeper cusps (typically \u221dr\u22122 to r\u22122.2, and r\u22123 in extreme cases), respectively (Bahcall & Wolf 1977; Freitag et al. 2006; Hopman & Alexander 2006b; Keshet et al. 2009; Aharon & Perets 2016). Recent observations of the stellar distribution in the Milky Way NSC identify a cusp with n(r) \u221d r\u22121.25 (Gallego-Cano et al. 2018; Sch\u00f6del et al. 2018), which is consistent with the profile after a Hubble time (Baumgardt et al. 2018). As BHs are heavier than typical stars, they are expected to relax into the steeper cusps. The relaxation time of BH populations is much shorter: 0.1\u20131 Gyr (O\u2019Leary et al. 2009), although it can become much longer than that in the case of a shallow stellar density profile (Dosopoulou & Antonini 2017). In this work, we assume that the BH number density follows a cusp with either n(r) \u221d r\u22122 or r\u22123 in our two sets of calculations.The BH mass in the two cases is set arbitrarily to 107 and 4 \u00d7 106 M\u2299, respectively, and we refer to the two models as Bahcall\u2013Wolf\u2013like (BW) and GC examples. Note that here GC refers to the assumed MBH mass (Ghez et al. 2008) and the observed stellar distribution (see below). In reality, the cusp distribution varies with the BH mass. Thus, we have also generated initial conditions for the BW case so that \u03b2 in the number density distribution, n(r) \u221d r\u2212(3\/2+\u03b2), is calculated by \u03b2(m) = m\/4M0, where m is the binary mass, and M0 is the weight average mass (e.g., Alexander & Hopman 2009; Keshet et al. 2009; Aharon & Perets 2016). We found that, due to the stability conditions, the initial condition distribution does not change significantly. For the GC case, the initial conditions\u2019 distribution will change more significantly if we allow \u03b2 to vary with the mass. However, we keep \u03b2 = 3 to investigate the effects of a steep number density distribution on the rates. As we will show later in this work, the choice of the number density distribution will have a very limited effect on the merger rate.","Citation Text":["O\u2019Leary et al. 2009"],"Functions Text":["As BHs are heavier than typical stars, they are expected to relax into the steeper cusps. The relaxation time of BH populations is much shorter: 0.1\u20131 Gyr"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[985,1004]],"Functions Start End":[[829,983]]}
{"Identifier":"2017ApJ...845...86E__Soler_&_Terradas_2015_Instance_2","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"],"Functions Text":["The background Alfv\u00e9n or slow continuum can be due to the variation of the plasma density"],"Functions Label":["Background"],"Citation Start End":[[1430,1451]],"Functions Start End":[[1254,1343]]}
{"Identifier":"2017ApJ...850L..40A__Yang_et_al._2017_Instance_1","Paragraph":"Aided by the tight localization constraints of the three-detector network and the proximity of the GW source, multiple independent surveys across the EM spectrum were launched in search of a counterpart beyond the sGRB (Abbott et al. 2017c). Such a counterpart, SSS17a (later IAU-designated AT 2017gfo), was first discovered in the optical less than 11 hours after merger, associated with the galaxy NGC 4993 (Coulter et al. 2017a, 2017b), a nearby early-type E\/S0 galaxy (Lauberts 1982). Five other teams made independent detections of the same optical transient and host galaxy all within about one hour and reported their results within about five hours of one another (Allam et al. 2017; Arcavi et al. 2017a, 2017b; Lipunov 2017b; Tanvir & Levan 2017; Yang et al. 2017; Soares-Santos et al. 2017; Lipunov et al. 2017a). The same source was followed up and consistently localized at other wavelengths (e.g., Corsi et al. 2017; Deller et al. 2017a, 2017b, 2017c; Goldstein et al. 2017; Haggard et al. 2017a, 2017b; Mooley et al. 2017; Savchenko et al. 2017; Alexander et al. 2017; Haggard et al. 2017c; Goldstein et al. 2017; Savchenko et al. 2017). The source was reported to be offset from the center of the galaxy by a projected distance of about 10\u2033 (e.g., Coulter et al. 2017a, 2017b; Haggard et al. 2017a, 2017b; Kasliwal et al. 2017; Yang et al. 2017; Yu et al. 2017). NGC 4993 has a Tully\u2013Fisher distance of \u223c40 Mpc (Freedman et al. 2001; NASA\/IPAC Extragalactic Database164\n\n164\nThe NASA\/IPAC Extragalactic Database (NED) is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration.\n), which is consistent with the luminosity distance measurement from gravitational waves (\n\n\n\n\n\n Mpc). Using the Tully\u2013Fisher distance, the \u223c10\u2033 offset corresponds to a physical offset of \u22432.0 kpc. This value is consistent with offset measurements of sGRBs in other galaxies, though below the median value of \u223c3\u20134 kpc (Fong et al. 2010; Fong & Berger 2013; Berger 2014).","Citation Text":["Yang et al. 2017"],"Functions Text":["Five other teams made independent detections of the same optical transient and host galaxy all within about one hour and reported their results within about five hours of one another"],"Functions Label":["Background"],"Citation Start End":[[756,772]],"Functions Start End":[[489,671]]}
{"Identifier":"2018ApJ...856..136P__Burkhart_et_al._2010_Instance_1","Paragraph":"Depending on the specific driver, the characteristics of turbulence will then be imprinted within the ISM mainly as three-dimensional density and velocity fluctuations, and these fluctuations have been traditionally studied via correlation functions such as the spatial power spectrum (SPS) (e.g., Crovisier & Dickey 1983), \u0394-variance (e.g., Stutzki et al. 1998), and structure function (e.g., Padoan et al. 2002; Burkhart et al. 2015b). In particular, the SPS approach has been applied to observations of various Galactic and extragalactic environments (e.g., Plume et al. 2000; Dickey et al. 2001; Elmegreen et al. 2001; Burkhart et al. 2010; Combes et al. 2012; Zhang et al. 2012; Pingel et al. 2013), showing power spectral slopes \u03b2 roughly ranging from \u22122.7 to \u22123.7 depending on the tracers used (e.g., H i, carbon monoxide (CO), and dust). These slopes essentially provide information on the relative amount of structure as a function of spatial scale and can be compared with theoretical models of turbulence (mainly numerical simulations) to characterize turbulence cascade (e.g., Burkhart et al. 2010), to determine the influence of shocks (e.g., Beresnyak et al. 2005), to reveal the injection and dissipation scales of turbulent energy (e.g., Kowal & Lazarian 2007; Federrath & Klessen 2013; Chen et al. 2015), and to trace the evolution of MCs (e.g., Burkhart et al. 2015a). The proximity and abundance of multi-wavelength observations make MCs in the solar neighborhood an ideal laboratory for probing the impact of turbulence on their formation and evolution. In this paper, we focus on the Perseus MC, which is a nearby (\u223c300 pc; e.g., Herbig & Jones 1983; \u010cernis 1990), low-mass (\u223c2 \u00d7 104 M\u2299; e.g., Sancisi et al. 1974; Lada et al. 2010) cloud. Its star formation activities, as well as atomic and molecular gas content, have been extensively examined over the past decade (e.g., Ridge et al. 2006; J\u00f8rgensen et al. 2007; Pineda et al. 2008; Lee et al. 2012, 2014, 2015; Mercimek et al. 2017), revealing that the cloud consists of several individual dark and star-forming regions (e.g., B5, B1, B1E, IC 348, and NGC 1333) and is actively forming low- to intermediate-mass stars (see Bally et al. 2008 for a review).","Citation Text":["Burkhart et al. 2010"],"Functions Text":["In particular, the SPS approach has been applied to observations of various Galactic and extragalactic environments (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[623,643]],"Functions Start End":[[438,560]]}
{"Identifier":"2015ApJ...806...15Z___2015_Instance_1","Paragraph":"Since the launch of Swift (Gehrels et al. 2004) in 2004, the nature of jet breaks in gamma-ray burst (GRB) afterglows has become increasingly puzzling. There is strong evidence that the ejecta from the GRB central engine must be jet-like (Zhang & M\u00e9sz\u00e1ros 2004). Thus a collimation correction factor, \n\n\n\n\n\n, where \n\n\n\n\n\n is the jet opening angle with typical values of 5\u00b0\u201310\u00b0, can be applied to relieve the energy budget problem, so that a typical GRB energy is \n\n\n\n\n\n \n\n\n\n\n\n erg, where \n\n\n\n\n\n is the isotropic equivalent energy release in gamma-rays. Such collimated ejecta expand outward relativistically with Lorentz factors \u0393 of several hundred initially. Internally, the ejecta release their energy through internal shocks (Rees & Meszaros 1994; Kobayashi et al. 1997; Daigne & Mochkovitch 1998), or magnetic dissipation processes (e.g., the ICMART model; Zhang & Yan 2011) or photospheric dissipation (e.g., Guiriec et al. 2010, 2011, 2015; Lazzati & Begelman 2010; Ryde et al. 2010, 2011; Pe\u2019er & Ryde 2011) and produce the prompt gamma-ray emission of GRBs. Externally, the ejecta are further decelerated by an ambient medium (e.g., a constant density interstellar medium (ISM); or a stellar-wind environment with density inversely proportional to distance squared) and produce long-term broadband afterglows through external shocks (see e.g., Gao et al. 2013, for a review). Due to relativistic beaming, only a portion of the radiation from the ejecta front surface, which is within a cone of half-opening angle \n\n\n\n\n\n, can be observed (Rhoads 1997, for reviews see Piran 2004; Granot 2007; van Eerten 2013). An unavoidable consequence of this general picture is that when the ejecta are decelerated to \n\n\n\n\n\n, the light curve should steepen because (1) the maximum observable portion of the ejecta (the cone of the whole jet with opening angle \n\n\n\n\n\n) is now smaller than that which is expected (a cone with half-opening angle \n\n\n\n\n\n and (2) the onset of lateral spreading of the ejecta, predicted to become noticeable in the observer frame around the same time (Rhoads 1999), causes the blast wave to decelerate further. Such a \u201cjet break\u201d in a GRB light curve is expected to behave achromatically because it only reflects the ejecta geometry, under the assumption that the afterglow emission regions and mechanisms do not change in different spectral regimes (Rhoads 1999; Sari 1999; Huang et al. 2000; Granot et al. 2002; see also reviews by M\u00e9sz\u00e1ros 2002; Piran 2004; Zhang & M\u00e9sz\u00e1ros 2004). The achromaticity was apparently confirmed in the optical and near-IR band in a few cases of pre-Swift GRBs (Kulkarni et al. 1999; Harrison et al. 2001; Klose et al. 2004).","Citation Text":["Guiriec et al.","2015"],"Functions Text":["Internally, the ejecta release their energy through","or photospheric dissipation (e.g.,","and produce the prompt gamma-ray emission of GRBs."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[915,929],[942,946]],"Functions Start End":[[661,712],[880,914],[1016,1066]]}
{"Identifier":"2018ApJ...864..160R__Bhat_et_al._2013_Instance_1","Paragraph":"Figure 9 shows the components required for such a time-domain survey, with the PC beamformer as specified above. The required components are shown in four different colors. Visibilities computed in the uGMRT backend (GWB; marked in blue) at 1 ms time resolution are transferred to the the PC beamformer nodes (marked in orange) with an aggregate data rate of \u223c3 GB s\u22121. We aim to implement in-field phasing (Kudale & Chengalur 2017) using a sky model derived from the time-averaged visibilities in order to improve the coherence in phasing up to a baseline length of several kilometers. This optimizes the GMRT PA sensitivity beyond a central compact core (most current PA observations use only the antennas in the central square). In addition to deriving the phasing model, a baseline-based flag masking the bad baselines will also be generated in real time from these time-averaged visibilities. Coherent additions of these visibilities will result in 128 such visibility beams. The multi-DM search for single pulses (colored in yellow) on each of these visibility beams would need to be executed on a separate FRB cluster, followed by coincidence filtering to remove spurious events (Bhat et al. 2013). We also propose recording these 128 beams with a 1 ms time resolution, giving a total data rate of 200 MB s\u22121 into a disk for a quasi-real-time search for pulsars using the same cluster. We note that the proposed 1 ms time resolution is sufficient to detect double neutron star systems, young pulsars, and normal pulsars, as well as objects like radio magnetars. Visibility buffers corresponding to candidate single-pulse events will be recorded at a 1 ms time resolution covering the full DM sweep time-range. For an event at a DM of 2000 pc cm\u22123, the total DM sweep time over 200 MHz band in uGMRT 300\u2013500 MHz band is \u223c50 s, which results in a 40 GB buffer size on each of the PC beam nodes. This means one can easily hold few buffers for accommodating the pipeline delay and flush them into storage based on the real-time triggers. These visibilities will be processed through the processing blocks (marked in green) for millisecond imaging localization at quasi-real-time. This block includes removal of dispersion delay, followed by a flagging and calibration pipeline and snapshot imaging. We note that part of this imaging pipeline for localizing time-domain events has already been demonstrated for the GHRSS Phase1 survey (Bhattacharyya et al. 2016).","Citation Text":["Bhat et al. 2013"],"Functions Text":["The multi-DM search for single pulses (colored in yellow) on each of these visibility beams would need to be executed on a separate FRB cluster, followed by coincidence filtering to remove spurious events"],"Functions Label":["Future Work"],"Citation Start End":[[1187,1203]],"Functions Start End":[[981,1185]]}
{"Identifier":"2017AandA...601A..72I__Kobayashi_&_Tanaka_2010_Instance_3","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"],"Functions Text":["In the obtained size distribution, erosive collisions are more important than catastrophic collisions (see Fig. 10 of","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."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1275,1298]],"Functions Start End":[[1157,1274],[1301,1890]]}
{"Identifier":"2020ApJ...897...94G__Stanway_&_Eldridge_2019_Instance_1","Paragraph":"It is interesting that the sources with suspected or confirmed Lyman continuum radiation at high redshifts are peculiar and rare, bright SFGs with rather hard ionizing spectra marked by high-ionization emission lines (e.g., N v, C iv, He ii, O iii, C iii). Their presence at high redshift can be hardly explained by pure SFGs, even requiring uncommon assumptions (e.g., large stellar rotation, binary stellar population, top-heavy initial mass function, extremely low metallicity) as discussed by a number of recent works (see Bowler et al. 2014; Bradley et al. 2014; Kehrig et al. 2015; Stark et al. 2015; Jaskot & Ravindranath 2016; Stark 2016; Senchyna et al. 2017, 2019, 2020; Berg et al. 2018; Nakajima et al. 2018; Chisholm et al. 2019; Jaskot et al. 2019; Le Fevre et al. 2019b; Nanayakkara et al. 2019; Schaerer et al. 2019; Stanway & Eldridge 2019). The majority of galaxies showing Lyman continuum emission (both at low z and at z \u223c 3\u20134) populate the upper end of the Baldwin, Phillips, and Terlevich (BPT) diagram (Baldwin et al. 1981) or occupy a region of high-ionization line ratio in between SFGs and AGNs, or mainly populated by AGNs (see, e.g., Figures 11 and 14 by Nakajima et al. 2018). Indeed, Le Fevre et al. (2019b) find a marginal 2\u03c3 detection in the X-ray stacking of strong C iii emitters at 2  z  4, consistent with the presence of low-luminosity AGNs. Interestingly, evidence recently emerged that local confirmed Lyman continuum emitters or low-z Green Peas, blue compact galaxies, Lyman break analogs, which are usually associated with reliable Lyman continuum candidates, are characterized by significant X-ray emission, not compatible with star formation activity but more plausibly powered by low-luminosity AGN activity or by a large population of high-mass X-ray binaries (Kaaret et al. 2017; Bao et al. 2019; Bluem et al. 2019; Latimer et al. 2019; Plat et al. 2019; Prescott & Sanderson 2019; Senchyna et al. 2019, 2020; Svoboda et al. 2019; Wu et al. 2019; Baldassare et al. 2020; Birchall et al. 2020; Dittenber et al. 2020). It is thus possible that pure stellar radiation from SFGs is a negligible source of H i ionizing radiation, and the bulk of Lyman continuum photons escaping at low and high z are produced instead by accretion onto supermassive black holes (SMBHs).","Citation Text":["Stanway & Eldridge 2019"],"Functions Text":["Their presence at high redshift can be hardly explained by pure SFGs, even requiring uncommon assumptions (e.g., large stellar rotation, binary stellar population, top-heavy initial mass function, extremely low metallicity) as discussed by a number of recent works (see"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[833,856]],"Functions Start End":[[257,526]]}
{"Identifier":"2017AandA...599A..55B__Shakura_1972_Instance_1","Paragraph":"When the characteristic time of variability of the mass flux along the accretion disk is longer than the relaxation time of the local disk equilibrium, it is possible to use the approximation of local equilibrium Shakura (1972), see also Bisnovatyi-Kogan (2011), to calculate the transient disk structure. The equilibrium along a radius of the accretion disk around a star with a mass M is determined by the Keplerian rotational velocity \u03a9K(1)\\begin{equation} \\Omega=\\Omega_{\\rm K}=\\left(\\frac{GM}{r^3}\\right)^{1\/2}\\cdot \\label{omega} \\end{equation}\u03a9=\u03a9K=GMr31\/2\u00b7Writing the equation of the vertical equilibrium in approximate algebraic form, we obtain (2)\\begin{equation} h=\\sqrt 2 \\frac{v_{\\rm s}}{\\Omega}, \\label{h} \\end{equation}h=2vs\u03a9,where \\hbox{$v_{\\rm s}=\\sqrt{P\/\\rho}$}vs=P\/\u03c1 is the speed proportional to the sound velocity, P and \u03c1 are the (gas + radiation) pressure and density at the symmetry plane of the accretion disk, and h is the semi-thickness of the accretion disk. The specific angular momentum l of the matter in the accretion disk is connected to the rotation velocity as(3)\\begin{equation} l=r\\,v_\\phi=r^2\\Omega. \\label{l} \\end{equation}l=r\u2009v\u03c6=r2\u03a9.The mass flux through the disk at radius r is connected to the radial velocity vr as (4)\\begin{equation} \\dot M=-4\\pi h\\rho r v_r, \\quad \\dot M>0,\\quad v_r<0. \\label{mflux} \\end{equation}M\u0307=\u22124\u03c0h\u03c1rvr,\u2001M\u0307>0,\u2001vr0.We use an \u03b1 approximation for the turbulent viscosity (Shakura 1972) when the (r\u03c6) component of the stress tensor tr\u03c6 is written as (5)\\begin{equation} t_{r\\phi}=\\alpha\\, P, \\label{trphi} \\end{equation}tr\u03c6=\u03b1\u2009P,where the phenomenological non-dimensional parameter \u03b1 \u2264 1. The condition of stationarity of the angular momentum, in which the outward viscous radial flux of the angular momentum is balanced by the angular momentum of the inward flux of the mass, is written as (see, e.g., Bisnovatyi-Kogan 2011) (6)\\begin{equation} r^2 h \\alpha P=\\frac{\\dot M}{4\\pi}(l-l_{\\rm in}). \\label{angmom} \\end{equation}r2h\u03b1P=M\u03074\u03c0(l\u2212lin).The main input into the time lag comes from the outer regions of the disk with l \u226b lin. Then we have from Eqs. (4) and (6) the expression for the radial velocity in the form (7)\\begin{equation} v_r=-\\alpha\\frac{v_{\\rm s}^2}{v_\\phi}\\cdot \\label{vr} \\end{equation}vr=\u2212\u03b1vs2v\u03c6\u00b7We also define the surface density \u03a3, and write Eq. (6) in light of Eq. (3), using condition l \u226b lin, in the form (8)\\begin{equation} \\Sigma=2\\rho h,\\quad \\dot M \\Omega=4\\pi \\alpha P h. \\label{sigma} \\end{equation}\u03a3=2\u03c1h,\u2001M\u0307\u03a9=4\u03c0\u03b1Ph.The equation of the local thermal balance in the accretion disk, when the heat produced by viscosity Q+ is entirely emitted through the sites of the optically thick accretion disk with a total flux Q\u2212, at l \u226b lin is written as (see, e.g., Bisnovatyi-Kogan 2011)(9)\\begin{equation} \\frac{3}{2}\\dot M \\Omega^2=\\frac{16\\pi ac T^4}{3\\varkappa \\Sigma}\\cdot \\label{theq} \\end{equation}32M\u0307\u03a92=16\u03c0acT43\u03f0\u03a3\u00b7Here T is the temperature in the symmetry plane of the accretion disk, a is the constant of the radiation energy density, c is the speed of light, and \u03f0 is the Thompson (scattering) opacity of the matter. ","Citation Text":["Shakura (1972)"],"Functions Text":["When the characteristic time of variability of the mass flux along the accretion disk is longer than the relaxation time of the local disk equilibrium, it is possible to use the approximation of local equilibrium","see also Bisnovatyi-Kogan (2011), to calculate the transient disk structure."],"Functions Label":["Uses","Uses"],"Citation Start End":[[213,227]],"Functions Start End":[[0,212],[229,305]]}
{"Identifier":"2019MNRAS.482..194B__Choi_et_al._2015_Instance_1","Paragraph":"Observationally, it is a big challenge to conclude whether AGN\u2019s feedback could regulate star formation or not, and how it regulates star formation. On one hand, if strong outflows emerge, they could clear out the star-forming gas to suppress star formation (Alexander & Hickox 2012; Garc\u00eda-Burillo et al. 2014; Alatalo et al. 2015; Hopkins et al. 2016; Wylezalek & Zakamska 2016). The heating by jets propagating through the galaxies could prevent gas from cooling and cut-off the gas supply for further star formation (Karouzos et al. 2014; Choi et al. 2015). On the other hand, the outflows and jets interact with the gas in host galaxies and compress it to trigger new star formation (Silk 2013; Zubovas et al. 2013; Zubovas & Bourne 2017). In fact observations show either no or positive relationships between star formation rates and SMBH accretion rates but no negative trends are seen (Shi et al. 2007, 2009; Baum et al. 2010; Xu et al. 2015; Zhang et al. 2016; Mallmann et al. 2018). Whether AGN\u2019s feedback plays the role may also depends on the spatial scale that observations could resolve and time-scale that the observed tracers could probe (Harrison et al. 2012; Cresci et al. 2015; Feruglio et al. 2015). For example, radiation from AGNs nearly instantaneously impact the surrounding ISM while the attenuation from ISM probably limits their impact to the nuclear regions (Roos et al. 2015). Outflows or jets travel slowly and may be decelerated after interactions with ambient gas, which delays their effects on star formation at large distances from the nuclei (Harrison 2017; Harrison et al. 2018). Feedback by jets or outflows on ISM also depends on their orientation relative to the dusty torus, making their effects on star formation to be anisotropic. The short duty cycle of AGNs could also make feedback by radiation from AGNs temporally variable in strength. Case studies of individual AGNs find evidence of coexistence of positive and negative feedback on star formation (Zinn et al. 2013), suggesting the complicated nature of AGN feedback.","Citation Text":["Choi et al. 2015"],"Functions Text":["The heating by jets propagating through the galaxies could prevent gas from cooling and cut-off the gas supply for further star formation"],"Functions Label":["Background"],"Citation Start End":[[543,559]],"Functions Start End":[[382,519]]}
{"Identifier":"2019MNRAS.490.1870L__Guo_et_al._2015_Instance_2","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)$$\\begin{eqnarray*}\r\n\\hat{\\xi }_\\ell &amp;=&amp; \\hat{n}_{\\rm gal}^{-2} \\sum \\limits _{i = {\\rm c}, {\\rm s}} \\sum \\limits _{j = {\\rm c}, {\\rm s}} \\sum \\limits _{k = 1}^{n_{\\rm bins}} \\sum \\limits _{l = 1}^{n_{\\rm bins}} \\Big [ N_{{\\rm h}, k} N_{{\\rm h}, l} \\langle N_i | M_k, c_k \\rangle \\nonumber \\\\\r\n&amp;&amp;\\times \\,\\langle N_j | M_l, c_l \\rangle \\xi _{\\ell , kl}^{ij} \\Big ] \\, ,\r\n\\end{eqnarray*}$$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"],"Functions Text":["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"],"Functions Label":["Uses"],"Citation Start End":[[2541,2556]],"Functions Start End":[[2385,2540]]}
{"Identifier":"2021AandA...652A.114P__Nakariakov_&_Verwichte_2005_Instance_1","Paragraph":"Multiple solar missions, such as the Solar Dynamics Observatory (SDO) and the Interface Region Imaging Spectrograph (IRIS), have shown that a diversity of waves occur in the solar atmosphere (Jess et al. 2009; McIntosh et al. 2011; Okamoto & De Pontieu 2011). The various wave types that occur include Alfv\u00e9n waves (Alfv\u00e9n 1942). These waves are transverse magnetohydrodynamic (MHD) waves that travel along the magnetic field lines. Alfv\u00e9n waves were reported to be present in both the photosphere and chromosphere (Srivastava et al. 2017; Baker et al. 2021). As they pass by, they modify the transverse magnetic field and velocity components but, do not alter the gas pressure or the mass density (Nakariakov & Verwichte 2005), at least in the linear limit and in a homogeneous background plasma. A thorough understanding of Alfv\u00e9n waves is essential because they could be a part of the solution to the major problems of heliophysics, such as the solar coronal heating and wind acceleration (Uchida & Kaburaki 1974; Ofman 2010). Recent theoretical research revealed that Alfv\u00e9n waves can carry enough energy to heat the solar corona (Yang & Xiang 2016). However, the details of the mechanism(s) of the thermal energy release related to their dissipation remain unknown. One potential candidate for that may be associated with ion\u2013neutral collisions (Soler et al. 2017). Piddington (1956), Osterbrock (1961), and Haerendel (1992) were the first to study ion\u2013neutral collisions, but they did not find that this interaction affects the chromospheric temperature. Ballester et al. (2018) showed that ambipolar diffusion leads to substantial chromospheric heating, and Zaqarashvili et al. (2013) derived a dispersion relation for two-fluid Alfv\u00e9n waves and confirmed that the damping of Alfv\u00e9n waves resulting from the ion\u2013neutral collisions is quite significant. Khomenko (2017), based on a two-fluid model, stated that the presence of neutrals affects the solar atmosphere. The effect of ion\u2013neutral interactions is expected to influence the energy balance of the chromosphere. Zaqarashvili et al. (2013) also confirmed that low- and high-frequency photospheric Alfv\u00e9n waves might not reach the solar corona because ion\u2013neutral collisions damp them very efficiently in the upper chromosphere. According to Song & Vasyli\u016bnas (2011), the rate of Alfv\u00e9n wave damping varies with magnetic field strength and wave frequency. For a strong magnetic field, wave damping is low. Low-frequency waves are also weakly damped, and so there is a chance to detect low-frequency Alfv\u00e9n waves in the solar corona under the condition of a strong magnetic field.","Citation Text":["Nakariakov & Verwichte 2005"],"Functions Text":["As they pass by, they modify the transverse magnetic field and velocity components but, do not alter the gas pressure or the mass density","at least in the linear limit and in a homogeneous background plasma."],"Functions Label":["Background","Background"],"Citation Start End":[[699,726]],"Functions Start End":[[560,697],[729,797]]}
{"Identifier":"2021MNRAS.504.5840F__Eriksen_et_al._2007_Instance_2","Paragraph":"The standard cosmological model stands on the shoulders of a fundamental assumption: that the universe is statistically homogeneous and isotropic on the largest scales. This assumption has been thoroughly tested over the last years both with cosmic microwave background (CMB) and Large-scale structure data. In particular, the analysis of CMB data, most notably from the Wilkinson Microwave Anisotropy Probe (WMAP; Bennett et al. 2013) and Planck (Planck Collaboration I 2020) experiments, has not yet provided conclusive evidence for the hypothesis of Cosmological Isotropy (Eriksen et al. 2004, 2007; Hajian, Souradeep & Cornish 2005; Land & Magueijo 2007; Hansen et al. 2009; Samal et al. 2009; see also Planck Collaboration VII 2020 and references therein). Moreover, Galactic foreground contamination or known systematic effects in the data alone can not explain the observed CMB \u2018anomalies\u2019, i.e. large-scale deviations from the concordance Lambda cold dark matter (\u039bCDM) model (see e.g. Rassat et al. 2014; see Planck Collaboration VII 2020 for a recent overview). Power asymmetry from CMB data has also been a matter of intense debate and scrutiny (Gazta\u00f1aga, Fosalba & Elizalde 1998; Eriksen et al. 2007; Lew 2008; Hoftuft et al. 2009; Paci et al. 2010; Axelsson et al. 2013; Shaikh et al. 2019, see also Dai et al. 2013 for a comprehensive discussion and references therein), and evidence has been reported that this could source deviations from isotropy on cosmological scales (Hansen et al. 2009). However, a more recent analysis based on Planck data finds no evidence for such power asymmetry when all scales are taken into account (Quartin & Notari 2015). This is in qualitative agreement with the latest results from the Planck Collaboration analysis (Planck Collaboration VII 2020) where they conclude that the observed power asymmetry is not robust to foreground contamination or systematic residuals. It is important to note that previous analysis have concentrated on quantifying potential deviations from statistical isotropy using a statistical prior. First analyses using WMAPdata looked for the direction of maximal asymmetry in the sky, thus quantifying anisotropy for a given preferred direction (Hansen et al. 2009). In turn, this led to proposing a particular angular distribution of power in the sky to simply capture the observed anisotropy, such as the so-called \u2018dipole anisotropy\u2019 modulation (Prunet et al. 2005; Gordon 2007). This same model has been further constrained with Planck data (Planck Collaboration XXIII 2014; Planck Collaboration XVI 2016; Planck Collaboration VII 2020; Aiola et al. 2015; Mukherjee et al. 2016). Alternatively, a recent analysis (Ho & Chiang 2018) focuses on quantifying possible CMB peak shifts across the sky, finding significant variations, but they attribute this behaviour to possible systematic effects or the solar dipole. Complementary evidence for cosmological anisotropy has been investigated using probes of the low-redshift universe (see Colin et al. 2011; Secrest et al. 2021 and references therein).","Citation Text":["Eriksen et al. 2007"],"Functions Text":["Power asymmetry from CMB data has also been a matter of intense debate and scrutiny"],"Functions Label":["Background"],"Citation Start End":[[1193,1212]],"Functions Start End":[[1072,1155]]}
{"Identifier":"2020ApJ...899..147F__H\u00f6rst_et_al._2018_Instance_1","Paragraph":"Recent observations of transit spectra of hot Jupiter atmospheres show limited spectral modulation due to H2O that has been largely interpreted as the indicator of the presence of aerosols (Barstow et al. 2016; Iyer et al. 2016; Sing et al. 2016; Pinhas et al. 2019). Whether these aerosols are condensate clouds of photochemical organic aerosols or other refractory materials remains unknown. Although thermochemical equilibrium models predict the formation of condensate clouds with various composition in these hot atmospheres (Lecavelier Des Etangs et al. 2008; Lee et al. 2015; Parmentier et al. 2016), recent laboratory works highlighted that photochemistry could strongly affect the composition of exoplanet atmospheres and lead to the formation of aerosols in a variety of conditions, including the ones encountered in hot Jupiters (H\u00f6rst et al. 2018; Fleury et al. 2019; He et al. 2019, 2018a, 2018b). These photochemical aerosols could represent another source of opacity to explain some of the observed transit spectra of hot Jupiter atmospheres, e.g., of HD 189733 b (Lavvas & Koskinen 2017). On the other hand, the bulk elemental ratio can also drastically affect the molecular composition of these atmospheres. In the external layers (region with pressure 1 bar) of atmospheres with temperatures higher than 1000 K, carbon preferentially bonds with oxygen to form CO, and the excess of oxygen bonds with hydrogen to form H2O when the C\/O ratio is 1. At a higher CO ratio \u2265 1, CO remains an abundant species but the water mixing ratio decreases (Lodders & Fegley 2002; Moses et al. 2013; Venot et al. 2015; Heng & Lyons 2016; Tsai et al. 2017; Goyal et al. 2018; Drummond et al. 2019). For this reason, another explanation for the low spectral modulation due to water observed in some hot Jupiter atmospheres is that these atmospheres have low H2O abundances presumably reflecting high C\/O ratios (Madhusudhan et al. 2011; Madhusudhan 2012). However, the existence of such \u201ccarbon-rich\u201d exoplanets continues to be debated. The first analysis of the hot Jupiter WASP-12b observations suggested a C\/O ratio > 1 (Madhusudhan et al. 2011), but another study found a C\/O ratio  1 using another approach (Kreidberg et al. 2015), leaving the question of the C\/O ratio in WASP-12b\u2019s atmosphere open. In addition, a recent survey suggests that the carbon enrichment of hot Jupiter atmospheres compared to their host stars may be common, but uncertainties on C\/O measurements in exoplanet atmospheres are large and prevent a firm conclusion from being reached (Brewer et al. 2017).","Citation Text":["H\u00f6rst et al. 2018"],"Functions Text":["Although thermochemical equilibrium models predict the formation of condensate clouds with various composition in these hot atmospheres","recent laboratory works highlighted that photochemistry could strongly affect the composition of exoplanet atmospheres and lead to the formation of aerosols in a variety of conditions, including the ones encountered in hot Jupiters"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[841,858]],"Functions Start End":[[394,529],[608,839]]}
{"Identifier":"2020MNRAS.494.5134L__Taylor_1922_Instance_1","Paragraph":"The gas velocity is unknown since no exact analytic solution for turbulence in a disc \u2013 and turbulence in general \u2013 are known. However, statistical properties of turbulence can be inferred from laboratory, numerical experiment or theory, and turbulent fluctuations can be modelled using stochastic processes, independently from the origin of the turbulence itself. In a seminal study, Thomson (1987) proved that the only expression of vg that is consistent with Kolmogorov turbulence and the hydrodynamical equations is\n(16)$$\\begin{eqnarray*}\r\n\\frac{\\mathrm{d}v_{\\rm g}}{\\mathrm{d}t} = -\\frac{v_{\\rm g}}{t_{\\rm e}} + \\frac{\\sqrt{D}}{t_{\\rm e}}\\dot{w} ,\r\n\\end{eqnarray*}$$where te denotes the Lagrangian time-scale of the turbulence, D is the turbulent diffusivity (in units m2s\u22121). w is a Wiener process, such that its derivative is a white noise such that\n(17)$$\\begin{eqnarray*}\r\n\\left\\langle \\dot{w} (t) \\right\\rangle = 0,\r\n\\end{eqnarray*}$$(18)$$\\begin{eqnarray*}\r\n\\left\\langle \\dot{w} (t) \\, \\dot{w} (t^{\\prime }) \\right\\rangle = \\delta (t - t^{\\prime }) ,\r\n\\end{eqnarray*}$$where \u03b4 denotes the Dirac distribution and the notation \u00b7 > is the expectation operator (see also Sawford 1984; Wilson & Sawford 1996). Equation (16) describes turbulent fluctuations from a Lagrangian point of view (Taylor 1922). From equation (16), the gas velocity can be rewritten\n(19)$$\\begin{eqnarray*}\r\nv_{\\rm g} = \\zeta \\left(t,t_{\\rm e},D \\right) ,\r\n\\end{eqnarray*}$$where \u03b6 is a stationary Ornstein\u2013Uhlenbeck process defined by\n(20)$$\\begin{eqnarray*}\r\n\\left\\langle \\zeta (t,t_{\\rm e},D) \\right\\rangle = 0,\r\n\\end{eqnarray*}$$(21)$$\\begin{eqnarray*}\r\n\\left\\langle \\zeta (t,t_{\\rm e},D) \\, \\zeta (t^{\\prime },t_{\\rm e},D) \\right\\rangle = \\frac{D}{2 t_{\\rm e}} \\rm {e}^{-\\frac{\\vert t - t^{\\prime } \\vert }{t_{\\rm e}}} .\r\n\\end{eqnarray*}$$Equation (16) defines a model of turbulence with two parameters, D and te. In discs, te is typically of order one orbital period, since turbulent vortices are stretched out by differential rotation in a few orbits (e.g. Beckwith, Armitage & Simon 2011). From equation (21), D is related to the autocorrelation of the turbulent noise according to\n(22)$$\\begin{eqnarray*}\r\nD = 2 \\int _{0}^{+\\infty } \\left\\langle v_{\\rm g}\\left(0 \\right) v_{\\rm g}\\left(t \\right)\\right\\rangle \\mathrm{d}t .\r\n\\end{eqnarray*}$$Equation (22) can alternatively be seen as a definition of the turbulent diffusivity, useful in practice to measure D in numerical simulations. The Wiener\u2013Khinchin theorem ensures that the power spectrum of the turbulent velocity field S(\u03c9) is the Fourier transform of this autocorrelation function, i.e.\n(23)$$\\begin{eqnarray*}\r\nS \\left(\\omega \\right) = \\frac{1}{2 \\pi } \\int _{-\\infty }^{+\\infty } \\mathrm{e}^{- i \\omega t} \\left\\langle v_{\\rm g}\\left(0 \\right) v_{\\rm g}\\left(t \\right)\\right\\rangle \\mathrm{d}t = \\frac{D}{2\\pi \\left(1 + \\omega ^2 t_{\\rm e}^2 \\right)} .\r\n\\end{eqnarray*}$$Thus, in the inertial subrange ($\\omega ^2 t_{\\rm e}^2 \\gg 1$), we have S(\u03c9) \u221d \u03c9\u22122, whose equivalent in the wavelength space is $\\tilde{S}(k)\\propto k^{-5\/3}$ (Batchelor 1950). From equation (23), the standard deviation of the velocity fluctuation \u03c3 is\n(24)$$\\begin{eqnarray*}\r\n\\sigma ^{2} \\equiv \\int _{-\\infty }^{+\\infty } S\\left(\\omega \\right) \\mathrm{d} \\omega = \\frac{D}{2t_{\\rm e}} .\r\n\\end{eqnarray*}$$Physically, equation (24) is a turbulent fluctuation\u2013dissipation theorem.","Citation Text":["Taylor 1922"],"Functions Text":["Equation (16) describes turbulent fluctuations from a Lagrangian point of view"],"Functions Label":["Uses"],"Citation Start End":[[1297,1308]],"Functions Start End":[[1217,1295]]}
{"Identifier":"2019ApJ...871...86K__Friesen_et_al._2014_Instance_1","Paragraph":"On the basis of the 1.3 mm continuum image made from the SMA archival data, Nakamura et al. (2012) found conspicuous substructures inside a prestellar core in the Oph B2 region, which had been previously identified as a single core (B2-N5) in single-dish molecular line (N2H+ (J = 1\u22120)) observations (Friesen et al. 2010). The substructures consist of several small condensations, and their typical mass and size are around 0.05 M\u2299 and 500 au, respectively. The mass is comparable to or larger than the local Jeans mass of 0.04 M\u2299, and thus the self-gravity of the condensations appears to play an important role in their dynamics. Similarly, Kamazaki et al. (2001) found small condensations of 1000 au scale inside the SM1 core in the Oph A region based on 3 mm dust continuum observations with the Nobeyama Millimeter Array (see also Nakamura et al. 2012; Friesen et al. 2014). The masses of the condensations are around 0.01\u22120.1 M\u2299, comparable to or larger than the local Jeans mass. Recently, Kirk et al. (2017) carried out comprehensive survey observations with ALMA Band 3 toward 60 dense cores that were identified in Oph with SCUBA at the James Clerk Maxwell telescope (JCMT). They found 38 compact emission structures of \u223c100 au size within the dense cores. On the other hand, on the basis of 3 mm continuum observations with CARMA, Schnee et al. (2010) found no significant substructures inside prestellar cores in Perseus (see also Dunham et al. 2016), although their target cores are in relatively isolated environments. These observations suggest that internal structures of prestellar cores and their physical properties may strongly depend on cloud environment. It still remains unclear, however, whether or not they are really different between clustered environments and isolated environments. As a step toward a more comprehensive understanding of the star formation process, it is important to characterize better the substructures inside dense cores in cluster-forming regions. In the present paper, we further investigate the internal structures of the Oph B2 region using Cycle 2 ALMA data at a spatial resolution of \u223c3\u2033, corresponding to \u223c410 au at Oph distance.","Citation Text":["Friesen et al. 2014)"],"Functions Text":["Similarly, Kamazaki et al. (2001) found small condensations of 1000 au scale inside the SM1 core in the Oph A region based on 3 mm dust continuum observations with the Nobeyama Millimeter Array (see also","The masses of the condensations are around 0.01\u22120.1 M\u2299, comparable to or larger than the local Jeans mass."],"Functions Label":["Background","Background"],"Citation Start End":[[858,878]],"Functions Start End":[[632,835],[880,986]]}
{"Identifier":"2021MNRAS.501.2522J__Mukherjee_&_Paul_2004_Instance_2","Paragraph":"GX 301-2 is an HMXB consisting of a highly magnetized (B \u223c 4 \u00d7 1012\u2009G, or even larger Doroshenko et al. 2010) pulsar and a B-type hyper-giant star Wray 977 (Vidal 1973; Kaper et al. 1995; Staubert et al. 2019). According to modelling of high-resolution optical spectra, Wray 977 has a mass of 43 \u00b1 10 $\\, \\mathrm{M}_{\\odot }$, a radius of 62 R\u2299, and looses mass through powerful stellar winds at a rate of $\\sim \\! 10^{-5}\\, \\mathrm{M}_\\odot \\, {\\rm yr}^{-1}$ with terminal velocity of 300\u2009$\\rm km\\, s^{-1}$ (Kaper, van der Meer & Najarro 2006). The system is highly eccentric (e \u223c 0.46), with an orbital period of \u223c41.5\u2009d, and exhibits strong variation of the X-ray flux with orbital phase (Koh et al. 1997; Doroshenko et al. 2010). In particular, periodic outbursts at the orbital phase \u223c1.4\u2009d before the periastron passage (Sato et al. 1986), and a fainter one near the apastron passage are observed (Pravdo et al. 1995). The broad-band X-ray spectrum is orbital phase-dependent and can be approximately described as a power law with a high-energy cutoff and a cyclotron resonant scattering feature around 40\u2009keV (Kreykenbohm et al. 2004; Mukherjee & Paul 2004; La Barbera et al. 2005; Doroshenko et al. 2010; Suchy et al. 2012; Islam & Paul 2014; F\u00fcrst et al. 2018; Nabizadeh et al. 2019). During the periastron flares, the source exhibits strong variability with an amplitude of up to a factor of 25, reaching a few hundreds mCrab in the energy band of 2\u201310\u2009keV (e.g. Rothschild & Soong 1987; Pravdo et al. 1995). The flares are accompanied by the variability of the equivalent hydrogen column density ($\\rm \\mathit{ N}_{\\rm H}$) and of the fluorescent iron lines, which is believed to be associated with clumpiness of the stellar wind, launched from the donor star (Mukherjee & Paul 2004). We note the clumpiness in this paper refers to any inhomogeneities in the stellar wind\/stream, which are higher density regions, regardless of its specific formation mechanisms. On the other hand, F\u00fcrst et al. (2011) reported a long XMM\u2013Newton observation in GX 301-2 around its periastron, which also exhibits systematic variations of the flux and $\\rm \\mathit{ N}_{H}$ at a time-scale of a few kiloseconds. Several wind accretion models, consisting of stellar winds and a gas stream, were proposed to explain the observed flares (e.g. Haberl 1991; Leahy 1991; Leahy & Kostka 2008; M\u00f6nkk\u00f6nen et al. 2020).","Citation Text":["Mukherjee & Paul 2004"],"Functions Text":["The flares are accompanied by the variability of the equivalent hydrogen column density ($\\rm \\mathit{ N}_{\\rm H}$) and of the fluorescent iron lines, which is believed to be associated with clumpiness of the stellar wind, launched from the donor star"],"Functions Label":["Background"],"Citation Start End":[[1772,1793]],"Functions Start End":[[1519,1770]]}
{"Identifier":"2015ApJ...809...44Z__Jang_et_al._2014_Instance_1","Paragraph":"Admittedly, although the prediction success rate of SPM3 is greatly improved compared with the previous SPM2 model, its improvements in the shock TT prediction are limited. Many factors could potentially cause this. On one hand, the sample events used in this study are CMEs from Solar Cycle 23 when SOHO was the only spacecraft tracking their movements in the sky-plane. The CME speed derived in this way is the projected speed, which does not represent the propagation speed of the CME along the Sun\u2013Earth direction. Large uncertainties in VCME restrict further improvements to the TT prediction of the shock. Possible solutions to this restriction could include adopting the CME\u2019s radial speed, which is derived from models (such as cone models) based on single spacecraft observations (Jang et al. 2014). Another solution involves estimating the initial geometry and three-dimensional (3D) speeds of CMEs based on observations from multiple spacecraft (STEREO, SOHO; Kilpua et al. 2012; Gopalswamy et al. 2013; Lee et al. 2013). On the other hand, the input parameters used in this study were obtained when the disturbances propagated near the Sun. Therefore, the lead time of SPM3's prediction is very long, nearly the whole TT of the disturbance from the Sun to Earth, as the model is analytic and thus requires no running time. Models based on heliospheric image data (STEREO HIs and SMEI) can provide more accurate predictions for arrival times but with shorter lead times (Colaninno et al. 2013; Mishra & Srivastava 2013; M\u00f6stl et al. 2014). For example, Webb (2013) applied the Tapping\u2013Howard model (Tapping & Howard 2009) to predict the arrival time at Earth of the 2011 February 15 CME event based on HI and\/or SMEI observations, and the corresponding prediction accuracy could be within an hour. However, the prediction\u2019s lead time was only several hours. Kilometric Type II radio burst emission can also be used to track shock dynamics in the inner heliosphere and provide shock arrive time predictions (Corona-Romero et al. 2013; Xie et al. 2013). Further prediction models considering these factors should be developed based on the events of Solar Cycle 24. This is the next goal of our research.","Citation Text":["Jang et al. 2014"],"Functions Text":["Large uncertainties in VCME restrict further improvements to the TT prediction of the shock. Possible solutions to this restriction could include adopting the CME\u2019s radial speed, which is derived from models (such as cone models) based on single spacecraft observations"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[790,806]],"Functions Start End":[[519,788]]}
{"Identifier":"2019MNRAS.490.5478W__Winter,_Booth_&_Clarke_2018c_Instance_1","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, Booth & Clarke 2018c"],"Functions Text":["The influence of tidal truncation is therefore limited to stellar multiples","or during the decay of higher order multiplicity"],"Functions Label":["Background","Background"],"Citation Start End":[[955,983]],"Functions Start End":[[764,839],[905,953]]}
{"Identifier":"2019MNRAS.490.5478W__Winter_et_al._2018b_Instance_3","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"],"Functions Text":["Many stars in the solar neighbourhood are born in regions where UV fields are sufficient to significantly shorten disc lifetimes"],"Functions Label":["Background"],"Citation Start End":[[1989,2008]],"Functions Start End":[[1837,1965]]}
{"Identifier":"2020MNRAS.494.3413T__Shidatsu_&_Done_2019_Instance_1","Paragraph":"The existence of winds is shown by blueshifted absorption lines from highly ionized ions. These are only seen in soft state but not in hard state (Ponti et al. 2012), anticorrelated with the radio jet which is seen in the hard state but not in the soft. This was thought to be evidence that the wind was magnetically driven by the same field as was responsible for the jet, but in a different geometric configuration (Miller et al. 2012). However, in Tomaru et al. (2019, hereafter Paper I) we show instead that thermally driven winds can explain this switch (see also Done, Tomaru & Takahashi 2018; Shidatsu & Done 2019). Thermal driving produces a wind by irradiation from the central source heating the surface of accretion disc up to the Compton temperature ($T_\\text{IC} \\sim 10^7 \\!-\\!10^8\\, \\text{K}$), which is hot enough for its thermal energy to overcome the gravity at large radii. The characteristic radius at which the wind can be launched is called the Compton radius, defined by RIC = \u03bcmpGM\/kTIC \u223c 105 \u2212 106Rg (Begelman, McKee & Shields 1983). Paper I show the first modern radiation hydrodynamic simulations of thermal (and thermal-radiative) winds designed to investigate the switch in wind properties between the hard and soft states changing illumination spectra. These simulations were tailored to the BHB system H1743\u2212332, where there is Chandra high-resolution data in both states giving detailed spectral information on the wind or its absence (Miller et al. 2012; Shidatsu & Done 2019). They incorporate radiation force on the electrons, both bound and free, as they show that this is important factor driving the escape of the thermal wind in the fairly high Eddington fraction (L\/LEdd \u223c 0.2\u20130.3), fairly low Compton temperature (TIC \u223c 0.1 \u00d7 108 K) characteristics of the soft state. The only other modern hydrodynamic simulation of thermal winds (e.g. Luketic et al. 2010; Higginbottom & Proga 2015; Higginbot et al. 2016) has not included radiation pressure, which is important in setting the velocity structure for L \u2265 0.3LEdd as required here (Paper I).","Citation Text":["Shidatsu & Done 2019"],"Functions Text":["However, in Tomaru et al. (2019, hereafter Paper I) we show instead that thermally driven winds can explain this switch (see also"],"Functions Label":["Background"],"Citation Start End":[[600,620]],"Functions Start End":[[439,568]]}
{"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"],"Functions Text":["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"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[2812,2829]],"Functions Start End":[[2536,2810]]}
{"Identifier":"2019AandA...626A..64H__Grinberg_et_al._2017_Instance_1","Paragraph":"Line driven winds are not expected to be smooth flows, but show strong density perturbations or \u201cclumps\u201d (Owocki et al. 1988; Feldmeier et al. 1997; Puls et al. 2006, 2008; Oskinova et al. 2012; Sundqvist & Owocki 2013). In X-ray binaries, thedensity contrast could even be further enhanced by the interaction between the wind and the strong X-rays from the compact object (Blondin 1994; Blondin & Woo 1995; Manousakis & Walter 2011, 2015, and references therein). For Vela X-1 and Cyg X-1 it has been estimated that more than 90% of the wind mass is contained in less than 10% of the wind volume (Sako et al. 1999; Rahoui et al. 2011). When the line of sight to the compact object passes through one of these clumps, X-rays are absorbed by the moderately ionized material in the clump, leading to a so-called dipping event. This is also observed for other sources (Hemphill et al. 2014; Grinberg et al. 2017). During the hard state of Cyg X-1, such short-term dipping events are observed predominantly during the upper conjunction of the black hole, i.e., when the line of sight passes through the densest region of the stellar wind and is most likely to pass through a clump (Li & Clark 1974; Mason et al. 1974; Parsignault et al. 1976; Pravdo et al. 1980; Remillard & Canizares 1984; Kitamoto et al. 1984; Ba\u0142uci\u0144ska-Church et al. 2000; Feng & Cui 2002; Poutanen et al. 2008; Hanke et al. 2009; Mi\u0161kovi\u010dov\u00e1 et al. 2016; Grinberg et al. 2015). The precise structure of the clumps, i.e., their density and ionization structure, is unknown. Most recent 2D simulations of such a stellar wind show a very complex evolution of velocity and density structures with the formation of characteristic small-scale clumps of various shapes embedded in areas with lower density (Sundqvist et al. 2018). Sundqvist et al. (2018) have found a typical clump mass of 1017 g and an average clump size of 1% of the stellar radius at a distance of two stellar radii. These results qualitatively confirm earlier theoretical models (e.g., Oskinova et al. 2012; Sundqvist & Owocki 2013) and observations (e.g., Grinberg et al. 2015). See also the review paper by Mart\u00ednez-N\u00fa\u00f1ez et al. (2017).","Citation Text":["Grinberg et al. 2017"],"Functions Text":["When the line of sight to the compact object passes through one of these clumps, X-rays are absorbed by the moderately ionized material in the clump, leading to a so-called dipping event. This is also observed for other sources"],"Functions Label":["Similarities"],"Citation Start End":[[888,908]],"Functions Start End":[[637,864]]}
{"Identifier":"2021ApJ...906...21J__Fabian_et_al._2009_Instance_1","Paragraph":"Ricci et al. (2017c, hereafter R17c) recently reported on a study of the relationship between obscuration and accretion rate in a large, relatively unbiased, and complete sample of local AGNs. Specifically, they investigated 836 AGNs with a median redshift of \u2329z\u232a = 0.037 selected by the hard X-ray (14\u2013195 keV) Swift Burst Alert Telescope (BAT; Gehrels et al. 2004; Barthelmy et al. 2005; Krimm et al. 2013) all-sky survey (Baumgartner et al. 2013; Koss et al. 2017; Oh et al. 2018), which is sensitive to sources with column densities up to NH \u2248 1024 cm\u22122.  Approximately one-half of the sources had robust measurements of column densities, intrinsic X-ray luminosities, and black hole masses, from which R17c was able to show that while unobscured AGNs are seen with Eddington fractions up to the Eddington limit, very few local, obscured AGNs are found with Eddington fractions above approximately 10%. This strengthened earlier results based on smaller samples (e.g., Fabian et al. 2009) and was interpreted as evidence for radiation-pressure-driven AGN feedback (e.g., King 2003; Murray et al. 2005) clearing the immediate BH environment of dusty gas (e.g., Fabian et al. 2006, 2008). For dusty gas (neutral or partially ionized), the effective cross section between matter and radiation (\u03c3dust) becomes larger than that between electrons and radiation for ionized gas (\u03c3T, for Thompson scattering), due to absorption of radiation by dust. This is given by the Eddington ratio for ionized gas,\n1\n\n\n\n\n\nwhere Lbol is the bolometric luminosity, LEdd is the Eddington luminosity with MBH the BH mass, and mp is the proton mass. For dusty gas, we use \u03c3dust instead of \u03c3T in Equation (1), where fEdd is redefined as the effective Eddington ratio (fEdd,dust = fEdd\u03c3dust\/\u03c3T). AGNs are strong ionizing sources, but they are fully ionized close to their accretion disks (e.g., Osterbrock 1979; Ballantyne et al. 2001), though the greater than parsec-scale environment starts to be composed of dusty gas (e.g., Kishimoto et al. 2011; Minezaki et al. 2019).","Citation Text":["Fabian et al. 2009"],"Functions Text":["This strengthened earlier results based on smaller samples (e.g.,"],"Functions Label":["Similarities"],"Citation Start End":[[973,991]],"Functions Start End":[[907,972]]}
{"Identifier":"2020ApJ...901...41S__Duval_et_al._2014_Instance_1","Paragraph":"Observations have shown that the shape of the Ly\u03b1 line is diverse. It includes broad damped absorption profiles, P-Cygni profiles, double-peak profiles, pure symmetric emission profiles, and combinations thereof (Kunth et al. 1998; Mas-Hesse et al. 2003; Shapley et al. 2003; M\u00f8ller et al. 2004; Noll et al. 2004; Tapken et al. 2004; Venemans et al. 2005; Wilman et al. 2005). This variety can be understood through a detailed radiative transfer calculation, which is analytically solvable only for simple cases (e.g., a static, plane-parallel slab by Harrington 1973 and Neufeld 1990, and a static uniform sphere by Dijkstra et al. 2006). Later, numerical algorithms based on Monte Carlo techniques were developed to solve radiative transfer for more general cases. Now theoretical studies mostly rely on them (e.g., Spaans 1996; Loeb & Rybicki 1999; Ahn et al. 2000, 2002; Zheng & Miralda-Escud\u00e9 2002; Richling 2003; Cantalupo et al. 2005; Dijkstra et al. 2006; Hansen & Oh 2006; Tasitsiomi 2006; Verhamme et al. 2006, 2015; Laursen et al. 2013; Behrens et al. 2014; Duval et al. 2014; Gronke et al. 2015; Smith et al. 2019; Lao & Smith 2020; Michel-Dansac et al. 2020). Meanwhile, a galaxy model needs to be constructed to perform such a radiative transfer calculation. One can adopt a realistic galaxy model from hydrodynamical simulations. Galaxies from such simulations can be useful for performing a statistical study of Ly\u03b1 properties, but they cannot be directly used to model individual galaxies in observations. Therefore it would be better to adopt a simple but manageable toy model for the purpose of reproducing observations. One example for such models is a shell model, in which a central Ly\u03b1 source is surrounded by a constantly expanding, homogeneous, spherical shell of H i medium with dust. Although this shell model has surprisingly well reproduced many observed Ly\u03b1 line profiles (e.g., Ahn 2004; Schaerer & Verhamme 2008; Verhamme et al. 2008; Schaerer et al. 2011; Gronke et al. 2015; Yang et al. 2016; Gronke 2017; Karman et al. 2017), because of its extreme simplicity and contrivance, there is still room for improvement (e.g., see Section 7.2 in Yang et al. 2016; Orlitov\u00e1 et al. 2018).","Citation Text":["Duval et al. 2014"],"Functions Text":["Later, numerical algorithms based on Monte Carlo techniques were developed to solve radiative transfer for more general cases. Now theoretical studies mostly rely on them (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[1069,1086]],"Functions Start End":[[640,817]]}
{"Identifier":"2022ApJ...929...11M__Prasad_et_al._2017_Instance_1","Paragraph":"In this context, it must also be noted here that these conclusions are specific to the geometry of the GCS model, which is an idealized geometrical figure that has its limitations and constraints (see Thernisien et al. 2009). Regarding the evolution of the legs, the identification of the two separate legs of the CMEs requires observation at the absolute lower heights. Thus the legs can be identified in the K-Cor FOV, while they are not seen in the COR-1 FOV at the same time, as shown in Figure 1; but it should also be noted that despite the promising FOV of K-Cor, the poor image quality due to the challenges faced from it being a ground-based coronagraph makes it difficult to fit (refer to the discussion in Section 2.2). In this regard, the upcoming ADITYA-L1 mission (Seetha & Megala 2017), with the Visible Emission Line Coronagraph (VELC; FOV: 1.05\u20133 R\n\u2299; Banerjee et al. 2017; Prasad et al. 2017) on board, and PROBA-3 (FOV : 1.02\u20133 R\n\u2299; Renotte et al. 2014), with the giant Association de Satellites pour l\u2019Imagerie et l\u2019Interferom\u00e9trie de la Couronne Solaire (ASPIICS; Lamy et al. 2017), will provide much better data and hence will help in arriving at much stronger conclusions on the evolution of CME legs. Having said that, it must also be noted that a true estimation of the volume of CME legs will require the CME to be seen FO, as an FO view will help in identifying the inner edges of the CME and hence the volume of its legs. The studied CMEs in this work are all seen FO in the K-Cor FOV (please see Figure 1). Thus, in future, for a larger statistical study, the appearance of the CME (whether FO or EO) should also be considered in the estimation of the volume of the CME legs. Apart from that, around one-third of CMEs have been reported as having a flux-rope morphology (see Vourlidas et al. 2013), which happens to be the bedrock of the GCS model, thus a study of the three separate sections of the flux-rope model of the CME will help us to have a much better understanding of the validity of self-similar expansion, and thus provide more precise constraints to models that study flux-rope initiation and evolution.","Citation Text":["Prasad et al. 2017"],"Functions Text":["In this regard, the upcoming ADITYA-L1 mission","with the Visible Emission Line Coronagraph","will provide much better data and hence will help in arriving at much stronger conclusions on the evolution of CME legs."],"Functions Label":["Future Work","Future Work","Future Work"],"Citation Start End":[[891,909]],"Functions Start End":[[731,777],[802,844],[1104,1224]]}
{"Identifier":"2015AandA...584A..75V__Essen_et_al._(2014)_Instance_6","Paragraph":"The data presented here comprise quasi-simultaneous observations during secondary eclipse of WASP-33 b around the V and Y bands. The predicted planet-star flux ratio in the V-band is 0.2 ppt, four times lower than the accuracy of our measurements. Therefore, we can neglect the planet imprint and use this band to measure the stellar pulsations, and most specifically to tune their current phases (see phase shifts in von Essen et al. 2014). Particularly, our model for the light contribution of the stellar pulsations consists of eight sinusoidal pulsation frequencies with corresponding amplitudes and phases. Hence, to reduce the number of 24 free parameters and the values they can take, we use prior knowledge about the pulsation spectrum of the star that was acquired during von Essen et al. (2014). As the frequency resolution is 1\/\u0394T (Kurtz 1983), 3.5 h of data are not sufficient to determine the pulsations frequencies. Therefore, during our fitting procedure we use the frequencies determined in von Essen et al. (2014) as starting values plus their derived errors as Gaussian priors. As pointed out in von Essen et al. (2014), we found clear evidences of pulsation phase variability with a maximum change of 2 \u00d7 10-3 c\/d. In other words, as an example after one year time a phase-constant model would appear to have the correct shape with respect to the pulsation pattern of the star, but shifted several minutes in time. To account for this, the eight phases were considered as fitting parameters. The von Essen et al. (2014) photometric follow-up started in August, 2010, and ended in October, 2012, coinciding with these LBT data. We then used the phases determined in von Essen et al. (2014) during our last observing season as starting values, and we restricted them to the limiting cases derived in Sect. 3.5 of von Essen et al. (2014), rather than allowing them to take arbitrary values. The pulsation amplitudes in \u03b4 Scuti stars are expected to be wavelength-dependent (see e.g. Daszy\u0144ska-Daszkiewicz 2008). Our follow-up campaign and these data were acquired in the blue wavelength range. Therefore the amplitudes estimated in von Essen et al. (2014), listed in Table 1, are used in this work as fixed values. This approach would be incorrect if the pulsation amplitudes would be significantly variable (see e.g., Breger et al. 2005). Nonetheless, the short time span of LBT data, and the achieved photometric precision compared to the intrinsically low values of WASP-33\u2019s amplitudes, make the detection of any amplitude variability impossible. ","Citation Text":["von Essen et al. (2014)"],"Functions Text":["We then used the phases determined in","during our last observing season as starting values,"],"Functions Label":["Uses","Uses"],"Citation Start End":[[1684,1707]],"Functions Start End":[[1646,1683],[1708,1760]]}
{"Identifier":"2016MNRAS.461..248S__Munari_et_al._2013_Instance_1","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"],"Functions Text":["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"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[953,971]],"Functions Start End":[[604,951]]}
{"Identifier":"2018ApJ...855...23I__Yang_et_al._2014_Instance_1","Paragraph":"Cosmic rays (CRs) represent a crucial ingredient in the dynamical and chemical evolution of interstellar clouds. Interaction of CRs with molecular clouds is accompanied by various processes generating observable radiation signatures, such as ionization of molecular hydrogen (see, e.g., Oka et al. 2005; Dalgarno 2006; Indriolo & McCall 2012) and iron (e.g., Dogiel et al. 1998, 2011; Tatischeff et al. 2012; Yusef-Zadeh et al. 2013; Nobukawa et al. 2015; Krivonos et al. 2017), as well as production of neutral pions whose decay generates gamma-rays in the GeV (e.g., Yang et al. 2014, 2015; Tibaldo et al. 2015) and TeV (e.g., Aharonian et al. 2006; Abramowski et al. 2016; Abdalla et al. 2017) energy ranges. Being a unique source of ionization in dark clouds, where the interstellar radiation cannot penetrate, CRs provide a partial coupling of the gas to magnetic field lines, which could slow down or prevent further contraction of the cloud (e.g., Shu et al. 1987). CRs are fundamental to the beginning of astrochemistry because they promote the formation of \n\n\n\n\n\n ions, which can easily donate a proton to elements such as C and O, and thus eventually form molecules containing elements heavier than H (e.g., Yamamoto 2017). Through the ionization of H2 molecules and the consequent production of secondary electrons, CRs are an important heating source of dark regions (e.g., Goldsmith 2001). Their interaction with H2 can also result in molecular excitation, followed by fluorescence producing a tenuous UV field within dark clouds and dense cores (Cecchi-Pestellini & Aiello 1992; Shen et al. 2004; Ivlev et al. 2015a); this UV field can photodesorb molecules from the icy dust mantles and help to maintain a non-negligible amount of heavy molecules (such as water) in the gas phase (e.g., Caselli et al. 2012). Furthermore, CRs can directly impinge on dust grains and heat up the icy mantles, causing catastrophic explosions of these mantles (L\u00e9ger et al. 1985; Ivlev et al. 2015b) and activating the chemistry in solids (Shingledecker et al. 2017). Finally, CRs play a fundamental role in the charging of dust grains and the consequent coagulation of dust (Okuzumi 2009; Ivlev et al. 2015a, 2016), which is particularly important for the formation of circumstellar disks (e.g., Zhao et al. 2016) and of planets in more evolved protoplanetary disks (e.g., Testi et al. 2014).","Citation Text":["Yang et al. 2014"],"Functions Text":["Interaction of CRs with molecular clouds is accompanied by various processes generating observable radiation signatures, such as","as well as production of neutral pions whose decay generates gamma-rays in the GeV (e.g.,"],"Functions Label":["Background","Background"],"Citation Start End":[[569,585]],"Functions Start End":[[113,241],[479,568]]}
{"Identifier":"2015AandA...578L...8B__Bern\u00e9_et_al._2009_Instance_1","Paragraph":"Gomez\u2019s Hamburger (IRAS 18059-3211; hereafter GoHam) is an A star surrounded by a dusty disk. When first studied by Ruiz et al. (1987), it was classified as an evolved object (post-AGB star) on the basis of its spectral type and the presence of dust. However, all recent studies (Bujarrabal et al. 2008, 2009; Wood et al. 2008; De Beck et al. 2010) clearly indicate that it is a young A star surrounded by a protoplanetary disk. The distance to GoHam is not known precisely, but a value d = 250 \u00b1 50 pc is required to satisfy all the existing observational constraints (Wood et al. 2008; Bern\u00e9 et al. 2009; Bujarrabal et al. 2009). We here adopt this value with the uncertainty. GoHam presents intense CO emission; SMA maps of 12CO and 13CO J = 2 \u2212 1 lines very clearly show the Keplerian dynamics of the disk (Bujarrabal et al. 2008, 2009). The lower limit for the disk mass derived from these CO observations is of about 10-2M\u2299, while the mass upper limit is estimated to be ~0.3 M\u2299  based on dust emission (Bujarrabal et al. 2008; Wood et al. 2008) and assuming an interstellar dust-to-gas mass ratio of 0.01. Overall, GoHam appears to be similar to isolated Herbig stars (Meeus et al. 2001) such as HD 100546 and HD 169142, in a more massive version, but still smaller than the recently discovered disk around CAHA J23056+6016 (Quanz et al. 2010). GoHam is seen almost perfectly edge on, which offers the possibility to study this class of objects from a new and complementary perspective, in particular, with improved constraints on the vertical structure of the disk. Using a radiative transfer model to predict line emission from a Keplerian flaring disk, Bujarrabal et al. (2009) derived a large-scale description of the physical conditions throughout the disk. After subtracting the model that best fit the observations, these authors found a significant residual emission situated about 1.3\u2032\u2032 (330 \u00b1 70 AU) south of the central star, which they identified as a gas condensation, containing a mass between one and few times that of Jupiter. Hence, this source was proposed to be a candidate protoplanet, possibly resulting from a GI collpapse. ","Citation Text":["Bern\u00e9 et al. 2009"],"Functions Text":["The distance to GoHam is not known precisely, but a value d = 250 \u00b1 50 pc is required to satisfy all the existing observational constraints","We here adopt this value with the uncertainty."],"Functions Label":["Uses","Uses"],"Citation Start End":[[588,605]],"Functions Start End":[[429,568],[632,678]]}
{"Identifier":"2022AandA...666A..28S__Rutherford_1903_Instance_1","Paragraph":"The velocity distribution of the plasma motion is shown in Fig. 12, where the magnitude is scaled with respect to the Alfv\u00e9n velocity, vA, which is measured based on the magnetic field strength, B0\u2004=\u20042 G, and the mass density of the equilibrium current sheet, \u03c1c\u2004=\u20042.81\u2005\u00d7\u200510\u221215 g cm\u22123. The Alfv\u00e9n timescale is measured by \n\n\n\n\nt\nA\n\n=\n\n\nL\n\n\n\u00af\n\n\n\/\n\nv\nA\n\n\n\n$ t_A = \\bar{L}\/\\mathit{v}_A $\n\n\n, where, \n\n\n\n\nL\n\n\n\u00af\n\n\n\n$ \\bar{L} $\n\n\n is the unit length of 109 cm. It is evident from Figs. 12a, b and c that the velocity, vx remains localized in the vicinity of the current sheet (y\u2004=\u20040). We note that in line with our initial single-island magnetic field perturbation, we see a pronounced rightward motion in the right half of the domain, and a leftward motion in the left half. We later see typical Petschek-like signatures in the flow fields between islands, especially about the middle x\u2004=\u20040, with super-Alfv\u00e9nic outflow speeds bounded by slow shocks. Figure 13a represents the evolutionary nature of a current sheet in adiabatic and non-adiabatic conditions. It is clear from the figure that the instantaneous maximum velocity growth for the non-adiabatic case is more rapid than for the adiabatic conditions. The evolutionary behavior of the current sheet configuration due to thermal and tearing instabilities is shown by the black curve in Fig. 13a for plasma-\u03b2\u2004=\u20040.2, and a given resistivity value, \u03b7\u2004=\u20040.001, while the evolution for different \u03b7 values are shown in Fig. 13b. As a diagnostic measurement of the instability, we determine the evolution of the instantaneous maximum absolute velocity, |vx|max. Figure 13a shows that this evolution exhibits three distinct phases: (i) the early phase (between t\u2004=\u20040 and 250 s), where the velocity growth occurs exponentially (linearly on the logarithmic-linear scale), which is called the linear growth regime; (ii) the middle phase between t\u2004=\u2004250 and 665 s, where the growth rate is slower compared to the linear phase, which is called the Rutherford regime (Rutherford 1903), and (iii) the final phase, which starts at t\u2004=\u2004665 s, where the instability suddenly develops in an explosive way, and finally saturates at a later time, which we call the post-Rutherford regime. To infer the evolution rates quantitatively for all the different phases, we calculate the growth rates by scaling them with respect to the Alfv\u00e9n timescale, tA. We define the growth rate as \u03b3\u2004=\u2004d(ln(|vx|max))\/dt. To estimate the linear growth rate, \u03b3lin, we calculate the growth rate in the linear regime by taking the mean value of the slope, which gives \n\n\n\n\u03b3\n\nl\ni\nn\n\n\n=\n3.76\n\u00d7\n\n10\n\n\u2212\n1\n\n\n\nt\nA\n\n\u2212\n1\n\n\n\n$ \\gamma_{\\rm lin}=3.76 \\times 10^{-1} t_A^{-1} $\n\n\n. This value is larger by an order of magnitude compared to the studies of the double current sheet problem (Otto & Birk 1992; Zhang & Ma 2011; Akramov & Baty 2017; Paul & Vaidya 2021), where the radiative cooling effect (or other non-adiabatic effects, e.g., thermal conduction) is not incorporated. This implies that the higher linear growth rate can be ascribed to the non-adiabatic effects of the radiative cooling and background heating. This is also in agreement with our own study for a single current layer model reflected in Fig. 13a, where the average growth rate for the adiabatic medium is lower than the non-adiabatic case. Similarly, we estimate the average growth rates for the Rutherford regime (\u03b3Ruth) and the post-Rutherford regime (\u03b3PR) for different resistivity values within the range of \u03b7\u2004=\u20040.0001 to 0.005. The velocity evolution for some selected resistivity values are shown in Fig. 13b. This shows that the explosive phase of the evolution starts at later times for higher resistivity values, and converges at the final stage. For a Sweet\u2013Parker type current sheet (where the inverse aspect ratio of the current sheet follows the scaling relation \n\n\n\n\nl\ns\n\n\/\nL\n\u223c\n\nS\n\nL\n\n\n\u2212\n1\n\/\n2\n\n\n\n\n$ l_s\/L \\sim S_{\\mathrm{L}}^{-1\/2} $\n\n\n), the thickness of the current sheet increases with the resistivity (Loureiro et al. 2007), which reduces the growth rate of the tearing mode when it is normalized with respect to the Alfv\u00e9n crossing time along the length of the current sheet (in the x-direction in our case). Hence, the explosive phase of the evolution in our simulation starts at later times for higher resistivity values. We estimated the absolute current density, |Jz| (normalized to unity) before the fragmentation stage of the current sheet (t\u2004=\u2004214 s) by taking a vertical cut along the y-direction at x\u2004=\u20040 for two different resistivities, \u03b7\u2004=\u20040.0001 and 0.001, to confirm that the thickness of the current sheet increases with resistivity (see Fig. 14). The resistivity dependence for the different evolution phases is shown in Fig. 15. Figure 15a shows that \u03b3Ruth follows a power-law dependence with the resistivity, \u03b3Ruth\u2004\u2248\u2004\u03b7\u22120.1, with a correlation coefficient (CC) of \u221264.1%. The resistivity scaling relation for the post-Rutherford and the entire nonlinear regimes are shown in Figs. 15b and c respectively. We estimate the growth rate scaling relations for the post-Rutherford regime, \u03b3PR\u2004\u2248\u2004\u03b70.03 (with CC = 59.9%), and the entire non-linear regime, \u03b3avg\u2004\u2248\u2004\u03b70.017 (with CC = 66.7%). Previous studies by Zhang & Ma (2011), Akramov & Baty (2017), and Guo et al. (2017 and references therein) have reported the resistivity scaling relation of the non-linear growth rates for the DTM setup in the adiabatic environment, which have larger power-law indices compared to our estimation. Hence, our study infers that the resistivity dependence on the nonlinear growth rates is weaker when the thermal instability reinforces the tearing mode.","Citation Text":["Rutherford 1903"],"Functions Text":["Figure 13a shows that this evolution exhibits three distinct phases: (i) the early phase (between t\u2004=\u20040 and 250 s), where the velocity growth occurs exponentially (linearly on the logarithmic-linear scale), which is called the linear growth regime; (ii) the middle phase between t\u2004=\u2004250 and 665 s, where the growth rate is slower compared to the linear phase, which is called the Rutherford regime","and (iii) the final phase, which starts at t\u2004=\u2004665 s, where the instability suddenly develops in an explosive way, and finally saturates at a later time, which we call the post-Rutherford regime."],"Functions Label":["Uses","Uses"],"Citation Start End":[[2006,2021]],"Functions Start End":[[1607,2004],[2024,2219]]}
{"Identifier":"2020MNRAS.495..758H__Wang_et_al._2014_Instance_1","Paragraph":"In all our simulations independent of the inflow Mach number, the radial velocity dispersion at the filament boundary amounts to about 0.85 times the total equilibrium velocity dispersion of the non-self-gravitational case for which a functional form can be found in the appendix. Thus, we can calculate the theoretical radius and central density of the filament at every line-mass and therefore we can make predictions on the fragmentation length and time-scales of cores forming in an accreting filament using the gravitational fragmentation model. This model was successfully applied to explain several observed core distances (Jackson et al. 2010; Miettinen 2012; Busquet et al. 2013; Beuther et al. 2015; Contreras et al. 2016; Heigl et al. 2016; Kainulainen et al. 2016) however it is not able to explain all observations (Andr\u00e9 et al. 2010; Kainulainen et al. 2013; Takahashi et al. 2013; Lu et al. 2014; Wang et al. 2014; Henshaw et al. 2016; Teixeira et al. 2016; Kainulainen et al. 2017; Lu et al. 2018; Palau et al. 2018; Williams et al. 2018; Zhou et al. 2019). It predicts that small density perturbations in the linear regime along the filament axis of the form:\n(32)$$\\begin{eqnarray*}\r\n\\rho (r, x, t) = \\rho _0(r) \\left(1 + \\epsilon \\exp (ikx -i\\omega t)\\right)\r\n\\end{eqnarray*}$$will grow for values of k where the dispersion relation \u03c92(k) is negative. Here \u03c10 is the unperturbed initial density, k = 2\u03c0\/\u03bb is the wave vector with \u03bb being the perturbation length, x is the filament axis, \u03c9 = 1\/\u03c4 is the growth rate with \u03c4 being the growth time-scale, t the time variable, and \u03f5 the perturbation strength. The fastest growing, or dominant, fragmentation length scale \u03bbdom as well as the growth time-scale of the dominant mode \u03c4dom depend on the current line-mass as well as the current central density of the filament and are given by the pre-calculated (Nagasawa 1987) and interpolated values in Fischera & Martin (2012), shown by their Table E.1. We use these values to determine the length scale of the fastest growing mode at every line-mass for the same mass accretion rate but for different inflow Mach numbers as shown in Fig. 10. As one can see, the dominant fragmentation length changes over the evolution of the line-mass. At the boundary values it vanishes to zero and it has a maximum at about fcyl = 0.4. The figure is self-similar for different mass accretion rates, with a lower rate leading to a larger dominant fragmentation length. For a constant accretion rate, the fragmentation length does not vary much for different inflow Mach numbers. Only for large and for very low inflow Mach numbers, the fragmentation length is slightly larger. As the dominant fragmentation length constantly changes as fcyl grows, it is hard to make predictions of what will be the final distance between forming cores. But the curves have a maximum which allows us to make a prediction about the minimum number of cores that will form. For instance, a filament with an inflow Mach number of 4.0 and a length of 0.2\u2009pc will form at least one core. As soon as the first core forms, the further evolution of the filament is also influenced by the gravitational attraction of the core. This makes the formation of additional cores even more unpredictable.","Citation Text":["Wang et al. 2014"],"Functions Text":["This model was successfully applied to explain several observed core distances","however it is not able to explain all observations"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[912,928]],"Functions Start End":[[551,629],[777,827]]}
{"Identifier":"2018MNRAS.473.3810Y__Mitrushchenkov_et_al._2017_Instance_2","Paragraph":"The lack of data on inelastic processes due to collisions with neutral hydrogen atoms has been a major limitation on modelling of F-, G- and K-star spectra in statistical equilibrium, and thus to reliably proceeding beyond the assumption of local thermodynamic equilibrium (LTE) in analysis of stellar spectra and the determination of elemental abundances. This problem has been well documented, e.g. see Lambert (1993); Barklem (2016a) and references therein. Significant progress has been made in recent times through detailed full-quantum scattering calculations, based on quantum chemical data, for the cases of simple atoms such as Li, Na, Mg and Ca (Belyaev & Barklem 2003; Barklem, Belyaev & Asplund 2003; Belyaev et al. 2010; Barklem et al. 2010; Belyaev et al. 2012; Barklem et al. 2012; Mitrushchenkov et al. 2017). These calculations have demonstrated the importance of the ionic-covalent curve crossing mechanism leading naturally to charge transfer processes (mutual neutralization and ion-pair production), in addition to excitation and de-excitation processes. The importance of this mechanism has allowed various simplified model approaches to be developed, which may be used in cases where suitable quantum chemistry data are not been available. In particular a semi-empirical model has been employed for Al, Si, Be and Ca (Belyaev 2013a,b; Belyaev, Yakovleva & Barklem 2014b; Yakovleva, Voronov & Belyaev 2016; Belyaev et al. 2016), and a theoretical model based on a two-electron asymptotic linear combinations of atomic orbitals (LCAO) approach, has also been employed for Ca (Barklem 2016b, 2017). Comparisons of the two methods show quite good agreement and reasonable agreement with the full quantum calculations is found, particularly for the most important processes with the largest rates (Barklem 2016b, 2017; Mashonkina, Sitnova & Belyaev 2017; Mitrushchenkov et al. 2017). Thus, the model approaches provide a useful route for obtaining estimates of the rates for these processes for many elements of astrophysical interest.","Citation Text":["Mitrushchenkov et al. 2017"],"Functions Text":["Comparisons of the two methods show quite good agreement and reasonable agreement with the full quantum calculations is found, particularly for the most important processes with the largest rates"],"Functions Label":["Similarities"],"Citation Start End":[[1873,1899]],"Functions Start End":[[1619,1814]]}
{"Identifier":"2021MNRAS.507.1138M__Parfenov_&_Sobolev_2014_Instance_1","Paragraph":"Several mechanisms have been proposed to explain the periodicity of methanol masers. van der Walt, Goedhart & Gaylard (2009) and van der Walt (2011) proposed a colliding wind binary (CWB) system for the periodic methanol masers in G9.62+0.20E and other periodic methanol maser sources with similar light curves. Using a more realistic model van den Heever et al. (2019) also showed that the CWB scenario can explain the light curves of the periodic methanol masers in G9.62+0.20E, G22.357+0.066, G37.55+0.20, and G45.473+0.134. Other mechanisms proposed are the pulsational instabilities of very young accreting high-mass stars (Inayoshi et al. 2013), spiral shocks associated with very young binary systems orbiting within a circumbinary disc (Parfenov & Sobolev 2014), periodic accretion in a very young binary system (Araya et al. 2010), and outflows in a binary system (Singh & Deshpande 2012). Maswanganye et al. (2015) evoked an eclipsing binary system involving a high-mass (primary) star and a bloated low-mass (companion) star in explaining the observed periodicity of the methanol masers in G358.460-0.391. While each of the suggested mechanisms may be sufficient for explaining periodicity in individual star-forming regions, the complexity in the nature of massive star-forming environments and the uniqueness of each region, make it difficult to evoke any one of the mechanisms for all cases. In fact, different mechanisms may be responsible for the different observed flare profiles (see van der Walt et al. 2016, for a discussion). However, the fact that it is possible to identify at least two groups of periodic sources (as explained above), each for which the light curves are very similar, suggests that there are some sources for which the underlying mechanism is the same. It is therefore reasonable to argue that there might be other similarities in star-forming regions that host periodic methanol masers with same light curves. In other words, the similarity in light curves might also manifest in other properties of the star-forming regions in addition to the maser emission.","Citation Text":["Parfenov & Sobolev 2014"],"Functions Text":["Other mechanisms proposed are","spiral shocks associated with very young binary systems orbiting within a circumbinary disc"],"Functions Label":["Background","Background"],"Citation Start End":[[745,768]],"Functions Start End":[[528,557],[652,743]]}
{"Identifier":"2019AandA...628A.118B__Feruglio_et_al._2017_Instance_2","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"],"Functions Text":["Recent studies increasing the statistics of sources with detection of molecular outflows have widened the range of measured energetics (e.g."],"Functions Label":["Background"],"Citation Start End":[[1731,1751]],"Functions Start End":[[1540,1680]]}
{"Identifier":"2020ApJ...903L..22T__Vuitton_et_al._2007_Instance_1","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)"],"Functions Text":["While the Loison et al. (2015) CH3C3N model corroborates the upper atmospheric abundance of C4H3N inferred by","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."],"Functions Label":["Similarities","Differences"],"Citation Start End":[[110,131]],"Functions Start End":[[0,109],[237,508]]}
{"Identifier":"2020ApJ...892..110C__Saladino_et_al._2018_Instance_1","Paragraph":"Asymptotic-giant-branch (AGB) stars have a significantly larger size (\u223c1 au) than their main-sequence (MS) counterparts. They have pulsating atmospheres (Vlemmings et al. 2017; Khouri et al. 2019) and may exhibit variability with long periods ranging from 200 to 1000 days (Mowlavi et al. 2018; Karambelkar et al. 2019). AGB stars are one of the major sites in galaxies that produce metals. Metals can be carried away from the AGB stars by radiation-driven AGB winds when dust forms. The speed of the AGB wind varies from 4 to 20 km s\u22121 (H\u00f6fner & Olofsson 2018), and a companion star may capture the wind with its gravity. In the case that there is an MS star close to an AGB star, a substantial fraction of the mass loss may be accreted onto the MS companion (Chen et al. 2017; Saladino et al. 2018, 2019). As a result, the metallicity of the companion may change. Such early stage low-mass stars become chemically peculiar, and their future evolution will be strongly affected. Carbon-enhanced-metal-poor (CEMP) stars (Beers & Christlieb 2005; Abate et al. 2013, 2015), Barium stars (Bidelman & Keenan 1951; Escorza et al. 2019), CH stars (Keenan 1942; McClure & Woodsworth 1990), and dwarf carbon stars (Dahn et al. 1977; Roulston et al. 2019) are common examples of the chemically peculiar stars. Their existence could be evidence of mass transfer during the previous AGB binary phase. The binarity of CH stars and CEMP stars has been studied (McClure & Woodsworth 1990; Starkenburg et al. 2014; Jorissen et al. 2016), confirming that many of them have companions. A number of recent studies show that the eccentricity of some of the aforementioned chemically peculiar stars may be large (Hansen et al. 2016; Jorissen et al. 2016, 2019; Van der Swaelmen 2017; Oomen et al. 2018), and their orbital periods range from hundreds to thousands of days. The nonzero eccentricity in these close binary stars indicates that some intense interactions that can pump the eccentricity may happen during their AGB binary phases. A strong correlation between a circumstellar disk and binarity has also been established in Galactic RV Tauri stars (Manick et al. 2017). Furthermore, many RV Tauri stars show a lack of refractory elements, which is called \u201cdepletion\u201d (Giridhar et al. 1994; Van Winckel et al. 1998). Some researches suggest that the reaccretion of gas from a circumstellar disk around the post-AGB star (Gezer et al. 2019; Oomen et al. 2019) may induce the \u201cdepletion.\u201d Besides the \u201csmoking gun\u201d evidence, observations also reveal that dusty circumbinary disks exist in binary systems with evolved stars (Kervella et al. 2015; Hillen et al. 2016; Homan et al. 2017; Ertel et al. 2019). The UV excess of some AGB stars also suggests that there could be accreting MS companions near them (Sahai et al. 2008; Ortiz & Guerrero 2016).","Citation Text":["Saladino et al. 2018"],"Functions Text":["In the case that there is an MS star close to an AGB star, a substantial fraction of the mass loss may be accreted onto the MS companion","As a result, the metallicity of the companion may change. Such early stage low-mass stars become chemically peculiar, and their future evolution will be strongly affected."],"Functions Label":["Background","Background"],"Citation Start End":[[779,799]],"Functions Start End":[[623,759],[808,979]]}
{"Identifier":"2021AandA...647A.140C__Gianninas_et_al._2016_Instance_1","Paragraph":"In recent years, numerous low-mass and ELM WDs have been detected in the context of relevant surveys, such as the SDSS, ELM, SPY and WASP (see, e.g., Koester et al. 2009; Brown et al. 2010, 2016, 2020; Kilic et al. 2011, 2012; Gianninas et al. 2015; Kosakowski et al. 2020). The discovery of their probable precursors, namely, the so-called low-mass pre-WDs, has triggered an interest in these types of objects because of the possibility of studying the evolution of the progenitors that lead to the WD phase. Even more interestingly, the detection of multi-periodic brightness variations in low-mass WDs (Hermes et al. 2012, 2013a,b; Kilic et al. 2015, 2018; Bell et al. 2017, 2018; Pelisoli et al. 2018), and low-mass pre-WDs (Maxted et al. 2013, 2014; Gianninas et al. 2016; Wang et al. 2020) has brought about new classes of pulsating stars known as ELMVs and pre-ELMVs, respectively (ELM and pre-ELM variables, respectively). It has allowed for the study of their stellar interiors using the tools of asteroseismology, similarly to the case of other pulsating WDs such as ZZ Ceti stars or DAVs \u2013pulsating WDs with H-rich atmospheres \u2013 and V777 Her or DBVs \u2013 pulsating WDs with He-rich atmospheres (Winget & Kepler 2008; Fontaine & Brassard 2008; Althaus et al. 2010; C\u00f3rsico et al. 2019). The pulsations observed in ELMVs are compatible with global gravity (g)-mode pulsations. In the case of pulsating ELM WDs, the pulsations have large amplitudes mainly at the core regions (Steinfadt et al. 2010; C\u00f3rsico et al. 2012; C\u00f3rsico & Althaus 2014a), allowing for the study of their core chemical structure. According to nonadiabatic computations (C\u00f3rsico et al. 2012; Van Grootel et al. 2013; C\u00f3rsico & Althaus 2016), these modes are probably excited by the \u03ba\u2005\u2212\u2005\u03b3 (Unno et al. 1989) mechanism acting at the H-ionization zone. In the case of pre-ELMVs, the nonadiabatic stability computations for radial (Jeffery & Saio 2013) and nonradial p- and g-mode pulsations (C\u00f3rsico et al. 2016; Gianninas et al. 2016; Istrate et al. 2016b) revealed that the excitation is probably due to the \u03ba\u2005\u2212\u2005\u03b3 mechanism, acting mainly in the zone of the second partial ionization of He, with a weaker contribution from the region of the first partial ionization of He and the partial ionization of H. The presence of He in the driving zone is crucial to having the modes destabilized by the \u03ba\u2005\u2212\u2005\u03b3 mechanism (C\u00f3rsico & Althaus 2016; Istrate et al. 2016b).","Citation Text":["Gianninas et al. 2016"],"Functions Text":["Even more interestingly, the detection of multi-periodic brightness variations in low-mass WDs","has brought about new classes of pulsating stars known as ELMVs and pre-ELMVs, respectively (ELM and pre-ELM variables, respectively)."],"Functions Label":["Background","Background"],"Citation Start End":[[755,776]],"Functions Start End":[[510,604],[796,930]]}
{"Identifier":"2020ApJ...892L...3A___2019e_Instance_1","Paragraph":"Because of its large mass, the discovery of GW190425 suggests that gravitational-wave analyses can access densities several times above nuclear saturation (see, e.g., Figure 4 in Douchin & Haensel 2001) and probe possible phase transitions inside the core of a neutron star (NS) (Oertel et al. 2017; Essick et al. 2019; Tews et al. 2019). However, binaries comprised of more massive stars are described, for a fixed EoS, by smaller values of the leading-order tidal contribution to the gravitational-wave phasing \n\n\n\n\n\n (Flanagan & Hinderer 2008). These are intrinsically more difficult to measure. For GW190425, this is exacerbated by the fairly low S\/N of the event compared to GW170817. Overall, we find that constraints on tides, radius, possible p\u2013g instabilities (Venumadhav et al. 2013; Weinberg et al. 2013; Weinberg 2016; Zhou & Zhang 2017), and the EoS from GW190425 are consistent with those obtained from GW170817 (Abbott et al. 2017b, 2019e). However, GW190425 is less constraining of NS properties, limiting the radius to only below 15 km, \n\n\n\n\n\n to below 1100 and only ruling out phenomenological p\u2013g amplitudes above 1.3 times the 90% upper limit obtained from GW170817 at the same confidence level. The p\u2013g constraints were obtained with a different high-spin prior than the rest of the results (see Appendix F.5) but the difference does not significantly change our conclusions. Spin priors can affect the inference of tidal and EoS parameters, and we note that the low-spin results are generally more constrained. Following Agathos et al. (2020), we estimate the probability of the binary promptly collapsing into a black hole (BH) after merger to be 96%, with the low-spin prior, or 97% with the high-spin prior. Repeating the analyses of Chatziioannou et al. (2017) and Abbott et al. (2019d), we find no evidence of a postmerger signal in the 1 s of data surrounding the time of coalescence. We obtain 90% credible upper limits on the strain amplitude spectral density and the energy spectral density of \n\n\n\n\n\n and \n\n\n\n\n\n, respectively, for a frequency of 2.5 kHz. Similar to GW170817, this upper limit is higher than any expected post-merger emission from the binary (Abbott et al. 2019d). More details on all calculations and additional analyses are provided in Appendix F.7.","Citation Text":["(Abbott et al","2019e"],"Functions Text":["Overall, we find that constraints on tides, radius, possible p\u2013g instabilities","and the EoS from GW190425 are consistent with those obtained from GW170817"],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[926,939],[948,953]],"Functions Start End":[[690,768],[851,925]]}
{"Identifier":"2016ApJ...821..107G__Schwadron_et_al._2011_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":["Schwadron et al. 2011"],"Functions Text":["For the density of neutral hydrogen in the inner heliosheath a constant nH = 0.1 cm\u22123 is assumed"],"Functions Label":["Uses"],"Citation Start End":[[1151,1172]],"Functions Start End":[[1053,1149]]}
{"Identifier":"2020MNRAS.497.3943M__Eckert_et_al._2012_Instance_1","Paragraph":"The surface brightness profile for the mosaicked image of A2199 was then extracted in concentric annuli centred at the cluster centre (RA, Dec.) = (16:28:38.21,  +39:33:02.31). This corresponds to the location of the peak X-ray flux in the cluster. The radial profile of the surface brightness is shown in the upper panel of Fig. 2. In the lower panel of this figure, we present the radial profile of the azimuthally averaged electron density of A2199. The profile is recovered from the deprojection of the median surface brightness profile using the onion peeling technique (Ettori et al. 2010), and assuming that the ICM plasma is spherical symmetry. To convert from surface brightness to density, we use the following widely used approach (Eckert et al. 2012; Tchernin et al. 2016; Ghirardini et al. 2018; Ghirardini et al. 2019; Walker et al. 2020). Using xspec and the response files for XMM\u2013Newton, the conversion factor between APEC normalization and X-ray count rate in the 0.7\u20131.2 keV band is found. In the 0.7\u20131.2 keV band, this is largely independent of the temperature of the gas, so this allows a direct conversion from the deprojected X-ray surface brightness profile to a deprojected profile of APEC normalization. The APEC normalization is related to the gas density by the equation:\n(1)$$\\begin{eqnarray*}\r\n{\\rm Norm} = \\frac{10^{-14}}{4 \\pi [d_\\mathrm{A}(1+z)]^2 }\\int n_\\mathrm{ e} n_\\mathrm{H} \\mathrm{d}V,\r\n\\end{eqnarray*}$$where the electron and ion number densities are ne and nH, respectively (for a fully ionized plasma these are related by ne = 1.17nH, Asplund et al. 2009), the angular diameter distance to the cluster is dA, and z is the cluster redshift. The deprojected APEC normalizations are then converted to deprojected density in the usual fashion, assuming spherical symmetry and a constant density in each shell, by calculating the projected volumes, V, of each shell in the 2D annuli. When performing this conversion, we used a column density of 0.08 \u00d7 1022 cm\u22122 (from Kalberla et al. 2005), and using the abundance tables of Asplund et al. (2009).","Citation Text":["Eckert et al. 2012"],"Functions Text":["To convert from surface brightness to density, we use the following widely used approach"],"Functions Label":["Uses"],"Citation Start End":[[743,761]],"Functions Start End":[[653,741]]}
{"Identifier":"2022MNRAS.514.2407S__Matsunaga_et_al._2011_Instance_1","Paragraph":"The age distribution of the inner Galaxy is less well known. Traditionally, from photometry, the bulge has been viewed as an old structure (e.g. Zoccali et al. 2003) but this was thrown into question by spectroscopic ages of microlensed dwarfs (Bensby et al. 2013), many of which are young. Recent work by Bernard et al. (2018) constrained the age distribution of the bulge ($-5\\, \\mathrm{deg}\\lesssim b\\lt -2\\, \\mathrm{deg}$) from Hubble Space Telescope photometry of the main sequence turn-off stars, concluding that, although the bulge is predominantly old, approximately $10\\, \\mathrm{per\\, cent}$ of stars are younger than $5\\, \\mathrm{Gyr}$. This fraction increases to $\\sim 20\\, \\mathrm{per\\, cent}$ for more metal-rich ($[\\mathrm{Fe}\/\\mathrm{H}]\\gtrsim 0.2\\, \\mathrm{dex}$) stars, consistent with the Bensby et al. (2013) work. Further evidence for a predominantly old ($\\gtrsim 8\\, \\mathrm{Gyr}$) bulge comes (indirectly) from [C\/N] measurements of giant stars (Bovy et al. 2019; Hasselquist et al. 2020), although, as highlighted by Hasselquist et al. (2020), age appears to correlate with both metallicity and Galactic height of the populations. Nogueras-Lara et al. (2020a) have used the luminosity of red clump stars to conclude the majority ($\\sim 95\\, \\mathrm{per\\, cent}$) of the nuclear stellar disc formed more than $8\\, \\mathrm{Gyr}$ ago with some evidence of a more recent ($\\lt 1\\, \\mathrm{Gyr}$ ago) star formation burst (Matsunaga et al. 2011). This is consistent with ongoing\/recent star formation within the central molecular zone (Morris & Serabyn 1996) and is broadly consistent with the conclusions of Bernard et al. (2018) on the wider bar\/bulge but possibly suggesting the nuclear stellar disc is on average older than the surrounding bulge. We have fitted by-eye a very simple star formation history to the \u2018cleanest\u2019 combined fit from Bernard et al. (2018) of the form $\\mathrm{sech}^2((13.5\\, \\mathrm{Gyr}-\\tau)\/4.7\\, \\mathrm{Gyr})$ with a truncation at $14\\, \\mathrm{Gyr}$. Combining the metallicity distributions and star formation histories with the PARSEC isochrones and adopting a Kroupa (2001) initial mass function, we have computed the luminosity function of the giant branch stars in the inner bulge region. We show the results in the lower panel of Fig. 4 along with a simple double Gaussian plus quadratic fit to represent the red clump stars, the red giant branch bump stars and the red giant branch stars respectively. We find that the lowest latitude bin has a red clump magnitude of $M_{K_s,\\mathrm{RC}}=-1.61\\, \\mathrm{mag}$. This agrees well with the mean solar neighbourhood result from Chan & Bovy (2020) of $M_{K_s,\\mathrm{RC}}=-1.622\\, \\mathrm{mag}$ and more specifically using their relations adopting the mean (J \u2212 Ks) = 0.647 (see Appendix B1) and mean metallicity $-0.18\\, \\mathrm{dex}$ gives $M_{K_s,\\mathrm{RC}}=-1.595\\, \\mathrm{mag}$. The metallicity gradient with latitude produces a red clump magnitude gradient of $0.032\\, \\mathrm{mag}\\, \\mathrm{deg}^{-1}$ whilst using the change in mean metallicity in combination with the results of Chan & Bovy (2020) we would expect $0.024\\, \\mathrm{mag}\\, \\mathrm{deg}^{-1}$. At all latitudes the red clump distribution is well reproduced by a Gaussian with standard deviation $\\sim 0.11\\, \\mathrm{mag}$. Chan & Bovy (2020) measured the solar neighbourhood red clump to have an intrinsic standard deviation of $0.097\\, \\mathrm{mag}$ which combined in quadrature with that arising from the metallicity variance predicts a standard deviation of $\\sim 0.13\\, \\mathrm{mag}$, similar to the PARSEC models. The red clump peaks from the PARSEC isochrones have a slight bimodal structure arising from the bimodal metallicity distributions such that the mode typically peaks $\\sim 0.03\\, \\mathrm{mag}$ fainter than the Gaussian mean.","Citation Text":["Matsunaga et al. 2011"],"Functions Text":["Nogueras-Lara et al. (2020a) have used the luminosity of red clump stars to conclude the majority ($\\sim 95\\, \\mathrm{per\\, cent}$) of the nuclear stellar disc formed more than $8\\, \\mathrm{Gyr}$ ago with some evidence of a more recent ($\\lt 1\\, \\mathrm{Gyr}$ ago) star formation burst","This is consistent with ongoing\/recent star formation within the central molecular zone","and is broadly consistent with the conclusions of Bernard et al. (2018) on the wider bar\/bulge but possibly suggesting the nuclear stellar disc is on average older than the surrounding bulge."],"Functions Label":["Background","Similarities","Similarities"],"Citation Start End":[[1444,1465]],"Functions Start End":[[1157,1442],[1468,1555],[1580,1771]]}
{"Identifier":"2020MNRAS.494.5576P__Pastorello_et_al._2018_Instance_1","Paragraph":"Another interesting type of transient to compare DES17X1boj and DES16E2bjy with are the SN impostors. As shown in Fig. 3, SN2009ip has a short phase of re-brightening around the same phase as the secondary peak of the DES-SN transients, and its peak brightness (MV = \u221217.7; see e.g. Fraser et al. 2013) is similar to DES16E2bjy. However, several other features distinguish our double-peaked DES-SN transients from the SN impostors. While SN2009ip does show re-brightening, its light-curve evolution is clearly different from the DES transients. Additionally, other impostor candidates such as SN2015bh (see e.g. Elias-Rosa et al. 2016) and SN2016bdu (Pastorello et al. 2018) have very similar light curves with SN2009ip, but do not exhibit rebrightening. Furthermore, our photometric data also constrain the long-term variability of DES17X1boj to a level below what was seen in SN2009ip (Pastorello et al. 2013) and SN2016bdu (Pastorello et al. 2018) in the years before the brightest event (MV in range \u221213 to \u221214). For the more distant event DES16E2bjy, such outbursts would have been below our detection threshold. Regarding the spectroscopic data, the impostors exhibit strong, narrow H and He lines around peak brightness (see e.g. Fraser et al. 2013; Mauerhan et al. 2013; Pastorello et al. 2013). No such features are seen in either of the DES transients (see Figs 6 and 9), but it is possible that the lines are hidden in the noise.To investigate this, we estimated the limiting equivalent width (EW) for a Gaussian-shaped narrow H\u2009\u03b1 line with $v$FWHM = 500 \u2009km\u2009s\u22121 in our spectra. H lines with similar widths are often seen in both SN impostors (Smith et al. 2011) and in SNe IIn (Taddia et al. 2013) where $v$FWHM \u223c 100\u22121000 \u2009km\u2009s\u22121 are typically measured. For the given configuration, we found limits of EW \u2272 5 \u00c5 for DES17X1boj and EW \u2272 14 \u00c5 for DES16E2bjy. In the case of SNe IIn, the line strengths are typically measured in several tens to hundreds of \u00c5ngstroms (EW\u2273 40 \u00c5; Smith, Mauerhan & Prieto 2014), and thus it is unlikely that narrow H\u2009\u03b1 lines are hiding in the spectra. Due to both photometric and spectroscopic differences, it is unlikely that DES17X1boj and DES16E2bjy are events similar to SN impostors.","Citation Text":["Pastorello et al. 2018"],"Functions Text":["Furthermore, our photometric data also constrain the long-term variability of DES17X1boj to a level below what was seen in SN2009ip","and SN2016bdu","in the years before the brightest event (MV in range \u221213 to \u221214)."],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[927,949]],"Functions Start End":[[755,886],[912,925],[951,1016]]}
{"Identifier":"2016ApJ...833....7Y__Owen_&_Wu_2013_Instance_2","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"],"Functions Text":["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"],"Functions Label":["Uses"],"Citation Start End":[[1966,1980]],"Functions Start End":[[1792,1964]]}
{"Identifier":"2019MNRAS.485.3715B__Reid_et_al._2014_Instance_1","Paragraph":"The presumed WD cooling age is \u22735 Gyr. If it has a low mass, this age limit can be larger since after the Roche lobe detachment a proto-WD goes through the contraction phase until it reaches its cooling track (Istrate et al. 2014, 2016). The duration of this phase increases as the mass of the proto-WD decreases, and may last as long as \u223c2 Gyr. The J0740 characteristic age of 3.75 Gyr (Table 1) is smaller than the WD age estimate. However, the observed pulsar period derivative and consequently its characteristic age can be biased by kinematic effects, i.e. the effects of the pulsar proper motion (Shklovskii effect; Shklovskii 1970), the acceleration towards the Galactic plane and the acceleration due to differential Galactic rotation (Nice & Taylor 1995). Using the J0740 proper motion value from Table 1, the Sun\u2019s Galactocentric velocity and the distance (240\u2009km s\u22121 and 8.34 kpc, respectively; Reid et al. 2014), we calculated these corrections to the pulsar period derivative: $\\dot{P}_{\\rm S}=3.0\\times 10^{-21}$, $\\dot{P}_{\\rm G,\\perp }=-1.6\\times 10^{-22}$, $\\dot{P}_{\\rm G,p}=3.8\\times 10^{-23}$ for the minimum and $\\dot{P}_{\\rm S}=6.9\\times 10^{-21}$, $\\dot{P}_{\\rm G,\\perp }=-2.2\\times 10^{-22}$, $\\dot{P}_{\\rm G,p}=8.9\\times 10^{-23}$ for the maximum pulsar distance estimates. The corresponding intrinsic characteristic ages are \u03c4i \u223c 5 and \u223c8.5 Gyr, which are compatible with the cooling age.8 Thus, the considered binary system indeed can be very old and the presumed WD can be ultracool. This is not a unique situation. There are other examples of the objects with similar characteristics. One of them is the WD companion of PSR J0751+1807 (Bassa et al. 2006). Its colours (see Figs 2 and 3) indicate that the WD has a pure helium or mixed H\/He atmosphere with a temperature T \u223c 3500\u20134300 K. Other examples are isolated ultracool white dwarfs WD J1102 (Hall et al. 2008; Kilic et al. 2012) and WD 0346+246 (Oppenheimer et al. 2001). These WDs have temperatures of about 3650 and 3300 K, respectively, and are best explained by the mixed atmosphere models (Gianninas et al. 2015).","Citation Text":["Reid et al. 2014"],"Functions Text":["Using the J0740 proper motion value from Table 1, the Sun\u2019s Galactocentric velocity and the distance (240\u2009km s\u22121 and 8.34 kpc, respectively;","we calculated these corrections to the pulsar period derivative: $\\dot{P}_{\\rm S}=3.0\\times 10^{-21}$, $\\dot{P}_{\\rm G,\\perp }=-1.6\\times 10^{-22}$, $\\dot{P}_{\\rm G,p}=3.8\\times 10^{-23}$ for the minimum and $\\dot{P}_{\\rm S}=6.9\\times 10^{-21}$, $\\dot{P}_{\\rm G,\\perp }=-2.2\\times 10^{-22}$, $\\dot{P}_{\\rm G,p}=8.9\\times 10^{-23}$ for the maximum pulsar distance estimates."],"Functions Label":["Uses","Uses"],"Citation Start End":[[906,922]],"Functions Start End":[[765,905],[925,1298]]}
{"Identifier":"2021MNRAS.507.2012B__Vogelsberger_et_al._2014a_Instance_1","Paragraph":"Our simulations were run using the AREPO (Springel 2010; Pakmor, Bauer & Springel 2011; Pakmor et al. 2016; Weinberger, Springel & Pakmor 2020) moving-mesh magnetohydrodynamics (MHD) code. The code solves for gravity coupled with MHD. The gravity solver uses the PM-tree method (Barnes & Hut 1986) and the MHD solver uses a non-static unstructured grid formed by performing a Voronoi tesselation of the domain. AREPO has been used to produce simulations of the Universe at a wide range of scales. At the largest scales, we have uniform volume cosmological simulations such as the Illustris (Genel et al. 2014; Vogelsberger et al. 2014a; Nelson et al. 2015; Sijacki et al. 2015) and Illustris-TNG  (Marinacci et al. 2018; Naiman et al. 2018; Nelson et al. 2018; Pillepich et al. 2018b; Springel et al. 2018; Nelson et al. 2019a,b; Pillepich et al. 2019) suites. These simulations have box sizes ranging from \u223c50 to \u223c300\u2009Mpc and baryonic mass resolutions ranging from \u223c105 to 107\u2009M\u2299. They have been largely successful in producing galaxy and SMBH populations consistent with observations, in Illustris (Vogelsberger et al. 2014b; Sales et al. 2015; Sijacki et al. 2015) and in TNG (Genel et al. 2018; Weinberger et al. 2018; Pillepich et al. 2018b; Donnari et al. 2019; Rodriguez-Gomez et al. 2019; Torrey et al. 2019; Habouzit et al. 2021; \u00dcbler et al. 2021). At the smallest scales, we have cosmological zoom simulation suites such as AURIGA (Grand et al. 2017) for individual milky-type galaxies, and HESTIA (High-resolutions Environmental Simulations of The Immediate Area) (Libeskind et al. 2020) for the Local Group. These simulations have been successful in reproducing observational results for the internal structures of galaxies  (Bl\u00e1zquez-Calero et al. 2020; Cautun et al. 2020; Grand et al. 2020). All these developments make AREPO an ideal tool for the development of black hole models, which require a reliable modelling of the necessary physics over a large dynamic range.","Citation Text":["Vogelsberger et al. 2014a"],"Functions Text":["AREPO has been used to produce simulations of the Universe at a wide range of scales. At the largest scales, we have uniform volume cosmological simulations such as the Illustris","suites. These simulations have box sizes ranging from \u223c50 to \u223c300\u2009Mpc and baryonic mass resolutions ranging from \u223c105 to 107\u2009M\u2299."],"Functions Label":["Background","Background"],"Citation Start End":[[610,635]],"Functions Start End":[[411,589],[853,981]]}
{"Identifier":"2021MNRAS.504..228S__Lyne_&_Manchester_1988_Instance_1","Paragraph":"In addition, broad-band polarization observations provide essential information about the pulsar radio emission mechanism, beam geometry, and the Galactic magneto-ionic ISM. Pulsars are among the most highly polarized radio sources known (e.g. Lyne & Smith 1968; Gould & Lyne 1998), and the polarization varies with observing frequency (e.g. Manchester, Taylor & Huguenin 1973; Johnston et al. 2008; Dai et al. 2015), providing insight into the magnetospheric emission and propagation mechanisms. In addition, the linear polarization P.A.s across pulse phase can constrain the beam size and inclination angles, with respect to the pulsar\u2019s rotation axis and our line of sight (LoS). For example, the rotating vector model (RVM) predicts a smooth \u2018S\u2019-shape, due to the projected vectors of the magnetic field lines as they sweep across our LoS (e.g. Radhakrishnan & Cooke 1969; Lyne & Manchester 1988; Johnston et al. 2007). Many pulsars show more complex P.A. curves, particularly discontinuities with rapid jumps of \u224890\u00b0, which suggests the presence of two orthogonal polarization modes (OPM; e.g. Manchester 1975; Backer, Rankin & Campbell 1976). Furthermore, circular polarization across the pulse is observed to either remain in the same hand or change sense (e.g. Radhakrishnan & Rankin 1990), which may be intrinsic to the emission mechanism or due to propagation effects (e.g. Han et al. 1998; Kennett & Melrose 1998). Additional diagnostics of magnetospheric effects have also been investigated, including variations in Faraday rotation measure (RM) and circular polarization across the pulse (e.g. Ramachandran et al. 2004; Karastergiou 2009; Noutsos et al. 2009; Ilie, Johnston & Weltevrede 2019). Although pulsars were discovered over 50 yr ago (Hewish et al. 1968), it is clear that current understanding and models of the emission mechanism are far from replicating this wide range of observed behaviour, as well as additional emission phenomena such as nulling and mode changing.","Citation Text":["Lyne & Manchester 1988"],"Functions Text":["In addition, the linear polarization P.A.s across pulse phase can constrain the beam size and inclination angles, with respect to the pulsar\u2019s rotation axis and our line of sight (LoS). For example, the rotating vector model (RVM) predicts a smooth \u2018S\u2019-shape, due to the projected vectors of the magnetic field lines as they sweep across our LoS (e.g."],"Functions Label":["Background"],"Citation Start End":[[877,899]],"Functions Start End":[[497,848]]}
{"Identifier":"2018MNRAS.476.3631P__Shi_&_Sheth_2018_Instance_1","Paragraph":"The left-hand panel of Fig. 3 shows the scatter plot of b1 and mass, coloured by $1+\\delta _{5\\,h^{-1}\\,{\\rm Mpc}}$. There is an obvious correlation visible, with a largely vertical trend in which b1 increases monotonically with $\\delta _{5\\,h^{-1}\\,{\\rm Mpc}}$. The symbols with errors show the median bias as a function of mass, in four bins of $\\delta _{5\\,h^{-1}\\,{\\rm Mpc}}$ and averaged over 10 realizations of the default box. It is clear that, at fixed $\\delta _{5\\,h^{-1}\\,{\\rm Mpc}}$, the trend of bias with halo mass is weak. This trend is consistent with previous results in the literature, which have shown that large-scale bias is more strongly correlated with halo-centric overdensity than it is with halo mass (see e.g. Abbas & Sheth 2007; Shi & Sheth 2018). The right-hand panel of the figure explores this further, showing the Spearman rank correlation coefficient between b1 and \u03b4R as a function of halo mass, for R = 2, 3, 5\u2009h\u22121\u2009Mpc. We see that the strength of the correlation is only a weak function of mass for each smoothing scale, but monotically increases with R. This increase with R is not surprising, since our estimator for b1 itself is ultimately measuring a large-scale halo-centric overdensity, so that b1 and \u03b4R are measuring essentially the same quantity for large R. To appreciate this point better, Fig. 4 shows a visualization of the haloes in a subvolume of our high-resolution box, with haloes shown as circles whose radii scale with R200b and whose colour scales with halo bias b1 as indicated by the colour bar. The panels focus on massive (top) and low mass haloes (bottom). We discuss some connections between halo-by-halo bias and gravitational redshift measurements (see e.g. Wojtak, Hansen & Hjorth 2011; Croft 2013; Alam et al. 2017) in Appendix D. In Appendix B1, we present analytical arguments that explain the size of the scatter in b1 at fixed mass and also qualitatively reproduce the trends seen in Fig. 3.","Citation Text":["Shi & Sheth 2018"],"Functions Text":["This trend is consistent with previous results in the literature, which have shown that large-scale bias is more strongly correlated with halo-centric overdensity than it is with halo mass (see e.g."],"Functions Label":["Similarities"],"Citation Start End":[[756,772]],"Functions Start End":[[537,735]]}
{"Identifier":"2021MNRAS.500.3368S__Dubinski_1998_Instance_1","Paragraph":"Like normal elliptical galaxies, the analysis of the stellar populations in the inner regions of BCGs indicates that the bulk of their stars was formed rapidly in a very intense starburst at redshift z > 2 (Renzini 2006). However, despite sharing similar morphologies with normal massive ellipticals, in addition to red colours, old and metal-rich stellar populations and alpha-enhancements (Brough et al. 2008; Loubser et al. 2009; Donahue et al. 2010; Loubser & S\u00e1nchez-Bl\u00e1zquez 2012; Barbosa et al. 2016; Edwards et al. 2020), BCGs constitute a special category of objects with peculiar star formation histories (SFHs) seen from both observations (Tran et al. 2008; Barbosa et al. 2016) and models (Dubinski 1998; De Lucia & Blaizot 2007). The evolution of BCGs is significantly affected by their surrounding environments. Due to their central positions in the gravitational potential well of their host clusters, central cluster galaxies accrete stars and gas from satellite galaxies that orbit around them and fall in, developing extended light profiles. The more representative SFHs of their dominant stellar population include components from in situ star formation, and from the interaction with other galaxies and with the intracluster medium, where the outer regions in BCGs are continually assembling mass through minor mergers (Cooke et al. 2019). The stellar populations of a large sample of observed BCGs have been recently studied in Edwards et al. (2020), from the galaxy core into the intracluster light (ICL) out to 4 effective radii, finding old stellar populations of \u223c13 Gyr and high metallicities [Fe\/H] \u223c0.3 in the galaxy cores, whereas the average age in the ICL is estimated to be slightly younger, \u223c9.2 Gyr with lower metallicities \u22120.4  [Fe\/H]  0.2 at 40 kpc. This broadly supports the idea of two-phase galaxy formation, with the BCG cores and inner regions formed faster and earlier than the outer regions that have formed more recently, or have accreted mass afterwards through galaxy mergers and thereby also increasing in size (Oser et al. 2010; Kubo et al. 2017; Cooke et al. 2019).","Citation Text":["Dubinski 1998"],"Functions Text":["BCGs constitute a special category of objects with peculiar star formation histories (SFHs) seen from both observations","and models"],"Functions Label":["Background","Background"],"Citation Start End":[[702,715]],"Functions Start End":[[530,649],[690,700]]}
{"Identifier":"2017AandA...606A..17M__Conselice_et_al._2003_Instance_1","Paragraph":"Studies of SMGs over the past few tens of years have provided valuable insights into their properties. These include the redshift distribution (e.g. Chapman et al. 2005; Aretxaga et al. 2007; Wardlow et al. 2011; Yun et al. 2012; Smol\u010di\u0107 et al. 2012; Simpson et al. 2014, 2017; Zavala et al. 2014; Miettinen et al. 2015a; Chen et al. 2016a; Strandet et al. 2016; Danielson et al. 2017; Brisbin et al. 2017), spatial clustering and environment (e.g. Ivison et al. 2000; Blain et al. 2004; Aravena et al. 2010; Hickox et al. 2012; Miller et al. 2015; Chen et al. 2016b; Wilkinson et al. 2017; Smol\u010di\u0107 et al. 2017), merger incidence (e.g. Conselice et al. 2003), and circumgalactic medium (Fu et al. 2016). Regarding the intrinsic physical characteristics of SMGs, the properties studied so far include the sizes and morphologies (e.g. Swinbank et al. 2010; Men\u00e9ndez-Delmestre et al. 2013; Aguirre et al. 2013; Targett et al. 2013; Chen et al. 2015; Simpson et al. 2015; Ikarashi et al. 2015; Miettinen et al. 2015b, 2017b,c; Hodge et al. 2016), panchromatic spectral energy distributions (SEDs; e.g. Micha\u0142owski et al. 2010; Magnelli et al. 2012; Swinbank et al. 2014; da Cunha et al. 2015; Miettinen et al. 2017a), stellar masses (e.g. Dye et al. 2008; Hainline et al. 2011; Micha\u0142owski et al. 2012; Targett et al. 2013), gas masses (e.g. Greve et al. 2005; Tacconi et al. 2006, 2008; Engel et al. 2010; Ivison et al. 2011; Riechers et al. 2011; Bothwell et al. 2013; Huynh et al. 2017), gas kinematics (e.g. Alaghband-Zadeh et al. 2012; Hodge et al. 2012; Carilli & Walter 2013; Olivares et al. 2016), and active galactic nucleus (AGN) incidence (Alexander et al. 2003, 2005; Laird et al. 2010; Johnson et al. 2013; Wang et al. 2013). The role played by SMGs in a broader context of galaxy formation and evolution has also been investigated through models (e.g. Baugh et al. 2005; Fontanot et al. 2007; Dav\u00e9 et al. 2010; Gonz\u00e1lez et al. 2011; Hayward et al. 2013) and observational approach (e.g. Swinbank et al. 2006; Toft et al. 2014; Simpson et al. 2014). ","Citation Text":["Conselice et al. 2003"],"Functions Text":["Studies of SMGs over the past few tens of years have provided valuable insights into their properties. These include","merger incidence (e.g."],"Functions Label":["Background","Background"],"Citation Start End":[[636,657]],"Functions Start End":[[0,116],[613,635]]}
{"Identifier":"2021MNRAS.500.1817L__Abbott_et_al._2020b_Instance_2","Paragraph":"Since the errors of the LIGO-estimated rates are dominated by Poisson statistics (Abbott et al. 2020a,b), we approximate the PDF for the expected number of detections $\\mathcal {N}=\\mathcal {R}VT$ (from the surveyed space\u2013time volume VT) by $\\mathrm{d}P\/\\mathrm{d}\\mathcal {N}\\propto \\mathcal {N}^{k-1\/2}\\mathrm{e}^{-\\mathcal {N}}\/k!$, where k = 1 for each of the relevant cases ($\\mathcal {R}_{190814}$, $\\mathcal {R}_{170817}$, and $\\mathcal {R}_{190425}$), and the factor of $\\mathcal {N}^{-1\/2}$ is from Jeffrey\u2019s prior (Abbott et al. 2020a). From the median values of $\\bar{\\mathcal {R}}_{190814}=7\\rm \\, Gpc^{-3}\\, yr^{-1}$ (Abbott et al. 2020a), $\\bar{\\mathcal {R}}_{\\rm 170817}=760\\rm \\, Gpc^{-3}\\, yr^{-1}$, and $\\bar{\\mathcal {R}}_{\\rm 190425}=460\\rm \\, Gpc^{-3}\\, yr^{-1}$ (Abbott et al. 2020b), we obtain the effective surveyed space\u2013time volumes $VT=1.2\/\\bar{\\mathcal {R}}$ for each of these three events (\u20181.2\u2019 is the median of $\\mathrm{d}P\/\\mathrm{d}\\mathcal {N}$). We consider both GW170817 and GW190425 as bNS mergers, because the component masses of GW190425 are not far from those of GW170817 and the nature of the merging objects makes little practical difference in our model. Thus, the PDF of the total bNS merger rate from the sum of the two is given by a convolution of the two individual PDFs\n(1)$$\\begin{eqnarray*}\r\n{\\mathrm{d}P\\over \\mathrm{d}\\mathcal {R}_{\\rm bns}} = \\int _0^{\\mathcal {R}_{\\rm bns}} \\mathrm{d}\\mathcal {R}_1 {\\mathrm{d}P\\over \\mathrm{d}\\mathcal {R}_1} \\left.{\\mathrm{d}P\\over \\mathrm{d}\\mathcal {R}_2}\\right|_{\\mathcal {R}_{\\rm bns}-\\mathcal {R}_1},\r\n\\end{eqnarray*}$$where we have written $\\mathcal {R}_{1} = \\mathcal {R}_{170817}$, $\\mathcal {R}_{2} = \\mathcal {R}_{190425}$ for brevity. We then calculate the PDF for the inverse of the total bNS merger rate $\\mathrm{d}P\/\\mathrm{d}\\mathcal {R}_{\\rm bns}^{-1}=\\mathcal {R}_{\\rm bns}^2\\mathrm{d}P\/\\mathrm{d}\\mathcal {R}_{\\rm bns}$. Finally, the PDF of the rate ratio $\\beta =\\mathcal {R}_{190814}\/\\mathcal {R}_{\\rm bns}$ is given by\n(2)$$\\begin{eqnarray*}\r\n{\\mathrm{d}P\\over \\mathrm{d}\\beta } = \\int _0^\\infty {\\mathrm{d}\\mathcal {R}_{3}\\over \\mathcal {R}_3} {\\mathrm{d}P\\over \\mathrm{d}\\mathcal {R}_{3}} \\left.{\\mathrm{d}P\\over \\mathrm{d}\\mathcal {R}_{\\rm bns}^{-1}}\\right|_{\\beta \/\\mathcal {R}_3},\r\n\\end{eqnarray*}$$where we have written $\\mathcal {R}_{3} = \\mathcal {R}_{190814}$ for brevity. We find the 90 per cent confidence interval for the rate ratio to be in the range $0.064\\, \\rm {per\\, cent}\\lt \\beta \\lt 2.8\\, \\rm {per\\, cent}$.","Citation Text":["Abbott et al. 2020b"],"Functions Text":["From the median values of $\\bar{\\mathcal {R}}_{190814}=7\\rm \\, Gpc^{-3}\\, yr^{-1}$","$\\bar{\\mathcal {R}}_{\\rm 170817}=760\\rm \\, Gpc^{-3}\\, yr^{-1}$, and $\\bar{\\mathcal {R}}_{\\rm 190425}=460\\rm \\, Gpc^{-3}\\, yr^{-1}$","we obtain the effective surveyed space\u2013time volumes $VT=1.2\/\\bar{\\mathcal {R}}$ for each of these three events (\u20181.2\u2019 is the median of $\\mathrm{d}P\/\\mathrm{d}\\mathcal {N}$)."],"Functions Label":["Uses","Uses","Uses"],"Citation Start End":[[785,804]],"Functions Start End":[[547,629],[653,783],[807,980]]}
{"Identifier":"2021MNRAS.504.1939G__Zhang_&_Yan_2011_Instance_1","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"],"Functions Text":["Alternatively, such a change in the PA can be obtained due to magnetic reconnection, e.g. in the ICMART model","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"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1158,1174]],"Functions Start End":[[1047,1156],[1177,1376]]}
{"Identifier":"2015MNRAS.446.1799O__Fruscione_et_al._2006_Instance_1","Paragraph":"A262 (RA = 01:52:46.299, Dec. = +36:09:11.80) is a bright, nearby poor cluster at z = 0.0162 (Struble & Rood 1999) with mean ICM temperature \u22482 keV( see e.g. Vikhlinin et al. 2005, 2006; Sato, Matsushita & Gastaldello 2009; Sanders et al. 2010). Due to its low mass and temperature, it may be considered as an intermediate between clusters and groups. A262 was observed for 110.7 ks in ACIS-S and a blank-sky observation of 450 ks was used for the background fitting. We used dmcopy tool in CIAO: Chandra's data analysis system (Fruscione et al. 2006) to restrict the energy range to 0.7\u22127\u2009keV for both imaging and spectral analysis in all of the four data product files: event file, blank-sky observation, RMF and ARF. This reduced the number of PI channels to 433. Also, since the output of CIAO is in fits format, we used ftools-fv6 to export the event and background files as ASCII files. We then performed a 2\u2009arcsec binning to both the event and the background files to generate the 3D data cubes for the spectral analysis without having to re-group the energy and PI columns. We used our in-house binning software tool for binning the data; this software is also available in the online package. This reduced the size of the data to a manageable level without adversely affecting the subsequent inference of the cluster. The output is a photon counts in a grid of 256 \u00d7 256 \u00d7 433 to be read in by bayes-x. Our Bayesian framework also allows us to analyse the data from one CCD. The X-ray images (Fig. 2) were then generated by summing up the counts at each pixel. To illustrate the large-scale features in the image we also binned the events with a cell size of 16\u2009arcsec (right-hand panel of Fig. 2). It should be noted that the images are for illustration purposes only. We applied the rmfimg tool in CIAO to convert and expand RMF and ARF files into 2D images (matrices) for the spectral analysis. Similarly, we used ftools-fv to export the 2D RMF and ARF as ASCII files to be read in by bayes-x.","Citation Text":["Fruscione et al. 2006"],"Functions Text":["We used dmcopy tool in CIAO: Chandra's data analysis system","to restrict the energy range to 0.7\u22127\u2009keV for both imaging and spectral analysis in all of the four data product files: event file, blank-sky observation, RMF and ARF."],"Functions Label":["Uses","Uses"],"Citation Start End":[[529,550]],"Functions Start End":[[468,527],[552,719]]}
{"Identifier":"2015ApJ...808..157M__Nayfeh_1981_Instance_1","Paragraph":"As we have seen, the asymptotic reduction of the original CR propagation problem, given by Equation (9), to its isotropic part cannot proceed to higher orders of approximation using a simple asymptotic series in Equation (10) and requires a multi-time asymptotic expansion. In the Chapman\u2013Enskog method, the operator \n\n\n\n\n\n is expanded instead. Its purpose is to avoid unwanted higher time derivatives to appear in higher orders of approximation. This is very similar to, e.g., a secular growth in perturbed oscillations of dynamical systems. To eliminate the secular terms, one seeks to alter (also expand in a small parameter) the frequency of the zero-order motion, which is similar to the \n\n\n\n\n\n expansion. One example of such an approach may be found in a derivation of hydrodynamic equations for strongly collisional, but magnetized plasmas, starting from the Boltzmann equation (Mikhailovsky 1967). The classical monograph by Chapman & Cowling (1991; Ch.VIII) gives another example of a subdivision of the \n\n\n\n\n\n operator for solving the transport problem in a non-uniform gas-mixture. Expanding \n\n\n\n\n\n operators eliminates secular terms, such as the telegraph term. Perhaps more customary today, and equivalently, is to introduce a hierarchy of formally independent time variables (e.g., Nayfeh 1981) \n\n\n\n\n\n, so that\n14\n\n\n\n\n\nInstead of Equation (13), from Equation (9), we have\n15\n\n\n\n\n\nwhere the conditions \n\n\n\n\n\n are implied. The solution of this equation should be sought in the following form\n16\n\n\n\n\n\nwhere \n\n\n\n\n\n and \n\n\n\n\n\n are chosen to satisfy, respectively, the following two equations.\n17\n\n\n\n\n\nand\n18\n\n\n\n\n\nThe solution for \n\n\n\n\n\n is as follows\n19\n\n\n\n\n\nand it can be evaluated for arbitrary n by expanding both sides of Equation (17) in a series of eigenfunctions of the diffusion operator on its lhs:\n\n\n\n\n\nFor D = 1, for example, \n\n\n\n\n\n are the Legendre polynomials with \n\n\n\n\n\n, \n\n\n\n\n\n. The time dependent coefficients \n\n\n\n\n\n are determined by the initial values of \n\n\n\n\n\n (the anisotropic part of the initial CR distribution) and the rhs of Equation (17), that depends on \n\n\n\n\n\n, obtained at the preceding step. It is seen, however, that \n\n\n\n\n\n exponentially decay in time for \n\n\n\n\n\n and we may ignore them3\n\n3\nIn fact, we must do so because our asymptotic method has a power accuracy in \n\n\n\n\n\n, but not the exponential accuracy.\n because we are primarily interested in evolving the system over time \n\n\n\n\n\n and even longer. Starting from n = 0 and using Equation (15), for the slowly varying part of f, we have\n20\n\n\n\n\n\nThe solubility condition for \n\n\n\n\n\n (obtained by integrating both sides of Equation (15) in \u03bc) also gives a trivial result\n21\n\n\n\n\n\nso the last two conditions are consistent with the suggested decomposition in Equation (16), since from Equation (18) with n = 1, we have\n22\n\n\n\n\n\nand, thus both \n\n\n\n\n\n and \n\n\n\n\n\n are, indeed, independent of \n\n\n\n\n\n and \n\n\n\n\n\n. We have introduced the function \n\n\n\n\n\n here by the following two relations\n23\n\n\n\n\n\nThe solubility condition for \n\n\n\n\n\n yields the nontrivial and well-known (e.g., Jokipii 1966) result, which is actually the leading term of the \n\n\n\n\n\n expansion in \n\n\n\n\n\n\n\n24\n\n\n\n\n\nwhere\n\n\n\n\n\nThe solubility conditions for \n\n\n\n\n\n, \n\n\n\n\n\n will generate the higher order terms of our expansion, which, after some algebra, can be manipulated into the following expressions for the third and fourth orders of approximation\n25\n\n\n\n\n\n\n\n26\n\n\n\n\n\nWe have denoted\n\n\n\n\n\nand \n\n\n\n\n\n. The pitch-angle diffusion coefficient \n\n\n\n\n\n and magnetic focusing \u03c3 are considered z-independent for simplicity, a limitation that can be easily relaxed by rearranging the operators containing \n\n\n\n\n\n in Equation (26). We can proceed to higher orders of approximation ad infinitum since terms containing \n\n\n\n\n\n can be expressed through \n\n\n\n\n\n. According to Equations (20)\u2013(21), of interest is the evolution of \n\n\n\n\n\n on the timescales \n\n\n\n\n\n or \n\n\n\n\n\n; thus, as we already mentioned, the contributions of \n\n\n\n\n\n to all of the solubility conditions, similar to those given by Equations (24)\u2013(26), have to be dropped (because they become exponentially small) and only \n\n\n\n\n\n-contributions should be retained. Using Equations (20)\u2013(21) and (24)\u2013(26) to form the combinations \n\n\n\n\n\n and summing up both sides, on the lhs of the resulting equation, we simply obtain \n\n\n\n\n\n (see Equation (14)). Therefore, the evolution of \n\n\n\n\n\n up to the fourth order in \u03f5 takes the following form.\n27\n\n\n\n\n\nwhere \n\n\n\n\n\n, \n\n\n\n\n\n, and\n28\n\n\n\n\n\n\n","Citation Text":["Nayfeh 1981"],"Functions Text":["Perhaps more customary today, and equivalently, is to introduce a hierarchy of formally independent time variables (e.g.,",", so that"],"Functions Label":["Uses","Uses"],"Citation Start End":[[1296,1307]],"Functions Start End":[[1174,1295],[1315,1324]]}
{"Identifier":"2016ApJ...826..137M__Jewitt_et_al._2013_Instance_1","Paragraph":"We performed a preliminary, zeroth-order analysis of the images by constructing a syndyne\u2013synchrone map for each observing date. From those maps, we inferred that the activation time of the asteroid should be close in time to the discovery date, owing to the absence of dust features that could have shown up at the corresponding locations of synchrones approximately two months before discovery or earlier. In particular, no neck-line or trail features appear in the image from 2016 January 7 (PlAng \u223c 0\u00b0), which could have indicated past activity. In addition, there are no dust condensations along the direction of isolated synchrones, which could have indicated short bursts of activity (e.g., the case of P\/2012 F5 (Gibbs), Moreno et al. 2012a), or several separated short bursts, as in the case of P\/2013 P5 (Jewitt et al. 2013; Moreno et al. 2014). According to this, it is reasonable to start the search for a minimum in the function \u03c7 defined above by placing the activation time (t0) between a few days before the discovery date (102.5 days before perihelion) and about 60 days before. Regarding the duration of the activity, the smooth variation in absolute magnitudes (from Hv = 17.88 to Hv = 18.16, see Equation (1)) over the \u223c40 days period of observation and the aforementioned lack of single-synchrone dust features would suggest a long-lasting process and not an impulsive, short-duration event, such as a collision with another body. In any case, we considered both long- and short-duration events by varying HWHM in a wide range between a few days and several months in the starting simplex of the search of the five-dimensional parameter space. For the peak dust mass loss, we imposed a wide range between a minimum of 0.1 kg s\u22121 and 100 kg s\u22121, while for the velocities we set broad limits for the parameters v0, v1, and v2, so that the velocities ranged from 0 to 5 \u00d7 103 m s\u22121 (the mean velocity in the asteroid belt), and the parameter \u03b3 from 0.5 to 0, i.e., from typical gas drag to a nearly flat distribution of velocities.","Citation Text":["Jewitt et al. 2013"],"Functions Text":["In addition, there are no dust condensations along the direction of isolated synchrones, which could have indicated short bursts of activity","or several separated short bursts, as in the case of P\/2013 P5"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[815,833]],"Functions Start End":[[550,690],[751,813]]}
{"Identifier":"2017MNRAS.472.1152R__Cenko_et_al._2010_Instance_2","Paragraph":"Alternatively, if a magnetar is the central engine powering GRBs, we might expect to see periodic features in the emission. Known magnetars have clear periodic signals in their emission caused by their rotation periods (e.g. Mazets et al. 1979; Kouveliotou et al. 1998). The X-ray pulsations typically contribute to 30\u2009per\u2009cent of the signal, with a range of 10\u201380\u2009per\u2009cent (Israel et al. 1999; Kargaltsev et al. 2012; Kaspi & Beloborodov 2017). There is an energy dependence on the pulsed fraction of the signal, where low energies tend to have smaller pulsed fractions (Vogel et al. 2014). Detection of a periodic signal during the plateau phase in the X-ray light curve would provide excellent supporting evidence for the magnetar central engine model. There have been searches for a periodic signal in the prompt emission of GRBs with a number of instruments with no success, for example: Burst And Transient Source Experiment (BATSE) GRBs ( Deng & Schaefer 1997), INTErnational Gamma-Ray Astrophysics Laboratory (INTEGRAL) GRBs (Ryde et al. 2003), GRB 051103 (an extragalactic Soft Gamma-ray Repeater giant flare candidate detected by the Inter Planetary Network; Hurley et al. 2010) and Burst Alert Telescope (BAT) GRBs (Cenko et al. 2010; de Luca et al. 2010; Guidorzi et al. 2012). Dichiara et al. (2013) searched the prompt emission of a number of short GRBs for evidence of a precessing jet (predicted by Stone, Loeb & Berger 2013). However, these searches typically target the prompt emission and have not probed the regime where we might expect periodic signals from a magnetar central engine (i.e. during the plateau phase). Only two GRBs have been searched for periodic emission during the X-ray observations when the magnetar central engine may dominate the emission, GRB 060218 (Mirabal & Gotthelf 2010) and GRB 090709A (Mirabal & Gotthelf 2009; de Luca et al. 2010). The prompt emission of GRB 090709A possibly showed evidence of a periodic signal (Golenetskii et al. 2009; Gotz et al. 2009; Markwardt et al. 2009; Ohno et al. 2009), however this was ruled out with a more careful analysis of the prompt data from BAT, X-ray Telescope (XRT) and X-ray Multi-mirror Mission (XMM) observations of the X-ray afterglow (Cenko et al. 2010; de Luca et al. 2010). However, in the majority of these studies, the authors have targeted a constant spin period whereas a magnetar central engine is expected to have a rapidly decelerating spin period which would be very difficult to detect in standard searches for periodic signals. Dichiara et al. (2013) did conduct a deceleration search, however they were targeting signals in the prompt emission where we do not expect the signal from a spinning down magnetar.","Citation Text":["Cenko et al. 2010"],"Functions Text":["The prompt emission of GRB 090709A possibly showed evidence of a periodic signal","however this was ruled out with a more careful analysis of the prompt data from BAT, X-ray Telescope (XRT) and X-ray Multi-mirror Mission (XMM) observations of the X-ray afterglow","However, in the majority of these studies, the authors have targeted a constant spin period whereas a magnetar central engine is expected to have a rapidly decelerating spin period which would be very difficult to detect in standard searches for periodic signals."],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[2232,2249]],"Functions Start End":[[1884,1964],[2051,2230],[2273,2536]]}
{"Identifier":"2020ApJ...898L..33P__Delrez_et_al._2018_Instance_1","Paragraph":"For the TRAPPIST-1 system, data obtained by HST provide initial constraints on the extent and composition of the planet\u2019s atmospheres, suggesting that the four innermost planets do not have a cloud\/haze-free H2-dominated atmosphere (de Wit et al. 2016, 2018). However, follow-up work by Moran et al. (2018) have shown that HST data can also be fit to a cloudy\/hazy H2-dominated atmosphere. Complementary to HST, NASA\u2019s Spitzer Space Telescope\u2014which played a major role in the discovery and orbital determination of TRAPPIST-1d, e, f, and g (Gillon et al. 2017)\u2014has also allowed us to put additional constraints on the atmospheric composition of TRAPPIST-1b. Transit observations with Spitzer (Delrez et al. 2018) have found a +208 \u00b1 110 ppm difference between the 3.6 and 4.2 \u03bcm bands, suggesting CO2 absorption. Spitzer also showed that transit depth measurements do not show any hint of significant stellar contamination in the 4.5 \u03bcm spectral range. Morris et al. (2018) reached the same conclusion using a \u201cself-contamination\u201d approach based on the Spitzer data set. Spitzer's \u201cRed Worlds\u201d Program encompassed over 1000 hours of observations of the TRAPPIST-1 system, whose global results have been presented (Ducrot et al. 2020). HST and Spitzer measurements have also been combined with transit light curves obtained from space with K2 (Luger et al. 2017) and from the ground with the SPECULOOS-South Observatory (Burdanov et al. 2018; Gillon 2018) and Liverpool Telescope (Steele et al. 2004) where Ducrot et al. (2018) produced featureless transmission spectra for the planets in the 0.8\u20134.5 \u03bcm wavelength range, showing an absence of significant temporal variations of the transit depths in the visible. Additional ground-based observations with the United Kingdom Infra-Red Telescope, Anglo-Australian Telescope, and Very Large Telescope also show no substantial temporal variations of transit depths for TRAPPIST-1 b, c, e, and g (Burdanov et al. 2019). While the K2 optical data set detected a 3.3 day periodic 1% photometric modulation, it is not present in the Spitzer observations (Delrez et al. 2018). Further constraints on the molecular weight and presence\/absence of atmospheres on the TRAPPIST-1 planets will require additional observations with future facilities.","Citation Text":["Delrez et al. 2018"],"Functions Text":["Transit observations with Spitzer","have found a +208 \u00b1 110 ppm difference between the 3.6 and 4.2 \u03bcm bands, suggesting CO2 absorption."],"Functions Label":["Motivation","Background"],"Citation Start End":[[693,711]],"Functions Start End":[[658,691],[713,812]]}
{"Identifier":"2021ApJ...907...55M__Wang_2016_Instance_1","Paragraph":"Els\u00e4sser variables (Els\u00e4sser 1950) are usually employed in solar wind studies to separate the outward-propagating waves (denoted by z+) and the reflected or inward propagating waves (denoted by z\u2212). The separation is exact for even fully nonlinear, unidirectionally propagating waves in homogeneous and incompressible plasma, i.e., Alfv\u00e9n waves, and it even holds for radially inhomogeneous (along a purely radial magnetic field) but otherwise homogeneous plasma without nonlinear interactions (Hollweg & Isenberg 2007; Magyar et al. 2019b). However, beyond pure Alfv\u00e9n wave dynamics, it is often overlooked that transverse inhomogeneities, compressibility, and the nonlinear interaction of waves renders the separation of fluctuations into inward and outward-propagating waves inexact. For example, inhomogeneity and compressibility allows for waves (e.g., fast, slow MHD waves, surface Alfv\u00e9n waves, kink waves, etc.) that are mostly described by both Els\u00e4sser variables as they propagate (Magyar et al. 2019a, 2019b). In fact, waves other than pure Alfv\u00e9n waves generally perturb both Els\u00e4sser variables as they propagate. While kink waves, both propagating and standing, are routinely observed in the corona (e.g., Nakariakov et al. 1999; Tomczyk et al. 2007; Anfinogentov et al. 2015; Wang 2016; Nechaeva et al. 2019), evidence of surface Alfv\u00e9n waves in the solar wind is as of yet inconclusive (e.g., Horbury et al. 2001; Vasquez et al. 2001; Paschmann et al. 2013). Besides waves that are not pure Alfv\u00e9n waves, structures (inhomogeneities) advected by the solar wind also perturb both Els\u00e4sser variables (Tu & Marsch 1990, 1995; Zank et al. 2012; Adhikari et al. 2015). The nonlinear interaction of Alfv\u00e9n waves can generate purely magnetic fluctuations, 2D modes (k\u2225 = 0) or condensates which as well perturb both Els\u00e4sser variables (Boldyrev & Perez 2009; Howes & Nielson 2013). Indeed, the nature of the inward z\u2212 Els\u00e4sser variable is often not clear (Wang et al. 2018). Previous studies on Alfv\u00e9n wave dynamics in radially inhomogeneous models often mention the existence of an \u201canomalous\u201d z\u2212 component that is co-propagating with z+ (Velli et al. 1989; Verdini et al. 2009; Perez & Chandran 2013). The issue of anomalous waves is solved by Hollweg & Isenberg (2007), who showed that, while the continuously generated, reflected z\u2212 components might show up as co-propagating in a harmonic analysis, their impulse response analysis shows that these reflected Alfv\u00e9n waves still follow sunward characteristics, i.e., that there are no truly co-propagating Els\u00e4sser variables in these studies. Nevertheless, the coherence of the Els\u00e4sser variables resulting from this linear coupling of Alfv\u00e9n waves seems to influence their spectrum (Verdini et al. 2009).","Citation Text":["Wang 2016"],"Functions Text":["While kink waves, both propagating and standing, are routinely observed in the corona (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[1290,1299]],"Functions Start End":[[1126,1218]]}
{"Identifier":"2020AandA...641A.139D__Dvorak_et_al._(2015)_Instance_2","Paragraph":"The use of N-body simulations that include fragmentation allows us to perform a more detailed study of the final composition of the planets formed. In particular, we can study the water loss and\/or accretion of the final planets more realistically than in the classic models of accretion. Marcus et al. (2010) presented two empirical models for the mantle stripping in differentiated planetary embryos after a collision. The authors set a simple planet structure of two layers, assuming differentiation in core and mantle, where the mantle could be composed by silicate or ice. In this work, the authors concluded that the more energetic the collision, the more mass from the mantle is lost. Therefore, for violent collisions, water could be more easily removed. Dvorak et al. (2015) performed SPH (smoothed particle hydrodynamics) simulations and studied water loss in planetary embryos and water retained in significant fragments after a collision. They concluded that the impact velocity and the impact angle play a key role in the water loss of a planetary embryo after a collision. The investigations developed by Marcus et al. (2010) and Dvorak et al. (2015) suggest that incorporating a realistic model of volatile transport and removal in an N-body code, may lead to reduced water contents on the resulting terrestrial-like planets, in comparison with those derived from classical models that assume perfect mergers. Burger et al. (2018) studied the volatile loss and transfer. The authors focused on hit-and-run encounters using SPH simulations. They concluded that the cumulative effect of several hit-and-run collisions could efficiently strip off volatile layers of protoplanets. Driven by this, Dugaro et al. (2019) studied the water delivery in planets formed in the habitable zone (HZ), using the mantle stripping models derived by Marcus et al. (2010) in their N-body simulationswith fragmentation. The authors showed that fragmentation is not a barrier for the surviving of water worlds in the HZ, and fragments may be important in the final water content of the potentially habitable terrestrial planets formed in situ.","Citation Text":["Dvorak et al. (2015)"],"Functions Text":["The investigations developed by Marcus et al. (2010) and","suggest that incorporating a realistic model of volatile transport and removal in an N-body code, may lead to reduced water contents on the resulting terrestrial-like planets, in comparison with those derived from classical models that assume perfect mergers."],"Functions Label":["Background","Background"],"Citation Start End":[[1144,1164]],"Functions Start End":[[1087,1143],[1165,1424]]}
{"Identifier":"2021ApJ...908...95H__D\u00edaz-S\u00e1nchez_et_al._2017_Instance_1","Paragraph":"Here we outline our sample of strongly lensed Planck-selected, dusty star-forming galaxies, hereafter \u201cLPs\u201d (Table 1). Our sample of 24 LPs began with a Planck and Herschel cross-match identification of eight objects (8\/24) with continuum detections at 857 GHz (Harrington et al. 2016) greater than 100 mJy. The remaining 16\/24 LPs were selected based on continuum detections by Planck, at 857, 545, and\/or 353 GHz in the maps of all the available, clean extragalactic sky. These bright Planck point sources were then analyzed through a filtering process using a WISE color selection for the four WISE bands (3.4, 4.6, 12, 22 \u03bcm; Yun et al. 2008; D. Berman et al. 2021, in preparation). Other methods to identify strong gravitational lenses using (sub)millimeter data were independently verified by other teams using Planck and Herschel color criteria (Ca\u00f1ameras et al. 2015). The 24 LPs presented in these analyses include eight systems identified by Ca\u00f1ameras et al. (2015). The use of Planck and WISE data resulted in the discovery of the brightest known, dusty starburst galaxy at z > 1, the \u201cCosmic Eyebrow\u201d (D\u00edaz-S\u00e1nchez et al. 2017; Dannerbauer et al. 2019), which has also been independently recovered as one of the LPs presented in this survey work. Note that LPsJ1329 corresponds to the location on the sky associated with the Cosmic Eyebrow-A lens component (Dannerbauer et al. 2019). Table 1 shows the size of the lensed emission for each of the LPs, in which there are 21\/24 with lens sizes \u226410\u2033. Half of the LPs are galaxy\u2013galaxy lenses, while the other half are a mix of cluster or group lensing. The foreground lens galaxies have a negligible contribution to the observed FIR emission of the lensed galaxy (Harrington et al. 2016). The LPs have CO-based spectroscopic redshifts ranging from zCO \u223c 1.1 to 3.6 (Harrington et al. 2016, 2018; Ca\u00f1ameras et al. 2018b; this work). They are comparable or brighter in CO and FIR luminosity than other strongly lensed SPT- (Wei\u00df et al 2013; Strandet et al. 2016, 2017) or Herschel-selected; dusty star-forming galaxies (Harris et al. 2012; Bussmann et al. 2013, 2015; Yang et al. 2017). The Planck and Herschel wavelength selections preferentially target z \u223c 2\u20133 galaxies, versus the millimeter-selected SPT sources with a median closer to z \u223c 4, although with a wide range between z \u223c 2 and 7 (Wei\u00df et al 2013; Spilker et al. 2016; Strandet et al. 2016; Reuter et al. 2020).","Citation Text":["D\u00edaz-S\u00e1nchez et al. 2017"],"Functions Text":["The use of Planck and WISE data resulted in the discovery of the brightest known, dusty starburst galaxy at z > 1, the \u201cCosmic Eyebrow\u201d","which has also been independently recovered as one of the LPs presented in this survey work."],"Functions Label":["Background","Similarities"],"Citation Start End":[[1114,1138]],"Functions Start End":[[977,1112],[1166,1258]]}
{"Identifier":"2016AandA...588A..42S__Hopkins_et_al._2010_Instance_1","Paragraph":"Classical bulges (hereafter ClBs) are the central building blocks in many early-type spiral galaxies. Classical bulges might have formed as a result of major mergers during the early phase of cosmic evolution (Kauffmann et al. 1993; Baugh et al. 1996; Hopkins et al. 2009; Naab et al. 2014), or through a number of other mechanisms such as the monolithic collapse of primordial gas clouds (Eggen et al. 1962), the coalescence of giant clumps in gas-rich primordial galaxies (Noguchi 1999; Immeli et al. 2004; Elmegreen et al. 2008), violent disk instability at high-redshift (Ceverino et al. 2015), multiple minor mergers (Bournaud et al. 2007; Hopkins et al. 2010), and accretion of small companions or satellites (Aguerri et al. 2001). Although most of these studies do not provide quantitative predictions for the bulge kinematics, it is generally believed that ClBs formed through these processes have low rotation compared to the random motion. For example, Naab et al. (2014) showed that spheroids produced by minor and major mergers (which include ClBs) in full cosmological hydrodynamical simulations have a wide range of rotational properties; the massive ones have V\/\u03c3 less than 0.5. Elmegreen et al. (2008) reported dispersion dominated clump-origin ClBs with upper limit on V\/\u03c3 ~ 0.4\u22120.5, where V is the rotation velocity and \u03c3 is the central velocity dispersion. A similar study by Inoue & Saitoh (2012) suggests that clump-origin bulges have exponential surface density profiles and rotate rapidly with V\/\u03c3 ~ 0.9, resembling pseudobulges (Kormendy & Kennicutt 2004). However, using cosmological hydrodynamical simulations with continuous gas accretion, Ceverino et al. (2015) showed that massive classical bulges with non-zero angular momenta are produced at high redshift, but provided no estimate on the bulge V\/\u03c3. Overall, there is no clear quantitative picture of the rotational motion induced during the formation of classical bulges in numerical simulations. ","Citation Text":["Hopkins et al. 2010"],"Functions Text":["Classical bulges might have formed as a result of","multiple minor mergers"],"Functions Label":["Background","Background"],"Citation Start End":[[645,664]],"Functions Start End":[[102,151],[599,621]]}
{"Identifier":"2021AandA...648A...5M__Windhorst_et_al._(1990)_Instance_2","Paragraph":"Another important consistency check regards the angular size distribution of the sources. Figure 6 shows the cumulative size distributions of the final catalogs combined together, in four flux density bins (yellow solid lines). Such distributions can be considered reliable only down to a flux-dependent minimum intrinsic size (see vertical gray lines), below which most of the sources cannot be reliably deconvolved and they are conventionally assigned \u0398 = 0. The observed distributions are compared with various realizations of the cumulative distribution function described by Eq. (6), obtained by varying either the function exponent q (left and right columns respectively) or the assumed median size \u2013 flux relations (see various black lines).The original function proposed by Windhorst et al. (1990) (Eq. (6) with q = 0.62, see left column) does provide a good approximation of the observed distributions, when assuming the original \u0398med \u2212 S relation described by Eq. (7), only at flux densities S150 MHz\u227310 mJy (see long-dashed lines). This is perhaps not surprising considering that this relation was calibrated at 1.4 GHz down to a few mJy fluxes. At the lowest flux densities (S150 MHz\u22721 mJy) we need to assume a steepening of the parameter m (see Eq. (8)), to get a good match with observations (dotted line in the top left panel). This is consistent with what proposed for higher frequency deep surveys (as discussed earlier in this section). At intermediate fluxes (S150 MHz ~ 1\u221210) mJy, on the other hand, none of the discussed median size \u2013 flux relations can reproduce the observed size distribution (see second-row panel on the left). It is interesting to note, however, that if we assume a steeper exponent for the distribution function described by Eq. (7) (i.e., q = 0.80), we get a very good match with observations at all fluxes, when assuming a flux-dependent scaling factor (k = k(S); see Eq. (9)) for the Windhorst et al. (1990) median size \u2013 flux relation (black solid lines on the right). The median sizes derived from the T-RECS simulated catalogs (Bonaldi et al. 2019) also provide good results for q = 0.80 (dot-dashed lines on the right), except again at intermediate fluxes (S150 MHz ~ 1\u221210), where they show strong discrepancies with observations also in Fig. 5. This seems to indicate that the number density of extended radio galaxies in this flux density range is over-estimated in the T-RECS simulated catalogs.","Citation Text":["Windhorst et al. (1990)"],"Functions Text":["It is interesting to note, however, that if we assume a steeper exponent for the distribution function described by Eq. (7) (i.e., q = 0.80), we get a very good match with observations at all fluxes, when assuming a flux-dependent scaling factor (k = k(S); see Eq. (9)) for the","median size \u2013 flux relation (black solid lines on the right)."],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[1930,1953]],"Functions Start End":[[1652,1929],[1954,2015]]}
{"Identifier":"2017ApJ...849..123M__Ansdell_et_al._2016_Instance_1","Paragraph":"Until recently, only \u03b2 Pic and 49 Cet were known as CO-bearing debris disks (Vidal-Madjar et al. 1994; Zuckerman et al. 1995; Roberge et al. 2000). Later, more such objects were discovered, and the first statistical studies could be done (Greaves et al. 2016; Lieman-Sifry et al. 2016; P\u00e9ricaud et al. 2017). Our present list partly incorporates these earlier samples and extends them to a full list of 17 dust-rich debris disks in the solar neighborhood (Section 2). In Figure 3(a), we display the 12CO 2\u20131 (or 12CO 3\u20132 for HR 4796) line luminosities (or upper limits), as a function of the fractional luminosity for our sample. Line luminosities of detected CO-bearing disks span almost two orders of magnitude, in which the brightest disks have luminosities that are comparable to those of fainter Herbig Ae and T Tauri disks (Ansdell et al. 2016; P\u00e9ricaud et al. 2017). Since the sensitivity of the HR 4796 observation is nearly two orders of magnitude worse than that of the other measurements, we discard this object from the following analysis, reducing the size of our statistical sample to 16. For the other objects, by adopting the highest upper limit (source No. 7) in Figure 3(a), we derived a detectability threshold of \u223c1.4 \u00d7104 Jy km s\u22121 pc2 for the 12CO 2\u20131 line luminosity. With our 3 new discoveries, we found 11 disks in this sample that harbor CO gas, resulting in a very high detection rate of 68.8\n\n\n\n\n\n\n\n\u2212\n13.1\n\n\n+\n8.9\n\n\n%\n\n\n. Because of the small sample, we computed the uncertainties (corresponding to 68% confidence interval) using the binomial distribution approach proposed by Burgasser et al. (2003). Our result indicates that the presence of CO gas in dust-rich debris disks around young A-type stars is more likely the rule than the exception. The obtained incidence rate of 11\/16 is valid above our detectability threshold for the 12CO J = 2\u20131 line luminosities. Nevertheless, we cannot rule out that all of our targets harbor CO gas at some level. Remarkably, as Figure 3(a) shows, apart from HR 4796, all of the disks with \n\n\n\n\n\n\nL\n\n\ndisk\n\n\n\n\/\n\n\n\nL\n\n\nbol\n\n\n>\n2\n\u00d7\n\n\n10\n\n\n\u2212\n3\n\n\n\n\n contain detectable levels of CO gas.","Citation Text":["Ansdell et al. 2016"],"Functions Text":["Line luminosities of detected CO-bearing disks span almost two orders of magnitude, in which the brightest disks have luminosities that are comparable to those of fainter Herbig Ae and T Tauri disks"],"Functions Label":["Background"],"Citation Start End":[[830,849]],"Functions Start End":[[630,828]]}
{"Identifier":"2022MNRAS.516.3900A__Calmonte_et_al._2016_Instance_1","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":["Calmonte et al. 2016"],"Functions Text":["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"],"Functions Label":["Uses"],"Citation Start End":[[794,814]],"Functions Start End":[[643,792]]}
{"Identifier":"2019MNRAS.482.3288G__Mayer_2013_Instance_1","Paragraph":"The orbital decay of BSBHs may slow down or stall at \u223cpc scales (e.g. Begelman et al. 1980; Milosavljevi\u0107 & Merritt 2001; Zier & Biermann 2001; Yu 2002; Vasiliev, Antonini & Merritt 2014; Dvorkin & Barausse 2017; Tamburello et al. 2017), or the barrier may be overcome in gaseous environments (e.g. Gould & Rix 2000; Escala et al. 2004; Hayasaki, Mineshige & Sudou 2007; Hayasaki 2009; Cuadra et al. 2009; Lodato et al. 2009; Chapon, Mayer & Teyssier 2013; Rafikov 2013; del Valle et al. 2015), in triaxial or axisymmetric galaxies (e.g. Yu 2002; Berczik et al. 2006; Preto et al. 2011; Khan et al. 2013, 2016; Vasiliev, Antonini & Merritt 2015; Gualandris et al. 2017; Kelley, Blecha & Hernquist 2017a), and\/or by interacting with a third SMBH in hierarchical mergers (e.g. Valtonen 1996; Blaes, Lee & Socrates 2002; Hoffman & Loeb 2007; Kulkarni & Loeb 2012; Tanikawa & Umemura 2014; Bonetti et al. 2018). The accretion of gas and the dynamical evolution of BSBHs are likely to be coupled (Ivanov, Papaloizou & Polnarev 1999; Armitage & Natarajan 2002; Haiman, Kocsis & Menou 2009; Bode et al. 2010, 2012; Farris, Liu & Shapiro 2010, 2011; Kocsis, Haiman & Loeb 2012; Shi et al. 2012; D\u2019Orazio, Haiman & MacFadyen 2013; Shapiro 2013; Farris et al. 2014, 2015) such that the occurrence rate of BSBHs depends on the initial conditions and gaseous environments at earlier phases (e.g. thermodynamics of the host galaxy interstellar medium; Dotti et al. 2007, 2009; Dotti, Sesana & Decarli 2012; Fiacconi et al. 2013; Mayer 2013; Tremmel et al. 2018). Quantifying the occurrence rate of BSBHs at various merger phases is therefore important for understanding the associated gas and stellar dynamical processes. This is a challenging problem for three main reasons. First, BSBHs are expected to be rare (e.g. Foreman, Volonteri & Dotti 2009; Volonteri, Miller & Dotti 2009), and only a fraction of them accrete enough gas to be \u2018seen\u2019. Secondly, the physical separations of BSBHs that are gravitationally bound to each other (\u2272a few pc) are too small for direct imaging. Even VLBI cannot resolve BSBHs except for in the local universe (Burke-Spolaor 2011). CSO 0402+379 (discovered by VLBI as a double flat-spectrum radio source separated by 7 pc) remains the only secure case known (Rodriguez et al. 2006; Bansal et al. 2017, see Kharb, Lal & Merritt 2017; however, for a possible 0.35-pc BSBH candidate in NGC 7674). Thirdly, various astrophysical processes complicate their identification such as bright hot spots in radio jets (e.g. Wrobel, Walker & Fu 2014b). Until recently, only a handful cases of dual active galactic nuclei (AGNs) \u2013 galactic-scale progenitors of BSBHs \u2013 were known (Owen et al. 1985; Junkkarinen et al. 2001; Komossa et al. 2003; Ballo et al. 2004; Hudson et al. 2006; Max, Canalizo & de Vries 2007; Bianchi et al. 2008; Guidetti et al. 2008). While great strides have been made in identifying dual AGNs at kpc scales (e.g. Gerke et al. 2007; Comerford et al. 2009, 2012, 2015; Green et al. 2010; Liu et al. 2010, 2013, 2018; Fabbiano et al. 2011; Fu et al. 2011, 2012, 2015a,b; Koss et al. 2011, 2012, 2016; Rosario et al. 2011; Teng et al. 2012; Woo et al. 2014; Wrobel, Comerford & Middelberg 2014a; McGurk et al. 2015; M\u00fcller-S\u00e1nchez et al. 2015; Shangguan et al. 2016; Ellison et al. 2017; Satyapal et al. 2017), there is no confirmed BSBH at sub-pc scales (for recent reviews, see e.g. Popovi\u0107 2012; Burke-Spolaor 2013; Bogdanovi\u0107 2015; Komossa & Zensus 2016).","Citation Text":["Mayer 2013"],"Functions Text":["The accretion of gas and the dynamical evolution of BSBHs are likely to be coupled","such that the occurrence rate of BSBHs depends on the initial conditions and gaseous environments at earlier phases (e.g. thermodynamics of the host galaxy interstellar medium;"],"Functions Label":["Background","Background"],"Citation Start End":[[1516,1526]],"Functions Start End":[[908,990],[1262,1438]]}
{"Identifier":"2015ApJ...806....1M__Ahn_et_al._2012_Instance_1","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"],"Functions Text":["The","BOSS project","obtained additional imaging data of about 3000 deg2"],"Functions Label":["Background","Background","Background"],"Citation Start End":[[606,621]],"Functions Start End":[[554,557],[592,604],[643,694]]}
{"Identifier":"2021AandA...650A.133J__Goss_&_Field_(1968)_Instance_1","Paragraph":"Collisions with charged particles, namely with electrons and heavier ions are not considered in the radiative-transfer calculations presented in this paper, but may play an important role in the excitation of the CH \u039b-doublet, particularlyin regions with high electron fractions, \n\n$x_{\\textrm{e}} = n_{\\textrm{e}}\/(n_{\\textrm{H}}  2n_{\\textrm{H}_{2}}) >10^{-5}$xe=ne\/(nH+2nH2)>10\u22125\n\u201310\u22124. Such high electron fractions are ubiquitous in the diffuse molecular clouds present along the lines of sight studied here and may even be prevalent in the clouds surrounding the observed HII regions themselves, making electrons an important collision partner at low gas temperatures T \u2264 100 K. Bouloy & Omont (1977, 1979) studied the impact of collisional excitation by electrons on \u039b-doublet transitions with particular emphasis on the ground-state \u039b-doublet transitions of OH. These authors compute the collisional rate coefficients for collisions with electrons, either using perturbation methods such as those used by Goss & Field (1968) or using the Born approximation, both of which yield comparable results. These authors concluded that collisions with electrons, while incapable of inducing level inversion in the ground-state lines of OH at 18 cm, are responsible for thermalising them. Bouloy et al. (1984) further studied the excitation conditions of the ground-state lines of CH, and model the excitation by considering the radiative and collisional (de-) excitation of CH with H, H2, and electrons. Their results once again point to the role played by the collisions with electrons in thermalising the CH lines rather than inverting them. However, the excitation temperature of OH is found to be a 1\u20132 K above TCMB as derived from the resolved optical spectra of the OH A \u2212 X band transitions (Felenbok & Roueff 1996) or when measured by comparing the emission and absorption profiles of the radio L-band transitions of OH (Liszt & Lucas 1996, and references therein). This implies that densities much higher than the critical density are needed for thermalisation, which might also be the case for CH. More recently, Goldsmith & Kauffmann (2017) examined the impact of electron excitation on high-dipole-moment molecules like HCN, HCO+, CS, and CN in various interstellar environments. As long-range forces dominate the collisional cross-sections for electron excitation, the cross-sections and, in turn, the collisional rate coefficients scale with the square of the electric dipole moment, \u03bce. Hence, the electron collisional rate coefficients for CH with \u03bce = 1.46 D (Phelps & Dalby 1966) are \u2243 25 % of those of HCN with \u03bce = 2.985 D (Ebenstein & Muenter 1984). Scaling the value of the HCN\u2013e\u2212 collisional rate coefficient at T \u2264 100 K from Faure et al. (2007), we find the CH\u2013e\u2212 collisional rates to be of the order of ~ 2 \u00d7 10\u22125 cm3 s\u22121. From this we can compute the critical electron fractional abundance, x* (e\u2212), which defines the fractional abundance of electrons required for the collisional rate coefficients with electrons to be the same as that with H2 such that x* (e\u2212) = ncrit(e\u2212)\u2215ncrit(H2). Under the validity of these assumptions, x*(e\u2212) for CH is ~10\u22126, making CH likely to be affected by electron excitation. Therefore, a complete treatment of the radiative\u2013collisional (de-)excitation of the CH ground state would still require the availability of accurate collisional rate coefficients for collisional excitation by electrons. However, to our knowledge, these are currently not available.","Citation Text":["Goss & Field (1968)"],"Functions Text":["These authors compute the collisional rate coefficients for collisions with electrons, either using perturbation methods such as those used by"],"Functions Label":["Uses"],"Citation Start End":[[1012,1031]],"Functions Start End":[[869,1011]]}
{"Identifier":"2022MNRAS.516.5289M__Thompson_et_al._2015_Instance_3","Paragraph":"Given the number densities within the mass-dissociation index plane of Fig. 8, we now ask ourselves whether known dissociated clusters, such as the Bullet cluster, are expected in L210N1024NR? The Bullet Cluster has a mass of $\\sim 1.5 \\times 10^{15} \\, {\\rm M}_{\\odot }$ (e.g. Clowe et al. 2004; Brada\u010d et al. 2006; Clowe et al. 2006) and we estimated a dissociation index of SBullet \u223c 0.335 \u00b1 0.06. As seen in Fig. 8 there are no Bullet cluster analogues (structures of approximate mass and dissociation) in L210N1024NR, this is unsurprising as a simulation requires a significantly larger volume than that of L210N1024NR ((210cMpc\u2009h\u22121)3) to expect such an object (e.g. Lee & Komatsu 2010; Thompson & Nagamine 2012; Bouillot et al. 2015; Kraljic & Sarkar 2015; Thompson et al. 2015). From the distribution presented in Fig. 8, it is trivial to estimate the required cosmological volume (the effective volume, Veff) to expect structures of a given mass and dissociation index. By separating the 2D distribution on the mass-dissociation index planes into the component 1D distributions of mass and dissociation the effective volume is computed as\n(12)$$\\begin{eqnarray}\r\nV_\\text{eff}~^{-1} &amp;=&amp;\\int \\int \\,{\\rm{ d}} S \\, {\\rm{ d}} M \\phi (S, M) \\\\\r\n&amp;=&amp; \\int _{S_\\text{a}}^{S_\\text{b}} \\, {\\rm{ d}} S \\phi _S(S) \\int _{M_\\text{a}}^{M_\\text{b}} \\, {\\rm{ d}} M \\phi _M(M)~,\r\n\\end{eqnarray}$$where \u03d5S(S) is the number density function associated with S and $\\phi _\\mathit {M}(\\mathit {M})$ is the mass function presented in Fig. 7. Assuming a probable range of S = 0.335 \u00b1 0.06 and $1 \\lt M \\lt 2 \\times 10^{15} \\, {\\rm M}_{\\odot }$ we estimate a number density \u223c4.92 \u00d7 10\u221210 Mpc\u22123 or that an effective volume of \u223c2.03 Gpc3 would be required to observe a single Bullet-like cluster. This result is inline with the number density estimate of the order of \u223c10\u221210 Mpc\u22123 by Thompson et al. (2015), which improves on previous estimates (e.g. Lee & Komatsu 2010; Thompson & Nagamine 2012; Bouillot et al. 2015) due to more sophisticated halo finding methods (e.g. Behroozi, Wechsler & Wu 2013). Conversely, it was estimated by Kraljic & Sarkar (2015) (utilizing the same halo finder as Thompson et al. 2015) that given an effective volume of \u223c14.6 Gpc3, no Bullet cluster analogues are expected, however as indicated by a pairwise velocity distribution it would be expected that present binary halo\u2013halo orbits have the potential to form a Bullet-like object.","Citation Text":["Thompson et al. 2015"],"Functions Text":["utilizing the same halo finder as"],"Functions Label":["Similarities"],"Citation Start End":[[2191,2211]],"Functions Start End":[[2157,2190]]}
{"Identifier":"2022MNRAS.516.1539O__Fujita,_Ohira_&_Yamazaki_2013_Instance_1","Paragraph":"In this scenario, the present \u03b3-ray emission from the Fermi bubbles arises predominantly through inverse Compton scattering of an energetic non-thermal cosmic ray (CR) electron population in the remnant structures with ambient radiation supplied by the interstellar radiation field (ISRF) and the cosmological microwave background (CMB). A subdominant component due to non-thermal bremsstrahlung may also be present, emitted primarily from regions of high gas and non-thermal electron density within the bubble (OY22). Models of this nature are broadly referred to as \u2018leptonic\u2019 models (see also Su et al. 2010; Zubovas, King & Nayakshin 2011; Su & Finkbeiner 2012; Fujita, Ohira & Yamazaki 2013 for similar configurations). Other similar approaches invoking Sgr A* activity, but where the CR composition is not specifically required to be leptonic, invoke a pair of jet-driven outflowing bubbles assuming constant AGN activity or continuous energy injection over Myr time-scales (e.g. Zhang & Guo 2020). Notably, these models have been able to account for the bi-conical X-ray structures observed near the GC as part of the same phenomenon as the Fermi bubbles. Alternative proposals have also been discussed, where the bubbles arise from the confluence of a number of processes operating more gradually within the inner part of the Milky Way (Thoudam 2013). These could include tidal disruption events (TDEs) occurring at regular intervals of 10s to 100s kyr (Cheng et al. 2011; Ko et al. 2020), or the action of a bipolar galactic outflow driven by the ongoing intense GC star-formation activity and\/or the processes associated with Sgr A* (Lacki 2014), with the resulting \u03b3-ray glow instead arising from a hadronic CR population interacting with an advected supply of entrained gas in the wind (the \u2018hadronic\u2019 models \u2013 see  Crocker & Aharonian 2011; Cheng et al. 2014, 2015; Crocker et al. 2014, 2015; Mou et al. 2014, 2015; Razzaque & Yang 2018).","Citation Text":["Fujita, Ohira & Yamazaki 2013"],"Functions Text":["In this scenario, the present \u03b3-ray emission from the Fermi bubbles arises predominantly through inverse Compton scattering of an energetic non-thermal cosmic ray (CR) electron population in the remnant structures with ambient radiation supplied by the interstellar radiation field (ISRF) and the cosmological microwave background (CMB). A subdominant component due to non-thermal bremsstrahlung may also be present, emitted primarily from regions of high gas and non-thermal electron density within the bubble (OY22). Models of this nature are broadly referred to as \u2018leptonic\u2019 models (see also"],"Functions Label":["Background"],"Citation Start End":[[666,695]],"Functions Start End":[[0,595]]}
{"Identifier":"2017MNRAS.464L..26F__O'Sullivan_et_al._2001_Instance_2","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'Sullivan et al. 2001"],"Functions Text":["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"],"Functions Label":["Uses"],"Citation Start End":[[892,914]],"Functions Start End":[[655,891]]}
{"Identifier":"2022MNRAS.517.4327M__Indebetouw_et_al._2014_Instance_1","Paragraph":"Supernovae (SNe) play a dual role in the evolution of interstellar dust. On one hand, they are the most important source of dust production in galaxies, but on the other had had also the most important source of grain destruction. Theoretical models show that most of the heavy elements produced can precipitate out of the gas and form refractory grains (Sarangi & Cherchneff 2013, 2015; Sluder, Milosavljevi\u0107 & Montgomery 2018; Sarangi, Matsuura & Micelotta 2019). Infrared and submilimetre observations of Cassiopeia\u2009A (Barlow et al. 2010; Arendt et al. 2014; De Looze et al. 2017), SN\u20091987A (Matsuura et al. 2011; Indebetouw et al. 2014; Matsuura et al. 2015), Crab Nebula (Gomez et al. 2012), and young Galactic (up to \u223c2000 yr old) SN remnants (Chawner et al. 2019) confirm the presence of \u223c0.1\u20131.0 M\u2299 of dust, indicating that a substantial fraction of refractory elements in their ejecta went to dust grains. If the majority of dust in SNe can survive the shock interactions, SNe could be an important source of dust production in the ISM (Dwek & Cherchneff 2011). The fate of this newly-formed dust is still a subject of active studies. The reverse shock traveling through the ejecta can destroy newly-formed dust (Dwek, Foster & Vancura 1996; Schneider, Ferrara & Salvaterra 2004; Nozawa et al. 2007; Biscaro & Cherchneff 2014, 2016; Silvia, Smith & Shull 2010; Micelotta, Dwek & Slavin 2016; Kirchschlager et al. 2019). Any grains surviving the reverse shock may also be destroyed during the injection phase into the interstellar medium (ISM; Slavin et al. 2020) Thereafter, ISM dust will be subject to destruction as it encounters the SN remnant shocks. The grain destruction efficiency and ISM dust lifetimes are highly uncertain since they depend on a long list of parameters. Macroscopic parameters include the energy of the SN explosion, the morphology of the medium surrounding the SN (Slavin et al. 2020), and that of the general ISM. Microscopic parameters include the composition and size distribution of the SN condensates, and the detailed interaction of the dust with the shocked gas and other grains (Dwek & Arendt 1992; Jones, Tielens & Hollenbach 1996; Slavin, Dwek & Jones 2015; Kirchschlager, Mattsson & Gent 2021; Priestley, Chawner & of 2021). Because the evolution of dust in the ISM is a fine balance between dust production and destruction, intense investigations are currently underway into dust production and destruction by SNe. In this paper, we examine the latter point of view, and investigate how dust grains are impacted by SN shocks over time.","Citation Text":["Indebetouw et al. 2014"],"Functions Text":["Infrared and submilimetre observations of","SN\u20091987A","confirm the presence of \u223c0.1\u20131.0 M\u2299 of dust, indicating that a substantial fraction of refractory elements in their ejecta went to dust grains."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[617,639]],"Functions Start End":[[466,507],[585,593],[771,914]]}
{"Identifier":"2019MNRAS.482.4290H__Hofmann_2017_Instance_1","Paragraph":"The only non-standard term involves the factor $\\frac{\\mathrm{d} q}{\\mathrm{d} \\eta }$. As usual, the perturbation \u03a8 of the unperturbed phase-space distribution f0 = 1\/(exp\u2009(\u03b5\/T0) \u2212 1) is introduced as\n(36)\r\n\\begin{eqnarray*}\r\nf = f_0(\\epsilon) \\left(1 + \\Psi \\right)\\, ,\r\n\\end{eqnarray*}\r\nwhere the co-moving energy \u03b5 reads\n(37)\r\n\\begin{eqnarray*}\r\n\\epsilon \\equiv \\left(q^2 + a^2 \\frac{m_{V_\\pm }^2}{{\\cal S}^2\\, T_0^2}\\right)^{1\/2}\\, .\r\n\\end{eqnarray*}\r\nThe scaling function $\\mathcal {S}$ is defined in equation (5). Since the term $\\frac{\\mathrm{d} q}{\\mathrm{d} \\eta }$ is determined by the geodesic equation for a massive point particle, we may write\n(38)\r\n\\begin{eqnarray*}\r\n\\frac{\\mathrm{d} q}{\\mathrm{d} \\eta } \\frac{\\partial f}{\\partial q} &amp;=&amp; \\left(q \\dot{\\phi } - \\epsilon n_i \\partial _i \\psi - \\frac{a^2 m_{V_\\pm } \\dot{m}_{V_\\pm }}{q} \\right)\\nonumber\\\\\r\n&amp;&amp; \\times \\,\\left(\\frac{\\partial f_0}{\\partial q} (1 + \\Psi) + f_0 \\frac{\\partial \\Psi }{\\partial q}\\right).\r\n\\end{eqnarray*}\r\nThe use of the geodesic equation for a quasi-particle must be questioned, if this particle associates with pure quantum fluctuations (Hofmann 2017). If at all, temperature fluctuations in the V\u00b1 sector can thus only be coherently propagated via the low-frequency regime of \u03b3 fluctuations in terms of classical electromagnetic waves (Hofmann 2016b). To do this, a residual interaction between V\u00b1 and \u03b3 is required. Albeit when such an interaction occurs (Hofmann 2016a), its efficiency in conveying the coherent propagation of V\u00b1 temperature fluctuations must be questioned, especially at high temperatures (Falquez et al. (2010). To ignore the V\u00b1 Boltzmann equations and associated source terms in linearized Einstein equations thus is a physically motivated option. On the other hand, considering the evolution of V\u00b1 temperature fluctuations via the coherently propagating low-frequency sector in \u03b3 implies that \u03b5 = q in the V\u00b1 geodesic equation. At the same time, $m_{V_\\pm }\\gt 0$ is required in f0. Since the structure of temperature fluctuations is mainly imprinted before and during recombination, setting $m_{V_\\pm }=0$ in the geodesic equation does not influence the prediction for the power spectra in practice. Note that due to equation (37) an explicit dependence of f0 on \u03b7 needs to be considered via a = a(\u03b7). Transforming equation (35) into k space and otherwise following the standard procedure of linear perturbation theory (Ma & Bertschinger 1995), one arrives at the following hierarchy:\n(39)\r\n\\begin{eqnarray*}\r\n\\dot{\\Psi }_0 = - k\\Psi _1 - \\frac{\\mathrm{d} \\ln f_0}{\\mathrm{d} \\ln q}\\dot{\\phi } - \\frac{1+\\Psi _0}{f_0} \\frac{\\partial f_0}{\\partial \\eta }\\, ,\r\n\\end{eqnarray*}\r\n(40)\r\n\\begin{eqnarray*}\r\n\\dot{\\Psi }_1 = \\frac{k}{3 }(\\Psi _0 - 2 \\Psi _2) - \\frac{1}{3} \\frac{\\mathrm{d} \\ln f_0}{\\mathrm{d} \\ln q} k \\psi - \\frac{\\Psi _1}{f_0} \\frac{\\partial f_0}{\\partial \\eta },\r\n\\end{eqnarray*}\r\n(41)\r\n\\begin{eqnarray*}\r\n\\dot{\\Psi _l} = \\frac{k}{(2l+1)}\\left[l \\Psi _{l-1}-(l+1)\\Psi _{l+1} \\right] -\\frac{\\Psi _l}{f_0} \\frac{\\partial f_0}{\\partial \\eta },\r\n\\end{eqnarray*}\r\nwhere the $\\Psi _l(\\vec{k},q,\\eta)$ are the expansion coefficients for $\\Psi (\\vec{k},\\hat{n},q,\\eta)$ into Legendre polynomials.","Citation Text":["Hofmann 2017"],"Functions Text":["The use of the geodesic equation for a quasi-particle must be questioned, if this particle associates with pure quantum fluctuations"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1148,1160]],"Functions Start End":[[1014,1146]]}
{"Identifier":"2022AandA...663A..50B__Buat_et_al._2005_Instance_1","Paragraph":"In the absence of dust, the spectral emission of a normal star-forming galaxy is dominated by stellar populations of different ages with superimposed nebular emission, mainly in the form of recombination lines as well as continuum. The interaction with dust has a dramatic effect, both dimming and reddening the emission from stars and ionized gas. This negatively impacts our ability to measure star formation as energetic photons produced by massive young stars are far more easily attenuated than longer wavelength photons, and even a small quantity of dust can lead to a significant attenuation in the ultraviolet (UV). In the case of particularly dust-rich galaxies, it can render their detection in the rest-frame UV especially difficult. However, as the FUV emission vanishes due to dust attenuation, this dust re-emits the absorbed energy in the mid-infrared (MIR) and far-infrared (FIR), which can in turn be exploited to trace star formation. Except for the most extreme cases (e.g., when the dust content is negligible or, conversely, when almost all of the UV photons are absorbed by dust), an attenuation correction must be carried out to retrieve the star formation rate (SFR). One of the most direct ways is to simply apply a hybrid SFR estimator combining the rest-frame UV with the IR (e.g., Hao et al. 2011; Boquien et al. 2016). The obvious downside is that this requires observations of the dust emission that are costly and difficult to obtain, and even more so at increasing redshifts, where they tend to be limited to vanishingly small samples. With the rest-frame UV emission being relatively easy to obtain from the ground from z\u2004\u223c\u20042 and beyond, techniques have been developed to relate the UV slope (\u03b2) to the UV attenuation (the IRX-\u03b2 relation). While this approach initially appeared to work remarkably well in the case of starburst galaxies (Meurer et al. 1999), there is now ample evidence that there is no tight universal relation between the UV slope and the attenuation (e.g., Buat et al. 2005; Seibert et al. 2005; Howell et al. 2010; Casey et al. 2014). In fact, this relation relies on two strong underlying assumptions: the intrinsic UV slope of the stellar populations in the absence of dust and the exact shape of the attenuation curve. Numerous studies have analyzed their respective impact in an attempt to understand why and when such relations fail and build more reliable ones (e.g., Kong et al. 2004; Boquien et al. 2009, 2012; Popping et al. 2017, and many others). In particular, the recent study of Salim & Boquien (2019) found that the diversity of attenuation curves is a strong driver of the scatter around the IRX-\u03b2 relation. This finding, which is consistent with simulations (Narayanan et al. 2018b; Liang et al. 2021), is especially important in that we can observe a broad variety of attenuation curves at all redshifts (e.g., Salmon et al. 2016; Buat et al. 2018; Salim et al. 2018). With the shape of the attenuation curve being strongly dependent on the relative geometry of stars, ionized gas, and dust (Salim & Narayanan 2020), from the disturbed morphologies observed at higher redshifts, we can only expect important variations there as well (e.g., Faisst et al. 2017). However, due to the great difficulty in measuring them and given the sparsity of the data available, our knowledge of attenuation curves beyond z\u2004=\u20044 remains limited. In effect, most observational studies on the attenuation properties of distant galaxies tend to concentrate on redshifts between 2 and 4 (e.g., Noll et al. 2009b; Buat et al. 2012, 2019; Reddy et al. 2012, 2015; Shivaei et al. 2015; \u00c1lvarez-M\u00e1rquez et al. 2016; Salmon et al. 2016; Fudamoto et al. 2017, 2020b; Lo Faro et al. 2017; \u00c1lvarez-M\u00e1rquez et al. 2019; Reddy et al. 2018; Koprowski et al. 2020). There is only a handful of examples at higher redshift (Capak et al. 2015; Scoville et al. 2015; Bouwens et al. 2016; Barisic et al. 2017; Koprowski et al. 2018). Because of the inherent limits of the observations, studies based on numerical simulations of galaxies at very high redshift (e.g., Mancini et al. 2016; Cullen et al. 2017; Di Mascia et al. 2021) are an important source of information. However, they lead to contrasted results, finding both flat (Cullen et al. 2017) and steep (Mancini et al. 2016) attenuation curves.","Citation Text":["Buat et al. 2005"],"Functions Text":["While this approach initially appeared to work remarkably well in the case of starburst galaxies","there is now ample evidence that there is no tight universal relation between the UV slope and the attenuation (e.g.,"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[2010,2026]],"Functions Start End":[[1773,1869],[1892,2009]]}
{"Identifier":"2019AandA...622A..62A__Aviles_et_al._(2018)_Instance_2","Paragraph":"On the other hand, PT has experienced many developments in recent years (Matsubara 2008a; Baumann et al. 2012; Carlson et al. 2013), in part because it can be useful to analytically understand different effects in the power spectrum and correlation function for the dark matter clustering. These effects can be confirmed or refuted, and further explored with simulations to ultimately understand the outcomes of present and future galaxy surveys, such as eBOSS (Zhao 2016), DESI (Aghamousa et al. 2016), EUCLID (Amendola et al. 2013), and LSST (LSST Dark Energy Science Collaboration 2012). Two approaches have been used to study PT: the Eulerian standard PT (SPT) and Lagrangian PT (LPT), which both have advantages and drawbacks, but they are complementary in the end (Tassev 2014). The nonlinear PT for MG was developed initially in (Koyama et al. 2009), and has been further studied in several other works (Taruya et al. 2014a,b; Brax & Valageas 2013; Bellini & Zumalacarregui 2015; Taruya 2016; Bose & Koyama 2016, 2017; Barrow & Mota 2003; Akrami et al. 2013; Fasiello & Vlah 2017; Aviles & Cervantes-Cota 2017; Hirano et al. 2018; Bose et al. 2018; Bose & Taruya 2018; Aviles et al. 2018). The LPT for dark matter fluctuations in MG was developed in Aviles & Cervantes-Cota (2017), and further studies for biased tracers in Aviles et al. (2018). The PT for MG has the advantage that it allows us to understand the role that these physical parameters play in the screening features of dark matter statistics. We here study some of these effects through screening mechanisms by examining them at second- and third-order perturbation levels using PT for some MG models. To this end, we build on the LPT formalism developed in Aviles & Cervantes-Cota (2017), which was initially posited for MG theories in the Jordan frame, in order to apply it to theories in the Einstein frame. Because of the direct coupling of the scalar field and the dark matter in the Klein\u2013Gordon equation, the equations that govern the screening can differ substantially from those in Jordan-frame MG theories. In general, screening effects depend on the type of nonlinearities that are introduced in the Lagrangian density. We present a detailed analysis of screening features and identify the theoretical roots of their origin. Our results show that screenings possess peculiar features that depend on the scalar field effective mass and couplings, and that may in particular cases cause anti-screening effects in the power spectrum, such as in the symmetron model. We perform this analysis by separating the growth functions into screening and non-screened parts. We note, however, that we do not compare the perturbative approach with a fully nonlinear simulation. We refer to (Koyama et al. 2009), for instance, for investigations like this at the level of the power spectrum.","Citation Text":["Aviles et al. (2018)"],"Functions Text":["The LPT for dark matter fluctuations in MG was developed in Aviles & Cervantes-Cota (2017), and further studies for biased tracers in"],"Functions Label":["Extends"],"Citation Start End":[[1331,1351]],"Functions Start End":[[1197,1330]]}
{"Identifier":"2016ApJ...833...51Y__Liu_et_al._2012_Instance_1","Paragraph":"The past two decades have seen rapid progress in the field of solar magnetoseismology (SMS; for recent reviews, see e.g., Nakariakov & Verwichte 2005; Banerjee et al. 2007; De Moortel & Nakariakov 2012; Nakariakov et al. 2016; Wang 2016). Among the rich variety of low-frequency waves observed in the Sun\u2019s atmosphere, flare-related quasi-periodic fast propagating wave trains have received much attention (see Liu & Ofman 2014, for a recent review). Their quasi-periods usually range from 25 to 400 s. These wave trains were discovered (Liu et al. 2011) and extensively observed in images acquired with the Atmospheric Imaging Assembly on board the Solar Dynamics Observatory (SDO\/AIA; Liu et al. 2012; Shen et al. 2013; Yuan et al. 2013; Nistic\u00f2 et al. 2014; also see Lemen et al. 2012 for a description of the instrument). On the other hand, quasi-periodic signals in coronal emissions presumably from density-enhanced loops have been known since the 1960s (e.g., Frost 1969; Parks & Winckler 1969; Rosenberg 1970; McLean & Sheridan 1973; see Table 1 of Aschwanden et al. 1999 for a comprehensive compilation of measurements prior to 2000). The quasi-periods P of a considerable fraction of these signals were of the order of seconds. While these measurements were largely spatially unresolved, more recent high-cadence instruments imaging the corona at total eclipses indicated the presence in coronal loops of quasi-periodic signals both with \n\n\n\n\n\n s (Williams et al. 2001, 2002; Katsiyannis et al. 2003) and with \n\n\n\n\n\n s (Samanta et al. 2016). In addition, quasi-periodic pulsations in the lightcurves of solar flares with similar periods have also been measured using imaging instruments such as the Nobeyama Radioheliograph (e.g., Asai et al. 2001; Nakariakov et al. 2003; Melnikov et al. 2005; Kupriyanova et al. 2013), SDO\/AIA (e.g., Su et al. 2012), and more recently with the Interface Region Imaging Spectrograph (IRIS; Tian et al. 2016; see also De Pontieu et al. 2014 for a description of IRIS).","Citation Text":["Liu et al. 2012"],"Functions Text":["Among the rich variety of low-frequency waves observed in the Sun\u2019s atmosphere, flare-related quasi-periodic fast propagating wave trains have received much attention","Their quasi-periods usually range from 25 to 400 s. These wave trains were","and extensively observed in images acquired with the Atmospheric Imaging Assembly on board the Solar Dynamics Observatory (SDO\/AIA;"],"Functions Label":["Background","Background","Background"],"Citation Start End":[[687,702]],"Functions Start End":[[239,405],[451,525],[555,686]]}
{"Identifier":"2018AandA...619A..13V__Saviane_et_al._2012_Instance_5","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"],"Functions Text":["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"],"Functions Label":["Uses"],"Citation Start End":[[2364,2383]],"Functions Start End":[[2182,2363]]}
{"Identifier":"2020MNRAS.498..464F__Alam_et_al._2017_Instance_1","Paragraph":"The late-time matter density PDF at a given smoothing scale is mostly sensitive to the skewness of the primordial density field at that scale and to the running of that skewness around the smoothing scale. As such \u2013 unless the PDF is measured on a wide range of smoothing scales \u2013 it can only poorly distinguish between different primordial bispectrum shapes. Any model that produces mainly one of the possible bispectrum template can however be successfully tested with PDF measurements. In this paper, we consider an analysis of the PDF at redshift z = 1 in a survey volume of V = 100(Gpc\u2009h\u22121)3, which is smaller than the effective volume of upcoming surveys such as Spherex with Veff \u2248 150(Gpc\u2009h\u22121)3 and somewhat larger than existing surveys such as BOSS with Veff \u2248 55(Gpc\u2009h\u22121)3 (Dor\u00e9 et al. 2014; Alam et al. 2017). At a smoothing scale of 30\u2009Mpc\u2009h\u22121 we find our PDF model to agree with the high-resolution run of the Quijote N-body simulations (Villaescusa-Navarro et al. 2019) to $\\lesssim 0.2{{\\ \\rm per\\ cent}}$ accuracy over a range of \u03b4[30\u2009Mpc\u2009h\u22121] \u2208 [\u22120.3, 0.4]. This is within cosmic variance of the considered volume of 100(Gpc\u2009h\u22121)3 (which is also the combined volume of the Quijote high-resolution boxes). Restricting to this smoothing scale and to these mildly non-linear densities we find that a PDF based analysis can measure the amplitude of different primordial bispectrum shapes to an accuracy of $\\Delta f_{\\mathrm{NL}}^{\\mathrm{loc}} = \\pm 7.4\\ ,\\ \\Delta f_{\\mathrm{NL}}^{\\mathrm{equi}} = \\pm 22.0\\ ,\\ \\Delta f_{\\mathrm{NL}}^{\\mathrm{ortho}} = \\pm 46.0$ \u2013 even when marginalizing over the non-linear variance of the density field as a free parameter. When pushing to smaller scales and assuming a joint analysis of the PDF with smoothing radii of 30\u2009 and 15\u2009Mpc\u2009h\u22121 (\u03b4[15\u2009Mpc\u2009h\u22121] \u2208 [\u22120.4, 0.5]) this improves to $\\Delta f_{\\mathrm{NL}}^{\\mathrm{loc}} = \\pm 3.3\\ ,\\ \\Delta f_{\\mathrm{NL}}^{\\mathrm{equi}} = \\pm 11.0\\ ,\\ \\Delta f_{\\mathrm{NL}}^{\\mathrm{ortho}} = \\pm 17.0$ \u2013 even when marginalizing over the non-linear variances at both scales as two free parameters. Especially, such an analysis can simultaneously measure fNL and the amplitude and slope of the non-linear power spectrum. Note that any dependence of these forecasts on \u03c38 is completely mitigated by this marginalization. We do not consider the impact of \u03a9m on our signals (see Uhlemann et al. 2019, for an investigation of the general cosmology dependence of the PDF) though Friedrich et al. (2018) and Gruen et al. (2018) have demonstrated that parameters of the \u039bCDM model and higher order moments of the density field can be measured simultaneously from what they call lensing-around-cells. Ultimately, we are working towards a combination of a late-time PDF analysis with the early-universe results of Planck Collaboration IX (2019). These two analyses have the potential to complement each other: the CMB providing information about the background \u039bCDM space\u2013time, the late-time density PDF providing information about non-linear structure growth and both of them containing independent information about the imprint of primordial non-Gaussianities on the large-scale structure.","Citation Text":["Alam et al. 2017"],"Functions Text":["In this paper, we consider an analysis of the PDF at redshift z = 1 in a survey volume of V = 100(Gpc\u2009h\u22121)3, which is smaller than the effective volume of upcoming surveys such as Spherex with Veff \u2248 150(Gpc\u2009h\u22121)3 and somewhat larger than existing surveys such as BOSS with Veff \u2248 55(Gpc\u2009h\u22121)3"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[802,818]],"Functions Start End":[[489,782]]}
{"Identifier":"2019ApJ...871..176X__Eldridge_et_al._2013_Instance_1","Paragraph":"The progenitors of SNe Ib\/c have been thought to be Wolf-Rayet (W-R) stars with high initial masses (MZAMS \u2273 25 M\u2299; Crowther 2007). Before core collapse, these stars usually have experienced severe mass loss through strong stellar winds or due to interaction with companion stars (van der Hucht 2006; Paxton et al. 2015). As the evolution of massive stars is usually dominated by binary evolution (Heger et al. 2003) and also depends largely on metallicity, rotation, and so on (Heger et al. 2003; Georgy et al. 2013, 2012), this makes the direct identification of their progenitors complicated (Smartt 2015). However, there are increasing studies suggesting that a lower-mass binary scenario is more favorable for most SNe Ib\/c, considering the measured low ejecta masses (Eldridge et al. 2013; Lyman et al. 2016). In addition, the H\/He envelopes of the progenitor stars are stripped by binary interaction. There are many detections of progenitor stars for SNe II. For example, most SNe IIP are found to originate from red supergiants (Smartt et al. 2009), while SNe IIL are typically from progenitors with somewhat warmer colors (see Smartt 2015, for a review), and SNe IIb are from those with higher effective temperatures such as yellow supergiants that have had their H\/He envelopes partially stripped through binary interaction (e.g., SN 1993J; Podsiadlowski et al. 1993; Maund et al. 2004; Fox et al. 2014). Until recently, there has been only one report of the possible identification of a progenitor star for SNe Ib, namely, iPTF 13bvn, which was proposed to spatially coincide with a single W-R-like star identified on the pre-explosion Hubble Space Telescope (HST) images (Cao et al. 2013; Groh et al. 2013). But such an identification is still controversial (e.g., Bersten et al. 2014; Fremling et al. 2014; Eldridge et al. 2015; Eldridge & Maund 2016). Direct detection of progenitor stars is still elusive for SNe Ic, which prevents us from further testing the theoretical evolution of massive stars (Eldridge et al. 2013).","Citation Text":["Eldridge et al. 2013"],"Functions Text":["However, there are increasing studies suggesting that a lower-mass binary scenario is more favorable for most SNe Ib\/c, considering the measured low ejecta masses"],"Functions Label":["Background"],"Citation Start End":[[774,794]],"Functions Start End":[[610,772]]}
{"Identifier":"2022MNRAS.513..232N__Hayden_et_al._2015_Instance_1","Paragraph":"There are a plethora of data available in the form of spectra, astrometric, and photometric information, as well as multiwavelength maps with the advent of large-scale spectroscopic (Apache Point Observatory Galactic Evolution Experiment\/APOGEE: Eisenstein et al. 2011, RAdial Velocity Experiment\/RAVE: Steinmetz et al. 2006, Gaia-ESO: Gilmore et al. 2012, Large Sky Area Multi-Object Fiber Spectroscopic Telescope\/LAMOST: Cui et al. 2012, Galactic Archaeology with HERMES\/GALAH: De Silva et al. 2015, Abundances and Radial velocity Galactic Origins Survey\/ARGOS: Ness et al. 2012), astrometric (Hipparcos: Perryman et al. 1997, Gaia: Gaia Collaboration 2016), and photometric surveys (Two-Micron All Sky Survey\/2MASS: Skrutskie et al. 2006, Sloan Digital Sky Survey\/SDSS: Stoughton et al. 2002, Vista Variables in the V\u00eda L\u00e1ctea\/VVV: Minniti et al. 2010, the SkyMapper Southern Survey : Wolf et al. 2018). These surveys have enabled the chemo-dynamic characterization of stellar populations in the Milky way that constitute different Milky Way components like thin disc, thick disc, halo, bulge, etc. For example, star count observations in the solar neighbourhood (Yoshii 1982; Gilmore & Reid 1983) led to the discovery of the thick disc, followed by its characterization as the old \u03b1-enhanced population in the double sequence exhibited by the solar neighbourhood stars in the [\u03b1\/Fe] versus [Fe\/H] plane (Fuhrmann 1998; Bensby, Feltzing & Lundstr\u00f6m 2003; Reddy, Lambert & Allende Prieto 2006; Adibekyan et al. 2012; Haywood et al. 2013). At present, data from large-scale spectroscopic surveys (Anders et al. 2014; Hayden et al. 2015; Weinberg et al. 2019) have led to the discovery of this trend at different galactocentric radius, R, and average height, |Z|, across the Galaxy shedding light on the disc formation and evolution scenarios. In addition, many age determination methods have been developed that uses these survey data to provide valuable information about the star formation histories and age metallicity relation of disc stellar populations (Casagrande et al. 2011; Bedell et al. 2018; Lin et al. 2020; Nissen et al. 2020). Secular processes such as radial migration (Sellwood & Binney 2002; Sch\u00f6nrich & Binney 2009; Minchev & Famaey 2010), which leads to the mixing of stars across the Galaxy, are also being explored using a combination of accurate phase space information from Gaia (Gaia Collaboration 2018) and chemistry and age information of stars from large-scale spectroscopic surveys (Buder et al. 2019). The discovery of streams and dynamically different stellar populations in the Milky Way halo, considered to be the result of past accretion\/merger events (Belokurov et al. 2018; Helmi et al. 2018; Ibata, Malhan & Martin 2019; Myeong et al. 2019) using the Gaia data and their further exploration with chemistry from large-scale spectroscopic surveys (Buder et al. in preparation) is another example. Multiple components in the Bulge metallicity distribution function discovered by multiple individual and large-scale spectroscopic observations, are being studied in detail to understand the origin of the Bulge and its connection with the Milky Way bar and Galaxy evolution (Ness et al. 2013; Rojas-Arriagada et al. 2017, 2020). There are many upcoming surveys [4-metre Multi-Object Spectroscopic Telescope\/4MOST: de Jong et al. (2019), Sloan Digital Sky Survey\/SDSS-V: Kollmeier et al. (2017), WEAVE: Dalton et al. (2018)] that will further improve our understanding of the formation and evolution of the Milky Way and its components.","Citation Text":["Hayden et al. 2015"],"Functions Text":["At present, data from large-scale spectroscopic surveys","have led to the discovery of this trend at different galactocentric radius, R, and average height, |Z|, across the Galaxy shedding light on the disc formation and evolution scenarios."],"Functions Label":["Background","Background"],"Citation Start End":[[1618,1636]],"Functions Start End":[[1541,1596],[1660,1843]]}
{"Identifier":"2021AandA...651A.111P__Herrera-Camus_et_al._2018_Instance_2","Paragraph":"Irrespective of its origin, the [C\u202fII] emission is linked to the presence of stellar far-ultraviolet (FUV) photons (E 13.6 eV). As FUV photons are tied to the presence of massive O and B stars that have short lifetimes, the [C\u202fII] 158 \u03bcm line is also astar formation rate (SFR) indicator. Indeed, ISO, Herschel and SOFIA observations have demonstrated the good correlation between the [C\u202fII] luminosity and the SFR in the Milky Way and in regions of massive star formation in other galaxies (e.g., Kramer et al. 2013, 2020; Pineda et al. 2014, 2018; Herrera-Camus et al. 2015, 2018; De Looze et al. 2011). With ALMA and NOEMA, ground-based observations of the [C\u202fII] 158 \u03bcm line in high redshift galaxies have come into reach and such data are routinely used to infer SFRs (e.g., Walter et al. 2012; Venemans et al. 2012; Knudsen et al. 2016; Bischetti et al. 2018; Khusanova et al. 2021) based upon validations of this relationship in the nearby Universe (Herrera-Camus et al. 2018; De Looze et al. 2011). However, it is well-understood that the intensity of the [C\u202fII] line depends on the local physical conditions (Hollenbach & Tielens 1999). Observationally, the presence of the so-called [C\u202fII]-deficit \u2013 a decreased ratio of [C\u202fII] 158 \u03bcm luminosity to FIR dust continuum with an increasing dust color temperature and also with FIR luminosity \u2013 is well established (Malhotra et al. 2001; D\u00edaz-Santos et al. 2013; Magdis et al. 2014; Smith et al. 2017). This deficit is particularly pronounced in (local) ultraluminous infrared galaxies (ULIRGs), very dusty galaxies characterized by vigorous embedded star formation (e.g., Luhman et al. 2003; Abel et al. 2009; Graci\u00e1-Carpio et al. 2011). This deficit, however, does not necessarily hold in the early Universe at high redshift (e.g., Stacey et al. 2010; Brisbin et al. 2015; Capak et al. 2015). Some studies have indicated that not only [C\u202fII] emission is deficient in some sources, but other FIR cooling lines ([O\u202fI], [O\u202fIII], [N\u202fII]), as well (e.g., Graci\u00e1-Carpio et al. 2011; Herrera-Camus et al. 2018). These deficits must be linked to the global ISM properties and star-formation characteristics in these galaxies.","Citation Text":["Herrera-Camus et al. 2018"],"Functions Text":["Some studies have indicated that not only [C\u202fII] emission is deficient in some sources, but other FIR cooling lines ([O\u202fI], [O\u202fIII], [N\u202fII]), as well (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[2035,2060]],"Functions Start End":[[1851,2007]]}
{"Identifier":"2020ApJ...903L..12H__Petschek_1964_Instance_1","Paragraph":"Magnetic reconnection (MR) may occur in various space and astrophysical plasma environments, among which the planetary magnetopause boundaries separating the solar wind and magnetospheric origins of plasmas and magnetic field are some of the most likely sites for the occurrence of MR. Due to the easy access to the in situ spacecraft observations the Earth\u2019s magnetopause is the most widely studied space plasma environment for MR (Paschmann et al. 1979; Vaivads et al. 2004; Graham et al. 2014). In particular, the Magnetospheric Multiscale (MMS) mission has contributed greatly to the kinetic physics of magnetopause reconnection (Burch et al. 2016; Hasegawa et al. 2017; Zhong et al. 2020). Many studies have shown that an initial Harris type equilibrium profile with constant total pressures \n\n\n\n\n\n and antiparallel magnetic field with or without a guide field (Harris 1962) may tend to develop MR geometry. In particular, two major categories of MR have been proposed: the steady state model with a single X line and the outflow approaching the Alfv\u00e9n speed (Petschek 1964), and the tearing mode instability with a series of X and O lines and mild plasma velocity (Furth et al. 1963). Numerous fluid and kinetic simulations have been carried out to examine the various aspects of MR processes for the past 50 yr (Hau & Chiou 2001; Guo et al. 2015; Landi et al. 2015). In particular, the effects of pressure or temperature anisotropy on MR have been examined by a number of authors (Chen & Palmadesso 1984; Shi et al. 1987; Birn & Hesse 2001; Chiou & Hau 2002, 2003; Hung et al. 2011). In the MHD models the double-polytropic (DP) laws are widely adopted as the energy closures to study the effects of temperature anisotropy and energy closures on MR and tearing mode instability (Chiou & Hau 2002, 2003; Hung et al. 2011). It is shown that the mirror type temperature anisotropy of \n\n\n\n\n\n may greatly enhance the growth rate of tearing mode instability and the merging rate of single X-line reconnection. In particular, the coupling of tearing and mirror instabilities may lead to relatively larger magnetic islands as compared to the cases with isotropic pressure and the mirror waves with anticorrelated density and magnetic field may be present in the vicinity of X lines.","Citation Text":["Petschek 1964"],"Functions Text":["In particular, two major categories of MR have been proposed: the steady state model with a single X line and the outflow approaching the Alfv\u00e9n speed"],"Functions Label":["Background"],"Citation Start End":[[1065,1078]],"Functions Start End":[[913,1063]]}
{"Identifier":"2018AandA...616A.173K__Pickett_1991_Instance_1","Paragraph":"The room-temperature millimeter wave rotational spectrum of ethyl isocyanate in Fig. 1 presents an exceptionally high line density. Its analysis was started using predictions from the spectroscopic constants obtained in the previous microwave works (Heineking et al. 1994; Kasten et al. 1983; Sakaizumi et al. 1976) for the stable cis configuration. It is characterized by a relatively large dipole moment along the a inertial axis (|\u00b5a| = 2.81 D and |\u00b5b| = 0.03 D (Sakaizumi et al. 1976), Cs symmetry, and dihedral angle \u03c4(C\u2013C\u2013N = C) = 0\u25e6 (see Fig. 1). The identification of Ka = 0 and lower-frequency Ka = 1 transitions, originating from J0 J and J1 J energy levels, was relatively straightforward and their assignment and analysis could be easily expanded up to 340 GHz (J\u02dd = 64) with the help of the Loomis-Wood-type plot technique from the AABS package (Kisiel et al. 2005) and SP-FIT\/SPCAT program suite (Pickett 1991). The upper-frequency Ka = 1 transitions (J1 J\u22121 levels), which became degenerate with Ka = 2 transitions (J2 J\u22121 levels) for J\u02dd > 45, were subsequently localized, however, their analysis quickly ran into problems. As shown in Fig. 2, significant departures from the predicted positions were observed at high J\u02dd even though it was clear from the Loomis-Wood-type plots that the assignments were correct. Similar situations also occurred for higher Ka transitions and with increasing value of Ka, the limit of J up to which the rotational transitions were amenable to the simple semirigid-rotor Hamiltonian analysis, was decreasing (see Fig. 2). Only those transitions that supported the semirigid-rotor treatment were retained in the analysis. These transitions were finally combined with hyperfine-free transitions (5\u201324 GHz, J\u02dd = 0\u22127, K\u02dda = 0\u22123) from Heineking et al. (1994) and globally fitted using Watson\u2019s S -reduced Hamiltonian in Ir representation. The adjusted rotational and centrifugal distortion constants are given in Table 1. The list of experimental frequencies is provided in the Table 4. Finally, the deviation trends observed in the Loomis\u2013Wood type plots, such as those in Fig. 2, were advantageously followed to measure the frequencies of more than 200 transitions that could not be encompassed in the semirigid rotor fit. These transitions are collected in the Table 5.","Citation Text":["Pickett 1991"],"Functions Text":["The identification of Ka = 0 and lower-frequency Ka = 1 transitions, originating from J0 J and J1 J energy levels, was relatively straightforward and their assignment and analysis could be easily expanded up to 340 GHz (J\u02dd = 64) with the help of the Loomis-Wood-type plot technique from the AABS package","and SP-FIT\/SPCAT program suite"],"Functions Label":["Future Work","Future Work"],"Citation Start End":[[911,923]],"Functions Start End":[[554,857],[879,909]]}
{"Identifier":"2021MNRAS.500.5009M__Bono_et_al._2003_Instance_2","Paragraph":"RR Lyrae are old low-mass stars that, during the central helium-burning phase, show mainly radial pulsation while crossing the classical instability strip in the colour\u2013magnitude diagram. From the observational point of view, they represent the most numerous class of pulsating stars in the Milky Way and, being associated with old stellar populations, are typically found in globular cluster and abundant in the Galactic halo and bulge. The investigation of RR Lyrae properties is motivated by their important role both as distance indicators and tracers of old stellar populations. In particular, evolving through the central helium-burning phase, they represent the low-mass, Population II counterparts of Classical Cepheids, as powerful standard candles and calibrators of secondary distance indicators. In particular, they can be safely adopted to infer distances to Galactic globular clusters (see e.g. Coppola et al. 2011; Braga et al. 2016, 2018, and references therein), the Galactic centre (see e.g. Contreras Ramos et al. 2018; Marconi & Minniti 2018; Griv, Gedalin & Jiang 2019), and Milky Way satellite galaxies (see e.g. Coppola et al. 2015; Mart\u00ednez-V\u00e1zquez et al. 2019; Vivas et al. 2019, and references therein). Being associated with old stellar populations, they represent the basis of an alternative Population II distance scale (see e.g. Beaton et al. 2016, to the traditionally adopted Classical Cepheids), more suitable to calibrate secondary distance indicators that are not specifically associated with spiral galaxies (e.g. the globular cluster luminosity function, see Di Criscienzo et al. 2006, and references therein). The properties that make RR Lyrae standard candles are (i) the well-known relation connecting the absolute visual magnitude MV to the metal abundance [Fe\/H] (see e.g. Sandage 1993; Caputo et al. 2000; Cacciari & Clementini 2003; Catelan, Pritzl & Smith 2004; Di Criscienzo, Marconi & Caputo 2004; Federici et al. 2012; Marconi 2012; Marconi et al. 2015, 2018; Muraveva et al. 2018, and references therein); (ii) the period\u2013luminosity relation in the near-infrared (NIR) filters and in particular in the K 2.2 \u03bcm band (see e.g. Longmore et al. 1990; Bono et al. 2003; Dall\u2019Ora et al. 2006; Coppola et al. 2011; Ripepi et al. 2012; Coppola et al. 2015; Marconi et al. 2015; Muraveva et al. 2015; Braga et al. 2018; Marconi et al. 2018, and references therein). In spite of the well-known advantage of using NIR filters (see e.g. Marconi 2012; Coppola et al. 2015, and references therein), in the last decades there has been a debate on the coefficient of the metallicity term of the KB and PL relation (see e.g. Bono et al. 2003; Sollima, Cacciari & Valenti 2006; Marconi et al. 2015, and references therein). On the other hand, it is interesting to note that many recent determinations (see e.g. Sesar et al. 2017; Muraveva et al. 2018) seem to converge towards the predicted coefficient by Marconi et al. (2015), with values in the range 0.16-0.18 mag\u2009dex\u22121. As for the optical bands, our recently developed theoretical scenario (Marconi et al. 2015) showed that, apart from the MV\u2212[Fe\/H] relation that is affected by a number of uncertainties (e.g. a possible non-linearity, the metallicity scale with the associated \u03b1 elements enhancement and helium abundance variations, as well as evolutionary effects, see Caputo et al. 2000; Marconi et al. 2018, for a discussion), the metal-dependent Period\u2013Wesenheit (PW) relations are predicted to be sound tools to infer individual distances. In particular, for the B-, V- band combination, Marconi et al. (2015) demonstrated that the inferred PW relation is independent of metallicity. In order to test this theoretical tool, we need to compare the predicted individual distances with independent reliable distance estimates, for example, the astrometric ones recently obtained by the Gaia satellite (Gaia Collaboration 2016). To this purpose, in this paper we transform the predicted light curves derived for RR Lyrae models with a wide range of chemical compositions (Marconi et al. 2015, 2018) into the Gaia bands, derive the first theoretical PW relations in these filters and apply them to Gaia Data Release 2 Data base (hereinafter Gaia DR2; Gaia Collaboration 2018; Clementini et al. 2019; Ripepi et al. 2019). The organization of the paper is detailed in the following. In Section 2, we summarize the adopted theoretical scenario, while in Section 3 we present the first theoretical light curves in the Gaia filters. From the inferred mean magnitudes and colours, the new theoretical PW relations are derived in Section 4 that also includes a discussion of the effects of variations in the input chemical abundances. In Section 5, the obtained relations are applied to Gaia Galactic RR Lyrae with available periods, parallaxes, and mean magnitudes to infer independent predictions on their individual parallaxes, to be compared with Gaia DR2 results. The conclusions close the paper.","Citation Text":["Bono et al. 2003"],"Functions Text":["in the last decades there has been a debate on the coefficient of the metallicity term of the KB and PL relation (see e.g."],"Functions Label":["Motivation"],"Citation Start End":[[2658,2674]],"Functions Start End":[[2535,2657]]}
{"Identifier":"2022MNRAS.511.2105K__McElroy_et_al._2015_Instance_1","Paragraph":"AGN feedback can exist in several forms such as radiation, thermal, or non-thermal (cosmic rays) pressure-driven winds, jet-mode feedback, and via magnetic forces on accretion disc scales. AGN feedback can explain several observed properties such as the presence of high velocity (>1000 km s\u22121) multiphase gas outflows in low and high redshift galaxies and observations of bubbles or cavities in X-ray observations of galaxy clusters (e.g. Blanton et al. 2011; Fabian 2012; Sanders et al. 2014; Feruglio et al. 2015; Laha et al. 2021). High velocity outflows from AGN host galaxies have been reported in numerous studies in the literature (see Veilleux et al. 2020 for a review and the references therein) using optical spectroscopy (e.g. Greene et al. 2011; McElroy et al. 2015; Sun, Greene & Zakamska 2017; Durr\u00e9 & Mould 2018; Manzano-King, Canalizo & Sales 2019; Perna et al. 2020; Santoro et al. 2020; Trindade Falc\u00e3o et al. 2021), near-infrared spectroscopy (e.g. Kakkad et al. 2016; Zakamska et al. 2016; Bischetti et al. 2017; Diniz et al. 2019; Riffel, Zakamska & Riffel 2020a; Riffel et al. 2020b) and sub-mm spectroscopy (e.g. Michiyama et al. 2018; Zschaechner et al. 2018a; Audibert et al. 2019; Impellizzeri et al. 2019; Garc\u00eda-Bernete et al. 2021). One of the key quantities that is not well understood through these observations is how efficiently does the outflow couple with the ISM (e.g. Harrison et al. 2018). The coupling efficiency i.e. the ratio between the kinetic power of the outflow ($\\dot{E}_{\\rm kin}$) and the bolometric luminosity of the AGN (Lbol) or the star formation rate (SFR) of the host galaxy is critical to quantify the true impact of AGN feedback on host galaxies \u2013 the higher the efficiency, the easier it is for these outflows to heat the gas or propagate the outflows to the galaxy outskirts. An accurate measurement of mass outflow rate and kinetic energy is therefore necessary to estimate the true coupling efficiency, which can also be used as constraints in cosmological simulations.","Citation Text":["McElroy et al. 2015"],"Functions Text":["High velocity outflows from AGN host galaxies have been reported in numerous studies in the literature","using optical spectroscopy"],"Functions Label":["Background","Background"],"Citation Start End":[[759,778]],"Functions Start End":[[536,638],[706,732]]}
{"Identifier":"2015ApJ...808...56M__Beaulieu_et_al._2011_Instance_3","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"],"Functions Text":["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.,"],"Functions Label":["Background"],"Citation Start End":[[2482,2502]],"Functions Start End":[[2193,2481]]}
{"Identifier":"2017AandA...608A..67B__Luu_(1991)_Instance_1","Paragraph":"The large Jovian irregular satellites have been observed and analyzed using their light curves, colours, and reflectance spectra but the reported measurements are sometimes contradictory. The multicolour observation of some retrograde and prograde Jovian irregular satellites by Tholen & Zellner (1984) suggested C-class surface composition for prograde and more diverse colours for the retrograde families with a mixture of C- and D-type spectra. They noted that Carme had a flat visible wavelength reflectance spectrum, but with a strong upturn in the ultraviolet. Tholen & Zellner (1984) suggested that Carme might be showing low-level cometary activity with CN emission at 0.388 \u03bcm. Luu (1991) identified C- and D-type asteroid spectral features for both prograde and retrograde families based on spectroscopic observations of JV-JXIII and suggested them to be similar to Jupiter\u2019s Trojan asteroids. Based on 1.3\u20132.4 \u03bcm near-IR (NIR) observations, Brown (2000) reported their compositions as being similar to P- and D-class asteroids from the outer asteroid belt, while their visible spectra resemble C-class asteroids. Cruikshank (1977) and Degewij et al. (1980) observed Himalia in the NIR and confirmed that its surface composition is similar to that of C-type asteroids. Brown (2000) concluded that NIR spectra of Himalia and Elara are featureless between 1.4 and 2.5 \u03bcm and do not contain any water-ice absorption features. Subsequently, Brown & Rhoden (2014) supported these findings and suggested that these objects lacked aqueously altered phyllosilicates based on the absence of a 3 \u03bcm absorption band. Brown et al. (2003) and Chamberlain & Brown (2004) studied Himalia using data acquired by the Visual and Infrared Mapping Spectrometer (VIMS) on-board Cassini spacecraft during Jupiter\u2019s fly-by and found that its spectrum (0.3\u20135.1 \u03bcm) has low reflectance, a slight red slope, and an apparent absorption near 3-\u03bcm suggesting the presence of water in some form. In addition, Jarvis et al. (2000) reported a weak absorption at 0.7 \u03bcm in Himalia\u2019s spectrum and attributed it to oxidized iron. Contrary to this result, Brown & Rhoden (2014) found no evidence for aqueously altered phyllosilicates in the 2.2\u20133.8 \u03bcm region. ","Citation Text":["Luu (1991)"],"Functions Text":["identified C- and D-type asteroid spectral features for both prograde and retrograde families based on spectroscopic observations of JV-JXIII and suggested them to be similar to Jupiter\u2019s Trojan asteroids."],"Functions Label":["Background"],"Citation Start End":[[687,697]],"Functions Start End":[[698,903]]}
{"Identifier":"2022AandA...661A..10B__Ghirardini_et_al._2021a_Instance_1","Paragraph":"It is also possible that these clusters have a smaller extent and can just be missed by our extent selection as our detection algorithm sets the extent to zero if it is smaller than 6 (Brunner et al. 2022). Following the method presented in Ghirardini et al. (2021a), we estimated several dynamical properties of the clusters in the point source sample and compared them with the extent-selected sample presented in Ghirardini et al. (2021a). In Fig. 5 we compare the distributions of the core radii (Rcore) constrained by the V06 model and the concentration parameter (cSB) between these two samples. The concentration parameter is defined as the ratio of the surface brightness within 0.1 R500 to the surface brightness within R500 (Ghirardini et al. 2021a; Santos et al. 2008; Maughan et al. 2012). Intuitively, the expectation is that the smaller the core radius, the more compact the cluster. The left panel of Fig. 5 clearly shows that the clusters in the point source sample have relatively smaller core radii, hence the emission is more concentrated in a smaller area. Consistently, the concentration of the point source sample shows a clear excess in higher values than the extent-selected sample, indicating that a significantly larger fraction of cool-core clusters and clusters host a central AGN. We performed the same experiment by applying cuts in flux 1.5 \u00d7 10\u221214 ergs s\u22121 cm\u22122 and in detection likelihood to test whether the clusters are missed by the extent selection because they are fainter and\/or more compact than the extent-selected clusters. The distribution of number density, core radius, and concentration parameter remains the same, indicating that the population of clusters in the point source catalog is more compact than the extent-selected sample. The extent-selected sample does not show a clear bias toward cool-core clusters or clusters with a central AGN, but contains the fraction of cool-cores is similar to that of SZ surveys (Ghirardini et al. 2021a). In this sample, we observe the opposite trend.","Citation Text":["Ghirardini et al. (2021a)"],"Functions Text":["Following the method presented in"],"Functions Label":["Uses"],"Citation Start End":[[241,266]],"Functions Start End":[[207,240]]}
{"Identifier":"2018AandA...615L..16F__Hily-Blant_et_al._2013b_Instance_1","Paragraph":"The observational and theoretical studies of nitrogen isotope fractionation in star-forming regions can help to constrain nitrogen chemistry. Nitrogen has two stable isotopes, 14N and 15N. The elemental abundance ratio [14N\/15N]elem in the local interstellar medium (ISM) has been estimated to be ~200\u2013300 from the absorption line observations of N-bearing molecules toward diffuse clouds (Lucas & Liszt 1998; Ritchey et al. 2015). L1544 is a prototypical prestellar core located in the Taurus molecular cloud complex. In L1544, the 14N\/15N ratio of several different molecules has been measured: \n\n\n\n\n\n\n14\n\n\n\n\nN\n\n\n2\n\n\n\n\nH\n\n\n+\n\n\n=\n\n\n920\n\n\n\u2212\n200\n\n\n+\n300\n\n\n,\n\n\n\n14\n\n\n\n\nN\n\n\n2\n\n\n\n\nH\n\n\n+\n\n\n\/\n\n\nNNH\n\n\n+\n\n\n=\n\n\n1000\n\n\n\u2212\n220\n\n\n+\n260\n\n\n\n\n$ {}^{14}{\\mathrm N}_2\\mathrm H^+=920_{-200}^{+300},^{14}{\\mathrm N}_2\\mathrm H^+\/\\mathrm{NNH}^+=1000_{-220}^{+260} $\n\n\n (Bizzocchi et al. 2013; Redaelli et al. 2018), NH2D\/15NH2D > 700 (G\u00e9rin et al. 2009), CN\/C15N = 500 \u00b1 75 (Hily-Blant et al. 2013b), and HCN\/HC15N = 257 (Hily-Blant et al. 2013a). Among the measurements, the significant depletion of 15N in N2H+ is the most challenging for the theory of 15N fractionation. In general, molecules formed at low temperatures are enriched in 15N through gas-phase chemistry triggered by isotope exchange reactions (e.g., Terzieva & Herbst 2000). A 15N-bearing molecule has a slightly lower zero-point energy than the corresponding 14N isotopolog. This results in endothermicity for the exchange of 15N for 14N, which inhibits this exchange at low temperature enabling the concentration of 15N in molecules. Astrochemical models for prestellar cores, which consider a set of nitrogen isotope exchange reactions, have indeed predicted that atomic N is depleted in 15N, while N2 (and thus N2H+) is enriched in 15N (e.g., Charnley & Rodgers 2002). The model prediction clearly contradicts the observation of the N2H+ isotopologs in L1544. The 15N depletion in N2H+ was recently found in other prestellar cores as well, such as L183, L429, and L694-2 (Redaelli et al. 2018). Furthermore, Roueff et al. (2015) recently pointed out the presence of activation barriers for some key nitrogen isotope exchange reactions, based on their quantum chemical calculations. Then 15N fractionation triggered by isotope exchange reactions may be much less efficient than previously thought (Roueff et al. 2015, but see also Wirstr\u00f6m & Charnley 2018).","Citation Text":["Hily-Blant et al. 2013b"],"Functions Text":["In L1544, the 14N\/15N ratio of several different molecules has been measured:","CN\/C15N = 500 \u00b1 75"],"Functions Label":["Uses","Uses"],"Citation Start End":[[955,978]],"Functions Start End":[[519,596],[935,953]]}
{"Identifier":"2022MNRAS.515L..39Z__Koefoed_et_al._2016_Instance_1","Paragraph":"Currently, this 53Mn\u201353Cr age of 4566.6 \u00b1 0.6\u2009Ma for crystallization for EC 002 represents the oldest record of volcanism in the Solar system. For example, the oldest crust formation of Earth and Moon only dates back to \u223c4.3\u20134.4\u2009Ga (O\u2019Neil & Carlson 2017; Borg et al. 2019), and Mars, Vesta, and the angrite and main-group aubrite parent bodies show ages of mantle\u2013crust differentiation at \u223c4547\u2009Ma (Bouvier et al. 2018), 4564.8 \u00b1 0.6\u2009Ma (Trinquier et al. 2008), 4563.2 \u00b1 0.2\u2009Ma (Zhu et al. 2019b), and 4562.5 \u00b1 1.1\u2009Ma (Zhu et al. 2020b), respectively. The crystallization age of EC 002 also predates all those of the other dated achondrites, such as angrites (Amelin 2008a, b; Connelly et al. 2008), ureilites (Goodrich et al. 2010; Bischoff et al. 2014), NWA 8704\/6693, NWA 11119 (Srinivasan et al. 2018), and NWA 7325 (Koefoed et al. 2016). The result strongly supports the notion that advanced silicate differentiation occurred and evolved planetary crust formation very early in the Solar system, i.e. within the first 1\u2009Ma after CAI formation (4567.3 \u00b1 0.1\u2009Ma; Amelin et al. 2010; Connelly et al. 2012). The crystallization of andesitic crust must post-date both accretion and core formation on the EC 002 parent body, which is also consistent with evidence for early core formation for some asteroids derived from some iron meteorites (Kruijer et al. 2014; Anand et al. 2021). The age for EC 002 is older than some of the chondrule formation ages (Connelly et al. 2012; Bollard et al. 2017; Zhu et al. 2020b). This observation supports previous suggestions that many chondrites and their components reflect younger nebular processes, post-dating the oldest differentiated planetesimals, such as the EC 002 parent body. Thus, chondrules may not necessarily reflect an important ingredient in the accretion history of terrestrial planets (Johansen et al. 2015), although this cannot be excluded for earlier chondrule precursors with older generations (Zhu et al. 2019a). Considering its very old age and the short half-life of 0.7\u2009Ma of 26Al, the heat source for melting of the EC 002 parent body must have been the decay of 26Al. The reason why EC 002 cooled and crystallized so early might have been that its parent body was of a much smaller size than the terrestrial planets, since small bodies cannot retain their heat well. The size of the EC 002 parent body may have been smaller than the size of asteroids like Vesta (with mean radius of 262.7\u2009km; Russell et al. 2012) and the angrite and aubrite (main-group) parent bodies, which differentiated later, at 2.5\u20135\u2009Ma after CAIs (Amelin 2008a; Trinquier et al. 2008; Zhu et al. 2019b, 2021b).","Citation Text":["Koefoed et al. 2016"],"Functions Text":["The crystallization age of EC 002 also predates all those of the other dated achondrites, such as","and NWA 7325"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[822,841]],"Functions Start End":[[553,650],[808,820]]}
{"Identifier":"2017MNRAS.464..968S__Tacconi_et_al._2006_Instance_2","Paragraph":"Comparison of apparent effective diameters of these sources to direct size measurements supports a similar conclusion. Simpson et al. (2015) present ALMA observations of 23 SCUBA-2-selected SMGs with a median physical half-light diameter of 2.4 \u00b1 0.2 kpc, while Ikarashi et al. (2015) show ALMA observations of 13 AzTEC-selected SMGs with a median physical half-light diameter of $1.34^{+0.26}_{-0.28}$ kpc. ALMA observations of four SPT-selected lensed SMGs give a mean physical half-light diameter of 2.14 kpc (Hezaveh et al. 2013b). This measurement is consistent with a recent lensing analysis of a significantly expanded SPT-selected DSFG sample (Spilker et al. 2016). These high-resolution ALMA observations constrain the FIR sizes of the sources to be 1.0\u20132.5 kpc. Earlier observations of the physical sizes of SMGs by CO detection and 1.4 GHz imaging suggest larger sizes (e.g. Tacconi et al. 2006; Biggs & Ivison 2008; Younger et al. 2008). However, Simpson et al. (2015) point out that the submillimetre sizes are consistent with resolved 12CO detections, while the sizes derived from 1.4 GHz imaging are about two times larger because of the cosmic ray diffusion, which can explain the results before higher frequency observations at ALMA were possible (Chapman et al. 2004; Tacconi et al. 2006; Biggs & Ivison 2008; Younger et al. 2008). Similarly, Ikarashi et al. (2015) reveal that the 12CO detected sizes and the 1.4 GHz imaging sizes of similar sources are greater than their submillimetre sizes as well. Furthermore, observations of local galaxies also show the submillimetre sizes are smaller than the CO detected sizes (e.g. Sakamoto et al. 2006, 2008; Wilson et al. 2008) and the 1.4 GHz continuum sizes (e.g. Elbaz et al. 2011). Our photometrically derived $\\sqrt{\\mu }d$ is best compared to the submillimetre continuum sizes. With a median apparent effective diameter of $4.2^{+1.7}_{-1.0}$ kpc, the $\\sqrt{\\mu }d$ of our sample is one to six times the observed intrinsic diameters (1.0\u20132.5 kpc). Lensing (or multiplicity) increases the apparent effective size of a source, so this comparison favours a lensing (or multiplicity) interpretation for the ACT-selected sources.","Citation Text":["Tacconi et al. 2006"],"Functions Text":["However, Simpson et al. (2015) point out that the submillimetre sizes are consistent with resolved 12CO detections, while the sizes derived from 1.4 GHz imaging are about two times larger because of the cosmic ray diffusion, which can explain the results before higher frequency observations at ALMA were possible"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1286,1305]],"Functions Start End":[[950,1263]]}
{"Identifier":"2021ApJ...908..164K__Beckwith_et_al._2006_Instance_1","Paragraph":"While these S\/Ns may not look promising, there is an interesting question that one can ask: How much LUVOIR-15 m time is needed to detect the present Earth-level concentration of NO2 around a Sun-like star at 10 pc? Figure 6(a) shows the spectrum of the difference in geometric albedo with and without NO2. Also shown as dashed lines are the noise levels for 10 (blue), 300 (red), and 1200 (black) hr of LUVOIR-A observation times. The present Earth level of NO2 seems to be well above the noise level after 300 hr of observation time (compare the solid green curve with dashed red line), indicating that it might be detectable. To find out with what S\/N it would be detectable, Figure 6(b) shows the \u201cnet S\/N\u201d to detect present Earth-level NO2 as a function of observation time. To achieve a net S\/N of 5 (dashed red line), it would take LUVOIR-15 m about 400 hr. For comparison, to obtain the Hubble Ultra Deep Field (UDF) image, \u223c400 hr of actual observation time (\u223c1 yr in real time) was needed (Beckwith et al. 2006). In fact, Hubble has run even larger programs, such as the CANDLES galaxy evolution survey (Grogin et al. 2011) with 902 orbits (\u223c900 hr of observation time assuming \u223c1 hr per orbit). This took about 3 yr in real time. However, these large programs also obtained data on a huge sample size with thousands of galaxies. LUVOIR is envisaged to be 100% community-competed time, and the final report of the LUVOIR team laid out a Design Reference Mission (DRM) in which comparable allocations of time were spent on general astrophysics observations and exoplanet detection and characterization observations during the first 5 yr of the mission. So, over the course of the nominal LUVOIR mission lifetime of about 5 yr, it may be possible to take data with \u223c400 hr observation time on a prime HZ planet candidate(s) within 10 pc, to potentially obtain an S\/N \u223c 5 for an NO2 feature at the present Earth level on an Earth\u2013Sun system at 10 pc. An even more interesting aspect is that we can place upper limits on the amount of NO2 available on that planet as we spend more observation time on a prime HZ candidate. This could potentially indicate the presence or absence or the level of technological civilization on that planet.","Citation Text":["Beckwith et al. 2006"],"Functions Text":["For comparison, to obtain the Hubble Ultra Deep Field (UDF) image, \u223c400 hr of actual observation time (\u223c1 yr in real time) was needed"],"Functions Label":["Background"],"Citation Start End":[[1000,1020]],"Functions Start End":[[865,998]]}
{"Identifier":"2022MNRAS.512.4136C__P\u00e9rez-Montero_&_Contini_2009_Instance_1","Paragraph":"If we recall the tight, monotonic dependence of the position of galaxies along the SF sequence in the diagram with metallicity (as outlined in Section 3.1), we can interpret our global results of Figs 4 and 5 as a manifestation of the existence of an O\/H versus N\/O relation for SDSS star-forming galaxies, whose intrinsic scatter is reflected and, to some extent, translated into the observed distribution of emission line ratios within the [N\u2009ii]-BPT. A tight relationship between O\/H and N\/O abundances is indeed observed in both H\u2009ii regions and local galaxies, especially at M\u22c6 \u2273 109.5M\u2299 (Vila Costas & Edmunds 1993; van Zee et al. 1998; P\u00e9rez-Montero & Contini 2009; Pilyugin et al. 2012; Andrews & Martini 2013; Hayden-Pawson et al. 2021), and it is set by the predominant nucleosynthetic origin of nitrogen from CNO burning of pre-existing stellar carbon and oxygen in low- and intermediate-mass stars experiencing the AGB phase (i.e. the \u2018secondary\u2019 nitrogen production mechanism, Kobayashi, Karakas & Umeda 2011; Ventura et al. 2013; Vincenzo et al. 2016); alternatively, Vincenzo & Kobayashi (2018) reproduced the observed N\/O-O\/H relation introducing failed supernovae (SNe) in massive stars within their cosmological simulations. Recently, such relationship between O\/H and N\/O has been suggested as even tighter than the one between M\u22c6 and N\/O (Hayden-Pawson et al. 2021), in contrast to what claimed by previous studies (e.g. Andrews & Martini 2013; Masters et al. 2016). In light of our results, this would confirm that deviations in N\/O at fixed O\/H are more likely to be related to the offset from the SF sequence in the [N\u2009ii]-BPT than relative variations in M\u22c6, although the two are clearly physically correlated. The connection between the two diagrams is also readily evident if we look at the distribution of our galaxy sample in the N\/O versus O\/H diagram, as shown in Fig. 8 (where [N\u2009ii] \u03bb6584\/[O\u2009ii] \u03bb3727, 29 is converted to N\/O following the Te-based calibrations presented in Hayden-Pawson et al. 2021); here, each hexagonal bin is colour-coded by the average distance D of galaxies from the best-fitting line of the [N\u2009ii]-BPT, almost perfectly tracing the scatter around the median N\/O versus O\/H relation.","Citation Text":["P\u00e9rez-Montero & Contini 2009"],"Functions Text":["A tight relationship between O\/H and N\/O abundances is indeed observed in both H\u2009ii regions and local galaxies, especially at M\u22c6 \u2273 109.5M\u2299"],"Functions Label":["Similarities"],"Citation Start End":[[643,671]],"Functions Start End":[[454,592]]}
{"Identifier":"2020AandA...633A..34C__Martell_&_Shetrone_2013_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":["Martell & Shetrone 2013"],"Functions Text":["Other works report, however, that LiRG can be randomly located in the HRD","The distinction is crucial to understanding the processes that may provide an explanation for the phenomenon,"],"Functions Label":["Differences","Motivation"],"Citation Start End":[[526,549]],"Functions Start End":[[365,438],[571,680]]}
{"Identifier":"2018ApJ...868..139W__Schlickeiser_&_Jenko_2010_Instance_1","Paragraph":"By radio continuum surveys of interstellar space and direct in situ measurements in the solar system, it is well established that for many scenarios the background magnetic fields are spatially varying. However, the above research about parallel and perpendicular diffusion only explored the uniform mean magnetic field. One can show that the spatially varying background magnetic fields lead to the adiabatic focusing effect of charged energetic particle transport and introduces correction to the particle diffusion coefficients (see, e.g., Roelof 1969; Earl 1976; Kunstmann 1979; Beeck & Wibberenz 1986; Bieber & Burger 1990; K\u00f3ta 2000; Schlickeiser & Shalchi 2008; Shalchi 2009b, 2011; Litvinenko 2012a, 2012b; Shalchi & Danos 2013; Wang & Qin 2016; Wang et al. 2017b). To explore the influence of adiabatic focusing on particle transport, the perturbation method is frequently used (see, e.g., Beeck & Wibberenz 1986; Bieber & Burger 1990; Schlickeiser & Shalchi 2008; Schlickeiser & Jenko 2010; Litvinenko & Schlickeiser 2013; He & Schlickeiser 2014). To use the perturbation method, since the anisotropic distribution function is an implicit function, by using the iteration method, one can find that the anisotropic distribution function becomes an infinite series of the spatial and temporal derivatives of the isotropic distribution function. Therefore, the governing equation of the isotropic distribution function derived from the Fokker\u2013Planck equation contains infinite terms because of the infinite series of the anisotropic distribution function. By using the truncating method to neglect the higher-order derivative terms, the approximate correction formulas of parallel or perpendicular diffusion coefficients were obtained (see, e.g., Schlickeiser & Shalchi 2008; Schlickeiser & Jenko 2010; He & Schlickeiser 2014). However, the higher-order derivative terms probably also make the correction to the parallel and perpendicular diffusion much like the lower-order derivative ones do. The magnitude of the correction from higher-order derivative term might not necessarily be a higher-order small quantity than the magnitude of the lower-order derivative terms. Therefore, the correction obtained by the previous authors is likely to contain significant errors. In this paper, by considering the higher-order derivative terms, we derive the parallel and perpendicular diffusion coefficients and obtain the correction formulas coming from all order derivative terms by using the improved perturbation method (He & Schlickeiser 2014) and the iteration operation. And for the weak adiabatic focusing limit we evaluate the correction to the parallel diffusion coefficient and compare it with the correction obtained in the previous papers.","Citation Text":["Schlickeiser & Jenko 2010"],"Functions Text":["One can show that the spatially varying background magnetic fields lead to the adiabatic focusing effect of charged energetic particle transport and introduces correction to the particle diffusion coefficients (see, e.g.,"],"Functions Label":["Background"],"Citation Start End":[[974,999]],"Functions Start End":[[321,542]]}
{"Identifier":"2018ApJ...854...26L__Tian_2017_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":["Tian 2017"],"Functions Text":["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"],"Functions Label":["Uses"],"Citation Start End":[[2299,2308]],"Functions Start End":[[2079,2279]]}
{"Identifier":"2022MNRAS.512..439C__Lian_et_al._2021_Instance_1","Paragraph":"It is still unclear whether this incompatibility is evidence against the spatially flat \u039bCDM model or is caused by unidentified systematic errors in one of the established cosmological probes or by evolution of the parameters themselves with the redshift (Dainotti et al. 2021b, 2022). Newer, alternate cosmological probes could help alleviate this issue. Recent examples of such probes include reverberation-mapped quasar (QSO) measurements that reach to redshift z \u223c 1.9 (Czerny et al. 2021; Khadka et al. 2021a,b; Yu et al. 2021; Zaja\u010dek et al. 2021), H\u2009ii starburst galaxy measurements that reach to z \u223c 2.4 (Mania & Ratra 2012; Ch\u00e1vez et al. 2014; Gonz\u00e1lez-Mor\u00e1n et al. 2019, 2021; Cao, Ryan & Ratra 2020, 2022a; Cao et al. 2021a; Johnson, Sangwan & Shankaranarayanan 2022; Mehrabi et al. 2022), QSO angular size measurements that reach to z \u223c 2.7 (Cao et al. 2017, 2020, 2021a; Ryan, Chen & Ratra 2019; Lian et al. 2021; Zheng et al. 2021), QSO flux measurements that reach to z \u223c 7.5 (Risaliti & Lusso 2015, 2019; Khadka & Ratra 2020a,b, 2021, 2022; Lusso et al. 2020; Yang, Banerjee & \u00d3 Colg\u00e1in 2020; Li et al. 2021; Lian et al. 2021; Luongo et al. 2021; Rezaei, Sol\u00e0 Peracaula & Malekjani 2021; Zhao & Xia 2021),1 and the main subject of this paper, gamma-ray burst (GRB) measurements that reach to z \u223c 8.2 (Amati et al. 2008, 2019; Cardone, Capozziello & Dainotti 2009; Cardone et al. 2010; Samushia & Ratra 2010; Dainotti et al. 2011, 2013a,b; Postnikov et al. 2014; Wang, Dai & Liang 2015; Wang et al. 2016, 2022; Fana Dirirsa et al. 2019; Khadka & Ratra 2020c; Hu, Wang & Dai 2021; Dai et al. 2021; Demianski et al. 2021; Khadka et al. 2021c; Luongo et al. 2021; Luongo & Muccino 2021; Cao et al. 2021a). Some of these probes might eventually allow for a reliable extension of the Hubble diagram to z \u223c 3\u20134, well beyond the reach of Type Ia supernovae. GRBs have been detected to z \u223c 9.4 (Cucchiara et al. 2011), and might be detectable to z = 20 (Lamb & Reichart 2000), so in principle GRBs could act as a cosmological probe to higher redshifts than 8.2.","Citation Text":["Lian et al. 2021"],"Functions Text":["Newer, alternate cosmological probes could help alleviate this issue. Recent examples of such probes include","QSO angular size measurements that reach to z \u223c 2.7"],"Functions Label":["Motivation","Background"],"Citation Start End":[[909,925]],"Functions Start End":[[286,394],[801,852]]}
{"Identifier":"2017AandA...601A...4A__Cernicharo_et_al._1999_Instance_1","Paragraph":"In addition to thermal excitation through collisions with H2 and He, absorption of infrared photons and pumping to excited vibrational states, followed by radiative decay to rotational levels in the ground-vibrational state, is an important excitation mechanism of molecules in IRC\u2009+10216 (Deguchi & Uyemura 1984; Ag\u00fandez & Cernicharo 2006; Gonz\u00e1lez-Alfonso et al. 2007; Ag\u00fandez et al. 2008, 2015; Cordiner & Millar 2009; Daniel et al. 2012; De Beck et al. 2012). Here, we have included excitation through infrared pumping for all studied species, mostly through bands lying in the mid-infrared, where the flux in IRC\u2009+10216 is large (Cernicharo et al. 1999). To facilitate the excitation and radiative transfer calculations, we have collapsed the fine rotational structure of the radicals and simply treated these species as linear molecules with a 1\u03a3 electronic state. For C2H, we have included the first four vibrationally excited states of the bending mode (\u03bd2 = 1, 2, 3, 4), and the first vibrationally excited states of the stretching modes (\u03bd1 = 1 and \u03bd3 = 1). The vibrationally excited state that plays the most important role, through infrared pumping, in the excitation of C2H in IRC\u2009+10216 is \u03bd2 = 1, which lies 371 cm-1 above the ground-vibrational state. The wavelengths and strengths of the vibrational bands have been taken from Tarroni & Carter (2004). For the radical CN, we have included the v = 0 \u2192 1 band, lying at 2042 cm-1 (H\u00fcbner et al. 2005; Brooke et al. 2014), which plays a minor role on the excitation of the \u03bb 3 mm lines in IRC\u2009+10216, however. For HC3N, we have included the first excited states of the vibrational bending modes \u03bd5 and \u03bd6, which have strong fundamental bands at 663 and 498 cm-1. The wavelengths and strengths of the vibrational bands are from the compilation by J. Crovisier1, which are based on extensive laboratory work (e.g., Uyemura et al. 1982; Jolly et al. 2007). For cyanodiacetylene, we have included the first excited states of the vibrational bending modes \u03bd7 and \u03bd8, whose calculated fundamental bands, lying at 566 and 685 cm-1, have been found to be important for the rotational excitation of HC5N in IRC\u2009+10216 (Deguchi & Uyemura 1984). For the radicals C4H, C6H, and C3N there is little information on the wavelengths and strengths of vibrational bands. For these species we have instead included a generic vibrationally excited state lying at 15 \u03bcm above the ground-vibrational state, with an Einstein coefficient of spontaneous emission of 5 s-1 for the P(1) transition of the vibrational band. A similar treatment, with slightly different parameters, was adopted for C4H and C6H by Cordiner & Millar (2009). ","Citation Text":["Cernicharo et al. 1999"],"Functions Text":["Here, we have included excitation through infrared pumping for all studied species, mostly through bands lying in the mid-infrared, where the flux in IRC\u2009+10216 is large"],"Functions Label":["Uses"],"Citation Start End":[[635,657]],"Functions Start End":[[464,633]]}
{"Identifier":"2020AandA...640L..11B__Segretain_1996_Instance_1","Paragraph":"Another possibly important cooling delay may arise from the phase separation of 22Ne during crystallization (Isern et al. 1991; Althaus et al. 2010). Our current best understanding is that at the small 22Ne concentrations typical of C\/O white dwarfs (\u223c1% by number), the presence of 22Ne should not affect the phase diagram, except near the azeotropic point of the C\/O\/Ne phase diagram. Thus, the crystallization of the C\/O core initially proceeds as in the case without 22Ne with no redistribution of neon ions between the solid and liquid phases. After a significant fraction of the core has crystallized, the temperature approaches the azeotropic point and the existing calculations indicate that the liquid phase is enriched in 22Ne relative to the solid (Segretain 1996; Garc\u00eda-Berro et al. 2008). The 22Ne-poor solid is lighter than the surrounding liquid and floats upward where it eventually melts. This gradually displaces the 22Ne-rich liquid downward toward the solid\u2013liquid interface until the azeotropic composition is reached, thereby releasing a considerable amount of gravitational energy. Given our very limited knowledge of the ternary C\/O\/Ne phase diagram (Segretain 1996; Hughto et al. 2012), this effect cannot be quantitatively implemented in our evolution models. However, we note that our current understanding of 22Ne phase separation is remarkably consistent with the missing cooling delay. In Fig. 2 we show the luminosity function obtained by adding an artificial 0.6 Gyr delay when 60% of the core is crystallized. These parameters are entirely consistent with those found in preliminary studies (Segretain 1996; Garc\u00eda-Berro et al. 2008) and yield an excellent fit to the crystallization pile-up3. Based on the current (albeit limited) knowledge of the C\/O\/Ne phase diagram, we propose that the phase separation of 22Ne in the advanced stage of crystallization significantly contributes to the pile-up in the luminosity function of 0.9\u22121.1\u2006M\u2299 white dwarfs (Fig. 2).","Citation Text":["Segretain 1996"],"Functions Text":["After a significant fraction of the core has crystallized, the temperature approaches the azeotropic point and the existing calculations indicate that the liquid phase is enriched in 22Ne relative to the solid"],"Functions Label":["Background"],"Citation Start End":[[760,774]],"Functions Start End":[[549,758]]}
{"Identifier":"2018AandA...616A..99K__Narang_et_al._2016_Instance_1","Paragraph":"The high-resolution imaging observations of TR from IRIS reveal the ubiquitous presence of network jets. We have used three different IRIS observations of the quiet sun, which are located near the disk center. On the basis of careful inspection, 51 network jets are identified from three QS observations and used for further analysis. These 51 network jets are very well resolved and are not affected by the dynamics of other jets. The study is focused on the rotating motion of network jets along with the estimation of their other properties (speed, height, and lifetime). The mean speed, as predicted by statistical distributions of the speed, is 140.16 km s\u22121 with a standard deviation of 39.41 km s\u22121. The mean speed of network jets is very similar, as reported in previous works (e.g., Tian et al. 2014; Narang et al. 2016). However, in case of their lifetimes, we found a value that is almost double (105.49 s) that of the previously reported mean lifetime of the network jets (49.6 s; Tian et al. 2014). As mentioned above, we took only those network jets that are very well resolved in space and in the time; these criteria exclude short lifetime network jets. Therefore, our statistical distribution of the lifetime predicts a higher mean lifetime. The mean length of the network jets is 3.16 Mm with a standard deviation of 1.18 Mm. In the case of CH network jets, Tian et al. (2014) have reported that most of the network jets have lengths from 4.0 to 10.0 Mm. However, the mean length for QS network jets is smaller (3.53 Mm; Narang et al. 2016). So, the mean length for QS network jets from the present work is in good agreement with Narang et al. (2016). In addition, the apparent speed and length of these network jets are positively correlated, which is very similar to what has already been reported in previous works (Narang et al. 2016). Finally, we can say that these networks jets are very dynamic features of the solar TR, as revealed by their estimated properties.","Citation Text":["Narang et al. 2016"],"Functions Text":["The mean speed of network jets is very similar, as reported in previous works (e.g.,"],"Functions Label":["Similarities"],"Citation Start End":[[810,828]],"Functions Start End":[[707,791]]}
{"Identifier":"2016AandA...589A..73R__Husser_et_al._2013_Instance_1","Paragraph":"Single-burst stellar population (SSP) models mimic uniform stellar populations of fixed age and metallicity, and are an important tool to study unresolved stellar clusters and galaxies. They are created by populating theoretical stellar evolutionary tracks with stars of a stellar library, according to a prescription given by a chosen initial mass function (IMF). Thus, the quality of the resulting SSP models depends significantly on the completeness of the used input stellar library in terms of evolutionary phases represented by the atmospheric parameters temperature, Teff, surface gravity, log\u2009(g), and metallicity. A sufficiently large spectral coverage is equally crucial when constructing reasonable SSP models. Theoretical stellar libraries like, e.g. BaSeL (Kurucz 1992; Lejeune et al. 1997, 1998; Westera et al. 2002), or PHOENIX (Allard et al. 2012; Husser et al. 2013) are generally available for both a large range in wavelength and in stellar parameters, whereas empirical libraries are found to be more incomplete in both respects. However, the advantage of the latter ones is that they are not hampered by the still large uncertainties in the calculation of model atmospheres. Examples of empirical stellar libraries in the optical wavelenth range encompass the Pickles library (Pickles 1998), ELODIE (Prugniel & Soubiran 2001), STELIB (Le Borgne et al. 2003), Indo-US (Valdes et al. 2004), MILES (S\u00e1nchez-Bl\u00e1zquez et al. 2006), and CaT (Cenarro et al. 2001, 2007). In the near-infrared (NIR) and mid-infrared (MIR)1, only very few empirical libraries have been observed so far (e.g. Lan\u00e7on & Wood 2000; Cushing et al. 2005; Rayner et al. 2009). The NASA Infrared Telescope Facility (IRTF) spectral library, described in the latter two papers, is to date the only empirical stellar library in the NIR and MIR which offers a sufficiently complete coverage of the stellar atmospheric parameter space to construct SSP models. In the future, the X-Shooter stellar library, which contains around 700 stars, and which covers the whole optical (see Chen et al. 2014) and NIR wavelength range until 2.5 \u03bcm, will clearly improve the current situation in the NIR. ","Citation Text":["Husser et al. 2013"],"Functions Text":["Theoretical stellar libraries like, e.g.","or PHOENIX","are generally available for both a large range in wavelength and in stellar parameters, whereas empirical libraries are found to be more incomplete in both respects. However, the advantage of the latter ones is that they are not hampered by the still large uncertainties in the calculation of model atmospheres."],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[864,882]],"Functions Start End":[[722,762],[832,842],[884,1195]]}
{"Identifier":"2022AandA...659A..41E__Hobbs_et_al._2005_Instance_1","Paragraph":"The age of a neutron star is difficult to measure, as for many other astronomical sources. The most robust way to do it is by identifying the birth supernova of the neutron star. However, this can be done precisely only for a very small number of objects, as 5\u201310 supernovae have historically been observed in our galaxy (Stephenson & Green 2005), and neutron stars are faint sources\u2013practically undetectable at distances beyond the Magellanic Clouds. The explosions, however, leave imprints in the interstellar medium that can remain visible at radio wavelengths for 10\u2005\u2212\u2005100 kyr (Sarbadhicary et al. 2017), thereby allowing the association of pulsars with supernova remnants (SNRs). However, pulsars are rarely found at the centre of SNRs (Frail et al. 1994), as most are expelled like bullets during the explosions possibly due to asymmetries in the process (e.g. Socrates et al. 2005). The transverse velocities of pulsars (based on proper motion and distance estimates) are particularly large, with a mean close to 310 km s\u22121 (Hobbs et al. 2005), which is at least ten times larger than the average velocities for stars in the solar neighbourhood (e.g. Gaia Collaboration 2018). Moreover, some measured velocities range as high as 1000 km s\u22121 (Chatterjee et al. 2005; Deller et al. 2019). Thus, associations between SNRs and pulsars are not always straightforward to make (e.g. see the chapter on young pulsars in Lyne & Graham-Smith 2012). The farther the pulsar is from the explosion site, the higher the possibility that the pulsar and SNR are unrelated. In order to confirm an association, it could be necessary to account for up to 100 kyr of evolution of the SNR (that we assume as the maximum possible age of a SNR), and movement across the Galaxy of the pulsar (e.g. Suzuki et al. 2021). In some situations, proper motion measurements for the pulsars can shed light on the matter. For an association to be secure, the pulsar must be moving away from where the explosion took place (usually adopted as the centre of the SNR), and the time necessary to move the pulsar to its current position must match the age of the system. If such time coincided with an independent age measurement of the SNR or the pulsar, or both, then the association would be concretely confirmed. However, this is rarely possible as SNR and pulsar ages are hard to obtain.","Citation Text":["Hobbs et al. 2005"],"Functions Text":["The transverse velocities of pulsars (based on proper motion and distance estimates) are particularly large, with a mean close to 310 km s\u22121"],"Functions Label":["Background"],"Citation Start End":[[1032,1049]],"Functions Start End":[[890,1030]]}
{"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)"],"Functions Text":["Since for this present UM, the estimation of H0 from CMB alone is compatible with the local estimation of H0 by","thus, we can safely add both the data sets to see whether we could have something interesting."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[409,428]],"Functions Start End":[[297,408],[430,524]]}
{"Identifier":"2022AandARv..30....6M__Blanton_et_al._2001_Instance_1","Paragraph":"In a similar fashion to what done by the CARLA survey, the COBRA (Clusters Occupied by Bent Radio AGN) program (Paterno-Mahler et al. 2017) searches for overdense regions around radio-AGN with double-lobed structures which are not aligned with each other, but bent by forming angles 180\u2218\\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}$$180^\\circ$$\\end{document}. The rationale behind this search is that the radio lobes of these AGNs are most likely bent because of the ram pressure that occurs due to the relative motion of the AGN host galaxy and the ICM (e.g., Feretti et al. 1992; Blanton et al. 2001; Giacintutti and Venturi 2009; Wing and Blanton 2011), which makes these sources good tracers for finding galaxy clusters. Indeed, out of 646 bent radio-AGN, 530 (corresponding to \u223c\\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}$$\\sim$$\\end{document} 82% of the original sample) are associated with over-densities\u2014mostly at high, z=1-3\\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}$$z=1-3$$\\end{document}, redshifts\u2014in the Spitzer\/IRAC maps, and 190 are associated with galaxy cluster candidates. By following up on the previous work, Golden-Marx et al. (2021) also show for a subsample of 36 high-z (0.35z2.2\\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}$$0.35z2.2$$\\end{document}) cluster candidates that radio-AGN with narrower (\u227280\u2218\\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}$$\\lesssim 80^\\circ$$\\end{document}) opening angles reside in richer clusters (cf. Fig. 26), clearly indicating that the cluster environment impacts radio morphology.","Citation Text":["Blanton et al. 2001"],"Functions Text":["The rationale behind this search is that the radio lobes of these AGNs are most likely bent because of the ram pressure that occurs due to the relative motion of the AGN host galaxy and the ICM (e.g.,",", which makes these sources good tracers for finding galaxy clusters."],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[802,821]],"Functions Start End":[[580,780],[875,944]]}
{"Identifier":"2017AandA...608A...8L__Momose_et_al._(2014)_Instance_1","Paragraph":"At high redshift, the mapping of the extended Ly\u03b1 haloes around galaxies (non-AGN) is however a lot more difficult because of sensitivity and resolution limitations. Detections of extended Lyman alpha emission at high redshift have been obtained in the past. While some large Ly\u03b1 blobs have been observed (e.g. Steidel et al. 2000; Matsuda et al. 2004, 2011), most of these studies were forced to employ stacking analyses because of sensitivity limitations. The first tentative detections of Ly\u03b1 haloes around normal star-forming galaxies emitting Ly\u03b1 emission using narrowband (NB) imaging methods were reported by M\u00f8ller & Warren (1998) and Fynbo et al. (2001). Later, Hayashino et al. (2004) observed 22 Lyman break galaxies (LBG) and detected extended Ly\u03b1 emission by stacking the NB images. These authors were followed six years later by Ono et al. (2010) who detected Ly\u03b1 haloes in their composite NB images of 401 Ly\u03b1 emitters (LAEs) at z = 5.7 and 207 at z = 6.6. Matsuda et al. (2012) and Momose et al. (2014) significantly increased the size of LAEs samples used by stacking \u22482000 and \u22484500 LAEs at redshift z \u2243 3 and 2.2 \u2264 z \u2264 6.6, respectively. Momose et al. (2014) found typical Ly\u03b1 halo exponential scale lengths of 5\u201310 physical kpc. Matsuda et al. (2012) found that Ly\u03b1 halo sizes are dependent on environmental density; these halo sizes extend from 9 to up to 30 physical kpc towards overdense regions. More recently, Xue et al. (2017) studied \u22481500 galaxies in two overdense regions at z \u2248 3 and 4. Using stacking methods these authors reported Ly\u03b1 halo exponential scale lengths of 5\u20136 physical kpc and found that Ly\u03b1 halo sizes correlate with the UV continuum and Ly\u03b1 luminosities, but not with overdensity. Steidel et al. (2011) stacked 92 brighter (RAB \u2243 24.5) and more massive LBGs at z = 2.3\u22123, finding large Ly\u03b1 extents of \u224880 physical kpc beyond the mean UV continuum size at a surface brightness level of ~10-19 erg\u2009s-1\u2009cm-2 arcsec-2. Put together, all these studies showed that Lyman alpha emission is on average more spatially extended than the UV stellar continuum emission from galaxy counterparts. ","Citation Text":["Momose et al. (2014)","Momose et al. (2014)"],"Functions Text":["significantly increased the size of LAEs samples used by stacking \u22482000 and \u22484500 LAEs at redshift z \u2243 3 and 2.2 \u2264 z \u2264 6.6, respectively.","found typical Ly\u03b1 halo exponential scale lengths of 5\u201310 physical kpc."],"Functions Label":["Background","Background"],"Citation Start End":[[998,1018],[1157,1177]],"Functions Start End":[[1019,1156],[1178,1248]]}
{"Identifier":"2021MNRAS.507..904N__Castor,_Abbott_&_Klein_1975_Instance_1","Paragraph":"In equations (2) and (3), $\\boldsymbol{f}_{\\rm rad}=(f_{{\\rm rad},\\, r},\\, f_{{\\rm rad},\\theta })$ is the radiation force described as\n(6)$$\\begin{eqnarray*}\r\n\\boldsymbol{f}_{\\rm rad}=\\frac{\\sigma _{\\rm e} \\boldsymbol{F}_{\\rm D}}{c}+\\frac{\\sigma _{\\rm e} \\boldsymbol{F}_{\\rm line}}{c}M,\r\n\\end{eqnarray*}$$where \u03c3e is the mass-scattering coefficient for free electrons, $\\boldsymbol{F}_{\\rm D}$ is the radiation flux emitted from the accretion disc integrated by the wavelength throughout the entire range, and $\\boldsymbol{F}_{\\rm line}$ is the line-driving flux, which is the same as $\\boldsymbol{F}_{\\rm D}$ but integrated across the UV band of 200\u20133200\u2009\u00c5. The second term of equation (6) is the line force. As mentioned in Section 1, the line transitions depend on the wavelength of the radiation. The line force is exerted mainly by the radiation flux in the UV band (200\u20133200\u2009\u2009\u00c5), because the metal lines are densely distributed (e.g. Castor, Abbott & Klein 1975; Stevens & Kallman 1990). Thus, in this paper, we evaluate the line force using the UV (200\u20133200\u2009\u2009\u00c5) radiation flux and the corresponding force multiplier M same as Proga & Kallman (2004). Here M is the numerical factor indicating how much the spectral lines enhance the radiation force compared to the Thomson scattering. The radial components of the fluxes are estimated as $F^r_{\\rm D}= F^r_{\\rm D,\\, thin}e^{-\\tau _{\\rm e}}$ and $F^r_{\\rm line}= F^r_{\\rm line,\\, thin}e^{-\\tau _{\\rm e}}$, respectively, where \u03c4e is the electron-scattering optical depth estimated as $\\tau _{\\rm e} =\\int ^r_{r_{\\rm in}} \\sigma _{\\rm e} \\rho (r^{\\prime },\\, \\theta) \\mathrm{ d}r^{\\prime }$, where rin is the inner boundary of the computational box. We ignore the attenuation in the \u03b8-direction as $F^\\theta _{\\rm D}= F^\\theta _{\\rm D,\\, thin}$ and $F^\\theta _{\\rm line}= F^\\theta _{\\rm line,\\, thin}$. We calculate $\\boldsymbol{F}_{\\rm D,\\, thin}$ and $\\boldsymbol{F}_{\\rm line,\\, thin}$ by integrating intensity transferred in the optically thin media from the grids on the disc to the point of interest. Here we employ the standard disc model (Shakura & Sunyaev 1973). We divide the hot region of the disc where the effective temperature is larger than $3\\times 10^3\\, {\\rm K}$ into the grids. In contrast to the previous method where we prepared 4096 uniform grids both in the r- and \u03c6-directions, we here prepare 12\u2009800 grids whose sizes are determined by \u0394ri\/\u0394ri \u2212 1 = 1.0005 in the r-direction. In the \u03c6-direction, we set 1600 uniform grids in the range of 0 \u2264 \u03c6  2\u03c0. In order to resolve the hot region of the disc ($T_{\\rm eff}\\gt 3\\times 10^3 \\, {\\rm K}$) for the IMBHs (MBH = 103\u2013$10^6 \\, M_\\odot$), a large number of grids is required in the r-direction. This is because the size of the hot region normalized by Schwarzschild radius RS increases with the decrease of the BH mass. For example, the outer radius of the hot region is \u223c1000RS for $M_{\\rm BH}=10^8\\, M_\\odot$ while it reaches as far as 3 \u00d7 104RS for $M_{\\rm BH}=10^3\\, M_\\odot$.","Citation Text":["Castor, Abbott & Klein 1975"],"Functions Text":["The line force is exerted mainly by the radiation flux in the UV band (200\u20133200\u2009\u2009\u00c5), because the metal lines are densely distributed (e.g."],"Functions Label":["Background"],"Citation Start End":[[940,967]],"Functions Start End":[[801,939]]}
{"Identifier":"2020ApJ...901...10D__Raddi_et_al._2015_Instance_1","Paragraph":"In order to calculate oxygen fugacities, we follow the methods described by Doyle et al. (2019). From the element abundance ratios, we assign oxygen to Mg, Si, Ca, and Al in the necessary proportions to obtain the relative abundances of the charge-balanced rock-forming oxide components MgO, SiO2, CaO, and Al2O3. The remaining excess oxygen (\n\n\n\n\n\n), is assigned to Fe to make FeO until either O or Fe is exhausted. The excess oxygen available to make FeO is obtained using\n4\n\n\n\n\n\nwhere \n\n\n\n\n\n is the amount of oxygen needed to form the metal oxide, i, and OTotal is the total abundance of oxygen in the system. Other studies have used similar methods for budgeting oxygen (Klein et al. 2010, 2011; Farihi et al. 2011, 2013, 2016; Raddi et al. 2015). Once the relative abundances of the oxides are obtained, they are normalized to 1, yielding mole fractions and permitting application of Equation (3). In principle, if insufficient oxygen exists to pair with Fe to make FeO, then the Fe that remains should have been present as a metal in the accreted parent body. We emphasize that oxygen fugacity is recorded by the mole fraction of FeO, which depends on all of the oxides (FeO, SiO2, MgO, \n\n\n\n\n\n, CaO) and not simply the FeO\/Fe ratio for the body. It is possible for metal and water to have coexisted in the parent body if it was undifferentiated, meaning that oxygen which is attributed to FeO in this calculation may have existed as H2O in the parent body. However, during the differentiation of a rocky body, the oxygen from ices will oxidize metallic Fe to form FeO. We are assuming here that the bodies we are observing in the WDs were either differentiated themselves, or they are the building blocks of differentiated bodies (chondrite meteorites would be the appropriate analog). Where a parent body was composed in part of Fe metal and H2O, our calculation is a measure of the prospects for FeO, and thus the \u0394IW expected for the body taken as a whole, including accreted rock and ices.","Citation Text":["Raddi et al. 2015"],"Functions Text":["Other studies have used similar methods for budgeting oxygen"],"Functions Label":["Background"],"Citation Start End":[[732,749]],"Functions Start End":[[613,673]]}
{"Identifier":"2022AandA...666A.153D__Dartois_2005_Instance_1","Paragraph":"In the first stages of star formation, protostars are still embedded in their parental cloud, where an active gas-grain chemistry is at work. Using either (i) background stars for dense clouds or (ii) a nascent protostellar object once it is able to emit sufficient light flux in the vibrational infrared wavelength range, or in a few protoplanetary disks well inclined towards the observer, the infrared pencil beam allows the probing of the composition of the cloud or circumstellar dust. The low-temperature ice mantles formed on top of or mixed with refractory dust (silicates and\/or organics) can be retrieved. A harvest of astronomical observations from ground-based telescopes (e.g. UKIRT, IRTF, CFHT, and VLT) or satellites (e.g. IRAS, ISO, Akari, and Spitzer) of such lines of sight has led, since the late seventies, to the deciphering of the chemical compositions, column densities, and variations associated with these ice mantles (e.g. Boogert et al. 2008, 2015; \u00d6berg et al. 2011; Dartois 2005; van Dishoeck 2004; Gibb et al. 2004; Keane et al. 2001; Dartois et al. 1999b; Brooke et al. 1999, and references therein). The interpretation of these observed spectra is mainly based on their comparison with the infrared spectra of laboratory-produced ice films of well-controlled composition and cryogenic temperatures (e.g. Hudson et al. 2014, 2021; Palumbo et al. 2020; Rachid et al. 2020; Terwisscha van Scheltinga et al. 2018; \u00d6berg et al. 2007; Dartois et al. 2003, 1999a,b; Moore & Hudson 1998; Ehrenfreund et al. 1997; Gerakines et al. 1995; Hudgins et al. 1993). The routes investigated are the influence of the ice mixture on the line width and position, temperature modifications, segregation (phase separation), and\/or intermolecular interactions (polar or apolar ices and molecular complexes). The impact of a distribution of grain shapes, mainly in the Rayleigh regime, is also explored in some cases. The literature is dominated by analyses based on the decomposition of the observed astronomical profiles into principal components from different ice mixtures. When dust grains evolve from the diffuse interstellar medium (ISM) to the dense phase and the protoplanetary phases, grains grow. This will affect the observed profiles and is expected to be, at least partly, responsible for enhanced scattering effects in dense cloud evolution, often referred to as cloudshine or coreshine effects (Ysard et al. 2016, 2018; Saajasto et al. 2018; Jones et al. 2016; Steinacker et al. 2015; Lef\u00e8vre et al. 2014). The growth can also be inferred from the evolution of the silicate-to-K band ratio (\u03c49,7\/AK; e.g. Madden et al. 2022; van Breemen et al. 2011; Chiar et al. 2007).","Citation Text":["Dartois 2005"],"Functions Text":["A harvest of astronomical observations from ground-based telescopes (e.g. UKIRT, IRTF, CFHT, and VLT) or satellites (e.g. IRAS, ISO, Akari, and Spitzer) of such lines of sight has led, since the late seventies, to the deciphering of the chemical compositions, column densities, and variations associated with these ice mantles (e.g."],"Functions Label":["Background"],"Citation Start End":[[995,1007]],"Functions Start End":[[616,948]]}
{"Identifier":"2021ApJ...919...33C__Staubert_et_al._2019_Instance_1","Paragraph":"The variability of the cyclotron line centroid energy in the spectra of XRPs is considered to be related to the geometry of accretion flow in close proximity to the surface of an NS. The geometry of the emitting region is related to the mass accretion rate. At low mass accretion rates, the radiation pressure is small, and one expects hot spots at the NS surface. At high mass accretion rates, the radiation pressure is high enough to stop accreting material above the NS surface. In this case, the accretion column supported by radiation pressure and confined by a strong magnetic field arises above the stellar surface (Basko & Sunyaev 1976; Wang & Frank 1981; Mushtukov et al. 2015a). The luminosity separating these two accretion regimes is called the \u201ccritical\u201d luminosity Lcrit. The critical luminosity was shown to be dependent on the magnetic field strength (Mushtukov et al. 2015b). The dynamics of the cyclotron line was shown to be dependent on the luminosity state of XRPs (see Staubert et al. 2019 for review). In particular, the positive correlation between the cyclotron line centroid energy and accretion luminosity is considered to be typical for subcritical XRPs (Staubert et al. 2007; Klochkov et al. 2012; F\u00fcrst et al. 2014), while the negative correlation was detected in bright supercritical sources (Mihara et al. 2004; Tsygankov et al. 2006; Boldin et al. 2013). At the same time, there are some sources without any observed correlation between the cyclotron line energy and luminosity (Caballero et al. 2007). Several theoretical models are aiming to explain the variability of a cyclotron line. The positive correlation was explained by the Doppler effect in the accretion channel (Mushtukov et al. 2015c) and alternatively by the onset of collisionless shock above hot spots at low mass accretion rates (Shapiro & Salpeter 1975; Rothschild et al. 2017). The negative correlation was explained by the variations of accretion column height above the NS surface. Different models consider different locations of a line-forming region at supercritical mass accretion rates, which might be a radiation-dominated shock on top of an accretion column (Becker et al. 2012) or NS surface (see e.g., Poutanen et al. 2013; Lutovinov et al. 2015; Mushtukov et al. 2018), which reprocesses a large fraction of beamed radiation from the accretion column. Alternatively, Nishimura (2014) argued that some variations of cyclotron line centroid energy could be related to the changes of a beam pattern. By considering the structure of an accretion column in two dimensions, Nishimura (2019) suggested that the line-forming region is a region around an accretion mound in which the bulk velocity in the line-forming region can be considerably different from that in the continuum-forming region, which is assumed to be inside an accretion mound, so that the variation of Ecyc results from the motion of the accretion mound in the different luminosity ranges.","Citation Text":["Staubert et al. 2019"],"Functions Text":["The dynamics of the cyclotron line was shown to be dependent on the luminosity state of XRPs (see","for review)."],"Functions Label":["Background","Background"],"Citation Start End":[[991,1011]],"Functions Start End":[[893,990],[1012,1024]]}
{"Identifier":"2022MNRAS.512.1499R__LeVeque_1992_Instance_1","Paragraph":"Let ui be the evolved quantity at the coordinate position xi. Then, THC_M1 approximates the derivative of the flux f(u) at the location xi as\n(31)$$\\begin{eqnarray}\r\n\\partial _x f (u) \\simeq \\frac{F_{i - 1 \/ 2} - F_{i + 1 \/ 2}}{\\Delta x},\r\n\\end{eqnarray}$$where Fi \u2212 1\/2 and Fi + 1\/2 are numerical fluxes defined at $x_i \\mp \\frac{\\Delta x}{2}$, respectively. The fluxes are constructed as linear combination of a non-diffusive second order flux $F^{\\operatorname{HO}}$ and a diffusive first order correction $F^{\\operatorname{LO}}$:\n(32)$$\\begin{eqnarray}\r\nF_{i + 1 \/ 2} = F_{i + 1 \/ 2}^{\\operatorname{HO}} - A_{i + 1 \/ 2} \\varphi _{i + 1 \/ 2} \\left(F_{i + 1 \/ 2}^{\\operatorname{HO}} - F_{i + 1 \/ 2}^{\\operatorname{LO}}\\right) .\r\n\\end{eqnarray}$$The term \u03c6i + 1\/2 is the so-called flux limiter (LeVeque 1992), while Ai + 1\/2 is a coefficient introduced to switch off the diffusive correction at high optical depth (more below). The role of the flux limiter is to introduce numerical dissipation in the presence of unresolved features in the solution u and ensure the non-linear stability of the scheme. In particular, if Ai + 1\/2\u03c6i + 1\/2 = 0 the second-order flux is used, while if Ai + 1\/2\u03c6i + 1\/2 = 1, then the low order flux is used. A standard second order non-diffusive flux is used for $F^{\\operatorname{HO}}$, while the Lax\u2013Friedrichs flux is used for $F^{\\operatorname{LO}}$:\n(33)$$\\begin{eqnarray}\r\nF^{\\operatorname{HO}}_{i + 1 \/ 2} = \\frac{f (u_i) + f (u_{i + 1})}{2},\r\n\\end{eqnarray}$$(34)$$\\begin{eqnarray}\r\nF^{\\operatorname{LO}}_{i + 1 \/ 2} = \\frac{1}{2} [f (u_i) + f (u_{i + 1})] - \\frac{c_{i + 1 \/ 2}}{2} [u_{i + 1} - u_i] .\r\n\\end{eqnarray}$$The characteristic speed in the Lax\u2013Friedrichs flux ci is taken to be the maximum value of the speed of light between the right and left cells\n(35)$$\\begin{eqnarray}\r\nc_{i + 1 \/ 2} = \\max _{a \\in \\lbrace i, i + 1 \\rbrace } \\left\\lbrace \\left| \\alpha _a \\sqrt{\\gamma _a^{x x}} \\pm \\beta _a^x \\right| \\right\\rbrace .\r\n\\end{eqnarray}$$We remark that it is known that the M1 system can, in some circumstances, lead to acausal (faster than light) propagation of neutrinos in GR (Shibata et al. 2011). For this reason, one might argue that a better choice of the characteristic velocity for the Lax\u2013Friedrichs formula would have been given by the eigenvalue of the Jacobian of $\\boldsymbol{F}$. These values are known analytically (Shibata et al. 2011), however in our preliminary tests we found that the use of the full eigenvalues resulted did not improve on the stability or accuracy of the M1 solver.","Citation Text":["LeVeque 1992"],"Functions Text":["The term \u03c6i + 1\/2 is the so-called flux limiter"],"Functions Label":["Uses"],"Citation Start End":[[796,808]],"Functions Start End":[[747,794]]}
{"Identifier":"2020MNRAS.494.2948P__Lyne_et_al._1990_Instance_1","Paragraph":"Spider pulsar systems are characterized by having a low-mass companion star in a compact orbit with an energetic millisecond pulsar (MSP) resulting in heavy irradiation of the companion by the pulsar\u2019s wind. The spider pulsar population has been observed to have a clearly bimodal distribution of companion star masses (Roberts 2011; Strader et al. 2019) made up of two distinct subgroups: black widows (BW) with companion star masses \u223c0.01\u20130.05\u2009M\u2299, and redbacks (RB) with companion star masses \u223c0.1\u20131\u2009M\u2299. A large proportion of the spiders, whether BWs or RBs, have been observed to exhibit (quasi-)periodic eclipses of the pulsars\u2019 radio emission (e.g. Fruchter, Stinebring & Taylor 1988; Lyne et al. 1990) that are generally attributed to excess material in the orbits \u2013 that has been driven from the companion stars by the pulsar irradiation (Podsiadlowski 1991; van den Heuvel & van Paradijs 1988; Phinney et al. 1988; Kluzniak et al. 1988) \u2013 interfering with the propagation of the radio emission. Studies of such eclipses are key for understanding mass loss from the irradiated companion stars, the properties of the medium causing the eclipses, interactions between the pulsar wind and the eclipse medium, and the mechanisms responsible for the apparent attenuation of pulsar radio emission during the eclipse. In the years after the initial BW discovery (Fruchter et al. 1988), there were a number of excellent early works (e.g. Ryba & Taylor 1991; Stappers et al. 2001a) investigating the observed radio eclipses. However, unfortunately, a lack of further in-depth eclipse analyses \u2013 largely as a result of difficult observing requirements and (until recently) a low number of known spider pulsars \u2013 has meant slow progress in reaching an understanding in any of these topics. However, the last few years have marked a revival of the field with detailed and novel studies beginning to give important insight into the nature of eclipsing pulsar systems (e.g. Broderick et al. 2016; Main et al. 2018; Li et al. 2019).","Citation Text":["Lyne et al. 1990"],"Functions Text":["A large proportion of the spiders, whether BWs or RBs, have been observed to exhibit (quasi-)periodic eclipses of the pulsars\u2019 radio emission (e.g.","that are generally attributed to excess material in the orbits","interfering with the propagation of the radio emission.","Studies of such eclipses are key for understanding mass loss from the irradiated companion stars, the properties of the medium causing the eclipses, interactions between the pulsar wind and the eclipse medium, and the mechanisms responsible for the apparent attenuation of pulsar radio emission during the eclipse."],"Functions Label":["Background","Background","Background","Motivation"],"Citation Start End":[[690,706]],"Functions Start End":[[506,653],[708,770],[947,1002],[1003,1317]]}
{"Identifier":"2016ApJ...833...76B__Klimchuk_et_al._2008_Instance_2","Paragraph":"A significant limitation of the model is that it ignores the well-established hydrodynamic evolution of the loop during the cooling process, involving the substantial transfer of mass between the chromosphere and the corona. For large downward heat fluxes, the transition region is unable to radiate the supplied energy, resulting in the deposition of thermal energy in the dense chromosphere. The resulting two to three orders-of-magnitude temperature enhancements create a large pressure gradient that drives an upward enthalpy flux of \u201cevaporating\u201d plasma. However, as the loop cools, the decreased heat flux becomes insufficient to sustain the radiation emitted in the now-dense transition region and hence an inverse process of downward enthalpy flux starts to occur. It has been suggested (Klimchuk et al. 2008) that the enthalpy fluxes associated with both evaporating and condensing plasma are at all times in approximate balance with the excess or deficit of the heat flux relative to the transition region radiation loss rate. This basic idea has allowed the development of global \u201cEnthalpy-Based Thermal Evolution of Loops\u201d (EBTEL) models that describe the evolution of the average temperature and density in the coronal part of the loops; these models are generally in good agreement with one-dimensional hydrodynamic simulations (Klimchuk et al. 2008; Cargill et al. 2012a, 2012b). It is, in principle, possible to include the effects of a turbulence-controlled heat flux in EBTEL (or 1D hydrodynamic) models. If this heat flux is reduced sufficiently relative to its collisional value, then, for the reasons explained above, there will be a significant impact on the thermal evolution of the loop. Doing so, however, would still require a numerical treatment, which is beyond the scope of the present work (but which it is our intention to carry out in a future work). Instead, we adopt a simpler approach that allows a systematic and fairly transparent quantitative analysis of the impact of turbulence on the thermodynamics of post-flare loops.","Citation Text":["Klimchuk et al. 2008"],"Functions Text":["This basic idea has allowed the development of global \u201cEnthalpy-Based Thermal Evolution of Loops\u201d (EBTEL) models that describe the evolution of the average temperature and density in the coronal part of the loops;","these models are generally in good agreement with one-dimensional hydrodynamic simulations"],"Functions Label":["Background","Similarities"],"Citation Start End":[[1343,1363]],"Functions Start End":[[1037,1250],[1251,1341]]}
{"Identifier":"2015MNRAS.449.1018Y__Essey_et_al._2010_Instance_1","Paragraph":"VHE gamma-rays from distant blazars suffer serious EBL absorption. For instance, the attenuation factor of the flux at \u223c1\u2009TeV emitted at z = 0.6 due to EBL absorption is \u223c10\u22124 (Dom\u00ednguez et al. 2011). If VHE gamma-rays from a blazar are produced in its jet, the intrinsic VHE spectrum after de-absorption would be very hard and not a simple power law (e.g. Archambault et al. 2014), which is hardly explained plausibly in leptonic models. In such a case, a leptohadronic jet model is used to explain VHE spectra of distant blazars (e.g. B\u00f6ttcher, Reimer & Marscher 2009; Yan & Zhang 2015), in which VHE emission is attributed to synchrotron emissions of relativistic protons and pair cascades created in proton\u2013photon (p\u03b3) interaction. However, the jet model cannot explain the VHE emission from PKS 1424+240 if its redshift z > 0.7\u20130.8 (Yan & Zhang 2015). Alternatively, it is recently proposed that VHE gamma-rays from distant blazars may be the secondary gamma-rays produced in the rectilinear propagation of the UHECRs escaping from these blazars (e.g. Essey & Kusenko 2010; Essey et al. 2010). In the latter case, UHECRs interact with background photons, i.e. EBL photons and microwave background (CMB) photons, which creates UHE electrons and photons through Bethe\u2013Heitler (BH) pair production and photo-meson production. These UHE electrons and photons would interact with background photons again, and then pair cascades are induced; the secondary photons are inverse-Compton-scattered (ICS) CMB photons by the pair cascades. Because of the large mean interaction path of UHECRs, these secondary photons are produced relatively close to the Earth (e.g. Lee 1998; Essey et al. 2011b; Murase et al. 2012). The UHECR induced cascade model is successful in explaining the observed VHE spectra of extreme high-synchrotron-peaked BL Lacertae objects (HBLs; e.g. Essey & Kusenko 2010, 2014; Murase et al. 2012; Aharonian et al. 2013; Takami et al. 2013). In particular, UHECR induced cascade model is able to explain the VHE emission from a blazar with redshift z > 1 (Aharonian et al. 2013; Essey & Kusenko 2014). However, in the previous works the primary emission produced in the jet is either simply assumed (e.g. Aharonian et al. 2013; Essey & Kusenko 2014) or neglected (e.g. Takami et al. 2013).","Citation Text":["Essey et al. 2010"],"Functions Text":["Alternatively, it is recently proposed that VHE gamma-rays from distant blazars may be the secondary gamma-rays produced in the rectilinear propagation of the UHECRs escaping from these blazars (e.g.","In the latter case, UHECRs interact with background photons, i.e. EBL photons and microwave background (CMB) photons, which creates UHE electrons and photons through Bethe\u2013Heitler (BH) pair production and photo-meson production. These UHE electrons and photons would interact with background photons again, and then pair cascades are induced; the secondary photons are inverse-Compton-scattered (ICS) CMB photons by the pair cascades."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1079,1096]],"Functions Start End":[[857,1056],[1099,1533]]}
{"Identifier":"2018ApJ...853...12M__Weaver_et_al._1978_Instance_1","Paragraph":"There are numerous models which lay down nucleosynthesis yields from CCSN explosions (Burbidge et al. 1957; Woosley & Weaver 1986, 1995; Thielemann et al. 1996; Heger & Woosley 2002, 2010; Rauscher et al. 2002; Kobayashi et al. 2006; Nomoto et al. 2006; Woosley & Heger 2007; Pignatari et al. 2015; Sukhbold et al. 2016). In this work, we utilize the zonal yield sets from Sukhbold et al. (2016; hereafter, S16), which used the modified 1D hydrodynamic code KEPLER4\n\n4\n\nhttps:\/\/2sn.org\/kepler\/doc\/Introduction.html\n\n along with P-HOTB. P-HOTB stands for Prometheus-Hot Bubble and was used to study core collapse (Janka & Mueller 1996; Kifonidis et al. 2003), whereas KEPLER was used to evolve the star along zero age main sequence and calculate nucleosynthesis yields and light curves (Weaver et al. 1978). For models that exploded, isotopic yields were generated post explosion. The zonal yields were obtained for three particular models, 15.2 M\u2299, 20.1 M\u2299, and 25.2 M\u2299, \u223c200 s after the explosion, before any mixing could take place.5\n\n5\nT. Sukhbold (2017, private communication).\n Although all models in S16 data set assume solar metallicity and do not take into account the effects of rotation, they can account for (1) detailed neutrino transport calculations using an improved explosion mechanism, as compared to Rauscher et al. (2002) and Woosley and Heger (2007); (2) a central engine that considers matter inside the collapsed core, unlike certain other models that investigated only the matter exterior to the central engines used; and (3) unlike previous nucleosynthesis models, all models6\n\n6\nEach model has a particular progenitor mass.\n used here are not exploded by injecting artificial energy because (a) models below 15 M\u2299 almost always explode, (b) models in 20\u201330 M\u2299 rarely explode, and (c) most models above 30 M\u2299 implode and become black holes (see Figure 14 in S16 for the probability of explosion of different progenitor masses). In fact, the few models above 30 M\u2299 in which explosion does take place is due to their core being ripped apart by winds to sizes comparable to \u223c15 M\u2299. The decimals in the progenitor masses of these models might seem bizarre; the reason is that the authors have tried to explode all possible progenitor masses in steps of 0.1 M\u2299 between 12 and 30 M\u2299; however, 15.0, 15.1, 20.0, 25.0, and 25.1 M\u2299 imploded in their simulations. This apparently small change in progenitor mass, which leads to an altogether different end scenario, is due to small but significant variations in the progenitor compactness (O\u2019Connor & Ott 2011) rather than the central engine characteristics (Pejcha & Thompson 2015). This effect is more pronounced near progenitor masses of \u223c20 M\u2299 because the carbon burning stage changes to the radiative pathway from a convective mechanism. In fact, it has been recently shown that two similar progenitors with identical masses but slightly different input physics can lead to totally different scenarios (Sukhbold et al. 2017). Thus it is not unusual for such stark differences to show up between two similar progenitor stars. Throughout this paper, we frequently approximate 15.2\u201315, 20.1\u201320, and 25.2\u201325 M\u2299 models for the sake of simplicity.","Citation Text":["Weaver et al. 1978"],"Functions Text":["whereas KEPLER was used to evolve the star along zero age main sequence and calculate nucleosynthesis yields and light curves"],"Functions Label":["Uses"],"Citation Start End":[[786,804]],"Functions Start End":[[659,784]]}
{"Identifier":"2018AandA...617A..94L__Joblin_et_al._2018_Instance_1","Paragraph":"In an interstellar cloud the spatial transition from atomic to molecular gas takes place in photon dominated regions (PDRs; see reviews by Hollenbach & Tielens 1997, 1999), which are also a source of a significant fraction of the far-infrared emission from the Milky Way and other galaxies. Exterior to the PDRs the gas makes the transition from neutral to ionized hydrogen. The ionized gas can take the form of a low density ionized boundary layer (IBL) in the case of weak UV fields, or a dense H\u202fII region in the proximity to a strong UV field arising from massive star formation. PDRs and H\u202fII regions are the boundary regions where the effects of star formation on molecular clouds manifest themselves. They have been the focus of a considerable modeling effort (see Tielens & Hollenbach 1985; Sternberg & Dalgarno 1989; Kaufman et al. 1999; Abel et al. 2005; Le Petit et al. 2006; Bron et al. 2018, and references therein). The observational analysis of PDRs, H\u202fII regions, and IBLs has improved considerably since the availability of far-infrared spectroscopic data from the Herschel Space Observatory (see Ossenkopf et al. 2013; K\u00f6hler et al. 2014; Stock et al. 2015; Joblin et al. 2018; Wu et al. 2018, and references therein) and the Stratospheric Observatory for Infrared Astronomy (SOFIA; e.g., Schneider et al. 2012; P\u00e9rez-Beaupuits et al. 2015; Pabst et al. 2017; Mookerjea et al. 2018). Most of these studies of the ionized and PDR layers have focused on very bright H\u202fII regions where high UV flux, density, and temperature produce strong far-infrared emission, making such regions easily observable in key gas tracers such as the fine-structure lines of C+, N+, and O. Less is known about the IBL\u2013PDR conditions for typical molecular clouds where the UV field is smaller and, thus, the lines are weaker. The Herschel Space Observatory HIFI GOT C+ survey (Langer et al. 2010; Pineda et al. 2013) took a step in studying molecular cloud PDRs and IBLs in that it sampled [C\u202fII] along several hundred lines of sight (LOS) in the Galaxy producing an unbiased database of a few thousand clouds of various evolutionary stages with most LOS not containing H\u202fII regions as indicated by weak [C\u202fII] emission. However, because [C\u202fII] samples both weakly and highly ionized regions, there remains some uncertainty about the relative contributions of the ionized and PDR regions. Furthermore, because [C\u202fII] has only one fine-structure transition one cannot solve uniquely for the properties of the gas. For the GOT C+ survey Langer et al. (2014) derived the column density of material traced by [C\u202fII] by assuming a thermal pressure and its Galactic gradient.","Citation Text":["Joblin et al. 2018"],"Functions Text":["The observational analysis of PDRs, H\u202fII regions, and IBLs has improved considerably since the availability of far-infrared spectroscopic data from the Herschel Space Observatory (see"],"Functions Label":["Background"],"Citation Start End":[[1176,1194]],"Functions Start End":[[930,1113]]}
{"Identifier":"2019AandA...630A.123K__Kohutova_&_Verwichte_2016_Instance_1","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"],"Functions Text":["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"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[464,489]],"Functions Start End":[[127,425]]}
{"Identifier":"2022MNRAS.512.3828B__Dokkum_&_Franx_2001_Instance_1","Paragraph":"Understanding the redshift dependence of galaxy evolution requires disentangling the evolution of an individual galaxy within a population from the evolution in the average properties of the population as a whole (due to the continual addition of newly formed galaxies with different properties to the extant population). A key method of studying how both individual and populations of galaxies evolve is by analysing their growth in mass and size. An individual star-forming galaxy is expected to evolve along the relation of star-forming galaxies in the mass\u2013size plane (Lilly et al. 1998; Ravindranath et al. 2004; Trujillo et al. 2006; Pezzulli et al. 2015; van Dokkum et al. 2015), while the average size of the population as a whole increases with decreasing redshift due to processes such as mergers (Hopkins et al. 2009; Naab, Johansson & Ostriker 2009) and \u2019puffing up\u2019 due to strong active galactic nuclei (AGNs) feedback (Fan et al. 2008, 2010), but also new galaxies forming with larger radii (van Dokkum & Franx 2001; Carollo et al. 2013). In other words, star-forming galaxies at both low and intermediate redshifts follow parallel tracks in the mass\u2013size plane (Speagle et al. 2014; van Dokkum et al. 2015), with their starting location in the plane depending on redshift. Therefore, when analysing star-forming galaxies in the mass\u2013size plane at $z$ \u223c 0, we are observing the combined effect of both the evolution of individual galaxies throughout their lifetimes and the evolution of the population due to the addition of new members and loss of old members (as they quench). We want to disentangle these two effects to understand how the redshift range over which a galaxy formed and evolved influences its evolutionary path. We therefore need to understand how the processes influencing a galaxy may be regulated by the broader conditions of the Universe and how these conditions change with redshift. An important tool to measure the impact of various processes is the analysis of scaling relations, which quantify the link between different galaxy parameters to determine their dependence. Specifically, scaling relations between stellar population parameters and galaxy structure allow us to quantify how processes involved in star formation and stellar mass assembly interrelate with processes dominating structural and dynamical changes.","Citation Text":["van Dokkum & Franx 2001"],"Functions Text":["An individual star-forming galaxy is expected to evolve along the relation of star-forming galaxies in the mass\u2013size plane","while the average size of the population as a whole increases with decreasing redshift due to processes such as mergers","and \u2019puffing up\u2019 due to strong active galactic nuclei (AGNs) feedback","but also new galaxies forming with larger radii"],"Functions Label":["Background","Background","Background","Background"],"Citation Start End":[[1006,1029]],"Functions Start End":[[449,571],[687,806],[862,931],[957,1004]]}
{"Identifier":"2020MNRAS.499.5562Z__Miller_2015_Instance_1","Paragraph":"One explanation for the low effective temperatures is that the TDE thermal emission does not originate from the accretion disc, but from an outflow supported by radiation pressure from the disc\u2019s super-Eddington accretion rate (e.g. Loeb & Ulmer 1997; Strubbe & Quataert 2009; Lodato & Rossi 2011; Metzger & Stone 2016; Roth et al. 2016; Curd & Narayan 2019). The TDE\u2019s high photosphere radius and low temperature can then be explained by the increased emitting area from the optically thick, expanding outflow, launched from the accretion disc or SMBH. These outflows can explain the nearly constant temperatures inferred from the spectrum of optically bright TDEs (Strubbe & Quataert 2009; Miller 2015), and may lead to observable emission or absorption line features in the TDE\u2019s spectrum (Strubbe & Quataert 2011; Roth et al. 2016; Roth & Kasen 2018). Hydrodynamical simulations show that the outflow can be supported not only by radiation pressure from the compact disc (Dai et al. 2018; Curd & Narayan 2019) but also by shocks driven by stream\u2013stream collisions during the circularization of stellar debris (Liptai et al. 2019; Lu & Bonnerot 2020). However, to power the outflow, a significant fraction of the tidally disrupted star\u2019s rest-mass energy must be liberated ($0.05 \\, {\\rm M}_\\odot \\, c^2 \\sim 10^{53} \\, {\\rm erg}$), much larger than the typical energy liberated by an optically bright TDE\u2019s early emission ($\\sim \\! 10^{49}{\\!-\\!}10^{51} \\, {\\rm erg}$; e.g. Komossa 2015; van Velzen et al. 2020). This so-called missing energy problem has a number of proposed solutions. For instance, some argue most of the rest-mass energy is radiated in the unobservable far-UV wavelength bands (e.g. Lu & Kumar 2018; Jonker et al. 2020), while others propose this energy is carried away by a jet whose emission is unobservable for most TDE viewing angles (Dai et al. 2018). Some authors suggest this energy is never emitted in the first place, but rather becomes trapped due to the TDE disc and outflow\u2019s high optical depth (photon trapping; e.g. Curd & Narayan 2019). The wind model has yet to conclusively address the missing energy problem.","Citation Text":["Miller 2015"],"Functions Text":["These outflows can explain the nearly constant temperatures inferred from the spectrum of optically bright TDEs"],"Functions Label":["Similarities"],"Citation Start End":[[692,703]],"Functions Start End":[[554,665]]}
{"Identifier":"2015MNRAS.452.2731S__Stroe_et_al._2013_Instance_2","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"],"Functions Text":["There is evidence for a few additional smaller shock fronts throughout the cluster volume"],"Functions Label":["Background"],"Citation Start End":[[1362,1379]],"Functions Start End":[[1271,1360]]}
{"Identifier":"2019AandA...632A..40D__S\u00e1nchez-Fern\u00e1ndez_et_al._2017_Instance_1","Paragraph":"An extreme example of a rapid decay is V404 Cyg. The parameters of the long Porb system were chosen to be close to those of V404 Cyg, whose last outburst lasted only a couple of weeks and showed a pronounced disc outflow, conjectured to be a thermal wind (Mu\u00f1oz-Darias et al. 2016). The mass in the disc at the onset of the outburst is \u223c3\u2005\u00d7\u200510\u22127\u2006M\u2299 in our models1. Blowing away most of this mass would require a sustained outflow rate of \u2248 10\u22125 M\u2299 yr\u22121 \u2248 30 \u1e40Edd over 15 days for a 9\u2006M\u2299 black hole. Such very high mass outflow rates may be reached close to the Eddington luminosity as electron scattering contributes to the driving force of the wind. Figure 7 shows the mass outflow rate in the wind diverges near L\u2004\u2248\u20040.7\u2006LEdd due to the estimated radiation driving correction (Eq. (4)). In principle, it might thus be possible to shorten the outburst of V404 Cyg to a couple of weeks by fine-tuning the model parameters to sample this high luminosity region. In support, observations of V404 Cyg do indicate the source likely reached LEdd (Kimura et al. 2016) and was enshrouded by rapidly varying Compton-thick outflowing material (S\u00e1nchez-Fern\u00e1ndez et al. 2017) with an estimated Mw\u2004\u2248\u20044\u2005\u00d7\u200510\u22126\u2006M\u2299 lost to the wind (Casares et al. 2019). The lower effective gravity due to the high radiation should enhance the wind (Proga & Kallman 2002) but the mass loss must saturate at some level as the outflow becomes optically thick. Higginbottom et al. (2019) do not find a significantly enhanced \u1e40w near LEdd in their radiation-hydrodynamic simulations of thermal winds, but these neglect radiative driving by electron scattering. If winds are boosted near Eddington, a puzzle is why GRS 1915+105 has not been affected as much as V404 Cyg despite its luminosity also being close to Eddington and its disc size even greater. If the short duration of the V404 Cyg was due to a thermal wind, then this wind likely required very specific conditions. Instead, we speculate that the angular momentum transport was instead dominated by the jet. The system likely stayed in the (very) bright hard state during the outburst, where it has a strong jet which is almost certainly coupled to the accretion flow via the magnetic fields and could be responsible for angular momentum transport through the hot flow in this state (e.g. Ferreira et al. 2006).","Citation Text":["S\u00e1nchez-Fern\u00e1ndez et al. 2017"],"Functions Text":["In support, observations of V404 Cyg do indicate the source","and was enshrouded by rapidly varying Compton-thick outflowing material"],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[1134,1163]],"Functions Start End":[[960,1019],[1061,1132]]}
{"Identifier":"2015ApJ...813..109D__G\u00f3rski_et_al._2005_Instance_1","Paragraph":"Photometric Calibration: Photometric calibration was performed using the stellar locus regression technique (SLR: Ivezi\u0107 et al. 2004; MacDonald et al. 2004; High et al. 2009; Gilbank et al. 2011; Coupon et al. 2012; Desai et al. 2012; Kelly et al. 2014). Our reference stellar locus was empirically derived from the globally calibrated DES Y1A1 stellar objects in the region of the Y1A1 footprint with the smallest E(B \u2212 V) value from the Schlegel et al. (SFD; 1998) interstellar extinction map. We performed a 1\u2033 match on all Y1 and Y2 objects with S\/N > 10 observed in r-band and at least one other band. We then applied a high-purity stellar selection based on the weighted average of the spread_model quantity for the matched objects (\n\n\n\n\n\n; see below). The average zero point measured in Y1A1, ZPgrizY = {30.0, 30.3, 30.2, 29.9, 28.0}, was assigned to each star as an initial estimate. Starting from this coarse calibration, we began an iterative procedure to fix the color uniformity across the survey footprint. We segmented the sky into equal-area pixels using the HEALPix scheme (G\u00f3rski et al. 2005). For each \u223c0.2 deg2 (resolution nside = 128) HEALPix pixel, we chose the DES exposure in each band with the largest coverage and ran a modified version of the Big MACS SLR code (Kelly et al. 2014)41\n\n41\n\nhttps:\/\/code.google.com\/p\/big-macs-calibrate\/\n\n to calibrate each star from the reference exposure with respect to the empirical stellar locus. These stars became our initial calibration standards. We then adjusted the zero points of other CCDs so that the magnitudes of the matched detections agreed with the calibration set from the reference exposure. CCDs with fewer than 10 matched stars or with a large dispersion in the magnitude offsets of matched stars (\u03c3ZP > 0.1 \n\n\n\n\n\n) were flagged. For each calibration star, we computed the weighted-average magnitude in each band using these new CCD zero points; this weighted-average magnitude was used as the calibration standard for the next iteration of the SLR. In the first iteration, we assigned SLR zero points to the calibration stars based on the HEALPix pixel within which they reside. In subsequent iterations, we assigned SLR zero points to the calibration stars based on a bi-linear interpolation of their positions onto the HEALPix grid of SLR zero points. After the second iteration, the color zero points were stable at the 1\u20132 \n\n\n\n\n\n level. The absolute calibration was set against the 2MASS J-band magnitude of matched stars (making use of the stellar locus in color-space), which were de-reddened using the SFD map with a reddening law of AJ = 0.709 \u00d7 E(B \u2212 V)SFD from Schlafly & Finkbeiner (2011). The resulting calibrated DES magnitudes are thus already corrected for Galactic reddening by the SLR calibration.","Citation Text":["G\u00f3rski et al. 2005"],"Functions Text":["We segmented the sky into equal-area pixels using the HEALPix scheme"],"Functions Label":["Uses"],"Citation Start End":[[1090,1108]],"Functions Start End":[[1020,1088]]}
{"Identifier":"2015ApJ...801..103G___2014_Instance_2","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"],"Functions Text":["The derived central engine parameters are sometimes ad hoc or inconsistent. For instance, the analyses","and for GRB 130427A","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."],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[832,851]],"Functions Start End":[[670,772],[811,830],[853,1099]]}
{"Identifier":"2021AandA...655A.111K__Rojas-Arriagada_et_al._(2019)_Instance_1","Paragraph":"Over the last decade, the radial and vertical dependences of the metallicity-alpha-element distribution have been studied in more and more detail with increasingly larger samples (e.g., Bensby et al. 2011; Anders et al. 2014; Nidever et al. 2014; Hayden et al. 2015; Queiroz et al. 2020). Figure 6 is mostly consistent with similar plots shown in the above papers. In the inner 10 kpc, it displays two over-densities, a high alpha-element (here [Mg\/Fe]), and a low one. Between Rg\u2004=\u20046 and 10 kpc, the two over-densities define two different sequences. In Appendix E, we note that when the sample is restricted to a \u00b1500 pc layer around the Galactic plane, two close but separated sequences are observed in the Rg\u2004\u2208\u2004[4,\u20066] kpc interval. Because of their scale height (Bovy et al. 2012), kinematics (Bensby et al. 2003), and age properties (Haywood et al. 2013), these two sequences are associated with the thick disc (high-alpha) and thin disc (low-alpha), respectively. Moving inward of Rg\u2004=\u20044\u2005\u2212\u20056 kpc, Fig. 6 shows that the two over-densities connect through a zone of lower density to form a single sequence. This is in agreement with the observations of Hayden et al. (2015), Bensby et al. (2017), Zasowski et al. (2019), Bovy et al. (2019), and Lian et al. (2020a,b), who also report a single sequence in the inner disc and\/or in the bulge\/bar area. Conversely, Rojas-Arriagada et al. (2019) and Queiroz et al. (2020) observe two sequences in the inner regions. In Appendix F, we compare the distributions of different APOGEE DR16 alpha elements in the ([Fe\/H], [\u03b1\/Fe]) plane (restricting the sample to the stars contained in the Rg\u2004\u2208\u2004[0,\u20062] kpc interval). The different elements produce different patterns: the global alpha-element abundance5 and oxygen show a double sequence, while magnesium, silicon, and calcium present a single sequence. This could explain, at least partly, why Queiroz et al. (2020), who use a combined \u03b1-element abundance, observe a double sequence, while we see a single one with magnesium. However, this does not explain the discrepancy with Rojas-Arriagada et al. (2019), who also used magnesium. Beyond Rg\u2004=\u200410 kpc, the high-alpha sequence gradually vanishes. This is in agreement with the finding that the thick disc has a shorter scale length than the thin disc (Bensby et al. 2011; Cheng et al. 2012; Bovy et al. 2012). It should be emphasised that in this paragraph the term \u2018sequence\u2019 is used in the geometrical sense. It does not presuppose the number of chemical tracks that form the sequence or sequences. In particular, based on Fig. 6, it can not be excluded that the single geometrical sequence observed in the inner disc be made of two chemical tracks, with the low-alpha one restricted to a narrow metallicity range. We discuss and propose an interpretation of the inner disc sequence in Sect. 5.","Citation Text":["Rojas-Arriagada et al. (2019)"],"Functions Text":["Conversely,","observe two sequences in the inner regions."],"Functions Label":["Differences","Differences"],"Citation Start End":[[1366,1395]],"Functions Start End":[[1354,1365],[1422,1465]]}
{"Identifier":"2018ApJ...862..150H__Cudlip_et_al._1982_Instance_1","Paragraph":"In Figure 1, we show the archival 850 \u03bcm polarization data from the legacy program of SCUPOL with the JCMT (Matthews et al. 2009) (effective 20\u2033 beam \u223c0.76 pc). Matthews et al. (2009) collected the JCMT data toward the GC, where the B-field is sampled on a 10\u2033 grid. The linearly polarized light from dust grains is frequently used to probe the integrated plane-of-sky B-field morphology. Interstellar dust grains are elongated, with their minor axes parallel to the B-field. The thermal emission from the aligned dust grains is then polarized with polarization segments perpendicular to the field lines (Cudlip et al. 1982; Hildebrand et al. 1984; Hildebrand 1988; Lazarian 2000; Andersson et al. 2015). The dust polarization can, therefore, reveal the plane-of-sky projected B-field orientations. In Figure 1, the SCUPOL B-field segments are overlaid on the JCMT 850 \u03bcm map (Di Francesco et al. 2008). The segments are plotted with p\/dp \u2265 2, 3, 4. The 450 \u03bcm Submillimeter Polarimeter for Antarctic Remote Observing (SPARO) B-field (Novak et al. 2003) is also overlaid with a resolution of 6\u2032 (corresponding to a linear scale of 13.7 pc). The low-resolution SPARO map traces the large-scale B-field which is parallel to the plane of our Galaxy. This alignment of field orientations with the Galactic plane is attributed to a large-scale toroidal B-field configuration (i.e., azimuthal field). The SCUPOL 850 \u03bcm B-field as well as the continuum are clearly detected and resolved along and across the sub-features in the GC at a 20\u2033 resolution (=0.76 pc). The detected B-field orientations (\u03a6B) vary enormously over the entire map, ranging from \u221290\u00b0 to 90\u00b0 (0\u00b0 is west, positive is counterclockwise). Nevertheless, the \u03a6B varies smoothly and systematically along certain substructures, revealing organized patches in, e.g., the CND, the giant molecular cloud 20\/50 MC (e.g., G\u00fcsten et al. 1981), and the HVCC CO 0.02\u20130.02 (Oka et al. 1999). The azimuthal correlation seen in the CND is providing the link between the model proposed by Wardle & K\u00f6nigl (1990) and the SCUPOL 850 \u03bcm polarization data. We are exploring this link in the following sections. A comparison of polarization data is presented in Appendix A. We find that the morphology of the B-field is consistent with different threshold cuts in p\/dp. Hence, in order to work with the maximum of independent data points, we present data with p\/dp \u2265 2, and we will focus on the B-field structure of the CND.","Citation Text":["Cudlip et al. 1982"],"Functions Text":["The thermal emission from the aligned dust grains is then polarized with polarization segments perpendicular to the field lines"],"Functions Label":["Uses"],"Citation Start End":[[605,623]],"Functions Start End":[[476,603]]}
{"Identifier":"2021ApJ...920..147Z__Schneider_1959_Instance_1","Paragraph":"Mainly four mechanisms are considered to be causing the observed solar radio wave emissions: two causing incoherent and two causing coherent emissions. Incoherent emissions can be due to bremsstrahlung and gyrosynchrotron radiation. In them, every electron radiates independent on the others. The total emission is simply the sum of the emissions of every single electron (Rybicki & Lightman 1979; Dulk 1985; Melrose 2017; Nindos 2020). Comparing with coherent emission mechanisms, the incoherent emission mechanisms are better understood. The two coherent emission mechanisms are plasma (Ginzburg & Zhelezniakov 1958; Melrose 1970a, 1970b; Zheleznyakov & Zaitsev 1970a, 1970b) and electron cyclotron maser (ECM) emissions (Twiss 1958; Gaponov 1959; Schneider 1959; Pritchett 1984a). In contrast to the incoherent emission mechanisms, coherent emission can explain (a) the high brightness temperatures, (b) the short eruption timescales, (c) the narrow frequency bands, and (d) the strong polarization of Type I, II, III solar radio bursts (SRBs) and solar radio spikes (Aschwanden 2005). And coherent radio emission mechanisms involve plasma instabilities. The plasma emission follows a beam or bump-on-tail instability, the source of free energy of which is related to an electron velocity distribution \n\n\n\n\n\n\nu\n\n\n\u2225\n\n\n\u00b7\n\u2202\n\n\nf\n\n\n\n\nu\n\n\n\u2225\n\n\n\n\n\n\/\n\n\u2202\n\n\nu\n\n\n\u2225\n\n\n>\n0\n\n\n. On the other hand, the ECM instability requires a positive gradient in the electron velocity distribution along the direction perpendicular to the ambient magnetic field (\n\n\n\n\n\u2202\n\n\nf\n\n\n\n\nu\n\n\n\u22a5\n\n\n\n\n\n\/\n\n\u2202\n\n\nu\n\n\n\u22a5\n\n\n>\n0\n\n\n) (see Melrose 1986, 2017). Here, f is the electron distribution function and u\u2225, u\u22a5 are the electron velocities parallel and perpendicular to the ambient magnetic field, respectively. The expressions \n\n\n\n\n\n\nu\n\n\n\u2225\n\n\n\u00b7\n\u2202\n\n\nf\n\n\n\n\nu\n\n\n\u2225\n\n\n\n\n\n\/\n\n\u2202\n\n\nu\n\n\n\u2225\n\n\n>\n0\n\n\n and \n\n\n\n\n\u2202\n\n\nf\n\n\n\n\nu\n\n\n\u22a5\n\n\n\n\n\n\/\n\n\u2202\n\n\nu\n\n\n\u22a5\n\n\n>\n0\n\n\n indicate an electron beam and a ring (or a loss cone or a horseshoe) distribution in the direction parallel and perpendicular to the ambient magnetic field, respectively. The existence of electron velocity distributions \n\n\n\n\n\n\nu\n\n\n\u2225\n\n\n\u00b7\n\u2202\n\n\nf\n\n\n\n\nu\n\n\n\u2225\n\n\n\n\n\n\/\n\n\u2202\n\n\nu\n\n\n\u2225\n\n\n>\n0\n\n\n has been conjectured based on observations of, e.g., SRBs and hard X-ray bursts as well as confirmed by observations of solar energetic particles (Chen et al. 2015, 2018; Cairns et al. 2018). Numerical simulations have also considered possible formation mechanisms of velocity distributions with \n\n\n\n\n\u2202\n\n\nf\n\n\n\n\nu\n\n\n\u22a5\n\n\n\n\n\n\/\n\n\u2202\n\n\nu\n\n\n\u22a5\n\n\n>\n0\n\n\n, e.g., by magnetic reconnection (e.g., B\u00fcchner & Kuska 1996; Bessho et al. 2014; Shuster et al. 2014; Zhou et al. 2016; Treumann & Baumjohann 2017; Mu\u00f1oz & B\u00fcchner 2018a; Voitcu & Echim 2018; Yao et al. 2021a). It is, however, still under debate which coherent mechanism dominates and which role different features of the electron distribution play for the emission of coherent radio waves from the solar corona.","Citation Text":["Schneider 1959"],"Functions Text":["The two coherent emission mechanisms are","and electron cyclotron maser (ECM) emissions"],"Functions Label":["Background","Background"],"Citation Start End":[[750,764]],"Functions Start End":[[540,580],[678,722]]}
{"Identifier":"2021ApJ...909...65K__Marsh_et_al._2016_Instance_1","Paragraph":"As mentioned in the previous section, if a magnetized WD rotates with a misalignment between its magnetic field and rotation axes (similar configuration to a pulsar), it can emit a continuous GW. We already provided a detailed discussion on GWs emitted from WDs with different magnetic field geometries and strengths in GR (Kalita & Mukhopadhyay 2019b; Kalita et al. 2020). Figure 7 shows an illustrative diagram of a magnetized WD where the magnetic field is along the z\u2032-axis and rotation is along the z-axis, with \u03c7 being the angle between these two axes. We calculate the amplitude of GW using the set of Equations (34) assuming the difference in radii of the WD between those along x- and z-axes to be 0.01%, i.e., \n\n\n\n\n\n, due to the presence of a very weak magnetic field and slow rotation. The choice of weak fields and slow rotation assures that the underlying WD mass\u2013radius solutions do not practically differ from the solutions based on the f(R) gravity without magnetic fields and rotation. In future, we plan to check rigorously by solving the set of equations, if indeed such \u03f5 is possible in the presence of weak magnetic fields and rotation keeping the mass and radius practically intact. As we will show below, however, the chosen \u03f5 appears to be the minimally required value to have any appreciable effect. Nevertheless, there are examples of weakly magnetized WD pulsars, which can be explained even in the GR framework, e.g., AE Aquarii (Bookbinder & Lamb 1987) and AR Scorpii (Marsh et al. 2016), where magnetic fields hardly affect their mass\u2013radius relations. Figure 8 shows the PSD as a function of frequency for various detectors along with \n\n\n\n\n\n over 5 s integration time for various f(R)-gravity-induced WD pulsars with different i assuming \u03c7 = 90\u00b0 and r = 100 pc. It is evident that while DECIGO and BBO can immediately detect such weakly magnetized super-Chandrasekhar WDs, the Einstein Telescope can detect them in \n\n\n\n\n\n minutes with S\/N \u2248 5 (see Figure 5(b)). However, for ALIA and LISA, the corresponding integration time respectively turns out to be \n\n\n\n\n\n days and \n\n\n\n\n\n yr3\n\n3\nNote that even if the threshold S\/N for detection increases slightly (say, from 5 to 20), many of these sources can still be detected in a few seconds to a few days of integration time depending on the type of the detectors.\n. Hence, it is also possible to detect such weakly magnetized WDs using ALIA, whereas for LISA it is quite impossible. Figure 5(b) depicts \n\n\n\n\n\n for these WDs with different integration times to show that S\/N increases if the integration time increases so that various detectors can detect them eventually. For such a system, the GW luminosity is given by (Zimmermann & Szedenits 1979)\n43\n\n\n\n\n\nIt is expected that a source can emit electromagnetic radiation in the presence of a magnetic field, and it is the dipole radiation in the case of a WD pulsar. However, because of the presence of a weak magnetic field, the dipole radiation emitted from such a WD is minimal, and the corresponding dipole luminosity is negligible as compared to LGW. Hence, the spin-down timescale is mostly governed by LGW, given by (Kalita et al. 2020)\n44\n\n\n\n\n\nFigure 9 shows the variation of LGW and P with respect to M for various WDs with \u03c7 = 90\u00b0. The maximum LGW in the case of a WD is \u223c1037 erg s\u22121. The empirical relations of LGW and P, in various branches, are same as in the previous case provided in Table 1. It is also clear from the figure that the massive WD pulsars are short-lived as compared to the lighter ones.","Citation Text":["Marsh et al. 2016"],"Functions Text":["Nevertheless, there are examples of weakly magnetized WD pulsars, which can be explained even in the GR framework, e.g.","and AR Scorpii","where magnetic fields hardly affect their mass\u2013radius relations."],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1498,1515]],"Functions Start End":[[1325,1444],[1482,1496],[1518,1582]]}
{"Identifier":"2015MNRAS.448..822X__Steinmetz_et_al._2006_Instance_1","Paragraph":"Various methods have been developed in the past to derive stellar atmospheric parameters from large number of medium-to-low resolution spectra (Recio-Blanco, Bijaoui & de Laverny 2006; Lee et al. 2008a, Wu et al. 2011). The approaches generally fall into two main categories of method (Wu et al. 2011): the minimum distance method (MDM) and non-linear regression method. Both categories of method have been applied to large stellar spectroscopic surveys, including the SEGUE (Yanny et al. 2009), RAVE (Steinmetz et al. 2006), APOGEE (Majewski et al. 2010) and LAMOST (Zhao et al. 2012). The MDM is usually based on spectral template matching, and searches for the template spectrum that has the shortest distance in parameter space from the target spectrum. The \u03c72 minimization, cross-correlation, weighted mean algorithm and the k-nearest neighbour are thought to be specific cases of MDM (Wu et al. 2011). Software and pipelines developed based on those algorithms include the tgmet (Katz et al. 1998), matisse (Recio-Blanco et al. 2006), sspp (Lee et al. 2008a), ulyss (Koleva et al. 2009), that of Allende Prieto et al. (2006) and of Zwitter et al. (2008). The non-linear regression method is sometimes also referred to as the artificial neural network. The method constructs a functional mapping between the spectra and stellar atmospheric parameters by training a library of template spectra with non-linear algorithms such as the principal component analysis, and then apply the mapping to target spectra. Related work can be found in Re Fiorentin et al. (2007) and Lee et al. (2008a). In addition to the above two categories of method, other approaches have been developed, for example, the line-index method based on the relations between the stellar atmospheric parameters and the equivalent widths (EWs) of spectral features and\/or photometric colours (Beers et al. 1999; Wilhelm, Beers & Gray 1999; Cenarro et al. 2002). More recently, a Bayesian approach to determine stellar atmospheric parameters combing spectral and photometric measurements has been developed by Sch\u00f6nrich & Bergemann (2014).","Citation Text":["Steinmetz et al. 2006"],"Functions Text":["Both categories of method have been applied to large stellar spectroscopic surveys, including","RAVE"],"Functions Label":["Background","Background"],"Citation Start End":[[502,523]],"Functions Start End":[[371,464],[496,500]]}
{"Identifier":"2016ApJ...832...41M__Helled_et_al._2011_Instance_1","Paragraph":"We first discuss the relative bulk enrichment. From internal structure models one can derive the MZ necessary to reproduce the observed mass and radius, and\u2014for the Solar System planets\u2014the gravitational moments. Studies inferring in this way eZ,rel,int of transiting exoplanets have found that eZ,rel,int decreases with increasing mass (left panel of Figure 3). The planetary mass where \n\n\n\n\n\n defines the parity mass M1 (see Appendix B). It is shown in the right panel of Figure 3. It is extrapolated to be between \u223c13 and 60 MJup (Miller & Fortney 2011; Thorngren et al. 2015). The planets analyzed in these studies have equilibrium temperatures of less than \u223c1000 K (corresponding to an orbital distance of about 0.08 au for a solar-like star) so that they are not affected by the aforementioned bloating mechanisms. A similar decrease of eZ,rel,int with increasing mass is found for the bulk metal content of Solar System giants (Saumon & Guillot 2004; Helled et al. 2011), where the mass where eZ,rel,int reaches 1 is extrapolated to be at about 11 MJup. From theoretical planet population syntheses based on the core accretion theory one finally finds that the parity mass is at about 10 to 18 MJup (Mordasini et al. 2014). Considering that of the 255 extrasolar giant planets (M sin i > 0.1 M) inside of 0.1 au currently listed on www.exoplanets.org (Han et al. 2014) only 4 have a mass exceeding 10 M (which is not an observational bias). We thus deduce that at least based on their masses regarding the bulk composition of hot Jupiters, it appears that almost all of them should be dominated by planetesimal enrichment. We add the caveat that the bulk heavy element content cannot be inferred directly for typical hot Jupiters at equilibrium temperatures of Teq \u2273 1500 K because of bloating mechanisms. But the fact that both the planets analyzed by Miller & Fortney (2011), Thorngren et al. (2015) (a = 0.03\u20131 au, Teq \u2272 1000 K) and the solar system planets (a \u2248 5\u201330 au) follow the same trend, makes it appear unlikely\u2014even though in principle not excluded\u2014that the hot Jupiters at a \u223c 0.04 au do not follow the same enrichment pattern.","Citation Text":["Helled et al. 2011"],"Functions Text":["A similar decrease of eZ,rel,int with increasing mass is found for the bulk metal content of Solar System giants","where the mass where eZ,rel,int reaches 1 is extrapolated to be at about 11 MJup."],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[958,976]],"Functions Start End":[[821,933],[979,1060]]}
{"Identifier":"2020ApJ...888...46C__Singh_et_al._1995_Instance_2","Paragraph":"The theoretical models mentioned above have made assumptions. The validity of these assumptions needs to be examined. Besides, all these models involve parameters to be determined. These problems could be better understood through numerical simulations. Hurlburt et al. (1994) performed a two-dimensional numerical simulation of compressible flow on downward overshooting. They investigated the dependence of subadiabatic extent \u03b4 on stability parameter S. They revealed a scaling law of \u03b4 \u221d S\u22121 in the penetrative layer and \u03b4 \u221d S\u22121\/4 in the overshooting layer, respectively. This result is in good agreement with Zahn\u2019s analytic model (Zahn 1991). Freytag et al. (1996) performed two-dimensional numerical simulations with a realistic description of radiation and ionization on A-type stars and DA white dwarf stars. They described the overshooting as a diffusive process, and derived an exponential decay parameter for the diffusion. Early attempts of low-resolution three-dimensional numerical simulations on overshooting were made by Singh et al. (1994; upward overshooting), Singh et al. (1995, 1998), and Saikia et al. (2000) (downward overshooting). The scaling laws derived from the numerical simulations of downward overshooting (Singh et al. 1995, 1998; Saikia et al. 2000) agree well with Zahn\u2019s analytical model. Only a qualitative result was given in the numerical simulations of upward overshooting (Singh et al. 1994). High-resolution numerical simulations of downward overshooting across a wide range of parameters were presented by Brummell et al. (2002). They confirmed the \u22121\/4 scaling law of the thermal adjustment overshooting layer, while the \u22121 scaling law of the nearly adiabatic penetrative layer was absent in the simulations. Based on a semianalytic model, Rempel (2004) argued that the absence of the nearly adiabatic penetrative layer is caused by the large energy flux specified in the numerical simulations. Numerical experiments on Boussinesq flow were performed by Korre et al. (2019). They reported steeper scaling laws of \u03b4 \u221d S\u22121\/3 or \u03b4 \u221d S\u22121\/2, depending on the steepness of the background radiative temperature gradient. Simulations with realistic physical variables on stellar core convection were performed by Browning et al. (2004) and Brun et al. (2005). The effects of rotation and magnetic field are considered. They found that the penetrative convection yields a prolate shape of a nearly adiabatic region. Kitiashvili et al. (2016) performed 3D radiative hydrodynamic simulations of the outer layer of a moderate-mass star (1.47 solar mass). Their result discovered a nearly adiabatic layer and a deeper subadiabatic layer. The recent work of Brun et al. (2017) simulated the differential rotation and overshooting in solar-like stars. Their result indicated that slow rotators favor a wider overshooting region near the poles and fast rotators at mid-to-low latitude. Hotta (2017) performed numerical simulations on the solar overshooting with very low energy fluxes F. He found that the overshooting distance obeys a scaling law of \u03b4 \u221d F0.31. K\u00e4pyl\u00e4 (2019) conducted numerical experiments on downward overshooting by considering the effect of the smoothness of the heat conduction profiles. He discovered that the power-law index of the overshooting distance on the energy flux is smaller in the smooth heat conduction profile than in the step profile. Efforts on prediction of 321D turbulent theory were made by Arnett et al. (2015) and Arnett & Moravveji (2017). They separated the overshooting region into three layers: a fully mixed layer, a partially mixed wave layer, and an extra diffusive mixing layer. With the scale analysis of turbulent plumes and eddies, Viallet et al. (2015) discussed the three possible regimes of turbulent overshooting: a diffusion-dominated regime (only mix composition), a penetrative regime (transition within the boundary layer), and an entrainment regime (mix both entropy and composition). The selection criterion of different regimes during a stellar evolution calculation is not well defined yet.","Citation Text":["Singh et al. 1995"],"Functions Text":["The scaling laws derived from the numerical simulations of downward overshooting","agree well with Zahn\u2019s analytical model."],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[1239,1256]],"Functions Start End":[[1157,1237],[1284,1324]]}
{"Identifier":"2020MNRAS.496.1718E__Wong_et_al._2019_Instance_2","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"],"Functions Text":["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"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[2258,2274]],"Functions Start End":[[2092,2256]]}
{"Identifier":"2015MNRAS.446.2468E__Choi,_Shlosman_&_Begelman_2013_Instance_1","Paragraph":"The equations of motions and hydrodynamics are solved by the ramses code (Teyssier 2002). The gas follows a piecewise polytropic equation of state (EoS) fitting the heating\/cooling equilibrium (see Kraljic et al. 2014, and references therein). A Jeans polytrope sets a pressure floor in the most refined volumes, to prevent artificial fragmentation. We refer the reader specifically to sections 2.2, 2.3 and 4.2 of R+13 and Kraljic et al. (2014) for further discussions on this and specific EoS (see Robertson & Kravtsov 2008; Tasker & Bryan 2008; Dobbs, Burkert & Pringle 2011; Bonnell, Dobbs & Smith 2013, for a few other implementation schemes). The resulting gas density probability distribution function (PDF; McKee & Ostriker 2007, and references therein) in the present simulation follows a classic log-normal shape (Nordlund & Padoan 1999; Padoan & Nordlund 2002) with an additional few\u2009per\u2009cent of the mass in a power-law tail at high density (Choi, Shlosman & Begelman 2013; R+13), as expected from gravity and observed in real molecular clouds and galaxies (Lombardi, Alves & Lada 2010; Druard et al. 2014). While changes in the AMR grid refinement can locally bias the velocity dispersion, the density and velocity power spectra are thus clearly dominated by a single turbulence cascade with a well-identified injection scale at the average Jeans length (Bournaud et al. 2010; R+13). The very high resolution of the present simulation allows to resolve the turbulent cascade with a realistic power spectrum (Combes et al. 2013; R+13) and density distribution (Druard et al. 2014) down to the parsec scale. The simulation comprises conversion of gas into stellar particles (down to 160 M\u2299) where the volume density \u03c10 exceed 2000 cm\u22123, assuming that the local star formation rate depends on the free-fall time and with the star formation efficiency set at 3\u2009per\u2009cent (R+13). This recipe does not take into account additional physics that may impact on the formation of molecules, and we thus rely on the EoS to follow the cloud collapse, the low temperatures and the high density which triggers the formation of new stars. These newly formed \u2018stars\u2019 are evolved on the AMR grid, i.e. with gravitational softening down to 0.05 pc, much smaller than the dark matter and primordial stellar components. The implementation of stellar feedback includes photoionization through heating, radiative pressure via injection of momentum and supernova explosions in the kinetic form (see R+13 for more details). A more thorough study of the impact of resolution or metallicity in such simulations was conducted by Kraljic et al. (2014), who have shown that the artificial density threshold \u03c10 does not tune the efficiency of star formation, which mostly depends on the turbulence level (e.g. Mach number; see Klessen 2000; Li et al. 2004; Audit & Hennebelle 2010; Renaud, Kraljic & Bournaud 2012).","Citation Text":["Choi, Shlosman & Begelman"],"Functions Text":["The resulting gas density probability distribution function","with an additional few\u2009per\u2009cent of the mass in a power-law tail at high density"],"Functions Label":["Uses","Uses"],"Citation Start End":[[953,978]],"Functions Start End":[[649,708],[872,951]]}
{"Identifier":"2019AandA...631A..35B__Bridges_et_al._(1996)_Instance_2","Paragraph":"The collision velocity dependence of the coefficient of restitution between particles was observed in experiments (Bridges et al. 1996; Higa et al. 1996) and is discussed in the literature (e.g., Ram\u00edrez et al. 1999; Zhang & Vu-Quoc 2002). However, the experiments by Hei\u00dfelmann et al. (2010), used in the present paper to support our assumption of a constant coefficient of restitution, do not see a variation of the coefficient of restitution between particles at low collision velocities (\u2264 1 cm s\u22121). This discrepancy in results might originate in the nature of the collisions studied in these different experiments: Bridges et al. (1996) and Higa et al. (1996) performed collisions of a particle with a flat surface, while Hei\u00dfelmann et al. (2010) observed particle-particle collisions in a free-floating environment. The latter is an experimental environment very similar to NanoRocks. In such inter-particle collisions in free-floating environments, other physical effects lead to a different behavior of the energy dissipation during collisions. In particular, the damping behavior of a large plate or surface is expected to differ from that of a same-sized particle, so that the velocity dependence of the coefficient of restitution might be an effect of the experimental setup in Bridges et al. (1996) and Higa et al. (1996). Colwell et al. (2016) and Brisset et al. (2018) studied collisions between a round cm-sized particle and a flat surface of fine grains. They also observed an increase of the coefficient of restitution with decreasing collision velocity. While the composition of the target surface was different than in Bridges et al. (1996) and Higa et al. (1996) (fine granular material vs. solid ice), the similar behavior of the coefficient of restitution supports the fact that particle-surface collisions are very different from particle-particle collisions, and coefficients of restitution are only velocity dependent for collisions with particles with very different sizes (a much larger particle can be approximated as a target surface).","Citation Text":["Bridges et al. (1996)"],"Functions Text":["This discrepancy in results might originate in the nature of the collisions studied in these different experiments:","and Higa et al. (1996) performed collisions of a particle with a flat surface, while Hei\u00dfelmann et al. (2010) observed particle-particle collisions in a free-floating environment."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[621,642]],"Functions Start End":[[505,620],[643,822]]}
{"Identifier":"2021MNRAS.501.2934C__Dong_&_Fung_2017_Instance_1","Paragraph":"Understanding how the diverse populations of protoplanetary discs in young stellar regions results in the range of exoplanet types and architectures found in the Galaxy is one of the major goals of planet-formation theory. This is an extremely challenging task due in part to the limited observational constraints available. The Atacama Large Millimetre\/submillimetre Array (ALMA) is providing truly transformational images of protoplanetary discs with unprecedented sensitivity and resolution (Andrews 2020). However, millimetre wavelength images reveal the locations of small dust grains but provide little information on the presence of larger particles, beyond centimetre scales. Gas giant planets are mostly made of hydrogen and helium, which ALMA cannot directly observe; therefore, the information on the gas content relies on the observations of less abundant molecules, such as CO and its isotopologues, that are subjected to uncertain depletion processes in the gas-phase (e.g. Miotello et al. 2016). Planets might be detectable by ALMA, although indirectly, by the effects they have on the gas and\/or dust in the disc. When planets become massive enough, they can carve gaps (e.g. Rice et al. 2006; Pinilla, Benisty & Birnstiel 2012; Zhu et al. 2012) and disturb the dynamics of the gas (Teague et al. 2018; Casassus & P\u00e9rez 2019; Pinte et al. 2019). The minimum gap-opening mass depends on the viscosity and scale-height of the disc (Crida, Morbidelli & Masset 2006; Duffell & MacFadyen 2013), but mini-Neptune-mass (P\u00e9rez et al. 2019) or even Earth-mass planets (Rosotti et al. 2016; Dong & Fung 2017) could produce detectable gaps. Gaps consistent with fully formed planets have been imaged by ALMA in discs with estimated ages ranging from  1 Myr (HL Tau and Elias 2\u201324; ALMA Partnership et al. 2015; Cieza et al. 2017) to \u223c10 Myr (TW Hydra; Andrews et al. 2016). However, the origin of these gaps still remains to be established and several alternative explanations have been proposed, including the effect of snow-lines on the dust\/gas evolution of different volatiles (Zhang, Blake & Bergin 2015), magneto-hydrodynamic effects (Flock et al. 2015), secular gravitational instability (e.g. Youdin 2011; Takahashi & Inutsuka 2014), and viscous ring-instabilities (Dullemond & Penzlin 2018). Each one of the proposed mechanisms has their merits and shortcomings, and it is possible that different mechanisms operate together or dominate in different objects or in the same object at at different times. For a recent review on disc (sub)structures, see Andrews (2020). Substructures are also expected to be ubiquitous in protoplanetary discs from a theoretical point of view. Without substructures to halt the migration of mm-size grains at large radii, dust particles should migrate towards the innermost part of the disc in time-scales shorter than 0.1 Myr (e.g. Brauer et al. 2007), which is inconsistent with the observations showing significant mm emission at large radii (\u227310 au) at much older ages. Understanding the origin and evolution of substructures in protoplanetary discs and their implications for planet formation is currently one the main challenges in the field. To better understand the incidence and properties of disc substructures in any given molecular cloud, here we present 1.3 mm\/230 GHz continuum ALMA long-baseline observations at 3\u20135 au resolution of the 10 brightest targets of the \u2018Ophiuchus DIsc Survey Employing ALMA\u2019 (ODISEA) project (Cieza et al. 2019) that were not included in \u2018The disc Substructures at High Angular Resolution Project\u2019 (DSHARP) ALMA Cycle-4 Large Program (Andrews et al. 2018). Our new observations result in the largest sample of disc images at \u223c3\u20135 au resolution in any star-forming region observed so far at mm wavelengths (15 objects when combined with the brightest Ophiuchus objects in DSHARP). In Section 2, we discuss the sample selection, the long-baseline observations, and the data reduction. In Section 3, we characterize the observed substructures, including gaps, rings, inner discs, and cavities. In Section 4, we discuss individual objects and use the full sample of 15 bright Ophiuchus discs observed at high-resolution to construct a tentative evolutionary sequence in which the observed substructures are mostly driven by dust evolution and the formation of giant planets. We also discuss possible connections between the substructures observed in primordial discs and those seen in more evolved debris disc systems. A summary of our results and conclusions is presented in Section 5.","Citation Text":["Dong & Fung 2017"],"Functions Text":["The minimum gap-opening mass depends on the viscosity and scale-height of the disc","or even Earth-mass planets","could produce detectable gaps."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[1597,1613]],"Functions Start End":[[1362,1444],[1548,1574],[1615,1645]]}
{"Identifier":"2020ApJ...889L..10M__McKay_et_al._2018_Instance_1","Paragraph":"As stated earlier, during review of this manuscript Croviser et al. announced in a CBET a tentative water production rate approximately five times larger than our reported value. While the brief nature of the CBET precludes a detailed comparison, we discuss some possible reasons for this discrepancy. At the high airmass of these observations and the small dimensions of the ARCES slit, differential refraction can result in wavelength-dependent slit loss, which can skew flux measurements. However, this is not expected for [O i] 6300 \u212b emission because this feature is close to the guiding wavelength (\u223c5500 \u212b). We confirmed that this is indeed negligible for [O i] 6300 \u212b emission based on observations of comet C\/2012 S1 (ISON) that were performed at a similarly high airmass with ARCES, and found that the production rates derived from the ISON [O i] 6300 \u212b measurements were consistent with values determined using other methods (McKay et al. 2018). Therefore, we do not consider this or other airmass-dependent phenomena as the reason for the discrepancy. At certain geocentric velocities the cometary [O i] 6300 \u212b emission sits on top of a strong telluric absorption, and at high airmass inaccurate removal of this feature can result in a decrease in the measured flux and therefore production rate. This was observed for C\/2012 S1 (ISON) (McKay et al. 2018). However, the geocentric velocity of 2I\/Borisov during our observations was \u223c\u221235 km s\u22121, while the effect on observed [O i] 6300 \u212b line fluxes in comet ISON was only observed at geocentric velocities of \u223c\u221250 km s\u22121. Therefore, this is also not a likely candidate to explain the discrepancy. It is also possible that the activity is highly variable, and we observed Borisov at a minimum in activity, while the Nan\u00e7ay observations, which were coadded over three weeks of observations, provide a long-term average production rate. However, no such variability is observed for CN, with the CN production rate being fairly constant over a several week period (Kareta et al. 2019; Opitom et al. 2019).","Citation Text":["McKay et al. 2018"],"Functions Text":["We confirmed that this is indeed negligible for [O i] 6300 \u212b emission based on observations of comet C\/2012 S1 (ISON) that were performed at a similarly high airmass with ARCES, and found that the production rates derived from the ISON [O i] 6300 \u212b measurements were consistent with values determined using other methods"],"Functions Label":["Similarities"],"Citation Start End":[[937,954]],"Functions Start End":[[615,935]]}
{"Identifier":"2021MNRAS.501...50S__Gupta_et_al._2019_Instance_2","Paragraph":"There have been rather strong claims of AGN QPOs in different bands of the electromagnetic spectrum, ranging from minutes through days through months and years (e.g. Gierli\u0144ski et al. 2008; Lachowicz et al. 2009; Gupta, Srivastava & Wiita 2009; Gupta et al. 2018, 2019; King et al. 2013; Gupta 2014, 2018; Ackermann et al. 2015; Pan et al. 2016; Zhou et al. 2018; Bhatta 2019; and references therein). However, many of the claimed QPOs, particularly those made earlier, were marginal detections (Gupta 2014), lasting only a few cycles, and the originally quoted statistical significances are probably overestimates (Gupta 2014; Covino, Sandrinelli & Treves 2019). Among the better recent claims of QPOs in the gamma-ray band are of \u223c34.5 d in the blazar PKS 2247\u2013131 (Zhou et al. 2018) and of \u223c71 d in the blazar B2 1520+31 (Gupta et al. 2019) found as part of a continuing analysis of blazar Fermi\u2013LAT observations. A recent claim of a \u223c44 d optical band QPO in the narrow-line Seyfert 1 galaxy KIC 9650712 from densely sampled Kepler data has been made by Smith et al. (2018); it was supported by an independent analysis of the same data, indicating a QPO contribution at 52 \u00b1 2 d (Phillipson et al. 2020). Some possibly related QPOs of a few hundred days in two widely separated bands have been reported (Sandrinelli et al. 2016a; Sandrinelli, Covino & Treves 2016b; Sandrinelli et al. 2017). However, an analysis of the Fermi\u2013LAT and aperture photometry light curves by Covino et al. (2019) argued that some multiwaveband QPOs, along with many earlier claims of gamma-ray QPOs, are not significant. Among the gamma-ray QPOs with month-like periods, none showed simultaneous oscillations in a different wavebands. Evidence for related QPOs in multiple wavebands was observed in PG 1553+113, where a QPO was detected in the 0.1\u2013300 GeV and the optical waveband (Ackermann et al. 2015). The observed QPO had a dominant period of \u223c754 d and the source showed strong inter-waveband cross-correlations.","Citation Text":["Gupta et al. 2019"],"Functions Text":["Among the better recent claims of QPOs in the gamma-ray band are of","and of \u223c71 d in the blazar B2 1520+31","found as part of a continuing analysis of blazar Fermi\u2013LAT observations."],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Background"],"Citation Start End":[[825,842]],"Functions Start End":[[664,731],[786,823],[844,916]]}
{"Identifier":"2018ApJ...856...94Z__Bieber_et_al._1991_Instance_2","Paragraph":"Figures 8 and 9 show the effects of solar activity on the CR parallel \u03bb\u2225 (blue line), perpendicular \u03bb\u22a5 (red line), and radial mean free path \u03bbrr (gray line) for a proton with rigidity 445 MV (corresponding to a 100 MeV proton) for the inwardly and outwardly directed IMF, respectively. As described in Zank et al. (1998), the parallel mean free path (mfp) based on standard QLT and assuming magnetostatic turbulence is approximated by\n13\n\n\n\n\n\nwhere \n\n\n\n\n\n, \n\n\n\n\n\n, and \n\n\n\n\n\n. RL is the particle Larmor radius, P is the particle rigidity, and B0 is the mean magnetic field strength. The analytic form of the perpendicular mfp based on NLGC theory is given by (Zank et al. 2004; Shalchi et al. 2010)\n14\n\n\n\n\n\nwhere a2 = 1\/3 is a factor related to the gyrocenter velocity. \n\n\n\n\n\n is a constant such that \u03bd = 5\/6 yields a Kolmogorov (1941) spectrum. Note that Equation (14) was derived under the assumption of a specific form of 2D wave spectrum, which is a constant at large turbulence scales. It means that the 2D turbulence spectrum is independent of wavenumber in the energy range in Equation (14). Observation of magnetic fluctuations in the SW indicates that omnidirectional power spectra approach a k\u22121 wavenumber dependence and that low-frequency turbulence exhibits some sunspot cycle variability (Bieber et al. 1991). Based on this, Engelbrecht & Burger (2015) derived the perpendicular mfp by specifying the energy range spectral index of 2D turbulence power spectra as \u22121. A more general form of the 2D power spectrum with an energy range spectral index q was proposed by Shalchi et al. (2010). They show that the spectral index has a strong influence on the perpendicular diffusion coefficient. In their model, negative values of q correspond to a decreasing spectrum in the energy range, q = 0 corresponds to the constant spectrum we use here, and positive values of q correspond to an increasing spectrum. Matthaeus et al. (2007) presented a similar spectrum in different regimes: energy range, inertial range, and intermediate regime where the spectrum is proportional to k\u22121 to coincide with observations (Bieber et al. 1991; Goldstein & Roberts 1999). However, Shalchi (2013) argues that a spectrum that behaves like k\u22121 does not provide a different perpendicular diffusion coefficient (see also Shalchi et al. 2010), since the field lines for such length scales behave superdiffusively as in the inertial range (Shalchi & Kourakis 2007). In view of this uncertainty, we do not take into account a more elaborate spectrum in the present paper. The behavior of the 2D wave spectrum in the energy range, which may also be correlated with the sunspot cycle, is an important factor in deriving the CR perpendicular mfp. A general form (e.g., Shalchi et al. 2010; Shalchi 2013) should be employed in future studies of CR diffusion.","Citation Text":["Bieber et al. 1991"],"Functions Text":["Matthaeus et al. (2007) presented a similar spectrum in different regimes: energy range, inertial range, and intermediate regime where the spectrum is proportional to k\u22121 to coincide with observations"],"Functions Label":["Background"],"Citation Start End":[[2119,2137]],"Functions Start End":[[1917,2117]]}
{"Identifier":"2019MNRAS.482.5430B__Salmonson_2003_Instance_1","Paragraph":"In light of this, the allowed structure of gamma-ray burst (GRB) jets and the efficiency at which it produces gamma-rays at large angles remains a topic of major importance, and it is useful to consider what types of jet structures are consistent with GRB observations (see also Beniamini et al. 2018b). Previous studies have considered the implications of structure models on the true energetics and rates of GRBs (Frail et al. 2001; Lipunov, Postnov & Prokhorov 2001; Rossi, Lazzati & Rees 2002; Zhang & M\u00e9sz\u00e1ros 2002; Eichler & Levinson 2004; van Eerten & MacFadyen 2012; Pescalli et al. 2015), on the shape of the afterglow light curve (Granot & Kumar 2003; Kumar & Granot 2003; Salmonson 2003) or on detectability of orphan afterglows (Lamb & Kobayashi 2017). Here, we propose a novel way to test the allowed structure of GRBs (in terms of both the energy and Lorentz factor angular distributions), by applying three independent techniques. We focus on long GRBs for which more detailed observations are available. First, we compare the predictions of these models regarding the EX\/E\u03b3 distribution (i.e. the isotropic equivalent early X-ray afterglow to prompt gamma-ray energy ratio) to the observations. We show that a variety of structure models predict large variations in this quantity, in contrast with results from GRB observations. Secondly, we reconsider the effect of the structure on the observed luminosity function and show that a large family of models can be ruled out as they lead to an overproduction of bursts with gamma-ray luminosities below the peak of the observed luminosity function. Both these considerations imply that while the energy angular profile may be steep, the Lorentz factor of GRBs must remain large at any region that produces gamma-rays efficiently. However, even such models typically lead to very peculiar light curves that can be ruled out by observations. The most likely implication is that efficient gamma-ray emission must be confined to a narrow opening angle around the jet\u2019s core, where the isotropic equivalent energy is not much lower than that of the core. This will naturally resolve all the problems mentioned above.","Citation Text":["Salmonson 2003"],"Functions Text":["Previous studies have considered the implications of structure models","on the shape of the afterglow light curve"],"Functions Label":["Background","Background"],"Citation Start End":[[683,697]],"Functions Start End":[[304,373],[598,639]]}
{"Identifier":"2022MNRAS.513.4361M__Vasudevan_&_Fabian_2007_Instance_1","Paragraph":"The photoionization of the disc surface is characterized by two main parameters: irradiating X-ray continuum flux and disc density. Thus the measurement of the disc ionization parameter can address various aspects of the disc\/corona interplay. According to Ballantyne, McDuffie & Rusin (2011), the ionization parameter of a radiation pressure supported disc illuminated by a geometrically thick corona in the SS73 model can be approximated as\n(11)$$\\begin{eqnarray*}\r\n\\xi &amp;\\approx &amp; 4.33\\times 10^{9}\\left(\\frac{\\eta }{0.1}\\right)^{-2}\\left(\\frac{\\alpha }{0.1}\\right)\\left(\\frac{r}{r_{\\rm g}}\\right)^{-\\frac{7}{2}} \\nonumber\\\\\r\n&amp;&amp;\\times \\,f_{\\rm c}(1-f_{\\rm c})^{3}\\Big (\\frac{L_{{\\rm bol}}}{L_{{\\rm E}}}\\Big)^{3}GR_{{\\rm corr}},\r\n\\end{eqnarray*}$$where radiative efficiency \u03b7 \u2248 0.1 (Davis & Laor 2011); viscosity parameter \u03b1 = 0.1 (Shakura & Sunyaev 1973); coronal dissipation fraction fc \u2248 0.45 (Vasudevan & Fabian 2007); r is the disc radius in units of rg; $GR_{{\\rm corr}}=R_{{\\rm R}}^{3}R_{{\\rm z}}^{-2}R_{{\\rm T}}^{-1}$ is a general relativistic correction factor and is solely dependent on the dimensionless black hole spin a (e.g. Novikov & Thorne 1993; Krolik 1999). Therefore, the dependence of disc ionization log\u2009\u03be on the Eddington ratio (Lbol\/LE) is predominantly determined by two parameters: r and a:\n(12)$$\\begin{eqnarray*}\r\n\\log \\xi \\simeq \\varPsi (r,a)\\Big (\\frac{L_{\\rm bol}}{L_{\\rm E}}\\Big).\r\n\\end{eqnarray*}$$Fig. 4 shows the dependence of the disc ionization parameter on the Eddington ratio for the low-mass sample. We show the SS73 model predicted log\u2009\u03be\u2013Lbol\/LE relationships for five different combinations of BH spin and inner disc radius: (a = 0.998, r = 2rg), (a = 0.998, r = 6rg), (a = 0.75, r = 6rg), (a = 0.5, r = 6rg), and (a = 0.25, r = 6rg). We noticed that the inferred ionization states of low-mass AGN discs are consistent with the SS73 model predicted solutions. The derived log\u2009\u03be\u2013Lbol\/LE plane suggests that if the relativistic reflection originated from within 6rg of the inner accretion disc, then the measured ionization parameters of the low-mass sample require spins to be in the range of a \u2208 [0.25, 0.998] with a median value of \u223c0.75.","Citation Text":["Vasudevan & Fabian 2007"],"Functions Text":["coronal dissipation fraction fc \u2248 0.45"],"Functions Label":["Uses"],"Citation Start End":[[914,937]],"Functions Start End":[[874,912]]}
{"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"],"Functions Text":["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","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."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[837,854]],"Functions Start End":[[681,835],[856,1023]]}
{"Identifier":"2016ApJ...819L...7N___2015b_Instance_1","Paragraph":"The gap and ring resemble those in the HL Tau system, recently found by the ALMA long baseline campaign (ALMA Partnership et al. 2015). Our result shows that gaps and rings in the (sub)millimeter dust continuum can exist, not only in relatively young disks (0.1\u20131 Myr) but also in relatively old disks (3\u201310 Myr). One possible mechanism for opening a gap is the gravitational interaction between a planet and the gas (e.g., Lin & Papaloizou 1979; Goldreich & Tremaine 1980; Fung et al. 2014). Such an interaction may also produce the spiral density waves recently found in optical and near-infrared scattered light imaging of dust grains in protoplanetary disks (e.g., Muto et al. 2012). According to recent theoretical analyses of gap structure around a planet (Kanagawa et al. 2015a, 2015b, 2016), the depth and width of the gap are controlled by the planetary mass, the turbulent viscosity, and the gas temperature. The shape of the gap is strongly influenced by angular momentum transfer via turbulent viscosity and\/or instability caused by a steep pressure gradient at the edges of a gap. The observed gap has an apparent width and depth of \n\n\n\n\n\n au and \n\n\n\n\n\n, respectively. This is too shallow and too wide compared with that predicted by theory. However, the observations are limited to an angular resolution of \u223c15 au, and the depth and width could be deeper and narrower in reality. For instance, if we assume that the gap depth times the gap width retains the value derived from the observations, it is possible for the gap to have a width and depth of \n\n\n\n\n\n 6 au and \n\n\n\n\n\n, which is similar to the GPI result (Rapson et al. 2015). Such a gap could be opened by a super-Neptune-mass planet, depending on the parameters of the disk, such as the turbulent viscosity (Kanagawa et al. 2015a, 2015b, 2016). If the gap in the larger dust grains is deeper than that in the gas, the planet could be lighter than super-Neptune mass. We note that a planet of even a few Earth masses, although it cannot open a gap in the gas, can open a gap in the dust distribution if a certain amount of pebble-sized particles, whose motions are not perfectly coupled to that of gas, are scattered by the planet and\/or the spiral density waves excited by the planet (Paardekooper & Mellema 2006; Muto & Inutsuka 2009).","Citation Text":["Kanagawa et al.","2015b"],"Functions Text":["According to recent theoretical analyses of gap structure around a planet","the depth and width of the gap are controlled by the planetary mass, the turbulent viscosity, and the gas temperature."],"Functions Label":["Background","Background"],"Citation Start End":[[763,778],[786,791]],"Functions Start End":[[688,761],[800,918]]}
{"Identifier":"2018MNRAS.478.1884F__Widing_&_Feldman_2001_Instance_1","Paragraph":"Several studies indicate that the solar wind produced by the WTD mechanism has higher AHe and v\u03b1p, while the solar wind may have lower AHe and v\u03b1p when the RLO mechanism is at work. While it is challenging to obtain direct observations of the AHe in the solar corona, first ionization potential (FIP) bias measurements are more readily available. It is found that generally the FIP bias is higher in AR and QS regions (mainly occupied by closed loops) than in CH regions (generally taken up by open magnetic field lines) (Widing & Feldman 2001; Feldman et al. 2005; Brooks & Warren 2011; Baker et al. 2013). It is believed that the reason for the enrichment of the low FIP ions in the corona and the solar wind is that they are ionized earlier in comparison to high FIP elements. The helium has the highest FIP and remains neutral longest. This results in the enrichment\/depletion of low FIP elements\/helium because only ions interact with waves (Laming 2012, 2015, 2017). It means that the helium abundance should be inversely proportional to the low FIP bias elements in the corona. Thus, the AHe is higher in open magnetic field structures and lower in closed loops if the above mechanism is valid. The helium abundance for the fast SW coming from large CHs is higher and remarkably stable (Schwenn 2006). In contrast, Rakowski & Laming (2012) suggested that the helium is depleted in closed loops and the depletion efficiency is higher in larger loops, and lower in smaller loops. Furthermore, from simulations, Laming (2017) showed that the AHe is higher in open magnetic field regions and it is lower in closed loops. Suess et al. (2009) found that the solar wind that comes from big streamers has lower helium abundance. The solar wind can be produced by interchange reconnection in the streamers (Huang et al. 2016). This gives the observational support to the notion that the AHe is lower in the closed loops as streamer structures are composed by very large closed loops.","Citation Text":["Widing & Feldman 2001"],"Functions Text":["It is found that generally the FIP bias is higher in AR and QS regions (mainly occupied by closed loops) than in CH regions (generally taken up by open magnetic field lines)"],"Functions Label":["Background"],"Citation Start End":[[522,543]],"Functions Start End":[[347,520]]}
{"Identifier":"2021MNRAS.508.4512L__Haiman_2017_Instance_1","Paragraph":"The GW sources that LISA will observe at cosmological distances can be used as standard sirens. These include MBHBs, EMRIs, and SOBHBs. Unfortunately, only for the first of these types of sources are EM counterparts plausibly expected to be produced and observed by future EM facilities (Tamanini et al. 2016). MBHBs are in fact expected to emit a large amount of EM radiation in different bands at merger or during long-lasting (\u223c weeks\/months) afterglows (see, e.g. Palenzuela, Lehner & Liebling 2010; Dotti, Sesana & Decarli 2012; Giacomazzo et al. 2012; Moesta et al. 2012), and possibly even through pre-merger signals (Kocsis, Haiman & Menou 2008; Kaplan et al. 2011; O\u2019Shaughnessy et al. 2011; Haiman 2017; Dal Canton et al. 2019). If sufficiently accurate sky localization can be attained from the GW parameter estimation analysis and if the EM counterpart is sufficiently powerful to be spotted by EM telescopes, then we expect to identify the host galaxy of up to a few LISA MBHB mergers per year (Tamanini et al. 2016; Tamanini 2017). These golden sources can then be used as high-redshift standard sirens to map the expansion of the Universe up to z \u223c 10. Although the low number of expected EM counterparts and the high redshift of MBHB mergers are not ideal to test standard cosmological models such as \u039bCDM or to place constraints on late-time dark energy (DE) (Tamanini et al. 2016; Tamanini 2017; Belgacem et al. 2019b), they can efficiently be used to probe deviations from \u039bCDM at earlier cosmological epochs, specifically in the interval 3 \u2272 z \u2272 10 (Caprini & Tamanini 2016; Cai, Tamanini & Yang 2017; Belgacem et al. 2019b; Speri et al. 2020). Standard siren analyses with MBHBs would moreover definitely benefit from a network of space-based detectors, e.g. LISA and Taiji, which would greatly improve the sky location accuracy of each MBHBs and thus provide better chances to spot the EM counterpart (see, e.g. Shuman & Cornish 2021; Wang et al. 2021; Yang 2021).","Citation Text":["Haiman 2017"],"Functions Text":["MBHBs are in fact expected to emit a large amount of EM radiation in different bands","and possibly even through pre-merger signals"],"Functions Label":["Background","Background"],"Citation Start End":[[701,712]],"Functions Start End":[[311,395],[579,623]]}
{"Identifier":"2015MNRAS.453.3414A__Filippenko_&_Chornock_2001_Instance_2","Paragraph":"Filippenko & Chornock (2001) first presented the dynamical estimate of mass of the source to be around 7.4 \u00b1 1.1\u2009M\u2299. Recently, Radhika & Nandi (2014) claimed that the mass of XTE J1859+226 is perhaps in between 6.58 and 8.84\u2009M\u2299 which is similar to the prediction of Shaposhnikov & Titarchuk (2009), although the lower mass limit is estimated as 5.4\u2009M\u2299 by Corral-Santana et al. (2011). However, we consider the typical mass of the source as 7\u2009M\u2299. The distance of this source is around d \u223c 11\u2009kpc (Filippenko & Chornock 2001). Steiner et al. (2013) measured the spin as ak \u223c 0.4; however, Motta et al. (2014b) recently reported that the spin of the source is ak \u223c 0.34. Since the spin predictions are quite close, we use ak \u223c 0.4 for this analysis. We estimate the fluxes Fx (see Table 1) of LHS and HIMS of the 1999 outburst of the source (Radhika & Nandi 2014). The corresponding disc luminosities are calculated as $L_{\\rm disc}^{{\\rm LHS}}=8.26 \\times 10^{37}\\ {\\rm erg\\ s^{-1}}$ and $L_{\\rm disc}^{{\\rm HIMS}}=1.85 \\times 10^{38}\\ {\\rm erg\\ s^{-1}}$, respectively. Now, it is reasonable to assume the accretion efficiency for rotating BH as \u03b7 = 0.3 which corresponds to the accretion rate of the inflowing matter as ${\\dot{M}}_{{\\rm acc}}^{{\\rm LHS}} = 0.304 {\\dot{M}}_{{\\rm Edd}}$ in LHS and ${\\dot{M}}_{{\\rm acc}}^{{\\rm HIMS}} = 0.680 {\\dot{M}}_{{\\rm Edd}}$ in HIMS. For LHS, we use $R_{\\dot{m}}=9.83$\u2009per\u2009cent following our theoretical estimate where xs = 64.6rg for ak = 0.4, ${\\mathcal {E}}=0.001\\,98$ and \u03bb = 3.18. Incorporating these inputs in equation (15), we obtain the jet kinetic power as $L^{{\\rm LHS}}_{{\\rm jet}} = 2.52\\times 10^{37}\\ {\\rm erg\\ s^{-1}}$. The maximum mass outflow rate for HIMS corresponding to ak = 0.4 is obtained from Fig. 9 as $R^{\\rm max}_{\\dot{m}}=17.5$\u2009per\u2009cent for ${\\mathcal {E}}=0.005\\,47$ and \u03bb = 3.1, where the shock transition occurs at 21.9rg. Using these values in equation (15), we obtain the maximum jet kinetic power as $L^{{\\rm HIMS}}_{{\\rm jet}} = 1.08\\times 10^{38}\\ {\\rm erg\\ s^{-1}}$ which we regard to be associated with the HIMS of this source.","Citation Text":["Filippenko & Chornock 2001"],"Functions Text":["The distance of this source is around d \u223c 11\u2009kpc"],"Functions Label":["Uses"],"Citation Start End":[[496,522]],"Functions Start End":[[446,494]]}
{"Identifier":"2016MNRAS.455.1905H__Shields_et_al._2003_Instance_1","Paragraph":"Galaxies hosting active galactic nuclei (AGN) are not only promising candidates for providing direct evidence for such a co-evolution, they can also be used to study the scaling relations over cosmic times. Many studies suggest that BH growth precedes spheroid assembly (Treu, Malkan & Blandford 2004; McLure et al. 2006; Peng et al. 2006a,b; Shields et al. 2006; Woo et al. 2006, 2008; Salviander et al. 2007; Treu et al. 2007; Gu, Chen & Cao 2009; Jahnke et al. 2009; Bennert et al. 2010, 2011; Decarli et al. 2010; Merloni et al. 2010; Wang et al. 2010); however, other studies find no significant evolution with redshift (e.g. Shields et al. 2003; Shen et al. 2008a; Salviander & Shields 2013; Schramm & Silverman 2013; Salviander, Shields & Bonning 2015; Shen et al. 2015; Sun et al. 2015). The likely reason for the disagreement is that the studies are differently affected by intrinsic scatter in the relation, selection effects, and observational biases (see e.g. Lauer et al. 2007; Volonteri & Stark 2011; Schulze & Wisotzki 2014). Another issue is that the MBH\u2013\u03c3* relation is not well constrained even for local AGN host galaxies at the high-mass end. Especially for the AGN appearing as quasi-stellar objects (QSOs), the bright nuclear point source often outshines its host galaxy. Measurements of the spheroid properties are therefore difficult, in particular the stellar velocity dispersion for which high signal-to-noise (S\/N) spectra are needed. Thus, many studies focus on Seyfert galaxies often using aperture spectra which integrates over the central few kpc of the host galaxy (e.g. Greene & Ho 2006; Woo et al. 2006, 2008; Treu et al. 2007; Shen et al. 2008a; Matsuoka et al. 2015). The side effect is that, if present, a significant disc contribution may be included, questioning the definition of the spheroid stellar velocity dispersion in these cases. Bennert et al. (2015) study the effect of different definitions of \u03c3* in the literature for a sample of 66 local Seyfert-1 galaxies and find that it can vary by up to 40 per cent.","Citation Text":["Shields et al. 2003"],"Functions Text":["Many studies suggest that BH growth precedes spheroid assembly","however, other studies find no significant evolution with redshift (e.g."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[631,650]],"Functions Start End":[[207,269],[558,630]]}
{"Identifier":"2022ApJ...936..102A__Williams_et_al._2006_Instance_1","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"],"Functions Text":["Electrostatic solitary waves have been observed in Saturn\u2019s magnetosphere"],"Functions Label":["Background"],"Citation Start End":[[1699,1719]],"Functions Start End":[[1624,1697]]}
{"Identifier":"2021ApJ...907...47L__Lee_et_al._2019_Instance_3","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"],"Functions Text":["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","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."],"Functions Label":["Future Work","Future Work"],"Citation Start End":[[2439,2454]],"Functions Start End":[[2212,2422],[2457,2617]]}
{"Identifier":"2017MNRAS.472.1361B__Salvo_&_Stella_2002_Instance_1","Paragraph":"Historically, the explanation of soft state spectra of NSs demanded the presence of a blackbody emission from the boundary layer of an NS (Mitsuda et al. 1984). For the harder states with a power-law tail in the energy spectrum, the need of Compton scattering became evident (White et al. 1986; Mitsuda et al. 1989). The difference between these two models was that while the former assumed a cooler boundary layer, the latter assumed a hotter one, compared to the accretion disc. Sunyaev and his collaborators (Inogamov & Sunyaev 1999, hereafter IS99; Popham & Sunyaev 2001, hereafter PS01; Gilfanov & Sunyaev 2014 , hereafter GS14) assume that the KD reaches all the way to the NS and is connected with the boundary layer where the thickness increases due to higher temperature. Most of these studies were done to address the soft state spectra of NSs. The state transition of NSs in Low Mass X-ray Binaries (LMXBs), presented another problem. The fact that disc accretion rate was not the single factor that controlled the size or temperature of the Compton cloud, used to model the hard state spectra, lead to the conclusion that some unknown parameter, related to the truncation radius of the disc, is responsible for the hard X-ray tail (Barret 2001; Barret & Olive 2002; Di Salvo & Stella 2002). Paizis et al. (2006) found a systematic positive correlation between the X-ray hard tail and the radio luminosity, inferring that the Compton cloud might serve as the base of radio jets (see Chakrabarti 2016, and references therein). Recent phenomenological works places a transition layer (TL) or Compton cloud between the KD and the boundary layer (Farinelli et al. 2008; Titarchuk, Seifina & Shrader 2014, hereafter TSS14). It has been argued in the past (Chakrabarti 1989; C96; Chakrabarti & Sahu 1997) that while in black hole accretion, passing of the flow through the inner sonic point ensures that the flow becomes sub-Keplerian just outside the horizon, in the case of NSs, the Keplerian flow velocity must slow down to match with the sub-Keplerian surface velocity. Numerical simulations clearly showed that jumping from a KD to a sub-Keplerian disc is mediated by a super-Keplerian region (Chakrabarti & Molteni 1995). In Titarchuk, Lapidus & Muslimov (1998, hereafter TLM98), a super-Keplerian TL was invoked to explain the kHz quasi-periodic oscillations (QPOs) and in TSS14 the TL was expanded several fold to explain the spectral properties. In reality, there are two such layers simultaneously present in an NS accretion: One is similar to the NBOL and the other is similar to the CENBOL in a black hole accretion (CT95). In a black hole accretion, only CENBOL is present. All these approaches clearly point to the existence of a CENBOL type hot electron reservoir, which naturally occurred in black hole accretion, confirming Chakrabarti & Sahu (1997) conclusions that the solutions of the transonic flows are modified only in the last few Schwarzschild radii as per the boundary condition of the gravitating object.","Citation Text":["Di Salvo & Stella 2002"],"Functions Text":["The fact that disc accretion rate was not the single factor that controlled the size or temperature of the Compton cloud, used to model the hard state spectra, lead to the conclusion that some unknown parameter, related to the truncation radius of the disc, is responsible for the hard X-ray tail"],"Functions Label":["Background"],"Citation Start End":[[1278,1300]],"Functions Start End":[[946,1242]]}
{"Identifier":"2018MNRAS.479.3254V___2000_Instance_2","Paragraph":"The lifetime of molecular clouds (MCs) remains an active research topic in the study of the interstellar medium and star formation, and most recent studies, both observational and theoretical, place this lifetime at a few times 107 yr for clouds in the 105\u2013106M\u2299 mass range (e.g. Blitz & Shu 1980; Kawamura et al. 2009; Zamora-Avil\u00e9s, V\u00e1zquez-Semadeni & Col\u00edn 2012; Zamora-Avil\u00e9s & V\u00e1zquez-Semadeni 2014; Lee, Miville-Desch\u00eanes & Murray 2016). In addition, several observational studies have suggested that the star formation rate (SFR) of the clouds appears to increase over their lifetimes. For example, studies of young clusters embedded in moderate-mass MCs (\u223c104M\u2299) (e.g. Palla & Stahler 1999, 2000; Da Rio et al. 2010) have shown that their age histograms contain a large majority of young (1\u20132 Myr) objects, but also a tail of older (up to several Myr) ones suggesting an accelerating star formation activity, sometimes followed by a subsequent decline (see also Povich et al. 2016; Schneider et al. 2018). In addition, Kawamura et al. (2009) reported a clear evolutionary process over the lifetime of giant molecular clouds (GMCs; of masses \u223c105\u2013106M\u2299) in the Large Magellanic Cloud, evidenced by the increasing number of massive stars across the sequence of GMC \u2018classes\u2019 proposed by those authors. Finally, on the basis of the large scatter in the observed star formation efficiency (SFE) in Milky Way GMCs, Lee et al. (2016) have concluded that the SFR in those clouds must also be time variable. Numerical simulations of MC formation and evolution also exhibit time varying, increasing SFRs during their early stages (e.g. V\u00e1zquez-Semadeni et al. 2007; Hartmann, Ballesteros-Paredes & Heitsch 2012). Also, in the presence of stellar feedback, at late times the SFRs reach a maximum and begin to decrease again (e.g. V\u00e1zquez-Semadeni et al. 2010; Col\u00edn, V\u00e1zquez-Semadeni & G\u00f3mez 2013). V\u00e1zquez-Semadeni, Gonz\u00e1lez-Samaniego & Col\u00edn (2017) have recently shown that the simulations of Col\u00edn et al. (2013) in fact produce stellar age histograms highly resemblant of the observed ones (Palla & Stahler 1999, 2000; Da Rio et al. 2010), and reproduce observed radial age gradients in clusters (Getman et al. 2014) as well as bottom-heavy stellar initial mass functions (IMFs) in scattered regions of massive star formation (Povich et al. 2016).","Citation Text":["Palla & Stahler","2000"],"Functions Text":["V\u00e1zquez-Semadeni, Gonz\u00e1lez-Samaniego & Col\u00edn (2017) have recently shown that the simulations of Col\u00edn et al. (2013) in fact produce stellar age histograms highly resemblant of the observed ones"],"Functions Label":["Similarities"],"Citation Start End":[[2092,2107],[2114,2118]],"Functions Start End":[[1897,2090]]}
{"Identifier":"2017AandA...602A..82D__Kirkpatrick_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":["Kirkpatrick et al. 2012"],"Functions Text":["On the other hand, the discovery of isolated L dwarfs in young clusters","and very cold very nearby Y dwarf objects (e.g.,","show that very low-mass isolated brown dwarfs exist and overlap with the planetary masses."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[752,775]],"Functions Start End":[[488,559],[703,751],[790,880]]}
{"Identifier":"2020MNRAS.498.1496M__Dopita_&_Sutherland_1996_Instance_1","Paragraph":"The spectra from Sloan Digital Sky Survey (SDSS) have enabled the detection of the relatively faint He\u2009ii\u2009\u03bb4686 line in large samples of star-forming galaxies (e.g. Shirazi & Brinchmann 2012). These studies find that the observed He\u2009ii\u2009\u03bb4686\/H\u2009\u03b2 intensity ratio does not drop at low metallicities. In fact, recent studies find the ratio to be increasing as the metallicity decreases (Schaerer, Fragos & Izotov 2019). Furthermore, these low-metallicity He\u2009ii\u2009\u03bb4686-emitting galaxies often show only a weak or no evidence of the presence of WR stars (Shirazi & Brinchmann 2012). Thus, questions have been raised on the WR stars as the sole source of ionization of He+ (Plat et al. 2019). Alternative mechanisms such as hard radiation from high-mass stars in binaries (Eldridge et al. 2017), shocks from supernova remnants (Garnett et al. 1991; Dopita & Sutherland 1996), and high-mass X-ray binaries (HMXBs; Schaerer et al. 2019; Kojima et al. 2020) are often invoked. Nearby low-metallicity systems offer an opportunity to address the He+ ionization problem by enabling study of individual star-forming knots. In a detailed study of the metal-poor (Z = 3\u20134 per cent Z\u2299) galaxy SBS\u20090335\u2212052E using MUSE, Kehrig et al. (2018) discard WR stars as the source of ionization and instead propose rotating metal-free stars or a binary population with Z = 10\u22125 and an extremely top-heavy initial mass function (IMF) as the only plausible way of getting around the problem of the ionization budget. In a recent study, Schaerer et al. (2019) find that the observed He\u2009ii\u2009\u03bb4686 intensity in metal-poor star-forming galaxies can be naturally reproduced if the bulk of the He+ ionizing photons is emitted by the HMXB, whose number is found to increase with decreasing metallicity. X-ray binaries in a cluster appear only after the death of the most massive stars, and hence this scenario cannot explain the He+ ionization in young systems [H\u2009\u03b2 equivalent widths (EWs) \u2265 200 \u00c5], as illustrated by Plat et al. (2019).","Citation Text":["Dopita & Sutherland 1996"],"Functions Text":["Alternative mechanisms such as","shocks from supernova remnants","are often invoked."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[842,866]],"Functions Start End":[[686,716],[789,819],[948,966]]}
{"Identifier":"2018MNRAS.478L..18T__Zhang_et_al._2004_Instance_1","Paragraph":"The discovery of GW170817 and its electromagnetic counterparts (GRB170817A and AT2017gfo; Abbott et al. 2017a; Coulter et al. 2017; Goldstein et al. 2017) ushered in a new era of multimessenger astrophysics, in which both gravitational waves and photons provide complementary views of the same source. While observations at optical and infrared wavelengths unveiled the onset and evolution of a radioactive-powered transient, known as kilonova, observations at X-rays and, later, radio wavelengths probed a different component of emission, likely originated by a relativistic outflow launched by the merger remnant. Troja et al. (2017) explained the observed X-ray and radio data as the onset of a standard short GRB (sGRB) afterglow viewed at an angle (off-axis). However, as already noted in Troja et al. (2017) and Kasliwal et al. (2017), a standard top-hat jet model could explain the afterglow data set collected at early times, but failed to account for the observed gamma-ray emission. Based on this evidence, Troja et al. (2017) suggested that a structured jet model (e.g. Zhang et al. 2004; Kathirgamaraju, Barniol Duran & Giannios 2018) provided a coherent description of the entire broad-band data set. Within this framework, the peculiar properties of GRB170817A\/AT2017gfo could be explained, at least in part, by its viewing angle (see also Lamb & Kobayashi 2017; Lazzati et al. 2017). An alternative set of models invoked the ejection of a mildly relativistic wide-angle outflow, either a jet-less fireball (Salafia, Ghisellini & Ghirlanda 2018) or a cocoon (Nagakura et al. 2014; Hallinan et al. 2017). In the latter scenario, the jet might be chocked by the merger ejecta (Mooley et al. 2017), and the observed gamma-rays and broad-band afterglow emission are produced by the expanding cocoon. The cocoon may be energized throughout its expansion by continuous energy injection. In this paper detailed models of structured jet and cocoon, from its simplest to more elaborate version, are compared with the latest radio to X-ray data. Predictions on the late-time evolution are derived, and an unambiguous measurement capable of disentangling the outflow geometry, jet versus cocoon, is presented.","Citation Text":["Zhang et al. 2004"],"Functions Text":["Based on this evidence, Troja et al. (2017) suggested that a structured jet model (e.g.","provided a coherent description of the entire broad-band data set. Within this framework, the peculiar properties of GRB170817A\/AT2017gfo could be explained, at least in part, by its viewing angle"],"Functions Label":["Uses","Uses"],"Citation Start End":[[1081,1098]],"Functions Start End":[[993,1080],[1147,1343]]}
{"Identifier":"2022ApJ...924...42N__Cheng_et_al._1990_Instance_1","Paragraph":"It is generally thought that the emission from radio to medium energy gamma rays is generated by the injected electrons through the synchrotron radiative mechanism. The high-energy photon emission mainly comes from inverse Compton (IC) scattering of the high-energy electrons on the background seed photons, which include the synchrotron background, the cosmic microwave background, and infrared photons in the PWNe (see, e.g., Zhang et al. 2008; Fang & Zhang 2010; Torres et al. 2013; Lu et al. 2020). On the other hand, it is also suggested that the gamma rays could be emitted by the hadronic processes. The relativistic protons accelerated in the Crab pulsar outer gap interact with the matter inside the nebula. and this process may contribute in the high-energy gamma-ray range (see, e.g., Cheng et al. 1990; Bednarek & Protheroe 1997; Khangulyan et al. 2020; Cao et al. 2021). Therefore, it has been long debated whether the high-energy emission from the PWNe is the leptonic or hadronic origin. The details of the high-energy radiation produced by leptonic process have been discussed for the Crab Nebula (see, e.g., Venter & de Jager 2007; Zhang et al. 2008; Mart\u00edn et al. 2012), and that of the gamma-ray emission about the hadronic process have been also investigated (see, e.g., Cheng et al. 1990; Bednarek & Protheroe 1997; Bednarek 2003, 2007). However, with the establishment of more and more high-energy observatories, some telescopes have possessed the performance of observing photons of exceeding to the PeV from the astronomic objects. An increasing number of observational data has been reported by the different experiments. For example, the Amenomori et al. (2019) reported that the Tibet air shower array with the underground water-Cerenkov-type muon detector array observed the highest energy photons of exceeding 100 TeV with a 5.6\u03c3 statistical significance and pointed the measured spectrum with energy extended to the sub-PeV from the Crab Nebula have an absence of high-energy cutoff. Recently, more than 530 photons at energies above 100 TeV and up to 1.4 PeV from the 12 ultra-high-energy gamma-ray sources with a statistical significance greater than seven standard deviations were reported again by LHAASO (Cao et al. 2021). Together with the earlier investigations about the leptonic scenario, the radiative spectrum from the leptons has a cutoff around the sub-PeV region (see, e.g., Zhang et al. 2008; Mart\u00edn et al. 2012; Zhang et al. 2020). It seems that the other components of gamma rays have a significant contribution.","Citation Text":["Cheng et al. 1990"],"Functions Text":["On the other hand, it is also suggested that the gamma rays could be emitted by the hadronic processes. The relativistic protons accelerated in the Crab pulsar outer gap interact with the matter inside the nebula. and this process may contribute in the high-energy gamma-ray range (see, e.g.,"],"Functions Label":["Motivation"],"Citation Start End":[[796,813]],"Functions Start End":[[503,795]]}
{"Identifier":"2021AandA...645A..99C__analysis,_Uttley_et_al._(2011)_Instance_1","Paragraph":"X-ray reverberation in black hole X-ray binaries was first robustly detected in GX 339\u20134 by Uttley et al. (2011) when the source was in its hard state. Previous studies of GX 339\u20134 pointed to the approximate central mass being \u22656\u2006M\u2299 (e.g. Hynes et al. 2003) and a small disc inclination angle (De Marco et al. 2015). Miller et al. (2008) fitted the Suzaku spectra and found that the central black hole has a very high spin, a\u2004\u223c\u20040.998. The X-ray spectroscopic analysis of the hard state spectra from the RXTE archive carried out by Garc\u00eda et al. (2015) suggested the black hole spin to be a\u2004\u223c\u20040.95. Spectral fitting of GX 339\u20134 during its very high flux state using NuSTAR and Swift also suggested a high spin of a\u2004\u223c\u20040.95 (Parker et al. 2016). According to the time-lag analysis, Uttley et al. (2011) found that the disc thermal emission (\u223c0.3\u20130.7 keV, soft band) leads the power-law variations (\u223c0.7\u20131.5 keV, hard band) on long timescales (> 1s). Mahmoud et al. (2019) assumed that the soft component that leads the power-law emission is a soft Comptontized component. Rapisarda et al. (2016) and Rapisarda et al. (2017) instead modelled it as a variable inner region of the thin disc. However, the disc black-body variations lag behind the power-law variations by a few milliseconds on short timescales ( 1s). This switch from low-frequency hard to high-frequency soft lags is thought to be produced by two distinct mechanisms. While the hard lags are likely due to inward propagating fluctuations (e.g. Kotov et al. 2001; Ar\u00e9valo & Uttley 2006), the soft lags can be explained by thermal reverberation associated with the longer light-travel time the hard photons take from the central power-law X-ray source to the disc where they are reprocessed into relatively soft black-body emission. The thermal reverberation lags then provide clues to the geometry of the X-ray source and the inner accretion flow close to the event horizon of the central black hole.","Citation Text":["Uttley et al. (2011)"],"Functions Text":["X-ray reverberation in black hole X-ray binaries was first robustly detected in GX 339\u20134 by","when the source was in its hard state."],"Functions Label":["Background","Background"],"Citation Start End":[[92,112]],"Functions Start End":[[0,91],[113,151]]}
{"Identifier":"2020AandA...635A..60D__Castelli_&_Kurucz_2003_Instance_1","Paragraph":"We further constrain the stellar parameters for HD 85628 using the Virtual Observatory SED Analyzer (VOSA)5 and fitting the H\u03b1 profile of the star\u2019s CHIRON spectrum. We fit the spectral energy distribution (SED; Fig. 8) of HD 85628 using published photometry from Tycho (BT VT; ESA 1997), APASS (BV; Henden et al. 2016), Gaia DR2 (BpRpG; Gaia Collaboration 2018), DENIS (JK; Epchtein et al. 1994), 2MASS (JHKs; Skrutskie et al. 2006), and WISE (W1W2W3W4; Wright et al. 2010), and fit Kurucz ATLAS9 stellar atmosphere models (Castelli & Kurucz 2003). Given that the star is clearly main sequence and relatively young ( 1 Gyr), we constrain the surfacegravity to be within \u00b10.5 dex of log g = 4.0 and metallicity within \u00b10.5 dex of solar, and include the extinction and 1\u03c3 uncertainty presented previously as a constraint, but allow the effective temperature to float. We find the best-fit parameters to be for a Kurucz model with Teff, = 7844\n\n$^{57}_{-285}$\u2212285+57\n K, log g = 4.0, and solar metallicity. We also fitted the shape of the H\u03b1 line from the average high-resolution CHIRON spectra with a grid of Kurucz spectra (as used above). While we fail to meaningfully constrain its metallicity and surface gravity in this way, the spectral line shape is consistent with a temperature of 7700 \u00b1 300 K. Based on these independentanalyses, we adopt Teff \u2243 7800 \u00b1 200 K. We note that this is somewhat cooler than expected given the A3V classification by Houk & Cowley (1975), since typical A3V stars have Teff \u2243 8550 K (Pecaut & Mamajek 2013), however we cannot reconcile such a hot temperature given the available data, leading us to believe that the initial classification was incorrect, and that this star is in fact an A7V star. The best-fit luminosity to SED fit using the VO Sed Analyzer (VOSA) tool, and adopting the distance based on the Gaia DR2 trigonometric parallax, is L = 12.23 \u00b1 0.0655 L\u2299 or log (L\u2215L\u2299) \u2243 1.087 \u00b1 0.023 dex. For the adopted effective temperature, this implies that HD 85628 has a radius of 1.92 \u00b1 0.11 R\u2299 6.","Citation Text":["Castelli & Kurucz 2003"],"Functions Text":["and fit Kurucz ATLAS9 stellar atmosphere models"],"Functions Label":["Uses"],"Citation Start End":[[525,547]],"Functions Start End":[[476,523]]}
{"Identifier":"2017MNRAS.469S.731L__Hearn_et_al._2011_Instance_1","Paragraph":"The optical, spectrocopic and infrared remote imaging system (OSIRIS) scientific imaging cameras on the Rosetta spacecraft have been monitoring the coma activity of comet 67P Churyumov\u2013Gerasimenko (67P hereafter) since their orbital rendezvous in 2014 August (Lara et al. 2015; Lin et al. 2015; Sierks et al. 2015; Lin et al. 2016; Shi et al. 2016; Vincent et al. 2016a,b). The solar heating of the sunlit side of the nucleus surface leads to sublimation of the volatiles and to the formation of dust jets. On 2015 March 12, a small outburst was first detected from a part of the Imhotep region on the night side (Knollenberg et al. 2016). Such mini-outbursts or night-side activities have been observed before at comet 9P\/Tempel 1 by the Deep Impact mission (Farnham et al. 2007, 2013) and comet 103P\/Hartley 2 by the EPOXI mission (A\u2019Hearn et al. 2011; Bruck Syal et al. 2013). Shortly before the close approach to comet 9P\/Tempel 1, the high-resolution camera on the Deep Impact spacecraft found a number of small, well-defined jets whose bases were rooted at the nucleus surface. Some of these, called limb jets, appeared to come from the darker regions and appeared to be associated with the ice patches reported by Sunshine et al. (2006). A later mission of the Stardust\u2013New Exploration of comet Tempel 1 (NExT) imaging of 9P\/Tempel 1 allowed us to connect the jets with cliffs (Farnham et al. 2013). Comet 103P\/Hartley 2 also displayed several narrow jet features emitting from the un-illuminated regions beyond the terminator at the time of the flyby observations (Bruck Syal et al. 2013). Unlike the less certain identification of the source regions on Tempel 1, the source region of the night-side jets of 103P\/Hartley 2 could be clearly traced to some rough surface topography. However, the mechanism for this type of activity is still unknown. Fortunately, unlike the snap shots from the previous flyby observations, the OSIRIS measurements can provide precise information on the timing and location of the outbursts via a time series of high-resolution images. After the first detection in 2015 March, the OSIRIS wide-angle camera (WAC) and narrow-angle camera (NAC) captured another outburst in mid-July of 2015. Since then, many more outbursts from the night-side and sunlit regions have been detected (Feldman et al. 2016; Gr\u00fcn et al. 2016), with most of their source regions located in the Southern hemisphere of comet 67P (Vincent et al. 2016a). The detected outburst events show a variety of morphological features that can be classified into three different types: broad fans, narrow jets and complex plumes. In this work, we investigate the morphology of these events and characterize their physical properties in detail, including the surface brightness profiles, ejected mass and speed if there are two or more sequential images acquired by the same filter in short duration during the time frame of the outburst.","Citation Text":["A\u2019Hearn et al. 2011"],"Functions Text":["Such mini-outbursts or night-side activities have been observed before at","and comet 103P\/Hartley 2 by the EPOXI mission"],"Functions Label":["Background","Background"],"Citation Start End":[[834,853]],"Functions Start End":[[640,713],[787,832]]}
{"Identifier":"2017AandA...601A..72I__Kobayashi_&_Tanaka_2010_Instance_2","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"],"Functions Text":["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"],"Functions Label":["Uses"],"Citation Start End":[[1131,1154]],"Functions Start End":[[888,1129]]}
{"Identifier":"2021MNRAS.502.2859N__Evans_&_Howarth_2008_Instance_2","Paragraph":"It is harder to evaluate the behaviour of the young stellar population along the line of sight, since radial velocity measurements for our sample of Cepheids, needed for a thorough study, do not exist. Given this deficit, we provide only a simplified estimate using radial velocities of OBA-type stars from Evans & Howarth (2008). Since they belong to the same young population, we assume that they have a similar distance distribution and kinematics as the Cepheids. Fig. 16 shows the massive star sample in the plane of the sky. Except for the northernmost region (\u03b4 \u2265 \u221272\u25cb), where no data exist, these stars cover a comparable area to the Cepheids (indicated as grey dots in the figure for comparison). The radial velocities show a distinct and well-known gradient across the SMC with higher velocities in the eastern part (see also fig. 5 of Evans & Howarth 2008). Such a gradient in radial velocity is also present in older (few\u2009Gyr) RGB stars (see fig. 9 of Dobbie et al. 2014). This gradient is commonly attributed to rotation of the SMC. Based on our results obtained for the Cepheids, we propose a different interpretation: this line-of-sight velocity gradient may instead be caused by the fact that the nearest parts of the galaxy, in the region of the SMC Wing, move with a higher radial velocity compared with the main body of the galaxy. Given the additional differences in tangential velocities, these outer parts might be in the process of being stripped from the SMC. Diaz & Bekki (2012) show in their simulations that tidal effects can produce a velocity gradient that is similar to that of a rotating disc. We stress again that this interpretation is based on the assumption that the Cepheid sample and the OBA-type stellar sample trace a similar three-dimensional distribution. For any conclusive answer, radial velocities of the Cepheid stars are required. Such measurements will be provided by the One Thousand and One Magellanic Fields (1001MC) survey (Cioni et al. 2019), which is a consortium survey with the forthcoming multi-object spectrograph 4MOST that will be mounted on the VISTA telescope.","Citation Text":["Evans & Howarth 2008"],"Functions Text":["The radial velocities show a distinct and well-known gradient across the SMC with higher velocities in the eastern part (see also fig. 5 of"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[846,866]],"Functions Start End":[[706,845]]}
{"Identifier":"2017MNRAS.465..248J__Tristram_et_al._2007_Instance_1","Paragraph":"First direct evidence for the geometrical distribution of the dust was found by Jaffe et al. (2004) with the help of the mid-infrared interferometric instrument (MIDI; Leinert et al. 2003) at the Very Large Telescope Interferometer (VLTI). One of the best-studied parsec-sized dust distribution is harboured by the Circinus galaxy, which is the closest (4.2 Mpc, 1 arcsec \u224820 pc; Freeman et al. 1977) and the second brightest Seyfert galaxy in the MIR. High-resolution mid-infrared (MIR) interferometric observations have revealed a two-component morphology in the brightness distribution: (i) a larger scale, elongated component in direction of the ionization cone (along PA \u2248 107\u00b0) with a full-width at half-maximum (FWHM) size of roughly 0.8 \u00d7 1.9 pc, which is responsible for 80 per cent of the total MIR flux and was interpreted either as the directly illuminated funnel walls of the dusty molecular torus or as part of the (filamentary\/clumpy) outflow; (ii) a disc-like component with an FWHM size of roughly 0.2 \u00d7 1.1 pc and elongated along PA \u2248 46\u00b0 (Tristram et al. 2007, 2014). Both components have a dust temperature of roughly 300 K. The disc-like component was modelled by two Gaussian emitters (a highly elongated emitter plus an unresolved source) in the simple model used by Tristram et al. (2014). This was interpreted as signs for a more complex structure as well as for asymmetry in the second component. The orientation and size of the disc-like component coincides with an edge-on warped disc seen in maser emission with the help of very long baseline interferometry traced by H2O masers (Greenhill et al. 2003). It has an outer radius of roughly 0.4 pc. A similar picture emerges for NGC 1068: A hot parsec-scale disc (FWHM \u22481.35 \u00d7 0.45\u2009pc, T \u2248 800 K) was found with similar orientation and extent as the H2O maser disc (Greenhill et al. 1996; Greenhill & Gwinn 1997). It is surrounded by warm dust extended in polar direction with an FWHM \u22483 \u00d7 4\u2009pc and a temperature T \u2248 300 K (Jaffe et al. 2004; Raban et al. 2009). However, in a systematic study of a larger sample of objects a large range of properties of the dust distribution has been found (Burtscher et al. 2013).","Citation Text":["Tristram et al. 2007"],"Functions Text":["High-resolution mid-infrared (MIR) interferometric observations have revealed a two-component morphology in the brightness distribution: (i) a larger scale, elongated component in direction of the ionization cone (along PA \u2248 107\u00b0) with a full-width at half-maximum (FWHM) size of roughly 0.8 \u00d7 1.9 pc, which is responsible for 80 per cent of the total MIR flux and was interpreted either as the directly illuminated funnel walls of the dusty molecular torus or as part of the (filamentary\/clumpy) outflow; (ii) a disc-like component with an FWHM size of roughly 0.2 \u00d7 1.1 pc and elongated along PA \u2248 46\u00b0"],"Functions Label":["Background"],"Citation Start End":[[1058,1078]],"Functions Start End":[[453,1056]]}
{"Identifier":"2016MNRAS.456..512C__Kronberg_et_al._2004_Instance_2","Paragraph":"Extended radio emission in galaxies is associated with both radio jets and lobes and with outflows, seen often as aligned radio sources in the opposite directions with respect to the central compact radio core. Giant radio galaxies (GRG) are extreme cases of this phenomenology with jets and lobes extending on \u223c Mpc scales suggesting that they are either very powerful or very old site for electron acceleration. In this respect, GRGs have a crucial role in the acceleration of cosmic rays over large cosmic scales (e.g. Kronberg et al. 2004), in the feedback mechanism of AGNs into the intergalactic and intracluster medium (e.g. Subrahmanyan et al. 2008) and in the seeding of large-scale magnetic fields in the universe (e.g. Kronberg et al. 2004) and they are excellent sites to determine the total jet\/lobe energetics in AGN-dominated structures (see e.g. Colafrancesco 2008, Colafrancesco & Marchegiani 2011). To date our knowledge of GRGs (see e.g. Ishwara-Chandra & Saikia 1999, 2002; Lara et al. 2001; Machalski, Jamrozy & Zola 2001; Schoenmakers et al. 2001; Kronberg et al. 2004; Saripalli et al. 2005; Malarecki et al. 2013; Butenko et al. 2014) is limited by their sparse numbers and by the difficulty of detecting them over large areas of the sky. Low-frequency radio observations have an enhanced capacity to detect the extended old electron population in these objects (see e.g. the recent Low Frequency Array \u2013 LOFAR \u2013 observation of the GRG UGC095551), but high-frequency radio observations are less efficient in this task due to the steep-spectra of giant radio lobes. In this context these sources will be ideal targets for the next coming deep, wide-field surveys like, e.g. the ATLAS survey of the Australia Telescope Network Facility (ATNF; see Norris et al. 2009) or the Square Kilometre Array (SKA) deep surveys that will have the potential to study their population evolution up to high redshifts and thus clarifying their role on the feedback for the evolution of non-thermal processes in large-scale structures.","Citation Text":["Kronberg et al. 2004"],"Functions Text":["and in the seeding of large-scale magnetic fields in the universe (e.g."],"Functions Label":["Background"],"Citation Start End":[[730,750]],"Functions Start End":[[658,729]]}
{"Identifier":"2018ApJ...853...50F__Bernard_et_al._2015b_Instance_1","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"],"Functions Text":["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"],"Functions Label":["Similarities"],"Citation Start End":[[755,775]],"Functions Start End":[[576,733]]}
{"Identifier":"2021MNRAS.500.5009M__Bono_et_al._2003_Instance_1","Paragraph":"RR Lyrae are old low-mass stars that, during the central helium-burning phase, show mainly radial pulsation while crossing the classical instability strip in the colour\u2013magnitude diagram. From the observational point of view, they represent the most numerous class of pulsating stars in the Milky Way and, being associated with old stellar populations, are typically found in globular cluster and abundant in the Galactic halo and bulge. The investigation of RR Lyrae properties is motivated by their important role both as distance indicators and tracers of old stellar populations. In particular, evolving through the central helium-burning phase, they represent the low-mass, Population II counterparts of Classical Cepheids, as powerful standard candles and calibrators of secondary distance indicators. In particular, they can be safely adopted to infer distances to Galactic globular clusters (see e.g. Coppola et al. 2011; Braga et al. 2016, 2018, and references therein), the Galactic centre (see e.g. Contreras Ramos et al. 2018; Marconi & Minniti 2018; Griv, Gedalin & Jiang 2019), and Milky Way satellite galaxies (see e.g. Coppola et al. 2015; Mart\u00ednez-V\u00e1zquez et al. 2019; Vivas et al. 2019, and references therein). Being associated with old stellar populations, they represent the basis of an alternative Population II distance scale (see e.g. Beaton et al. 2016, to the traditionally adopted Classical Cepheids), more suitable to calibrate secondary distance indicators that are not specifically associated with spiral galaxies (e.g. the globular cluster luminosity function, see Di Criscienzo et al. 2006, and references therein). The properties that make RR Lyrae standard candles are (i) the well-known relation connecting the absolute visual magnitude MV to the metal abundance [Fe\/H] (see e.g. Sandage 1993; Caputo et al. 2000; Cacciari & Clementini 2003; Catelan, Pritzl & Smith 2004; Di Criscienzo, Marconi & Caputo 2004; Federici et al. 2012; Marconi 2012; Marconi et al. 2015, 2018; Muraveva et al. 2018, and references therein); (ii) the period\u2013luminosity relation in the near-infrared (NIR) filters and in particular in the K 2.2 \u03bcm band (see e.g. Longmore et al. 1990; Bono et al. 2003; Dall\u2019Ora et al. 2006; Coppola et al. 2011; Ripepi et al. 2012; Coppola et al. 2015; Marconi et al. 2015; Muraveva et al. 2015; Braga et al. 2018; Marconi et al. 2018, and references therein). In spite of the well-known advantage of using NIR filters (see e.g. Marconi 2012; Coppola et al. 2015, and references therein), in the last decades there has been a debate on the coefficient of the metallicity term of the KB and PL relation (see e.g. Bono et al. 2003; Sollima, Cacciari & Valenti 2006; Marconi et al. 2015, and references therein). On the other hand, it is interesting to note that many recent determinations (see e.g. Sesar et al. 2017; Muraveva et al. 2018) seem to converge towards the predicted coefficient by Marconi et al. (2015), with values in the range 0.16-0.18 mag\u2009dex\u22121. As for the optical bands, our recently developed theoretical scenario (Marconi et al. 2015) showed that, apart from the MV\u2212[Fe\/H] relation that is affected by a number of uncertainties (e.g. a possible non-linearity, the metallicity scale with the associated \u03b1 elements enhancement and helium abundance variations, as well as evolutionary effects, see Caputo et al. 2000; Marconi et al. 2018, for a discussion), the metal-dependent Period\u2013Wesenheit (PW) relations are predicted to be sound tools to infer individual distances. In particular, for the B-, V- band combination, Marconi et al. (2015) demonstrated that the inferred PW relation is independent of metallicity. In order to test this theoretical tool, we need to compare the predicted individual distances with independent reliable distance estimates, for example, the astrometric ones recently obtained by the Gaia satellite (Gaia Collaboration 2016). To this purpose, in this paper we transform the predicted light curves derived for RR Lyrae models with a wide range of chemical compositions (Marconi et al. 2015, 2018) into the Gaia bands, derive the first theoretical PW relations in these filters and apply them to Gaia Data Release 2 Data base (hereinafter Gaia DR2; Gaia Collaboration 2018; Clementini et al. 2019; Ripepi et al. 2019). The organization of the paper is detailed in the following. In Section 2, we summarize the adopted theoretical scenario, while in Section 3 we present the first theoretical light curves in the Gaia filters. From the inferred mean magnitudes and colours, the new theoretical PW relations are derived in Section 4 that also includes a discussion of the effects of variations in the input chemical abundances. In Section 5, the obtained relations are applied to Gaia Galactic RR Lyrae with available periods, parallaxes, and mean magnitudes to infer independent predictions on their individual parallaxes, to be compared with Gaia DR2 results. The conclusions close the paper.","Citation Text":["Bono et al. 2003"],"Functions Text":["The properties that make RR Lyrae standard candles are","the period\u2013luminosity relation in the near-infrared (NIR) filters and in particular in the K 2.2 \u03bcm band (see e.g."],"Functions Label":["Background","Background"],"Citation Start End":[[2197,2213]],"Functions Start End":[[1648,1702],[2060,2174]]}
{"Identifier":"2021AandA...656A..16C__Bruno_&_Carbone_2013_Instance_1","Paragraph":"Investigations of the turbulent nature of solar wind fluctuations have been ongoing for more than half a century (see, e.g., Bruno & Carbone 2016). Advances have been made consistently thanks to the increasingly accurate measurements of several dedicated space mission as well as to the enormous improvement of numerical calculation, new detailed models and theoretical frameworks, and the development of specific data analysis techniques. Nevertheless, the extremely complex nature of the system and the coexistence of multiple actors, scales, and dynamical regimes have led to a number of questions that remain open (Viall & Borovsky 2020). Among these, the very nature of the turbulent cascade of the solar wind flow and its relationship with the small-scale processes still need to be described in full (Tu & Marsch 1995; Bruno & Carbone 2013; Matthaeus & Velli 2011; Chen 2016). Magnetic field fluctuations have been characterized with great detail at magnetohydrodynamic and kinetic scales, for example, through spectral and high-order moments analysis (Tu & Marsch 1995; Bruno & Carbone 2013). The anisotropic nature of magnetic turbulence has also been addressed, and is still being debated, due to the limited access to three-dimensional measurements in space (see, e.g., Horbury et al. 2008, 2012; Sorriso-Valvo et al. 2010; Yordanova et al. 2015; Verdini et al. 2018; Telloni et al. 2019a; Oughton & Matthaeus 2020). Velocity fluctuations have been studied thoroughly (see, e.g., Sorriso-Valvo et al. 1999; Bruno & Carbone 2013), although the kinetic scales still remain quite unexplored for instrumental limitations, most notably in the sampling time resolution. Both the velocity and magnetic field show highly variable turbulence properties, with well developed spectra, strong intermittency (Sorriso-Valvo et al. 1999), anisotropy, and linear third-order moments scaling (Sorriso-Valvo et al. 2007; Carbone et al. 2011). The level of Alfv\u00e9nic fluctuations (mostly but not exclusively found in fast streams, see e.g., D\u2019Amicis et al. 2011; Bruno et al. 2019) are believed to be associated with the state of the turbulence. In particular, solar wind samples containing more Alfv\u00e9nic fluctuations are typically associated with less developed turbulence, as inferred from both shallower spectra and reduced intermittency (see Bruno & Carbone 2013, and references therein). This is consistent with the expectation that uncorrelated Alfv\u00e9nic fluctuations contribute to reduce the nonlinear cascade by sweeping away the interacting structures (Dobrowolny et al. 1980), as also confirmed by the observed anticorrelation between the turbulent energy cascade rate and the cross-helicity (Smith et al. 2009; Marino et al. 2011a,b).","Citation Text":["Bruno & Carbone 2013"],"Functions Text":["Among these, the very nature of the turbulent cascade of the solar wind flow and its relationship with the small-scale processes still need to be described in full"],"Functions Label":["Future Work"],"Citation Start End":[[826,846]],"Functions Start End":[[643,806]]}
{"Identifier":"2021ApJ...917...24Z__Coughlin_et_al._2020b_Instance_1","Paragraph":"Our simulation results show that the median detectable distances of targeted GW events from BH\u2013NS mergers for a single 2nd generation GW detector and a network of such detectors are \u223c300 Mpc and \u223c700 Mpc, respectively (see Table 4). For comparison, Figure 12 shows that the detection rate and detectable distance for HLV (O3) are approximately the same as those for the case when only bKAGRA is running. This is basically consistent with the detection rate and the distance distribution of BH\u2013NS merger candidates detected during LVC O3 (e.g., Anand et al. 2020; Antier et al. 2020b; Coughlin et al. 2020b; Gompertz et al. 2020; Kasliwal et al. 2020). In Section 3, we have shown that the kilonova absolute magnitude at 0.5 days after a BH\u2013NS merger is mainly distributed in the range \u223c\u201310 to \u223c\u201315.5. In view of the fact that the limiting magnitude of the follow-up wide-field survey projects is almost \u227221 mag (e.g., Antier et al. 2020b; Gompertz et al. 2020; Coughlin et al. 2020b; Kasliwal et al. 2020; Wyatt et al. 2020), the maximum detectable distance for BH\u2013NS kilonovae would be \u2272200 Mpc, which can hardly cover the horizon of GW-triggered BH\u2013NS merger events that O3 found (as shown in Figure 10). However, although BH\u2013NS merger kilonovae can hardly be detected for the present search depths, Figures 9 and 11 reveal that there are great opportunities to discover on-axis afterglows associated with sGRBs or orphan afterglows if the BH components have a high-spin distribution. In order to cover the distance range for searching for BH\u2013NS kilonovae for the network of 2nd generation GW detectors as completely as possible, a search limiting magnitude mlimit \u2273 23\u201324 is required as shown in Figure 10. Present survey projects could reach this search limiting magnitude by increasing exposure times and the number of simultaneous exposures. However, the GW candidates during O3 had very large localization areas with an average of thousands of square degrees (Antier et al. 2020b). Increasing exposure times makes it hard for the present survey projects to cover such large localization areas. Therefore, during the HLVK era, we recommend that survey projects may search for jet afterglows after GW triggers with a relatively shallow search limiting magnitude. If BH\u2013NS mergers have a high location precision, a limiting magnitude of mlimit \u2273 23\u201324 can be reached, which gives a higher probability of discovering associated kilonovae.","Citation Text":["Coughlin et al. 2020b"],"Functions Text":["This is basically consistent with the detection rate and the distance distribution of BH\u2013NS merger candidates detected during LVC O3 (e.g.,"],"Functions Label":["Similarities"],"Citation Start End":[[584,605]],"Functions Start End":[[404,543]]}
{"Identifier":"2020AandA...641A..85S__Dallacasa_et_al._2000_Instance_1","Paragraph":"Compact symmetric objects are very powerful (P1.4 GHz > 1025 W Hz\u22121), compact (linear size of 50\u2013100 pc), and young (\u2272104 yr) objects with a rather symmetric radio structure and convex synchrotron radio spectra (Wilkinson et al. 1994; Murgia 2003; Polatidis & Conway 2003). The characterized convex synchrotron radio spectrum peaks at around 100 MHz in the case of compact steep spectrum (CSS) sources, at about 1 GHz in the case of GHz-peaked spectrum (GPS) objects, and up to a few GHz (e.g., 5 GHz) in the case of high-frequency peakers (HFPs; Dallacasa et al. 2000). The turnover is explained as synchrotron self-absorption (SSA) affecting in a small radio-emitting region or the free-free absorption (FFA) by the dense ambient medium. In a \u201cyoung scenario\u201d, as the source grows, the inner region (possibly a tiny radio lobe) expands, and as a result, the turnover frequency moves to lower frequencies. In this scenario, the HFPs are newborn radio sources that develop into extended radio sources (e.g., FR I, FR II) after evolving through GPS and CSS stages. It is possible that the activity is recurrent in at least some sources: there have been observations of faint extended emission around a few GPS sources (e.g., Baum et al. 1990; Stanghellini et al. 1990; Marecki et al. 2003). The extended emission could be a relic of an earlier active period, into which the reborn radio jets are expanding. Another popular explanation for the turnover and compact natures of CSS, GPS, and HFP is the \u201cfrustration\u201d hypothesis. This theory argues that these sources are confined in small spatial scale and high-density environments, and as a consequence the radio emission is frustrated by the abundant nuclear plasma (van Breugel et al. 1984; Peck et al. 1999; Tingay & de Kool 2003; Callingham et al. 2015; Tingay et al. 2015). In addition, An & Baan (2012) also proposed that young sources with strong constant AGN power breaking through the dense inner region of the host galaxy could result in the compact morphology and the turnover properties of CSOs.","Citation Text":["Dallacasa et al. 2000"],"Functions Text":["The characterized convex synchrotron radio spectrum peaks at around 100 MHz in the case of compact steep spectrum (CSS) sources, at about 1 GHz in the case of GHz-peaked spectrum (GPS) objects, and up to a few GHz (e.g., 5 GHz) in the case of high-frequency peakers (HFPs;"],"Functions Label":["Background"],"Citation Start End":[[547,568]],"Functions Start End":[[274,546]]}
{"Identifier":"2019AandA...630A.131M__Chartas_et_al._2009_Instance_1","Paragraph":"Comptonisation Monte Carlo code (MoCA; see Tamborra et al. 2018 for a detailed description of the code) is based on a single photon approach, working in a fully special relativistic scenario. MoCA allows for various and different physical and geometrical conditions of the accretion disc and of the Comptonising corona. In this paper, the corona is assumed to have either a spherical or a slab-like geometry, and to be as extended as the disc, whose radii have been set to be Rout\u2004=\u2004500 rg and Rin\u2004=\u20046 rg, respectively. Even though arguments (e.g. variability, Uttley et al. 2014, and references therein, microlensing Chartas et al. 2009; Morgan et al. 2012 and timing Kara et al. 2016; De Marco et al. 2013) exist that favour a compact corona, we used extended coronae. In fact, as discussed by Marinucci et al. (2019), Comptonised spectra emerging from compact corona (Rout\u2004=\u2004100 rg\u2013Rin\u2004=\u20046) do not deviate significantly from those produced in more extended corona; see their Fig. 3. The adoption of even more compact coronae (Rout\u2004=\u200420 rg, Rin\u2004=\u20046) results only in the need for higher optical depths to recover the same spectral shape for a given temperature. However, in such cases, general relativity (GR) effects are not negligible (see Tamborra et al. 2018, for a detailed discussion on this topic), and the present version of MoCA does not include GR. For the slab-like geometry case, MoCA allows the user to set up the corona height above the accretion disc (set to 10\u2006rg in our simulations). We use synthetic spectra computed assuming the source BH mass and accretion rate to be the same as those of Ark 120 (e.g. Marinucci et al. 2019, and references therein), namely MBH\u2004=\u20041.5\u2005\u00d7\u2005108\u2006M\u2299 and \u1e41 = Lbol\/LEdd = 0.1. For both the slab and spherical hot electron configurations, we simulated the Comptonised spectra using a wide range of values for electron temperature and optical depth: 0.1\u2004 \u2004\u03c4\u2004 \u20047 and 20  kT  200 keV, and in Fig. 1 we show a sample of spectra obtained by MoCA. Moreover, spectra are computed from 0.01 keV up to 700 keV using 1000 logarithmic energy bins, and a Poissonian error accompanies each spectral point. The obtained spectra are averaged over the inclination angle and in Fig. 1 we show some exemplificative spectra normalised at 1 keV accounting for the two geometries considered in this work.","Citation Text":["Chartas et al. 2009"],"Functions Text":["Even though arguments","microlensing","exist that favour a compact corona, we used extended coronae."],"Functions Label":["Differences","Differences","Differences"],"Citation Start End":[[618,637]],"Functions Start End":[[520,541],[605,617],[709,770]]}
{"Identifier":"2022MNRAS.517.4119T___2017_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":["Kundu et al. 2017"],"Functions Text":["Stringent non-detections in radio","observations exclude nearby CSM at high significance."],"Functions Label":["Background","Background"],"Citation Start End":[[960,977]],"Functions Start End":[[884,917],[1032,1085]]}
{"Identifier":"2015AandA...580A.135D__Hunt_et_al._2010_Instance_1","Paragraph":"How does the propagation of radiation and the ISM composition affect ISM observables in low-metallicity galaxies? Addressing this question is important to understand the evolution of low-metallicity galaxies, which undergo more bursty star formation than normal galaxies. Nearby star-forming dwarf galaxies present distinct observational signatures compared to well-studied disk galaxies. Dwarfs are usually metal poor, H\u2009i rich, and molecule poor as a result of large-scale photodissociation (e.g., Kunth & \u00d6stlin 2000; Hunter et al. 2012; Schruba et al. 2012). Mid-IR (MIR) and far-IR (FIR) observations have revealed bright atomic lines from H\u2009ii regions ([S\u2009iii], [Ne\u2009iii], [Ne\u2009ii], [O\u2009iii], etc.) and PDRs ([C\u2009ii], [O\u2009i]) (e.g., Hunter et al. 2001; Madden et al. 2006; Wu et al. 2008; Hunt et al. 2010; Cormier et al. 2015). Their spectral energy distributions (SEDs) are also different from spiral and elliptical galaxies and indicative of altered dust properties, with a relatively low abundance of polycyclic aromatic hydrocarbons (PAHs) and perhaps a different dust composition (e.g., Madden et al. 2006; Galliano et al. 2008; R\u00e9my-Ruyer et al. 2013). It is still unknown, however, whether these differences between dwarf and disk galaxies are the direct result of recent star formation activity shaping the ISM or instead a consequence of the low-metallicity ISM that is independent of star formation activity. To answer this, one needs to observe tracers of the interplay between the ISM and various stages of star formation activity. While there are now a number of important studies available on PDR properties modeling FIR lines on large scales in various extragalactic environments (e.g., Kaufman et al. 2006; Vasta et al. 2010; Graci\u00e1-Carpio et al. 2011; Cormier et al. 2012; Parkin et al. 2013) or in our Galaxy under solar-metallicity conditions (e.g., Cubick et al. 2008; Bernard-Salas et al. 2012, 2015), only a few studies are published on individual extragalactic regions (Mookerjea et al. 2011; Lebouteiller et al. 2012). Of particular interest are dwarf galaxies, where the effect due to radiative feedback is expected to be most significant. The goal of this paper is to investigate how the low-metallicity ISM reacts under the effects of star formation in regions that have undergone different histories. The nearby low-metallicity galaxy NGC\u20094214 provides an excellent environment to perform this experiment because it has well-separated star-forming centers, one hosting a super star cluster, which allows us to study the effects of extreme star-forming conditions on the surrounding ISM. ","Citation Text":["Hunt et al. 2010"],"Functions Text":["Mid-IR (MIR) and far-IR (FIR) observations have revealed bright atomic lines from H\u2009ii regions ([S\u2009iii], [Ne\u2009iii], [Ne\u2009ii], [O\u2009iii], etc.) and PDRs ([C\u2009ii], [O\u2009i]) (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[790,806]],"Functions Start End":[[563,733]]}
{"Identifier":"2018ApJ...854...26L__Tian_2017_Instance_1","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":["Tian 2017"],"Functions Text":["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.,"],"Functions Label":["Background"],"Citation Start End":[[1116,1125]],"Functions Start End":[[694,1022]]}
{"Identifier":"2018ApJ...862....8T__Hull_et_al._2017a_Instance_1","Paragraph":"Interstellar turbulence is supposed to be one of the most important factors to regulate star formation activities. Shocks induced by supersonic turbulence dramatically increase the density and temperature in the post-shock layer and promote the structure formation, such as dense cores and protostars (e.g., Padoan 1995; Padoan & Nordlund 2002; Klessen et al. 2005; Matsumoto et al. 2015a; Inoue et al. 2018). Observational studies have been attempting to search for dense cores that originated from turbulent phenomena. Recent systematic surveys with ALMA found that there are no or a few starless cores with their internal substructures originating from turbulent fragmentation (cf., Padoan & Nordlund 2002; Offner et al. 2010) in each star-forming region (Dunham et al. 2016; Kirk et al. 2017). On the contrary, complex spatial\/velocity gas structures in the protostellar envelopes with the spatial scale from 0.1 pc down to a few tens of astronomical units are also revealed both by the single-dish observations and the interferometric observations (e.g., Tobin et al. 2011; Tokuda et al. 2014, 2016; Maureira et al. 2017). Some multiscale polarization observations and numerical simulations suggest that turbulent motions, rather than magnetic fields, are dynamically important to form complex gas morphologies of protostellar envelopes (Hull et al. 2017a, 2017b). To understand such diversities\/complexities at early stages of star formation, it is important to investigate physical properties of structures originating from the turbulent gas kinematics. In early phases of star formation, the shock waves can be formed by interactions among different density gas with the different velocities because the gas motions, such as infalling\/outflowing gas, are supersonic. Detections of high-J transitions of molecular lines (e.g., CO) excited by magnetohydrodynamic (MHD) shocks have been expected by theoretical modelings (e.g., Pon et al. 2012; Lehmann & Wardle 2016). Recent submillimeter observations (e.g., Shinnaga et al. 2009; van Kempen et al. 2009a, 2009b; Kristensen et al. 2013) have detected warm envelopes seen in high-J CO lines from protostellar sources. However, the origin and heating mechanisms of warm gas in low-mass star-forming dense cores are still under debate. For example, van Kempen et al. (2009b) suggested several origins to produce the high-J CO lines; (1) inner envelope heated by protostellar luminosity, (2) shocked gas in the outflow, and (3) quiescent gas heated by UV photons. Shinnaga et al. (2009) also detected the extended CO (J = 6\u20135, 7\u20136) emission with the size scale of a few thousand astronomical units in our current target, MC27\/L1521F (see also later descriptions in this section). They argued that the warm gas may be coming from shock regions created by interactions between the collapsing envelopes and the internal disk-like materials around the protostar. Although their observations were not able to resolve the internal substructures of the dense core due to the lack of the angular resolution, they pointed out the importance for investigating such warm gas formed in an early phase of star formation to understand the evolution from dense cores to protostars.","Citation Text":["Hull et al. 2017a"],"Functions Text":["Some multiscale polarization observations and numerical simulations suggest that turbulent motions, rather than magnetic fields, are dynamically important to form complex gas morphologies of protostellar envelopes"],"Functions Label":["Background"],"Citation Start End":[[1343,1360]],"Functions Start End":[[1128,1341]]}
{"Identifier":"2022ApJ...939..103R__Markwardt_2009_Instance_1","Paragraph":"We analysed the change as a function of time, by using annual groups of line intensities for each instrumental setup as independent data points d\n\ni\n. The baseline is taken as time t\n0 = 2012.45, which is the mean time point of our original data, 16.31 years after the discovery of the event (Nakano et al. 1996). The total process follows an exponential decline (see Equation (7)). As the total timescale \u03c4\nrec is so much longer than the epoch of our investigation, a free parameter, giving the curvature of the exponential, cannot be derived unambiguously and numerically stably. Thus we use the Taylor series linear approximation here. We derive independent model regression points m\n\ni\n of the type\n2\n\n\n\nmi=c(ti\u2212t0)+n,where\u2211idi\u2212mi2\u03c3i2\u2192Min,\n\nwith c being the average annual change. As mentioned above the errors \u03c3\n\ni\n vary strongly for some lines (see also Table 2). Thus Equation (2) does not resemble the \u03c7\n2 definition where \n\n\n\n\u03c3i\u221dmi\n\n anymore. Thus standard regression algorithms used widely do not apply (York 1966; Giordano & Iavernaro 2021; Lecuna et al. 2020). The regression analysis follows Tellinghuisen (2020). The derived value of n normalizes each of the data sets with respect to the line strengths at t\n0. This is slightly different from using just a weighted mean for the time t\n0, but is numerically more stable against the strong year-on-year variations of the errors, which are seen especially for the helium lines. For the regression the mpfit library (Markwardt 2009) was used, with the variable errors handled according to York (1966) in the implementation of Tellinghuisen (2020). The errors given by the mpfit library for the c parameter represent the statistical error with two parameters for the calculation of the degrees of freedom. However, we are primarily interested in the significance of any slope, rather than the parameter c (respectively, its normalized counterpart C\n\nk\n from Equation (3)) and its potential contribution to the statistical error budget. To derive the significance of the slope, a C program was written to perform a Monte Carlo (MC) simulation. The value n from Equation (2) defines the normalized values \n\n\n\nd\u00afi\u2254di\/n\n\n and \n\n\n\n\u03c3\u00afi\u2254\u03c3i\/n\n\n given in Table 2. Each data point \n\n\n\nd\u00afi\n\n was varied independently 10 million times in agreement with its individual Gaussian error distribution, and a new regression for the normalized change C\n\nk\n \u2200 k \u2208 [1, 107] was calculated with model points M\n\ni\n types as\n3\n\n\n\nMi=Ck(ti\u2212t0)+Ni,where\u2211id\u00afi\u2212Mi2\u03c3\u00afi2\u2192Min.\n\nMoreover, a similar set of parameters from the same MC simulated data points, assuming that there was no change in time-defining model points M, was derived\n4\n\n\n\nMi=Ni,where\u2211id\u00afi\u2212Mi2\u03c3\u00afi2\u2192Min.\n\nThe fractional area of overlap A between the \u201csloping\u201d and \u201cnon-sloping\u201d distribution functions yields the statistical significance (1 \u2212 A) of the slope as a single parameter. This significance is lower than what would be derived from the standard deviation of the inclination (relative change per year) given by the fit with two free parameters. Figure 7 shows an example of such a pair of histograms. A priori this solution of the MC simulation, caused by the wide spread of errors between the individual data points, does not have to be distributed as a Gaussian. Tests indeed showed that it deviates from solutions with large slopes. As the result for our cases only shows very small slopes, there is only a marginal deviation from a Gaussian. As the assumption of a regular standard deviation holds we are able to use the error function erf from the integral of the Gaussian according to Equation (5) to derive the size of the intersect\n5\n\n\n\nA=1\u2212erf\u2212x2,\n\nand from this the level x \u00d7 \u03c3 of significance.","Citation Text":["Markwardt 2009"],"Functions Text":["For the regression the mpfit library","was used, with the variable errors handled according to York (1966) in the implementation of Tellinghuisen (2020)."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1478,1492]],"Functions Start End":[[1440,1476],[1494,1608]]}
{"Identifier":"2017ApJ...839...83W__Orlando_et_al._2016_Instance_1","Paragraph":"A closer look at the known ejecta-dominated SNRs supports the role of surrounding winds and\/or shells, and may provide additional insight into why such remnants are so rare. Several, and perhaps all of them, are ones where the ejecta are expanding into a pre-SN stellar wind from the progenitor, or into a cavity carved out by such winds. The clearest case is for Cas A, where the light-echo spectrum of the actual SN that shows it to have been a Type IIb event (Rest et al. 2008; Krause et al. 2008). These are produced from the collapse of the helium core of a red supergiant that had lost most of its hydrogen envelope before exploding, so the ejecta expand into the stellar wind from the pre-SN star (Chevalier & Oishi 2003; Orlando et al. 2016). The extreme luminosity of NGC4449-1 has been best explained by its expansion into a dense and extensive circumstellar environment produced by winds from its massive progenitor, possibly with additional contributions due to winds from other massive stars in the surrounding dense OB cluster (Milisavljevic & Fesen 2008). For G292.0+1.8 and E0102\u20137219, the fact that fast knots of ejecta are expanding ballistically 2000\u20133000 yr after the explosion requires that they must be expanding into low-density cavities\u2014ones evacuated by pre-SN winds. In both cases, there is an outer shell of X-ray emission where the SN blast wave is interacting with the CSM. As with Cas A, this interaction also leads to the reverse shock that excites the dense fragments of ejecta, producing the optical emission. For both Puppis A and N132D, spectra of the outer radiative filaments show them to be very high in nitrogen; these too are likely overtaking winds enriched by dredged-up nitrogen. These stripped-envelope SNe\u2014types Ib and Ic\u2014are relatively rare compared to their cousins, types II and IIL, that explode with their envelopes more or less intact; the recent review by Smartt (2009) indicates that together SN Ib and Ic comprise only \u223c20%\u201330% of core-collapse events.","Citation Text":["Orlando et al. 2016"],"Functions Text":["These are produced from the collapse of the helium core of a red supergiant that had lost most of its hydrogen envelope before exploding, so the ejecta expand into the stellar wind from the pre-SN star"],"Functions Label":["Background"],"Citation Start End":[[729,748]],"Functions Start End":[[502,703]]}
{"Identifier":"2017ApJ...835L...1W__Peixoto_&_Oort_1992_Instance_1","Paragraph":"In this work, we investigate how the climate of an Earth-like planet is affected when its orbit is perturbed by the presence of a nearby giant planet. For the first time, a GCM coupled with analytical equations that describe the orbital evolution of a terrestrial planet are used. An additional major difference between our work and previous studies is that we utilize a fully coupled ocean model and an Earth continental layout. This is in contrast to WP2002 who used a 50 m \u201cthermodynamic slab\u201d ocean model without horizontal ocean heat transport or Linsenmeier et al. (2015), who used an aquaplanet model and a 50 m slab ocean, but again with no horizontal ocean heat transport. We use a fully coupled ocean model because alongside atmospheric heat transport, ocean heat transport plays a vital role in the climate of Earth (Peixoto & Oort 1992). In particular, the work of Hu & Yang (2014) has shown that the effects of a fully coupled ocean versus a shallow slab ocean can be significant when looking at synchronously rotating worlds around M-dwarf stars. Godolt et al. (2015) demonstrated stark differences for planets orbiting F-type stars when changing ocean heat transport, while Rose (2015) nicely demonstrated the climatic effects of changing ocean heat transport equations for aqua-type and ridge-type worlds. The downside of a fully coupled ocean approach is that it can take hundreds of model years for a fully coupled ocean to come into equilibrium with the atmosphere. Yet, it will provide a more accurate picture of the climate of the world being modeled. In this study we focus on the effects of the terrestrial planet\u2019s orbital eccentricity on the planet\u2019s climate, which is an under-researched area in 3D GCM studies. At the same time, we keep \u03b8p = 235, following for modern Earth. The latter is a necessary requirement for comparing with past and future work in the literature since obliquity plays such an important role in the possible climate states of terrestrial planets.","Citation Text":["Peixoto & Oort 1992"],"Functions Text":["We use a fully coupled ocean model because alongside atmospheric heat transport, ocean heat transport plays a vital role in the climate of Earth"],"Functions Label":["Uses"],"Citation Start End":[[828,847]],"Functions Start End":[[682,826]]}
{"Identifier":"2022ApJ...936..102A__Williams_et_al._2006_Instance_4","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"],"Functions Text":["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"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[3638,3658]],"Functions Start End":[[3370,3636]]}
{"Identifier":"2021MNRAS.500.2564O__Ziurys_1987_Instance_1","Paragraph":"Phosphorus, which is isoelectronic with both nitrogen and the CH group, is a key element of living systems as a major component of nucleic acids and nucleotides, performing many relevant biochemical functions (Pasek & Lauretta 2005). Conversely, despite significant progress in the recent years (Jim\u00e9nez-Serra et al. 2018; Rivilla et al. 2018; Fontani et al. 2019; Chantzos et al. 2020) and its importance to astrobiology, the interstellar chemistry of phosphorus is still poorly understood. The first identified phosphorus compound in the ISM was PN through observations of its rotational transitions in Orion-KL, W51M, and Sgr B2 (Turner & Bally 1987; Ziurys 1987). This observation was followed by the detection of free CP in the carbon star envelope IRC +10216 (Guelin et al. 1990). The other phosphorus-bearing molecules detected so far are HCP (Ag\u00fandez, Cernicharo & Gu\u00e9lin 2007), CCP (Halfen, Clouthier & Ziurys 2008), PO (Tenenbaum, Woolf & Ziurys 2007), and PH3 (Ag\u00fandez et al. 2014b), while a tentative detection of NCCP is also reported (Ag\u00fandez, Cernicharo & Gu\u00e9lin 2014a). Phosphorus was also detected in the coma of the comet 67P\/Churyumov\u2013Gerasimenko (Altwegg et al. 2016), with recent analysis indicating PO as its main carrier (Rivilla et al. 2020). In addition, PH3 was recently detected in the cloud decks of Venus through millimetre waveband observations (Greaves et al. 2020). Larger phosphorus-bearing molecules, such as phosphorus oxoacids (Turner et al. 2018a) and alkylphosphonic acids (Turner et al. 2018b) \u2013 the latter detected in the Murchison meteorite (Cooper, Onwo & Cronin 1992) \u2013 were identified in PH3-doped interstellar ice analogues exposed to ionizing radiation. These results, in combination with the increasing literature on phosphorus-bearing molecules in circumstellar envelopes (Ziurys, Schmidt & Bernal 2018) and solar-type star-forming regions (Lefloch et al. 2016) and protostars (Lefloch et al. 2018; Bergner et al. 2019), suggest that their role in the ISM might be greater than previously thought.","Citation Text":["Ziurys 1987"],"Functions Text":["The first identified phosphorus compound in the ISM was PN through observations of its rotational transitions in Orion-KL, W51M, and Sgr B2"],"Functions Label":["Background"],"Citation Start End":[[654,665]],"Functions Start End":[[492,631]]}
{"Identifier":"2018ApJ...866...15N__Collet_et_al._2007_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":["Collet et al. 2007"],"Functions Text":["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"],"Functions Label":["Similarities"],"Citation Start End":[[1010,1028]],"Functions Start End":[[820,1008]]}
{"Identifier":"2016ApJ...823...20K__Xue_et_al._2011_Instance_1","Paragraph":"Ouchi et al. (2008) find that there is a possible excess of the Ly\u03b1 LFs at z = 3.1 and 3.7 similar to the bright-end hump, and they claim that 100% of LAEs host AGNs at the bright ends of \n\n\n\n\n\n and 43.4 erg s\u22121, respectively, based on the large-area LAE survey with the multiwavelength data set. Thus, the bright-end hump of our z = 2.2 Ly\u03b1 LF may be produced by AGNs. To examine whether our LAEs at the bright end include AGNs, we use the multiwavelength data of X-ray, UV, and radio available in the SXDS, COSMOS, CDFS, HDFN, and SSA22 fields. For the X-ray data, we use the XMM-Newton source catalog in the SXDS field (Ueda et al. 2008), the Chandra1.8 Ms catalog in the COSMOS field (Elvis et al. 2009), the Chandra 4 Ms source catalog in the CDFS field (Xue et al. 2011), and the Chandra 2 Ms catalog in the HDFN field (Alexander et al. 2003). The typical sensitivity limits of these X-ray data are \u223c10\u221216 to 10\u221215 erg cm \u22122 s\u22121 for the SXDS and COSMOS fields and \u223c10\u221217 to 10\u221216 erg cm \u22122 s\u22121 for the CDFS and HDFN fields. We use GALEX far-UV (FUV) and near-UV (NUV) images for the UV data and obtain these images from the Multimission Archive at STScI (see also Zamojski et al. 2007 for the COSMOS field). The GALEX images reach the 3\u03c3 detection limit of \u223c25\u201326 mag. The Very Large Array 1.4 GHz source catalogs of Simpson et al. (2006) (SXDS), Schinnerer et al. (2007) (COSMOS), and Miller et al. (2013) (CDFS) are used for the radio data. These radio data reach an rms noise level of \u223c10 \u03bcJy beam\u22121. We find that a majority of our bright LAEs are detected in the multiwavelength data, and we summarize the numbers of these LAEs in Table 2. Under the column of \u201cculled sample\u201d in Table 2, we show the numbers of LAEs with no counterpart detection(s) in the X-ray, UV, and radio data. As shown in Table 2, the SXDS and COSMOS fields have the data that cover all of the X-ray, UV, and radio wavelengths. Moreover, the X-ray, UV, and radio data spatially cover the entire fields of SXDS and COSMOS with the similar sensitivities. We make a subsample that is composed of all 1576 LAEs found in the SXDS and COSMOS fields, and we refer to this subsample as SXDS+COSMOS\/All. We then make another subsample consisting of 1538 LAEs with no multiwavelength counterpart detection(s) in the SXDS and COSMOS fields, which is dubbed SXDS+COSMOS\/Culled.","Citation Text":["Xue et al. 2011"],"Functions Text":["For the X-ray data, we use","the Chandra 4 Ms source catalog in the CDFS field","The typical sensitivity limits of these X-ray data are","\u223c10\u221217 to 10\u221216 erg cm \u22122 s\u22121 for the CDFS and HDFN fields."],"Functions Label":["Uses","Uses","Uses","Uses"],"Citation Start End":[[760,775]],"Functions Start End":[[547,573],[709,758],[850,904],[970,1029]]}
{"Identifier":"2019MNRAS.485.4841R__Creminelli_et_al._2010_Instance_1","Paragraph":"Although, the standard form of Press\u2013Schechter mass function with $f(\\nu)=\\sqrt{{2}\/{\\pi }} \\nu \\mathrm{ e}^{-\\frac{\\nu }{2}}$ which discussed in Press & Schechter (1974) and Bond et al. (1991) can provide a good approximation of the predicted number density of haloes, it fails by predicting approximation too many low-mass haloes and too few high-mass ones (Sheth & Tormen 1999, 2002; Lima & Marassi 2004). Thus, in this study we apply another well-known fitting formula which first proposed in Sheth & Tormen (1999): \n(24)\r\n\\begin{eqnarray*}\r\nf(\\nu)=0.2709\\sqrt{\\dfrac{2}{\\pi }}(1+1.1096\\nu ^{0.6})\\mathrm{ exp}(-\\dfrac{0.707 \\nu ^2}{2})\\,\\, .\r\n\\end{eqnarray*}\r\nIn a Gaussian density field, \u03c3 is given by \n(25)\r\n\\begin{eqnarray*}\r\n\\sigma ^2(R)=\\dfrac{1}{2 \\pi ^2}{\\int _0}^\\infty k^2 P(k) W^2(kR) \\, {\\rm d}k\\,\\, ,\r\n\\end{eqnarray*}\r\nwhere R = (3M\/4\u03c0\u03c1m0)1\/3 is the radius of the spherical overdense region, W(kR) is the Fourier transform of a spherical top-hat profile with radius R and P(k) is the linear power spectrum of density fluctuations (Peebles 1993). To obtain the value of \u03c3, we follow the procedure presented in Abramo et al. (2007a). Following on Ade et al. (2016), we use the normalization of matter power spectrum \u03c38 = 0.815 for \u039bCDM cosmology. The number density of virialized haloes above a certain value of mass M at zc, the collapse redshift obtained by \n(26)\r\n\\begin{eqnarray*}\r\nN(\\: M,z)={\\int _0}^\\infty \\dfrac{{\\rm d}n(z)}{{\\rm d}M^{\\prime }}\\, {\\rm d}M^{\\prime }\\,\\,. \r\n\\end{eqnarray*}\r\nThe above limit of integration in equation (26) is $M=10^{18}\\, \\mathrm{ M}_{\\rm \\odot}\\, \\mathrm{ h}^{-1}$ which such gigantic structures could not in practice be observed. Now we can calculate the number density of virialized haloes in both homogeneous and clustered DE scenarios using equations (23) and (26). In this way the total mass of a halo is equal to the mass of pressureless matter perturbations. However, the virialization of dark matter perturbations in the non-linear regime cannot be independent from the properties of DE (Lahav et al. 1991; Maor & Lahav 2005; Creminelli et al. 2010; Basse, Bjlde & Wong 2011). Thus, in clustered DE scenarios, we should consider the contribution of perturbated DE components to the total mass of the haloes (Creminelli et al. 2010; Basse et al. 2011; Batista & Pace 2013). Based on the behaviour of wde(z), DE can reduce or enhance the total mass of the virialized halo. One can obtain \u03f5(z), the ratio of DE mass to be taken into account with respect to the mass of dark matter, from: \n(27)\r\n\\begin{eqnarray*}\r\n\\epsilon (z)=\\dfrac{m_{\\rm DE}}{m_{\\rm DM}}\\,\\, ,\r\n\\end{eqnarray*}\r\nwhere the value of mDE depends on what we consider as the mass of DE component. When one only considers the contribution of the perturbations of DE, the mDE takes the form \n(28)\r\n\\begin{eqnarray*}\r\n{m_{\\rm DE}}^{\\mathrm{ Perturbed}}=4 \\pi \\bar{\\rho }_{\\rm DE}{\\int _0}^{R_{\\rm vir}} \\, {\\rm d}R R^2 \\delta _{\\rm DE}(1+3{c_{\\rm eff}}^2)\\,\\,. \r\n\\end{eqnarray*}\r\nIn the other hand, if we assume both DE contributions of perturbation and background level, the total mass of DE in virialized haloes takes this new form\n(29)\r\n\\begin{eqnarray*}\r\n{m_{\\rm DE}}^{\\mathrm{ Total}}=4 \\pi \\bar{\\rho }_{\\rm DE}{\\int _0}^{R_{\\rm vir}} {\\rm d}R R^2 [(1+3 w_{\\rm DE})+ \\delta _{\\rm DE}(1+3{c_{\\rm eff}}^2)]. \r\n\\end{eqnarray*}\r\nThe quantities inside a spherical collapsing region in the framework of the top-hat profile, evolve only with cosmic time. Thus from equation (28) one can find \n(30)\r\n\\begin{eqnarray*}\r\n\\epsilon (z)=\\dfrac{\\Omega _{\\rm DE}}{\\Omega _{\\rm DM}}\\dfrac{\\delta _{\\rm DE}}{1+\\delta _{\\rm DM}}\\,\\, \r\n\\end{eqnarray*}\r\nand from equation (29) we can obtain \n(31)\r\n\\begin{eqnarray*}\r\n\\epsilon (z)=\\dfrac{\\Omega _{\\rm DE}}{\\Omega _{\\rm DM}}\\dfrac{1+3 w_{\\rm DE}+\\delta _{\\rm DE}}{1+\\delta _{\\rm DM}}\\,\\, .\r\n\\end{eqnarray*}\r\nThe mass of dark matter also is obtained from (see also Batista & Pace 2013): \n(32)\r\n\\begin{eqnarray*}\r\n{m_{\\rm DM}}=4 \\pi \\bar{\\rho }_{\\rm DM}{\\int _0}^{R_{\\rm vir}} \\, {\\rm d}R R^2 (1+ \\delta _{\\rm DM})\\,\\,. \r\n\\end{eqnarray*}\r\nIn Fig. 5 we plot the evolution of \u03f5(z) using equation (30) as the definition of DE mass. We observe that, at high redshift, where the role of DE is less important, \u03f5 for both of parametrizations becomes negligible. This parameter has a greater value in the case of parametrization (2).","Citation Text":["Creminelli et al. 2010"],"Functions Text":["However, the virialization of dark matter perturbations in the non-linear regime cannot be independent from the properties of DE"],"Functions Label":["Uses"],"Citation Start End":[[2090,2112]],"Functions Start End":[[1922,2050]]}
{"Identifier":"2016ApJ...829..120A__Perron_et_al._1988_Instance_1","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"],"Functions Text":["This correlates with the analysis","which demonstrated that the preliminary track annealing led to unpredictable changes in track lengths, resulting in a lower accuracy of nuclear charge determination."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1987,2005]],"Functions Start End":[[1952,1985],[2008,2173]]}
{"Identifier":"2022AandA...666A..95H__Hartman_et_al._2022_Instance_2","Paragraph":"Scaling relations for the core radii rc, core densities \u03b4c, and core masses Mc as functions of the total halo mass M200 were fitted to the simulated halo populations, which largely agree with hydrostatic considerations of the halo cores where rc is nearly constant, as well as velocity dispersion tracing in the halo envelope, \n\n\n\n\nv\nc\n2\n\n\u2248\nv\n\n200\n2\n\n\n\n$ v^2_{\\rm c} \\approx v^2_{200} $\n\n\n\n. However, these trends do not agree with those obtained by fitting the Burkert profile to nearby galaxies in the SPARC dataset and the classical Milky Way dSphs. This poses an issue for SIBEC-DM with Rc\u2004\u2273\u20041 kpc and a largely CDM-like matter power spectrum at late times (Harko 2011; Harko & Mocanu 2012; Velten & Wamba 2012; Freitas & Gon\u00e7alves 2013; Bettoni et al. 2014; de Freitas & Velten 2015; Hartman et al. 2022), as was used in our simulations, although these scenarios are not well-motivated. For SIBEC-DM given by the field Lagrangian in Eq. (1), the self-interaction is constrained to Rc\u2004\u20041 kpc, otherwise an early radiation-like period and a large comoving Jeans\u2019 length washes out too much structure to be consistent with observations (Shapiro et al. 2021; Hartman et al. 2022). In fact, Shapiro et al. (2021) found by using constraints on FDM as a proxy for SIBEC-DM, and matching their transfer function cut-offs and HMFs, that the SIBEC-DM self-interaction should be as low as Rc\u2004\u223c\u200410 pc to not be in conflict with observations. We were unable to probe SIBEC-DM with initial conditions and parameters consistent with the Lagrangian in Eq. (1), since the large gap between the halo cores and the cut-off scale requires both a large simulation box and very high spatial resolution. It should be noted that our SIBEC-DM-only simulations do provide a better agreement with the slopes in observed scaling relations than FDM. In particular, FDM simulations generally find \n\n\n\n\nM\nc\n\n\u223c\n\nM\n\n200\n\n\u03b3\n\n\n\n$ M_{\\mathrm{c}}\\sim M_{200}^{\\gamma} $\n\n\n with 1\/3\u2004\u2004\u03b3\u2004\u20040.6, while we find \u03b3\u2004\u2248\u20040.75, which is closer to the observed \u03b3\u2004\u2248\u20041.1. Additionally, FDM halos have core radii that generally decrease with the halo mass, while we find a slightly increasing trend due to larger halos experiencing more thermal heating, although not as steep as in the SPARC dataset and the Milky Way dSphs.","Citation Text":["Hartman et al. 2022"],"Functions Text":["For SIBEC-DM given by the field Lagrangian in Eq. (1), the self-interaction is constrained to Rc\u2004\u20041 kpc, otherwise an early radiation-like period and a large comoving Jeans\u2019 length washes out too much structure to be consistent with observations"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1160,1179]],"Functions Start End":[[892,1137]]}
{"Identifier":"2016ApJ...826...54D__Schlickeiser_1984_Instance_1","Paragraph":"In both models, the particles interact with magnetohydrodynamic (MHD) Alfv\u00e9n waves in the plasma. If the Doppler-shifted wave frequency is a constant multiple of the particle gyrofrequency in the particle guiding center frame, then a resonant interaction between the particle and the transverse component of the electric field of the MHD wave will occur (Dermer et al. 1996; Becker et al. 2006; Dermer & Menon 2009). The particle will experience either an accelerating or decelerating electric field in the transverse direction of motion over a fraction of the cyclotron period, resulting in an increase or decrease in energy. The accelerating or decelerating electric field causes the particle distributions to diffuse in energy, pushing particles to higher or lower energies in a diffusion pattern. This stochastic acceleration process typically causes the particle distributions to have a pronounced curvature in the energy spectrum (Schlickeiser 1984). The strength of the particle diffusion depends on the spectral index of the MHD turbulence, p. A Kolmogorov, p = 5\/3, or a Kraichnan, p = 3\/2, spectrum are most often used to model MHD turbulence. In this study, we restrict the spectral index of the turbulence to p = 2 to simulate hard sphere scattering between the MHD waves and the particle spectra. The stochastic acceleration timescale can be expressed as\n3\n\n\n\n\n\nwhere tdyn represents the dynamical timescale over the region in which turbulence generated (which may be smaller than the entire emission region), \u03b2A represents the Alfv\u00e9n velocity of the plasma normalized to the speed of light and \u03bei represents the ratio of the magnetic field fluctuations relative to the background magnetic field, \n\n\n\n\n\n. The stochastic acceleration timescale is independent of particle mass and will therefore be the same for all charged-particle species (protons, electrons\/positrons, pions, and muons). The diffusion term in Equation (1) describes the stochastic acceleration of particles in the quasi-linear approximation (Dermer et al. 1996). For gyro-resonant interactions to occur in the quasi-linear regime, the magnetic field fluctuations must be much smaller than the background magnetic field, \n\n\n\n\n\n. If the energy density in the plasma waves starts to approach the energy density of the magnetic field, then the field becomes disordered and there exists no well defined gyrofrequency. In both models, we use a ratio between the acceleration timescale and the escape timescale as an input parameter. The ratio between the acceleration and escape timescales constrain the maximum size in which turbulence is injected for stochastic acceleration to occur in the quasi-linear regime; see Section 5.","Citation Text":["Schlickeiser 1984"],"Functions Text":["This stochastic acceleration process typically causes the particle distributions to have a pronounced curvature in the energy spectrum"],"Functions Label":["Uses"],"Citation Start End":[[937,954]],"Functions Start End":[[801,935]]}
{"Identifier":"2018MNRAS.478.3890B__Heckman_et_al._2017_Instance_5","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"],"Functions Text":["see fig. 5 from"],"Functions Label":["Uses"],"Citation Start End":[[1212,1231]],"Functions Start End":[[1196,1211]]}
{"Identifier":"2020AandA...643A.128K__Younes_et_al._2015_Instance_1","Paragraph":"In order to allow us to compare the continuum shape with earlier analyses, we use phenomenological continuum models rather than more physically motivated models such as those by Becker & Wolff (2007)2 or Farinelli et al. (2016). As discussed by M\u00fcller et al. (2013), among others, phenomenological spectral models typically used to describe the continua of accreting neutron stars are the exponentially cutoff power law (cutoffpl), the power law with Fermi-Dirac cutoff (FDcut, Tanaka 1986), a negative-positive cutoff power law (NPEX, Mihara 1995), and a model consisting of a blackbody disk (diskbb, Mitsuda et al. 1984) and thermally comptonized continuum (nthcomp, Zdziarski et al. 1996; \u017bycki et al. 1999). The residuals of the cutoffpl, FDcut, NPEX, and diskbb+nthcomp models are shown in Fig. 2, and the best-fit parameters are given in Table 1. The NPEX and FDcut residuals look very similar because they are driven to parameters that effectively mimic the cutoffpl solution. All tested continuum models describe the data similarly well. Due to its simplicity and in order to allow comparison with previous work (e.g., Younes et al. 2015), we used the cutoffpl model for all subsequent analysis. Photoelectric absorption in the interstellar medium is accounted for with the tbnew model (TBabs in XSPEC) with cross sections and abundances according to Verner et al. (1996) and Wilms et al. (2000), respectively. The iron fluorescence line complex can formally be described by a slightly broadened (\n\n\n\n\u03c3\n=\n0\n.\n\n23\n\n\u2212\n0.04\n\n\n+\n0.05\n\n\n\n\n$ \\sigma=0.23^{+0.05}_{-0.04} $\n\n\n keV) Gaussian component at 6.59\u2005\u00b1\u20050.04 keV. This is most likely a blend of different ionization states that cannot be resolved with NuSTAR. The strongest fluorescence lines are often produced by neutral (6.4 keV), He-like (6.7 keV), and H-like iron (7.0 keV), and the structure seen in the data is also consistent with a set of narrow K\u03b1 lines from these ions, as well as neutral K\u03b2 (7.1 keV) with a K\u03b2\/K\u03b1 flux ratio of 13% (Palmeri et al. 2003). With fixed energies and widths, this approach is also statistically valid and has the same degrees of freedom as using one broad emission feature, but shows less interference with the continuum modeling because all line energies and widths are fixed and broadening is only due to the detector response. Using both approaches, slight residuals still remain at the iron K edge. These residuals are due to a combination of a gain-shift in NuSTAR energy calibration and the fact that the tbnew model only includes neutral iron.","Citation Text":["Younes et al. 2015"],"Functions Text":["Due to its simplicity and in order to allow comparison with previous work (e.g.,","), we used the cutoffpl model for all subsequent analysis."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1127,1145]],"Functions Start End":[[1046,1126],[1145,1203]]}
{"Identifier":"2018MNRAS.478.4357S__Peebles_&_Ratra_2003_Instance_1","Paragraph":"Observations over the years seem to firmly support the current acceleration of the Universe and therefore the possible existence of a generic cause responsible for it which we call dark energy (DE; see e.g. Riess et al. 1998; Perlmutter et al. 1999; WMAP collaboration 2013; Planck collaboration XVI 2014; Planck collaboration XIII 2016; Planck collaboration XIV 2016, and references therein). Cosmologists have worked hard to decipher the dark energy code, but we still ignore the physical nature of the DE and hence of the ultimate cause of the observed acceleration of the Universe. Such theoretical conundrum is the so-called Cosmological Constant Problem (CCP) (Weinberg 1989; Sahni & Starobinsky 2000; Padmanabhan 2003; Peebles & Ratra 2003; Copeland, Sami & Tsujikawa 2006; Sol\u00e0 2013). In fact, the cosmological constant (CC), \u039b, or equivalently the vacuum energy density associated to it, $\\rho _\\Lambda =\\Lambda \/(8\\pi G)$ (G being Newton\u2019s gravitational coupling), is usually regarded as the simplest possible explanation for the DE. Historically, the CC was introduced by Einstein in the gravitational field equations 101 yr ago (Einstein 1917). A positive, constant, tiny value (in particle physics units) of order $\\rho _\\Lambda \\sim 2.7\\times 10^{-47}$ GeV4 \u223c (2.3 \u00d7 10\u22123 eV)4 can explain the needed speed up of our cosmos according to the observations. The standard or \u2018concordance\u2019 cosmological model embodies such an assumption as a fundamental built-in principle, together with the hypothesis of dark matter (DM), and for this reason is called the \u039bCDM model. Formulated in terms of the current cosmological parameters, the \u039bCDM assumes that $\\rho _\\Lambda =$const. throughout the history of the Universe, with $\\Omega _\\Lambda \\simeq 0.7$ and \u03a9m \u2243 0.3 at present. Unfortunately, no convincing theoretical explanation is provided about the measured value of $\\rho _\\Lambda$. At the end of the day, no fundamental theory, not even quantum field theory (QFT), can explain this value; and, what is worse, the typical prediction is preposterously large as compared to the measured value. The difficulties inherent to this concept were recognized as of the time when Y.B. Zeldovich first observed (Zeldovich 1967) that the contribution from QFT to the vacuum energy density should be of the order of \u223cm4 for any quantum field of mass m, and therefore many orders of magnitude bigger than the existing upper bound on $\\rho _\\Lambda$ in those days.","Citation Text":["Peebles & Ratra 2003"],"Functions Text":["Cosmologists have worked hard to decipher the dark energy code, but we still ignore the physical nature of the DE and hence of the ultimate cause of the observed acceleration of the Universe. Such theoretical conundrum is the so-called Cosmological Constant Problem (CCP)"],"Functions Label":["Background"],"Citation Start End":[[726,746]],"Functions Start End":[[394,665]]}
{"Identifier":"2020MNRAS.492.3021R__Machacek,_Bryan_&_Abel_2001_Instance_1","Paragraph":"In Fig. 4, we plot the mass growth of each candidate DCBH halo as a function of redshift. In both the panels, we plot the mass of the halo versus the redshift. The left-hand panel contains haloes from the Normal simulation while the right-hand panel contains haloes from the Rarepeak simulation. The grey region in each panel below 106$\\rm {M_{\\odot }}~$ signifies the region below which the mass resolution of Renaissance becomes insufficient to confidently model haloes. Generally, we are able to track haloes below this threshold and into the grey region but below 106$\\rm {M_{\\odot }}~$ results should be treated with caution. The dashed blue line is the limit above which a halo must grow in order to overwhelm the impact of LW radiation, Mmin, LW (Machacek, Bryan & Abel 2001; O\u2019Shea & Norman 2008; Crosby et al. 2013, 2016). The dashed red line is the approximate atomic cooling threshold, Matm, at which point cooling due to atomic hydrogen line emission becomes effective.5 Focusing first on the Normal region in the left-hand panel, we plot the growth rate of the three DCBH candidate haloes identified in the left-hand panel of Fig. 2. The DCBH candidate haloes are rapid growers but are not necessarily the fastest growing haloes in the Normal region. To emphasize this comparison, we also plot the growth of three rapidly growing haloes that contain stars. We select the three star-forming haloes from the final output of the Normal region but haloes at other redshifts do of course exist, which are rapidly growing and contain stars. In this case, we see that haloes with high dM\/dz (i.e. the mass as a function of redshift) values can be star free or star forming and hence having a high dM\/dz does not necessarily discriminate between DCBH halo candidates by itself. Rapidly growing haloes can become metal enriched through external enrichment processes. The enrichment allows the halo interior to cool and to form stars even in the presence of dynamical heating. Therefore, any semi-analytical model or subgrid prescription that uses dM\/dz alone as a predictor for DCBH candidates will inevitably overestimate the number of candidates.","Citation Text":["Machacek, Bryan & Abel 2001"],"Functions Text":["The dashed blue line is the limit above which a halo must grow in order to overwhelm the impact of LW radiation, Mmin, LW"],"Functions Label":["Uses"],"Citation Start End":[[754,781]],"Functions Start End":[[631,752]]}
{"Identifier":"2018AandA...612A..77M__Gromadzki_&_Miko\u0142ajewska_(2009)_Instance_2","Paragraph":"\u201cWiggling\u201d outflows are often observed among young stellar jets and protostellar molecular outflows (Eisloffel et al. 1996; Terquem et al. 1999). Terquem et al. (1999) investigated such binary systems where the accretion disk, from which the jet originates, is inclined to the binary orbital plane. They concluded that the observed jet \u201cwiggling\u201d is a consequence of the jet precession caused by tidal interactions in such non-coplanar binary systems. Nichols & Slavin (2009) as well as Hollis & Michalitsianos (1993) suggested that the precession of the accretion disk around the WD may be responsible for the bending of the wide-angle outflow found in the previous studies. In analogy to these observations of young stellar jets, we suggest that the \u201cwiggling\u201d that we also find here for the R Aqr jet may result from disk precession as well. We estimated the precession period using the theory from Terquem et al. (1999) and the binary parameters taken from Gromadzki & Miko\u0142ajewska (2009) \u2013 Mh = 0.8M\u2299 (the mass of the hot WD companion), Mp\u2215Mh = 1.65 (where Mp is the mass of the primary companion), D = 15.5 AU (the mean semi-major axis of the system), and e = 0.25 (eccentricity). The value of Rd is unknown in our case and we used Rd = 5 AU giving D\u2215Rd \u2248 3 which corresponds to the average value of 2 \u2264 D\u2215Rd \u2264 4 (the range taken from Terquem et al. 1999). For this calculation, we assumed that the angle \u03b4 between the disk plane and that of the binary orbit is small enough (10\u00b0) and we adopted cos\u03b4 = 1. Using Eq. (1) from Gromadzki & Miko\u0142ajewska (2009), we derived the precession time of T \u2248 530 yr. This value is quite large for the wiggling waves that we see. We estimated the projected spatial wavelength \u03bbproj of the \u201cwiggling\u201d wave according to \u03bb = \u03bbproj\u2215sin i, where i is the angle between the jet symmetry axis and the line of sight, and T = \u03bb\u2215\u03c5, where \u03c5 is the jet velocity, from Gromadzki & Miko\u0142ajewska (2009). Using i = 72\u00b0 and \u03c5 ~ 100 km\/s, we derive \u03bbproj \u2248 10 500 AU which is more than 20 times larger than the projected length of the observed wiggling outflow (2\u2032\u2032 \u2248 440 AU). However, we should note that the precessing time strongly depends on the D\u2215Rd; the T decreases significantly with increasing R. It may also be the case that the \u201cwiggling\u201d model developed for YSO jets is not fitting for R Aqr which consists of evolved objects, and both the WD and the disk where the jet probably forms are much hotter than YSO systems. Furthermore, we cannot exclude that the steady wiggling might be a sequence of dynamical interactions of the two collimated flows tilted to each other.","Citation Text":["Gromadzki & Miko\u0142ajewska (2009)"],"Functions Text":["Using Eq. (1) from",", we derived the precession time of T \u2248 530 yr.","This value is quite large for the wiggling waves that we see."],"Functions Label":["Uses","Uses","Compare\/Contrast"],"Citation Start End":[[1531,1562]],"Functions Start End":[[1512,1530],[1562,1609],[1610,1671]]}
{"Identifier":"2016ApJ...819L...7N___2015b_Instance_2","Paragraph":"The gap and ring resemble those in the HL Tau system, recently found by the ALMA long baseline campaign (ALMA Partnership et al. 2015). Our result shows that gaps and rings in the (sub)millimeter dust continuum can exist, not only in relatively young disks (0.1\u20131 Myr) but also in relatively old disks (3\u201310 Myr). One possible mechanism for opening a gap is the gravitational interaction between a planet and the gas (e.g., Lin & Papaloizou 1979; Goldreich & Tremaine 1980; Fung et al. 2014). Such an interaction may also produce the spiral density waves recently found in optical and near-infrared scattered light imaging of dust grains in protoplanetary disks (e.g., Muto et al. 2012). According to recent theoretical analyses of gap structure around a planet (Kanagawa et al. 2015a, 2015b, 2016), the depth and width of the gap are controlled by the planetary mass, the turbulent viscosity, and the gas temperature. The shape of the gap is strongly influenced by angular momentum transfer via turbulent viscosity and\/or instability caused by a steep pressure gradient at the edges of a gap. The observed gap has an apparent width and depth of \n\n\n\n\n\n au and \n\n\n\n\n\n, respectively. This is too shallow and too wide compared with that predicted by theory. However, the observations are limited to an angular resolution of \u223c15 au, and the depth and width could be deeper and narrower in reality. For instance, if we assume that the gap depth times the gap width retains the value derived from the observations, it is possible for the gap to have a width and depth of \n\n\n\n\n\n 6 au and \n\n\n\n\n\n, which is similar to the GPI result (Rapson et al. 2015). Such a gap could be opened by a super-Neptune-mass planet, depending on the parameters of the disk, such as the turbulent viscosity (Kanagawa et al. 2015a, 2015b, 2016). If the gap in the larger dust grains is deeper than that in the gas, the planet could be lighter than super-Neptune mass. We note that a planet of even a few Earth masses, although it cannot open a gap in the gas, can open a gap in the dust distribution if a certain amount of pebble-sized particles, whose motions are not perfectly coupled to that of gas, are scattered by the planet and\/or the spiral density waves excited by the planet (Paardekooper & Mellema 2006; Muto & Inutsuka 2009).","Citation Text":["Kanagawa et al.","2015b"],"Functions Text":["Such a gap could be opened by a super-Neptune-mass planet, depending on the parameters of the disk, such as the turbulent viscosity"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1779,1794],[1802,1807]],"Functions Start End":[[1646,1777]]}
{"Identifier":"2019ApJ...885..165M__Wilman_et_al._2009_Instance_1","Paragraph":"Morphology, colors, and star formation rate (SFR) primarily depend on the small-scale (1 Mpc) environment (Hogg et al. 2004; Kauffmann et al. 2004; Wetzel et al. 2012). This result has been extended by studying galaxy group samples; they show that colors and SF history most directly depend on the properties of the host dark matter (DM) halo (Blanton & Berlind 2007; Wei et al. 2010; Tinker et al. 2012), in agreement with the results of smooth particle hydrodynamic (SPH) simulations by Mazzei & Curir (2003, hereafter MC03). In this context, the investigation of the evolution of group members in the nearby universe acquires a great cosmological interest because more than half of galaxies reside in such environments. Furthermore, since the velocity dispersion of galaxies is significantly lower in groups than in clusters, the merger probability and the effects of interaction on galaxy evolution are much higher. Consequently, groups provide a zoom-in on phenomena driving the galaxy evolution before the galaxies fall into denser environments (e.g., Wilman et al. 2009; Just et al. 2010). Starting from the pioneering works of Toomre & Toomre (1972, and references therein), several papers contributed to shed light on the important role of mergers\/interactions in galaxy evolution\u2014Toomre (1977), Combes et al. (1990), Mihos & Hernquist (1994, 1996), Barnes & Hernquist (1996), Naab & Burkert (2003), Bournaud et al. (2005), and Di Matteo et al. (2007), to name a few\u2014up to the most recent papers of Eliche-Moral et al. (2018) and Martin et al. (2018, and references therein). The dissipative merger simulations of Eliche-Moral et al. (2018) start from systems just formed, composed of a spherical nonrotating DM halo, and by a disk of gas particles with or without the presence of a stellar bulge. These simulations explore about 3\u20133.5 Gyr of evolution. Martin et al. (2018) focused on processes triggering galaxy transformations of massive galaxies (M > 1010 M\u2299) exploiting cosmological hydrodynamic simulations by Kaviraj et al. (2017). These simulations, based on an adaptive mesh refinement code (RAMSES) and including the baryon treatment with stellar and active galactic nucleus (AGN) feedback, are able to resolve baryonic physics on kiloparsec scales, larger than we use in this paper (Section 3). They derived important statistical assessments about the processes that drive morphological transformation across cosmic time.","Citation Text":["Wilman et al. 2009"],"Functions Text":["In this context, the investigation of the evolution of group members in the nearby universe acquires a great cosmological interest because more than half of galaxies reside in such environments. Furthermore, since the velocity dispersion of galaxies is significantly lower in groups than in clusters, the merger probability and the effects of interaction on galaxy evolution are much higher. Consequently, groups provide a zoom-in on phenomena driving the galaxy evolution before the galaxies fall into denser environments (e.g.,"],"Functions Label":["Motivation"],"Citation Start End":[[1058,1076]],"Functions Start End":[[528,1057]]}
{"Identifier":"2021ApJ...920..145H__Damone_et_al._2018_Instance_1","Paragraph":"Over the past decade, many attempts to address this issue have been carried out, such as from the perspective of conventional nuclear physics and even exotic physics beyond the standard BBN framework (Angulo et al. 2005; Cyburt et al. 2008, 2016; Boyd et al. 2010; Pospelov & Pradler 2010; Fields 2011; Kirsebom & Davids 2011; Wang et al. 2011; Broggini et al. 2012; Coc et al. 2012, 2013, 2014; Cyburt & Pospelov 2012; Kang et al. 2012; Voronchev et al. 2012; Bertulani et al. 2013; Hammache et al. 2013; He et al. 2013; Kusakabe et al. 2014; Pizzone et al. 2014; Yamazaki et al. 2014; Hou et al. 2015, 2017; Famiano et al. 2016; Damone et al. 2018; Hartos et al. 2018; Luo et al. 2019; Rijal et al. 2019; Clara & Martins 2020). However, despite the fact some solutions using exotic physics have succeeded in resolving this issue, it appears there is still no universally accepted solution in the academic community since validations of these mysterious exotic physics are beyond the capabilities of current science. Conversely, it seems more worthwhile to exclude any potential possibility of resolving the 7Li discrepancy from the perspective of nuclear physics. It is known that the majority of the primordial 7Li production arises from the decay of 7Be by electron capture during the 2 months after BBN stops. Thus, for the solution of the Li problem, reactions involving 7Be could be more significant than those involving 7Li. Therefore, many reactions that potentially destroy 7Be were investigated to solve this discrepancy over past 10 yr (Kirsebom & Davids 2011; Broggini et al. 2012; Hammache et al. 2013; Hou et al. 2015; Hartos et al. 2018). Meanwhile, enormous efforts have been made to refine the reaction rates of key BBN reactions in the past 20 yr (Smith et al. 1993; Descouvemont et al. 2004; Serpico et al. 2004; Cyburt & Davids 2008; Neff 2011; Pizzone et al. 2014; Tumino et al. 2014; Hou et al. 2015; Barbagallo et al. 2016; Iliadis et al. 2016; Kawabata et al. 2017; Lamia et al. 2017, 2019; Damone et al. 2018; Rijal et al. 2019; Mossa et al. 2020), but the probability of solving or alleviating the 7Li problem by improving our knowledge of relevant nuclear reaction rates still cannot be eliminated. Recent experiments for key nuclear reactions like 7Be(n,p)7Li and 7Be(d,p)24He allow for a reduction of the 7Li production by about 12% (Damone et al. 2018; Rijal et al. 2019) compared to previous calculations. At present, nuclear uncertainties cannot rule out that some of the reactions destroying 7Li are indeed more efficient than those currently used (Boyd et al. 2010; Chakraborty et al. 2011).","Citation Text":["Damone et al. 2018"],"Functions Text":["Over the past decade, many attempts to address this issue have been carried out, such as from the perspective of conventional nuclear physics and even exotic physics beyond the standard BBN framework"],"Functions Label":["Background"],"Citation Start End":[[631,649]],"Functions Start End":[[0,199]]}
{"Identifier":"2022AandA...663A.105P__Brunetti_et_al._2008_Instance_2","Paragraph":"Regardless of the cluster orientation, the spectral index observed for the halo at all available frequencies suggests that it is a USSRH. Despite the number of detected USSRH is still low, radio halos with steep indices are being discovered more and more frequently in the last years thanks to the improved observational capabilities of low-frequency instruments such as GMRT, MWA (Murchison Widefield Array) and LOFAR (Shimwell et al. 2016; Wilber et al. 2018; Bruno et al. 2021; Di Gennaro et al. 2021; Duchesne et al. 2022). An in-depth analysis of all radio halos hosted in Planck clusters and observed in LoTSS, including A1550, has recently been presented in Botteon et al. (2022). USSRH are a prediction of turbulent re-acceleration models (Cassano et al. 2006; Brunetti et al. 2008), in which particles are re-accelerated by turbulence (Brunetti et al. 2001, 2017; Petrosian 2001; Brunetti & Lazarian 2011). On the other hand, the detection of such steep indices is not expected from hadronic (or secondary) models, in which the emission of halos comes from the production of secondary electrons from hadronic collisions between thermal and CR protons (Blasi & Colafrancesco 1999; Dolag & En\u00dflin 2000; Pfrommer et al. 2008). Given that the integrated spectral index observed for the USSRH with LOFAR is \n\n\n\n\n\u03b1\n\n54\n\nMHz\n\n\n144\n\nMHz\n\n\n\u223c\n\u2212\n1.6\n\n\n$ \\alpha_{54\\,\\rm MHz}^{144\\,\\rm MHz} \\sim -1.6 $\n\n\n, we expect an index for the spectral energy distribution8\u03b4\u2004=\u20042\u03b1\u2005\u2212\u20051\u2004=\u2004\u22124.2. If there is no break in the spectrum, the energy budget for these particles would be untenable (Brunetti et al. 2008). Therefore, a break at low energies (\u223cGeV) should exist, suggesting a possible interplay between radiative losses and turbulent re-acceleration during the lifetime of emitting electrons (Brunetti & Jones 2014). Moreover, re-acceleration models predict that a large fraction of halos associated with clusters of masses between 4 and 7\u2005\u00d7\u20051014\u2006M\u2299 should exhibit steep spectra (Cassano et al. 2010, 2012; Brunetti & Jones 2014; Cuciti et al. 2021). The mass of A1550 of \u223c6\u2005\u00d7\u20051014\u2006M\u2299 estimated from Planck Collaboration XXVII (2016) falls in this range9.","Citation Text":["Brunetti et al. 2008"],"Functions Text":["If there is no break in the spectrum, the energy budget for these particles would be untenable"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1575,1595]],"Functions Start End":[[1479,1573]]}
{"Identifier":"2015ApJ...811...57A__Choi_et_al._2014_Instance_1","Paragraph":"Shown in Figures 2(a) and (b) are the z-component of the electric field, Ez, and y-component of the magnetic field, By, respectively. Panels 2(c)\u2013(e) show the transversally averaged (in the yz-plane) electric and magnetic field components, \n\n\n\n\n\n. The energy distribution (total of jet+ambient) and average energy along the x-direction for electron and ion species are shown in Figures 2(f)\u2013(i). All panels are at t\u2217 = 500. Where high-speed jet particles interact with the ambient medium (behind the RS at x\u2217 \u2264 340) or scattered ambient particles blend with the upstream ambient (in front of the FS at 430 \u2264 x\u2217), particles distribution becomes strongly anisotropic. Anisotropies result in the Weibel instability which generates current filaments in these regions with currents in the x-direction. According to Ampere\u2019s law, these current filaments are encircled by transverse magnetic fields, and we see that \n\n\n\n\n\n in Figure 2(c). The transverse electric fields are related to the magnetic fields via \n\n\n\n\n\n where \n\n\n\n\n\n is the velocity of the electron (ion) carrier. The carriers move roughly at the speed of light in the x-direction, \n\n\n\n\n\n. Therefore, the transverse electric field components are Ey = Bz, and Ez = \u2212By, as are observed in the simulation results for \n\n\n\n\n\n (Figures 2(a)\u2013(e)). Additionally, there is a longitudinal ambipolar electric field within the RS transition region, 140 \u2264 x\u2217 \u2264 340 for t\u2217 = 500 (Figure 2(c)). This electric field is generated by the density gradient and different mobilities of electrons and ions (Forslund & Shonk 1970; Forslund & Freidberg 1971; Hoshino 2001; Choi et al. 2014). The magnetic fields act to isotropize the momentum distribution, while the electric fields function to thermalize, and accelerate the particles afterwards. In Figures 2(f)\u2013(i), the shocked region lies between x\u2217 = 340 and x\u2217 = 430. Within the RS transition region (140 \u2264 x\u2217 \u2264 340) jet electrons are trapped by the ambipolar electric field and effectively accelerated up to \u03b3e = 200 by the transverse electric fields (Figures 2(f) and (g)). A tenuous population of these electrons convect upstream due to reflection by the magnetic fields in the shocked region (ellipse in Figure 2(f)). On the other hand, jet ions are slowed in the RS transition region (140 \u2264 x\u2217 \u2264 340) by 40% from the initial Lorentz factor \u03b3i = 10, due to the effect of the ambipolar electric field. In the shocked region, jet electrons have been fully thermalized and are well merged with the thermalized ambient electrons. Thus, only a single electron population is present in the hot shocked region (Figure 2(f)). On the other hand, the kinetic energy of jet ions is transferred to the heating of ambient particles by means of the electromagnetic fields generated by the ion Weibel instability (Figure 2(h)). Full thermalization of the two ion populations (jet and ambient) has not yet occured (demands a longer simulation time), i.e., the two populations are distinguishable in Figure 2(h). Electrons located in the FS transition region (430 \u2264 x\u2217 \u2264 500), predominantly ambient electrons, also undergo the Weibel instability and are thermalized by the jet upstream kinetic energy. In this region, penetrating jet ions interact with ambient particles and are slowed down gradually by 50% from the jet front Lorentz factor \u03b3i = 10 to a minimum value of \u03b3i = 5 (Figures 2(h) and (i)).","Citation Text":["Choi et al. 2014"],"Functions Text":["This electric field is generated by the density gradient and different mobilities of electrons and ions"],"Functions Label":["Uses"],"Citation Start End":[[1605,1621]],"Functions Start End":[[1436,1539]]}
{"Identifier":"2021ApJ...908..248D__Guseva_&_Martynenko_1981_Instance_1","Paragraph":"Where \n\n\n\n\n\n is the straggle or standard deviation of the implantation distribution, D is the diffusion coefficient, and t is the irradiation time. In the case of gas implantation via ion irradiation, it is often assumed that irradiation damage produces enough defects to serve as trapping sites that the diffusion coefficient is negligible (Martynenko 1977; Scherzer 1983). In Section 2.5, an estimate of the effective diffusion in the presence of damage produced by relativistic light ion impacts is made. Lattice defects caused by ion irradiation also provide nucleation for fixed helium bubbles (Kornelsen 1972). We can assume additional trapping due to the interaction of hydrogen with helium in the material, since helium bubbles serve as trapping sites for hydrogen atoms through synergistic effects (Hayward & Deo 2012), leading to a lower critical dose in the case of mixed hydrogen and helium exposure (Guseva & Martynenko 1981). At temperatures well below the melting point, such as that expected for a relativistic interstellar spacecraft, high-flux hydrogen irradiation damage effects have been reproduced in low-flux experiments (Gao et al. 2019). However, for an interstellar probe, the flux is sufficiently low that diffusion effects may play a non-negligible role. Classical diffusion of individual gas atoms could lead to lower local gas atom concentrations as the implantation distributions widen, or increased gas atom concentrations around trapping sites such as grain boundaries; the effect of diffusion is discussed in Section 2.5. Helium bubble nucleation and migration could lead to increased gas atom concentrations and damage near surfaces, grain boundaries, and other defects (Nakamura et al. 1977; Goodhew 1983; Lane & Goodhew 1983). No single theoretical framework exists to summarily treat these effects. However, simple models such as that presented here immediately offer compelling mitigation strategies. To perform this analysis, we find implantation profiles of ISM gas atoms at relativistic speeds using an ion\u2013material interactions code, calculate critical concentrations for blistering onset for hydrogen and helium individually assuming a worst-case scenario of negligible diffusion, and show the effect of non-negligible diffusion on local gas concentrations.","Citation Text":["Guseva & Martynenko 1981"],"Functions Text":["We can assume additional trapping due to the interaction of hydrogen with helium in the material, since helium bubbles serve as trapping sites for hydrogen atoms through synergistic effects",", leading to a lower critical dose in the case of mixed hydrogen and helium exposure"],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[913,937]],"Functions Start End":[[617,806],[827,911]]}
{"Identifier":"2016MNRAS.460..590F__Miralda-Escud\u00e9_et_al._1996_Instance_1","Paragraph":"It is important to note that the above procedure differs from the observational one, whereby Voigt profiles are fit to the absorption features in transmission spectra (e\u2212\u03c4) in order to extract column densities and Doppler broadening parameters. This approach enables deblending of multiple-component absorption, and takes into account the broadening of the lines due to the instrumental profile, enabling accurate recovery of their column densities. The technique was originally devised under the premise that the intervening absorption lines in QSO spectra arose from discrete absorbing clouds, but this picture was challenged early by the smoothly distributed IGM captured in cosmological hydrodynamical simulations (e.g. Cen et al. 1994; Hernquist et al. 1996; Miralda-Escud\u00e9 et al. 1996; Theuns et al. 1998; Dav\u00e9 et al. 1999). Absorbers that have a large spatial extent take part in the Hubble expansion, which leads to a line profile that deviates from a Voigt profile. Voigt profile fitting the transmission spectra will therefore glean slightly different results to simply taking peaks in the \u03c4 distribution. In particular, we might expect a larger number of absorbers, and some differences in the derived Doppler broadening parameters and column densities. As a check, we derived Voigt profiles for lines recovered in some of the sight-lines extracted from the simulation using our \u03c4 peak method, and plotted these on top of the transmission spectrum. An example is shown in Fig. 3. The transmission spectrum is shown as the black dashed line, and the predicted Voigt profiles are shown in green. More structure is apparent in the real spectrum, which would yield a larger number of Voigt components in Voigt profile fitting, and column densities and Doppler broadening parameters that differ slightly from those we have recovered via the \u03c4 peak method. This is an obvious caveat to our approach, which we bear in mind when interpreting our results later on.","Citation Text":["Miralda-Escud\u00e9 et al. 1996"],"Functions Text":["The technique was originally devised under the premise that the intervening absorption lines in QSO spectra arose from discrete absorbing clouds, but this picture was challenged early by the smoothly distributed IGM captured in cosmological hydrodynamical simulations (e.g."],"Functions Label":["Background"],"Citation Start End":[[764,790]],"Functions Start End":[[450,723]]}
{"Identifier":"2017AandA...601A.143F___2015_Instance_1","Paragraph":"Although wind observations are very common in AGN (see Elvis 2000; Veilleux et al. 2005; and Fabian 2012, for reviews), most studies concern ionised gas and uncertain spatial scales. In the past few years the situation changed drastically. Several fast (vOF of the order of 1000 km\u2009s-1), massive outflows of ionised, neutral and molecular gas, extended on kpc scales, have been discovered thanks to three techniques: 1) deep optical\/NIR spectroscopy, mainly from integral field observations (IFU, e.g. Nesvadba et al. 2006, 2008; Alexander et al. 2010; Rupke & Veilleux 2011; Riffel & Storchi-Bergmann 2011; Cano-Diaz et al. 2012; Greene et al. 2012; Harrison et al. 2012, 2014; Liu et al. 2013a,b; Cimatti et al. 2013; Tadhunter et al. 2014; Genzel et al. 2014; Brusa et al. 2015a; Cresci et al. 2015; Carniani et al. 2015; Perna et al. 2015a,b; Zakamska et al. 2016); 2) interferometric observations in the (sub)millmetre domain (e.g. Feruglio et al. 2010, 2013a,b, 2015; Alatalo et al. 2011; Aalto et al. 2012; Cicone et al. 2012, 2014, 2015; Maiolino et al. 2012, Krips et al. 2011; Morganti et al. 2013a,b; Combes et al. 2013; Garcia-Burillo et al. 2014); and 3) far-infrared spectroscopy from Herschel (e.g. Fischer et al. 2010; Sturm et al. 2011; Veilleux et al. 2013; Spoon et al. 2013; Stone et al. 2016; Gonzalez-Alfonso et al. 2017). In addition, AGN-driven winds from the accretion disk scale up to the dusty torus are now detected routinely both in the local and in the distant Universe, as blue-shifted absorption lines in the X-ray spectra of a substantial fraction of AGN (e.g. Piconcelli et al. 2005; Kaastra et al. 2014). The most powerful of these winds, observed in 20\u201340% of local AGN (e.g. Tombesi et al. 2010) and in a handful of higher redshift objects (e.g. Chartas et al. 2009; Lanzuisi et al. 2012), have extreme velocities (ultra-fast outflows, UFOs, v 0.1-0.3c) and are made by highly ionised gas which can be detected only at X-ray energies. ","Citation Text":["Feruglio et al.","2015"],"Functions Text":["interferometric observations in the (sub)millmetre domain"],"Functions Label":["Background"],"Citation Start End":[[937,952],[968,972]],"Functions Start End":[[873,930]]}
{"Identifier":"2021MNRAS.503.4387A__Neronov_&_Vovk_2010_Instance_1","Paragraph":"Observations of large-scale magnetic fields offer clear insights about regular ordered patterns of the field lines, suggesting that mean-field dynamo processes are responsible for their order and structure, as well as the existence of additional transport processes carrying magnetic energy into huge regions of space (Clarke, Kronberg & B\u201dohringer 2001; Bonafede et al. 2010; Arlen et al. 2012; Govoni, F. et al. 2017; Han 2017). Magnetic fields cannot be directly observed, so their impact on radiation processes need to be considered (see, for instance, Rybicki & Lightman 1979, and references therein). In addition, observations of magnetic fields in voids provide bounds on their strength, depending on the analytical model: From the simplest ones, it is possible to obtain lower bounds, although when improving such models a bounded range of magnetic field strength values can be provided, ranging from 10\u221225 to 10\u221215 nG (Neronov & Vovk 2010; Tavecchio et al. 2010; Essey, Ando & Kusenko 2011; Takahashi et al. 2013). Other authors argue magnetic field strengths between 10\u221216 and 10\u221215 G (Einstein 1915; Einstein 1916; Hubble 1929; Hubble & Humason 1931; Bull et al. 2016). Magnetic fields from astrophysical voids are relevant since they could evidence truly cosmological magnetic fields, which could have served as seeds for magnetic fields in lower scales (such as galactic fields). Observations in galaxy clusters yield values of the order of 10\u22126 G (Boulanger et al. 2018). Interest in primordial magnetic fields generated during inflation has recently increased. This scenario has driven the search for phenomenologically viable mechanisms to explain the observed magnetic fields in a broad set of scales (Grasso & Rubinstein 2001; Brandenburg & Subramanian 2005; Demozzi, Mukhanov & Rubinstein 2009; Ferreira, Jain & Sloth 2013; Green & Kobayashi 2016). The latest data on reionization and the observed UV luminosity function of high-redshift galaxies place limits on the magnetic field strength due to its impact on the reionization process.","Citation Text":["Neronov & Vovk 2010"],"Functions Text":["In addition, observations of magnetic fields in voids provide bounds on their strength, depending on the analytical model: From the simplest ones, it is possible to obtain lower bounds, although when improving such models a bounded range of magnetic field strength values can be provided, ranging from 10\u221225 to 10\u221215 nG"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[928,947]],"Functions Start End":[[607,926]]}
{"Identifier":"2015MNRAS.448...42L__Shetrone_et_al._2003_Instance_1","Paragraph":"GCs formed in dwarf galaxies may differ from those found in the Galactic halo, depending on their age and metallicity. Dwarf galaxies show a wide variety of star formation histories (Hidalgo et al. 2011, 2013; Weisz et al. 2014) that are predicted to lead to variations in their metallicity distribution functions and chemical abundances. It has also been suggested these variations could be attributed to differences in the IMFs of these galaxies (McWilliam, Wallerstein & Mottini 2013). If the IMFs are the root cause of these differences, then this would also lead to differences in the age\u2013metallicity relationship, which is observed by both Forbes & Bridges (2010) and Leaman et al. (2013). From observations, field stars in dwarf galaxies do exhibit different abundance ratios from MW field stars, e.g. lower [\u03b1\/Fe] ratios and variations in neutron-capture element ratios at intermediate metallicities. However, these typically do not show up until [Fe\/H] \u223c \u22121.5 (Shetrone, Bolte & Stetson 1998; Shetrone, C\u00f4t\u00e9 & Sargent 2001; Shetrone et al. 2003; Venn et al. 2004; Okamoto et al. 2012; Tolstoy, Hill & Tosi 2009; Frebel 2010). At metallicities below [Fe\/H] = \u22121.5, the abundance variations between field and GC stars become less pronounced in dwarfs and the MW (Hill et al. 2000; Pritzl, Venn & Irwin 2005; Carretta et al. 2010; Letarte et al. 2010); a good example of this is M54, located at the heart of the Sagittarius (Sgr) dwarf accretion remnant. M54 has a much lower metallicity than the Sgr field stars (e.g. Carretta et al. 2010) and the [\u03b1\/Fe] ratios resemble the field stars in the MW halo and its detailed chemical abundance ratios resemble the patterns seen in other GC systems (e.g. the Na\u2013O anticorrelation; Carretta et al. 2009). Therefore, other than its physical association with the Sgr remnant, M54 does not stand out from other GCs in terms of its chemical abundance patterns, similar to the metal-poor GCs Terzan 8 and Arp 2 (both also kinematically and spatial associated with the Sgr stream; Mottini, Wallerstein & McWilliam 2008). On the other hand, Hodge 11 in the large magellanic cloud at [Fe\/H] = \u22122.0 does have lower [\u03b1\/Fe] than MW field and GC stars(Mateluna et al. 2012), and Ruprecht 106 has an anomalously low [\u03b1\/Fe] ratio for a MW GC (Villanova et al. 2013).","Citation Text":["Shetrone et al. 2003"],"Functions Text":["From observations, field stars in dwarf galaxies do exhibit different abundance ratios from MW field stars, e.g. lower [\u03b1\/Fe] ratios and variations in neutron-capture element ratios at intermediate metallicities. However, these typically do not show up until [Fe\/H] \u223c \u22121.5"],"Functions Label":["Background"],"Citation Start End":[[1033,1053]],"Functions Start End":[[696,968]]}
{"Identifier":"2016ApJ...827...75L__Gal-Yam_et_al._2006_Instance_1","Paragraph":"The coalescence of a binary compact object system (either a neutron star (NS) binary or a stellar-mass black hole (BH) and NS binary) has been widely suggested to account for short-duration gamma-ray burst (SGRB) events (Eichler et al. 1989; Narayan et al. 1992; Nakar 2007; Berger 2014) that last typically less than 2 s in the soft \u03b3\u2013ray band (Kouveliotou et al. 1993). Since 2006, it has been suspected that mergers of compact objects could also produce the so-called long\u2013short GRBs (also known as the supernova-less long GRBs, which are apparently long-lasting but do not show any signal of supernovae down to very stringent limits), which share some properties of both long- and short-duration GRBs (Della Valle et al. 2006; Gal-Yam et al. 2006; Gehrels et al. 2006; Zhang et al. 2007). Compact binary coalescence (CBC) is generally expected to be a strong source of gravitational wave (GW) radiation and such events are prime targets for some GW detectors like advanced Laser Interferometer Gravitational-Wave Observatory (LIGO)\/VIRGO (Abadie et al. 2015; Acernese et al. 2015; Belczynski et al. 2010, 2016; see also the latest LSC\u2013Virgo white paper at https:\/\/dcc.ligo.org\/LIGO-T1400054\/public). On 2015 September 14, the two detectors of the LIGO simultaneously detected a transient gravitational wave signal from the merger of two BHs (GW150914, Abbott et al. 2016). GW150914 is the first direct detection of gravitational waves and the first identification of a binary BH merger (Abbott et al. 2016). Surprisingly, the observations from the Fermi Gamma-ray Burst Monitor (GBM) at the time of GW150914 claimed a detection of a weak gamma-ray transient (i.e., GBM transient 150914) 0.4 s after GW150914 with a false alarm probability of 0.0022 (Connaughton et al. 2016). If true, this is the first GW\/SGRB association (see, however, Savchenko et al. 2016 for some arguments). Li et al. (2016) compared GBM transient 150914 with other SGRBs and found that such an event is remarkably different in its prompt emission properties. The binary BH merger origin as well as its property of \u201cdistinguished\u201d prompt emission suggest that GW150914\/GBM transient 150914 is not a typical GW\/SGRB association.","Citation Text":["Gal-Yam et al. 2006"],"Functions Text":["Since 2006, it has been suspected that mergers of compact objects could also produce the so-called long\u2013short GRBs (also known as the supernova-less long GRBs, which are apparently long-lasting but do not show any signal of supernovae down to very stringent limits), which share some properties of both long- and short-duration GRBs"],"Functions Label":["Background"],"Citation Start End":[[731,750]],"Functions Start End":[[372,704]]}
{"Identifier":"2018AandA...611A..85S__Schleicher_&_Dreizler_(2014)_Instance_1","Paragraph":"After the detection of V391 Peg b, some other planet or brown dwarf (BD) candidates orbiting sdB stars were proposed using different detection methods. From eclipse timing, about one-third of the known detached sdB\/sdO + dM (dM = M-dwarf) post-common-envelope binaries (PCEB) are suspected to host planets\/BDs: HW Vir (Beuermann et al. 2012 and references therein), HS 0705+6700 (alias V470 Cam, Qian et al. 2013 and references therein), HS 2231+2441 (Qian et al. 2010 and references therein; but see also Lohr et al. 2014), NSVS 14256825 (Almeida et al. 2013; Hinse et al. 2014 and references therein), NY Vir (Lee et al. 2014 and references therein), and 2M 1938+4603 (Baran et al. 2015). Interesting explorations on the origin of PCEB (and specifically sdB+MS\/BD) circumbinary planets can be found in Zorotovic & Schreiber (2013), Schleicher & Dreizler (2014), Bear & Soker (2014), and V\u00f6lschow et al. (2016). Very different planets or planetary remnants with terrestrial radii have been proposed from tiny reflection effects detected by the Kepler spacecraft in KIC 05807616 (Charpinet et al. 2011) and KIC 10001893 (Silvotti et al. 2014). However, none of these sdB planet\/BD candidates has been confirmed with at least two independent detection methods. More robust detections of a few brown dwarfs (BDs) in eclipsing sdB binaries (also called HW Vir systems from the sdB+dM protoptype) were obtained by combining stellar radial velocities (RVs) with photometric measurements: J08205+0008, J1622+4730 and V2008-1753 have companion masses of about 71, 67, and 69 MJup, respectively (Geier et al. 2011; Schaffenroth et al. 2014a, 2015). At least two more sdB+BD eclipsing systems were recently found from the OGLE survey (Schaffenroth, in prep., priv. comm.). Finally, two more BD candidates in sdB binaries were found by combining radial velocities (RVs) with photometric reflection effects: CPD-64\u00b06481 and PHL 457, with minimum masses of 50 and 28 MJup, respectively (Schaffenroth et al. 2014b).","Citation Text":["Schleicher & Dreizler (2014)"],"Functions Text":["Interesting explorations on the origin of PCEB (and specifically sdB+MS\/BD) circumbinary planets can be found in Zorotovic & Schreiber (2013),","Bear & Soker (2014), and V\u00f6lschow et al. (2016)."],"Functions Label":["Background","Background"],"Citation Start End":[[834,862]],"Functions Start End":[[691,833],[864,912]]}
{"Identifier":"2020ApJ...889...29C__Kalapotharakos_et_al._2014_Instance_1","Paragraph":"The global magnetospheric structures for the oblique rotator are very similar to the aligned one. We show the structure of magnetic field lines and the distribution of the accelerating electric field E0 in the x\u2013z plane for a 60\u00b0 rotator with the pair multiplicity \u03ba = {0, 1, 3} in Figure 4. As the pair multiplicity \u03ba increases, the field structure tends to the force-free solution with an equatorial current sheet outside the LC. We observer a dissipative region where E > B outside the LC. The spatial extension of the dissipative region decreases with increasing pair multiplicity and the E0 region is more confined to the equatorial current sheet outside the LC as the pair multiplicity \u03ba increases. In fact, the E0 distribution for the high \u03ba solution is qualitatively similar to the FIDO one (see, e.g., Kalapotharakos et al. 2014; Cao & Yang 2019). We also compare the field structures for \u03ba = 0 with Figure 1 of Contopoulos (2016) for \u03b1 = 0\u00b0 and \u03b1 = 60\u00b0 respectively. We find that the field structures are qualitatively very similar to those of Contopoulos (2016). For comparison, we also show the magnetic field lines and the E0 distributions for a 60\u00b0 rotator with the pair multiplicity \u03ba = 0 by implementing the AE formulation everywhere in Figure 5. The magnetospheric structure is very similar to the aligned one with a force-free zone bounded by a radiation zone. We observe a strong E0 distribution inside the LC, which is very different from those in the SG and OG models. A strong E0 region with E > B also appears outside the LC. We show the distributions of magnetic field lines and the accelerating electric field E0 in the x\u2013z plane for a 30\u00b0 rotator with the pair multiplicity \u03ba = 3 in Figure 6. We see that the field structure is very close to the force-free one and the E0 region is restricted to only near the current sheet outside the LC for this high \u03ba value. We also show the normalized Poynting flux L\/Laligned as a function of radius r for a 90\u00b0 rotator with different pair multiplicities in Figure 7. We see that the Poynting flux increases with increasing \u03ba values and approaches the force-free solution for the high \u03ba value. Our simulation shows a more than 1% dissipation rate outside the LC for a 90\u00b0 dissipative rotator. A similar dissipation rate is also found by the PIC simulation for the aligned and perpendicular rotator (Philippov et al. 2015). In fact, the spectral numerical methods present an unphysical dissipation beyond the LC due to discontinuity in the current sheet. A higher resolution is necessary to catch the discontinuity in the current sheet and reduce the unphysical dissipation.","Citation Text":["Kalapotharakos et al. 2014"],"Functions Text":["In fact, the E0 distribution for the high \u03ba solution is qualitatively similar to the FIDO one"],"Functions Label":["Similarities"],"Citation Start End":[[811,837]],"Functions Start End":[[705,798]]}
{"Identifier":"2015AandA...579A.132P__Simha_et_al._(2009)_Instance_2","Paragraph":"A common feature of all previous models is that the relation between the central galaxy stellar mass and the halo mass reaches a maximum at halo masses ~1012 M\u2299. According to Yang et al. (2012), below this threshold the mass accretion of the central galaxy is dominated by star formation. Thus, when the halo mass reaches ~1012 M\u2299 a process takes place that quenches the star formation. Interestingly, this mass scale is very similar to the cold-mode to hot-mode transition scale (Birnboim & Dekel 2003; Kere\u0161 et al. 2005) in the theory of gas accretion, as derived in hydrodynamic simulations, whereas large halos primarily accrete hot gas and low mass halos cold gas. This would suggest that the quenching of central galaxies coincides with the formation of a hot gaseous halo, and thus with a lack of cold gas supply. What would be the fate of satellites? According to Simha et al. (2009), the subhalos also retain their identity for quite some time after accreting a larger halo, so satellites in subhalos less massive than ~1012 M\u2299 do not immediately see the effect of the hot gas in the larger halo and accrete in cold mode. Thus, consistent with the results of Yang et al. (2012) and B\u00e9thermin et al. (2013), satellite galaxies continue to accrete gas and convert it to stars over a rather long period, which according to Simha et al. (2009) is about of 0.5\u22121 Gyr after the merger. The gas accretion declines steadily over this period. Since star formation follows mass accretion with a short delay, satellites should experience quenching in a similar amount of time. This scenario would be consistent with our observations. Indeed, at z ~ 1 when massive halos are just forming via merger, the SF activity in the accreted subhalos is still high. At later epochs, instead, the transition to the hot mode accretion of the satellites and the consequent progressive quenching of their SF activity would lead to the faster decline of their contribution to the CSFH with respect to lower mass halos, which evolve in a cold mode accretion phase maintaining a high SFR. ","Citation Text":["Simha et al. (2009)"],"Functions Text":["which according to","is about of 0.5\u22121 Gyr after the merger"],"Functions Label":["Uses","Uses"],"Citation Start End":[[1329,1348]],"Functions Start End":[[1310,1328],[1349,1387]]}
{"Identifier":"2017AandA...598A..66P___2016_Instance_1","Paragraph":"In both models, as we will see, the thick disc scale length is about a factor of two shorter than that of the thin disc, in agreement with the results by Bovy et al. (2012a). The choice of presenting two mass models for the mass distribution of our Galaxy is mainly dictated by two reasons. First, the need to add a central bulge to the global gravitational potential to reproduce the rotation curve in the inner kpcs of the Milky Way strongly depends on the observational data with which one compares the theoretical curve: to reproduce the rise observed in the inner kpcs (see the observational data adopted by Caldwell & Ostriker 1981), Allen & Santillan (1991) introduced a central mass concentration, whose mass is about 15% of the disc mass. However, the central rise observed in the rotation of the molecular gas in the inner Galaxy (for more recent estimates see, for example,  Sofue 2012) may be an effect of non circular motions generated by large scale asymmetries like the bar, as has been shown recently by Chemin et al. (2015). Moreover, this feature is not reported in all the observational studies (see, for example, Reid et al. 2014). Secondly, there is growing evidence that the mass of any classical bulge, if present in the Milky Way, must be small (Shen et al. 2010; Kunder et al. 2012, 2016; Di Matteo et al. 2014, 2015). For these reasons, we prefer to present a second model, our Model II, which does not include any spherical central component, and which is still compatible with the rotation curve of the Galaxy, as given by Reid et al. (2014). Because it has been widely used in the last decades, and due to the facility of its implementation, we explicitly aim at generating Galactic models similar to the Allen & Santillan (1991) model, so to make any implementation of these new models, and any comparison with Allen & Santillan (1991), straightforward. As for the model proposed by Allen & Santillan (1991), Models I and II are axisymmetric and time-independent, and do not include stellar asymmetries such as a bar or spiral arms. No truncation is assumed for the discs, while the halo is truncated at 100 kpc, in agreement with the choice of Allen & Santillan (1991). As we describe in the following section, the analytic forms for the discs, halo, and bulge potentials are the same as those adopted by Allen & Santillan (1991). To allow an easy comparison with the Allen & Santillan (1991) model, in the following we will make use of the same system of units adopted by these authors: the potential is given in units of 100 km2\/ s2, lengths are in kpc, masses in units of 2.32 \u00d7 107M\u2299, time in units of 0.1 Gyr, velocities in units of 10 km s-1 and the vertical force in units of 10-9 cm s2. In these units, the gravitational constant G is equal to 1 and the mass volume density is in units of 2.32 \u00d7 107M\u2299\/ kpc3. ","Citation Text":["Kunder et al.","2016"],"Functions Text":["Secondly, there is growing evidence that the mass of any classical bulge, if present in the Milky Way, must be small","For these reasons, we prefer to present a second model, our Model II, which does not include any spherical central component, and which is still compatible with the rotation curve of the Galaxy, as given by Reid et al. (2014)."],"Functions Label":["Background","Motivation"],"Citation Start End":[[1288,1301],[1308,1312]],"Functions Start End":[[1152,1268],[1344,1570]]}
{"Identifier":"2021MNRAS.501.3781R___2017_Instance_1","Paragraph":"While spatially extended optical jets and bipolar CO molecular outflows have been observed in numerous Class 0\/I protostars (e.g. Reipurth & Bally 2001; Bally 2016, and references therein), near-infrared high-resolution spectroscopy and spectroimaging observations in the past two decades have made it possible to study the kinematics of the outflowing gas and physical properties at the base of the jet within a few hundred au of the driving source in Class 0\/I protostars (e.g. Davis et al. 2001, 2003, 2011; Nisini et al. 2005, 2016; Caratti o Garatti et al. 2006; Takami et al. 2006; Antoniucci et al. 2008, 2011, 2017; Garcia Lopez et al. 2008, 2013). These microjets are bright in [Fe\u2009ii] forbidden and H2 rovibrational emission lines, hence showing the presence of forbidden emission-line (FEL) regions and molecular hydrogen emission-line (MHEL) regions in low-mass Class 0\/I protostars. While multiple low- and high-velocity components are observed in both MHELs and FELs, the higher velocity gas is slightly further offset from the driving source than the slower gas, and the kinematics of the H2 emission differs from [Fe\u2009ii] emission, revealing complicated kinematic structures. Evidence of H2 emission from cavity walls is also seen in some protostars, suggesting the presence of a wide-angled wind. Strong emission in the well-known accretion diagnostics of Paschen and Brackett hydrogen recombination lines is observed in protostars, with the ratio of the accretion luminosity to bolometric luminosity spanning from \u223c0.1 to \u223c1. The mass accretion and loss rates for Class 0\/I low-mass protostars span the range of 10\u22126\u201310\u22128 M\u2299 yr\u22121, and the derived jet efficiencies (ratio between mass ejection and accretion rates) range between \u223c1 per\u2009cent and 10 per\u2009cent (e.g. Davis et al. 2001, 2003, 2011; Nisini et al. 2005, 2016; Caratti o Garatti et al. 2006; Takami et al. 2006; Antoniucci et al. 2008, 2011, 2017; Garcia Lopez et al. 2008, 2013). These measurements are within the range predicted by the magnetohydrodynamic jet launching models (e.g. Frank et al. 2014).","Citation Text":["Antoniucci et al.","2017"],"Functions Text":["near-infrared high-resolution spectroscopy and spectroimaging observations in the past two decades have made it possible to study the kinematics of the outflowing gas and physical properties at the base of the jet within a few hundred au of the driving source in Class 0\/I protostars (e.g."],"Functions Label":["Background"],"Citation Start End":[[588,605],[618,622]],"Functions Start End":[[190,479]]}
{"Identifier":"2019ApJ...871..243Y__Yen_et_al._2011_Instance_1","Paragraph":"There are two possibilities resulting in the different magnetic field strengths inferred from the polarimetric and molecular-line observations: (1) the rotational-to-gravitational energy \u03b2rot is overestimated, and (2) there are additional contributions in the polarized intensity from other mechanisms, such as dust scattering. In our MHD simulations, \u03b2rot is adopted to be 0.4% based on the observational estimates of the core mass of \u223c1 M and the angular speed of the core rotation of 4 \u00d7 10\u221214 s\u22121. The angular speed was estimated based on the global velocity gradient along the major axis of the dense core observed with single-dish telescopes (Saito et al. 1999; Yen et al. 2011; Kurono et al. 2013). Numerical simulations of dense cores including synthetic observations show that the specific angular momentum derived from the synthetic images of the dense cores can be a factor of 8\u201310 higher than their actual specific angular momentum computed by the sum of the angular momenta contributed by the individual gas parcels in the dense cores (Dib et al. 2010). In addition, if there are filamentary structures in the dense core in B335, which could not be resolved with the single-dish observations, infalling motions along the filamentary structures could also contribute to the observed velocity gradient, leading to an overestimated angular speed of the core rotation (Tobin et al. 2012). We have also performed our simulations with a lower \u03b2rot, and we find that the rotational velocity on a 100 au scale in the simulations decreases with decreasing \u03b2rot. Thus, the discrepancy in the magnetic field strengths inferred from the field structures and the gas kinematics can be reconciled, if the core rotation in B335 is overestimated by a factor of a few in the observations, and these results would suggest a weak magnetic field of initial \u03bb of 9.6 in B335. Further observations combining single dishes and interferometers to have a high spatial dynamical range and to map the velocity structures of the entire dense core in B335 at a high angular resolution are needed to study coherent velocity features and provide a better estimate of the core rotation.","Citation Text":["Yen et al. 2011"],"Functions Text":["In our MHD simulations, \u03b2rot is adopted to be 0.4% based on the observational estimates of the core mass of \u223c1 M and the angular speed of the core rotation of 4 \u00d7 10\u221214 s\u22121. The angular speed was estimated based on the global velocity gradient along the major axis of the dense core observed with single-dish telescopes"],"Functions Label":["Uses"],"Citation Start End":[[668,683]],"Functions Start End":[[328,647]]}
{"Identifier":"2022AandA...662A..42M__V\u00e1zquez_2007_Instance_3","Paragraph":"A number of fundamental results have been rigorously proved in the mathematical literature concerning the asymptotic behaviour in time of some of the solutions of the porous medium equation and related equations (e.g. Kamin & V\u00e1zquez 1991; Bernis et al. 1993; Hulshof et al. 2001). What is of interest for us here is, primarily, the results that can be applied to the cylindrically symmetric case with diffusion coefficient which is proportional to the square of the dependent variable (n\u2004=\u20042, m\u2004=\u20043 in the notation of Eq. (7)). The most basic result, already mentioned in Sect. 4.4.1, is that initial conditions which have a finite nonzero flux integral (called \u2018the mass\u2019 in the mathematical literature for the PME) converge toward the ZKBP solution with the same flux integral (\u2018mass\u2019) asymptotically in time (V\u00e1zquez 2007, Theorem 18.2); here, allowance is made for either a positive or negative flux integral by globally changing the sign of the ZKBP solution; also, \u2018convergence\u2019 is meant in the sense that the Lp norm of the difference between the actual solution and the ZKBP function tends to zero as t\u2004\u2192\u2004\u221e faster than a negative power of the time with an exponent which is a function of n, m, and p (e.g. \u22121\/3 for n\u2004=\u20042 and m\u2004=\u20043 in the L2 norm; see details in the book by V\u00e1zquez 2007). A complementary result is the following: when the initial condition has positive net flux and its negative part has compact support, then the whole solution evolves into a positive function after a finite time (V\u00e1zquez 2007, Theorem 18.29). Since we are dealing with signed functions which have zero flux integral, these results are of interest mainly because they impose a strict condition on the possible flux imbalance caused by numerical errors (as discussed in Sect. 4.4.1, final paragraph): if it is not small, the numerical solutions will approach the ZKBP solution in a comparatively short time. However, the flux imbalance in all the Bifrost experiments discussed in the present paper is small enough that they have not shown this behaviour even though they have been run until a very long diffusive time.","Citation Text":["V\u00e1zquez 2007"],"Functions Text":["A complementary result is the following: when the initial condition has positive net flux and its negative part has compact support, then the whole solution evolves into a positive function after a finite time","Theorem 18.29"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1509,1521]],"Functions Start End":[[1298,1507],[1523,1536]]}
{"Identifier":"2021MNRAS.506.3313G__Surana_et_al._2020_Instance_1","Paragraph":"Deep learning (DL) is the paradigm of machine learning which uses multilayer neural networks. Neural networks are ML models inspired from the network of brain cells or neurons. The differentiating factor between DL and conventional ML is in the process of feature selection. In conventional ML, performance strongly depends on the features used. More often than not, new features are created for a task to better capture correlations in the data. This is called feature engineering. ML models perform poorly if not given good features to learn from. Deep learning eliminates this dependence by defining its own features to learn from, which are relevant to the task at hand. This makes deep learning a versatile, high performing solution for a variety of supervised ML tasks. The Artificial Neural Networks (ANNs) have found application in not only the traditional tasks in extragalactic astronomy such as photometric redshift estimation (e.g. Firth, Lahav & Somerville 2003; Tagliaferri et al. 2003; Collister & Lahav 2004; Vanzella et al. 2004; Sadeh, Abdalla & Lahav 2016; Bilicki et al. 2018; Pasquet et al. 2019), but also in more specialized problems such as predicting infrared luminosity of galaxies (Ellison et al. 2016), estimation of star formation properties (Surana et al. 2020), and ranking the quenching parameters of galaxies (Teimoorinia, Bluck & Ellison 2016). In particular, Convolutional Neural Networks (CNN) are now becoming increasingly popular in studies using imaging data on galaxies. The CNNs are deep learning models designed to extract features from images. They have provided state-of-the-art performance in majority of computer vision tasks in recent times (Krizhevsky, Sutskever & Hinton 2012; He et al. 2017). CNNs have been utilized for galaxy morphological classification e.g. classifying the optical morphologies broadly into spheroidal, disc, and irregular types (Huertas-Company et al. 2015) and even for classification of various radio galaxy morphologies (Wu et al. 2019). Recently, Ribli, Dobos & Csabai (2019) have shown that CNNs can be used to predict galaxy shapes, needed for weak lensing studies, from a wide field but shallow sky survey images using the \u2018ground truth\u2019 images from a deeper but narrower field survey. CNNs have also found application in the automated detection of features in sky-survey images such as galactic bars (Abraham et al. 2018) and strong gravitational lenses (Canameras et al. 2020; Li et al. 2020). Further, ML and CNNs in particular have also been used to detect outliers in large area Sloan Digital Sky Survey (SDSS) data (e.g. Baron & Poznanski 2017; Sharma et al. 2019). In this study, we want to determine the bulge to total luminosity ratio (B\/T) of a galaxy using its optical multiband images as input. Due to the use of a data set with galaxy images labelled by $B\/T\\, \\in [0,1]$ in a continuous space $\\rm I\\!R$, we have performed CNN based regression in this work.","Citation Text":["Surana et al. 2020"],"Functions Text":["The Artificial Neural Networks (ANNs) have found application in not only the traditional tasks in extragalactic astronomy","but also in more specialized problems such as","estimation of star formation properties"],"Functions Label":["Background","Background","Background"],"Citation Start End":[[1272,1290]],"Functions Start End":[[776,897],[1119,1164],[1231,1270]]}
{"Identifier":"2019AandA...623A.156D__Cacciari_&_Clementini_2003_Instance_1","Paragraph":"Cepheids and RR Lyrae stars are primary standard candles of the cosmological distance ladder because they follow canonical relations linking their intrinsic luminosity to the pulsation period and\/or the metal abundance. Specifically, for Cepheids the intrinsic luminosity (L) at any passband depends on the period (P) of light variation. This is traditionally referred to as the Cepheid period\u2013luminosity relation or Leavitt law, after its discoverer Mrs Henrietta Swan Leavitt (Leavitt & Pickering 1912). Modern realisations of the Cepheid PL relations from optical to infrared passbands include, among others, the ground-based studies of Madore & Freedman (1991), Ripepi et al. (2012), and Gieren et al. (2013), works based on Hubble Space Telescope (HST) data such as those by Freedman et al. (2001), Saha et al. (2006), and Riess et al. (2011), and theoretical investigations such as those by Marconi et al. (2005). Among the most recent studies of the Cepheid period\u2013luminosity relations are those based on Gaia trigonometric parallaxes of Galactic Cepheids (e.g. Clementini 2017; hereafter Paper I and references therein; Riess et al. 2016, 2018). For RR Lyrae stars the intrinsic luminosity (L) in the infrared passbands depends on P and possibly stellar metallicity (Z; PL \u2013 metallicity relation \u2013 PL(Z)), as first shown by Longmore et al. (1986) and later confirmed by (i) empirical studies of field and cluster RR Lyrae stars (e.g. Sollima et al. 2006, 2008; Borissova et al. 2009), (ii) theoretical models by Marconi et al. (2015) and Neeley et al. (2017), and (iii) the Gaia parallax-calibrated relations of Sesar et al. (2017), Paper I and references therein, and Muraveva et al. (2018a,b). In the visual passband, the luminosity L depends on Z in the form of the so-called RR Lyrae luminosity\u2013metallicity relation (see e.g. Cacciari & Clementini 2003; Clementini et al. 2003; the pulsation models by Bono et al. 2003; the theoretical calibration by Catelan et al. 2004; or the Gaia-based relations in Paper I; Muraveva et al. 2018a and references therein). The predicted precision of the Gaia end-of-mission parallaxes for local Cepheids and RR Lyrae stars1 will allow us to determine the slope and zero-point of these fundamental relations with unprecedented accuracy, thus setting the basis for a global reassessment of the whole cosmic distance ladder. As a first anticipation of the Gaia potential in this field of the cosmic distance ladder and a first assessment of improved precision with respect to previous astrometric missions such as HIPPARCOS, and the dramatic increase in statistics compared to what is achievable, for instance, through measuring parallaxes with the HST, Gaia DR1 published parallaxes for more than 700 Galactic Cepheids and RR Lyrae stars, computed as part of the Tycho-Gaia Astrometric Solution (TGAS; Lindegren et al. 2016). A number of papers after Gaia intermediate data releases in 2016 and 2018 (Gaia Data Release 1 \u2013 DR1 and DR2, respectively) have discussed Gaia Cepheids and RR Lyrae stars, specifically presenting the released samples (Clementini et al. 2016, 2019), their parallaxes (e.g. Lindegren et al. 2016) and possible offsets affecting them (Arenou et al. 2017, 2018); and addressing in particular their use as standard candles (Casertano et al. 2017 and Paper I for Gaia DR1 and Riess et al. 2018; Muraveva et al. 2018a for Gaia DR2). In Paper I we have used TGAS parallaxes, along with literature photometry and spectroscopy, to calibrate the zero-point of the PL relations of classical and type II Cepheids, and the near-infrared PL and PL(Z) relations of RR Lyrae stars by fitting these relations through adopting different techniques that operate either in parallax or absolute magnitude space. In that paper, different sources of biases affecting the TGAS samples of Cepheids and RR Lyrae stars were discussed at some length, and the possible systematic errors caused in the inferred luminosity calibrations were analysed in detail.","Citation Text":["Cacciari & Clementini 2003"],"Functions Text":["In the visual passband, the luminosity L depends on Z in the form of the so-called RR Lyrae luminosity\u2013metallicity relation (see e.g."],"Functions Label":["Background"],"Citation Start End":[[1838,1864]],"Functions Start End":[[1704,1837]]}
{"Identifier":"2016MNRAS.458.3181C__Trujillo_et_al._2011_Instance_1","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"],"Functions Text":["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"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[866,886]],"Functions Start End":[[450,864]]}
{"Identifier":"2020ApJ...903....8H__Finn_1987_Instance_1","Paragraph":"We note that during these encounters, energy is also lost due to tidal oscillations in the neutron star excited by the black hole (e.g., Press & Teukolsky 1977), and contributes to \u03c3. To check whether we should include this effect in our calculations or whether it can be safely neglected, we approximate the tidal energy dissipated during a parabolic encounter according to the formalism presented in Press & Teukolsky (1977):\n6\n\n\n\n\n\nwhere RNS is the radius of the neutron star, Rmin is the periastron of the approach, and Tl are dimensionless values associated with each spherical harmonic l (see Press & Teukolsky 1977 for calculation of Tl). We only consider the quadrupole mode (l = 2), which dominates over the other modes (Press & Teukolsky 1977). We approximate the NS as a polytropic star of index n = 0.5 (e.g., Finn 1987), and use values from Table 1 of Kokkotas & Schafer (1995) to aid in the calculation of Tl. Note that since there is a minimum impact parameter, there is a minimum possible value of Rmin. For a parabolic encounter the relationship between the impact parameter and the periastron distance is:\n7\n\n\n\n\n\nThus, combining Equations (4) and (7), we find the minimum possible Rmin to be:\n8\n\n\n\n\n\nEncounters with Rmin  Rmin(bmin) will result in a direct collision between the BH and NS instead of a bound binary. In Figure 1 we plot \n\n\n\n\n\n\u2014the ratio of energy lost to tidal oscillations to the energy lost to gravitational waves\u2014as a function of Rmin, for a 5 \n\n\n\n\n\n BH and a 1.4 \n\n\n\n\n\n NS. We have also marked the region where \n\n\n\n\n\n. We see that in the region of interest where \n\n\n\n\n\n, i.e., where bound binaries can form, \n\n\n\n\n\n, an extremely small value. We have verified (not shown to avoid clutter) that for larger BH masses, \n\n\n\n\n\n is even smaller. This is consistent with previous studies about NS\u2013NS captures (e.g., Gold et al. 2012; Chirenti et al. 2017). Thus, we can safely neglect tides in our calculation of the capture cross section.","Citation Text":["Finn 1987"],"Functions Text":["We approximate the NS as a polytropic star of index n = 0.5 (e.g.,"],"Functions Label":["Uses"],"Citation Start End":[[822,831]],"Functions Start End":[[755,821]]}
{"Identifier":"2021AandA...655A..99D__Carigi_et_al._2005_Instance_3","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"],"Functions Text":["The turning point of increased relative production of carbon from massive stars takes place at A(O) ~ 8.7 dex (see Fig. 2 of","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)."],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Differences"],"Citation Start End":[[2266,2284]],"Functions Start End":[[2141,2265],[2286,2318],[2319,2469]]}
{"Identifier":"2022ApJ...924...42N__Cheng_et_al._1990_Instance_2","Paragraph":"It is generally thought that the emission from radio to medium energy gamma rays is generated by the injected electrons through the synchrotron radiative mechanism. The high-energy photon emission mainly comes from inverse Compton (IC) scattering of the high-energy electrons on the background seed photons, which include the synchrotron background, the cosmic microwave background, and infrared photons in the PWNe (see, e.g., Zhang et al. 2008; Fang & Zhang 2010; Torres et al. 2013; Lu et al. 2020). On the other hand, it is also suggested that the gamma rays could be emitted by the hadronic processes. The relativistic protons accelerated in the Crab pulsar outer gap interact with the matter inside the nebula. and this process may contribute in the high-energy gamma-ray range (see, e.g., Cheng et al. 1990; Bednarek & Protheroe 1997; Khangulyan et al. 2020; Cao et al. 2021). Therefore, it has been long debated whether the high-energy emission from the PWNe is the leptonic or hadronic origin. The details of the high-energy radiation produced by leptonic process have been discussed for the Crab Nebula (see, e.g., Venter & de Jager 2007; Zhang et al. 2008; Mart\u00edn et al. 2012), and that of the gamma-ray emission about the hadronic process have been also investigated (see, e.g., Cheng et al. 1990; Bednarek & Protheroe 1997; Bednarek 2003, 2007). However, with the establishment of more and more high-energy observatories, some telescopes have possessed the performance of observing photons of exceeding to the PeV from the astronomic objects. An increasing number of observational data has been reported by the different experiments. For example, the Amenomori et al. (2019) reported that the Tibet air shower array with the underground water-Cerenkov-type muon detector array observed the highest energy photons of exceeding 100 TeV with a 5.6\u03c3 statistical significance and pointed the measured spectrum with energy extended to the sub-PeV from the Crab Nebula have an absence of high-energy cutoff. Recently, more than 530 photons at energies above 100 TeV and up to 1.4 PeV from the 12 ultra-high-energy gamma-ray sources with a statistical significance greater than seven standard deviations were reported again by LHAASO (Cao et al. 2021). Together with the earlier investigations about the leptonic scenario, the radiative spectrum from the leptons has a cutoff around the sub-PeV region (see, e.g., Zhang et al. 2008; Mart\u00edn et al. 2012; Zhang et al. 2020). It seems that the other components of gamma rays have a significant contribution.","Citation Text":["Cheng et al. 1990"],"Functions Text":["and that of the gamma-ray emission about the hadronic process have been also investigated (see, e.g.,","However, with the establishment of more and more high-energy observatories, some telescopes have possessed the performance of observing photons of exceeding to the PeV from the astronomic objects."],"Functions Label":["Background","Motivation"],"Citation Start End":[[1291,1308]],"Functions Start End":[[1189,1290],[1359,1555]]}
{"Identifier":"2019MNRAS.489.2355D__Chevallard_et_al._2018a_Instance_1","Paragraph":"As a comparison, we also use the publicly available mock catalogue JAdes extraGalactic Ultradeep Artificial Realizations (JAGUAR; Williams et al. 2018) to derive the relation between UV and [O iii] + H\u2009\u03b2 luminosity of simulated z \u223c 8 galaxies. The JAGUAR mock catalogue has been produced by matching luminosity and stellar mass functions as well as the relation between the stellar mass and UV luminosity, mostly at z \u2264 4. The galaxy properties are then extrapolated up to z \u223c 15. The JAGUAR catalogue provides emission line fluxes and EWs for the main lines based on modelling with the beagle code (Chevallard & Charlot 2016; Chevallard et al. 2018a). We identify all galaxies from the fiducial JAGUAR mock in the redshift range 7.11  z  9.05 and we randomly select 1000 of them to match the absolute UV magnitude distribution of our sample, and then fit the UV-[O iii] + H\u2009\u03b2 luminosity data. The result is shown in red in Fig. 6. Similarly to our sample, the z \u223c 8 galaxies from the JAGUAR catalogue exhibit a tight relation between UV and [O iii] + H\u2009\u03b2 luminosity (Spearman rank correlation coefficient \u03c1 = 0.73, standard deviation from null hypothesis \u03c3 > 40). However, the mock galaxies exhibit a significantly lower [O iii] + H\u2009\u03b2 luminosity (\u223c0.5 dex) at a given LUV compared to the relation of our galaxies. The detailed reason for this discrepancy relative to the JAGUAR mock is unclear, but one possible reason is differences in the median physical properties. For instance, while the mock galaxies exhibit (3.6\u20134.5)$\\mu$m colour similar to the ones from our sample at a given UV luminosity, the average F125W-3.6$\\mu$m colour in JAGUAR is smaller by \u223c0.3 mag compared to the observed F125W-3.6$\\mu$m colour in our sample. This means that while (3.6\u20134.5)$\\mu$m colour and EW([O iii] + H\u2009\u03b2) are on average similar between JAGUAR and our sample, the absolute [O iii] + H\u2009\u03b2 line luminosity scales with the 3.6$\\mu$m flux which is larger in our sample compared to the JAGUAR mock catalogue. Furthermore, JAGUAR models a small field comparatively to our data, therefore the overlap in UV luminosity is small.","Citation Text":["Chevallard et al. 2018a"],"Functions Text":["The JAGUAR catalogue provides emission line fluxes and EWs for the main lines based on modelling with the beagle code"],"Functions Label":["Uses"],"Citation Start End":[[627,650]],"Functions Start End":[[481,598]]}
{"Identifier":"2022ApJ...926..151Z__Jennings_et_al._2020_Instance_1","Paragraph":"Unlike the CMB, the 21 cm signal is highly non-Gaussian, because patchy, bubble-like structures of ionized hydrogen (H ii) regions are produced surrounding the ionizing sources. Thus, there is potentially a wealth of information in the 21 cm signal that is not contained in the 21 cm power spectrum, a two-point statistics of 21 cm brightness temperature fluctuations that is traditionally well studied in the literature. It is therefore essential to develop new methods that maximally exploit the full information in the 3D 21 cm images obtained by the SKA. Toward this goal, conventionally, new summary statistics that can only be measured with imaging have been proposed. These include the three-point correlation function (Hoffmann et al. 2019; Jennings et al. 2020), bispectrum (Yoshiura et al. 2015; Shimabukuro et al. 2016, 2017; Majumdar et al. 2018, 2020; Hutter et al. 2020; Saxena et al. 2020; Kamran et al. 2021), one-point statistics (Harker et al. 2009; Shimabukuro et al. 2015; Gorce et al. 2021), topological quantities such as the Minkowski functionals (Gleser et al. 2006; Chen et al. 2019; Kapahtia et al. 2021) and Betti numbers (Giri & Mellema 2021), the cross correlation between the 21 cm line and other probes, such as the CO line (Gong et al. 2011; Lidz et al. 2011), the C ii line (Gong et al. 2012; Beane & Lidz 2018), the kinetic Sunyaev\u2013Zel\u2019dovich (kSZ) effect (Ma et al. 2018; La Plante et al. 2020), and novel techniques such as the antisymmetric cross correlation between the 21 cm line and CO line (Zhou et al. 2021). Since those summary statistics are fully determined by the parameters in the reionization models (hereafter \u201creionization parameters\u201d), in principle, Monte Carlo Markov Chain (MCMC) methods can be employed to constrain the reionization parameters from measurements of those statistics with futuristic 21 cm experiments (see, e.g., Watkinson et al. 2022), just as the MCMC analysis with the 21 cm power spectrum (Greig & Mesinger 2015, 2017, 2018).","Citation Text":["Jennings et al. 2020"],"Functions Text":["Toward this goal, conventionally, new summary statistics that can only be measured with imaging have been proposed. These include the three-point correlation function"],"Functions Label":["Background"],"Citation Start End":[[749,769]],"Functions Start End":[[559,725]]}
{"Identifier":"2019ApJ...871..176X__Eldridge_et_al._2013_Instance_2","Paragraph":"The progenitors of SNe Ib\/c have been thought to be Wolf-Rayet (W-R) stars with high initial masses (MZAMS \u2273 25 M\u2299; Crowther 2007). Before core collapse, these stars usually have experienced severe mass loss through strong stellar winds or due to interaction with companion stars (van der Hucht 2006; Paxton et al. 2015). As the evolution of massive stars is usually dominated by binary evolution (Heger et al. 2003) and also depends largely on metallicity, rotation, and so on (Heger et al. 2003; Georgy et al. 2013, 2012), this makes the direct identification of their progenitors complicated (Smartt 2015). However, there are increasing studies suggesting that a lower-mass binary scenario is more favorable for most SNe Ib\/c, considering the measured low ejecta masses (Eldridge et al. 2013; Lyman et al. 2016). In addition, the H\/He envelopes of the progenitor stars are stripped by binary interaction. There are many detections of progenitor stars for SNe II. For example, most SNe IIP are found to originate from red supergiants (Smartt et al. 2009), while SNe IIL are typically from progenitors with somewhat warmer colors (see Smartt 2015, for a review), and SNe IIb are from those with higher effective temperatures such as yellow supergiants that have had their H\/He envelopes partially stripped through binary interaction (e.g., SN 1993J; Podsiadlowski et al. 1993; Maund et al. 2004; Fox et al. 2014). Until recently, there has been only one report of the possible identification of a progenitor star for SNe Ib, namely, iPTF 13bvn, which was proposed to spatially coincide with a single W-R-like star identified on the pre-explosion Hubble Space Telescope (HST) images (Cao et al. 2013; Groh et al. 2013). But such an identification is still controversial (e.g., Bersten et al. 2014; Fremling et al. 2014; Eldridge et al. 2015; Eldridge & Maund 2016). Direct detection of progenitor stars is still elusive for SNe Ic, which prevents us from further testing the theoretical evolution of massive stars (Eldridge et al. 2013).","Citation Text":["Eldridge et al. 2013"],"Functions Text":["Direct detection of progenitor stars is still elusive for SNe Ic, which prevents us from further testing the theoretical evolution of massive stars"],"Functions Label":["Background"],"Citation Start End":[[2015,2035]],"Functions Start End":[[1866,2013]]}
{"Identifier":"2016ApJ...828...41S__Chatterjee_et_al._2004_Instance_1","Paragraph":"Figure 1 shows the main features of the butterfly diagram: (1) the onset of the cycle at mid-latitudes; (2) the sunspot drift toward the equator and its slowdown represented by a change in the slope of the butterfly wing (Maunder 1904; Li et al. 2001); (3) the tail-like attachment over the minimum phase that is more prominent when the activity is stronger, which might lead to the overlap of successive cycles; (4) the length of the overlap varies within 1\u20132 years. It characterizes only the minimum phase and it is confined at latitudes \u226415\u00b0 (Cliver 2014). This feature is also seen in torsional oscillations shown in the bottom panels of Figure 1 (Wilson et al. 1988; Howe et al. 2009). The rate of drift of sunspots toward the equator slows as the sunspot band approaches the equator, and halts at about 8\u00b0 latitude (Hathaway et al. 2003). The end of the migration does not correspond to the end of the activity because it produces the tail-like attachment. When the new cycle at mid-latitudes starts before the end of the old cycle at low latitudes, it causes successive cycles to overlap. FTD models driven only by the Babcock\u2013Leighton mechanism (Chatterjee et al. 2004), or along with the \u03b1-turbulent effect operating in the bulk of the convection zone, currently have the best agreement with observations (Passos et al. 2014), because the length of the simulated overlap is short and it occurs only during the minimum at low latitudes. Conversely, thin-shell dynamo wave models (Moss & Brooke 2000; Sch\u00fcssler & Schmitt 2004; Bushby 2006) or the thin-shell flux transport dynamo (Dikpati & Gilman 2001) tend to produce dynamo waves with too short a wavelength, leading to excessive overlap between adjacent cycles because this involves a wider range of latitudes. Furthermore they also fail to reproduce the tail-like attachment over the minimum phase. Moreover the direction of the migration of activity could also provide information on the nature of the \u03b1 mechanism. Both formalisms make strong assumptions to initiate the sunspot cycle at mid-latitudes. The Babcock\u2013Leighton FTD models assume that the deep equatorward meridional flow penetrates slightly below the convection zone to a greater depth than usually believed (Nandy & Choudhuri 2002), in order to prevent the onset and occurrence of a sunspot cycle above 45\u00b0 as well as any other kind of cyclic activity. The same result is achieved with the \u03b1\u03a9 dynamo wave by inhibiting the \u03b1-turbulent effect at higher latitudes (Sch\u00fcssler & Schmitt 2004). Based on these assumptions, the magnetic activity in any type of FTD model starts at higher latitudes and then propagates only equatorward, while in the thin-shell \u03b1\u03a9 dynamo wave the magnetic activity can propagate poleward as well as equatorward (Bushby 2006). These two branches are also clearly seen in the torsional oscillation pattern (e.g., Howe et al. 2009). This results from the solar-like differential profile, which is characterized by a sign change in \n\n\n\n\n\n in the tachocline at high and low latitudes (Ruediger & Brandenburg 1995). This sign change, however, has not yet been confirmed by helioseismic observations.","Citation Text":["Chatterjee et al. 2004"],"Functions Text":["FTD models driven only by the Babcock\u2013Leighton mechanism","currently have the best agreement with observations","because the length of the simulated overlap is short and it occurs only during the minimum at low latitudes."],"Functions Label":["Similarities","Similarities","Similarities"],"Citation Start End":[[1154,1176]],"Functions Start End":[[1096,1152],[1262,1313],[1336,1444]]}
{"Identifier":"2018AandA...619A.105T__Matt_et_al._2015_Instance_1","Paragraph":"We have presented MoCA, a Monte Carlo code for Comptonisation in Astrophysics which includes polarisation. To our knowledge MoCA is the first code operating with single photons and including all special relativity and quantum effects. The main disadvantage of this approach is the long computing time, which implies the need to parallelise the code on clusters of computers. The advantage with respect to pure analytical models such as those available in XSPEC is that we can explore the totality of the parameters space for the Comptonising medium (i.e. thermal energy and optical thickness of the corona) without any restriction and this approach will also allow a better understanding of the whole process. We also included all corrections such as Klein\u2013Nishina cross-section and scattering angle distribution. These effects, small below 100 keV, must nonetheless be taken into account when inferring the thermal energy of the corona from observations. In some sources it has been found that coronae can have extremely high energy cut-off (e.g. NGC 5506, Matt et al. 2015) and therefore thermal energy, which is inferred by measuring the curvature of NuSTAR spectra at high energy and in this context K\u2013N effects cannot be neglected. From the polarimetric point of view we did not seen any deviation due to K\u2013N effects, but this was to expected as we focussed our attention below 100 keV where these effects are small. However, one can imagine a scenario in which the thermal energy of the corona is few tens of keV and in that case we expect to see a difference both on the spectrum and the polarisation but we defer such investigation to future papers focussed on the exploration of the coronal parameters space. In its actual form the code is fast enough to explore different geometries of the corona with different parameters. Spectra can then be compared with those obtained by NuSTAR to derive coronal parameters, especially in the high optical depth regime where analytical models are not reliable. As already mentioned, much observational evidence points in the direction of compact coronae above or around the compact object. In order to properly treat such coronae, gravitational effects must be taken into account. We have recently included a ray-tracing routine to take into account GR effects: this new version of MoCA, and applications of the code to different astrophysical scenarios, will be discussed in future papers. Nonetheless we have shown the geometrical effect of more compact coronae on the spectra and the polarisation signal: the spectra become softer as the corona shrinks and the polarisation changes dramatically as we approach a more symmetrical shape of the corona. The study we performed will also be useful to quantify the impact of GR effects on compact coronae with respect to a purely geometrical effect.","Citation Text":["Matt et al. 2015"],"Functions Text":["In some sources it has been found that coronae can have extremely high energy cut-off (e.g. NGC 5506,","and therefore thermal energy, which is inferred by measuring the curvature of NuSTAR spectra at high energy and in this context K\u2013N effects cannot be neglected."],"Functions Label":["Background","Background"],"Citation Start End":[[1058,1074]],"Functions Start End":[[956,1057],[1076,1236]]}
{"Identifier":"2020ApJ...899..147F__Venot_et_al._2015_Instance_2","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":["Venot et al. 2015"],"Functions Text":["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","Therefore, they can be considered as better candidates for the formation of complex organic molecules with longer carbon chains."],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[1624,1641]],"Functions Start End":[[1277,1622],[1644,1772]]}
{"Identifier":"2021MNRAS.500.1772N__Fernandez_et_al._2015_Instance_1","Paragraph":"While these early studies demonstrated the viability of neutron star mergers as a major r-process site, they identified only one ejection channel: \u2018dynamical ejecta\u2019 that are tidally flung out by gravitational torques. Since they are never substantially heated, these ejecta carry their original \u03b2 \u2212equilibrium electron fraction from the original neutron star, Ye \u2248 0.05, and this enormous neutron-richness allows them to undergo a \u2018fission cycling\u2019 process (Goriely, Bauswein & Janka 2011; Korobkin et al. 2012), which produces a very robust r-process abundance distribution close to the solar pattern for A \u2265 130, but hardly any lighter r-process elements. Oechslin, Janka & Marek (2007) pointed out that there is a second channel of mass ejection that also happens on a dynamical time-scale: shock-heated matter from the interface where the stars come into contact. As of today, many more mass ejection channels have been discussed: matter that becomes unbound on secular time-scales (\u223c1 s) from the post-merger accretion torus (Beloborodov 2008; Metzger, Piro & Quataert 2008; Fernandez & Metzger 2013; Fernandez et al. 2015; Just et al. 2015; Siegel & Metzger 2017, 2018; Fernandez et al. 2019; Miller et al. 2019a), as MHD-driven winds (Siegel & Ciolfi 2015) and by viscous effects (Shibata, Kiuchi & Sekiguchi 2017; Radice et al. 2018a; Shibata & Hotokezaka 2019) from a long-lived neutron star merger remnant. Similar to the case of proto-neutron stars, the enormous neutrino luminosities (>1053 erg s\u22121) after a neutron star merger can also drive substantial matter outflows (Ruffert et al. 1997; Rosswog & Ramirez-Ruiz 2002; Dessart et al. 2009; Perego et al. 2014; Martin et al. 2015; Radice et al. 2018b). The secular torus ejecta contain approximately 40 per cent of the initial torus mass and, although the latter may vary substantially from case to case, they likely contribute the lion\u2019s share to the total ejecta mass. Due to their different thermal histories and exposure times to neutrinos, the ejecta channels can have different electron fractions Ye and therefore different nucleosynthesis yields.1 For electron fractions below a critical value, $Y_{\\rm e}^{\\rm crit}\\approx 0.25$ (Korobkin et al. 2012; Lippuner & Roberts 2015), lanthanides and actinides are efficiently produced, which, due to their open f-shells, have particularly high bound\u2013bound opacities (Barnes & Kasen 2013; Kasen, Badnell & Barnes 2013; Tanaka & Hotokezaka 2013; Tanaka et al. 2020) and therefore lead to red transients that peak days after the merger. Ejecta with electron fractions above $Y_{\\rm e}^{\\rm crit}$, in contrast, only produce \u2018lighter\u2019 elements with lower opacities and thus result in bluer transients that peak after about 1 d. Opaque, low-Ye ejecta blocking the view on high-Ye ejecta can lead to a \u2018lanthanide curtaining\u2019 effect (Kasen, Fern\u00e1ndez & Metzger 2015; Wollaeger et al. 2018), which will efficiently block blue light. Therefore, it is important to understand the layering, dynamics, interaction and potential mixing of different ejecta channels.","Citation Text":["Fernandez et al. 2015"],"Functions Text":["As of today, many more mass ejection channels have been discussed: matter that becomes unbound on secular time-scales (\u223c1 s) from the post-merger accretion torus"],"Functions Label":["Background"],"Citation Start End":[[1107,1128]],"Functions Start End":[[869,1030]]}
{"Identifier":"2022MNRAS.515.1795B__Yang_et_al._2021_Instance_2","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"],"Functions Text":["and do not fully rule out a spatially non-flat Universe"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1392,1408]],"Functions Start End":[[1258,1313]]}
{"Identifier":"2021ApJ...920..145H__Damone_et_al._2018_Instance_2","Paragraph":"Over the past decade, many attempts to address this issue have been carried out, such as from the perspective of conventional nuclear physics and even exotic physics beyond the standard BBN framework (Angulo et al. 2005; Cyburt et al. 2008, 2016; Boyd et al. 2010; Pospelov & Pradler 2010; Fields 2011; Kirsebom & Davids 2011; Wang et al. 2011; Broggini et al. 2012; Coc et al. 2012, 2013, 2014; Cyburt & Pospelov 2012; Kang et al. 2012; Voronchev et al. 2012; Bertulani et al. 2013; Hammache et al. 2013; He et al. 2013; Kusakabe et al. 2014; Pizzone et al. 2014; Yamazaki et al. 2014; Hou et al. 2015, 2017; Famiano et al. 2016; Damone et al. 2018; Hartos et al. 2018; Luo et al. 2019; Rijal et al. 2019; Clara & Martins 2020). However, despite the fact some solutions using exotic physics have succeeded in resolving this issue, it appears there is still no universally accepted solution in the academic community since validations of these mysterious exotic physics are beyond the capabilities of current science. Conversely, it seems more worthwhile to exclude any potential possibility of resolving the 7Li discrepancy from the perspective of nuclear physics. It is known that the majority of the primordial 7Li production arises from the decay of 7Be by electron capture during the 2 months after BBN stops. Thus, for the solution of the Li problem, reactions involving 7Be could be more significant than those involving 7Li. Therefore, many reactions that potentially destroy 7Be were investigated to solve this discrepancy over past 10 yr (Kirsebom & Davids 2011; Broggini et al. 2012; Hammache et al. 2013; Hou et al. 2015; Hartos et al. 2018). Meanwhile, enormous efforts have been made to refine the reaction rates of key BBN reactions in the past 20 yr (Smith et al. 1993; Descouvemont et al. 2004; Serpico et al. 2004; Cyburt & Davids 2008; Neff 2011; Pizzone et al. 2014; Tumino et al. 2014; Hou et al. 2015; Barbagallo et al. 2016; Iliadis et al. 2016; Kawabata et al. 2017; Lamia et al. 2017, 2019; Damone et al. 2018; Rijal et al. 2019; Mossa et al. 2020), but the probability of solving or alleviating the 7Li problem by improving our knowledge of relevant nuclear reaction rates still cannot be eliminated. Recent experiments for key nuclear reactions like 7Be(n,p)7Li and 7Be(d,p)24He allow for a reduction of the 7Li production by about 12% (Damone et al. 2018; Rijal et al. 2019) compared to previous calculations. At present, nuclear uncertainties cannot rule out that some of the reactions destroying 7Li are indeed more efficient than those currently used (Boyd et al. 2010; Chakraborty et al. 2011).","Citation Text":["Damone et al. 2018"],"Functions Text":["Meanwhile, enormous efforts have been made to refine the reaction rates of key BBN reactions in the past 20 yr","but the probability of solving or alleviating the 7Li problem by improving our knowledge of relevant nuclear reaction rates still cannot be eliminated."],"Functions Label":["Background","Motivation"],"Citation Start End":[[2016,2034]],"Functions Start End":[[1655,1765],[2075,2226]]}
{"Identifier":"2017MNRAS.469S.238L__Massironi_et_al._2015_Instance_1","Paragraph":"ESA's Rosetta spacecraft orbited closely around comet 67P\/Churyumov\u2013Gerasimenko (hereafter 67P) during its 2 yr mission. From 2014 August to 2016 September, the camera system Optical, Spectroscopic, and Infrared Remote Imaging System (OSIRIS; Keller et al. 2007) onboard Rosetta observed the comet nucleus down to few centimetres per pixel providing detailed images of the comet's surface and activity. The study of 67P morphology has been performed with unprecedented detail thanks to the OSIRIS image spatial coverage and resolution. A first OSIRIS analysis on the nucleus structure, morphology and jets activity is reported in Sierks et al. (2015) and Thomas et al. (2015). The comet's surface revealed a wide diversity in morphology such as layering (Massironi et al. 2015), pits (Vincent et al. 2015), boulders and fractures (Pajola et al. 2015, 2016c; El-Maarry et al. 2015a), and high reflectivity boulders (Pommerol et al. 2015). During the entire mission, surface changes have been detected thanks to images acquired before and after perihelion passage (Groussin et al. 2015; Pajola 2017; El-Maarry 2017). The global spectrophotometric properties of the comet were investigated in detail by Fornasier et al. (2015) identifying three different groups of terrains depending on spectral slope values. In addition to a global characterization of the comet, OSIRIS allowed a detailed analysis of specific regions identified on the comet by means of high resolution and multifilter images (La Forgia et al. 2015; Deshapriya et al. 2016; Lucchetti et al. 2016; Oklay et al. 2016a; Pajola et al. 2016d; Oklay et al. 2017). We decided to follow the approach of characterizing at the highest resolution possible specific regions of interest (ROIs) on 67P. We therefore performed a detailed analysis on an area belonging to the Seth region (El-Maarry et al. 2015b) that is characterized by flat-floored and steep-walled circular depressions (El-Maarry et al. 2015b; Giacomini et al. 2016). We focused our attention on the circular niches of the Seth area (Fig. 1) and performed a multidisciplinary study of these structures investigating their geomorphological and spectrophotometric properties. In addition, thanks to images acquired pre- and post-perihelion, we conducted a comparative analysis to find if this area has been subjected to surface changes. This can be useful to provide constraints about Seth's niches properties as well as their possible origin being one of the interesting features located on the surface of 67P. Specifically, the circular niches deposits can be considered the result of landslide events that occurred on the comet surface, as the recent Aswan cliff collapse (Pajola 2017), where it has been reported the occurring of falling material from the adjacent cliff after the Rosetta perihelion passage. Hence, with this work we plan to understand if the formation of these Seth's circular niches occurred recently or if it is correlated to older events that have shaped the comet's surface.","Citation Text":["Massironi et al. 2015"],"Functions Text":["The comet's surface revealed a wide diversity in morphology such as layering"],"Functions Label":["Background"],"Citation Start End":[[755,776]],"Functions Start End":[[677,753]]}
{"Identifier":"2018AandA...616A..34H__Mohamed_&_Podsiadlowski_(2012)_Instance_2","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)"],"Functions Text":["The morphology resulting from this particular type of wind\u2013binary interaction is shown in Fig. 3 in"],"Functions Label":["Background"],"Citation Start End":[[1510,1540]],"Functions Start End":[[1410,1509]]}
{"Identifier":"2020ApJ...898....4C__Shipp_et_al._2018_Instance_1","Paragraph":"Detecting the halo response to the LMC-induced DM wake would be an exciting advancement in testing our assumptions about the properties of DM, as well as providing key constraints on the potential of the MW and the mass and orbital history of the LMC. However, the GC19 simulations give predictions for the response in the context of smooth MW DM and stellar halos. In reality, the MW stellar halo contains a wealth of substructure that is not yet phase mixed, in the form of stellar streams (e.g., Odenkirchen et al. 2001; Newberg et al. 2002; Belokurov et al. 2006; Grillmair 2006; Shipp et al. 2018; also see Newberg & Carlin 2016 for a recent review) and stellar clouds (e.g., Newberg et al. 2002; Rocha-Pinto et al. 2004; Juri\u0107 et al. 2008; Li et al. 2016). In addition, using a sample of MW halo main-sequence turnoff stars from the HALO7D survey (Cunningham et al. 2019a), Cunningham et al. (2019b) observed that the estimated parameters of the velocity ellipsoid (i.e., \n\n\n\n\n\n) were different in the different survey fields; these differing estimates could be interpreted as evidence that the halo is not phase mixed over the survey range (\n\n\n\n\n\n kpc). They also showed maps of the halo velocity anisotropy \u03b2 in two halos from the Latte suite of FIRE-2 simulations (introduced in Wetzel et al. 2016), finding that the anisotropy can vary over the range \n\n\n\n\n\n across the sky. Some of the variation in the \u03b2 estimates appeared to correlate with stellar overdensities in the halos, indicating that galactic substructure is at least in part responsible for the different velocity distributions. While some substructure in the halo can be clearly identified as overdensities in phase space and removed from analysis, the presence of velocity substructure in the halo could complicate attempts to detect signatures of the LMC-induced DM wake. For example, Belokurov et al. (2019) recently argued that the Pisces Overdensity (Sesar et al. 2007; Watkins et al. 2009; Nie et al. 2015) might be stars in the wake trailing the LMC in its orbit, because of their net negative radial velocities. However, it remains difficult to conclusively argue this scenario given that these stars could also be in Galactic substructure (or, perhaps, stars that are in substructure and have been perturbed by the DM wake).","Citation Text":["Shipp et al. 2018"],"Functions Text":["In reality, the MW stellar halo contains a wealth of substructure that is not yet phase mixed, in the form of stellar streams (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[584,601]],"Functions Start End":[[366,498]]}
{"Identifier":"2019ApJ...883...53T__Armano_et_al._2016_Instance_1","Paragraph":"LISA Pathfinder (LPF; Antonucci et al. 2011), a European Space Agency (ESA) mission that operated near the first Sun\u2013Earth Lagrange point (L1) from 2016 January through 2017 July, is in an ideal orbit to make such measurements. However, LPF flew no instrumentation dedicated to micrometeoroid or dust detection. LPF\u2019s primary objective was to demonstrate technologies for a future space-based observatory of millihertz-band gravitational waves. The key achievement of LPF was placing two gold-platinum cubes known as \u201ctest masses\u201d into a freefall so pure that it was characterized by accelerations at the femto-g level (e.g., Armano et al. 2016, 2018b), the level required to detect the minute disturbances caused by passing gravitational waves. In order to reach this level of performance, the test masses were released into cavities inside the spacecraft and a control system was employed to keep the spacecraft centered on the test masses. This control system was designed to counteract disturbances on the spacecraft, including those caused by impacts from micrometeoroids. Shortly before LPF\u2019s launch, it was realized that data from the control system, if properly calibrated, could be used to detect and characterize these impacts and infer information about the impacting particles (e.g., Thorpe et al. 2016). While such impact events have been reported by other spacecraft, LPF\u2019s unique instrumentation makes it sensitive to much smaller and much more numerous impacts and allows the impact geometry to be more fully constrained. Early results from the first few months of LPF operations suggested that such events could indeed be identified and were roughly consistent with the pre-launch predictions of their effect on the control system (e.g., Thorpe et al. 2017). In this paper we present results from the first systematic search for micrometeoroid impacts in the LPF data set. Our data set consists of 4348 hr of data in both the nominal LPF configuration and the \u201cDisturbance Reduction System\u201d (DRS) configuration, in which a NASA-supplied controller and thruster system took over control of the spacecraft (Anderson et al. 2018). Our data set corresponds to the times when LPF was operating in a \u201cquiet\u201d mode, without any intentional signal injections or other disturbances. During this period, we have identified 54 impact candidates using our detection pipeline and manual vetoing. We have characterized the properties of this data set and compared it to several theoretical models for the underlying dust population.","Citation Text":["Armano et al. 2016"],"Functions Text":["The key achievement of LPF was placing two gold-platinum cubes known as \u201ctest masses\u201d into a freefall so pure that it was characterized by accelerations at the femto-g level (e.g.,","the level required to detect the minute disturbances caused by passing gravitational waves."],"Functions Label":["Background","Background"],"Citation Start End":[[626,644]],"Functions Start End":[[445,625],[654,745]]}
{"Identifier":"2020MNRAS.498.6013A__Maga\u00f1a_et_al._2015_Instance_1","Paragraph":"On the other hand, observational data are used to test these models. Among the most frequently used are the cosmic microwave background radiation (CMB; Planck Collaboration XIII 2016; Aghanim et al. 2018), baryonic acoustic oscillations (BAO; Eisenstein et al. 2005; Blake et al. 2012; Alam et al. 2017; Bautista et al. 2017), Type Ia supernovae (SNe Ia; Scolnic et al. 2018), and observational Hubble data (OHD; Jimenez & Loeb 2002; Moresco et al. 2016; Maga\u00f1a et al. 2018). Consistency in the cosmological parameters among different techniques, rather than more accurate measurements, is desirable to better understand the nature of DE. In the last years, several efforts have been made by the community to include gravitational lens systems in the study of the Universe\u2019s evolution. Some of the pioneers are Futamase & Yoshida (2001) and Biesiada (2006), who used only one strong-lens system to study some of the most popular cosmological models. Grillo, Lombardi & Bertin (2008) introduced a methodology to estimate cosmological parameters using strong-lensing systems (SLS; see also Jullo et al. 2010; Maga\u00f1a et al. 2015, 2018). They apply the relation between the Einstein radius and the central stellar velocity dispersion, assuming an isothermal profile for the total density distribution of the lens (elliptical) galaxy. Their simulations found that the method is accurate enough to obtain information about the underlying cosmology. They concluded that the stellar velocity dispersion and velocity dispersion of the isothermal lens model are very similar in the w cold dark matter (wCDM) model. Biesiada, Pi\u00f3rkowska & Malec (2010) used the same procedure comparing a distance ratio, Dobs, constructed from SLS observations such as the Einstein radius and spectroscopic velocity dispersion of the lens galaxy, with a theoretical counterpart, Dth. By using a sample containing 20 SLS, they demonstrated that this technique is useful to provide insights into DE. Cao et al. (2012) updated the sample to 80 systems and proposed a modification that takes into account deviations from sphericity, i.e. from the singular isothermal sphere (SIS). Later on, Cao et al. (2015) considered lens profile deviations due to the redshift evolution of elliptical galaxies by using spherically symmetric power-law mass distributions for the lenses and also increased the compilation up to 118 points. They also explore the consequences of using aperture-corrected velocity dispersions on the parameter estimations. Some authors have pointed out the need for a sufficiently large sample to test DE models with higher precision (Yennapureddy & Melia 2018). For instance, Melia, Wei & Wu (2015) have emphasized that a sample of \u223c200 SLS can discern the Rh = ct model from the standard one. Qi et al. (2018) simulated strong lensing data to constrain the curvature of the Universe and found that, by increasing the sample (16000 lenses) and combining with compact radio quasars, it could be constrained with an accuracy of \u223c10\u22123. Recently, Leaf & Melia (2018) have revisited this cosmological tool with the largest sample of SLS (158) until now, including 40 new systems presented by Shu et al. (2017). The authors proposed a new approach to improve this technique by introducing in the observational distance ratio error (\u03b4Dobs), a parameter \u03c3x to take into account the SIE scatter and any other source of errors in the measurements. In their analysis, they excluded 29 SLS that are outside the region 0  Dobs  1, and the system SL2SJ085019\u2212034710 (Sonnenfeld et al. 2013b), which seems to be an extreme outlier for their models. Their results show that a $\\sigma _x = 12.2{{\\ \\rm per\\ cent}}$ provides more statistically significative cosmological constraints. Finally, Chen, Li & Shu (2018) used 157 SLS to analyse the Lambda cold dark matter (\u039bCDM) model. They considered a lens mass distribution \u03c1(r) = \u03c10r\u2212\u03b3 and three possibilities for the \u03b3 parameter: a constant value, a dependence with the lens redshift (zl), and a dependence with both the surface mass density and the lens redshift. They concluded that although \u03a90m, used as the only free parameter in \u039bCDM scenario, is very sensitive to the lens mass model, it provides weak constraints that are also in tension with Planck measurements.","Citation Text":["Maga\u00f1a et al. 2015"],"Functions Text":["Grillo, Lombardi & Bertin (2008) introduced a methodology to estimate cosmological parameters using strong-lensing systems (SLS; see also"],"Functions Label":["Background"],"Citation Start End":[[1107,1125]],"Functions Start End":[[950,1087]]}
{"Identifier":"2018ApJ...856...51R__Reale_2014_Instance_2","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"],"Functions Text":["However, if a single stellar loop were flaring, the cooling time of the confined plasma would be proportional to the loop length"],"Functions Label":["Uses"],"Citation Start End":[[2197,2207]],"Functions Start End":[[2048,2176]]}
{"Identifier":"2018ApJ...856..144M__McIntosh_et_al._2011_Instance_1","Paragraph":"The solar corona is still enigmatic from a scientific point of view, with important unsolved questions about its nature, such as solar wind acceleration and coronal heating (McComas et al. 2007; Parnell & De Moortel 2012). Although finding solutions to these questions is important on its own, it is also expected to have major implications for related fields, such as space weather (Singh et al. 2010), which is set to grow in importance as we make advances in technology and space exploration. One impediment toward solving coronal mysteries is the notorious difficulty in determining its key physical parameters, such as the magnetic field, by direct spectroscopic or polarimetric measurements. Early observational evidence of waves in coronal structures using SOHO\/TRACE (Aschwanden et al. 1999; Berghmans & Clette 1999) paved the way for the previously theorized coronal seismology (Uchida 1970; Roberts et al. 1984) to be a tool for coronal plasma diagnostics. The first attempt to seismologically determine the magnetic field of transversely oscillating coronal loops was applied to the first such event observed by Nakariakov & Ofman (2001). Since then, coronal seismology has successfully been applied to numerous oscillation events (for reviews, see, e.g., De Moortel 2005; De Moortel & Nakariakov 2012; Stepanov et al. 2012). A common feature of all previously diagnosed events is their localization in time, i.e., single events. This obviously limits the applicability of seismology to the brief duration of the oscillation event. Moreover, these single-oscillation events tend to be rare, as they are mostly related to flaring events or eruptions, meaning that seismology is restricted to brief diagnostics highly localized in time, or in time and space for non-global oscillations. However, the discovery of ubiquitous propagating transverse oscillations with CoMP (Tomczyk et al. 2007; Tomczyk & McIntosh 2009) and SDO (McIntosh et al. 2011) or the more recently identified nearly ubiquitous, decay-less low-amplitude kink coronal loop oscillations (Anfinogentov et al. 2013, 2015; Nistic\u00f2 et al. 2013), as well as oscillations in plumes (Thurgood et al. 2014), led to the possibility of continuous diagnostics in time, i.e., dynamic coronal seismology. Consequently, very recently, the first seismologic \u201cmagnetic field image\u201d was obtained, based on the ubiquitous transverse waves (Long et al. 2017), using a methodology put forth by Morton et al. (2015). In this study, the authors use the magnetohydrodynamic (MHD) kink phase speed of a flux tube (Edwin & Roberts 1983) as the inversion tool for the magnetic field. These ubiquitous propagating transverse waves are now widely regarded as Alfv\u00e9nic waves (Goossens et al. 2012), although this reinterpretation of the nature of the waves does not modify the phase speed formula used in the inversion. However, it is still assumed that the observed ubiquitous waves are transverse oscillations of flux tubes, and while the fine structure in the corona is still unknown (Peter et al. 2013; Reale 2014; Aschwanden & Peter 2017), it is unlikely that the magnetic cylinder model is satisfactory. Structuring across the magnetic field, among other factors (De Moortel & Pascoe 2009; Pascoe & De Moortel 2014), is an important detail for seismology, as it can greatly influence the nature and propagation of MHD waves (Luna et al. 2010; Terradas et al. 2010; Verth et al. 2010), altering the dependence of the observed phase speed on physical properties such as the magnetic field or mass density (e.g., Verth et al. 2007; Arregui et al. 2013), which might lead to erroneous inversions if not considered. Some light was shed recently on the weaknesses of modeling the corona as a bundle of independent thin magnetic strands by Magyar & Van Doorsselaere (2016). In their simulations, a loop consisting of packed strands is quickly deformed and mixed when disturbed by propagating transverse waves, leading to a turbulent cross-section (Magyar et al. 2017). This result reiterates the above-mentioned need to move away from rigid cylindrical models in favor of more realistic descriptions that account for nonlinearities.","Citation Text":["McIntosh et al. 2011"],"Functions Text":["However, the discovery of ubiquitous propagating transverse oscillations with CoMP","and SDO","led to the possibility of continuous diagnostics in time, i.e., dynamic coronal seismology."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[1935,1955]],"Functions Start End":[[1796,1878],[1926,1933],[2177,2268]]}
{"Identifier":"2018MNRAS.478.2541F__Bertin_&_Arnouts_1996_Instance_1","Paragraph":"The distance between the peak of the main source and the transient source (blue POSS I) is \u22486\u2009pixels or \u22486\u2009arcsec (\u22480.3\u2009kpc) towards the North, at an angle of \u2248\u2212 10\u00b0 (Fig. 2; left-hand column; rows 4\u20135). From the PSF characteristics of the blue POSS I and blue subtracted POSS images (Section 2.1.1), the transient source is slightly resolved ( 2 \u00d7 PSF FWHM). Because of the faintness of the transient and its proximity to the larger, brighter main source, an automated detection algorithm such as SExtractor (Bertin & Arnouts 1996), implemented in GAIA36 2016A, was unsuccessful in deblending the two sources. Therefore, a more rudimentary method was employed to determine the magnitude of the transient: using ds9, the transient and the brightest North-western source (Sections 2.2 and 2.3; Fig. 2; left-hand column; rows 4\u20135; Fig. 3) were fitted with ellipses to perform relative aperture photometry on the photometrically calibrated unconvolved POSS images, after sky subtraction (see below). This procedure shows that the transient source corresponds to \u227396\u2009per\u2009cent of the flux of the North-western source (m \u2248 20.9\u2009mag; USNO B1.0; Monet et al. 2003), which is equivalent to the transient being \u0394m \u2272 0.05\u2009mag fainter than the North-western source (Fig. 2; left-hand column; rows 4\u20135). Similarly, the flux ratio between the main source and the transient is approximately a factor of 6 (Section 2.1.4). Given this very simplistic procedure to obtain the transient photometry, the magnitude of the transient is quite uncertain. A magnitude of m \u2248 21.0\u00b10.2\u2009mag will be adopted in the following discussion (Section 3), with the confidence interval provided assuming a 20\u2009per\u2009cent error in the flux (detection at a 5\u03c3 level; see below), corresponding to the estimated photometric error at the maximum intensity extended to the full source, and neglecting unquantifiable systematics errors; the error estimation is included in AppendixA. Favourably, changes in the magnitude as large as \u0394m \u2248 1.5 will not significantly alter the conclusions of this work (Section 3).","Citation Text":["Bertin & Arnouts 1996"],"Functions Text":["Because of the faintness of the transient and its proximity to the larger, brighter main source, an automated detection algorithm such as SExtractor","implemented in GAIA36 2016A, was unsuccessful in deblending the two sources."],"Functions Label":["Uses","Compare\/Contrast"],"Citation Start End":[[510,531]],"Functions Start End":[[360,508],[534,610]]}
{"Identifier":"2020MNRAS.498.4906B__Stinson_et_al._2013_Instance_1","Paragraph":"Throughout their short lifetimes, high-mass stars (>8\u2009M\u2299) inject large amounts of energy and momentum into their host environments through a variety of feedback processes (e.g. Krumholz et al. 2014). The most potentially disruptive of these feedback mechanisms occurs when the stars eventually die, exploding as supernovae (SNe). Indeed, SNe are thought to play a major role in the self-regulation of star formation in galaxies through their contribution to the total energy and momentum budget of the interstellar medium (ISM; McKee & Ostriker 1977; Mac Low & Klessen 2004; Klessen & Glover 2016). As the rate of cooling in the ISM is proportional to the gas density squared, the efficiency with which SNe inject energy and momentum into the local galactic environment strongly depends on the density distribution of the gas into which they explode (see Girichidis et al. 2016 and references therein). For example, SNe that explode within dense molecular clouds may be limited to disrupting their natal gas clouds, whilst SNe that explode into lower density environments can drive hot expanding bubbles to much larger distances (tens to hundreds of pc) and influence galactic scale processes (e.g. kpc-scale galactic outflows; Veilleux, Cecil & Bland-Hawthorn 2005; Agertz et al. 2013; Stinson et al. 2013; Keller, Kruijssen & Wadsley 2020; Veilleux et al. 2020). Feedback from the pre-SNe stages of high-mass stars plays a significant role in determining the environment into which SNe subsequently explode. Simulations have long predicted that this \u2018pre-processing\u2019 can potentially even destroy the host molecular cloud before the first SN explosion (e.g. Dale, Ercolano & Bonnell 2012, 2013), and observations of molecular clouds and H\u2009ii regions in nearby galaxies now show that pre-SN feedback is primarily responsible for the destruction of molecular clouds across the local galaxy population (Kruijssen et al. 2019b; Chevance et al. 2020b,c). Studying the effects of these earliest stages of stellar feedback on their environment is then crucial to quantifying the contribution of SNe in driving the galaxy-scale energy and momentum cycle of the ISM in galaxies. In light of this, significant observational effort has been invested to better disentangle and quantify the effect of various feedback mechanisms within young stellar systems (e.g. Oey 1996a,b; Pellegrini, Baldwin & Ferland 2010, 2011). More recent efforts have focused on measuring and comparing the internal pressure components from different feedback mechanisms in H\u2009ii regions located within the Small and Large Magellanic Clouds (SMC and LMC, respectively), such as the well-known 30 Doradus complex (e.g. Lopez et al. 2011, 2014; Chevance et al. 2016; McLeod et al. 2019), as well as other nearby galaxies (e.g. McLeod et al. 2020). These studies have provided important insights into early-stage feedback, but further work is needed to understand how pre-processing varies with environment, particularly to higher density, pressure, and metallicity regimes such as those in galactic nuclei and high-redshift galaxies.","Citation Text":["Stinson et al. 2013"],"Functions Text":["For example, SNe that explode within dense molecular clouds may be limited to disrupting their natal gas clouds, whilst SNe that explode into lower density environments can drive hot expanding bubbles to much larger distances (tens to hundreds of pc) and influence galactic scale processes (e.g."],"Functions Label":["Background"],"Citation Start End":[[1287,1306]],"Functions Start End":[[903,1198]]}
{"Identifier":"2016AandA...592A..74S__Roming_et_al._2005_Instance_1","Paragraph":"We observed our full sample of all 24 candidate highly variable AGN with Swift (Gehrels et al. 2004) for ~2\u2009ks each, between 2010 and 2014 as part of a fill-in programme. All XRT (Burrows et al. 2005) observations were made in photon counting mode with exposure times ranging from 1.6\u20133.7\u2009ks. The Swift-XRT data were obtained from the UK Swift Science Data Centre1 and reduced following the procedures of Evans et al. (2009) using the Swift software and calibration database available within HEASOFT v.6.12. Simultaneous observations were made with the Swift BAT (Barthelmy et al. 2005) at 14\u2013195\u2009keV and the Swift UV\/Optical Telescope (UVOT; Roming et al. 2005) with the u filter applied. For ten sources, additional archival Swift observations were available at the time of writing which we have included and analysed in an identical manner. Details of all the observations used in this paper are given in the appendix (see Table A.1). With the XRT we detected 16 (or two-thirds) of the sample sources in our fill-in observations. Widening our search, we looked at data stacks in the Swift XRT Point Source Catalogue (1SXPS; Evans et al. 2014) and other pointed XRT observations and found that a further five sources were detected. Two of the three XRT non-detected sources have only ever been detected in the XMM slew survey (Figs. 1b and f), and cannot be identified with any source detected in other wavelength surveys such as 2MASS, WISE, SDSS or 6dF in our searches. One of these, XMMSL1\u2009J015510.9-140028, lies at the detection threshold of the slew survey. The other, XMMSL1\u2009J113001.8+020007, has a higher significance, however the photons at the source location are aligned along a row which indicates that this might not be an astrophysical point source. We conclude therefore that those two detections in the slew survey are highly likely to be spurious. We discuss the spurious fraction further in Sect. 8. The remaining XRT-undetected source is XMMSL1\u2009J193439.3+490922, which has three detections in XMM slews and is hence likely real. All XRT-detected sources are also detected with UVOT. Three sources (Mrk 352, ESO 362-G018, ESO 139-G012) can be found in the Swift BAT 70-month All-Sky Hard X-ray Survey Source Catalog (Baumgartner et al. 2013). ","Citation Text":["Roming et al. 2005"],"Functions Text":["Simultaneous observations were made with","the Swift UV\/Optical Telescope (UVOT;","with the u filter applied."],"Functions Label":["Uses","Uses","Uses"],"Citation Start End":[[643,661]],"Functions Start End":[[508,548],[605,642],[663,689]]}
{"Identifier":"2019AandA...625A..12G__Wunderlich_et_al._2019_Instance_1","Paragraph":"To calculate the boundaries of the HZ of a stagnant-lid Earth around M, K, G, and F dwarfs, we applied a 1D, cloud-free, radiative-convective climate model, which has been described in detail by von Paris et al. (2010) and von Paris et al. (2015), and is based on Kasting et al. (1984) and Segura et al. (2003). The radiative transfer is split into a stellar and a thermal wavelength regime. The short wavelength regime treats the absorption and scattering of stellar irradiation using a \u03b4-two-stream method including Rayleigh scattering coefficients following the approach of Allen (1973) and four-term correlated-k exponential sums covering a wavelength regime from 273.5 nm to 4.545 \u03bcm. This wavelength coverage is optimized for solar irradiation. Especially for late M dwarfs the cut-off at 4.545 \u03bcm leads to non-negligible loss in incoming radiation of up to \u22485% (see also Wunderlich et al. 2019). Hence, HZ boundaries obtained with the models lie closer to the star than would be expected when accounting for this missing portion of irradiation. The long-wavelength regime treats the absorption by CO2 and H2O in the wavelength regime from 1 to 500 \u03bcm via correlated-ks computed based on HITEMP 1995 (Rothman et al. 1995). The ckd continuum (Clough et al. 1989), and the collision-induced absorption as described in Kasting et al. (1984) for CO2 and as described in von Paris et al. (2013) for N2 \u2013N2 are included. Convection is treated by applying a convective adjustment when the adiabatic lapse exceeds the radiative lapse rate, including latent heat release from H2O or CO2 where applicable. The water mixing ratio profile (\n\n$C_{\\mathrm{H}_2\\mathrm{O}}$\n\n\n\nC\n\n\nH\n2\n\nO\n\n\n\n) is calculated from the temperature profile, the saturation vapour pressure (psat), and by assuming a relative humidity (RH): \n\n$C_{\\mathrm{H}_2\\mathrm{O}}=\\textrm{RH}\\frac{p_{\\mathrm{sat}}}{p}$\n\n\n\nC\n\n\nH\n2\n\nO\n\n=RH\n\n\np\n\nsat\n\n\np\n\n\n\n, with p the pressure of the atmosphere. By making use of our 1D climate model we estimate global diurnal mean values without accounting for effects such as slow planetary rotation or an interactive hydrological cycle. A discussion on the potential influence of 3D processes is given in Sect. 4.","Citation Text":["Wunderlich et al. 2019"],"Functions Text":["Especially for late M dwarfs the cut-off at 4.545 \u03bcm leads to non-negligible loss in incoming radiation of up to \u22485% (see also","Hence, HZ boundaries obtained with the models lie closer to the star than would be expected when accounting for this missing portion of irradiation."],"Functions Label":["Uses","Uses"],"Citation Start End":[[878,900]],"Functions Start End":[[751,877],[903,1051]]}
{"Identifier":"2016ApJ...833..216G___2010_Instance_1","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. 2010"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[878,900]],"Functions Start End":[[717,877]]}
{"Identifier":"2017MNRAS.470..755H__Toomre_&_Toomre_1972_Instance_1","Paragraph":"Supermassive black holes (SMBHs) are believed to exist in the centres of all massive galaxies (Kormendy & Richstone 1995). A small proportion of these are growing, with gas accretion rates ranging from \u223c10\u22124 to 10 M\u2299 yr\u22121 and a proportionately wide range of bolometric luminosities (\u223c1042\u20131047 erg s\u22121). These are active galactic nuclei (AGNs) and may accrete large fractions of their mass in bursts of rapid accretion (Croton et al. 2006), requiring rapid inflow of gas from galaxy length-scales. Stripping the gas of enough angular momentum to allow for such rapid accretion, thereby powering the most luminous AGN, proves extremely challenging. Theoretical work suggests major mergers can provide the torque to displace such an overwhelming fraction of the angular momentum of the gas, allowing for the highest accretion rates on to the central black hole whilst transforming the galaxy morphology (Toomre & Toomre 1972; Barnes 1988; BarnesBarnes & Hernquist 1991; Di Matteo, Springel & Hernquist 2005; Cox et al. 2008). Gas rich mergers may trigger nuclear and global starbursts (Mihos & Hernquist 1994, 1996; Hopkins et al. 2006) and major mergers disrupt the morphologies of the colliding galaxies, often exhibiting long tidal tails or shells of expelled gas and stars soon after the merger has begun. Detecting this can be challenging however, since the single new galaxy has a relaxation time-scale after which morphological features of mergers fade (Tinsley 1978; Kennicutt et al. 1987; Ellison et al. 2013). Observational evidence suggesting a link between major mergers and SMBH accretion has been mixed (e.g. Gabor et al. 2009; Cisternas et al. 2011; Schawinski et al. 2011; Kocevski et al. 2012; Treister et al. 2012; Ellison et al. 2013; Villforth et al. 2014; Kocevski et al. 2015; Villforth et al. 2017). Alternatively, AGNs may be triggered secularly through, for example, disc instabilities (Bournaud et al. 2011), bars (Knapen, Shlosman & Peletier 2000; Oh, Oh & Yi 2012) or otherwise by minor mergers (Kaviraj 2013). It remains unclear whether alternatives to major merger triggering can drive several M\u2299 yr\u22121 of gas to the central SMBHs, as is necessary to power the most luminous AGN.","Citation Text":["Toomre & Toomre 1972"],"Functions Text":["Theoretical work suggests major mergers can provide the torque to displace such an overwhelming fraction of the angular momentum of the gas, allowing for the highest accretion rates on to the central black hole whilst transforming the galaxy morphology"],"Functions Label":["Background"],"Citation Start End":[[902,922]],"Functions Start End":[[648,900]]}
{"Identifier":"2021ApJ...909..173B__Bandiera_&_Petruk_2016_Instance_1","Paragraph":"Taking advantage of modern computer architecture, numerical simulations have become a powerful tool for investigating SNR dynamical evolution by means of testing various scenarios with specific purposes (Ferrand 2020). As relevant observational evidence is steadily growing, particular attention has been devoted to tackling with SNR temporal evolution in an inhomogeneous medium, aiming at deciphering the observed asymmetries of emission morphology or clarifying the effects on the synthetic radio polarization maps exerted by various configurations of the interstellar magnetic field (e.g., Orlando et al. 2007; Schneiter et al. 2015; Yang et al. 2015; Bandiera & Petruk 2016; Petruk et al. 2017). As a first attempt, magnetohydrodynamic (MHD) simulations of a benchmark type Ia SNR expanding into a turbulent background were carried out by Balsara et al. (2001), and several obtained results are different from those in a uniform medium: significant azimuthal variations observed in density and magnetic field profiles, a patchy and time-independent synchrotron shell which may in turn serve as a probe into the nature of ISM turbulence, and amplification of the magnetic field in the post-shock region due to interactions between the forward shock and the turbulent background. Under the assumption that both density and magnetic field fluctuations follow a Kolmogorov-like power spectrum, the temporal evolution of an SNR propagating into a turbulent medium was investigated by Guo et al. (2012), with a focus on the structures and amplification of the magnetic field in the shock downstream. Recently, two-dimensional cylindrical MHD simulations were implemented to investigate the dynamical properties of young type Ia SNRs undergoing shock acceleration in a turbulent medium by Peng et al. (2020), where an initial power-law density profile is adopted and the derived relative contact discontinuity positions are compared with the observed results of two typical type Ia SNRs: SN 1006 and Tycho.","Citation Text":["Bandiera & Petruk 2016"],"Functions Text":["As relevant observational evidence is steadily growing, particular attention has been devoted to tackling with SNR temporal evolution in an inhomogeneous medium, aiming at deciphering the observed asymmetries of emission morphology or clarifying the effects on the synthetic radio polarization maps exerted by various configurations of the interstellar magnetic field (e.g.,"],"Functions Label":["Motivation"],"Citation Start End":[[656,678]],"Functions Start End":[[219,593]]}
{"Identifier":"2020AandA...639A.116Y__Wenger_et_al._2000_Instance_1","Paragraph":"The main RSG sample contains 1405 candidates from the SMC source catalog. However, targets with Rank 4 and 5 are selected only in one CMD by either the MIST models or the theoretical J\u2005\u2212\u2005KS color cuts, and many of them reach down close to the tip of the red giant branch (TRGB) and AGB population (see also Figs. 13 and 14 of Yang et al. 2019). To be on the safe side we adopted only targets with ranks from 0 to 3 (targets identified in at least two CMDs) as our initial sample, which resulted in 1107 targets. Due to the photometric quality cuts and uncertainties of the Spitzer and Gaia data, the strict constraints on the astrometric solution, and the deblending applied during the construction of the source catalog, some of the spectroscopically confirmed RSGs were also rejected. In order to make the sample as complete as possible, we retrieved and added all known spectroscopic RSGs in both optical and mid-infrared (MIR) bands from Simbad (Wenger et al. 2000) and data taken by Spitzer Infrared Spectrograph (IRS; Houck et al. 2004), respectively. From Simbad, we selected 322 RSGs with RV \u2265 90 km s\u22121, spectral type later than G0, and luminosity class brighter than II by using criteria query (Levesque 2013; Gonz\u00e1lez-Fern\u00e1ndez et al. 2015), for which 192 targets were matched with our initial sample within 1\u2033. Additionally, a crossmatching with the main RSG sample of 1405 candidates indicated that three Rank 4 candidates were also matched within 1\u2033. Surprisingly, there are two spectroscopic RSGs matched with the source catalog within 1\u2033, but not selected as the RSG candidates by either the MIST models or the theoretical J\u2005\u2212\u2005KS color cuts. Visual inspection of Gaia and 2MASS (Two Micron All Sky Survey; Skrutskie et al. 2006) CMDs (shown below) indicated that these two targets were slightly off the blue and red boundaries of the RSG region, respectively, which was likely due to the intrinsic variability of the RSGs (Kiss et al. 2006; Yang & Jiang 2011; Ren et al. 2019). Consequently, in total, there are 127 unselected spectroscopic RSGs from Simbad. For data taken by Spitzer\/IRS, there were 22 RSGs from Ruffle et al. (2015), who classified 209 point sources observed by Spitzer\/IRS using a decision tree method, based on IR spectral features, continuum and spectral energy distribution shape, bolometric luminosity, cluster membership, and variability information (all the targets from Kraemer et al. 2017 were also included). Of these 22 RSGs, 16 of them were matched with our initial sample within 1\u2033, and 4 of them were matched with the previous unselected Simbad RSGs within 1\u2033. Thus, there are only two unselected spectroscopic RSGs from Spitzer\/IRS. In total, there are additional 129 spectroscopic RSGs from both Simbad and Spitzer\/IRS, for which we give them Rank \u22121.","Citation Text":["Wenger et al. 2000"],"Functions Text":["In order to make the sample as complete as possible, we retrieved and added all known spectroscopic RSGs in both optical and mid-infrared (MIR) bands from Simbad"],"Functions Label":["Uses"],"Citation Start End":[[950,968]],"Functions Start End":[[787,948]]}
{"Identifier":"2018AandA...619A..13V__Saviane_et_al._2012_Instance_5","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"],"Functions Text":["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"],"Functions Label":["Uses"],"Citation Start End":[[2364,2383]],"Functions Start End":[[2182,2363]]}
{"Identifier":"2015ApJ...809..117Y__Hayasaki_et_al._2008_Instance_1","Paragraph":"Considering a BBH system resulting from a gas rich merger, the BBH is probably surrounded by a circumbinary disk, and each of the two SMBHs is associated with a mini-disk (see Figure 1). In between the circumbinary disk and the inner mini-disks, a gap (or hole) is opened by the secondary SMBH, which is probably the most distinct feature of a BBH\u2013disk accretion system, in analogy to a system in which a gap or hole is opened by a planet migrating in the planetary disk around a star (Lin et al. 1996; Quanz et al. 2013). This type of geometric configurations for the BBH\u2013disk accretion systems has been revealed by many numerical simulations and analysis (Artymowicz & Lubow 1996; Escala et al. 2005; Hayasaki et al. 2008; Cuadra et al. 2009; D\u2019Orazio et al. 2013; Farris et al. 2014; Roedig et al. 2014).4\n\n4\nThe width of the gap (or hole) is roughly determined by, but could be somewhat larger than, the Hill radius \n\n\n\n\n\n. However, set a slightly large gap size, e.g., \n\n\n\n\n\n, does not affect the results presented in this paper significantly.\n The continuum emission from disk accretion onto a BBH may be much more complicated than that from disk accretion onto a single SMBH, since the dynamical interaction between the BBH and the accretion flow onto it changes the disk structure (G\u00fcltekin & Miller 2012; Sesana et al. 2012; Rafikov 2013; Roedig et al. 2014; Yan et al. 2014; Farris et al. 2015). Nevertheless, we adopt a simple model to approximate the continuum emission from a BBH\u2013disk accretion system as the combination of the emissions from an outer circumbinary disk and an inner mini-disk around the secondary SMBH, each approximated by multicolor blackbody radiation in the standard thin disk model (Novikov & Thorne 1973; Shakura & Sunyaev 1973). The emission from the mini-disk around the primary SMBH is insignificant for a BBH system with a small mass ratio (roughly in the range of a few percent to 0.25) due to its low accretion rate as suggested by the state of the art numerical simulations (Roedig et al. 2012; Farris et al. 2014), thus its emission can be neglected. Our analysis suggests that a large q cannot lead to a good fit to the observations.","Citation Text":["Hayasaki et al. 2008"],"Functions Text":["This type of geometric configurations for the BBH\u2013disk accretion systems has been revealed by many numerical simulations and analysis"],"Functions Label":["Background"],"Citation Start End":[[703,723]],"Functions Start End":[[523,656]]}
{"Identifier":"2021AandA...655A..12T__Tang_et_al._2017b_Instance_4","Paragraph":"Using the RADEX3 non local thermodynamic equilibrium (LTE) modeling program (van der Tak et al. 2007) with collisional rate coefficients from Wiesenfeld & Faure (2013), we modeled the relation between the gas kinetic temperature and the measured average of para-H2CO 0.5 \u00d7 [(322\u2013221 + 321\u2013220)\/303\u2013202] ratios, adopting a 2.73 K background temperature, an average observational linewidth of 4.0 km s\u22121, and column densities N(para-H2CO) = 2.7 \u00d7 1012 and 3.7 \u00d7 1012 cm\u22122 for N113 and N159W, respectively. The results are shown in Fig. 5. The values of the para-H2CO column density were obtained with APEX data (beam size ~30\u2033; Tang et al. 2017b), which cover similar regions. Different column densities of para-H2CO only weakly affect derived kinetic temperatures (see Fig. 3 in Tang et al. 2017b or Fig. 4 in Tang et al. 2018a; this was also shown in Fig. 13 and discussed in Sect. 4.3.1 of Mangum & Wootten 1993) as long as all lines are optically thin. Considering that the relation between the gas temperature and the para-H2CO line ratio may vary at different spatial densities (see Fig. 2 in Tang et al. 2017b), we modeled it at spatial densities 104, 105, and 106 cm\u22123 in Fig. 5. It appears that Tkin at n(H2) = 105 cm\u22123 is consistently lower than values at 104 and 106 cm\u22123 by \u227223% and \u227234%, respectively, for Tkin \u2272 100 K. Local thermodynamic equilibrium (LTE) is a good approximation for the H2CO level populations under optically thin and high-density conditions (Mangum & Wootten 1993; Tang et al. 2017a,b, 2018b). Following the method applied by Tang et al. (2017b) in their Eq. (2), we plot the relation between the LTE kinetic temperature, TLTE, and the para-H2CO (3\u20132) line ratio in Fig. 5. Apparently, TLTE agrees well with Tnon-LTE at volume densities n(H2) ~ 105 cm\u22123 as long as Tkin \u2272 100 K. Previous observations show that para-H2CO (3\u20132) is sensitive to gas temperature at density 105 cm\u22123 (Ginsburg et al. 2016; Immer et al. 2016; Tang et al. 2017b). The spatial density measured with para-H2CO (303\u2013202) and C18O (2\u20131) in N113 and N159W is n(H2) ~ 105 cm\u22123 on a size of ~30\u2033 (Tang et al. 2017b). Therefore, here we adopt 105 cm\u22123 as an averaged spatial gas density in the N113 and N159W regions.","Citation Text":["Tang et al. 2017","b"],"Functions Text":["Local thermodynamic equilibrium (LTE) is a good approximation for the H2CO level populations under optically thin and high-density conditions"],"Functions Label":["Uses"],"Citation Start End":[[1497,1513],[1515,1516]],"Functions Start End":[[1331,1472]]}
{"Identifier":"2018AandA...610A..38F__Bisterzo_et_al._2017_Instance_3","Paragraph":"Similarly to the [\u03b1\/Fe] ratio, the ratio of the slow (s-) neutron capture process elements to iron can be regarded as a cosmic clock. Ba, Sr, La, and Y are mainly s-process elements produced on long timescales by low mass AGB stars (Matteucci 2012). Since a low mass star must evolve to the AGB phase before the s-process can occur, the s-process elements are characterized by a delay in the production, much like the delay of iron production by SNe Ia relative to the \u03b1 elements production by core collapse SNe. Among the four s-process elements mentioned above, GES provides the abundances of Y II (the first s-process peak element) and Ba II (the second s-process peak element) for all our sample stars. Their abundances behave differently in the Galactic thick and thin discs (Bensby et al. 2005, 2014; Israelian et al. 2014; Bisterzo et al. 2017; Delgado Mena et al. 2017). Unlike the Galactic thick disc stars, which show an almost constant [Ba\/Fe] abundance close to the solar value, the Galactic thin disc stars have their [Ba\/Fe] abundances increasing with [Fe\/H] and reaching their maximum values around solar metallicity, after which a clear decline is seen (see also Cristallo et al. 2015a,b, for the most recent s-process calculation in AGB yields). The same trend is observed in our sample. In Fig. 13 we display the Li-[Ba\/Fe], [Ba\/Fe] as a function of [Fe\/H], and the evolution of absolute Ba abundance A(Ba), as derived from Ba II lines. Similar figures are also plotted for yttrium (Y II). [Ba\/Fe] and [Y\/Fe] values here are derived from MCMC simulations, taking into account the measurement uncertainties of A(Ba II)\/A(Y II) and [Fe\/H]. By applying the same MCMC setups used for [\u03b1\/Fe] (see Sect. 3.1), we calculate the mean values of [Ba\/Fe] and [Y\/Fe] for each star. These values, together with their corresponding 1\u03c3 uncertainties, are listed in Table 1. In the literature there are several theoretical works on the evolution of [Ba\/Fe] and [Y\/Fe] in the Galactic thin disc (e.g. Pagel & Tautvaisiene 1997; Travaglio et al. 1999, 2004; Cescutti et al. 2006; Maiorca et al. 2012; Bisterzo et al. 2017). For comparison, we show in Fig. 13 the predictions of the most recent one (Bisterzo et al. 2017) where the updated nuclear reaction network was used.","Citation Text":["Bisterzo et al. 2017"],"Functions Text":["For comparison, we show in Fig. 13 the predictions of the most recent one","where the updated nuclear reaction network was used."],"Functions Label":["Compare\/Contrast","Background"],"Citation Start End":[[2199,2219]],"Functions Start End":[[2124,2197],[2221,2273]]}
{"Identifier":"2019MNRAS.485L..78V__Chatterjee_et_al._2017_Instance_1","Paragraph":"The properties of the persistent radio source associated with FRB 121102 may be constrained independently of the Faraday-rotating medium. We assume equipartition between the relativistic gas and magnetic field as is common in synchrotron sources3 (Readhead 1994). The source becomes self-absorbed at $1.5$ GHz for radius $R_{\\rm per} < 0.05$ pc; this is thus the lower bound on the source size. European Very Long Baseline Interferometry (VLBI) Network observations of the source at 5 GHz set an upper bound on the source radius of Rper \u2272 0.35 pc (Marcote et al. 2017). This is consistent with the ${\\approx } 30\\, {{\\rm per\\, cent}}$ amplitude modulations observed in the source at 3 GHz (Chatterjee et al. 2017) being caused by refractive interstellar scintillation in the Milky Way interstellar medium (ISM; Walker 1998). For any radius within the allowed range (0.05  Rper\/pc  0.35), we can determine the equipartition magnetic field, Beq, and the column of relativistic electrons, Nrel, using the standard expressions for synchrotron emissivity and absorption coefficients (Rybicki & Lightman 1979, their equations 6.36 and 6.53). We assume a power-law energy distribution of radiating electrons with somewhat shallow index of b = \u22121.5 that can account for the relatively flat spectrum of the source (Chatterjee et al. 2017). The peak Lorentz factor of the distribution, \u03b3max, is chosen to correspond to the observed spectral break frequency of $\\nu _{\\rm max}=10$ GHz. If the lower Lorentz factor cut-off corresponds to emission at $\\nu _{\\rm min}=1$ GHz,4 then the equipartition magnetic field and electron column thus determined for minimum and maximum source sizes are $B_{\\rm eq}\\approx 140$ mG, $\\gamma _{ \\rm min}\\approx 50$, \u03b3max \u2248 160, $N_{\\rm rel} \\approx 0.95\\, {\\rm pc}\\, {\\rm cm}^{-3}$ for $R_{\\rm per}=0.05$ pc, and $B_{\\rm eq}\\approx 27$ mG, $\\gamma _{ \\rm min}\\approx 120$, \u03b3max \u2248 370, $N_{\\rm rel} \\approx 0.1\\, {\\rm pc}\\, {\\rm cm}^{-3}$ for $R_{\\rm per}=0.35$ pc. The reader can scale the equipartition field to other source sizes using Beq(R) \u221d R\u22126\/7. The total energy contained in the relativistic electrons and the magnetic field (\u2018equipartition energy\u2019) is \u223c1049.1 and \u223c1050.2 erg, respectively. If the relativistic electrons were injected in a one-off event, the synchrotron cooling rates at \u03b3max yield source ages of $14$ yr for R = 0.05 pc and 60 yr for $R=0.35$ pc. The corresponding expansion velocities are $0.011\\, c$ and $0.02\\, c$, respectively.","Citation Text":["Chatterjee et al. 2017"],"Functions Text":["This is consistent with the ${\\approx } 30\\, {{\\rm per\\, cent}}$ amplitude modulations observed in the source at 3 GHz","being caused by refractive interstellar scintillation in the Milky Way interstellar medium"],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[690,712]],"Functions Start End":[[570,688],[714,804]]}
{"Identifier":"2020AandA...643A...5D__Saintonge_et_al._2011_Instance_1","Paragraph":"There is significant scatter (larger than 1 dex) among the tdepl measurements in all redshift bins, even though we only consider MS galaxies with \u0394MS\u2004=\u2004\u00b10.3 dex around the MS parametrization of Speagle et al. (2014). This scatter at a fixed redshift is believed to be a product of the multi-functional dependence of tdepl on many physical parameters, such as the offset from the MS, the star formation rate, the stellar mass, and possibly the environment (e.g., Dessauges-Zavadsky et al. 2015; Scoville et al. 2017; Noble et al. 2017; Silverman et al. 2018; Tacconi et al. 2018; Tadaki et al. 2019; Liu et al. 2019b). Given the strong anti-correlation found between tdepl and the offset from the MS (Genzel et al. 2015; Dessauges-Zavadsky et al. 2015; Tacconi et al. 2018), we still expect tdepl variations for galaxies on the MS while in their evolutionary process they are transiting up and down across the MS band (e.g., Sargent et al. 2014; Tacchella et al. 2016). The previously reported anti-correlation between tdepl and sSFR (Saintonge et al. 2011; Dessauges-Zavadsky et al. 2015) is also further supported by our galaxies at z\u2004=\u20044.4\u2005\u2212\u20055.9 (Fig. 6, middle panel). This highlights comparable timescales for gas consumption and stellar mass formation. We find a Spearman rank coefficient of \u22120.49 and p-value of 4.5\u2005\u00d7\u200510\u221210 for the dependence of tdepl on sSFR when considering the MS SFGs at z\u2004\u223c\u20041\u2005\u2212\u20055.9. The observed offset of ALPINE galaxies with respect to the tdepl\u2013sSFR relation of MS SFGs at z\u2004=\u20040 and to a smaller extent to the relations at z\u2004\u223c\u20041 and z\u2004\u223c\u20042 is compatible with the displacement of the z\u2004=\u20040 relation along the sSFR-axis by factors derived from the sSFR evolution with redshift of MS SFGs out to z\u2004\u223c\u20045 (Speagle et al. 2014). Nevertheless, a less steep sSFR redshift evolution toward z\u2004\u223c\u20045 than parametrized by Speagle et al. (2014) is suggested by the ALPINE sample, in line with the sSFR(z) results of Khusanova et al. (2020a). On the other hand, with tdepl measurements achieved down to Mstars\u2004\u223c\u2004108.4\u2006M\u2299 for the ALPINE galaxies, we confirm that for MS SFGs at z\u2004\u223c\u20041\u2005\u2212\u20055.9 the tdepl dependence on Mstars, if any, must be weak as shown in Fig. 6 (right panel). This further supports the idea that the linear KS relation established for local galaxies (Kennicutt 1998b) might hold up to z\u2004\u223c\u20045.9 for MS SFGs.","Citation Text":["Saintonge et al. 2011"],"Functions Text":["The previously reported anti-correlation between tdepl and sSFR","is also further supported by our galaxies at z\u2004=\u20044.4\u2005\u2212\u20055.9 (Fig. 6, middle panel)."],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[1034,1055]],"Functions Start End":[[969,1032],[1089,1171]]}
{"Identifier":"2016MNRAS.461..344P__Bernardi_2009_Instance_2","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"],"Functions Text":["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"],"Functions Label":["Differences"],"Citation Start End":[[1510,1523]],"Functions Start End":[[1354,1489]]}
{"Identifier":"2017AandA...599A...4K__hand,_Hotta_et_al._(2016)_Instance_1","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"],"Functions Text":["Our results appear to stand apart from similar studies in full spherical shells (e.g.,","in that the differential rotation is strongly quenched as a function of the magnetic Reynolds number."],"Functions Label":["Differences","Differences"],"Citation Start End":[[107,124]],"Functions Start End":[[0,86],[126,227]]}
{"Identifier":"2022MNRAS.516.5712T__Krumholz_et_al._2015_Instance_1","Paragraph":"As discussed in Section 1, one of the primary motivations for this work is to attempt to understand the variations in apparent IMF that have been observed in early type galaxies (ETGs). While there are multiple lines of evidence for this variation, the most straightforward to extract from out simulations is the mass to light ratio of the stellar populations we produce. In this section we therefore look at the mass to light ratio in our simulations for the purpose of comparing to that in observed galaxies. We calculate this by using the slug stellar population synthesis code (da Silva, Fumagalli & Krumholz 2012; Krumholz et al. 2015) to generate isochrones at stellar population ages from 5 Gyr to 10 Gyr, using the MIST stellar evolution tracks (Choi et al. 2016) and Starburst99-style stellar atmosphere models (Leitherer et al. 1999). Each isochrone provides a prediction of present-day mass, bolometric luminosity, and luminosity in a range of photometric filters as a function of initial mass for stars with initial mass $\\ge 0.1\\, {\\rm M}_{\\odot }$; we assume that the luminosities of stars with initial masses less than 0.1 M\u2299 are negligible, and that these stars also experience negligible mass loss. We further assume that all stars with initial mass $\\lt 8\\, {\\rm M}_{\\odot }$ (which are all the stars formed in our simulations) that reach the end of their lives leave behind 0.7 M\u2299 white dwarf remnants. We use the isochrone to calculate the luminosity and present-day mass of all the stars formed in each of our simulations at ages from 5 \u2212 10 Gyr, and from these we calculate the mass to light ratio of the stellar population as a function of age for each of our simulations. For comparison, we use the same isochrones to calculate the mass to light ratio of Chabrier (2005) and Salpeter (1955, truncated at a lower mass limit of $0.1\\, {\\rm M}_{\\odot }$) IMFs at the same ages. We use the SDSS r band for this calculation, but results are qualitatively similar in other filters.","Citation Text":["Krumholz et al. 2015"],"Functions Text":["We calculate this by using the slug stellar population synthesis code"],"Functions Label":["Uses"],"Citation Start End":[[619,639]],"Functions Start End":[[511,580]]}
{"Identifier":"2020MNRAS.498.4906B__Mehringer_et_al._1992_Instance_1","Paragraph":"The large radiation field directly produced from young stellar objects can exert a significant pressure on the surrounding material. This radiation pressure, Prad, at a given position within an H\u2009ii region, is related to the bolometric luminosity, Lbol, of the stellar population and the distance, r, from each star to that position within the region:\n(3)$$\\begin{eqnarray*}\r\nP_\\mathrm{rad} = \\sum {\\frac{L_\\mathrm{bol}}{4 \\pi r^{2} c}},\r\n\\end{eqnarray*}$$where the summation is over all stars within the region. The volume-averaged direct radiation pressure, Pdir, is then given as (Lopez et al. 2014),\n(4)$$\\begin{eqnarray*}\r\nP_\\mathrm{dir} = \\frac{3 L_\\mathrm{bol}}{4 \\pi R^{2} c},\r\n\\end{eqnarray*}$$where R is the radius of the H\u2009ii region (or effective radius, Reff, that we define later and use throughout the rest of the paper), and Lbol is the bolometric luminosity from the population of massive stars within the H\u2009ii region. This form differs by a factor of three from McLeod et al. (2019, equation 4), as these authors calculate the radiation surface pressure rather than the volume average pressure. This expression is appropriate to compute the force balance at the surface of an empty shell. However, as this work aims at understanding the large-scale dynamics of the Galactic Centre H\u2009ii regions (e.g. the total energy and pressure budget for each source), the inclusion of a factor of three in the numerator of the above equation is required. We also note here that the higher metallicity within the Galactic Centre, or increasing the amount of dust, has no effect on the Pdir calculation. Direct radiation pressure is limited by the momentum supplied by the stellar radiation field, and, as long as there is enough dust around to absorb all the radiation, the momentum per unit time delivered is the same. Lopez et al. (2011) determine the bolometric luminosity of the H\u2009ii regions within the LMC and SMC from H\u2009$\\alpha$ emission. However, this is not possible for the H\u2009ii regions investigated here, due to the high optical extinction towards the Galactic Centre (Av >20\u2009mag), which completely obscures any H\u2009$\\alpha$ emission. We, therefore, adopt two alternative methods of calculating the bolometric luminosity using radio and infrared observations (i.e. wavelengths where the emission is much less affected by dust extinction). First, we can make the assumption that the bolometric luminosity is proportional to the flux of ionizing photons, $\\mathcal {N}_\\mathrm{LyC}$, such that $L_\\mathrm{bol} = \\mathcal {N}_\\mathrm{LyC} \\left\\langle h \\nu \\right\\rangle$, where \u3008h\u03bd\u3009 \u223c\u200915\u2009\u2009eV is the mean photon energy (Pellegrini et al. 2007). We use the $\\mathcal {N}_\\mathrm{LyC}$ for each H\u2009ii region as determined from the radio observations outlined in Table \u20092 (i.e. Gaume & Claussen 1990; Mehringer et al. 1992; Schmiedeke et al. 2016), and solve for the direct radiation pressure using equation (4). The second method assumes that the luminosity integrated over infrared wavelengths approximately corresponds to the total bolometric luminosity. This is a common assumption made for embedded star-forming regions, where the luminosity from massive stars produced at ultraviolet wavelengths is absorbed and remitted by the dust in the infrared. Barnes et al. (2017) have produced maps of the total infrared luminosity across the Galactic Centre. These authors fit a two-component modified blackbody function to extinction corrected 5.8\u201324\u2009$\\mu $m (Carey et al. 2009; Churchwell et al. 2009) and 160\u2009\u2013500\u2009$\\mu $m (Molinari et al. 2010) emission maps (referred to as the warm and cool component of the bolometric luminosity; see fig. 2 of Barnes et al. 2017). These infrared (i.e. bolometric) luminosity maps are used with the two methods outlined below to also determine the direct radiation pressure within each of the Galactic Centre H\u2009ii regions. In comparison with the first method for calculating the direct radiation pressure from radio observations, we choose to identify individual sources within the infrared maps. These can be considered as discrete H\u2009ii regions each with a single value of the direct radiation pressure. We choose to identify these H\u2009ii regions in the map of the warm component of the bolometric luminosity using a dendrogram analysis (Rosolowsky et al. 2008). We choose to use a structure finding algorithm, as opposed to by-eye identification, to give reproducibility within regions with particularly complex morphology (the warm bolometric luminosity map is given in fig. 2 from Barnes et al. 2017).3 We make use of the \u2018leaves\u2019 identified from the dendrogram analysis, which are the highest level (i.e. smallest) structures in the analysis and here represent distinct H\u2009ii regions. We take the mask of each H\u2009ii regions (dendrogram leaf), and apply this to both the warm and cool bolometric luminosity component maps, which we sum to then get the total bolometric luminosity. This is used with equation (4) to get the direct radiation pressure (Pdir) within each H\u2009ii region. The effective radius (Reff) of each H\u2009ii region is defined as the radius for a circle with the corresponding area (A) of each structure (i.e. $R_{\\rm eff}=\\sqrt{A\/\\pi }$). In addition to the dendrogram analysis, we also calculate the direct radiation pressure within apertures of increasing radius from the centre of each H\u2009ii region complex. To do so, we place circular masks for each source on to both the warm and cool bolometric luminosity component maps, and sum the enclosed values to then get the total bolometric luminosity. The circle is then increased in radius, and the process repeated. We again use equation (4) to determine the direct radiation pressure within these increasing circular apertures. This method differs from the dendrogram analysis, as it returns a continuous radial distribution from the source centre, as opposed to a distribution of distinct H\u2009ii region with various sizes.","Citation Text":["Mehringer et al. 1992"],"Functions Text":["We use the $\\mathcal {N}_\\mathrm{LyC}$ for each H\u2009ii region as determined from the radio observations outlined in Table \u20092 (i.e.","and solve for the direct radiation pressure using equation (4)."],"Functions Label":["Uses","Uses"],"Citation Start End":[[2806,2827]],"Functions Start End":[[2654,2782],[2854,2917]]}
{"Identifier":"2020AandA...637A..59A__Ziurys_et_al._(2018)_Instance_2","Paragraph":"Several molecules show a large discrepancy between the abundances derived from observations and calculated by chemical equilibrium, although it is not as severe as for the molecules discussed above. We refer to PN in O-rich stars and H2S in C-rich stars, which are indicated by hatched rectangles in Fig. 2. For PN in O-rich AGB atmospheres, the disagreement between the observed abundances, (1\u20132) \u00d7 10\u22128 (Ziurys et al. 2018), and the calculated maximum chemical equilibrium abundance is almost three orders of magnitude. However, uncertainties on the observational and theoretical sides mean that the true level of disagreement is unclear. For example, while Ziurys et al. (2018) derived a PN abundance of 10\u22128 relative to H2 in IK Tau, De Beck et al. (2013) and Velilla Prieto et al. (2017) derived higher abundances, (3\u20137) \u00d7 10\u22127, in this source. When we give preference to these latter abundances, the level of disagreement would be even higher. On the other hand, the formation enthalpy of PN is rather uncertain (see Lodders 1999), which directly translates into the calculated chemical equilibrium abundance. In this study we adopted the thermochemical data for PN from Lodders (1999), who gives preference to a formation enthalpy at 298.15 K of 171.5 kJ mol\u22121, while other compilations such as JANAF use lower values that would result in higher chemical equilibrium abundances for PN. This would reduce the level of disagreement. In the case of H2S in C-rich AGB stars, the calculated maximum chemical equilibrium abundance is 7 \u00d7 10\u221211, while the value derived from observations is ~50 times higher. In this case, the observed abundance is based on the detection of only one line in only one source (see Ag\u00fandez et al. 2012), and thus it has to be viewed with some caution. In summary, the main failures of chemical equilibrium to account for the observed abundances of parent molecules in circumstellar envelopes are NH3, HCN, CS, SO2, and possibly PN in M-type stars, H2O and NH3 in S-type stars, and the hydrides H2O, NH3, SiH4, PH3, and perhaps H2S as well in C-type stars. The large discrepancies between the abundances derived from observations and those calculated with chemical equilibrium necessarily imply that nonequilibrium chemical processes must be at work in AGB atmospheres. Any invoked nonequilibrium scenario must account for all these anomalously overabundant molecules, but must also reproduce the remaining molecular abundances that are reasonably well explained by chemical equilibrium. No scenario currently provides a fully satisfactory agreement with observations, although two mechanisms that can drive the chemical composition out of equilibrium have been proposed.","Citation Text":["Ziurys et al. (2018)"],"Functions Text":["However, uncertainties on the observational and theoretical sides mean that the true level of disagreement is unclear. For example, while","derived a PN abundance of 10\u22128 relative to H2 in IK Tau, De Beck et al. (2013) and Velilla Prieto et al. (2017) derived higher abundances, (3\u20137) \u00d7 10\u22127, in this source."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[660,680]],"Functions Start End":[[522,659],[681,849]]}
{"Identifier":"2021AandA...654A..89P__Poutanen_&_Svensson_1996_Instance_1","Paragraph":"We now investigate relativistic reflection models with a primary Comptonisation continuum shape, which is more physical and has a sharper high-energy rollover compared to an exponential cutoff power law. In addition, such models have the advantage of having the hot corona temperature (kThot) as a physical parameter rather than a phenomenological exponential cutoff energy. We first apply the RELXILLCP model, which uses the NTHCOMP Comptonisation model (Zdziarski et al. 1996; \u017bycki et al. 1999) as the incident spectrum. The other physical parameters are the same as those in the RELXILL model presented above. We find a good fit and infer similar hot corona temperatures of kThot \u223c 26 keV for all three epochs (Table 2). Then, we consider the REFLKERR where the hard X-ray Comptonisation spectrum is computed with the COMPPS model (Poutanen & Svensson 1996), which appears to be a better description of thermal Comptonisation when compared to Monte Carlo simulations (Zdziarski et al. 2020). Moreover, REFLKERR has, as physical parameter, either the Compton parameter (y) or the optical depth (\u03c4). We choose to perform the fit with the optical depth as the direct inferred parameter. The temperature of the thermal seed photons (kTbb) Comptonised by the hot corona is an explicit physical parameter of this model. Here, we assume that the seed photons are provided by the cold disc, then kTbb was fixed at 10 eV corresponding to the expected maximum temperature of the accretion disc around a black hole mass of 1.4 \u00d7 108 M\u2299 accreting at a \u223c10% Eddington rate. In addition, REFLKERR allows us to choose either a slab or a spherical geometry for the hot corona. The latter corresponds to numerous active sphere regions above the disc surface. The hard X-ray shape of the reflected component is calculated using IREFLECT convolved with COMPPS rather than using XILLVERCP (see Nied\u017awiecki et al. 2019 for detailed explanations). Both models give good fits (Table 2, Fig. 3), though a larger \u03c72 value for the spherical geometry. Similar values of the hot corona temperature are measured for the three epochs: kThot \u223c 30\u221231 keV and kThot \u223c 21\u221222 keV for the slab and spherical geometries, respectively. From \u03c4, we can infer the corresponding Compton-parameter of the hot corona (yhot) using the relation y = 4\u03c4(kT\/511 keV) (Beloborodov 1999). This correspond to yhot \u223c 0.5 and yhot \u223c 1.1 for the slab and spherical geometries, respectively.","Citation Text":["Poutanen & Svensson 1996"],"Functions Text":["Then, we consider the REFLKERR where the hard X-ray Comptonisation spectrum is computed with the COMPPS model","which appears to be a better description of thermal Comptonisation when compared to Monte Carlo simulations"],"Functions Label":["Uses","Compare\/Contrast"],"Citation Start End":[[836,860]],"Functions Start End":[[725,834],[863,970]]}
{"Identifier":"2022AandA...664A..45A__Bergh_2000_Instance_1","Paragraph":"We can compare Fig. 14 with the corresponding figure obtained by Mowlavi et al. (2019) in the Magellanic Clouds (see their Figs. 4 and 7 for the LMC and SMC, respectively) and in the Milky Way from DR2 data (their Fig. 10). Our conclusions here are similar: (i) There are far fewer C-rich than O-rich stars in the Galaxy than in the Clouds. The low number of carbon stars compared to O-rich stars between the Galaxy and the Clouds agrees with the decrease in TDU efficiency with increasing metallicity, and with a higher O abundance in the envelope of Galactic AGB stars on average. On the other hand, because some of the stars located in RGB and faint AGB region in Fig. 14 might be CH and\/or R-hot type stars (i.e. C-rich objects, see previous sections), we could derive a lower limit for the ratio between carbon to M (O-rich) stars. It is well known that this ratio increases with the decreasing (average) metallicity of the galaxy. The primary reason for this correlation is that less C needs to be dredged-up in a metal-poor star to enable atmospheric carbon atoms to exceed those of oxygen. We derive a ratio \u223c0.05, which is very similar to the average ratio derived in the disc of M 31 (see e.g. van den Bergh 2000; Hamren et al. 2015). (ii) Moreover, the distribution of O-rich stars of the Gaia-2MASS diagram covers at any given MKs magnitude a much wider range of WRP,\u2006BP\u2005\u2212\u2005RP\u2005\u2212\u2005WKs,\u2006J\u2005\u2212\u2005Ks in the Galaxy than in the Clouds. The wider distribution in both axes of the O-rich zones in Fig. 14 compared with the corresponding feature in the Clouds would result from the combined effect of O-rich AGB stars turning much later into C-rich stars in the Galaxy, and the existence of a more ample range of stellar metallicities in the Galactic sample. Moreover, Fig. 14 clearly shows the high dispersion existing in MKs for the C-rich objects at a given WRP,\u2006BP\u2005\u2212\u2005RP\u2005\u2212\u2005WKs,\u2006J\u2005\u2212\u2005Ks value. Part of this dispersion in MKs is compatible with the typical range in Teff (2500\u22123500 K) deduced for N-type stars (Bergeat et al. 2001), and to the circumstellar extinction (which we did not consider here) preferably for the objects in the extreme C-rich region. However, the mixing of carbon stars of different populations probably also contributes significantly to this dispersion in MKs (and also in Mbol). In fact, following the method outlined in Sect. 4, we have studied the kinematics of these stars and calculated the membership probability to the halo and to the thin and thick disc of the 2659 new carbon star candidates with a Vrad measurement according to EDR3 (\u223c40% of the sample, i.e. 1305 objects). Figure 15 shows the corresponding Toomre diagram for these carbon stars. Of this limited sample, roughly 50% belong to the thin disc (blue circles in Fig. 15), \u223c30% to the thick disc or halo (blue crosses and triangles, respectively), and the rest (\u223c20%) have an ambiguous membership according to our membership criteria (open blue circles). Nevertheless, we note that assuming a less strict likelihood percentage to assign membership to a population (see Sect. 4), about 25% of the stars with ambiguous membership would be thick-disc and\/or halo stars. Because thick-disc and halo stars are older than thin-disc stars on average, many of these C-rich objects are very probably not intrinsic AGB carbon stars, but extrinsic C-rich giants: stars with masses lower than \u223c1.5\u2006M\u2299 that have become carbon rich through the mass transfer in a binary system. An alternative to this would be the possibility that the minimum mass for the formation of an intrinsic AGB stars could be as low as 1 M\u2299. Some observational evidence for this can be found in the literature (Shetye et al. 2019). This conclusion is reinforced by the scale height onto the Galactic plane that can be estimated for all the C-rich stars in Fig. 14 similarly to what was done in Sect. 2: an exponential fit gives zo\u2004\u223c\u2004600 pc, which is much larger that the scale height derived for the intrinsic N-type AGB carbon stars that clearly belong to the thin disc (see Sect. 2).","Citation Text":["van den Bergh 2000"],"Functions Text":["We derive a ratio \u223c0.05, which is very similar to the average ratio derived in the disc of M 31 (see e.g."],"Functions Label":["Similarities"],"Citation Start End":[[1204,1222]],"Functions Start End":[[1098,1203]]}
{"Identifier":"2021MNRAS.507.2115M__Isanto_&_Polsterer_2018_Instance_1","Paragraph":"In astrophysics, the number of studies that apply ML techniques has risen substantially in the last years. Unsupervised learning algorithms have been used to identify different kinematic components of simulated galaxies (Obreja et al. 2018, 2019), to compare stellar spectra (Traven et al. 2017), to classify pulsars (Lee et al. 2012), and to find high-redshift quasars (Polsterer, Zinn & Gieseke 2013). Supervised learning has been used to classify variable stars (Richards et al. 2011), to classify galaxies morphologically (Huertas-Company et al. 2008), and to determine the redshift of galaxies (Hoyle et al. 2015; Hoyle 2016; D\u2019Isanto & Polsterer 2018). Recently, ML has also been used to connect the properties of galaxies and dark matter haloes using supervised learning techniques. Kamdar, Turk & Brunner (2016a), Kamdar, Turk & Brunner (2016b) use tree-based methods to predict several galaxy properties from a set of halo properties and train the models on galaxy catalogues obtained from semi-analytic models and the Illustris hydrodynamic simulation (Vogelsberger et al. 2014). Sullivan, Iliev & Dixon (2018) train a simple neural network with one hidden layer to predict the baryon fraction within a dark matter halo at high redshift, given several halo properties (features). As training data they use the results of a cosmological hydrodynamic simulation with Ramses-RT. Similarly, Agarwal, Dav\u00e9 & Bassett (2018) use several ML methods to link input halo properties to galaxy properties, training on the Mufasa cosmological hydrodynamical simulation. Taking a reverse approach, Calderon & Berlind (2019) train tree-based methods and a neural network to derive halo mass from galaxy properties, training on an SDSS group catalogue. The limitation of all these studies is the supervised training and the training data. As labelled galaxy-halo data is not available for observed systems, the data for supervised learning has to be taken from a model. Even if the ML algorithms learn to reproduce the training data perfectly, the connection between galaxy and halo properties is the same as in the simulations. If the simulations predict the true relations poorly, so will the ML method. Therefore ML algorithms should not be trained on simulated data, but on observed data directly.","Citation Text":["D\u2019Isanto & Polsterer 2018"],"Functions Text":["Supervised learning has been used","and to determine the redshift of galaxies"],"Functions Label":["Background","Background"],"Citation Start End":[[631,656]],"Functions Start End":[[404,437],[557,598]]}
{"Identifier":"2015MNRAS.450.4364N__Wu_et_al._2004_Instance_1","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"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[874,888]],"Functions Start End":[[651,785]]}
{"Identifier":"2019MNRAS.483.3022G__Mowlavi_et_al._2018_Instance_1","Paragraph":"LPVs are known to exist on sets of sequences in period\u2013luminosity space (e.g. Wood et al. 1999; Wood 2000) depending on their variability type and pulsation mode. The primary periods of Mira-like variables lie on the commonly called C and C\u2032 sequences (Ita et al. 2004; Spano et al. 2011) with C\u2032 lying at a lower period than C. Sequences A and B lie at lower periods still and are populated by the so-called OGLE small amplitude red giants or OSARGs (Soszynski et al. 2004a). Sequences D and E are populated by long secondary pulsators, whose nature is still unclear, as well as ellipsoidal and eclipsing binary systems (Soszynski et al. 2004b), with the latter two contained in a separate CRTS catalogue not utilized in this work. Soszy\u0144ski & Wood (2013) also found that there is a low-amplitude population of SRVs whose primary period lies between the C and C\u2032 sequences. Given that the selections of Fig. 1 produce a CRTS sample with a wide range of visual amplitudes, it is uncertain on which sequence our CRTS sources lie and to which class of LPVs they belong. The left-hand panel of Fig. 2 shows sources from Gaia\u2019s DR2 LPV catalogue (Gaia Collaboration et al. 2016, 2018; Mowlavi et al. 2018) chosen to lie within a 15\u00b0 aperture of the LMC. Cross-matching with 2MASS reveals three distinct sequences existing in the Ks band, with the middle sequence corresponding to the Mira C sequence. Belokurov et al. (2017) defined a variability parameter (the \u2018Gaia amplitude\u2019) based on Gaia\u2019s flux information as:\n(1)\r\n\\begin{eqnarray*}\r\n\\mbox{Amp} = \\mbox{log}_{10} \\left(\\sqrt{N_{\\mbox{obs}}} \\: \\frac{\\sigma _{\\overline{I_{\\mathrm{ G}}}}}{\\overline{I_{\\mathrm{ G}}}} \\right)\r\n\\end{eqnarray*}\r\nwhere $\\sqrt{N_{\\mbox{obs}}}$ is the number of observations, $\\sigma _{\\overline{I_{\\mathrm{ G}}}}$ is the mean flux error in the G band, and $\\overline{I_{\\mathrm{ G}}}$ is the mean flux in the G band. Requiring a Gaia amplitude greater than \u22120.55 isolates a sample of LPVs lying on a single period\u2013luminosity sequence, as evident in Fig. 2. We produce a secondary cleaned CRTS sample located on this middle sequence by cross-matching with the Gaia DR2 source catalogue and enforcing that the Gaia amplitude > \u22120.55. Further, we restrict the sample to visual amplitudes >1 to eliminate small amplitude variables. Application of these two cuts gives a cleaned CRTS sample with 225 members. We utilize this cleaned CRTS sample in Section 3.1, where we exploit the period\u2013age correlation of LPVs, the interpretation of which may be muddied if low visual amplitude variables were inadvertently retained. For example, AGB variables residing in a common Magellanic Cloud globular cluster have been observed by Kamath et al. (2010) both at high and low amplitudes. They found stars with low amplitudes also had lower stellar periods, implying differing period\u2013age relations.","Citation Text":["Mowlavi et al. 2018"],"Functions Text":["The left-hand panel of Fig. 2 shows sources from Gaia\u2019s DR2 LPV catalogue","chosen to lie within a 15\u00b0 aperture of the LMC."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1181,1200]],"Functions Start End":[[1068,1141],[1202,1249]]}
{"Identifier":"2019ApJ...883...88B__Kleint_et_al._2016_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":["Kleint et al. 2016"],"Functions Text":["Recent results from solar flares observed from space","demonstrate an NUV spectrum originating largely from Hydrogen Balmer continuum emission."],"Functions Label":["Background","Background"],"Citation Start End":[[1256,1274]],"Functions Start End":[[1179,1231],[1276,1364]]}
{"Identifier":"2020ApJ...892...53A__Dudas_et_al._2018_Instance_1","Paragraph":"These new limits, in conjunction with the inconsistency of isotropic flux interpretations, leave no room for an astrophysical interpretation of AAEs in the context of the standard model for time windows as short as 103 s. However, it has been shown that these events can be explained using physics beyond the standard model, as many models suggest that AAEs lend support for axionic dark matter, sterile neutrinos, supersymmetry, or heavy dark matter (Anchordoqui et al. 2018; Connolly et al. 2018; Dudas et al. 2018; Fox et al. 2018; Huang 2018; Abdullah et al. 2019; Anchordoqui & Antoniadis 2019; Borah et al. 2019; Chauhan & Mohanty 2019; Cherry & Shoemaker 2019; Chipman et al. 2019; Cline et al. 2019; Collins et al. 2019; Esmaili & Farzan 2019; Esteban et al. 2019; Heurtier et al. 2019a, 2019b; Hooper et al. 2019). Many of these models, excluding the axionic dark matter explanation (Esteban et al. 2019) or those heavy dark matter scenarios that are tuned to prevent signatures in IceCube (Hooper et al. 2019), can be constrained by this nonobservation at IceCube. Dedicated tests to quantify these constraints are beyond the scope of this work and may be the focus of a future study. In addition to explanations that incite new physics, it has recently been suggested that AAEs could be explained by downward-going CR-induced EASs that reflected off of subsurface features in the Antarctic ice (Shoemaker et al. 2019). Another possible explanation could be coherent transition radiation from the geomagnetically induced air shower current, which could mimic an upgoing air shower (Motloch et al. 2017; de Vries & Prohira 2019). Explaining these anomalous events with systematic effects or confirming the need for new physics requires a deeper understanding of ANITA\u2019s detection volume. Efforts such as the HiCal radio frequency pulser, which has flown alongside ANITA on the last two flights (Prohira et al. 2018), are already underway to try to characterize the various properties of the Antarctic ice surface.","Citation Text":["Dudas et al. 2018"],"Functions Text":["However, it has been shown that these events can be explained using physics beyond the standard model, as many models suggest that AAEs lend support for axionic dark matter, sterile neutrinos, supersymmetry, or heavy dark matter"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[499,516]],"Functions Start End":[[222,450]]}
{"Identifier":"2020ApJ...899..143C___2016_Instance_1","Paragraph":"Having constrained the estimated CME propagation by matching the observed and modeled fronts in this way, it was then possible to estimate the arrival time and velocity of each CME as it passed Venus and Earth. At Venus, the shock driven by CME-2 caught up with CME-1, showing a typical shock-ICME structure. According to observations from the Wind spacecraft, the shock driven by CME-2 had passed through CME-1 before arrival at Earth. By comparing the estimated and observed arrival times of each CME, our approach resulted in arrival time estimates within 2.5 hr of those observed at VEX and Wind. The longitudinal structure of the CMEs obtained by this approach has the potential to improve space weather forecasting. The accuracy of arrival times can be affected by uncertainties in the CME initial conditions, interaction between CMEs, distortion of the CME shape by solar wind structure, the presence of shocks, and the efficiency of solar wind drag on each CME. The relative importance of each such factor could be investigated through an ensemble approach. The drag-based ensemble model developed by Dumbovi\u0107 et al. (2018) is a possible option, but the method is not valid for CME\u2013CME interaction events. In order to test the sensitivity of CME propagation direction and width to the predicted arrival time and velocity, we summarize the CME initial parameters from Srivastava et al. (2018), Kilpua et al. (2019), and Scolini et al. (2019) and two conditions from ours (the time at which the fastest velocity was estimated and the last observation by the coronagraphs). Table 2 shows the initial parameters from those studies and the errors in predicted velocity and arrival time. Note that the minimum value of the drag parameters we used here is 1 \u00d7 10\u22129 km s\u22121, which was revealed as the lowest possible value (Rollett et al. 2014, 2016; Temmer & Nitta 2015; Kubicka et al. 2016). Based on that, there may exist a residual between the predicted profile and the observations. The residual values of nose and ghost fronts are shown in the last two columns of Table 2. For an Earth-directed CME, the residual between model and observed elongations of the CME nose will reflect the accuracy of the predicted arrival time and in situ velocity of the CME at Earth. The residual value between model and observed elongations of the CME flank provides information about how well the model reproduces the shape of the CME front, which is important for predicting the time and in situ speed at Earth for those CMEs whose flank encounters the Earth. The model with the smallest residual values between model and observed elongations for both CME nose and flank provides the most creditable estimate of the in situ CME parameters at Earth, suggesting that the use of this technique can add to our forecast skill. In our analysis, the CME initial parameters estimated using the time and location of the fastest CME velocity observed in the coronagraph data produced the smallest residual between subsequent observed and model CME elongations. This should therefore be considered to be the best estimate of CME kinematics, and, if we were doing a weighted ensemble, this would be given a much higher weighting than the other runs. It should be noted that many of the other runs did not produce a minimum residual within the likely range of drag parameters, indicating that these runs are somehow not capturing the physics. Table 2 shows that different initial parameters change the predicted arrival time of CME-1 at Venus by 8.5 hr and at Earth by 14.6 hr in total. For CME-2, the prediction errors are found as 6 hr at Venus and 8 hr at Earth. The prediction errors are comparable to the mean absolute error obtained from the drag-based ensemble model (Dumbovi\u0107 et al. 2018). In future work, we will try to use the heliospheric upwind extrapolation model in large ensembles to efficiently investigate the effect on the CME transit time of the uncertainty in the initial CME parameters and ambient solar wind (Owens et al. 2020a). Our results indicate that CME-1 and CME-2 interacted with each other before they arrived at Venus, according to the propagation distances derived from our modeling. The in situ observation from Venus shows a typical shock-ICME complex structure, which provides confirmation of this interpretation. Our results also suggest that a CME\u2013CME interaction is possibly involved in disrupting the propagation direction or geometry of each CME. Such a deflection and interaction is consistent with the fact that CME-1 (CME-2) reached VEX earlier (later) than expected (Shen et al. 2012, 2014; Wang et al. 2014, 2016). The difference between expected arrival time and observational arrival time is larger at Earth than at Venus. This implies that the interaction persisted between Venus and Earth.","Citation Text":["Rollett et al.","2016"],"Functions Text":["Note that the minimum value of the drag parameters we used here is 1 \u00d7 10\u22129 km s\u22121, which was revealed as the lowest possible value"],"Functions Label":["Uses"],"Citation Start End":[[1823,1837],[1844,1848]],"Functions Start End":[[1690,1821]]}
{"Identifier":"2017ApJ...837..130V__Pasquini_et_al._2008_Instance_1","Paragraph":"From the X-ray point of view, old open clusters are interesting for a number of reasons. First, X-ray observations efficiently detect different classes of close, interacting binaries, enabling the study of processes such as tidal coupling and the link between X-rays and rotation. The X-ray luminosity of late-type stars strongly depends on stellar rotation. As single stars age, they spin down due to magnetic braking (Pallavicini 1989). As a result, their X-ray emission decreases accordingly. An old star like our Sun (\u223c4.5 Gyr) has an X-ray luminosity of about 1026 to 1027 erg s\u22121 (0.1\u20132.4 keV; Peres et al. 2000). Even with the deepest exposures of a sensitive X-ray telescope like the Chandra X-Ray Observatory, this is nearly impossible to detect except for the nearest stars. Nevertheless, an early ROSAT observation of the old open cluster M 67, which lies at \u223c840 pc (Pasquini et al. 2008) and is about as old as the Sun (\n\n\n\n\n\n; Dinescu et al. 1995), revealed a large number of X-ray sources among the cluster members (Belloni et al. 1993). Many of these turned out to be close, tidally interacting binaries, where the stellar rotation is locked to the orbital period and therefore kept at a level that can sustain magnetically active coronae. Subsequent XMM-Newton (Gondoin 2005; Giardino et al. 2008; Gosnell et al. 2012) and Chandra (van den Berg et al. 2004, 2013; Giardino et al. 2008) observations of old open clusters have detected many such active binaries (ABs). ABs can be binaries of two detached stars, or they can have a contact or semi-detached configuration such as in W UMa and Algol binaries, respectively. In terms of the number of sources, ABs are the most prominent X-ray source class in old open clusters, but other classes of interacting binaries are represented as well. In cataclysmic variables (CVs), the X-rays are the result of accretion from a late-type main-sequence donor onto a white dwarf. In fact, the first ROSAT observation of M 67 was aimed at studying the X-rays from a CV that was discovered in the optical (Gilliland et al. 1991). The origin of the X-ray emission from more exotic open-cluster binaries, like blue stragglers, is less understood, but in X-rays they are more similar to the ABs than to the mass-transfer sources (van den Berg 2013).","Citation Text":["Pasquini et al. 2008"],"Functions Text":["Nevertheless, an early ROSAT observation of the old open cluster M 67, which lies at \u223c840 pc"],"Functions Label":["Background"],"Citation Start End":[[879,899]],"Functions Start End":[[785,877]]}
{"Identifier":"2017MNRAS.464..968S__Tacconi_et_al._2006_Instance_1","Paragraph":"Comparison of apparent effective diameters of these sources to direct size measurements supports a similar conclusion. Simpson et al. (2015) present ALMA observations of 23 SCUBA-2-selected SMGs with a median physical half-light diameter of 2.4 \u00b1 0.2 kpc, while Ikarashi et al. (2015) show ALMA observations of 13 AzTEC-selected SMGs with a median physical half-light diameter of $1.34^{+0.26}_{-0.28}$ kpc. ALMA observations of four SPT-selected lensed SMGs give a mean physical half-light diameter of 2.14 kpc (Hezaveh et al. 2013b). This measurement is consistent with a recent lensing analysis of a significantly expanded SPT-selected DSFG sample (Spilker et al. 2016). These high-resolution ALMA observations constrain the FIR sizes of the sources to be 1.0\u20132.5 kpc. Earlier observations of the physical sizes of SMGs by CO detection and 1.4 GHz imaging suggest larger sizes (e.g. Tacconi et al. 2006; Biggs & Ivison 2008; Younger et al. 2008). However, Simpson et al. (2015) point out that the submillimetre sizes are consistent with resolved 12CO detections, while the sizes derived from 1.4 GHz imaging are about two times larger because of the cosmic ray diffusion, which can explain the results before higher frequency observations at ALMA were possible (Chapman et al. 2004; Tacconi et al. 2006; Biggs & Ivison 2008; Younger et al. 2008). Similarly, Ikarashi et al. (2015) reveal that the 12CO detected sizes and the 1.4 GHz imaging sizes of similar sources are greater than their submillimetre sizes as well. Furthermore, observations of local galaxies also show the submillimetre sizes are smaller than the CO detected sizes (e.g. Sakamoto et al. 2006, 2008; Wilson et al. 2008) and the 1.4 GHz continuum sizes (e.g. Elbaz et al. 2011). Our photometrically derived $\\sqrt{\\mu }d$ is best compared to the submillimetre continuum sizes. With a median apparent effective diameter of $4.2^{+1.7}_{-1.0}$ kpc, the $\\sqrt{\\mu }d$ of our sample is one to six times the observed intrinsic diameters (1.0\u20132.5 kpc). Lensing (or multiplicity) increases the apparent effective size of a source, so this comparison favours a lensing (or multiplicity) interpretation for the ACT-selected sources.","Citation Text":["Tacconi et al. 2006"],"Functions Text":["Earlier observations of the physical sizes of SMGs by CO detection and 1.4 GHz imaging suggest larger sizes (e.g."],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[886,905]],"Functions Start End":[[772,885]]}
{"Identifier":"2019MNRAS.482..988C__Zhang_et_al._2016_Instance_1","Paragraph":"The EW(\u03bb8542)\/EW(\u03bb8498) values are obtained around \u223c1.4 in the first observing night, which are somewhat smaller than the values around \u223c1.6 derived in the second night, when a strong optical flare decay was detected (also see the Section 4.2). These low ratios are consistent with the values found for several other stars with strong chromospheric activity (e.g. Montes et al. 2000; Gu et al. 2002; L\u00f3pez-Santiago et al. 2003; Zhang & Gu 2008; G\u00e1lvez et al. 2009; Cao & Gu 2014, 2015, 2017; Zhang et al. 2016), which suggests that the $\\rm{Ca\\,{\\small II}}$ IRT line emission arises from plage-like regions. The EH\u03b1\/EH\u03b2 ratios have also been usually used as a diagnostic for discriminating the presence of different structures on the stellar surface. As Huenemoerder & Ramsey (1987) discussed, the low ratios in RS CVn-type stars are caused by plage-like regions, while prominence-like structures have high values. Similar results have also been reported by Hall & Ramsey (1992) who found that low ratios (\u223c 1\u20132) can be achieved both in plages and prominences viewed against the disc, but high values (\u223c 3\u201315) can only be obtained in extended prominence-like structures viewed off the stellar limb. We obtain the ratios on SZ Psc during the flare decay phase change from 2.40 to 3.40, which are not especially high and anticorrelated with the variation of activity emission shown in Fig. 3, in which the EWs of H\u03b1 and H\u03b2 line subtraction and the EH\u03b1\/EH\u03b2 values are plotted as a function of orbital phase. Especially when the EWs have an increasing oscillation during the gradually decrease, which is resulted from flare ejection (see the discussion in Section 4.2), the ratios have a similar anticorrelation feature. These facts suggest that the low EH\u03b1\/EH\u03b2 ratios and its variation might be dominantly associated with the flare decrease, and accompanied cool post-flare loops (see the Section 4.3) might play a part of role for the low ratios, which may have a state like absorbing prominence projected against the disc.","Citation Text":["Zhang et al. 2016"],"Functions Text":["These low ratios are consistent with the values found for several other stars with strong chromospheric activity"],"Functions Label":["Similarities"],"Citation Start End":[[492,509]],"Functions Start End":[[245,357]]}
{"Identifier":"2018MNRAS.473.4566P__Papaderos_et_al._2006_Instance_1","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"],"Functions Text":["For instance, shallow gradient in metallicity is seen in SBS 0335\u2212052 (","while no significant variations were seen in Mrk 35"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1173,1194]],"Functions Start End":[[1102,1173],[1196,1247]]}
{"Identifier":"2018ApJ...863..162M__Liu_et_al._2013_Instance_2","Paragraph":"NLFFF extrapolation provides the reconstructed coronal magnetic field for AR 11158 from 2011 February 13 \u2212 2011 February 15 (Figures 1(d)\u2013(f)). The field lines (yellow lines) within the core of the AR have arcade-like structure with a relatively strong twist mainly near the PIL. These figures show that the magnetic field evolved during this period. Although we did not quantitatively compare the field lines with the observation, in general, the reconstructed coronal field morphologies match with the observations in Figures 1(a)\u2013(c). The general morphologies and the locations of the high-twist fields are also in agreement with many previous studies (Jing et al. 2012; Sun et al. 2012; Dalmasse et al. 2013; Inoue et al. 2013, 2014a; Liu et al. 2013; Wang et al. 2013; Aschwanden et al. 2014; Malanushenko et al. 2014; Zhao et al. 2014). Unlike Zhao et al. (2014) who could identify the twisted flux rope from the topology of the reconstructed coronal field, we could not find an obvious topological signature of a flux rope existing in our NLFFF during our analysis time window. It might be due to the fact that there was little magnetic flux with twist higher than one turn in our NLFFF and it is difficult to topologically define it as a flux rope. However, our result is consistent with other NLFFF results (Jing et al. 2012; Sun et al. 2012; Liu et al. 2013; Wang et al. 2013; Inoue et al. 2014a; Malanushenko et al. 2014). The high-twist region in our result is also in agreement with the region with high helicity flux (Dalmasse et al. 2013) and the location of the flare ribbons (Bamba et al. 2013; Liu et al. 2013), as well as the high current density region (Janvier et al. 2014). Figures 1(g)\u2013(i) show the evolution of the twist distribution map, with the magnetic twist of the field lines plotted at the footpoints of field lines according to a color scale. This shows that the high-twist (strongly right-handed twist corresponding to Tw > 0.5) areas are concentrated in only a limited part of the AR. The high-twist area grew and became even more twisted just before the X2.2 flare (Figure 1(i)). Most parts of the AR have twist values less than 0.25, but near the PIL the twist can reach more than 0.5, even up to about a full turn. This is consistent with the results of Sun et al. (2012) and Inoue et al. (2014a). A high-twist (strong negative\/left-handed twist) area also developed in the eastern part of the AR, which did not exist initially on February 13. Both of these high-twist areas produced several flares. However, here we focus on the flares that resulted from the high-twist core region near the center of the AR, where the M6.6 and X2.2 flares occurred.","Citation Text":["Liu et al. 2013"],"Functions Text":["However, our result is consistent with other NLFFF results"],"Functions Label":["Similarities"],"Citation Start End":[[1352,1367]],"Functions Start End":[[1257,1315]]}
{"Identifier":"2021AandA...655A..25Z__Shimizu_et_al._2019_Instance_1","Paragraph":"Outflows are ubiquitous in both luminous AGN and in local Seyfert galaxies, and occur on a wide range of physical scales, from highly ionised semi-relativistic winds and jets in the nuclear region at subparsec scales to galactic scale outflows seen in mildly ionised, molecular, and neutral gas (Morganti et al. 2016; Fiore et al. 2017; Fluetsch et al. 2019; Lutz et al. 2020; Veilleux et al. 2020, and references therein). In some cases molecular and ionised winds have similar velocities and are nearly co-spatial, suggesting a cooling sequence scenario where molecular gas forms from the cooling of the gas in the ionised wind (Richings & Faucher-Giguere 2017; Menci et al. 2019). Other AGN show ionised winds that are faster than the molecular winds, suggesting a different origin of the two phases (Veilleux et al. 2020, and references therein). The molecular phase is a crucial element of the feeding and feedback cycle of AGN because it constitutes the bulk of the total gas mass and it is the site of star formation activity. On galactic scales massive molecular winds are common in local Seyfert galaxies (e.g. Feruglio et al. 2010; Cicone et al. 2014; Dasyra et al. 2014; Morganti et al. 2015; Garc\u00eda-Burillo et al. 2014, 2017, 2019); these winds likely suppress star formation (i.e. negative feedback) as they reduce the molecular gas reservoir by heating or expelling gas from the host-galaxy ISM. In late-type AGN-host galaxies, the gas kinematics appears complex at all scales, showing several components such as bars, rings, and (warped) discs, with high velocity dispersion regions (e.g. Shimizu et al. 2019; Feruglio et al. 2020; Fern\u00e1ndez-Ontiveros et al. 2020; Alonso-Herrero et al. 2020; Aalto et al. 2020; Audibert et al. 2020). Accurate dynamical modelling of the molecular gas kinematics reveals kinematically decoupled nuclear structures, high velocity dispersion at nuclei, trailing spirals, and evidence of inflows and AGN-driven outflows. (e.g. Combes et al. 2019; Combes 2019, 2021). The outflow driving mechanism (wind shock, radiation pressure, or jet), their multiphase nature, and their relative weights and impact on the galaxy ISM are still open problems (Faucher-Gigu\u00e8re & Quataert 2012; Zubovas & King 2012; Richings & Faucher-Giguere 2017; Menci et al. 2019; Ishibashi et al. 2021). To date, far different outflow phases have been observed only for a handful of sources. Atomic, cold, and warm molecular outflows have been observed in radio galaxies (e.g. Morganti et al. 2007; Dasyra & Combes 2012; Dasyra et al. 2014; Tadhunter et al. 2014; Oosterloo et al. 2017). The nuclear semi-relativistic phase and the galaxy-scale molecular phase have been observed simultaneously in less than a dozen sources, with varied results: in some cases data suggest energy driven flows (Feruglio et al. 2015; Tombesi et al. 2015; Longinotti et al. 2018; Smith et al. 2019), in other cases data suggest momentum driven flows (e.g. Garc\u00eda-Burillo et al. 2014; Feruglio et al. 2017; Fluetsch et al. 2019; Bischetti et al. 2019; Marasco et al. 2020). Fiore et al. (2017), using a compilation of local and high redshift winds, showed that there is a broad distribution of the momentum boost, suggesting that both energy- and momentum-conserving expansion may occur. Enlarging the sample of local AGN-host galaxies with outflows detected in different gas phases is important to understand the nature and driving mechanisms of galaxy-scale outflows.","Citation Text":["Shimizu et al. 2019"],"Functions Text":["In late-type AGN-host galaxies, the gas kinematics appears complex at all scales, showing several components such as bars, rings, and (warped) discs, with high velocity dispersion regions (e.g."],"Functions Label":["Background"],"Citation Start End":[[1604,1623]],"Functions Start End":[[1410,1603]]}
{"Identifier":"2019ApJ...885...79S__Burrows_et_al._1997_Instance_1","Paragraph":"However, this outward migration of the inner edge should stop at the corotation radius rco, where the Keplerian frequency of the disk equals the spin frequency of Jupiter. When rcav > rco, there will be two possibilities: the angular momentum will be transferred from Jupiter to the disk and then the gas accretion will stop, or otherwise the corotation radius and the disk edge will move outward together and then the accretion will continue (Takata & Stevenson 1996; Liu et al. 2017). Although the current corotation radius is rco \u2248 2.25 RJ, Jupiter at its time of formation was much larger than it is today (Burrows et al. 1997; Fortney et al. 2011), and this means that the corotation radius was also larger than the current one if the conservation of the angular momentum of Jupiter is assumed. Considering the transport of the angular momentum from Jupiter to the disk, the angular momentum should have been conserved since the disk disappeared. According to a formation model of Jupiter, the radius of the planet was \u22481.75 RJ after its rapid gas accretion and it decreased little by little (Lissauer et al. 2009). When the radius of Jupiter is 1.75 RJ, the corotation radius should be rco \u2248 4.7 RJ. We can then consider two scenarios of Io formation. In the first one, Io formed around r \u2248 4.7 RJ, slightly interior to the r = 5.89 RJ of our fiducial model, and then moved outward after the disk dissipated. The satellites, especially the inner ones, could move outward by the tidal force from Jupiter (Yoder & Peale 1981). The outer ones would be pushed by the inner ones and move outward with them because of the resonance. In this case, the position of the snow line should have been more inside than the fiducial case in this work, but this thermal condition could easily be reproduced by another parameter set. The second possibility is that Io was not the innermost satellite. If a body was present at r = 3.7 RJ, Io would have been situated at r = 5.7 RJ if they were trapped in a 2:1 resonance. This orbit is consistent with the corotation radius when the radius of Jupiter is \u22481.5 RJ, and this radius can be achieved during the contraction of Jupiter. The innermost body may have been broken by the tidal force of Jupiter when it has entered inside the Roche limit. Current Io, trapped in the Laplace resonance, actually moves inward little by little because of the tidal dissipation, and the innermost body may have also experienced such inward migration (Lainey et al. 2009).","Citation Text":["Burrows et al. 1997"],"Functions Text":["Although the current corotation radius is rco \u2248 2.25 RJ, Jupiter at its time of formation was much larger than it is today","and this means that the corotation radius was also larger than the current one if the conservation of the angular momentum of Jupiter is assumed. Considering the transport of the angular momentum from Jupiter to the disk, the angular momentum should have been conserved since the disk disappeared."],"Functions Label":["Uses","Uses"],"Citation Start End":[[611,630]],"Functions Start End":[[487,609],[654,951]]}
{"Identifier":"2021ApJ...911...89M__Mozer_et_al._2020a_Instance_1","Paragraph":"Time domain structures (TDSs; electrostatic or electromagnetic electron holes, ion holes, solitary waves, double layers, nonlinear whistlers, etc.) are \u223c1 ms pulses having significant electric fields parallel to the background magnetic field (Mozer et al. 2015). They are abundant through space, occurring along auroral zone magnetic field lines (Temerin et al. 1982; Mozer et al. 1997; Ergun et al. 1998), in the magnetospheric tail and plasma sheet (Cattell et al. 2005; Tong et al. 2018; Lotekar et al. 2020), at reconnection sites (Cattell et al. 2005; Steinvall et al. 2019; Lotekar et al. 2020), in the solar wind (Mangeney et al. 1999; Malaspina et al. 2013), in collisionless shocks (Wilson et al. 2010; Vasko et al. 2020; Wang et al. 2020), and in the magnetospheres of other planets (Pickett et al. 2015). TDSs are also expected along the Parker Solar Probe orbit (Mozer et al. 2020a). According to theoretical estimates and simulations (Cranmer & van Ballegooijen 2003; Valentini et al. 2011, 2014), these nonlinear structures can provide thermalization of electron and ion beams produced in the course of the turbulence cascade development at scales the order of the electron inertial length and down to the Debye length. This paper discusses such observations at a heliocentric distance of 35 solar radii. The electric field experiment on the Parker Solar Probe measures electric fields from DC to 20 MHz. A general description of the instrument and its electronics appears elsewhere (Bale et al. 2016). In this paper, the data from DC to 2 MHz are discussed. These measurements are obtained from the potentials of the four antennas, V1 through V4, that are located in the plane perpendicular to the Sun\u2013satellite line. They produce E12 = (V1\u2212V2)\/3.5 and E34 = (V3\u2212V4)\/3.5, which are then rotated into the spacecraft coordinate system to produce EX and EY. The direction, X, is perpendicular to the Sun\u2013spacecraft line, in the ecliptic plane, and pointing in the direction of solar rotation (against the ram direction), Y is perpendicular to the ecliptic plane, pointing southward, and Z points toward the Sun. The numerical factor, 3.5, is the effective antenna length (Mozer et al. 2020a). The uncertainties of the amplitudes of the waves and time domain structures reported in this paper are estimated to be about a factor of 2. These amplitudes are underestimated because the capacitive divider that couples the antennas to the electronics decreases the measured electric field relative to that on the antennas. They are often overestimated because short antennas produce overestimates of the electric field by factors of 2\u20134, as was observed during antenna deployment on the Cluster satellite and as is observed on the Parker Solar Probe from the ratio of the electric field to the magnetic field in whistlers.","Citation Text":["Mozer et al. 2020a"],"Functions Text":["TDSs are also expected along the Parker Solar Probe orbit"],"Functions Label":["Background"],"Citation Start End":[[875,893]],"Functions Start End":[[816,873]]}
{"Identifier":"2022AandA...662A..42M__N\u00f3brega-Siverio_et_al._2020a_Instance_1","Paragraph":"In this paper, we are interested not only in the mathematical properties of the ambipolar diffusion as a nonlinear diffusion process, but also in the inclusion of ambipolar diffusion terms in MHD codes. In astrophysics, over the past few decades, multidimensional MHD computer codes have been developed that model a variety of physical processes including ambipolar diffusion. Representative examples of such codes and simulations outside solar physics can be found in Basu & Mouschovias (1994), Mac Low et al. (1995), Padoan et al. (2000), Basu & Ciolek (2004), Kudoh & Basu (2008), Choi et al. (2009), Gressel et al. (2015), Tomida et al. (2015), O\u2019Sullivan & Downes (2007), Masson et al. (2012), Vigan\u00f2 et al. (2019), and Grassi et al. (2019). The consideration of ambipolar diffusion processes in solar physics started many decades ago (e.g. Parker 1963), but it has undergone a true explosion in terms of its use in large numerical models (e.g. Leake et al. 2005; Leake & Arber 2006; Arber et al. 2007; Cheung & Cameron 2012; Leake & Linton 2013; Mart\u00ednez-Sykora et al. 2012, 2017a,b, 2020a,b; Ni et al. 2015, 2016, 2021; Khomenko et al. 2017, 2018, 2021; Gonz\u00e1lez-Morales et al. 2018, 2020; N\u00f3brega-Siverio et al. 2020a,b; Popescu Braileanu & Keppens 2021). Such numerical calculations often encounter a problem: given the comparatively high values of \u03c7a in different cosmic environments, the advance in time may grind to a halt in magnetised regions when a standard Courant-Friedrichs-Lewy condition is adopted for the timestep based on \u03b7a. The recent papers by Gonz\u00e1lez-Morales et al. (2018) and N\u00f3brega-Siverio et al. (2020b) describe the construction of so-called super-time-stepping (STS) modules for the Mancha code and for the Bifrost code, respectively, designed with the aim to overcome that stiffness problem. In at least three of the papers cited above (Masson et al. 2012; Vigan\u00f2 et al. 2019; N\u00f3brega-Siverio et al. 2020b), the basic ZKBP solution in cylindrical coordinates was used to test the ambipolar diffusion module, given its simplicity and the analytical expression available for it. However, as already mentioned, the ambipolar diffusion problem tends to give rise to sharp current sheets and singularities. The ZKBP solution is comparatively smooth in that sense and it would be good to have some other canonical solutions on hand that include current sheets having higher degrees of singularity, which naturally occur in the ambipolar diffusion problem.","Citation Text":["N\u00f3brega-Siverio et al. 2020a"],"Functions Text":["The consideration of ambipolar diffusion processes in solar physics started many decades ago","but it has undergone a true explosion in terms of its use in large numerical models (e.g.","Such numerical calculations often encounter a problem: given the comparatively high values of \u03c7a in different cosmic environments, the advance in time may grind to a halt in magnetised regions when a standard Courant-Friedrichs-Lewy condition is adopted for the timestep based on \u03b7a."],"Functions Label":["Motivation","Motivation","Motivation"],"Citation Start End":[[1197,1225]],"Functions Start End":[[747,839],[860,949],[1264,1547]]}
{"Identifier":"2020AandA...643A..35P__Irwin_et_al._2007_Instance_2","Paragraph":"In order to achieve high photometric accuracy and be sensitive to low amplitude undulations, we adopted techniques from the exoplanet community, with the purpose of eliminating the systematic errors. When performing differential photometry (Sect. 3), accurate bias-subtraction and flat-fielding are of major importance. According to Irwin et al. (2007), the Poisson noise is 200 e\u2212 for a typical detector with a gain of a few e\u2212 ADU\u22121 and the flat illumination level is of 20 000 ADU pixel\u22121 = 40 000 e\u2212 pixel\u22121. Thus, a typical photometric aperture with a radius of 3 pixels contributes \u223c1 mmag photon noise. For this reason, we obtained a considerable amount of biases (150\u2212300 frames) and twilight flat-fields (25\u2212100 frames) each night to reduce the Poisson noise to less than 0.2 mmag (Irwin et al. 2007). The bias frames were averaged together using the minmax in the reject option of the zerocombine task in IRAF with a view of keeping radiation events out of the master bias frame. The master flat frame was the result of combining all the frames using a median mode. The median value is an excellent way of removing the effects of hot pixels and cosmic rays, so these extreme values do not affect the calculation, as they would, if they would averaged. The reject option was set to avsigclip, in which case the \u201ctypical\u201d sigma would have been determined from the data itself rather than an a priori knowledge of the noise characteristics of the CCD. Other related issues that can limit the photometric precision are: (i) the positioning of the telescope, (ii) fringing issues, and (iii) the differential variations on the quantum efficiency of the pixels. With the aim of minimizing the contribution of these effects, we repositioned each star almost on the same pixel of the detector using the autoguiding system of each telescope. The read-out-noise of the detectors are insignificant, as it can be as low as a few e\u2212 ( 10\u2006e\u2212 for RISE2 and Andor Zyla cameras).","Citation Text":["Irwin et al. 2007"],"Functions Text":["For this reason, we obtained a considerable amount of biases (150\u2212300 frames) and twilight flat-fields (25\u2212100 frames) each night to reduce the Poisson noise to less than 0.2 mmag"],"Functions Label":["Uses"],"Citation Start End":[[791,808]],"Functions Start End":[[610,789]]}
{"Identifier":"2022MNRAS.509.3339K___2011_Instance_1","Paragraph":"Semi-analytical disc models indicate that the frequency of giant planets must increase with the mass of the host star between 0.2\u20131.5 M\u2299 (Ida & Lin 2005; Kennedy & Kenyon 2008). However, this trend is expected to decrease above 1.5 M\u2299 due to a smaller growth rate, longer migration time-scale, and shorter lifetime of the protostellar disc (Reffert et al. 2015). Looking for planets around main sequence stars more massive than the sun can help shed some light on this aspect. These stars have few spectral lines for Doppler measurements and are often broadened by the rapid rotation of the star. This has been the reason for RV surveys to have traditionally targeted slow rotating FGK type stars. However, as rapidly rotating stars evolve off the main sequence they slow down considerably and become much cooler making it relatively easy to search for planets around them. This fact was exploited by dedicated planet searches around intermediate mass sub-giants leading to dozens of planet discoveries (Johnson et al. 2007, 2010a, 2011). An important result obtained from the survey of giant stars at the Lick observatory pointed out that the occurrence rate peaks at a stellar mass of $1.9^{+0.1}_{-0.5}$ M\u2299. However, many of the discovered planets around evolved stars were found at large orbital separations (Hatzes et al. 2003; Fischer et al. 2007; Robinson et al. 2007; Johnson et al. 2008). This is not a surprise as star\u2013planet interaction is largely governed by tidal forces. When the stellar rotational period is longer than the planet orbital period, the star experiences spinning up, leading to orbital decay. Synchronization and circularization of orbit occurs in systems where the total angular momentum exceeds a critical value. When this total angular momentum is small enough, the orbit of the planet can continue to shrink and be engulfed by the host star. This phenomenon entirely depends on the dissipation time-scales for the star (Mazeh 2008, and references therein). The role of tidal forces becomes increasingly important in the context of host stars being in an evolved state. There is a higher chance of the planet being destroyed by the evolved star (Kunitomo et al. 2011; Schlaufman & Winn 2013). However, there is no obvious way to estimate these tidal dissipation forces. The circularization time-scale for such planets can be used to quantify tidal dissipation inside planets (Hansen 2010; Socrates et al. 2012). Most of the discovered hot Jupiters with periods less than 3 d are found to be on circular orbits. We calculate the circularization time-scales for the orbit of TOI-1789 b, which is \u03c4cir  = 0.08 Gyr (for QP = 106, equation 3 of Adams & Laughlin 2006). This is less than the age of the star as calculated from our work (Section 3.3.1).","Citation Text":["Johnson et al.","2011"],"Functions Text":["However, as rapidly rotating stars evolve off the main sequence they slow down considerably and become much cooler making it relatively easy to search for planets around them. This fact was exploited by dedicated planet searches around intermediate mass sub-giants leading to dozens of planet discoveries"],"Functions Label":["Background"],"Citation Start End":[[1004,1018],[1032,1036]],"Functions Start End":[[698,1002]]}
{"Identifier":"2022MNRAS.515..185O__Gutcke_et_al._2022_Instance_1","Paragraph":"While gas cooling and stellar feedback can transform dark matter cusps to cores, it is energetically challenging for this process to create large dark matter cores (typically >500\u2009pc) in the very smallest \u2018ultra-faint\u2019 dwarfs, since they form so few stars (M*  105\u2009M\u2299; Pe\u00f1arrubia et al. 2012; Garrison-Kimmel et al. 2013; Di Cintio et al. 2014; Maxwell, Wadsley & Couchman 2015; O\u00f1orbe et al. 2015; Tollet et al. 2016). However, smaller dark matter cores may still form inside the half-light radius (R1\/2 \u223c 20\u2013200\u2009pc for UFDs2), where the gravitational potential fluctuations are strongest (e.g. O\u00f1orbe et al. 2015; Read et al. 2016). Whether this is expected to happen in a \u039bCDM cosmology remains an active area of debate. Most studies to date find that cusp-core transformations are challenging at the likely mass-scale of UFDs (M200c \u223c 109\u2009M\u2299; M* \u223c 105\u2009M\u2299; e.g. Chan et al. 2015; Wheeler et al. 2019a; Gutcke et al. 2022). However, there are some notable exceptions. In recent work, Orkney et al. (2021) studied cusp-core transformations for UFDs drawn from the \u2018Engineering Dwarfs at Galaxy formation\u2019s Edge\u2019 (EDGE) simulation project (with a mass, baryonic mass, and spatial resolution of 120\u2009M\u2299, 20\u2009M\u2299, and 3\u2009pc, respectively, sufficient to resolve dark matter cores in UFDs larger than \u223c20\u2009pc). They found that UFDs can lower their inner dark matter density by up to a factor \u223c2 through a combination of early heating due to star formation, followed by late-time heating from minor mergers. While none of their simulated dwarfs formed a completely flat central core, their small sample size left open the possibility that this combination of mechanisms could form flatter cores for some rarer assembly histories. It is also important to note that dark matter core formation on these mass scales is very sensitive to small changes in the star formation and feedback modelling. Pontzen et al. (2021) showed that a small increase in variability of the star formation rate on time-scales shorter than the local dynamical time is sufficient to form a full dark matter core in an ultra-faint, without significantly altering its stellar mass. As such, the question of whether complete dark matter core formation is expected in some or all UFDs in \u039bCDM remains open.","Citation Text":["Gutcke et al. 2022"],"Functions Text":["Most studies to date find that cusp-core transformations are challenging at the likely mass-scale of UFDs (M200c \u223c 109\u2009M\u2299; M* \u223c 105\u2009M\u2299; e.g."],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[905,923]],"Functions Start End":[[724,864]]}
{"Identifier":"2018ApJ...854..137S__Manuel_et_al._2014_Instance_1","Paragraph":"Ulysses was launched on 1990 October 6 and orbited the Sun with a latitude varying from \u221280\u00b0 to 80\u00b0 and a solar distance ranging from \u223c1 au to \u223c5 au (Heber et al. 2009). The Kiel Electron Telescope (KET) on board Ulysses measured electrons in the energy range from \u223c3 MeV to above 300 MeV, and protons and helium nuclei in the energy range from \u223c5 MeV\/nuc to above 2 GeV\/nuc (Simpson et al. 1992). The data of the KET coincidence channel K12, which measures protons with energy in 0.25\u20132.0 GeV, has been used to study the GCR modulation outside the solar ecliptic plane in many works (see, e.g., Ndiitwani et al. 2005; Vos & Potgieter 2016; Boschini et al. 2017a). However, different works used different mono-energetic bins to represent this channel, e.g., 1.08 GeV (Rastoin et al. 1996), 2.5 GV (i.e., 1.73 GeV, Ndiitwani et al. 2005; Heber et al. 2009; Manuel et al. 2014), and 2.2 GeV (Boschini et al. 2017a). As the KET observations are integrated over a large energy interval (de Simone et al. 2011), it is perhaps better to weight the model results of several energy bins with the Ulysses response function and then combine them (Boschini et al. 2017a). Heber et al. (2009) obtained the 1 au equivalent count rates for this channel by correcting the proton intensity for the global spatial gradients of GCR protons, and the 1 au equivalent GCR proton flux can be obtained with the corresponding response factor. In addition, the precise cosmic-ray spectra measured by the PAMELA instrument can help us to roughly estimate the effective energy of the KET coincidence channel K12. In Figure 8, the black solid line shows the monthly averaged 0.25\u20132.0 GeV proton flux observed by Ulysses. Note that the original daily count rates from the Ulysses Final Archive (ufa.esac.esa.int\/ufa) are divided by the corresponding response factor to obtain the proton flux. The dash-dotted line represents the 1 au equivalent 0.25\u20132.0 GeV proton flux, and the relevant count rates are digitized from Figure 5 in Heber et al. (2009). The 1.2 GeV proton flux observed by SOHO\/EPHIN (K\u00fchl et al. 2016) and PAMELA (Adriani et al. 2013) is shown as magenta triangles and green circles, respectively. The 1 au equivalent GCR proton flux roughly matches the 1.2 GeV GCR proton observations, and 1.2 GeV can be used to represent the Ulysses\/KET channel K12. Therefore, we compute the 1.2 GeV proton flux along the trajectory of Ulysses to compare it with the Ulysses K12 measurements.","Citation Text":["Manuel et al. 2014"],"Functions Text":["However, different works used different mono-energetic bins to represent this channel, e.g.,","2.5 GV"],"Functions Label":["Differences","Differences"],"Citation Start End":[[856,874]],"Functions Start End":[[665,757],[790,796]]}
{"Identifier":"2015MNRAS.452.2731S__Stroe_et_al._2013_Instance_4","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"],"Functions Text":["This is probably ram pressure caused by the relative motion of galaxies with respect to the ICM"],"Functions Label":["Background"],"Citation Start End":[[2017,2034]],"Functions Start End":[[1920,2015]]}
{"Identifier":"2019MNRAS.490.5478W__Kurtovic_et_al._2018_Instance_1","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":["Kurtovic et al. 2018"],"Functions Text":["The influence of tidal truncation is therefore limited to stellar multiples, either in bound systems"],"Functions Label":["Background"],"Citation Start End":[[883,903]],"Functions Start End":[[764,864]]}
{"Identifier":"2020MNRAS.494.5110B__Troja_et_al._2018_Instance_2","Paragraph":"Following the short gamma-ray burst (sGRB) associated with this event, GRB 170817A (Abbott et al. 2017a,b; Goldstein et al. 2017), radio emission was anticipated as the associated merger outflow interacted with the circum-merger medium. Monitoring the radio emission could therefore provide crucial information on the energetics and geometry of the outflow, as well as the ambient environment. At radio frequencies, telescopes were observing the Advanced LIGO\u2013Virgo probability region for GW170817 within 29\u2009min post-merger (Callister et al. 2017a), and subsequent monitoring of AT 2017gfo resulted in an initial radio detection 16\u2009d after the event (Abbott et al. 2017a; Hallinan et al. 2017). Further monitoring, predominantly at frequencies between 0.6 and 15\u2009GHz, has since taken place (e.g. Alexander et al. 2017, 2018; Corsi et al. 2018; Dobie et al. 2018; Margutti et al. 2018; Mooley et al. 2018a,b,c;Resmi et al. 2018; Troja et al. 2018, 2019). At these frequencies, a general picture emerged in which the radio light curve was first observed to steadily rise, before it turned over and began a more rapid decay. Using a compilation of 0.6\u201310\u2009GHz radio data from 17 to 298\u2009d post-merger, Mooley et al. (2018c) derived both a fitted time for the radio peak of 174$^{+9}_{-6}$ d and a fitted 3-GHz peak flux density of 98$^{+8}_{-9}\\, \\mu$Jy (also see similar analyses in Dobie et al. 2018 and Alexander et al. 2018). The fitted radio spectral index \u03b11 from this study is \u22120.53 \u00b1 0.04, consistent with broad-band spectral indices determined using radio, optical, and X-ray data at various epochs, where the typical value is approximately \u22120.58 (e.g. Alexander et al. 2018; Margutti et al. 2018; Troja et al. 2018, 2019; Hajela et al. 2019). Mooley et al. (2018c) also found power-law dependencies for the rise and decay phases of approximately t0.8 and t\u22122.4, respectively, where t is the time since the merger. Within the associated uncertainties, these results are consistent with the broad-band evolution of AT 2017gfo (e.g. Alexander et al. 2018; Hajela et al. 2019; Lamb et al. 2019; Troja et al. 2019).","Citation Text":["Troja et al. 2018"],"Functions Text":["The fitted radio spectral index \u03b11 from this study is \u22120.53 \u00b1 0.04, consistent with broad-band spectral indices determined using radio, optical, and X-ray data at various epochs, where the typical value is approximately \u22120.58 (e.g."],"Functions Label":["Similarities"],"Citation Start End":[[1702,1719]],"Functions Start End":[[1425,1656]]}
{"Identifier":"2021ApJ...922..224L__Jain_et_al._2015_Instance_1","Paragraph":"Theoretically, coronal loops are made up of magnetically confined denser and hotter plasma, and can thus support modes of oscillations (Aschwanden 1987; Roberts 2008). These oscillations are closely related to the physical properties of their host coronal loops and may open a new window for detecting the inhomogeneous corona (Nakariakov & Verwichte 2005; Yuan & Van Doorsselaere 2016a; Li & Liu 2018). Theories of magneto-acoustic oscillations in coronal loops have indicated that the fast kink mode possesses both a unique dispersion relation and zero asymptotic dispersion for a long-wave approximation, and could provide a potential diagnostic tool for determining physical properties of the inhomogeneous corona (Edwin & Roberts 1983; Roberts et al. 1984). Due to the specific role of coronal loops\u2019 fast kink oscillations in coronal seismology, the proposal of using the fast kink oscillations to diagnose the physical properties of their host coronal loops has attracted a great deal of attention since it was first introduced (Roberts et al. 1984). Along with an increasing number of observational cases, the study of fast kink oscillations in coronal loops has been significantly advanced (Aschwanden et al. 1999; Nakariakov et al. 1999; Verwichte et al. 2004; Aschwanden & Schrijver 2011; Wang et al. 2012; Jain et al. 2015; Yuan & Van Doorsselaere 2016b; Shen et al. 2017; Pascoe et al. 2018; Shen et al. 2018; Nechaeva et al. 2019; Anfinogentov & Nakariakov 2019). Numerous detailed observations, on the other hand, have also brought new challenges to the current theory (De Moortel & Brady 2007; Guo et al. 2015; Li et al. 2017; Su et al. 2018; Li et al. 2019). Among these challenges, a prominent one is the increasingly presented oscillating coronal loops whose frequencies show significant changes, which is also known as the frequency drift (Su et al. 2018). It indicates that the frequency drift for the oscillation in fast kink mode can occur in a quiet loop (Su et al. 2018) or a quiet fiber (Li et al. 2018). It also indicates that the brightness of the oscillating structures undergoes a significant change during the occurrence of the frequency drift, which may imply changes in the loop\u2019s thermal properties. A detailed investigation present in Table 1 and Figure 5 of Su et al. (2018) indicates that the period of fast kink oscillation has an \u224825% increase corresponding to an \u224840% increase in the loop\u2019s density. A decrease in the period of fast kink oscillation is also observed by Li et al. (2018). It is also worth noting that the oscillation reported by Li et al. (2018) shows a significant increase of the amplitude as the period decreases, which is very unusual.","Citation Text":["Jain et al. 2015"],"Functions Text":["Along with an increasing number of observational cases, the study of fast kink oscillations in coronal loops has been significantly advanced"],"Functions Label":["Background"],"Citation Start End":[[1318,1334]],"Functions Start End":[[1058,1198]]}
{"Identifier":"2017ApJ...837...97L__Grillo_et_al._2015_Instance_1","Paragraph":"The newly discovered arcs and new spectroscopic redshifts have been incorporated into updated HFF+ versions of the Abell 2744 and MACSJ0416.1-2403 lensing models (Table 5); many of these have \n\n\n\n\n\n (see Figure 8 for a comparison of the arc redshift distributions adopted by the pre-HFF and new HFF+ lensing models; Cypriano et al. 2004; Okabe & Umetsu 2008; Zitrin et al. 2009; 2013; Okabe et al. 2010a, 2010b; Merten et al. 2011; Christensen et al. 2012; Mann & Ebeling 2012; Jauzac et al. 2014; Lam et al. 2014; Richard et al. 2014; Balestra et al. 2015; Diego et al. 2015; Grillo et al. 2015; Jauzac et al. 2015; Rodney et al. 2015; Wang et al. 2015; Kawamata et al. 2016). The incorporation of these new multiple image systems often results in a reduction in the statistical uncertainty in the galaxy magnifications for a given model. All of the public HFF lensing models provide a range of possible realizations from which the statistical uncertainty of a given model set may be calculated (typically 100 but no fewer than 30). We plot the cumulative distribution of the galaxy magnification uncertainties \u03c3(model)\/\n\n\n\n\n\n, for the galaxies and photometric redshifts provided by the ASTRODEEP catalogs (Merlin et al. 2016; Castellano et al. 2016a) for Abell 2744 (Figure 9) and MACSJ0416.1-2403 (Figure 10). Generally, the statistical uncertainties are reduced for the models computed with the new HFF data sets, with more dramatic reductions for the methods that rely strongly upon the strong-lensing constraints. The parametric methods (CATS, Sharon, Zitrin, GLAFIC) report median statistical magnification errors of 0.2%\u20135%, while the non-parametric methods (Bradac Williams, Diego) report median statistical magnification errors of 2%\u201311% for the post-HFF calculations (green curves), versus 2%\u201322% and 2%\u201317% respectively for pre-HFF models (blue curves). (We note that the statistical errors for the MACSJ0416.1-2403 Bradac post-HFF models (Hoag et al. 2016) included additional uncertainties due to the photometric redshift uncertainties of the multiple images. These were not included in the pre-HFF Bradac model, and thus may explain why the post-HFF statistical errors are larger for this model.)","Citation Text":["Grillo et al. 2015"],"Functions Text":["see Figure 8 for a comparison of the arc redshift distributions adopted by the pre-HFF and new HFF+ lensing models;"],"Functions Label":["Uses"],"Citation Start End":[[577,595]],"Functions Start End":[[200,315]]}
{"Identifier":"2022AandA...666A.190S__Velichko_et_al._1995_Instance_1","Paragraph":"For our dataset of absolute magnitudes, we used data collected at the Institute of Astronomy of V. N. Karazin Kharkiv National University within the long-term observational programme to study asteroid magnitude-phase curves (Shevchenko et al. 2010, 2012, 2014a, 2016; Slyusarev et al. 2012). We also used some observational data obtained within several other programmes (Belskaya et al. 2010; Chiorny et al. 2007, 2011; Dotto et al. 2009; Hahn et al. 1989; Kaasalainen et al. 2004; Lagerkvist et al. 1998; Michalowski et al. 1995; Mohamed et al. 1994, 1995; Oszkiewicz et al. 2021; Shevchenko et al. 1992, 2003, 2009, 2014b, 2021; Velichko et al. 1995; Wilawer et al. 2022). All magnitudes were measured in the Johnson V band and extrapolated to zero phase angle using the HG1G2 system proposed by Muinonen et al. (2010), with some modifications presented by Penttil\u00e4 et al. (2016). For computations, the online calculator1 for the HG1G2 photometric system was used. Since we derived absolute magnitudes in our data from the light curve maxima, and the definition of H is based on the rotationally averaged brightness, we added a half of the light curve amplitude corrected to zero phase angle to our results. We used the average correction coefficients from Zappala et al. (1990) for low- and moderate-albedo asteroids. This correction is typically very small because our light curve observations covered small phase angles. Absolute magnitudes obtained at different aspects were averaged. In such a manner, we obtained a homogeneous dataset of absolute magnitudes of about 400 asteroids up to H = 16.5 mag. Our database includes the absolute magnitude data, the G1 and G2 parameters, and the albedo and diameter values from different databases (such as Tedesco et al. 2002; Masiero et al. 2011, 2012; Nugent et al. 2015; Usui et al. 2011). The database is available at the CDS. Figure 1 shows the correlations of the absolute magnitudes from the largest datasets (MPC (HMPC), Pan-STARRS (HPS), and ATLAS (HATLAS)) with those of the Kharkiv dataset (HKH). For the ATLAS dataset, we used the absolute magnitudes in a cyan filter, since this filter overlaps the Johnson V band (Mahlke et al. 2021).","Citation Text":["Velichko et al. 1995"],"Functions Text":["We also used some observational data obtained within several other programmes"],"Functions Label":["Uses"],"Citation Start End":[[631,651]],"Functions Start End":[[292,369]]}
{"Identifier":"2015MNRAS.451.4290S__Governato_et_al._2004_Instance_1","Paragraph":"Hydrodynamical simulations of evolving galaxies allow us to calibrate these diagnostics by measuring their observability given a set of formation scenarios and physical processes (e.g. Jonsson et al. 2006; Rocha et al. 2007; Lotz et al. 2008a; Bush et al. 2010; Narayanan et al. 2010; Hayward et al. 2013; Snyder et al. 2013; Lanz et al. 2014). The quality and breadth of these experiments are limited by the availability of computational resources and the fidelity of models for galaxy physics such as star formation, supernovae, and the interstellar medium (ISM). It has only recently become widespread to model the formation of galaxies ab initio (e.g. Governato et al. 2004; Agertz, Teyssier & Moore 2011; Guedes et al. 2011; Marinacci, Pakmor & Springel 2013; Ceverino et al. 2014), and the realism continues to improve (Stinson et al. 2012; Hopkins et al. 2014; Torrey et al. 2014), albeit with still widely varying physics models (e.g. Scannapieco et al. 2012; Kim et al. 2014). Prior to these advances, studies were limited to small numbers of isolated galaxies or mergers to inform common diagnostics of galaxy evolution, an approach with a significant limitation: they do not fully account for cosmological context, such as gas accretion and the breadth of assembly histories. In addition to mergers, models of high-redshift galaxy formation (e.g. Dekel, Sari & Ceverino 2009; Dekel et al. 2013) have recently appreciated the tight coupling between gas accretion and disc evolution (e.g. Cacciato, Dekel & Genel 2012; Danovich et al. 2012; Dekel & Krumholz 2013), as well as bulge and super-massive black hole (SMBH) growth mediated by turbulent motions or violent disc instability (e.g. Bournaud et al. 2011; Porter et al. 2014) and the evolution of giant clumps (Dekel & Burkert 2013). These important processes likely complicate interpretation of a given observation, and recent studies of galaxy morphology have begun to exploit simulations including them (e.g. Scannapieco et al. 2010; Pedrosa, Tissera & De Rossi 2014).","Citation Text":["Governato et al. 2004"],"Functions Text":["It has only recently become widespread to model the formation of galaxies ab initio (e.g."],"Functions Label":["Background"],"Citation Start End":[[656,677]],"Functions Start End":[[566,655]]}
{"Identifier":"2022MNRAS.515.5495M__Genel_2016_Instance_1","Paragraph":"The stellar metallicity in the Universe evolves with redshift (Mannucci et al. 2010; Sommariva et al. 2012; Krumholz & Dekel 2012; Dayal, Ferrara & Dunlop 2013; Madau & Dickinson 2014). The metallicity at a high redshift (z > 2) is much smaller in comparison to the low redshift Universe z  2. The first-generation stars contaminate the interstellar medium and cause a chemical evolution of the Universe. We can treat the metallicity evolution with redshift by a relation \n(2)$$\\begin{eqnarray*}\r\n\\log _{10}(Z(z))= \\gamma z +\\zeta ,\r\n\\end{eqnarray*}$$where \u03b3 captures the redshift dependence and \u03b6 captures the metallicity value at z = 0 (Mannucci et al. 2010; Madau & Dickinson 2014). This relation captures the metallicity of the parent star or the gas cloud from which a star has formed. It is written to express only a mean evolution of the metallicity. Along with the mean metallicity evolution of the Universe, there is going to be a scatter in the metallicity depending on the galaxy properties. Such a source of uncertainty brings additional stochasticity to the metallicity relation. Currently, a limited number of observations (Gallazzi et al. 2008; Mannucci et al. 2010; Krumholz & Dekel 2012) are available to explore the environment dependence of the metallicity, and most of our current understandings are based on simulations(Genel 2016; Torrey et al. 2019). These studies show that the overall median metallicity dependence of the galaxies at different redshifts can be explained by power form (Pei, Fall & Hauser 1999; Young & Fryer 2007; Torrey et al. 2019). Several studies of GW merger rates and mass distribution are performed (Belczynski et al. 2002; Dominik et al. 2012, 2015; Mapelli et al. 2017; Giacobbo, Mapelli & Spera 2018; Toffano et al. 2019; van Son et al. 2022) which are motivated by these studies and show that the black hole mass distribution can exhibit a redshift dependence. The existence of any stochasticity in the galaxy metallicity distribution will also influence the mass distribution but is currently not well known. However, as the relation given in equation (1) is in terms of the logarithm of metallicity, so the impact of fluctuation around the median value depending on the individual galaxy properties is going to be a small (logarithmic) change. As we are unable to measure the host of the BBH due to a large sky localization error of the BBH, we cannot directly associate the properties of galaxies with BBH source properties. So, we can only infer an ensemble average mass distribution from the GW data and the additional stochasticity (which will depend on the host properties) will appear as an additional uncertainty in the measurement of MPISN. As a result, we consider a median distribution of galaxy metallicity and the dependence of MPISN on it.","Citation Text":["Genel 2016"],"Functions Text":["Currently, a limited number of observations","are available to explore the environment dependence of the metallicity, and most of our current understandings are based on simulations"],"Functions Label":["Background","Background"],"Citation Start End":[[1341,1351]],"Functions Start End":[[1093,1136],[1205,1340]]}
{"Identifier":"2022AandA...664A.127S__Lefebvre_et_al._2008_Instance_1","Paragraph":"Efforts have been made to model the signals caused by different sources of stellar variability within the RV time series (e.g., Tuomi et al. 2013; Rajpaul et al. 2015; Davis et al. 2017; Simola et al. 2019). Several solutions have been successfully proposed in order to deal with stellar oscillations and granulation phenomena, such as: calculating stellar evolution sequences (e.g., Christensen-Dalsgaard et al. 1995); fitting a two-level structure tracking (TST) algorithm based on a two-level representation of granulation (Del Moro 2004); using daytime spectra of the Sun in order to measure the solar oscillations (e.g., Kjeldsen et al. 2008; Lefebvre et al. 2008); and characterizing the statistical properties of magnetic activity cycles focusing on HARPS observations (e.g., Pepe et al. 2011; Dumusque et al. 2011b). However, properly modeling the other sources of stellar activity remains extremely challenging (e.g., Nava et al. 2019). In the present work, we deal with the cross correlation function (CCF) that is derived from the stellar spectrum (e.g., Hatzes 1996; Hatzes & Cochran 2000; Fiorenzano et al. 2005). As it is well known, the CCF barycenter estimates the RV of the star. The asymmetry and the full width at half maximum (FWHM) of the CCF give a strong indication for stellar activity, meaning that variations in RV are caused by active regions rather than by an exoplanet (e.g., Hatzes 1996; Queloz et al. 2001; Boisse et al. 2011; Figueira et al. 2013; Simola et al. 2019). Several solutions have been successfully proposed for mitigating stellar activity perturbations when working with RV measurements, including: decorrelating RV data against activity indicators such as log ${{R'}_{{\\rm{HK}}}}$R\u2032HK (e.g., Wilson 1968; Noyes et al. 1984) or H\u03b1 (e.g., Robertson et al. 2014); modeling stellar activity by fitting Gaussian processes (GPs, Rasmussen & Williams 2005; Haywood et al. 2014; Rajpaul et al. 2015); or moving averages (e.g., Tuomi et al. 2013) to the RV data. A common statistic employed for identifying changes in the shape of the CCF is the bisector span (e.g., Hatzes 1996; Queloz et al. 2001).","Citation Text":["Lefebvre et al. 2008"],"Functions Text":["Several solutions have been successfully proposed in order to deal with stellar oscillations and granulation phenomena, such as:","using daytime spectra of the Sun in order to measure the solar oscillations (e.g.,"],"Functions Label":["Background","Background"],"Citation Start End":[[648,668]],"Functions Start End":[[208,336],[543,625]]}
{"Identifier":"2021MNRAS.507.5567D__Guo_et_al._2019_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":["Guo et al. 2019"],"Functions Text":["This picture is consistent with the observed low temperature, large radius, and super-Eddington luminosity (see"],"Functions Label":["Similarities"],"Citation Start End":[[1225,1240]],"Functions Start End":[[1053,1164]]}
{"Identifier":"2021MNRAS.500..786D__Delgado_&_Perez_1997_Instance_1","Paragraph":"\nxstar predicts O to be progressively ionized through all its possible ionization stages as the central UV-X-ray source is approached providing a natural explanation for the absence of broad forbidden ${[}{\\rm O\\, \\small{III}}{]}$ emission lines. No $[{\\rm O\\, \\small{II}}]\\, {\\lambda }\\, {\\lambda }$3727, 3729 emission lines appear in the HST spectra, consistent with the xstar model. Furthermore, both the observed reddening-insensitive ratio $[{\\rm O\\, \\small{III}}]\\, {\\lambda }\\, 5008\/[{\\rm O\\, \\small{III}}]\\, {\\lambda }$4960 and the value predicted by xstar are consistent with the theoretical value (Dimitrijevic et al. 2007). However, there are two notable discrepancies between the HST observations and the xstar model predictions. The first is that the $[{\\rm O\\, \\small{III}}]\\, {\\lambda }$5008 emission line is observed to be 230 per cent brighter than predicted by xstar. This excess emission could be attributed to the extended source of $[{\\rm O\\, \\small{III}}]\\, {\\lambda}$5008 seen in the nucleus of NGC 3327 (Arribas & Mediavilla 1994; Mundell et al. 1995; Delgado & Perez 1997; Garc\u00eda-Lorenzo, Mediavilla & Arribas 2001; Walsh et al. 2008). However, the second discrepancy that xstar predicts the $[{\\rm O\\, \\small{III}}]\\, {\\lambda }$4364 emission line to be \u223c25 times brighter than the upper limit estimated in D13 is not as easily addressed as the first. On the one hand, the $[{\\rm O\\, \\small{III}}]\\, {\\lambda }$5008\/\u03bb4364 ratio produced by xstar is entirely consistent with that expected (Osterbrock 1989), given the temperature \u223c104 K and electron density \u223c108 cm\u22123 in the O++ region. On the other hand, the observed $[{\\rm O\\, \\small{III}}]\\, {\\lambda }$4364 emission line would have to be as bright as H\u2009\u03b3 in order to match the xstar prediction. Observationally, this seems unlikely, but the resolution of the ${\\it {\\it HST}}$ observations is insufficient to adequately resolve the two lines in question (D13). Equally disconcerting, however, is that contrary to commonly accepted wisdom, xstar predicts the $[{\\rm O\\, \\small{III}}]\\, {\\lambda }$4364 emission line to be bright even though the electron density in the O++ region exceeds the critical density of the 1S0 level responsible for the transition by about one order of magnitude.","Citation Text":["Delgado & Perez 1997"],"Functions Text":["However, there are two notable discrepancies between the HST observations and the xstar model predictions. The first is that the $[{\\rm O\\, \\small{III}}]\\, {\\lambda }$5008 emission line is observed to be 230 per cent brighter than predicted by xstar. This excess emission could be attributed to the extended source of $[{\\rm O\\, \\small{III}}]\\, {\\lambda}$5008 seen in the nucleus of NGC 3327"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1076,1096]],"Functions Start End":[[635,1026]]}
{"Identifier":"2019AandA...621A..27F__DeGraf_et_al._2017_Instance_1","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"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[725,743]],"Functions Start End":[[548,701]]}
{"Identifier":"2021ApJ...919..133P__Ravishankar_et_al._2019_Instance_1","Paragraph":"The inversion of the Fourier transform from limited data is a well-known problem in several imaging domains like, for instance, medical imaging (McGibney et al. 1993; Bronstein et al. 2002; Sutton et al. 2003; Fessler 2007; Gallagher et al. 2008; Lustig et al. 2008), crystallography (Eisebitt et al. 2004; Marchesini et al. 2008; Brady et al. 2009), and geophysics (Brossier et al. 2009; Jin 2010). This image reconstruction problem inspired several computational approaches like nonuniform Fast Fourier Transform (FFT) (Bronstein et al. 2002; Fessler & Sutton 2003; Greengard & Lee 2004; Lee & Greengard 2005), compressed sensing (Donoho 2006; Lustig et al. 2008; Bigot et al. 2016), and machine learning (Wang et al. 2018; Ravishankar et al. 2019). In the case of astronomical imaging, the use of Fourier methods is mainly related to radio and optical interferometry (Le Besnerais et al. 2008; Thi\u00e9baut & Giovannelli 2009; Wiaux et al. 2009; Felli & Spencer 2012; Thompson et al. 2017; Ye et al. 2020), although a similar methodology also involves snapshot imaging spectroscopy (Culhane et al. 2007; Harra et al. 2017; Courrier & Kankelborg 2018; Winebarger et al. 2019). However, in the last three decades, this approach has been utilized also in the case of solar hard X-ray telescopes that have been conceived in order to provide spatial Fourier components of the photon flux emitted via either the bremsstrahlung or thermal processes during solar flares (Lin et al. 2002; Krucker et al. 2020). These Fourier components, named visibilities, are sampled by the hard X-ray instrument in the two-dimensional Fourier space, named the (u, v) plane, in a sparse way, according to a geometry dependent on the instrument design. For instance, NASA\u2019s Reuven Ramaty High Energy Spectroscopic Imager (RHESSI) relies on the use of a set of nine Rotating Modulation Collimators (RMCs) whose FWHM is logarithmically spaced between 23 and 183\u2033 (Hurford et al. 2002). Each RMC measures visibilities on a circle of points in the (u, v) space with a spatial frequency that corresponds to its angular resolution and a position angle that varies according to the spacecraft rotation (see Figure 1, left panel). On the other hand, the Spectrometer\/Telescope for Imaging X-rays (STIX) on board ESA\u2019s Solar Orbiter is based on the Moir\u00e9 pattern technology (Giordano et al. 2015; Massa et al. 2019), and its 30 collimators sample the (u, v) plane over a set of six spirals for an FWHM resolution coarser than 7\u2033 (see Figure 1, right panel).","Citation Text":["Ravishankar et al. 2019"],"Functions Text":["This image reconstruction problem inspired several computational approaches like","and machine learning"],"Functions Label":["Background","Background"],"Citation Start End":[[726,749]],"Functions Start End":[[400,480],[686,706]]}
{"Identifier":"2019MNRAS.485.3600A__Bignall_et_al._2015_Instance_1","Paragraph":"Having eliminated calibration or observational effects as the cause of the spectral changes, it must be that either (1) the configuration of the sources themselves has changed or (2) the interstellar scintillation has caused time-variable focusing and defocusing of the polarized substructure in the radio jets (e.g. Rickett 2001; de Bruyn & Macquart 2015), which can generate spectropolarimetric variability (e.g. Kedziora-Chudczer 2006). Interstellar scintillation (ISS) is difficult to distinguish from intrinsic variability without fully time-resolved data (see e.g. Bignall et al. 2015 and references therein), and can coexist with it (Koay et al. 2018). Never the less, ISS does not easily explain our results: Sources located more than \u223c25\u2009degrees from the Galactic plane (3\/4 of our polarized sources) and with line-of-sight H\u2009\u03b1 intensities less than \u223c a few Rayleighs (7\/8 of our polarized sources; see column 9 of Table 1) are typically not strongly scattered (Pushkarev & Kovalev 2015), and experience associated flux density modulations of typically less than a few per\u2009cent (Heeschen 1984; Quirrenbach et al. 1992; Rickett, Lazio & Ghigo 2006; Lovell et al. 2008). This is smaller than the changes observed in our integrated fractional polarization spectra, which should undergo less modulation than the Stokes I in any case. ISS modulation also has a distinct frequency and time dependence. At the mid-galactic latitudes inhabited by most of our sources (see column 8 of Table 1), its magnitude peaks between 4 and 6\u2009GHz (\u03bb2 between 0.0024 and 0.0056\u2009m2), and drops to less than 50 per\u2009cent of this value at 1.4 and 10\u2009GHz (i.e. \u03bb2 of 0.0008 and 0.046\u2009m2). Thus, changes due to ISS should \u2018spike\u2019 in this narrow \u03bb2 window, which we see little evidence of (though see perhaps PKS B0517\u2013726 and PKS B1903\u2013802; see Figs 3 and 9). The characteristic time-scale of ISS is \u223chours to days (Rickett et al. 2006; Gab\u00e1nyi et al. 2007, and references therein), and over multiyear time-scales, intrinsic effects tend to dominate the observed variability (e.g. Lazio et al. 2001; Rickett et al. 2006; Mooley et al. 2016). We therefore claim that the spectral variability most likely reflects intrinsic changes in the sources themselves.","Citation Text":["Bignall et al. 2015"],"Functions Text":["Interstellar scintillation (ISS) is difficult to distinguish from intrinsic variability without fully time-resolved data (see e.g.","and references therein), and can coexist with it"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[571,590]],"Functions Start End":[[440,570],[591,639]]}
{"Identifier":"2018AandA...610A..15J__Tregloan-Reed_et_al._2015_Instance_1","Paragraph":"Figure 2 shows the geometry used within our model. The center of the stellar sphere is located at the origin of the three-dimensional spherical coordinate system. PyTranSpot does not take into account a fractional area correction, as used within the WD code (Wilson & Devinney 1971; Wilson 1979, 1990, 2008, 2012). We argue that this effect is negligible, as the resulting loss of accuracy is much smaller than the noise currently present in observations. However, to obtain precise photometric transit and spot parameters, we recommend the use of a planetary pixel radius between 15 and 50 pixels (Tregloan-Reed et al. 2015). On the stellar sphere, every point is described by the two angles (longitude \u03b8, co-latitude \u03c6) and the distance to the stellar center (stellar radius rs in pixels). The longitude \u03b8 varies between \u2212 90\u00b0 and 90\u00b0, with the center of the stellar disk corresponding to a value of 0\u00b0. The co-latitude \u03c6 ranges from 0\u00b0 to 180\u00b0, with the stellar equator set at 90\u00b0. Figure 3 illustrates the projection of a spot on to the stellar sphere, as seen from a two-dimensional perspective. The observer is assumed to lie far along the z-axis. To determine the pixels on the stellar sphere, which correspond to the starspot, we implement the following boundary condition: If the angle \u0394\u03c3 between the pixel on the sphere and the spotcenter is greater than the angular radius of the spot \u03b1, then this pixel does not belong to the spot. The values for \u0394\u03c3 are derived by using the spherical law of cosines: (2)\\begin{equation} \\cos(\\Delta\\sigma) = \\cos(\\phi_{\\rm{spot}}) \\cdot \\cos(\\phi) + \\sin(\\phi_{\\rm{spot}}) \\cdot \\sin(\\phi) \\cdot \\cos(\\Delta\\theta),  \\label{eq:slaw} \\end{equation}cos(\u0394\u03c3)=cos(\u03c6spot)\u00b7cos(\u03c6)+sin(\u03c6spot)\u00b7sin(\u03c6)\u00b7cos(\u0394\u03b8),where \u03c6spot and \u03c6 are the co-latitudes of the spotcenter and the surrounding pixels, respectively. The value \u0394\u03b8 represents the absolute difference in longitude between the spotcenter and the pixel center. ","Citation Text":["Tregloan-Reed et al. 2015"],"Functions Text":["However, to obtain precise photometric transit and spot parameters, we recommend the use of a planetary pixel radius between 15 and 50 pixels"],"Functions Label":["Uses"],"Citation Start End":[[599,624]],"Functions Start End":[[456,597]]}
{"Identifier":"2022MNRAS.509.5340B__Masci_et_al._2019_Instance_1","Paragraph":"We run forced photometry at these SN locations using a pipeline, hereafter known as the zuds pipeline1 (Dhawan et al. 2021), which performs aperture photometry using the astropy affiliated package PhotUtils (Bradley et al. 2019), using a 6-pixel diameter aperture on the difference images. The reference images for the difference images are constructed by co-adding exposures from epochs at least 30 d or more before the initial estimate of the time of maximum from the alert photometry, using the software swarp (Bertin 2010). In order to build the co-add, we only take epochs with seeing between 1.7 and 3 arcsec and a magnitude limit deeper than 19.2 mag. For consistency, we use the same reference image for both SNe. In the zuds pipeline, difference images are obtained using hotpants (Becker 2015), an implementation of the image subtraction algorithm (Alard & Lupton 1998). The zero points for each epoch are computed by the Infrared Processing and Analysis Center (IPAC), corrected for a 6-pixel diameter aperture. For the i band, we use the images corrected for an observed fringing pattern, using the fringez software (Medford et al. 2021). From the IPAC forced-photometry service (Masci et al. 2019) at the same locations, we obtain the metadata for each observation, including the magnitude limit mlim of the observation, the seeing of the observation, and the standard deviation \u03c3pix on the background at the pixel on which the SN is located. We combine this information with the zuds pipeline results for data quality assessment. Specifically, we only use those observations that satisfy the following conditions: $1.0 \\lt \\mathrm{seeing} \\lt 4.0\\,\\mathrm{ arcsec}, \\ m_{\\mathrm{lim}} \\lt 19.2 \\mathrm{\\, mag}, \\ \\mathrm{ and} \\ \\sigma _{\\mathrm{pix}} \\lt 14.0$, where \u03c3pix is the robust \u03c3 per pixel in the science image and is used as a metric to remove non-photometric data. We then use a maximum likelihood method to fit the salt2 model to each of these two SN light curves. The low seeing values are removed to protect against undersampling during image subtraction. We then remove the epochs that have 5\u03c3 flux outliers relative to the best-fitting salt2 model (discussed later and summarized in Table 2) and use the remaining selected points as the light curves of the individual SNe. The final photometry data sets used are included as Tables C1 and C2. The resulting light curves are shown in Fig. 3. As stated earlier, this light curve clearly shows the presence of two transients with no detections for a period of over 100 d in between. For the unclassified transient AT 2019lcj, we notice that AT 2019lcj has a shoulder in the redder bands (r- and i band) as seen in Fig. 3 strongly indicating that the SN is of Type Ia. We will verify this shortly after discussing its redshifts and host properties.","Citation Text":["Masci et al. 2019"],"Functions Text":["From the IPAC forced-photometry service","at the same locations, we obtain the metadata for each observation, including the magnitude limit mlim of the observation, the seeing of the observation, and the standard deviation \u03c3pix on the background at the pixel on which the SN is located."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1192,1209]],"Functions Start End":[[1151,1190],[1211,1455]]}
{"Identifier":"2020AandA...634L...8K__Ramos_et_al._2018_Instance_1","Paragraph":"The second data release (DR2) of the ESA Gaia mission has provided the largest available 6D phase-space (positions and velocities) dataset for 7.2 million stars brighter than GRVS\u2004=\u200412 mag (Gaia Collaboration et al. 2018b), making possible precise studies of the Milky Way structure and kinematics on large scales. Gaia data have already revealed signatures of non-equilibrium and ongoing vertical phase mixing in the Milky Way disc (Antoja et al. 2018), likely induced by a previous pericentric passage of the Sagittarius dwarf galaxy (Laporte et al. 2019; Bland-Hawthorn et al. 2019; Haines et al. 2019) or by internally driven bending waves (Khoperskov et al. 2019; Darling & Widrow 2019). A large number of kinematic arches with various morphologies not known prior to Gaia were found (Ramos et al. 2018; Kushniruk & Bensby 2019; Monari et al. 2019) and large-scale wiggles were discovered in the V\u03d5\u2005\u2212\u2005R plane (Kawata et al. 2018; Antoja et al. 2018). These features can be interpreted as the signature of the impact of an external perturbation (D\u2019Onghia et al. 2016; Laporte et al. 2019) or as evidence of the spiral structure and the bar of the Milky Way (Hunt et al. 2018; Fragkoudi et al. 2019). At the same time, in the solar vicinity the Galactic spiral structure is believed to generate various patterns in the phase-space (Siebert et al. 2012; Williams et al. 2013), but uncertainties on the location and strength of the spiral arms have so far prevented us from explaining the observed patterns. In this work, we use the high-quality Gaia DR2 sample of stars with radial velocities (hereafter GRV2), together with a new method to highlight the stellar density structures, to explore the Milky Way spiral arms without relying on any specific stellar tracer. The outline of the Letter is as follows: in Sect. 2 we describe the data and the method we used; in Sect. 3 we examine the phase-space structure of the Milky Way and compare its various features with Milky Way-type galaxy simulations; and in Sect. 4 we summarize the conclusions that can be drawn from our study.","Citation Text":["Ramos et al. 2018"],"Functions Text":["A large number of kinematic arches with various morphologies not known prior to Gaia were found"],"Functions Label":["Background"],"Citation Start End":[[790,807]],"Functions Start End":[[693,788]]}
{"Identifier":"2016AandA...588A.132T__Padilla_et_al._2014_Instance_1","Paragraph":"Most of the recent observational efforts to understand galaxy evolution have been focused on determining the history of cosmic star formation, gas density evolution, metallicity evolution, and mass growth of the Universe (Daddi et al. 2004; Mannucci et al. 2010; Madau & Dickinson 2014; Tomczak et al. 2014; Bouwens et al. 2015). These multiwavelength observational constraints have usually been summarized as galaxy scaling relations that might or might not change with redshift (Mannucci et al. 2010; Elbaz et al. 2011; Bouwens et al. 2014; Troncoso et al. 2014), in high- or low-density environments, in extreme physical conditions (starburst, AGN galaxies), and in spatially resolved data due to internal variations of the galaxy properties (Sanchez et al. 2013). In parallel, theoretical works and simulations have tried to explain the physical mechanisms that reproduce the measured global properties  (Daddi et al. 2010; Dav\u00e9 et al. 2011; Lilly et al. 2013; Lagos et al. 2014; Padilla et al. 2014). Despite these efforts, the completeness and cleanness of the sample are still challenging problems that depend on the sample selection-method, instrument limits, and telescope time. These problems make the comparison between observational and theoretical works even more difficult. For example, Campbell et al. (2014) compared the stellar mass of GALFORM galaxies predicted by the model with those obtained through the fit of their predicted broad-band colors. They found that both quantities differ for an individual galaxy, hence the clustering of mass-selected samples can be affected by systematic biases. Therefore, mass-selected samples might provide erroneous conclusions regarding their progenitors and descendants. In addition, the evolution of scaling relations is constrained with observations of galaxy samples that are selected with luminosity or stellar-mass thresholds and are located at different redshifts, which does not necessarily constitute causally connected populations (i.e., they do not follow a progenitor-to-descendant relation). Clustering-selected samples overcome this problem because in a hierarchical clustering scenario, a correlation analysis allows us to estimate the bias and hence statistically determine the progenitors and descendants of galaxy samples. The bias parameter measures the clustering difference between the galaxy spatial distribution and underlying dark-matter distribution. Thus, it relates the typical mass of halos hosting the galaxies (Sheth et al. 2001). Hence, measuring it in galaxy samples at different redshifts determines whether we are following the evolution of baryonic processes occurring in halos of similar masses or not. This fact is of extreme importance because once it is determined, we can use the multiwavelength data to study the evolution of the baryonic processes at certain halo mass, establishing a direct link between observations and galaxy formation models. Padilla et al. (2011) selected early-type galaxies according to their clustering and luminosity function in the MUSYC survey. So far, no study that selects star-forming galaxies according their clustering and luminosity function has been reported. ","Citation Text":["Padilla et al. 2014"],"Functions Text":["In parallel, theoretical works and simulations have tried to explain the physical mechanisms that reproduce the measured global properties"],"Functions Label":["Background"],"Citation Start End":[[984,1003]],"Functions Start End":[[768,906]]}
{"Identifier":"2016ApJ...821..107G__Gloeckler_&_Fisk_2015_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":["Gloeckler & Fisk 2015"],"Functions Text":["In 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"],"Functions Label":["Uses"],"Citation Start End":[[1029,1050]],"Functions Start End":[[700,1004]]}
{"Identifier":"2019AandA...630A..37S__Behar_et_al._2017_Instance_1","Paragraph":"Solar wind velocity distribution moments are described in Behar et al. (2017). The ion density nsw is the moment of order 0, and the ion bulk velocity usw (a vector) appears in the moment of order 1, the flux density \n\n$n_{\\mathrm{sw}} \\ \\underline{\\mathbf{u}_{\\mathrm{sw}}}$\n\n\n\nn\n\nsw\n\n\u2009\n\n\nu\n\nsw\n\n\n_\n\n\n\n. The bulk speed can be defined as the norm of the bulk velocity, that is, \n\n$u_{\\mathrm{sw}} = |\\underline{\\vec{u}_{\\mathrm{sw}}}$\n\n\n\nu\n\nsw\n\n=|\n\n\nu\n\nsw\n\n\n_\n\n\n\n|. However, this bulk speed is representative of single-particle speeds as long as the velocity distribution function is compact (e.g., a Maxwellian distribution). Complex velocity distribution functions were observed by RPC-ICA within the atmosphere of 67P. For instance, partial ring distributions were frequently observed for solar wind protons at intermediate heliocentric distances, when the spacecraft approached the SWIC (Behar et al. 2017). To illustrate the effect of such distorted distributions, a perfect ring (or shell) distribution centered on the origin of the plasma reference frame can be imagined, in which all particles have the same speed of 400 km s\u22121. The norm of the bulk velocity in this case would be 0 km s\u22121, whereas the mean speed of the particles is 400 km s\u22121, which is the relevant speed for SWCX processes. This mean speed, noted Usw, of the particles is calculated by first summing the differential number flux over all angles, and then taking the statistical average (Behar 2018). Over the entire mission, the deceleration of the solar wind using the mean speed of the particles is much more limited than the deceleration shown by the norm of the bulk velocity (Behar et al. 2017): there is more kinetic energy in the solar wind than the bulk velocity vector would let us think. This is the main difference with the paradigm used at previously studied (and more active) comets (Behar et al. 2018b). These complex, nonthermal velocity distribution functions also prevent us from reducing the second-order moment (the stress tensor) to a single scalar value, which, for a Maxwellian distribution, could be identified with a plasma temperature. In the context of 67P and for an important part of the cometary orbit around the Sun, the temperature of the solar wind proton has no formal definition.","Citation Text":["Behar et al. (2017)"],"Functions Text":["Solar wind velocity distribution moments are described in"],"Functions Label":["Background"],"Citation Start End":[[58,77]],"Functions Start End":[[0,57]]}
{"Identifier":"2016MNRAS.459.3998L__Blanton_&_Roweis_2007_Instance_1","Paragraph":"In this subsection, we first use our measured CLFs to infer the conditional stellar mass functions and then use the results to study the stellar mass contents of dark matter haloes. To convert luminosity into stellar mass, one typically uses a mass-to-light relation based on galaxy colour (e.g. Bell et al. 2003). This requires robust colour estimates. In our case, galaxies with r \u223c 21 in the SDSS photometric sample have typical error in the (u \u2212 r) colour of about 1 mag, mostly due to uncertainty in the u-band photometry. This error will propagate into the stellar mass estimates and can bias the stellar mass function, leading to an overestimate at the high-mass end.7 To reduce such bias, we estimate stellar masses using the observed mean colour\u2013magnitude relations for blue and red galaxies separately. The details of this procedure are described in Appendix C. As our final goal is to estimate the global baryon fractions in stars, the use of average values as opposed to full colour distributions is not a severe limitation. Following Bell et al. (2003), we convert the observed luminosity and colour into stellar mass using\n\n(11)\n\n\\begin{eqnarray}\n\\log \\bigg [\\frac{M_{{\\ast }}}{\\mathrm{M}_{\\odot }}\\bigg ] = -0.223+0.299\\,(u-r) -0.4\\,(M_{\\rm r}-4.64) -0.1,\\nonumber\\\\\n\\end{eqnarray}\n\nwhere (u \u2212 r) is the mean colour of a blue or red galaxy at a given absolute magnitude Mr. The constant, 4.64, is the r-band magnitude of the Sun in the AB system (Blanton & Roweis 2007) and the \u22120.1 offset corresponds to the choice of the Kroupa initial mass function (Kroupa 2001). Using this light-to-mass relation, we convert the global best-fitting luminosity functions into the corresponding stellar mass functions. The two left-hand panels in Fig. 10 show the estimated conditional stellar mass functions for blue and red satellites as a function of halo mass, respectively. Since a fixed $M^{{\\ast }}_{{\\rm b}}$ is applied to satellite galaxies for all halo masses, a slight overestimate of the stellar mass occurs at the massive ends for small haloes. As the stellar masses of central galaxies are obtained using individual observed (u \u2212 r) colours, the overestimate of the stellar mass of satellites can sometimes cause the stellar mass of a satellite galaxy to exceed that of the central. The dashed lines in the left-hand panels indicate the ranges where such situation is present. In order to estimate the total stellar mass in haloes of a given halo mass, we integrate the inferred conditional stellar mass functions down to low masses. We find that using 107\u2009M\u2299, which is about the minimum stellar mass reachable by the sample used here or zero lead to similar results. The results are shown in the right-hand panel of Fig. 10. The coloured dashed lines show the stellar mass to halo mass ratios for blue and red satellites, respectively. The grey dashed line is the total stellar mass of satellite galaxies to halo mass ratio, while the grey solid line is the stellar to dark matter mass ratio of central galaxies. The total ratio is shown as the black line. For haloes with M200  1013\u2009M\u2299, the total stellar mass is dominated by the central galaxies; in contrast, for more massive haloes, it is dominated by red satellites. The contributions from red and blue satellites are comparable for haloes with M200 \u223c 1012\u2009M\u2299, and the contribution from blue satellites appears to increase towards lower halo masses. Note that, although there are marked upturns in the stellar mass functions at the low-mass ends for red galaxies, the low-mass galaxies in the upturns (M*  108\u2009M\u2299) contribute little to the total stellar mass. Our results are qualitatively consistent with estimates based on data with more limited dynamical ranges (e.g. Leauthaud et al. 2012a,b; Kravtsov, Vikhlinin & Meshscheryakov 2014).","Citation Text":["Blanton & Roweis 2007"],"Functions Text":["The constant, 4.64, is the r-band magnitude of the Sun in the AB system"],"Functions Label":["Uses"],"Citation Start End":[[1462,1483]],"Functions Start End":[[1389,1460]]}
{"Identifier":"2020MNRAS.491.5759H__Dodson-Robinson_&_Salyk_2011_Instance_1","Paragraph":"More recently, Long et al. (2018) performed an analysis on 32 discs in the Taurus star-forming region using ALMA. From this sample 12 discs containing axisymmetric structure were identified. This structure takes the form of dark band bright ring pairs, emission bumps and cavities. Overall 19 gap ring pairs were identified, indicating that a number of these systems contain multiple gaps. These gaps range in location from R = 10\u2013120 au with no preferred distance. The majority of these gaps are narrow, but the weak correlation between gap location and gap width potentially implies formation via planet\u2013disc interactions. In addition a significant number of these gaps cannot be explained by condensation fronts. Long et al. (2018) perform an analysis on the width of these gaps in order to determine possible masses of planets that could form them. Assuming the width of the gaps corresponds to 4RHill (Dodson-Robinson & Salyk 2011) they estimate the masses of the planets to be in the 0.1\u20130.5MJ (q = 1.0 \u00d7 10\u22124\u20135.0 \u00d7 10\u22124) range, however they stress that these masses have large uncertainties and that the gap widths may be as large as 7\u201310RHill (Pinilla, Benisty & Birnstiel 2012). Alternatively they compare the distance between the minimum of the gaps and the maximum of the rings to the results of hydrodynamic simulations (Rosotti et al. 2016) with an \u03b1 = 10\u22124 (\u03bd = 2.5 \u00d7 10\u22127), resulting in a predicted planet mass of 0.05MJ (q = 5.0 \u00d7 10\u22125) and an \u03b1 = 10\u22122 (\u03bd = 2.5 \u00d7 10\u22125), resulting in a predicted planet mass of 0.3MJ (q = 3.0 \u00d7 10\u22124). Using equation (7) and an estimated disc lifetime of 2\u2009Myr [based on stellar age estimates for the spectral type of the target stars (Baraffe et al. 2015; Feiden 2016)] we can see that in the higher viscosity case planets of 0.3MJ (q = 3.0 \u00d7 10\u22124) can form at any of these radii without creating vortices. The lower viscosity case is more difficult to predict as it is significantly lower than the viscosities we have explored, however for this planet mass it seems unlikely from our results that the predicted planet could form without creating vortices. Instead, we expect the maximum mass planets that could be found here to be roughly half the mass they predict. In addition, from our results planets of mass 0.5MJ (q = 5.0 \u00d7 10\u22124) would require the disc viscosity to be \u03bd \u2a86 3.0\u20135.0 \u00d7 10\u22126 to form without creating vortices, depending on the radius at which they are located. Conversely we find planets as small as 0.1MJ (q = 1.0 \u00d7 10\u22124) can even form at \u03bd = 2.0 \u00d7 10\u22126 without creating vortices. This can be seen in Fig. 8. Hence our results show that these gaps could be opened by planets of the masses predicted by Long et al. (2018) without forming vortices.","Citation Text":["Dodson-Robinson & Salyk 2011"],"Functions Text":["Assuming the width of the gaps corresponds to 4RHill"],"Functions Label":["Uses"],"Citation Start End":[[907,935]],"Functions Start End":[[853,905]]}
{"Identifier":"2017MNRAS.464.4534Q__Schmidt_et_al._2012_Instance_1","Paragraph":"Space missions have traditionally focused on performing spectropolarimetric observations measuring the four Stokes parameters in a narrow spectral window where one or two photospheric absorption lines of interest are present. For instance, see Hinode\/SP (Tsuneta et al. 2008; Lites et al. 2013), SDO\/HMI (Pesnell, Thompson & Chamberlin 2012; Schou et al. 2012), and Solar Orbiter\/PHI (Gandorfer et al. 2011; Solanki et al. 2015). On the contrary, ground-based telescopes usually have instruments which can cover several spectral lines simultaneously, e.g. THEMIS\/MTR (L\u00f3pez Ariste, Rayrole & Semel 2000; Paletou & Molodij 2001), SST\/CRISP (Scharmer et al. 2003, 2008), DST\/IBIS (Cavallini 2006) and DST\/SPINOR (Socas-Navarro et al. 2006), Gregor\/GRIS (Collados et al. 2012; Schmidt et al. 2012), or ZIMPOL (Povel 2001), among others. In most of the mentioned cases, the light beam occupies almost the full length of the camera in one of its directions due to the larger spectral coverage which directly increases the amount of data generated. Although the data rate is not usually a crucial factor for ground-based telescopes, space-based missions have an extremely limited telemetry. Thus, unless there is a strong reason to expand the wavelength coverage it is highly recommendable to keep focusing on narrow spectral windows which contain a high density of useful lines. In this regard, we find it extremely helpful to perform theoretical studies and observations to define the optimum spectral window for the purposes of a given mission. For example, Hinode\/SP was strongly supported by ASP (Elmore et al. 1992) observations and SDO\/HMI benefited from ground-based observations but also from specific supporting works like Norton et al. (2006). We performed a similar study in Quintero Noda et al. (2016) aiming to support future missions with chromospheric polarimetry as the main target, for instance Solar-C (Katsukawa & Solar-C Working Groups 2011; Katsukawa et al. 2012; Watanabe 2014; Suematsu & Solar-C Working Group 2016). We concluded that the Ca\u2009ii 8542\u2009\u00c5 is a unique spectral line which is sensitive to a large range of heights, from the photosphere to the chromosphere (see figs 4 and 6 of the mentioned work). However, its sensitivity to the magnetic field at photospheric layers is low which precludes examining quiet Sun magnetic fields at these heights. Therefore, there is still room for improvement and, for this reason, we decided to examine the solar spectrum at the vicinity of the Ca\u2009ii 8542\u2009\u00c5 line. In this regard, if we observe additional spectral lines we will increase the number of spectral points with valuable information, particularly if these lines have different heights of formation, which will enhance the accuracy of the inferred atmospheric parameters (for example Asensio Ramos et al. 2007). However, in order to achieve this purpose we have to expand the spectral window which will increase the data rate if we maintain the same spectral sampling. Therefore, it is essential that these additional lines provide valuable information complementing the Ca\u2009ii 8542\u2009\u00c5 spectral line.","Citation Text":["Schmidt et al. 2012"],"Functions Text":["On the contrary, ground-based telescopes usually have instruments which can cover several spectral lines simultaneously, e.g.","Gregor\/GRIS"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[774,793]],"Functions Start End":[[430,555],[739,750]]}
{"Identifier":"2019ApJ...875..129K__Smith_et_al._2013_Instance_1","Paragraph":"We used the \n\n\n\n\n\n band high-resolution spectra of Arcturus and \u03bc Leo, obtained with WINERED, to estimate the microturbulence and iron abundance with a precision similar to that of previous results from spectra at different wavelengths. Our lists of Fe i lines in the 0.91\u20131.33 \u03bcm range will be useful for obtaining the precise metallicities of stars obscured by severe interstellar extinction compared with the optical regime, for which the extinction is stronger. For many objects in the Galactic disk found in recent infrared surveys, this new wavelength window may be ideal for detailed abundance analyses. One of the major error sources is the uncertainty in \u03be in various studies, including ours, based on spectra at different wavelengths from the optical (e.g., Table 3 of Jofr\u00e9 et al. 2014) to the H band (e.g., Table 7 of Smith et al. 2013). Furthermore, how to determine the microturbulence and its error is not established or straightforward. The bootstrap method that we demonstrated in this paper can give quantitative estimates of the microturbulence and its error. The error in microturbulence is 0.11\u20130.24 \n\n\n\n\n\n for each combination of target and line list. The obtained microturbulences are consistent with those that were estimated or assumed in previous studies on the same targets. Note, however, that using different line lists (or different sets of lines) can result in slightly different microturbulences depending especially on the \n\n\n\n\n\n values of strong lines used in the analysis. The very strong lines (X > \u22126) were rejected because these lines are likely to introduce problems into a chemical abundance analysis due to severe saturation, non-LTE effects, contribution of EW from the damping wing, and so on. Considering the comparison of our estimates with previous ones in addition to the scatters of \n\n\n\n\n\n, we adopt the measurements with the Fe i lines selected from MB99 as our best estimates: \n\n\n\n\n\n and \n\n\n\n\n\n for Arcturus and \u03bc Leo, respectively.","Citation Text":["Smith et al. 2013"],"Functions Text":["One of the major error sources is the uncertainty in \u03be in various studies, including ours, based on spectra at different wavelengths","to the H band"],"Functions Label":["Uses","Uses"],"Citation Start End":[[830,847]],"Functions Start End":[[611,743],[798,811]]}
{"Identifier":"2017AandA...605A..96D__Go\u017adziewski_et_al._2016_Instance_1","Paragraph":"In this article, I present an analytical model of resonant chains. Analytical models have already been proposed, in particular to study the dynamics of the Laplace resonance (1:2:4 chain) between Io, Europa, and Ganymede (e.g., Henrard 1983). However, while several numerical studies have been dedicated to the capture of planets in various resonant chains (e.g., Cresswell & Nelson 2006; Papaloizou & Terquem 2010; Libert & Tsiganis 2011; Papaloizou 2016), general analytical models have not yet been proposed. Recently, Papaloizou (2015) proposed a semi-analytical model of three-planet resonances taking into account only the interactions between consecutive planets in the chain, with a particular focus on the Kepler-60 system (12:15:20 resonant chain, see also Steffen et al. 2013; Go\u017adziewski et al. 2016). This model is very similar to the studies of the Laplace resonance between the Galilean moons, but is not well suited in the general case. For instance, four-planet (or more) resonances are not considered. Moreover, for some three-planet resonances, the interactions between non-consecutive planets cannot be neglected. For instance, in a 3:4:6 resonant chain, each planet is locked in a first-order resonance with each of the other planets. In particular, the innermost and outermost planets are involved in a 2\/1 MMR that strongly influences the dynamics of the system. I describe here a general model of resonant chains, with any number of planets, valid for any resonance order. I particularly focus on finding the equilibrium configurations (eccentricities, resonant arguments, etc.) around which a resonant system should librate. While a real system may be observed with significant amplitude of libration around the equilibrium, or could even have some angles circulating, the position of the equilibria still provides useful insights into the dynamics of the system. In Sect. 2, I describe this analytical model, and the method I use to find the equilibrium configurations. In Sect. 3, I apply the model to Kepler-223. I show that six equilibrium configurations exist for this resonant chain, and that the system is observed to be librating around one of them. I also show that knowing the current configuration of the system allows for interesting constraints to be put on its migration scenario, and in particular on the order in which the planets have been captured in the chain. ","Citation Text":["Go\u017adziewski et al. 2016"],"Functions Text":["Recently, Papaloizou (2015) proposed a semi-analytical model of three-planet resonances taking into account only the interactions between consecutive planets in the chain, with a particular focus on the Kepler-60 system (12:15:20 resonant chain, see also"],"Functions Label":["Background"],"Citation Start End":[[788,811]],"Functions Start End":[[512,766]]}
{"Identifier":"2021AandA...654A..34B__Grassi_et_al._2014_Instance_1","Paragraph":"The simulations presented in this work have been performed with the publicly available hydrodynamic code\u202fGIZMO (Hopkins 2015), which is a descendant of\u202fGADGET2 (Springel 2005). The code evolves the magneto-hydrodynamics equations for the gas including a constrained-gradient divergence-cleaning method (Hopkins & Raives 2016; Hopkins 2016), together with the gas self-gravity. For the purpose of this study, we equipped the code with an on-the-fly non-equilibrium chemical network, which was implemented via the public chemistry library\u202fKROME (Grassi et al. 2014), similarly to Bovino et al. (2019). In particular, in our simulations, we assumed an isothermal equation of state for the gas, with the temperature set to 10 K or 15 K. These temperatures are in line with kinetic temperatures obtained from NH3 for the same regions (Friesen et al. 2009, 2017). The initial conditions of our simulations consist of a collapsing filament, modelled as a cylinder with a typical observed (Arzoumanian et al. 2011) Plummer-like density profile \n\n$n(R) = n_0\/[1+(R\/R_{\\textrm{flat}}){}^2]^{p\/2}$\n\n\nn(R)=\nn\n0\n\n\/\n\n[1+\n\n(R\/\nR\n\nflat\n\n)\n2\n\n]\n\np\/2\n\n\n\n, where R is the cylindrical radius, n0 is the ridge volume density and is constant along the filament axis, Rflat is the characteristic radius of the flat inner part of the density profile, and p is the characteristic exponent. Following observed estimates (Arzoumanian et al. 2011), we set p = 2 and Rflat = 0.033 pc, which gives a mean filament width of 3 \u00d7 Rflat ~ 0.1 pc. The setup follows previous works (Seifried & Walch 2015; K\u00f6rtgen et al. 2018). To avoid any spurious effect at the edges of the filament along its axis, we embed the cylinder in an exponentially decaying background, with the decay scale length set to the filament length Lfil = 1.6 pc. The box is large Lbox = 2.4 pc. We initialised the filament in a turbulent state, assuming a Burgers-like power spectrum that grows as \u221d k10 up to \u03bbpeak = Lbox\u22156 and then decays as \u221d k\u22122 (see also K\u00f6rtgen et al. 2018). Finally, we assumed the box is permeated by a constant magnetic field B0 = 40 \u03bcG (Seifried & Walch 2015; K\u00f6rtgen et al. 2018).","Citation Text":["Grassi et al. 2014"],"Functions Text":["For the purpose of this study, we equipped the code with an on-the-fly non-equilibrium chemical network, which was implemented via the public chemistry library\u202fKROME"],"Functions Label":["Uses"],"Citation Start End":[[544,562]],"Functions Start End":[[377,542]]}
{"Identifier":"2017ApJ...836..124D__Davenport_et_al._2015_Instance_1","Paragraph":"In Table 3, we list the physical parameters of our model (not including the normalizations for each telescope and each eclipse), and in Tables 4 and 5 we show the resulting best-fit model and the 16th and 84th percentiles (approximately 1\u03c3) for each parameter for each season. We are able to fit our measured light curves equally well regardless of which star we place star spots. However, if we place the star-spot signal on the secondary component, the out-of-eclipse model parameters vary significantly between observing seasons, whereas if the star-spot signal is originating on the primary component, the parameters are stable between seasons. Particularly, F, the ratio between the orbital and rotational frequency has different values for each season if the secondary component is responsible for the star-spot modulation signal. Since the rotation period of the secondary star should not change (and the effect of differential rotation for star spots at different latitudes is small for M dwarfs; Davenport et al. 2015), we conclude that the star-spot signal cannot be originating on the secondary component. Since we only detect one rotational frequency, and this frequency differs from the orbital frequency of the system, we have assumed that it is originating solely on one component of the system (the primary component); though, in reality, there is likely to be some star spots on the surface of both components and that rotation of the secondary component may also be contributing somewhat to this signal (though we do not have a significant detection). We will utilize the model for which star spots are located on the primary component for the rest of this paper. We note, however, that this choice does not significantly affect the values of the masses and radii of the components but does change the uncertainty. The light curves themselves are most directly sensitive to the sum of the component radii and their ratio. The eclipse duration measures a combination of (\n\n\n\n\n\n)\/a (where R1 and R2 are the component radii and a is the semimajor axis) and the inclination of the orbit. In the case of grazing eclipse (where we cannot break this degeneracy with a measurement of the duration of the total phase), this degeneracy is instead broken by the eclipse depth, which also depends on the limb darkening of the stars and the star spots on the stellar surfaces.","Citation Text":["Davenport et al. 2015"],"Functions Text":["Since the rotation period of the secondary star should not change (and the effect of differential rotation for star spots at different latitudes is small for M dwarfs","we conclude that the star-spot signal cannot be originating on the secondary component."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1005,1026]],"Functions Start End":[[837,1003],[1029,1116]]}
{"Identifier":"2019MNRAS.486.1781R__Fossati_et_al._1998_Instance_1","Paragraph":"Blazars are a peculiar class of active galactic nuclei (AGNs) that have their relativistic jets pointed close to the line of sight to the observer with angle \u2264 10\u00b0 (Antonucci 1993; Urry & Padovani 1995). They are classified as flat-spectrum radio quasars (FSRQs) and BL Lacerate (BL Lac) objects based on the strength of the emission lines in their optical\/infrared (IR) spectrum. Both classes of objects emit radiation over the entire accessible electromagnetic spectrum from low-energy radio to high-energy \u03b3-rays. As blazars are aligned close to the observer, the emission is highly Doppler boosted, causing them to appear as bright sources in the extragalactic sky. They dominate the extragalactic \u03b3-ray sky first hinted by the Energetic Gammma-ray Experiment Telescope (EGRET) observations onboard the Compton Gamma-Ray Observatory (CGRO; Hartman et al. 1999) and now made apparent by the Large Area Telescope (LAT) onboard the Fermi Gamma-ray Space Telescope (Atwood et al. 2009). The broad-band spectra of blazars are dominated by emission from the jet with weak or absent emission lines from the broad-line region (BLR). One of the defining characteristics of blazars is that they show flux variations (Wagner & Witzel 1995) over a wide range of wavelengths on time-scales ranging from months to days and minutes. In addition to flux variations they also show large optical and radio polarization as well as optical polarization variability. In the radio band they have flat spectra with the radio spectral index (\u03b1r)  0.5 ($S_{\\nu } \\propto \\nu ^{-\\alpha _r})$. The broad-band spectral energy distribution (SED) of blazars is characterized by a two-hump structure, one peaking at low energies in the optical\/IR\/X-ray region and the other one peaking at high energies in the X-ray\/MeV region (Fossati et al. 1998; Mao et al. 2016). In the one-zone leptonic emission models, the low-energy hump is due to synchrotron emission processes and the high-energy hump is due to inverse Compton (IC) emission processes (Abdo et al. 2010b). The seed photons for the IC process can be either internal to the jet (synchrotron self-Compton or SSC; Konigl 1981; Marscher & Gear 1985; Ghisellini & Maraschi 1989) or external to the jet (external Compton or EC; Begelman et al. 1987). In the case of EC, the seed photons can be from the disc (Dermer & Schlickeiser 1993; Boettcher, Mause & Schlickeiser 1997), the BLR (Sikora, Begelman & Rees 1994; Ghisellini & Madau 1996), and the torus (B\u0142aerrorzdotejowski et al. 2000; Ghisellini & Tavecchio 2008). Though leptonic models are found to fit the observed SED of a majority of blazars, for some blazars, their SEDs are also well fitted by either hadronic (M\u00fccke et al. 2003; B\u00f6ttcher et al. 2013) or lepto-hadronic models (Diltz & B\u00f6ttcher 2016; Paliya et al. 2016). In the hadronic scenario, the \u03b3-ray emission is due to synchrotron radiation from extremely relativisitic protons (M\u00fccke et al. 2003) or the cascade process resulting from proton\u2013proton or proton\u2013photon interactions (Mannheim 1993). Even during different brightness\/flaring states of a source, a single emission model is not able to fit the broad-band SED at all times. For example in the source 3C 279, while the flare during 2014 March\u2013April is well fitted by a leptonic model (Paliya, Sahayanathan & Stalin 2015b), the flare in 2013 December with a hard \u03b3-ray spectrum is well described by lepto-hadronic processes (Paliya et al. 2016). Thus, the recent availability of multiwavelength data coupled with studies of sources at different active states indicates that we still do not have a clear understanding of the physical processes happening close to the central regions of blazars.","Citation Text":["Fossati et al. 1998"],"Functions Text":["The broad-band spectral energy distribution (SED) of blazars is characterized by a two-hump structure, one peaking at low energies in the optical\/IR\/X-ray region and the other one peaking at high energies in the X-ray\/MeV region"],"Functions Label":["Background"],"Citation Start End":[[1801,1820]],"Functions Start End":[[1571,1799]]}
{"Identifier":"2018MNRAS.475.3419A__Davis_et_al._1999_Instance_1","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"],"Functions Text":["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","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."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1531,1548]],"Functions Start End":[[1365,1529],[1551,1815]]}
{"Identifier":"2019AandA...627A.173V__Epstein_et_al._(2014)_Instance_1","Paragraph":"Since the age determination using asteroseismology relies on the mass-age relation that red giants follow, this means that the method is biased by any event that changes the stellar mass, such as mass accretion from a companion or stellar mergers (blue stragglers or stars rejuvenated by accretion, e.g. Boffin et al. 2015) or mass loss. One way to look for mass accretion events from a companion is to look for radial velocity, photometric variations, or chemical signs of accretion (e.g. high carbon and s-process enhancements; Beers & Christlieb 2005; Abate et al. 2015). The effect of mass loss can be minimised by looking at stars in the low-RGB phase, in which the effect of mass loss is smaller compared to red clump stars (Anders et al. 2017). The first study to determine masses for a sample of metal-poor halo giants with both seismic information (from Kepler; Borucki et al. 2010) and chemistry from high-resolution APOGEE (Majewski et al. 2017) spectra, was by Epstein et al. (2014). The authors used scaling relations at face value and reported masses larger (M\u2004> \u20041 M\u2299) than what would be expected for a typical old population. Similar results were obtained by Casey et al. (2018), also using scaling relations for three metal-poor stars. These findings led to the need for further tests of the use of asteroseismology in the low metallicity regime. Miglio et al. (2016) analysed a group of red giants in the globular cluster M4 ([Fe\/H] = \u22121.10 dex and [\u03b1\/Fe] = 0.4 dex) with seismic data from K2 mission (Howell et al. 2014). These authors found low seismic masses compatible with the old age of the cluster, hence suggesting that seismic masses and radii estimates would be reliable in the metal-poor regime provided that a correction to the \u0394\u03bd scaling relation is taken into account for red giant branch (hereafter RGB) stars. The correction presented in Miglio et al. (2016) is a correction that is theoretically motivated based on the computation of radial mode frequencies of stellar modes.","Citation Text":["Epstein et al. (2014)"],"Functions Text":["The first study to determine masses for a sample of metal-poor halo giants with both seismic information","and chemistry from high-resolution APOGEE","spectra, was by","The authors used scaling relations at face value and reported masses larger (M\u2004> \u20041 M\u2299) than what would be expected for a typical old population.","Similar results were obtained by Casey et al. (2018), also using scaling relations for three metal-poor stars.","These findings led to the need for further tests of the use of asteroseismology in the low metallicity regime."],"Functions Label":["Background","Background","Background","Background","Similarities","Background"],"Citation Start End":[[973,994]],"Functions Start End":[[752,856],[892,933],[957,972],[996,1141],[1142,1252],[1253,1363]]}
{"Identifier":"2016AandA...589A..73R__Kurucz_1992_Instance_1","Paragraph":"Single-burst stellar population (SSP) models mimic uniform stellar populations of fixed age and metallicity, and are an important tool to study unresolved stellar clusters and galaxies. They are created by populating theoretical stellar evolutionary tracks with stars of a stellar library, according to a prescription given by a chosen initial mass function (IMF). Thus, the quality of the resulting SSP models depends significantly on the completeness of the used input stellar library in terms of evolutionary phases represented by the atmospheric parameters temperature, Teff, surface gravity, log\u2009(g), and metallicity. A sufficiently large spectral coverage is equally crucial when constructing reasonable SSP models. Theoretical stellar libraries like, e.g. BaSeL (Kurucz 1992; Lejeune et al. 1997, 1998; Westera et al. 2002), or PHOENIX (Allard et al. 2012; Husser et al. 2013) are generally available for both a large range in wavelength and in stellar parameters, whereas empirical libraries are found to be more incomplete in both respects. However, the advantage of the latter ones is that they are not hampered by the still large uncertainties in the calculation of model atmospheres. Examples of empirical stellar libraries in the optical wavelenth range encompass the Pickles library (Pickles 1998), ELODIE (Prugniel & Soubiran 2001), STELIB (Le Borgne et al. 2003), Indo-US (Valdes et al. 2004), MILES (S\u00e1nchez-Bl\u00e1zquez et al. 2006), and CaT (Cenarro et al. 2001, 2007). In the near-infrared (NIR) and mid-infrared (MIR)1, only very few empirical libraries have been observed so far (e.g. Lan\u00e7on & Wood 2000; Cushing et al. 2005; Rayner et al. 2009). The NASA Infrared Telescope Facility (IRTF) spectral library, described in the latter two papers, is to date the only empirical stellar library in the NIR and MIR which offers a sufficiently complete coverage of the stellar atmospheric parameter space to construct SSP models. In the future, the X-Shooter stellar library, which contains around 700 stars, and which covers the whole optical (see Chen et al. 2014) and NIR wavelength range until 2.5 \u03bcm, will clearly improve the current situation in the NIR. ","Citation Text":["Kurucz 1992"],"Functions Text":["Theoretical stellar libraries like, e.g. BaSeL","are generally available for both a large range in wavelength and in stellar parameters, whereas empirical libraries are found to be more incomplete in both respects."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[770,781]],"Functions Start End":[[722,768],[884,1049]]}
{"Identifier":"2020MNRAS.492..821C__Masters_et_al._2014_Instance_1","Paragraph":"As a general remark, whether the locally calibrated metallicity diagnostics are applicable to high-redshift galaxies is still a matter of great debate. Diagnostics that are expected to be little affected by the ionization conditions of the gas (see e.g. Dopita et al. 2016) have been suggested to be valuable for high-redshift galaxies, where strong variations in ionization parameter and excitation conditions compared to local galaxies have been invoked to explain the observed evolution in the emission-line ratios (as seen, for example, from the offset of high-z sources in the classical BPT diagrams with respect to the local sequence (Kewley et al. 2013; Nakajima et al. 2013; Steidel et al. 2014; Kashino et al. 2017; Strom et al. 2017). However, since such diagnostics usually involve the [N\u2009ii]\/[O\u2009ii] or the [N\u2009ii]\/[S\u2009ii] line ratios, they are strongly dependent on the assumed relation between the N\/O ratio as a function of the oxygen abundance O\/H, which is affected by a large scatter and whose evolution with cosmic time and\/or dependence on galaxy mass is also indicated as a possible origin of the observed evolution of the emission-line properties in high-zgalaxies (Masters et al. 2014, 2016; Shapley et al. 2015). Therefore, strong-line indicators based only on alpha elements (like, e.g. oxygen) have also been suggested as appropriate to high redshift studies, since galaxies at z \u223c 1.5\u20132.5 seem to show no appreciable offset from local trends in oxygen-based diagnostic diagrams (e.g. R23 versus O32, Shapley et al. 2015). However, the location on the abovementioned diagram could even be sensitive to a variation in the hardening of the radiation field at a fixed metallicity rather than a variation in abundances (Steidel et al. 2016; Strom et al. 2017). In any case, at redshifts \u223c1.5 (where the majority of KLEVER galaxies considered in this work lie), the lack of the [O\u2009ii] doublet in the NIR bands observable from KMOS prevents us from using purely oxygen diagnostics, thus forcing us to exploit the nitrogen-based ones. When the survey will be complete, we will investigate the spatially resolved behaviour of z \u223c 2 galaxies on the R23 versus O32 diagram in a more statistically robust manner. Recently, Patr\u00edcio et al. (2018) have shown that oxygen-based diagnostics z \u223c 2 provide metallicities comparable to those inferred from the electron temperature method; unfortunately, just a handful of robust auroral line detections have been reported so far in high-z sources (e.g. Jones, Martin & Cooper 2015b; Sanders et al. 2016; see also Patr\u00edcio et al. 2018, and references therein), due to the intrinsic observational challenges in detecting the faint auroral lines with current instrumentation. Only the advent of new facilities like JWST or the MOONS spectrograph on the VLT will ultimately allow us to tackle this issue in the next few years, allowing us to properly calibrate the metallicity diagnostics against fully Te-based abundance determination at high redshifts and providing the key to overcome all these potential discrepancies.","Citation Text":["Masters et al. 2014"],"Functions Text":["However, since such diagnostics usually involve the [N\u2009ii]\/[O\u2009ii] or the [N\u2009ii]\/[S\u2009ii] line ratios, they are strongly dependent on the assumed relation between the N\/O ratio as a function of the oxygen abundance O\/H, which is affected by a large scatter and whose evolution with cosmic time and\/or dependence on galaxy mass is also indicated as a possible origin of the observed evolution of the emission-line properties in high-zgalaxies"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1185,1204]],"Functions Start End":[[745,1183]]}
{"Identifier":"2020AandA...637A..44N__Kerszberg_et_al._2017_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":["Kerszberg et al. (2017)"],"Functions Text":["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 (2017) over broken power-law models derived from the fits to lower energy data."],"Functions Label":["Uses","Uses"],"Citation Start End":[[634,657]],"Functions Start End":[[486,633],[657,748]]}
{"Identifier":"2019ApJ...887...26O__Garon_et_al._2019_Instance_1","Paragraph":"In an effort to understand the character of RG encounters with ICM shocks and the observable consequences, we have initiated an MHD simulation study of such interactions. The simulations include passive, but energy-dependent transport of CRe, so that resulting radio synchrotron properties can be modeled appropriately. Nolting et al. (2019a, 2019b) explored interactions between ICM-strength shocks and lobed RGs. Here, we extend that work to include such shock interactions with so-called \u201cnarrow-angle tail\u201d (NAT) RGs. NATs, which are especially common in clusters, (e.g., Blanton et al. 2001; Garon et al. 2019) result from sustained relative motion between an RG and its ambient ICM; that is, the RG effectively evolves in a wind that deflects the RG jets and their discharge downwind. The NATs can extend multiple hundreds of kpc from their AGN sources, and, so long as the AGN jets remain \u201cpowered on,\u201d simulations suggest that NAT tails remain dynamic and highly inhomogeneous (O\u2019Neill et al. 2019). Their considerable lengths and elongated morphologies make them interesting candidates to supply CRe at ICM shocks to form radio relics. We refer readers to O\u2019Neill et al. (2019) and references therein for details of NAT formation dynamics. We do provide in Section 2 an outline of salient features of the simulated NAT presented in O\u2019Neill et al. (2019) that are especially relevant to the present work. Our focus in this paper is on the interactions of shocks with such a previously formed NAT. Specifically, we use new three-dimensional (3D) MHD simulations to see what happens when plane shocks moving transverse to the NAT major axis collide with the NAT. That geometry is both relatively simple, and, on the face of it, perhaps the best able to form a radio relic in such a manner. Our study includes cases with shock Mach numbers, \n\n\n\n\n\n, in order to span the range of shock strengths most often mentioned in this context. As in the aforementioned studies, we carry out synthetic, frequency-dependent radio synchrotron observations of our simulated objects in order to better understand the relationships between their dynamical and observable properties.","Citation Text":["Garon et al. 2019"],"Functions Text":["NATs, which are especially common in clusters, (e.g.,","result from sustained relative motion between an RG and its ambient ICM; that is, the RG effectively evolves in a wind that deflects the RG jets and their discharge downwind."],"Functions Label":["Background","Background"],"Citation Start End":[[597,614]],"Functions Start End":[[522,575],[616,790]]}
{"Identifier":"2016ApJ...826..106G__Mazzarella_et_al._2012_Instance_1","Paragraph":"The galaxy merging process can cause enhanced accretion onto the central SMBHs and thus initiate activity. One or both of the dual SMBHs may be active or re-activated. Simulations (e.g., Van Wassenhove et al. 2012) suggest that simultaneous activity is mostly expected at the late phases of mergers, at or below 10 kpc scale separations. Additionally, the recent work of Fu et al. (2015b) indicates that mergers trigger and tend to synchronize activity at separations of a few kiloparsecs. Therefore dual AGNs are expected to be observed, but they are not easily resolvable with current observing facilities. The number of convincing dual AGN systems, mostly detected by X-ray and radio observations (e.g., Beswick et al. 2001; Junkkarinen et al. 2001; Komossa et al. 2003; Hudson 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; M\u00fcller-S\u00e1nchez et al. 2015), are relatively few compared to the theoretically predicted abundance (see also, Komossa & Zensus 2016). While very long baseline interferometric (VLBI) radio observations currently provide the highest spatial resolutions, making them a promising tool for the detection of dual AGNs6\n\n6\nIn fact, thanks to its superior resolution, the first binary AGN was discovered using VLBI technique (Rodriguez et al. 2006).\n, only about 10% of AGNs are radio-loud, therefore dual radio-emitting active nuclei would be quite rare. Moreover, when trying to confirm the existence of dual AGNs in candidate sources, non-detection in the radio bands does not immediately falsify the dual AGN hypothesis, because radio-quiet nuclei may also be present. Burke-Spolaor (2011) searched the archival VLBI observations of 3114 radio-emitting AGNs for sources containing compact double sources. Only one source (B3 0402+379) was found with a double nucleus with separation of about 7 pc, which was already known from the observations of Rodriguez et al. (2006).","Citation Text":["Mazzarella et al. 2012"],"Functions Text":["The number of convincing dual AGN systems, mostly detected by X-ray and radio observations (e.g.,","are relatively few compared to the theoretically predicted abundance"],"Functions Label":["Differences","Differences"],"Citation Start End":[[856,878]],"Functions Start End":[[609,706],[949,1017]]}
{"Identifier":"2021MNRAS.503.2973C__Carvalho_et_al._2020_Instance_1","Paragraph":"As a first approximation, we consider that most part of the accretion luminosity is emitted in the soft X-ray spectrum (10 eV\u201310 keV), thus, soft X-ray luminosity is simply estimated by using the mass accretion parameter, $\\dot{m}$(1)$$\\begin{eqnarray*}\r\nL_{\\rm AC}= \\frac{GM\\dot{m}}{R},\r\n\\end{eqnarray*}$$where M and R are the mass and radius of the object, respectively, which can be either a white dwarf, neutron star, or black hole. Mass\u2013radius relation is employed to estimate radius in case of white dwarfs and neutron stars (see e.g. Barstow 1991, 1993; Lopes & Menezes 2012; Berti 2013; Carvalho, Marinho & Malheiro 2018; de Santi & Santarelli 2019; Carvalho et al. 2020), and Schwarzschild radius is taken for black holes. We reinforce that the model above is only a way to estimate the X-ray flux produced by compact objects due to accretion. The accretion rate, $\\dot{m}$, is the parameter which defines the luminosity for a given system. However, in this work, we will not study any particular object, instead we will take fiducial values for mass and radius of the compact object along with standard values for snow line distances to estimate the time-scales ices need to reach chemical equilibrium. So, the model does not define if the compact object is isolated or in a binary system. We assume that accretion rate, $\\dot{m}$, is constant on time, which means that the systems we are studying are well behaved, not presenting any sudden infall of matter, bursts, substantial outflows, or strong winds. The main aim, as stated, is to obtain an estimate of the time-scales for ices to reach chemical equilibrium. The assumption of constant accretion rate is feasible in most cases because the time-scale for ices close to compact systems to reach chemical equilibrium is smaller than the time-scale for changes on mass, as one can notice from power-law parametrizations (Chen et al. 2019), where we see that $\\dot{m}$ scales with Myr or Gyr, while the time-scales for chemical equilibrium are at most 0.1 Myr for distances smaller than 100 LY, as we are going to show.","Citation Text":["Carvalho et al. 2020"],"Functions Text":["Mass\u2013radius relation is employed to estimate radius in case of white dwarfs and neutron stars (see e.g."],"Functions Label":["Uses"],"Citation Start End":[[658,678]],"Functions Start End":[[437,540]]}
{"Identifier":"2019MNRAS.484.3785B__Huchra_&_Sargent_1973_Instance_1","Paragraph":"Once we have the completeness corrections, we can estimate the relative luminosity function of SNe Ia, which we show in Fig. 3. In the left panel we show the observed distribution of absolute magnitudes. The black histogram shows the full sample, and the red histogram shows the volume-limited sample; the dot--dashed blue histogram shows the distribution of the volume-limited LOSS sample (Li et al. 2011b). The volume-limited samples have a higher fraction of low-luminosity SNe Ia than the full magnitude-limited sample. We convert the ASAS-SN distributions into estimates of the true luminosity function using the V\/Vmax method (Schmidt 1968; Huchra & Sargent 1973; Felten 1976). For each SN, we compute the maximum volume (Vmax) in which the SN could be recovered by a survey with a limiting magnitude mV = 16.8. This is an empirical limiting magnitude; this value produces a median value of V\/Vmax close to 0.5, which is to be expected if sources uniformly populate the survey volume. We compute the relative luminosity function for each bin in absolute magnitude centred on M as \n(2)\r\n\\begin{eqnarray*}\r\n\\Phi (M) = \\sum _{i=1}^{N} \\frac{1}{V_{M,i}} \\times w_i \\times (1+z_i), \r\n\\end{eqnarray*}\r\nwhere the sum is over all the SNe within the bin. The weights $w$i correct for the incompleteness given the apparent peak magnitude of each SN, and the factor of (1 + $z$) accounts for time dilation. The results are shown in the right panel of Fig. 3. The black circles show the relative luminosity function computed from the full sample, the red squares show the results for the volume-limited sample, and the blue crosses show the control-time weighted counts from Li et al. (2011b). The luminosity functions are normalized to the bin at MV = \u221219. In this paper we are not aiming for an absolute rate calibration. The shape of the relative luminosity function is consistent with the volume-limited luminosity function presented in Li et al. (2011b). We fit a Schechter (1976) function \n(3)\r\n\\begin{eqnarray*}\r\n\\phi (L) \\propto \\left(\\frac{L}{L_*}\\right)^{\\alpha } \\exp \\left(-\\frac{L}{L_*}\\right) \r\n\\end{eqnarray*}\r\nto the relative luminosity function of both the full and volume-limited samples, where \u03b1 is the faint-end slope, and L* (alternatively M* in magnitude space) determines the \u2018knee\u2019 of luminosity function. Our fits are shown in Fig. 3 as dashed lines. We find (\u03b1, M*) corresponding to (1.3 \u00b1 0.4, \u221218.1 \u00b1 0.1) and (2.1 \u00b1 0.3, \u221217.8 \u00b1 0.1) for the full sample and the volume-limited sample, respectively.","Citation Text":["Huchra & Sargent 1973"],"Functions Text":["We convert the ASAS-SN distributions into estimates of the true luminosity function using the V\/Vmax method"],"Functions Label":["Uses"],"Citation Start End":[[647,668]],"Functions Start End":[[524,631]]}
{"Identifier":"2016AandA...592A..19C__Davies_et_al._1993_Instance_1","Paragraph":"The downsizing scenario is evident in several cases of galaxy evolution (Fontanot et al. 2009). In the case of ETGs at z ~ 0, one of the first pieces of observational evidence can be referred to the studies of Dressler et al. (1987), Faber et al. (1992) and Worthey et al. (1992), who found more massive elliptical galaxies to be more enriched in \u03b1-elements than less massive ones. These works suggested selective mass-losses, different initial mass functions (IMF) and\/or different star-formation timescales as possible explanations for the high level of [\u03b1\/Fe]. Subsequent studies found the same trend of [\u03b1\/Fe] with mass (Carollo et al. e.g 1993; Davies et al. 1993; Bender et al. 1993; Thomas et al. 2005; 2010; McDermid et al. 2015), leading to the dominant interpretation that in more massive ETGs, the duration of star formation was substantially shorter than in less massive ones, with timescales short enough (e.g. 0.5 Gyr) to avoid the dilution of the \u03b1 element abundance (produced by Type II supernovae) by the onset of Fe production by Type Ia supernovae. This is considered to be one of the main pieces of evidence for the shortening of star formation. The age of the ETG stellar populations at z ~ 0 also show evidence of downsizing, with more massive objects being older than less massive ones. These results have been derived both in clusters (Thomas et al. 2005, 2010; Nelan et al. 2005) and in the field (Heavens et al. 2004; Jimenez et al. 2007; Panter et al. 2007; see also Renzini 2006). Most of these studies are based on fitting individual spectral features with the Lick\/IDS index approach (Burstein et al. 1984; Worthey et al. 1994) which allows us to mitigate the problem of the age-metallicity degeneracy (Graves & Schiavon 2008; Thomas et al. 2005, 2010; Johansson et al. 2012a; Worthey et al. 2013). However, more recently, other approaches based on the full-spectrum fitting have been developed (e.g. STARLIGHT, Cid Fernandes et al. 2005; VESPA, Tojeiro et al. 2009, 2013, and FIREFLY, Wilkinson et al. 2015), and applied to samples of ETGs at z ~ 0  (Jimenez et al. 2007; Conroy et al. 2014; McDermid et al. 2015). The results based on Lick indices are in general quite consistent with those of full spectral fitting within 10\u221230% (e.g. Conroy et al. 2014) and support the downsizing evolutionary pattern. ","Citation Text":["Davies et al. 1993"],"Functions Text":["Subsequent studies found the same trend of [\u03b1\/Fe] with mass","leading to the dominant interpretation that in more massive ETGs, the duration of star formation was substantially shorter than in less massive ones, with timescales short enough (e.g. 0.5 Gyr) to avoid the dilution of the \u03b1 element abundance (produced by Type II supernovae) by the onset of Fe production by Type Ia supernovae. This is considered to be one of the main pieces of evidence for the shortening of star formation."],"Functions Label":["Background","Background"],"Citation Start End":[[650,668]],"Functions Start End":[[564,623],[739,1165]]}
{"Identifier":"2017AandA...604A.112G__Hansen_2010_Instance_1","Paragraph":"In stars, there are two components to describe the tidal interaction, equilibrium tides and dynamical tides. Equilibrium tides correspond to a large-scale hydrostatic adjustment of a body and the resulting flow due to the gravitational field of a given companion (Zahn 1966; Remus et al. 2012). It is usually employed in the framework of the constant time lag model (see Mignard 1979; Hut 1981; Eggleton et al. 1998; Bolmont et al. 2011, 2012), which allows a fast computation of the orbital evolution of the planet and works for all eccentricities (Hut 1981; Leconte et al. 2010). In this model, the dissipation of the kinetic energy of the equilibrium tide inside the star is often taken to be constant throughout the system evolution and calibrated on observations (Hansen 2010, 2012). While considering such a constant equilibrium tide dissipation is a sensible assumption, several studies have shown that this quantity might vary during the different phases of stellar evolution. For example, Zahn & Bouchet (1989) showed that the dissipation of the equilibrium tide by the turbulent friction in the convective envelope of late-type stars is strongest during their PMS. Using this theoretical framework, Villaver & Livio (2009, see also Verbunt & Phinney 1995 recalled that the variation of the semi-major axis of a planet induced by such friction can be expressed as a function of the ratio of the mass of the convective envelope to the total mass of the star, the ratio between the radius of the star and the orbital semi-major axis (to the power 8), and finally of a power of the ratio between the tidal period and the convective turnover timescale. This allows the loss of efficiency of tidal friction to be modelled for rapid tides (e.g. Zahn 1966; Goldreich & Keeley 1977). Because of the variations of these quantities during post-MS phases (e.g. Charbonnel et al. 2017), this could lead to a more efficient dissipation than during the MS. Finally, Mathis et al. (2016) demonstrated that the action of rotation on convection deeply modifies the turbulent friction it applies on the equilibrium tide. In the regime of fast rotation, which corresponds to the end of PMS and early MS phase, the friction is several orders of magnitude lower than in a model ignoring rotation. This may lead to a loss of efficiency of the dissipation of the equilibrium tide. This shows that care should be taken when assuming a calibrated constant dissipation of the equilibrium tide during the evolution of stars. ","Citation Text":["Hansen 2010"],"Functions Text":["In this model, the dissipation of the kinetic energy of the equilibrium tide inside the star is often taken to be constant throughout the system evolution and calibrated on observations"],"Functions Label":["Background"],"Citation Start End":[[769,780]],"Functions Start End":[[582,767]]}
{"Identifier":"2016ApJ...827..104S__Risaliti_&_Lusso_2015_Instance_1","Paragraph":"Using these quasar survey data, Clowes et al. (2013), Nadathur (2013), Einasto et al. (2014), and Park et al. (2015) found very large quasar groups and discussed the cosmological implications of the existence and properties of these extreme objects. Nadathur (2013), Einasto et al. (2014), and Park et al. (2015) pointed out that the cosmological interpretation by Clowes et al. (2013) of large quasar groups, which questions the validity of the cosmological assumption of homogeneity and isotropy, is misleading. They stressed the importance of a statistically precise analysis in order to draw conclusions on cosmological implications from the observation of one or a few extreme objects. Park et al. (2015) also emphasized that statistical comparison with cosmological simulations must be employed as well. The quasar survey data are also used for studies to constrain cosmological parameters (Han & Park 2015; Risaliti & Lusso 2015). Besides the studies directly related to cosmology, there are more studies exploring the clustering properties of quasars. Einasto et al. (2014) made a catalog of quasar groups with different linking lengths and examined their properties. They found that the characteristics of quasar groups such as number density, size, and richness, identified with linking lengths varying from 20 to 40 h\u22121 Mpc, are well correlated with those of galaxy superclusters. Therefore such quasar groups can be markers of galaxy superclusters. As a classical way to study clustering properties, correlation functions have been measured for quasars (AGNs in general) by many different groups (Shen et al. 2007b, 2009, 2013; Ross et al. 2009; Krumpe et al. 2010, 2012, 2015; Miyaji et al. 2011; Cappelluti et al. 2012; Richardson et al. 2012; Allevato et al. 2014; Eftekharzadeh et al. 2015). They have measured two-point cross-correlation functions (2PCCFs) between quasars and galaxies, and have found the typical mass of quasar-hosting dark matter halos (DMHs) and the dependence of the mass on quasar luminosity.","Citation Text":["Risaliti & Lusso 2015"],"Functions Text":["The quasar survey data are also used for studies to constrain cosmological parameters"],"Functions Label":["Background"],"Citation Start End":[[914,935]],"Functions Start End":[[810,895]]}
{"Identifier":"2018MNRAS.481..533L__Ili\u0107,_Kova\u010devi\u0107_&_Popovi\u0107_2009_Instance_1","Paragraph":"We refer to models for which the pressure P depends on the radial distance r from the central continuum source, \n(2)\r\n\\begin{eqnarray*}\r\nP(r)\\propto r^{-s} \\, , \r\n\\end{eqnarray*}\r\nas pressure-law models. In this work, we examine two limiting cases, s = 0 and s = 2, representing constant density and constant ionization parameter models, respectively. We here adopt a spherically symmetric BLR geometry spanning more than two decades in radial extent. This model is chosen for (i) its simplicity, and because (ii) we can compare our radial pressure-law models with the Local Optimally emitting Cloud model for this source presented by Korista & Goad (2001), which also adopts spherical symmetry. Here, we summarize the radial dependencies of various physical quantities for spherically symmetric pressure-law models; the derivations in this section follow Rees et al. (1989) and Goad et al. (1993). We make the simplifying assumption that the cloud temperature does not vary with radius; for solar composition, photoionization equilibrium is achieved at temperatures T \u223c 104 K across a wide range of ionization parameter (equation 5), and the gas temperature will therefore vary weakly with radius (e.g. Netzer 1990; Ili\u0107, Kova\u010devi\u0107 & Popovi\u0107 2009). For constant cloud temperatures, the cloud hydrogen gas density nH\u2009is proportional to the pressure, P, and so \n(3)\r\n\\begin{eqnarray*}\r\nn_\\mathrm{H}(r)\\propto r^{-s} \\, . \r\n\\end{eqnarray*}\r\nThus, s = 0 corresponds to a constant nH\u2009throughout the BLR. The ionization parameter U is defined as \n(4)\r\n\\begin{eqnarray*}\r\nU(r)=\\frac{Q_\\mathrm{H}}{n_\\mathrm{H}(r)4\\pi r^2c} \\, , \r\n\\end{eqnarray*}\r\nwhere QH is the number of hydrogen-ionizing photons emitted by the central continuum source per second. Thus, s = 2 corresponds to a constant ionization parameter model, since: \n(5)\r\n\\begin{eqnarray*}\r\nU(r)\\propto r^{s-2} \\, . \r\n\\end{eqnarray*}\r\nThe surface area per cloud, Ac, is proportional to $R_{\\rm c}^2$, where Rc denotes the radius of a cloud. In general, Rc depends on the pressure P, and is therefore constant for s = 0. If we demand that the mass of each cloud is conserved as the clouds move radially outwards (i.e. clouds do not break up or coalesce within our region of interest), mass conservation implies that $R_c^3n_H={\\rm constant}$. Thus, we obtain the relation \n(6)\r\n\\begin{eqnarray*}\r\nA_{\\rm c}(r) \\propto R_{\\rm c}^2(r) \\propto r^{2s\/3} \\, . \r\n\\end{eqnarray*}\r\nThe column density of each cloud, Ncol, depends on the gas density and cloud radius: \n(7)\r\n\\begin{eqnarray*}\r\nN_\\mathrm{col}(r)\\propto R_cn_\\mathrm{H}\\propto r^{-2s\/3} \\, . \r\n\\end{eqnarray*}\r\nThe above relations determine the local physical conditions, as parametrized by Ncol, nH\u2009and incident ionizing photon flux \u03a6H, at any radius in a spherically symmetric pressure-law BLR model. These conditions determine the local surface emissivity \u03b5(r) of a BLR cloud at radius r (as determined via cloudy modeling; Section 3.1). The total luminosity for an emission line is then found by integrating over the distribution in cloud properties such that \n(8)\r\n\\begin{eqnarray*}\r\nL_{\\mathrm{line}}=4\\pi \\int _{r_{\\mathrm{in}}}^{r_{\\mathrm{out}}}\\epsilon (r)A_{\\rm c}(r)n_{\\rm c}(r)r^2{\\rm d}r \\, , \r\n\\end{eqnarray*}\r\nwhere rin \u2009and rout\u2009 are the inner and outer BLR radii, respectively, Ac is the surface area of a single cloud, and nc is the local number density of clouds. As dust grains strongly absorb UV photons, routis chosen to approximately coincide with the distance at which dust grains can form and survive (\u2248140 light-days for NGC 5548).","Citation Text":["Ili\u0107, Kova\u010devi\u0107 & Popovi\u0107 2009"],"Functions Text":["We make the simplifying assumption that the cloud temperature does not vary with radius; for solar composition, photoionization equilibrium is achieved at temperatures T \u223c 104 K across a wide range of ionization parameter (equation 5), and the gas temperature will therefore vary weakly with radius (e.g."],"Functions Label":["Uses"],"Citation Start End":[[1217,1247]],"Functions Start End":[[899,1203]]}
{"Identifier":"2018AandA...614A..66S__Jiang_et_al._2008_Instance_1","Paragraph":"Virtually all formulas that are currently used to estimate the merger time, \u03c4mer, are based on the idealised Chandraseckhar (1943) description of the deceleration caused by dynamical friction on a point mass (representing the satellite) travelling through an infinite, uniform, and collisionless medium (representing the host). They are mostly prescriptions inferred from the analytic or semi-analytic modelling of mergers (e.g. Lacey & Cole 1993; van den Bosch et al. 1999; Taffoni et al. 2003; Gan et al. 2010; Petts et al. 2015; Silva et al. 2016) or parametric equations tuned by direct comparison with the outcome of numerical simulations of collisions of live galaxy pairs (e.g. Boylan-Kolchin et al. 2008; Jiang et al. 2008; Just et al. 2011; McCavana et al. 2012; Villalobos et al. 2013), or both (e.g. Colpi et al. 1999). However, although several decades of studies of galactic mergers have allowed reaching a general consensus on what probably are the most relevant parameters in any process of this type, the extent of their impact is still a matter of debate. Factors such as continued mass losses due to tidal interactions, the drag force exerted by tidal debris, the re-accretion of some of this material onto the colliding galaxies, or the mutual tidal distortion of their internal structures are elements that introduce nontrivial uncertainties in any attempt to calculate \u03c4mer using analytical expressions. To these difficulties, which are caused by physical complexity, one must add the lack of a 100% standard methodology regarding the way mergers are tracked. The starting point of this work is precisely the definition of a new metric for calculating \u03c4mer, which is an obligatory first step to compare the outcomes of different experiments under equal conditions and derive universally valid formulas, in a form that is suitable for major mergers. With this aim, we have run a suite of nearly 600 high-resolution N-body simulations of isolated major mergers, with which we further explored the dependence of the merger time on a range of orbital parameters and on the mass-ratios, spins, and morphologies of the progenitors that are representative of such systems.","Citation Text":["Jiang et al. 2008"],"Functions Text":["or parametric equations tuned by direct comparison with the outcome of numerical simulations of collisions of live galaxy pairs (e.g."],"Functions Label":["Background"],"Citation Start End":[[713,730]],"Functions Start End":[[551,684]]}
{"Identifier":"2018MNRAS.477L..80K__Colella_&_Woodward_1984_Instance_1","Paragraph":"The numerical schemes in this work are essentially the same as those in Kuroda et al. (2016). Regarding the metric evolution, we evolve the standard BSSN variables (Shibata & Nakamura 1995; Baumgarte & Shapiro 1999; Marronetti et al. 2008) with a finite-difference scheme in space and with a Runge\u2013Kutta method in time, both in fourth-order accuracy. The gauge is specified by the \u20181+log\u2019 lapse and by the Gamma-driver-shift condition. Regarding the radiation-hydrodynamic evolution, the conservation equation $\\Delta_\\alpha T^{\\alpha \\beta }_{\\rm (total)}=0$ is solved using the piecewise parabolic method (Colella & Woodward 1984; Hawke, L\u00f6ffler & Nerozzi 2005). $T^{\\alpha \\beta }_{\\rm (total)}$ is the total stress-energy tensor,\n(1)\\begin{equation*}T_{\\rm (total)}^{\\alpha \\beta } = T_{\\rm (fluid)}^{\\alpha \\beta } +\\int \\mathrm{ d}\\varepsilon \\sum _{\\nu \\in \\nu _e, \\bar{\\nu }_e, \\nu _x}T_{(\\nu , \\varepsilon )}^{\\alpha \\beta }, \\end{equation*}\r\nwhere $T_{\\rm (fluid)}^{\\alpha \\beta }$ and $T_{(\\nu , \\varepsilon )}^{\\alpha \\beta }$ are the stress-energy tensor of the fluid and the neutrino radiation field, respectively. We consider three-flavour of neutrinos ($\\nu \\in \\nu _e, \\bar{\\nu }_e, \\nu _x$) with \u03bdx denoting heavy-lepton neutrinos (i.e. \u03bd\u03bc, \u03bd\u03c4, and their antiparticles). \u03f5 represents the neutrino energy measured in the comoving frame which logarithmically covers from 1 to 300 MeV with 12 energy bins. Employing an M1 analytical closure scheme (Shibata et al. 2011), we solve spectral neutrino transport of the radiation energy and momentum, based on the truncated moment formalism (e.g. Kuroda et al. 2016; Roberts et al. 2016; Ott et al. 2018). We include the gravitational red- and Doppler-shift terms to follow the neutrino radiation field in highly curved space\u2013time around BH. Regarding neutrino opacities, the standard weak interaction set in Bruenn (1985) plus nucleon\u2013nucleon bremsstrahlung (Hannestad & Raffelt 1998) is taken into account.","Citation Text":["Colella & Woodward 1984"],"Functions Text":["Regarding the radiation-hydrodynamic evolution, the conservation equation $\\Delta_\\alpha T^{\\alpha \\beta }_{\\rm (total)}=0$ is solved using the piecewise parabolic method"],"Functions Label":["Uses"],"Citation Start End":[[608,631]],"Functions Start End":[[436,606]]}
{"Identifier":"2016MNRAS.457..212S__Singh,_Sami_&_Dadhich_2003_Instance_1","Paragraph":"One of the first and simplest proposed Friedmann\u2013Robertson\u2013Walker (FRW) cosmological model is the \u039b cold dark matter (\u039bCDM) universe, which involves Einstein's cosmological constant \u039b. This standard model of cosmology, which is also referred to as the concordance model, assumes that the total energy density \u03c1 of the universe is made up of three components, namely matter \u03c1m (baryonic and dark matter), radiation \u03c1r, and dark energy or vacuum energy \u03c1\u039b, which produces the necessary gravitational repulsion. In this model, dark energy which has an equation of state (EOS) \u03c9\u039b = p\u039b\/\u03c1\u039b = \u22121, is a property of the space itself and its density \u03c1\u039b = \u2212p\u039b = \u039bc4\/8\u03c0G is constant, such that as the universe expands the constant vacuum energy density will eventually exceed the matter density of the universe which is ever decreasing. The spatially flat \u039bCDM model dominated by vacuum energy with \u03a9\u039b \u223c 0.70, with the rest of the energy density being in the form of non-relativistic cold dark matter with \u03a9m \u223c 0.25 and non-relativistic baryonic matter with \u03a9b \u223c 0.05, fits observational data reasonably well (Riess et al. 1998; Permutter et al. 1999; Knop et al. 2003; Riess et al. 2004). However, the main problem in this model is the huge difference of about 10120 orders of magnitude between the observed value of the cosmological constant and the one predicted from quantum field theory; known as the cosmological constant problem (Weinberg 1989). Another issue is the so called coincidence problem which expresses the fact that although in this model the matter and dark energy components scale differently with redshift during the evolution of the universe, both components today have comparable energy densities, and it is unclear why we happen to live in this narrow window of time. Besides these main issues, there are other inherent problems faced by the \u039bCDM, some of which arose as a result of recent observations that are in disagreement with the model's predictions. For example, in order to account for the general isotropy of the cosmic microwave background (CMB), the standard model invokes an early period of inflationary expansion (Kazanas 1980; Guth 1981; Linde 1982). However, the latest observations by Planck (Planck Collaboration XXIII 2003) indicate that there may be some problems with such an inflationary scenario (Ijjas, Steinhardt & Loeb 2013; Guth, Kaiser & Nomura 2014). It was partly due to these issues of the standard \u039bCDM, that during the last decade several alternative dark energy models have been proposed and tested with observations. In these models the dark energy density component \u03c1de is not constant and in most cases \u03c9de = pde\/\u03c1de depends on time, redshift, or scale factor. For example in some of these so called dynamical dark energy models, late time inflation is achieved using a variable cosmological term \u039b(t) (Ray et al. 2011; Basilakos 2015) sometimes taken in conjunction with a time dependent gravitational constant G(t) (Ray, Mukhopadhyay & Dutta Choudhury 2007; Ibotombi Singh, Bembem Devi & Surendra Singh 2013). Other sources of dark energy include scalar fields such as quintessence (Peebles & Ratra 2003), K-essence (Armendariz-Picon et al. 2001) and phantom fields (Singh, Sami & Dadhich 2003). An alternative approach to the dark energy problem relies on the modification of Einstein's theory itself such that in these alternative theories of gravity, cosmic acceleration is not provided solely by the matter side T\u03bc\u03bd of the field equations, but also by the geometry of spacetime. These theories include the scalar-tensor theory with non-minimally coupled scalar fields (Barrow & Parsons 1997; Bertolami & Martins 2000), f(R) theory (Tsujikawa 2008), conformal Weyl gravity (Mannheim 2000) and higher dimensional theories such as the Randall\u2013Sundrum (RS) braneworld model (Randall & Sundrum 1999), and the braneworld model of Dvali\u2013Gabadadze\u2013Porrati (DGP) (Dvali, Gabadadze & Porrati 2000). Over the last few years considerable interest has been shown in the simple FRW linearly expanding (coasting) model in Einstein's theory with a(t)\u2009\u221d\u2009t, H(z) = H0(1 + z). Like the \u039bCDM the total energy density and pressure in this model are expressed in terms of matter, radiation and dark energy components, such that p = \u03c9\u03c1 with \u03c1 = \u03c1m + \u03c1r + \u03c1de and p = pr + pde (since pm \u2248 0), but it includes the added assumption \u03c9 = \u22121\/3, i.e. the cosmic fluid acting as the source has zero active gravitational mass. So this would definitely exclude a cosmological constant as the source of the dark energy component in this case. The model was first discussed by Kolb (1989) who referred to this zero active mass cosmic fluid as \u2018K-matter\u2019. Interest in this model has been revived recently after it was noted (Melia 2003) that in the standard model the radius of the gravitational horizon Rh(t0) (also known as the Hubble radius) is equal to the distance ct0 that light has travelled since the big bang, with t0 being the current age of the universe. In the \u039bCDM this equality is a peculiar coincidence because it just happens at the present time t0. It was then proposed (Melia 2007, 2009; Melia & Shevchuk 2012) that this equality may not be a coincidence at all, and should be satisfied at all cosmic time t. This was done by an application of Birkhoff's theorem and its corollary, which for a flat universe allows the identification of the Hubble radius Rh with the gravitational radius Rh = 2GM\/c2, given in terms of the Misner\u2013Sharp mass $M = (4\\pi \/3)R_{{\\rm h}}^{3}(\\rho \/c^2)$ (Misner & Sharp 1964). The added assumption of a zero active gravitational mass \u03c1 + 3p = 0 implies (Melia & Shevchuk 2012) that Rh = ct or H = 1\/t for any cosmic time t. This linear model became known as the Rh = ct universe. Unlike the \u039bCDM\/\u03c9CDM which contains at least the three parameters H0, \u03a9m and \u03c9de, the Rh = ct model depends only on the sole parameter H0, so that for example the luminosity distance used to fit Type Ia supernova data (Melia 2009) is given by the simple expression dL = (1 + z)Rh(t0)ln\u2009(1 + z). Also while the \u039bCDM would need inflation to circumvent the well-known horizon problem, the Rh = ct universe does not require inflation. One should also point out that the condition Rh = ct is also satisfied by other linear models such as the Milne universe (Milne 1933), which however has been refuted by observations. Unlike the Rh = ct model discussed here, the Milne universe is empty (\u03c1 = 0) and with a negative spatial curvature (k = \u22121). As a result of these properties its luminosity distance is given by $d_{L}^{\\rm {Milne}} = R_{{\\rm h}}(t_0)(1 + z)\\sinh [\\ln (1+z)]$, and it was shown that this is not consistent with observational data (Melia & Shevchuk 2012). In the last few years the Rh = ct universe received a lot of attention when it was shown (Melia & Maier 2013; Wei, Wu & Melia 2013, 2014a, 2014b, 2015; Melia, Wei & Wu 2015) that it is actually favoured over the standard \u039bCDM (and its variant \u03c9CDM with \u03c9 \u2260 \u22121) by most observational data. This claim has been contested by Bilicki & Seikel (2012) and Shafer (2015) who argued that measurement of H(z) as a function of redshift and the analysis of Type Ia supernovae favoured the \u039bCDM over the Rh = ct universe. However, this was later contested by Melia & McClintock (2015) who showed that the Rh = ct was still favoured when using model-independent measurements that are not biased towards a specific model. Others (see for example van Oirschot, Kwan & Lewis 2010; Lewis & van Oirschot 2012; Mitra 2014) have also criticized the model itself, particularly the validity of the EOS \u03c9 = \u22121\/3 (Lewis 2013). These and other criticisms have been addressed by Bikwa, Melia & Shevchuk (2012); Melia (2012) (see also Melia 2015 and references therein.) As pointed out above the Rh = ct model would still require a dark energy component \u03c1de, albeit not in the form of a cosmological constant. So the obvious question at this point would be: what are the possible sources for this component that together with the matter and radiation components will give the required total EOS, \u03c9 = \u22121\/3? The purpose of this paper is to answer this question by discussing the various possible sources of dark energy that are consistent with this EOS. Since the radiation component \u03c1r at the present time t0 is insignificant (at least for the \u039bCDM with which this model has been compared) we assume that the total energy density \u03c1 = \u03c1de + \u03c1m and the total pressure p = pde (pm \u2248 0), as is normally done in the other alternative dynamical dark energy models found in the literature. So in the next three sections we examine three possibilities for the source of dark energy in the Rh = ct model, namely a variable cosmological term \u039b(t), a non-minimally coupled scalar field in Brans\u2013Dicke theory which is equivalent to a variable gravitational constant G(t), and finally quintessence represented by a minimally coupled scalar field \u03d5. We show that although the first two sources are consistent with the model, they are both unphysical, which leaves the third source of quintessence as the viable source of dark energy in the Rh = ct universe. Results are then discussed in the Conclusion. Unless otherwise noted we use units such that G = c = 1.","Citation Text":["Singh, Sami & Dadhich 2003"],"Functions Text":["Other sources of dark energy include scalar fields such as","and phantom fields"],"Functions Label":["Background","Background"],"Citation Start End":[[3218,3244]],"Functions Start End":[[3061,3119],[3198,3216]]}
{"Identifier":"2020ApJ...893..124Z___2017_Instance_1","Paragraph":"The terrestrial magnetosheath, downstream of the bow shock generated by the interaction between the supersonic solar wind and the Earth\u2019s magnetosphere, is representative of turbulence downstream of collisionless plasma shock in the universe. Usually, the downstream sheath regions have larger density, temperature, mean magnetic intensity, and compressibility, and higher plasma \u03b2, as compared to the upstream solar wind (Sahraoui et al. 2006; Alexandrova 2008; Hadid et al. 2015; Huang et al. 2016, 2017). Abundant unstable kinetic wave modes and multiscale coherent structures exist in the magnetosheath. For instance, mirror mode, ion cyclotron wave, kinetic Alfv\u00e9n wave, and so forth are routinely observed (Sahraoui et al. 2006; Lucek et al. 2005; Alexandrova 2008; Zhao et al. 2018; Zhang et al. 2018). Hence, the turbulent energy in the kinetic range might not only consist of energy cascading from larger scales, but also comprise direct injection of energy due to local instability at these scales. This raises an interesting question: what are the differences and similarities of turbulence intermittency between solar wind and magnetosheath? This is one goal of our research. When going down to the ion scales, PSD(\u03b4B) steepens, while PSD(\u03b4E) is relatively flattened (Chen & Boldyrev 2017; Matteini et al. 2017). The coexistence of steep PSD(B) and flat PSD(E) is caused by the ion demagnetization, which gives rise to the Hall term and\/or quasi-electrostatic field (Cramer 2001; Schekochihin et al. 2009). If monofractal scaling is a common phenomenon for magnetic fluctuations at kinetic scales both in the solar wind and in the magnetosheath, what would happen to the electric field intermittent fluctuations? is the electric field intermittency monofractal as well, or is it totally different from the magnetic one? Studying the above questions would help us better understand the characteristics and nature of turbulence at kinetic scales, which is more electrostatic than that at magnetohydrodynamic (MHD) scales.","Citation Text":["Huang et al.","2017"],"Functions Text":["Usually, the downstream sheath regions have larger density, temperature, mean magnetic intensity, and compressibility, and higher plasma \u03b2, as compared to the upstream solar wind"],"Functions Label":["Background"],"Citation Start End":[[482,494],[501,505]],"Functions Start End":[[243,421]]}
{"Identifier":"2019ApJ...870...70S__Bosch_2016_Instance_1","Paragraph":"The large differences noted in Figure 9 are surely related to the different approximations underlying both the complex hydrodynamical simulations of ILLUSTRIS-1 and our simplified approach. As recently discussed by Chua et al. (2017), the distribution of noncollisional matter (i.e., stars and DM) can effectively cool down in the cores of galaxies formed in the ancient epochs, at variance with those born more recently. This provides them a much better chance of survival, or at least stronger resilience against potentially disruptive interactions that they would surely undergo throughout the cluster environment. A key role is also played by the host cluster itself, because the mass profile of this latter becomes steeper and steeper in the core due to dynamical relaxation. Consequently, the environment gets denser with time; this trend is found to be quantifiable in terms of the number of dynamical times elapsed since the cluster formation (e.g., Jiang & van den Bosch 2016). Therefore, each galaxy has to attain a sufficiently cold core in order to first survive in the cluster environment and then grow by accretion of matter from the ICM and mergers with other substructures. Given these preliminary considerations, which explain the evaporation of galaxies while falling into more massive halos, the inclusion of the BM physics in this context is crucial, and its effects are still not fully explored. Based on a systematic comparison of the full-physics and DM-only ILLUSTRIS suites, Chua et al. (2017) claimed that galactic winds and photoionization from UV radiation may effectively inhibit mass aggregation, at least at the low-mass end and in certain regimes of redshifts (see also Despali & Vegetti 2017 for similar reasoning for DM subhalos of intermediate masses inside ILLUSTRIS and EAGLE simulations; Schaye et al. 2015). In particular, photometric data inside ILLUSTRIS are bound to whether stellar particles, i.e., stellar populations, are found inside a bound subhalo. Recalling that in ILLUSTRIS-1, the average gas cell mass is 8.86 \u00d7 105 M\u2299, from which integrated stellar populations might stem, and that a prescription for subgrid physics is always needed to follow up the integrated properties of all kinds of particles, these could be a substantial limit on state-of-the-art hydrodynamical simulations in describing photometry, especially for galaxies inside more recently formed DM halos of lower masses. On the other side, our numerical approach, at the basis of the buildup of samples inside clusters, provides a population of galaxies that should reflect the observational evidence of the mass\u2013radius relation at z \u223c 0 (Chiosi et al. 2012). If observed galaxies are merely those that survived all disrupting interactions they underwent, then at least a minimum fraction of the sample (not as close to zero as the one of ILLUSTRIS-1) should have formed in an evidently large range of look-back time. However, acknowledging that our approach might not be the most refined one, we feel that the best interpretation might be a middle ground between the two models.","Citation Text":["Jiang & van den Bosch 2016"],"Functions Text":["The large differences noted in Figure 9 are surely related to the different approximations underlying both the complex hydrodynamical simulations of ILLUSTRIS-1 and our simplified approach.","A key role is also played by the host cluster itself, because the mass profile of this latter becomes steeper and steeper in the core due to dynamical relaxation. Consequently, the environment gets denser with time; this trend is found to be quantifiable in terms of the number of dynamical times elapsed since the cluster formation (e.g.,","Therefore, each galaxy has to attain a sufficiently cold core in order to first survive in the cluster environment and then grow by accretion of matter from the ICM and mergers with other substructures."],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[958,984]],"Functions Start End":[[0,189],[618,957],[987,1189]]}
{"Identifier":"2022AandA...663L...4L___2020_Instance_2","Paragraph":"While previous studies mainly focused on the optical-UV properties of a MAD in RLAGNs, for the first time we try to investigate their X-ray properties in this work. The origin of X-ray emission in RLAGNs is still under debate, which may come from a corona, jet, or both. In observations, there is a big difference between the X-ray properties of radio-quiet AGNs (RQAGNs) and RLAGNs. Firstly, the average X-ray flux in RLAGNs is found to be 2\u20133 times higher than that in RQAGNs (e.g., Zamorani et al. 1981; Wilkes & Elvis 1987; Li & Gu 2021). Secondly, Laor et al. (1997) reported that RLAGNs have harder 2\u201310 kev X-ray spectra than RQAGNs by compiling a sample of 23 quasars observed with ROSAT, which was subsequently confirmed by Shang et al. (2011) with a larger sample. Comparing the X-ray spectrum of 3CRR quasars and that of radio-quiet quasars, Zhou & Gu (2020) also gave a similar result. In addition, the X-ray reflection features of RLAGNs are weaker than those of RQAGNs (Wozniak et al. 1998). All of these results seem to indicate that the contribution of a jet to X-ray spectra cannot be neglected. However, several recent works suggested a totally different result. First, the slope of LUV\u2005\u2212\u2005LX is found to be consistent for RLAGNs and RQAGNs (Zhu et al. 2020, 2021; Li & Gu 2021). Second, Gupta et al. (2018, 2020) discovered that the distributions of X-ray photon spectral indices between RLAGNs and their radio-quiet counterpart are very similar (see Zhu et al. 2021 either). This opposite conclusion may be due to the effect of sample selection. The sample of Gupta et al. (2018, 2020) was X-ray selected (and optically selected for the sample of Zhu et al. 2021), which may lead to the radio jet power being very feeble compared to the bolometric luminosity in most of the RLAGNs. These weakly jetted RLAGNs can therefore have different X-ray photon indices compared to the strong jetted RLAGNs, such as the 3CRR quasars of Zhou & Gu (2020). P15 also indicated that the weakly jetted RLAGNs have similar \u03b1EUV as RQAGNs. However, interestingly, Markoff et al. (2005) demonstrated that both the corona model and the jet model can fit the X-ray data of some Galactic X-ray binaries well and that the jet base may be subsumed to corona in some ways. The 3CRR quasars are low frequency radio selected and have a strong jet on a large scale. However, it is still unclear whether all the objects with a strong jet harbor a MAD, or just containing MAD when jet is firstly launching millions of years ago. We focus on the RLAGNs with an EUV deficit in this work, which should possess a MAD in the inner disk region as suggested by P15. The presence of a MAD surrounding the black hole may bring a remarkable difference to the X-ray emissions since the structure of disk-corona greatly changes in the case of MAD (e.g., Tchekhovskoy et al. 2011; McKinney et al. 2012; White et al. 2019). In theory, it has been suggested that X-ray emission increases when an advection-dominated accretion flow (ADAF) becomes a MAD in its inner region (Xie & Zdziarski 2019). Nevertheless, how MAD affects the disk-corona corresponding to the X-ray emission of quasars is still an open issue. This work can constrain a future theoretical model for MAD in RLAGNs.","Citation Text":["Gupta et al.","2020"],"Functions Text":["The sample of","was X-ray selected (and optically selected for the sample of Zhu et al. 2021), which may lead to the radio jet power being very feeble compared to the bolometric luminosity in most of the RLAGNs."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1579,1591],[1599,1603]],"Functions Start End":[[1565,1578],[1605,1800]]}
{"Identifier":"2022MNRAS.509.3599T__Du_et_al._2015_Instance_2","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"],"Functions Text":["The dimensionless accretion rate","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","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."],"Functions Label":["Uses","Uses","Uses"],"Citation Start End":[[1696,1710]],"Functions Start End":[[1471,1503],[1521,1694],[1713,1899]]}
{"Identifier":"2022ApJ...935..127H__Pan_et_al._2021_Instance_1","Paragraph":"Here we take the gravitational waveform of the nonrotating model with an observer viewed from the equatorial direction as an example for our HHT analysis. The gravitational waveforms of other models with different parameters are shown in the appendix (Figures B10\u2013B20). Figure 6 (a) shows the simulated strains, followed by the wavelet spectra (b)\u2013(c) and the stacked HHT spectra (d)\u2013(e). The waveform starts from the core collapse and reaches the core bounce at t = 0 s. As this model is nonrotating and the proto-neutron star convection happens after the core bounce, the gravitational strain shows a loud bounce signal at time zero, and then it is followed by a low-frequency peak PNS oscillation, which could be caused by the g-mode oscillation although the nature remains controversial (M\u00fcller et al. 2013; Pan et al. 2021). Multiple low-frequency oscillation signals dominate the first \u223c0.15 s, and they can be well resolved with the HHT. In comparison, the limitation of the wavelet spectrum prevents us to identify them. As the frequency of the oscillation increases with time, another high-frequency mode oscillation signal appears with similar strength as the peak PNS mode. The low-frequency PNS oscillation vanishes at t \u223c 0.5 s. Then, interestingly, it followed by the quadruple radiation of the standing accretion shock instability (SASI Blondin et al. 2003) signal, which was suggested to be seen in the GW data (Kuroda et al. 2016; Pan et al. 2021). The detailed waveform of the SASI signal is model dependent, but its time-frequency property has been shown in a few analysis either with the spectrogram or the HHT (see, e.g., Andresen et al. 2017; Mezzacappa et al. 2020; Takeda et al. 2021). In our analysis, the SASI signal can be seen on both the wavelet and the HHT spectra. The wavelet and the HHT results for other simulated CCSN events are shown in the Appendix B. The Hilbert spectra of all the data sets show much better resolution than the wavelet spectra.","Citation Text":["Pan et al. 2021","Pan et al. 2021"],"Functions Text":["As this model is nonrotating and the proto-neutron star convection happens after the core bounce, the gravitational strain shows a loud bounce signal at time zero, and then it is followed by a low-frequency peak PNS oscillation, which could be caused by the g-mode oscillation although the nature remains controversial","Then, interestingly, it followed by the quadruple radiation of the standing accretion shock instability","signal, which was suggested to be seen in the GW data"],"Functions Label":["Compare\/Contrast","Similarities","Similarities"],"Citation Start End":[[812,827],[1448,1463]],"Functions Start End":[[472,790],[1242,1345],[1373,1426]]}
{"Identifier":"2018MNRAS.475.3613S__Stabenau,_Connolly_&_Jain_2008_Instance_1","Paragraph":"With the availability of such a wide range of photo-z codes and methods, comparisons of various implementations have been performed (Hildebrandt et al. 2010; Abdalla et al. 2011; Sanchez et al. 2014). No obvious best photo-z code was named since each code displays different strengths depending on the metrics used. The focus of recent photo-z analyses has turned to improving error estimation (Oyaizu et al. 2008; Hoyle et al. 2015; Wittman, Bhaskar & Tobin 2016), the use of new statistical techniques (Lima et al. 2008; Zitlau et al. 2016), improving existing algorithms (Cavuoti et al. 2015; Sadeh et al. 2016), and the addition of extra input information to get more precise and accurate photo-z\u2019s. With regards to the inclusion of extra information, a recent example in template methods includes using surface brightness as a prior in spectral energy distribution templates (Kurtz et al. 2007; Stabenau, Connolly & Jain 2008), motivated by the knowledge of surface brightness dimming (1 + z)4. In empirical methods this application is more straightforward, since algorithms are constructed such that it is not difficult to add extra input parameters. For example, Collister & Lahav (2004) and Wadadekar (2005) demonstrated that by including the 50 and 90\u2009per cent Petrosian flux radii (RP50, RP90), the photometric redshift root-mean-square errors improve by 3 and 15\u2009per cent respectively for the SDSS main galaxy sample. Tagliaferri et al. (2003) used Petrosian fluxes and radii in their work on galaxies from the SDSS early data release and calculated robust errors decreasing as much as 24\u2009per cent. Meanwhile, Vince & Csabai (2007) included the concentration of galaxy light profiles in their study and reported that the root-mean-square error of photo-z\u2019s on SDSS galaxies improved by 3\u2009per cent. Wray & Gunn (2008) included surface brightness and the S\u00e9rsic index, and found improvements in variance when compared to other template fitting methods applied to the SDSS main galaxy sample previously.","Citation Text":["Stabenau, Connolly & Jain 2008"],"Functions Text":["With regards to the inclusion of extra information, a recent example in template methods includes using surface brightness as a prior in spectral energy distribution templates",", motivated by the knowledge of surface brightness dimming (1 + z)4.","In empirical methods this application is more straightforward, since algorithms are constructed such that it is not difficult to add extra input parameters."],"Functions Label":["Background","Background","Motivation"],"Citation Start End":[[900,930]],"Functions Start End":[[704,879],[931,999],[1000,1156]]}
{"Identifier":"2020ApJ...904...20S__Woods_et_al._2008_Instance_1","Paragraph":"Comparison of the DAXSS two temperature (2T) model fits to DEM estimates from other instruments could help validate the DAXSS simple modeling approach. Moore et al. (2018) shows that the X123 response is primarily over the temperature range of 1 MK to 10 MK, so this comparison needs to be done over a similar range, such as is accessible with using DEMs derived from extreme ultraviolet (EUV) emissions. One of the comparisons is with the DEM derived with SDO (Pesnell et al. 2012) EVE (Hock et al. 2012; Woods et al. 2012) solar EUV spectral irradiance data in the 6\u201337 nm range. The derivation of EVE-based DEM estimates is being developed for improving the X-ray ultraviolet Photometer System (XPS) data processing for the Solar Radiation and Climate Experiment (SORCE) (Woods et al. 2005a) and Thermosphere, Ionosphere, Mesosphere, Energetics, and Dynamics (Woods et al. 2005b) missions. The primary DEMs for this analysis is the quiet Sun (QS) and active region (AR) DEMs as needed for estimating the daily variations of the solar XUV spectral irradiance for the XPS Level 4 product (Woods et al. 2008). The Fe viii to Fe xvi lines in the SDO EVE spectra were initially used to estimate the DEMs for the reference QS and AR spectra derived with EVE data between 2010 and 2013 using the technique described by Schonfeld et al. (2017). It was then found that fitting with DEM Gaussian profiles with logarithmic temperature (K, log(T)) peaks every 0.2 and with Gaussian width of 0.42 in log(T) (FWHM of 1.0) provided more robust solutions for the DEM (a similar technique was described by Warren et al. 2013). Furthermore, fitting just specific Fe lines was providing low irradiance estimates in the 6\u201315 nm range. Better spectral model values for this range was found when fitting the EVE spectra over the ranges of 10\u201314 nm and 26\u201330 nm. The DEM estimate using just the rocket EVE spectral data flown with DAXSS on 2018 June 18 and the DEM based on combining the QS DEM and the AR DEM with an AR scaling factor of 0.00806 are shown in Figure 12. This scaling factor for the AR EM was determined as the best fit for the DAXSS spectral irradiance. The conversion of the EVE-based DEM (\n\n\n\n\n\n) to EM (cm\u22123) for comparison to DAXSS 2T model solution is the multiplication by the solar hemisphere area (\n\n\n\n\n\n) and by the temperature bin size (0.23 * temperature in K).","Citation Text":["Woods et al. 2008"],"Functions Text":["The primary DEMs for this analysis is the quiet Sun (QS) and active region (AR) DEMs as needed for estimating the daily variations of the solar XUV spectral irradiance for the XPS Level 4 product"],"Functions Label":["Uses"],"Citation Start End":[[1090,1107]],"Functions Start End":[[893,1088]]}
{"Identifier":"2017MNRAS.468.4992P__Lister_et_al._2016_Instance_1","Paragraph":"The intrinsic jet opening angles can be calculated as tan\u2009(\u03b1int\/2) = tan\u2009(\u03b1app\/2)\u2009sin\u2009\u03b8, where \u03b8 is the viewing angle to the jet axis. The latter, as well as the bulk Lorentz factor \u0393, can be derived from apparent jet speed and Doppler factor using the following relations:\n\r\n\\begin{eqnarray*}\r\n\\theta = \\arctan \\frac{2\\beta _{\\rm app}}{\\beta ^2_{\\rm app}+\\delta _{\\rm var}^2-1}\\,, \\quad \\Gamma = \\frac{\\beta _{\\rm app}^2+\\delta _{\\rm var}^2+1}{2\\delta _{\\rm var}}\\,.\r\n\\end{eqnarray*}\r\nFor \u03b2app and \u03b4var we used the fastest measured radial, non-accelerating apparent jet speed from the MOJAVE kinematic analysis (Lister et al. 2016) and the variability Doppler factor from the Mets\u00e4hovi AGN monitoring programme (Hovatta et al. 2009), respectively. The corresponding overlap of the programmes comprises 55 sources, which are all members of the MOJAVE-1 sample. Variability Doppler-factors for 10 more MOJAVE-1 sources were measured within the F-GAMMA programme (Liodakis et al. 2017). The intrinsic opening angles calculated for the 65 sources range from 0$_{.}^{\\circ}$1 to 9$_{.}^{\\circ}$4, with a median of 1$_{.}^{\\circ}$3, reflecting a very high degree of jet collimation. The intrinsic opening angles show an inverse dependence on Lorentz factor (Fig. 13), as predicted by hydrodynamical (Blandford & K\u00f6nigl 1979) and magnetic acceleration models (Komissarov et al. 2007) of relativistic jets. The median value of the product3 \u03c1 = \u03b1int\u0393 is 0.35 rad, close to earlier estimates derived both from observations (Jorstad et al. 2005; Pushkarev et al. 2009) and from a statistical model approach (Clausen-Brown et al. 2013). The variability Doppler factors derived from variability can be underestimated due to a limited cadence of the observations. In this case, the intrinsic opening angle estimates would be smaller, while Lorentz factors would be higher if $\\delta >(\\beta _{\\rm app}^2+1)^{1\/2}$ and smaller otherwise, implying that if the variability Doppler factors are essentially underestimated, the majority of points in Fig. 13 would move downward and to the right.","Citation Text":["Lister et al. 2016"],"Functions Text":["For \u03b2app and \u03b4var we used the fastest measured radial, non-accelerating apparent jet speed from the MOJAVE kinematic analysis"],"Functions Label":["Uses"],"Citation Start End":[[613,631]],"Functions Start End":[[486,611]]}
{"Identifier":"2021MNRAS.501...50S__Gupta_et_al._2019_Instance_1","Paragraph":"There have been rather strong claims of AGN QPOs in different bands of the electromagnetic spectrum, ranging from minutes through days through months and years (e.g. Gierli\u0144ski et al. 2008; Lachowicz et al. 2009; Gupta, Srivastava & Wiita 2009; Gupta et al. 2018, 2019; King et al. 2013; Gupta 2014, 2018; Ackermann et al. 2015; Pan et al. 2016; Zhou et al. 2018; Bhatta 2019; and references therein). However, many of the claimed QPOs, particularly those made earlier, were marginal detections (Gupta 2014), lasting only a few cycles, and the originally quoted statistical significances are probably overestimates (Gupta 2014; Covino, Sandrinelli & Treves 2019). Among the better recent claims of QPOs in the gamma-ray band are of \u223c34.5 d in the blazar PKS 2247\u2013131 (Zhou et al. 2018) and of \u223c71 d in the blazar B2 1520+31 (Gupta et al. 2019) found as part of a continuing analysis of blazar Fermi\u2013LAT observations. A recent claim of a \u223c44 d optical band QPO in the narrow-line Seyfert 1 galaxy KIC 9650712 from densely sampled Kepler data has been made by Smith et al. (2018); it was supported by an independent analysis of the same data, indicating a QPO contribution at 52 \u00b1 2 d (Phillipson et al. 2020). Some possibly related QPOs of a few hundred days in two widely separated bands have been reported (Sandrinelli et al. 2016a; Sandrinelli, Covino & Treves 2016b; Sandrinelli et al. 2017). However, an analysis of the Fermi\u2013LAT and aperture photometry light curves by Covino et al. (2019) argued that some multiwaveband QPOs, along with many earlier claims of gamma-ray QPOs, are not significant. Among the gamma-ray QPOs with month-like periods, none showed simultaneous oscillations in a different wavebands. Evidence for related QPOs in multiple wavebands was observed in PG 1553+113, where a QPO was detected in the 0.1\u2013300 GeV and the optical waveband (Ackermann et al. 2015). The observed QPO had a dominant period of \u223c754 d and the source showed strong inter-waveband cross-correlations.","Citation Text":["Gupta et al.","2019"],"Functions Text":["There have been rather strong claims of AGN QPOs in different bands of the electromagnetic spectrum, ranging from minutes through days through months and years (e.g."],"Functions Label":["Background"],"Citation Start End":[[245,257],[264,268]],"Functions Start End":[[0,165]]}
{"Identifier":"2015MNRAS.450..943T__Patterson_1984_Instance_1","Paragraph":"In order to recover the intrinsic SED, we employ the values for interstellar extinction E(B \u2212 V) from Schlafly & Finkbeiner (2011), collected on NASA's Infrared Science Archive (IRSA) web pages, to deredden the spectra. We use the relations derived by Cardelli, Clayton & Mathis (1989) as implemented in iraf, and a standard value for the ratio of the total to the selective extinction R(V) = A(V)\/E(B \u2212 V) = 3.1. The corresponding values are available as the average extinction within a 2 \u00d7 2\u2009deg2 field. This represents a severe limitation as it does not take into account potential small-scale variations in the dust structure, nor the actual amount of absorbing material in the line of sight to the target. Since CVs are intrinsically rather faint (typically, MV > 4 mag; e.g. Patterson 1984), in general the correction for the interstellar extinction will represent an upper limit. An additional uncertainty regarding the intrinsic SED is introduced by obtaining the flux calibration on the basis of a single standard star that was observed on a different night than the targets. Furthermore, the observations were not conducted at a parallactic angle and thus might be affected by chromatic slit loss due to differential refraction, since Gemini-South is not equipped with an atmospheric dispersion corrector. As a measure for the SED, we derive the exponent \u03b1 of a power law F\u2009\u221d\u2009\u03bb\u2212\u03b1 that was fitted to the spectra, restricting the continuum to wavelengths 5000\u20137000\u2009\u00c5 and masking strong emission and absorption lines. During that analysis we found that most of the post-novae required a two-component fit while previously this represented an exception (Paper I, Paper IV). Furthermore, in every case the two slopes differed in the same direction, with the red part of the spectrum requiring a steeper slope than the blue part, and the points that separated the two slopes were found to be within 5850\u20135900\u2009\u00c5 which marks the region that is affected by the gap between two CCDs (see Section 2). All this suggests a systematic effect in the definition of the instrumental response function, and potentially a difference in the spectral efficiency of these two CCDs. We have thus further divided the fitting range into a blue (5000  \u03bb  5870\u2009\u00c5) and a red (5870  \u03bb  7000\u2009\u00c5) part that are fitted individually. The consequences for the interpretation of the SEDs are further discussed in Section 4.","Citation Text":["Patterson 1984"],"Functions Text":["Since CVs are intrinsically rather faint (typically, MV > 4 mag; e.g.",", in general the correction for the interstellar extinction will represent an upper limit."],"Functions Label":["Uses","Uses"],"Citation Start End":[[781,795]],"Functions Start End":[[711,780],[796,886]]}
{"Identifier":"2021ApJ...922...85T__Porquet_et_al._2004_Instance_1","Paragraph":"When the primary X-ray continuum is scattered by distant matter at several thousands of gravitational radii, the line is narrow (FWHM  10,000 km s\u22121) and is unlikely to carry any information about the conditions in the accretion disk and the strong-gravity regime near the SMBH. Narrow lines have been detected in the great majority of AGNs with luminosities L\nX,2.0\u221210.0 keV  1045 erg s\u22121. The mean FWHM reported from Chandra High Energy Transmission Grating (HETG) spectra is \u223c2000 km s\u22121 (Yaqoob & Padmanabhan 2004; Shu et al. 2010, 2011; see also Nandra 2006). Low equivalent width (EW) values from Suzaku observations (tens of eV; Fukazawa et al. 2011a) have also been interpreted as evidence of narrow lines originating in distant material (Ricci et al. 2014). In contrast, when the primary X-ray continuum is scattered close to the SMBH, Doppler and general relativistic effects combined may give rise to a significantly broader line (FWHM tens of thousands km s\u22121), reported in at least \u223c36% of AGNs (de la Calle P\u00e9rez et al. 2010; see also, e.g., Porquet et al. 2004; Jim\u00e9nez-Bail\u00f3n et al. 2005; Guainazzi et al. 2006; Nandra et al.2007; Brenneman & Reynolds 2009; Patrick et al. 2012; Liu & Li 2015; Mantovani et al. 2016; Baronchelli et al. 2018). In this case, contributions to the line\u2019s width become increasingly stronger as the primary continuum is scattered closer, and up to, the innermost stable circular orbit (ISCO), where the accretion disk\u2019s inner edge is located. Since the ISCO location depends directly on black hole spin, the latter leaves an imprint on the broad line\u2019s profile and can in principle be measured by means of relativistic modeling of the line via the X-ray reflection method of spin determination. As a result, there is a significant number of SMBH spin measurements (e.g., Brenneman 2013, and references therein). SMBH spin constraints have very significant implications for understanding both SMBHs and the way they affect their environment. Apart from potentially constituting a Kerr-metric-based test of general relativity in the strong-field regime, spin measurements can inform on jet-driving mechanisms, e.g., via extraction of SMBH rotational energy (Blandford & Znajek 1977), which in turn critically affect galactic environments and evolution.","Citation Text":["Porquet et al. 2004"],"Functions Text":["In contrast, when the primary X-ray continuum is scattered close to the SMBH, Doppler and general relativistic effects combined may give rise to a significantly broader line (FWHM tens of thousands km s\u22121), reported in at least \u223c36% of AGNs","see also, e.g.,"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1056,1075]],"Functions Start End":[[767,1007],[1040,1055]]}
{"Identifier":"2021MNRAS.504.6198M__McCoy_et_al._2017_Instance_1","Paragraph":"Here, we present the GMC properties within the molecular disc of the closest giant elliptical galaxy, NGC\u20095128, which is the host of the radio-source Centaurus A (hereafter Cen\u2009A). Cen A is at a distance of only D \u2243 3.8\u2009Mpc (Harris, Rejkuba & Harris 2010, 1 arcsec = 18\u2009pc) and it is therefore by far the most adequate target in the class of giant elliptical galaxies as well as powerful radio galaxies for studies of their molecular gas with high resolution. Indeed, Cen A is a peculiar case of an elliptical galaxy whose gaseous component has been supplied a few 0.1\u2009Gyr ago by the accretion of a HI rich galaxy (e.g. Struve et al. 2010). Along the dust lane of the elliptical galaxy, there is a molecular gas component of mass \u223c109\u2009M\u2299 as probed by various molecular lines (e.g. Phillips et al. 1987; Eckart et al. 1990; Rydbeck et al. 1993; Liszt 2001; Espada et al. 2009; Espada 2013; McCoy et al. 2017), partially seen in the form of kpc scale spiral features (Espada et al. 2012). The dust lane is along the minor axis, different to other ellipticals where discs are usually along the major axis (Young 2002). The molecular gas is associated with other components of the ISM, such as ionized gas traced by the H\u2009\u03b1 line (e.g. Nicholson, Bland-Hawthorn & Taylor 1992), near-infrared continuum (Quillen, Graham & Frogel 1993), submillimetre continuum (e.g. Hawarden et al. 1993; Leeuw et al. 2002), and mid-IR continuum emission (e.g. Mirabel et al. 1999; Quillen et al. 2006). In the inner hundreds of parsecs there is a circumnuclear disc (CND) of 400\u2009pc total extent (\u223c24 arcsec) and a P.A. = 155\u00b0, perpendicular to the inner jet, at least as seen in projection (Espada et al. 2009). The total gas mass in this component has been estimated to be 9 \u00d7 10$^7\\, \\mathrm{M}_{\\odot }$ (Israel et al. 2014, 2017). More detailed studies of the CND with higher resolutions of \u223c5\u2009pc in CO(3\u20132) and CO(6\u20135) have revealed the complexity of the molecular gas distribution and kinematics in that region, with multiple internal filaments and shocks (Espada et al. 2017).","Citation Text":["McCoy et al. 2017"],"Functions Text":["Along the dust lane of the elliptical galaxy, there is a molecular gas component of mass \u223c109\u2009M\u2299 as probed by various molecular lines"],"Functions Label":["Background"],"Citation Start End":[[889,906]],"Functions Start End":[[641,774]]}
{"Identifier":"2015ApJ...811L..32H__Hellinger_et_al._2013_Instance_1","Paragraph":"During the second phase, protons are heated. For negligible heat fluxes, collisions, and fluctuations, one expects the double adiabatic behavior or CGL (Chew et al. 1956; Matteini et al. 2012): the parallel and perpendicular temperatures (with respect to the magnetic field) are expected to follow \n\n\n\n\n\n and \n\n\n\n\n\n respectively. Tp\u22a5 decreases slower than \n\n\n\n\n\n during the whole simulation; protons are heated in the perpendicular direction while in the parallel direction the heating lasts till about t \u223c 0.25te0, whereas afterward protons are cooled. The parallel and perpendicular heating rates could be estimated as (see Verscharen et al. 2015)\n1\n\n\n\n\n\nA more detailed analysis indicates that between t = 0.1te0 and t = 0.7te0 the parallel heating rate \n\n\n\n\n\n smoothly varies from about 0.2 Qe and \n\n\n\n\n\n whereas \n\n\n\n\n\n is about constant at \u223c0.2 Qe; here, \n\n\n\n\n\n In total, protons are heated until t \u223c 0.7te0, and the heating reappears near the end of the simulation t \u2273 0.95te0. Note that the perpendicular heating rate is a nonnegligible fraction of that observed in the solar wind, where \n\n\n\n\n\n (Hellinger et al. 2013); however, the proton heating in 2D hybrid simulations is typically quite sensitive to the used electron equation of state (Parashar et al. 2014) and also to the resistivity and the number of particles per cell used (Franci et al. 2015a). The turbulent heating is, however, not sufficient to overcome the expansion-driven perpendicular cooling as in the solar wind (Matteini et al. 2007). During the third phase, t \u2273 0.7te0; there is an enhancement of the parallel cooling and perpendicular heating that cannot be ascribed to the effect of the turbulent activity. For a large parallel proton temperature anisotropy, a firehose instability is expected. The presence of such an instability is supported by the fact that the fluctuating magnetic field increases (with respect to the linear prediction), suggesting a generation of fluctuating magnetic energy at the expense of protons. To analyze the role of different processes in the system, we estimate their characteristic times (Matthaeus et al. 2014). The bottom panel of Figure 2 compares the turbulent nonlinear eddy turnover time \n\n\n\n\n\n at kdp = 1 (see Matthaeus et al. 2014; the expansion time te, and the linear time tl of the oblique firehose (Hellinger & Matsumoto 2000, 2001) estimated as \n\n\n\n\n\n where \u03b3m is the maximum growth rate calculated from the average plasma properties in the box assuming bi-Maxwellian proton velocity distribution functions (Hellinger et al. 2006). The expansion time te is much longer than tnl at kdp = 1 (as well as at the injection scales). The expanding system becomes theoretically unstable with respect to the oblique firehose around t \u223c 0.47te0 but clear signatures of a fast proton isotropization and of a generation of enhanced magnetic fluctuations appear later t \u2273 0.7te0. This is about the time when the linear time becomes comparable to the nonlinear time at ion scales. After that, \n\n\n\n\n\n slightly increases as a result of a saturation of the firehose instability, whereas \n\n\n\n\n\n at kdp = 1 is about constant (note that \u03a9p decreases as R\u22122). This may indicate that the instability has to be fast enough to compete with turbulence; however, the 2D system has strong geometrical constraints. Also the stability is governed by the local plasma properties. Figure 3 shows the evolution of the system in the plane \n\n\n\n\n\n During the evolution, a large spread of local values in the 2D space \n\n\n\n\n\n develops. Between t \u2243 0.1te0 and t \u2243 0.65te0, the average quantities evolve in time following \n\n\n\n\n\n This anticorrelation is qualitatively similar to in situ Helios observations between 0.3 and 1 AU (Matteini et al. 2007). During the third stage, when the strong parallel temperature anisotropy is reduced, both local and average values of \n\n\n\n\n\n and Ap appear to be bounded by the linear marginal stability conditions of the oblique firehose (Hellinger & Tr\u00e1vn\u00ed\u010dek 2008), although relatively large theoretical growth rates \n\n\n\n\n\n are expected.","Citation Text":["Hellinger et al. 2013"],"Functions Text":["Note that the perpendicular heating rate is a nonnegligible fraction of that observed in the solar wind, where"],"Functions Label":["Uses"],"Citation Start End":[[1103,1124]],"Functions Start End":[[984,1094]]}
{"Identifier":"2021MNRAS.503.2776Y__Ajith_et_al._2007_Instance_2","Paragraph":"In order to investigate the signal-to-noise ratio (SNR), \u03c1 of NS\u2013WD binaries for LISA-type space GW detectors, we calculate the averaged square SNR $\\overline{\\rho ^{2}}$ over the sky location, inclination, and polarization as \n(30)$$\\begin{eqnarray*}\r\n\\overline{\\rho ^{2}} = \\int _{f_{1}}^{f_{2}}\\frac{4\\cdot \\frac{4}{5}fA^{2}(f)}{(P_{\\rm n}(f)\/R(f))} \\rm d (\\ln \\it f),\r\n\\end{eqnarray*}$$(Moore, Cole & Berry 2015; Robson, Cornish & Liu 2019), where f1 and f2 are the lower and upper limits of the integral, respectively. The factor 4 in the numerator of the integrand comes from the addition of strain noise in the detector arms and the two-way noise in each arm (Larson, Hiscock & Hellings 2000). We calculate the GW amplitudes A(f) of NS\u2013WD binaries using the phenomenological (PhenomA) waveform model in the Fourier domain (Ajith et al. 2007; Robson et al. 2019). A(f) is expressed as \n(31)$$\\begin{eqnarray*}\r\nA(f) = \\sqrt{\\frac{5}{24}}\\frac{G^{5\/6}\\mathcal {M}^{5\/6}}{\\pi ^{2\/3}c^{3\/2}R_{\\rm b}}f^{-7\/6}\\, {\\rm Hz}^{-1},\\,\\,\\,\\it f\\lt f_{\\rm m},\r\n\\end{eqnarray*}$$\n (32)$$\\begin{eqnarray*}\r\nf_{\\rm m} = \\frac{0.2974\\zeta ^{2}+0.04481\\zeta +0.09556}{\\pi (GM\/c^{3})}\\, {\\rm Hz},\r\n\\end{eqnarray*}$$\n (33)$$\\begin{eqnarray*}\r\n\\zeta = m_{1}m_{2}\/M^{2},\r\n\\end{eqnarray*}$$where $\\mathcal {M} \\equiv m_1^{3\/5} m_2^{3\/5}\/(m_1+m_2)^{1\/5}$, and fm is the GW frequency at the point of merging. If f > fm, the index of the power-law relation between A(f) and f changes (Ajith et al. 2007) and is beyond the scope of this study. The power spectral density of total detector noise $P_{\\rm n}=\\frac{1}{L^{2}}\\left[P_{\\rm o}+2(1+\\cos ^{2}(f\/f_{\\ast }))\\frac{P_{\\rm a}}{(2\\pi f)^{4}}\\right]$, where f* = c\/(2\u03c0L), L = 2.5 \u00d7 109\u2009m is the armlength of the detector, $P_{\\rm o}=2.25\\times 10^{-22} \\,\\rm m^{2}\\left(1+(\\frac{2\\,mHz}{\\it f})^{4}\\right) \\,\\, \\rm Hz^{-1}$ is the single-link optical metrology noise, and $P_{\\rm a}=9.0\\times 10^{-30} \\,\\rm (m\\,s^{-2})^{2}\\left(1+(\\frac{0.4\\,mHz}{\\it f})^{2}\\right)\\left(1+(\\frac{\\it f}{8\\,\\rm mHz})^{4}\\right) \\,\\,Hz^{-1}$ is the single test mass acceleration noise (LISA Science Study Team 2018; Robson et al. 2019). R(f) is the transfer function numerically calculated from Larson et al. (2000). The effective noise power spectral density can be defined as Sn(f) = Pn(f)\/R(f). For Taiji and Tianqin, we use the sensitivity curve data in Ruan et al. (2020) and Wang et al. (2019), respectively.","Citation Text":["Ajith et al. 2007"],"Functions Text":["If f > fm, the index of the power-law relation between A(f) and f changes","and is beyond the scope of this study."],"Functions Label":["Uses","Future Work"],"Citation Start End":[[1466,1483]],"Functions Start End":[[1391,1464],[1485,1523]]}
{"Identifier":"2019MNRAS.484.2000D__Sandage_&_Fouts_1987_Instance_1","Paragraph":"To understand the distribution of Li-rich giants among different stellar components of the Galaxy, we have computed membership probability, based on the recipe given in Reddy, Lambert & Allende Prieto (2006) and references therein, for each of the stars in the selected RGB sample for belonging to one of the three main components of the Galaxy, namely; thin disc, thick disc, and halo. For this, the heliocentric velocities (U, V, W) for each of the sample giant is calculated using the Gaia astrometry (positions, parallax, and proper motions) and radial velocities (RV) from the GALAH DR2. The heliocentric velocities are corrected for the solar motion using Uo = 10, Vo = 5.3, Wo = 7.2 (km s\u22121) from Dehnen & Binney (1998) to get velocities with respect to the local standard of rest (ULSR, VLSR, WLSR). Entire sample is shown in Toomre Diagram of rotational velocity (VLSR) and sum of quadrature of radial and vertical velocities ([$\\rm {\\it U}_{\\rm LSR}^2 + {\\it W}_{\\rm LSR}^2]^{1\/2}$) (Sandage & Fouts 1987). Li-rich giants are identified by coloured symbols as shown in Fig. 4. We used kinematic boundaries for thin disc ($|V_{\\rm Total}|$  = $[U_{\\rm LSR}^2 + V_{\\rm LSR}^2 + W_{\\rm LSR}^2]^{1\/2}$ \u2264 80\u2009km s\u22121), thick disc (80 |VTotal| \u2264 200\u2009km s\u22121) and halo (|VTotal| > 200\u2009km s\u22121) which are in accordance with the results in Reddy et al. (2006). Further to quantify the number of giants belonging to different components of the Galaxy, we used probability of 70${{\\ \\rm per\\ cent}}$ or more for any star being considered as a member of particular component. We found 223 Li-rich thin disc stars, which is about 0.8${{\\ \\rm per\\ cent}}$ of total thin disc RGB stars with $P_{\\rm thin} \\ge 70 {{\\ \\rm per\\,cent}}$ and 69 Li-rich thick disc stars consisting of 0.5${{\\ \\rm per\\ cent}}$ of total thick disc RGB stars with $P_{\\rm thick} \\ge 70{{\\ \\rm per\\,cent}}$. We have found just three Li-rich giants among 1442 halo giants with $P_{\\rm halo} \\ge 70{{\\ \\rm per\\,cent}}$. This shows that Li-rich giants are rare ($\\lt 1{{\\ \\rm per\\,cent}}$) across stellar components but are relatively more prevalent among metal-rich thin disc component compared to thick disc and very metal poor halo.","Citation Text":["Sandage & Fouts 1987"],"Functions Text":["Entire sample is shown in Toomre Diagram of rotational velocity (VLSR) and sum of quadrature of radial and vertical velocities ([$\\rm {\\it U}_{\\rm LSR}^2 + {\\it W}_{\\rm LSR}^2]^{1\/2}$)"],"Functions Label":["Uses"],"Citation Start End":[[994,1014]],"Functions Start End":[[808,992]]}
{"Identifier":"2022AandA...659A.180G__Giannattasio_et_al._2013_Instance_1","Paragraph":"In the last few decades, the dynamic properties of the quiet Sun have been thoroughly investigated using a range of substantially different techniques, allowing us to elaborate a consistent picture of the photospheric dynamics by approaching the problem from different points of view. Particularly interesting and promising are the studies involving the tracking of small-scale magnetic fields in the quiet photosphere. Such investigations reveal features that still cannot be captured by theoretical models and\/or simulations because of the complexity of the system and the simultaneous coupling of a wide range of spatial and temporal scales. These studies are based on the hypothesis that magnetic fields are passively transported by the plasma flow and provide a characterisation of advection and diffusion processes in the quiet photosphere from granular to supergranular scales (see, e.g. Wang 1988; Berger et al. 1998; Cadavid et al. 1998, 1999; Hagenaar et al. 1999; Lawrence et al. 2001; S\u00e1nchez Almeida et al. 2010; Abramenko et al. 2011; Manso Sainz et al. 2011; Lepreti et al. 2012; Giannattasio et al. 2013, 2014a,b; Giannattasio et al. 2019; Keys et al. 2014; Caroli et al. 2015; Del Moro et al. 2015; Yang et al. 2015a,b; Roudier et al. 2016; Abramenko 2017; Kutsenko et al. 2018; Agrawal et al. 2018; Giannattasio & Consolini 2021). These studies reveal an anomalous scaling of magnetic field transport with a superdiffusive character consistent with a Levy walk inside supergranules and a more Brownian-like motion in their boundary. This has provided constraints on the magnetic flux emergence and evolution that models have to consider in order to fully explain the dynamics governing these environments. At the same time, the scaling laws affecting magnetic fields in the quiet Sun are crucial to understanding how such fields vary with scale size, despite the small scales at which dissipation occurs still being inaccessible with the currently available observations (see, e.g. Lawrence et al. 1994; Stenflo 2012, and references therein). For example, in the milestone work by Stenflo (2012), the spectrum of magnetic flux density in the quiet Sun was found to be consistent with a Kolmogorov power-law scaling. The scale at which scale invariance is broken lies below the current resolution limit. This latter author argued that the collapse of magnetic fields in Kilogauss flux tubes injects energy that is expected to cascade down because of the flux decay occurring via interchange instability, and to fragment into weaker \u2018hidden\u2019 fields at smaller scales (down to \u223c10 km). As far as we know, no other studies have focused on the scaling properties characterising the magnetic fields in a quiet Sun region within a range of spatial and temporal scales from (sub)granular to supergranular in the time domain. In this work, for the first time we apply the structure function analysis typical of complex systems (Frisch 1995) to fill this gap. This approach complements the studies based on feature tracking mentioned above. The main difference is that, while in those works statistical properties of the photospheric plasma flows are investigated via the transport of small-scale magnetic fields in a frozen-in condition, here we directly study the magnetic field variations emerging from magnetogram time-series. The paper is organised as follows. Section 2 describes the data set used and the analysis techniques applied. Section 3 describes the obtained results, while Sect. 4 is devoted to their discussion in the light of current literature. Finally, in Sect. 5, we present our conclusions and present future perspectives.","Citation Text":["Giannattasio et al. 2013"],"Functions Text":["These studies are based on the hypothesis that magnetic fields are passively transported by the plasma flow and provide a characterisation of advection and diffusion processes in the quiet photosphere from granular to supergranular scales (see, e.g."],"Functions Label":["Background"],"Citation Start End":[[1095,1119]],"Functions Start End":[[645,894]]}
{"Identifier":"2022ApJ...930..125L__Zimbardo_et_al._2017_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 et al. 2017"],"Functions Text":["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"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[627,647]],"Functions Start End":[[0,545]]}
{"Identifier":"2018ApJ...854...26L___2015a_Instance_4","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"],"Functions Text":["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",", 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"],"Functions Label":["Uses","Uses"],"Citation Start End":[[3057,3066],[3073,3078]],"Functions Start End":[[2530,2766],[2909,3055]]}
{"Identifier":"2015MNRAS.454.1117S___2005_Instance_1","Paragraph":"One of the first models of MRI-driven accretion in protoplanetary discs was constructed by Gammie (1996), who proposed that Ohmic resistivity (due to weak coupling of charged species to the magnetic field) quenches the MRI in a \u2018dead\u2019 mid-plane region surrounded by \u2018active layers\u2019, which are ionized by cosmic rays. This layered-accretion scenario has been refined as additional non-thermal ionization sources (X-rays, UV radiation) and non-ideal magnetohydrodynamic (MHD) effects (ambipolar diffusion and the Hall effect) have been included (e.g. Igea & Glassgold 1999; Sano et al. 2000; Fromang, Terquem & Balbus 2002; Salmeron & Wardle 2003, 2005, 2008; Ilgner & Nelson 2006; Bai & Goodman 2009; Bai & Stone 2011; Wardle & Salmeron 2012). Several theoretical studies, spanning both analytical (e.g. Blaes & Balbus 1994; Desch 2004; Kunz & Balbus 2004) and numerical work (e.g. Hawley & Stone 1998; Bai & Stone 2011, 2013b; Simon et al. 2013a,b), have shown that ambipolar diffusion plays a major role in determining the level of turbulence in protoplanetary discs. In the outer regions (radii R \u2273 30\u2009au), ambipolar diffusion acts to reduce the strength of MRI-driven turbulence close to the disc mid-plane; turbulence is stronger above this \u2018ambipolar damping zone\u2019 in a thin layer of gas strongly ionized by FUV photons (Perez-Becker & Chiang 2011; Simon et al. 2013a,b). To produce accretion rates consistent with observations, a large-scale magnetic field threading the disc perpendicular to the disc plane is required (Simon et al. 2013a,b). This net field enhances turbulence, despite the damping effect of ambipolar diffusion, and allows additional angular-momentum loss via a magnetic wind (akin to the Blandford & Payne 1982 mechanism). In the inner disc, both local (Bai & Stone 2013b) and global (Gressel et al. 2015) simulations have shown that ambipolar diffusion can quench the MRI in the purported active layers of Gammie's (1996) model. In those cases, a large-scale magnetic field is again necessary to drive accretion at the observationally inferred rates, this time solely via a magnetic wind in a manner akin to previous models of wind-driven accretion in protoplanetary discs (e.g. Wardle & Koenigl 1993; Shu et al. 1994; Salmeron, K\u00f6nigl & Wardle 2007).","Citation Text":["Salmeron & Wardle","2005"],"Functions Text":["This layered-accretion scenario has been refined as additional non-thermal ionization sources (X-rays, UV radiation) and non-ideal magnetohydrodynamic (MHD) effects (ambipolar diffusion and the Hall effect) have been included (e.g."],"Functions Label":["Background"],"Citation Start End":[[622,639],[646,650]],"Functions Start End":[[317,548]]}
{"Identifier":"2020ApJ...896...59A__Henry_1989_Instance_1","Paragraph":"PNs are expanding shells of the luminous gas expelled by dying stars of low and intermediate masses (LIMS). They stem from objects that have lifetimes up to gigayears. The ionized gas surrounding the central star shows emission lines of highly ionized species from which the abundances can be derived. The neon abundance can be compared and related to the results from stellar abundance analysis, as neon (and oxygen) originates from primary nucleosynthesis in massive stars (\u226510 M\u2299) and is therefore nearly independent of the evolution of LIMS, the progenitor stars of PNs (Henry 1989; Henry et al. 2004). For comparison, we collected several studies, where the neon abundances were obtained in PNs. We would not say that the abundances from PNs are very accurate, because the differences between abundances from collisionally excited lines and optical recombination lines can be much higher in PNs compared to those of H ii regions. For example, according to Wang & Liu (2007), in four Galactic disk PNs (He 2\u2013118, H 1-35, NGC 6567, and M1-61), the mean neon abundance calculated from collisionally excited lines is 7.78 \u00b1 0.23, while it is 8.72 \u00b1 0.60 from optical recombination lines. Figure 9 presents neon abundances from different sources, including five studies focused on PNs. In those studies, where it was mentioned, we adopted neon abundances obtained from collisionally excited lines, e.g., Tsamis et al. (2003) and Wang & Liu (2007). The mean value of log \u03f5Ne = 7.84 \u00b1 0.24 was calculated on the basis of neon abundances in six Galactic disk PNs (Wang & Liu 2007). The mean value of the neon abundance log \u03f5Ne = 7.99 \u00b1 0.22 dex was calculated from six PNs (Hu 1-2, IC 418, NGC 40, NGC 2440, NGC 6543, NGC 7662) with Galactocentric distances from 7.9 kpc to 8.9 kpc from the summary of Pottasch & Bernard-Salas (2006). The value of log \u03f5Ne = 7.76 \u00b1 0.24 dex was derived from 16 Galactic PNs located at Galactocentric distances from 8.0 to 8.9 kpc in Stanghellini et al. (2006). The log \u03f5Ne = 8.02 \u00b1 0.25 dex was obtained from Marigo et al. (2003) from three PNs (NGC 6543, NGC 7027, NGC 7662), which are located at distances no more than 1.0 kpc. We adopted the mean value log \u03f5Ne = 8.15 \u00b1 0.21 dex, which was calculated from 16 Galactic PNs from Tsamis et al. (2003).","Citation Text":["Henry 1989"],"Functions Text":["The neon abundance can be compared and related to the results from stellar abundance analysis, as neon (and oxygen) originates from primary nucleosynthesis in massive stars (\u226510 M\u2299) and is therefore nearly independent of the evolution of LIMS, the progenitor stars of PNs"],"Functions Label":["Background"],"Citation Start End":[[575,585]],"Functions Start End":[[302,573]]}
{"Identifier":"2020MNRAS.497..829F__Diaz-Miller,_Franco_&_Shore_1998_Instance_1","Paragraph":"In order to consider both the photoionization and radiation-pressure effects from the accreting protostar, we solve the frequency-dependent radiative transfer for stellar irradiation by the following method. We inject photons from the sink cell to the computational domain at the rate\n(10)$$\\begin{eqnarray*}\r\nL_{\\rm *}^{\\rm tot} = L_* + L_{\\rm acc},\r\n\\end{eqnarray*}$$where L* is the stellar luminosity and $L_{\\rm acc} \\equiv G M_* \\dot{M}_* \/ R_*$ is the accretion luminosity. The radiation spectrum is assumed to be thermal blackbody $L_{*, \\nu }^{\\rm tot} \\propto B_\\nu (T_{\\rm eff})$, where Teff is the effective temperature defined as\n(11)$$\\begin{eqnarray*}\r\nT_{\\rm eff} = \\left(\\frac{L_*^{\\rm tot}}{4 \\pi \\sigma R_*^2} \\right)^{1\/4} .\r\n\\end{eqnarray*}$$Assumption of the thermal blackbody radiation could overestimate the emissivity of ionizing photons at Z \u223c Z\u2299 owing to the line-blanketing effect, especially for stars with M* \u223c 10 M\u2299 (e.g. Diaz-Miller, Franco & Shore 1998). This approximation should be valid in our cases since we mainly consider the feedback from more massive stars exceeding 100 M\u2299. As discussed in Section 2.2.3, the effective temperature is higher for lower metallicity at a given mass: at M* = 100 M\u2299, for instance, Teff \u2243 105 K for Z = 0, Teff \u2243 7 \u00d7 104 K for Z = 10\u22122\u2009Z\u2299, and Teff \u2243 5 \u00d7 104 K for Z = Z\u2299. We use a hybrid method where the direct light emitted from the central star and diffuse light re-emitted from the accretion envelope are separately solved (e.g. Kuiper et al. 2010a, b). We solve the transfer of the stellar direct component by means of the ray-tracing method,\n(12)$$\\begin{eqnarray*}\r\nF_{\\nu }(r) = \\frac{L_{*,\\nu }^{\\rm tot}}{4 \\pi r^2} \\mathrm{ e}^{-\\tau _{\\nu }},\r\n\\end{eqnarray*}$$where F\u03bd(r) is the flux at the radial position r and \u03c4\u03bd is the optical depth\n(13)$$\\begin{eqnarray*}\r\n\\tau _{\\nu } = \\int _{r_{\\rm in}}^{r} [ n_{\\rm H\\,{\\small I}} \\sigma _{\\rm H\\,{\\small I}}(\\nu) + \\rho \\kappa _{\\rm \\mathrm{ d}} (\\nu)] \\mathrm{ d}r,\r\n\\end{eqnarray*}$$where \u03c3H\u2009i(\u03bd) is the H\u2009i photoionization cross-section (e.g. Osterbrock 1989), \u03bad(\u03bd) the dust opacity given by Laor & Draine (1993), and rin the sink radius. We assume that the dust-to-gas mass ratio decreases with decreasing metallicities, linearly scaling with Z. The dust opacity is determined by the scaled dust-to-gas mass ratio. We also assume that photons freely travel without attenuation for R* r rin. We set 200 logarithmically spaced wavelength bins between 0.03 cm and 1.5 nm, achieving the higher resolution for the shorter wavelengths. The hydrogen photoionization rate is calculated by integrating the contributions by photons below the Lyman limit, i.e. 91.2 nm (Appendix A2.3). We do not include the diffuse ionizing photons, as their contribution to the disc photoevaporation is negligible compared with the stellar direct component in the massive star formation (e.g. Hosokawa et al. 2011; Tanaka, Nakamoto & Omukai 2013). As in Hosokawa et al. (2016), we do not incorporate helium ionization for simplicity. We assume that helium is in the atomic state everywhere including H\u2009ii regions. We consider the transfer of He ionizing photons above 54.4 eV emitted from the protostar, but they are consumed to photoionize hydrogen in our simulations. Regarding the photodissociation of hydrogen molecules, we do not use F\u03bd(r) because our wavelength grids are too sparse to resolve the Lyman\u2013Werner bands. We instead consider another monochromatic direct component representing FUV (11.2 eV \u2264 h\u03bd \u2264 13.6 eV) photons\n(14)$$\\begin{eqnarray*}\r\nF_{\\rm FUV} (r) = \\frac{L_{\\rm FUV}}{4 \\pi r^2} \\exp (- \\tau _{\\rm d, FUV}),\r\n\\end{eqnarray*}$$where LFUV is the luminosity in the FUV wavelength range, and \u03c4d,FUV is the dust opacity represented by the value at \u03bb = 100\u2009nm. The photodissociation rate of hydrogen molecules is estimated with the flux FFUV, including the $\\rm H_2$ self-shielding. The details of the calculation are shown in Appendix A2.3","Citation Text":["Diaz-Miller, Franco & Shore 1998"],"Functions Text":["Assumption of the thermal blackbody radiation could overestimate the emissivity of ionizing photons at Z \u223c Z\u2299 owing to the line-blanketing effect, especially for stars with M* \u223c 10 M\u2299 (e.g."],"Functions Label":["Uses"],"Citation Start End":[[952,984]],"Functions Start End":[[762,951]]}
{"Identifier":"2021AandA...651A..71L__Samland_et_al._2017_Instance_1","Paragraph":"Both SHINE and GPIES surveys recently discovered three new exoplanets (Macintosh et al. 2015; Chauvin et al. 2017b; Keppler et al. 2018) and a few additional higher-mass brown dwarfs (Konopacky et al. 2016; Cheetham et al. 2018). Smaller surveys using SPHERE and GPI also discovered several substellar companions (Milli et al. 2017a; Wagner et al. 2020; Bohn et al. 2020). These surveys offer unprecedented detection, astrometric, and spectrophotometric capabilities that allow us to characterize fainter and closer giant planets, such as the recent discovery of 51 Eri b (2 MJup at 14 au, T5-type, of an age of 20 Myr; Macintosh et al. 2015; Samland et al. 2017), HIP 65426 b, a young, warm, and dusty L5-L7 massive Jovian planet located at about 92 au from its host star (Chauvin et al. 2017b), and the young solar analogue PDS 70, which is now known to actually host two planets PDS 70 b discovered during the SHINE campaign (Keppler et al. 2018; M\u00fcller et al. 2018) and PDS 70 c by MUSE (Haffert et al. 2020). Such surveys also provide key spectral and orbital characterisation data for known exoplanets (e.g., De Rosa et al. 2016; Samland et al. 2017; Chauvin et al. 2018; Wang et al. 2018; Cheetham et al. 2019; Lagrange et al. 2019; Maire et al. 2019). Despite these new discoveries, SHINE and GPIES have yielded significantly fewer exoplanet detections than predicted, with their use of extrapolations of radial velocity planet populations to larger semi-major axes (e.g., Cumming et al. 2008). This results in our setting strong statistical constraints on the distribution of giant exoplanets at separations of >10 au from their stars, as well as sub-stellar companions to young stars (Nielsen et al. 2019; Vigan et al. 2021). As these systems are young (100 Myr), and thus closer to their epoch of formation than, for instance, radial-velocity planets (typically 1\u201310 Gyr), statistical analyses of large direct-imaging surveys can provide hints for the potential formation mechanism responsible for producing giant exoplanets on wide orbits.","Citation Text":["Samland et al. 2017","Samland et al. 2017"],"Functions Text":["These surveys offer unprecedented detection, astrometric, and spectrophotometric capabilities that allow us to characterize fainter and closer giant planets, such as the recent discovery of 51 Eri b (2 MJup at 14 au, T5-type, of an age of 20 Myr;","Such surveys also provide key spectral and orbital characterisation data for known exoplanets (e.g.,"],"Functions Label":["Background","Background"],"Citation Start End":[[643,662],[1136,1155]],"Functions Start End":[[373,619],[1014,1114]]}
{"Identifier":"2019ApJ...883...73C__Engelbrecht_&_Burger_2013a_Instance_1","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"],"Functions Text":["were used as inputs for the turbulence quantities (e.g.,","Furthermore, mean free paths such as those discussed have,"],"Functions Label":["Background","Similarities"],"Citation Start End":[[889,915]],"Functions Start End":[[832,888],[966,1024]]}
{"Identifier":"2022MNRAS.516.3381J__Lindblom_&_Owen_2002_Instance_2","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":["Lindblom & Owen 2002"],"Functions Text":["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"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[2311,2331]],"Functions Start End":[[2048,2309]]}
{"Identifier":"2020AandA...642A..90M__Werner_et_al._2006_Instance_1","Paragraph":"Additional information can be derived from the abundance ratios measured in the ICM. For example, the Mn\/Fe and Ni\/Fe ratios are both sensitive to the electron capture rates during SNIa explosions, and are therefore crucial indicators of their progenitor channels (Seitenzahl et al. 2013a; Mernier et al. 2016a; Hitomi Collaboration 2017), while ratios of lighter elements can in principle provide constraints on the IMF and the initial metallicity of the SNcc progenitors (e.g. de Plaa et al. 2007; Mernier et al. 2016a). By pushing current observatories to their limit, recent studies derived constraints on the relative fraction of SN events that effectively contribute to the ICM enrichment. These measurements showed that SNIa and SNcc contribute relatively equally to the overall chemical enrichment in the ICM (Werner et al. 2006; de Plaa et al. 2006, 2017; Bulbul et al. 2012; Mernier et al. 2016a). The comprehensive study of Simionescu et al. (2019) compiled the most accurate abundance measurements of the Perseus cluster (taken with the XMM-Newton RGS and Hitomi SXS instruments) and compared them to state-of-the-art SNIa and SNcc yield models. Their surprising conclusion is that no current set of models is able to reproduce all the observed abundance ratios at once. In particular, the measured Si\/Ar ratio tends to be systematically overpredicted by models, even when taking calibration and atomic uncertainties into account. Whereas further improvement of stellar nucleosynthesis models is expected, the non-negligible systematic errors associated to these observations and the lack of highly sensitive, spatially resolved high-resolution spectroscopy prevents us from steering any considerable change in the paradigm (see also de Grandi & Molendi 2009). In fact, measurements from currently flying missions are performed with moderate collective area, either over the whole X-ray band (0.4\u201310 keV) with modest spectral resolution (> 100 eV), or with higher resolution dispersive spectroscopy but over the low E band (0.3\u20132 keV) and without any spatial resolution, which considerably limits interpretations.","Citation Text":["Werner et al. 2006"],"Functions Text":["By pushing current observatories to their limit, recent studies derived constraints on the relative fraction of SN events that effectively contribute to the ICM enrichment. These measurements showed that SNIa and SNcc contribute relatively equally to the overall chemical enrichment in the ICM"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[818,836]],"Functions Start End":[[523,816]]}
{"Identifier":"2016ApJ...825...10T__Brisken_et_al._2003_Instance_1","Paragraph":"Although this penalty is small, we can still provide a quantitative estimate based on a few assumptions. The observations of M15 that measured the parallax to VLA J2130+12 (K14) employed an interferometric technique where the position of a weaker potential in-beam calibrator source (M15 S1 and VLA J2130+12) can be measured against a brighter primary calibration source, with the potential for using the in-beam calibrator to transfer more accurate calibration solutions to other in-beam targets. Based on a literature review for papers where VLBI parallax measurements were determined using an in-beam calibrator, we find 24 similar measurements, the majority of which have been used to measure parallaxes to pulsars (Fomalont et al. 1999; Brisken et al. 2003; Chatterjee et al. 2005, 2009; Ng et al. 2007; Middelberg et al. 2011; Deller et al. 2012, 2013; Ransom et al. 2014; Liu et al. 2016). As these measurements represent additional opportunities to detect the parallax of a VLA J2130+12-like object, they could be considered as potential trials. However, all of the reported in-beam calibrators were significantly brighter (4\u201386 mJy) than VLA J2130+12 (\u223c0.1\u20130.5 mJy). Given our conservative assumption of a uniform volume density of VLA J2130+12-like objects in the Galaxy, we should down-weight the number of trials from brighter sources by (f\u03bd,in-beam\/f\u03bd,VLA J2130+12)\u22121.5. In that case, the number of trials penalty is only an additional \u223c0.2 trials. In addition, we note that the PSR\u03c0 parallax project (a large VLBA program) has reported15\n\n15\n\nhttps:\/\/safe.nrao.edu\/vlba\/psrpi\/\n\n 111 additional in-beam calibrator sources that they used to measure parallaxes. Although they do not provide the flux densities of individual sources, they note that their median in-beam calibrator source is 9.2 mJy. We have measured the flux density function of secure 1\u201320 mJy FIRST sources (Helfand et al. 2015; \n\n\n\n\n\n) to estimate the expected distribution of the flux densities of in-beam calibrators. We found that the minimum flux density is likely \u223c3.2 mJy, and using the same down-weighting we estimate an additional penalty of \u223c0.9 trials.","Citation Text":["Brisken et al. 2003"],"Functions Text":["Based on a literature review for papers where VLBI parallax measurements were determined using an in-beam calibrator, we find 24 similar measurements, the majority of which have been used to measure parallaxes to pulsars"],"Functions Label":["Similarities"],"Citation Start End":[[742,761]],"Functions Start End":[[498,718]]}
{"Identifier":"2016ApJ...833..192S__Champion_et_al._2008_Instance_1","Paragraph":"In recent years, several large-scale pulsar surveys have been undertaken to search for new pulsars (Cordes et al. 2006; Keith et al. 2010; Barr et al. 2013; Boyles et al. 2013; Deneva et al. 2013; Coenen et al. 2014; Stovall et al. 2014). One of the drivers for such surveys is the discovery of millisecond pulsars (MSPs). MSPs are formed through accretion from a companion during an X-ray binary phase (Alpar et al. 1982; Bhattacharya & van den Heuvel 1991) in which the pulsar is \u201crecycled.\u201d This accretion phase spins the pulsar up to very fast rotational rates (spin periods \n\n\n\n\n\n ms). Such pulsars are useful for a variety of physical applications. Examples include tests of theories of gravity using MSP\u2013white dwarf systems such as PSR J1738+0333 and PSR J0348+0432 (Freire et al. 2012; Antoniadis et al. 2013) and triple systems like PSR J0337+1715 (Ransom et al. 2014); tests of general relativity using double neutron star systems, such as J0737\u22123039 (Kramer et al. 2006) and PSR B1913+16 (Weisberg et al. 2010); the study of binary systems such as eccentric MSPs like PSRs J1903+0327 (Champion et al. 2008) and J1950+2414 (Knispel et al. 2015) which are interesting due to their peculiar binary evolution; and constraining the equation of state of dense matter using measurements of neutron star masses (Demorest et al. 2010; Antoniadis et al. 2013). Another major driver for the discovery of new MSPs is the effort to detect gravitational wave emission using an array of pulsars (NANOGrav Collaboration et al. 2015; Lentati et al. 2015; Reardon et al. 2016). The large-scale pulsar surveys mentioned above, combined with targeted searches of unidentified gamma-ray sources from the Fermi Gamma-Ray Space Telescope (e.g., Hessels et al. 2011; Keith et al. 2011; Ransom et al. 2011; Kerr et al. 2012), have resulted in the discovery of about 90 new MSPs in the past 5 years, an increase of 40% in the known Galactic MSP population. A subset of the newly discovered sources are eclipsing systems that appear to fall into two categories (e.g., Freire 2005; Roberts 2012). The first category, known as black widow systems, has very low mass, degenerate companions (\n\n\n\n\n\n) believed to be the result of ablation by the pulsar. The second, known as redback systems, has low to moderate mass, non-degenerate companions (\n\n\n\n\n\n).","Citation Text":["Champion et al. 2008"],"Functions Text":["the study of binary systems such as eccentric MSPs like PSRs J1903+0327","which are interesting due to their peculiar binary evolution"],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[1096,1116]],"Functions Start End":[[1023,1094],[1155,1215]]}
{"Identifier":"2020AandA...641A.123H__Mikal-Evans_et_al._(2019)_Instance_1","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"],"Functions Text":["To confirm the importance of VO, water and an inversion layer","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."],"Functions Label":["Background","Background","Compare\/Contrast"],"Citation Start End":[[82,105]],"Functions Start End":[[0,61],[113,240],[241,830]]}
{"Identifier":"2020ApJ...902...98G__Bournaud_et_al._2014_Instance_1","Paragraph":"On balance, a large abundance of baryon-dominated, dark matter cored galaxies at z \u223c 2, most strongly correlated with baryonic surface density, angular momentum, and central bulge mass, may be most naturally accounted for by the interaction of baryons and dark matter during the formation epoch of massive halos. Massive halos (log(Mhalo\/M\u2299) > 12) formed for the first time in large abundances in the redshift range z \u223c 1\u20133 (Press & Schechter 1974; Sheth & Tormen 1999; Mo & White 2002; Springel et al. 2005). At the same time, gas accretion rates were maximal (Tacconi et al. 2020). This resulted in high merger rates (Genel et al. 2008, 2009; Fakhouri & Ma 2009), very efficient baryonic angular momentum transport (Dekel et al. 2009; Zolotov et al. 2015), formation of globally unstable disks, and radial gas transport by dynamical friction (Noguchi 1999; Immeli et al. 2004; Genzel et al. 2008; Bournaud & Elmegreen 2009; Bournaud et al. 2014; Dekel & Burkert 2014). These processes enabled galaxy mass doubling on a timescale  0.4 Gyr at z \u223c 2\u20133, and massive bulge formation by disk instabilities and compaction events on 1 Gyr timescales. However, central baryonic concentrations would naturally also increase central dark matter densities through adiabatic contraction (Barnes & White 1984; Blumenthal et al. 1986; Jesseit et al. 2002). For adiabatic contraction to be ineffective requires the combination of kinetic heating of the central dark matter cusp by dynamical friction from in-streaming baryonic clumps (El-Zant et al. 2001; Goerdt et al. 2010; Cole et al. 2011), with feedback from winds, supernovae, and AGNs driving baryons and dark matter out again (Dekel & Silk 1986; Pontzen & Governato 2012, 2014; Martizzi et al. 2013; Freundlich et al. 2020; K. Dolag et al. 2020, in preparation). Using idealized Monte Carlo simulations, El-Zant et al. (2001) demonstrated that dynamical friction acting on in-spiraling gas clumps can provide enough energy to heat up the central dark matter component and create a finite dark matter core (see also A. Burkert et al. 2020, in preparation). They argue that dark matter core formation in massive galaxies would require that clumps be compact, such that they avoid tidal and ram-pressure disruption, and have masses of >108 M\u2299. Other idealized simulations (e.g., Tonini et al. 2006) confirm these results.","Citation Text":["Bournaud et al. 2014"],"Functions Text":["This resulted in","formation of globally unstable disks, and radial gas transport by dynamical friction"],"Functions Label":["Background","Background"],"Citation Start End":[[926,946]],"Functions Start End":[[584,600],[759,843]]}
{"Identifier":"2022MNRAS.512.4136C__Ventura_et_al._2013_Instance_1","Paragraph":"If we recall the tight, monotonic dependence of the position of galaxies along the SF sequence in the diagram with metallicity (as outlined in Section 3.1), we can interpret our global results of Figs 4 and 5 as a manifestation of the existence of an O\/H versus N\/O relation for SDSS star-forming galaxies, whose intrinsic scatter is reflected and, to some extent, translated into the observed distribution of emission line ratios within the [N\u2009ii]-BPT. A tight relationship between O\/H and N\/O abundances is indeed observed in both H\u2009ii regions and local galaxies, especially at M\u22c6 \u2273 109.5M\u2299 (Vila Costas & Edmunds 1993; van Zee et al. 1998; P\u00e9rez-Montero & Contini 2009; Pilyugin et al. 2012; Andrews & Martini 2013; Hayden-Pawson et al. 2021), and it is set by the predominant nucleosynthetic origin of nitrogen from CNO burning of pre-existing stellar carbon and oxygen in low- and intermediate-mass stars experiencing the AGB phase (i.e. the \u2018secondary\u2019 nitrogen production mechanism, Kobayashi, Karakas & Umeda 2011; Ventura et al. 2013; Vincenzo et al. 2016); alternatively, Vincenzo & Kobayashi (2018) reproduced the observed N\/O-O\/H relation introducing failed supernovae (SNe) in massive stars within their cosmological simulations. Recently, such relationship between O\/H and N\/O has been suggested as even tighter than the one between M\u22c6 and N\/O (Hayden-Pawson et al. 2021), in contrast to what claimed by previous studies (e.g. Andrews & Martini 2013; Masters et al. 2016). In light of our results, this would confirm that deviations in N\/O at fixed O\/H are more likely to be related to the offset from the SF sequence in the [N\u2009ii]-BPT than relative variations in M\u22c6, although the two are clearly physically correlated. The connection between the two diagrams is also readily evident if we look at the distribution of our galaxy sample in the N\/O versus O\/H diagram, as shown in Fig. 8 (where [N\u2009ii] \u03bb6584\/[O\u2009ii] \u03bb3727, 29 is converted to N\/O following the Te-based calibrations presented in Hayden-Pawson et al. 2021); here, each hexagonal bin is colour-coded by the average distance D of galaxies from the best-fitting line of the [N\u2009ii]-BPT, almost perfectly tracing the scatter around the median N\/O versus O\/H relation.","Citation Text":["Ventura et al. 2013"],"Functions Text":["A tight relationship between O\/H and N\/O abundances is indeed observed in both H\u2009ii regions and local galaxies,","and it is set by the predominant nucleosynthetic origin of nitrogen from CNO burning of pre-existing stellar carbon and oxygen in low- and intermediate-mass stars experiencing the AGB phase (i.e. the \u2018secondary\u2019 nitrogen production mechanism,"],"Functions Label":["Uses","Uses"],"Citation Start End":[[1023,1042]],"Functions Start End":[[454,565],[747,989]]}
{"Identifier":"2021ApJ...912..163B__Brennecka_et_al._2020_Instance_1","Paragraph":"Braukmuller et al. (2018) proposed that all elements fall into one of four categories based on their condensation temperature: refractory elements (50% condensation temperature, Tc,50 > 1400 K), which exhibit approximately uniform enrichments in their Si-normalized concentrations in CC chondrites compared to CI chondrites by a factor of \u223c1\u20131.4; main component elements (1300 K  Tc,50  1400 K), which have approximately the same Si-normalized elemental abundances in CC chondrites as CI chondrites (differ by a factor of \u223c0.8\u20131.1); slope-volatile elements (800 K  Tc,50  1300 K), which exhibit monotonically decreasing Si-normalized concentrations with decreasing Tc,50 compared to CI chondrites; and plateau volatile elements (Tc,50  800 K), which display uniform depletions in Si-normalized concentrations compared to CI chondrites by a factor of \u223c0.1\u20130.7 that are characteristic of each CC chondrite group. Given their uniform nature with Tc,50 and comparatively well-constrained isotopic and elemental compositions, we chose to focus on the concentrations of refractory, main component, and plateau volatile elements in this study. For the refractory and main component elements in CC chondrites, we examine the elemental and isotopic compositions of Ti and Cr, respectively, because these are lithophile elements whose isotopic compositions have been measured precisely for a large number of chondrites and their components (Trinquier et al. 2007, 2009; Qin et al. 2010; Olsen et al. 2016; Van Kooten et al. 2016; Gerber et al. 2017; Davis et al. 2018; Zhu et al. 2019; Schneider et al. 2020; Williams et al. 2020). For CC iron meteorites, we examine the isotopic compositions of Mo and Ni, respectively, because these are siderophile elements (so are therefore present in appreciable concentrations in iron meteorites, unlike Ti and Cr) whose compositions have also been relatively well studied in a number of iron meteorites as well as chondrites and their components (Burkhardt et al. 2011; Budde et al. 2016; Kruijer et al. 2017; Bermingham et al. 2018; Nanne et al. 2019; Budde et al. 2019; Worsham et al. 2019; Brennecka et al. 2020; Spitzer et al. 2020). For the plateau volatile elements, we examine the elemental compositions of six elements (Bi, Ag, Pb, Zn, Te, and Sn) that exhibit a number of desirable properties: their concentrations have been relatively well constrained in CC chondrites; they show a range of lithophile, siderophile, and chalcophile behaviors; their concentrations do not appear to be strongly dependent on redox state; they show minimal variability among NC chondrite groups. Our reasoning for not considering the isotopic compositions of these elements is discussed in Section 2.3. The adopted isotopic and chemical composition of each element used in this study in CC chondrites, CC iron meteorites, CAIs, CI chondrites, and NC chondrites are included in Table 1. Uncertainties on elemental concentrations have not been routinely reported throughout the literature, although these values are typically \u00b15 wt% (e.g., Lodders 2003; Palme et al. 2014). CAIs can be categorized into six groups based on their compositions (Stracke et al. 2012). For the purposes of this study, we adopt the composition of type I CAIs as the representative value of refractory objects because they are seemingly the most abundant type and lack the characteristic elemental depletions of other CAI groups (e.g., Stracke et al. 2012; Brennecka et al. 2020). We also focus largely on ordinary chondrites (OC) as representative NC meteorites rather than enstatite chondrites (EC) or Rumuruti chondrites (RC). This is because EC chondrites formed under more reducing conditions than OC and RC chondrites, which introduced a compositional signature for some elements to EC chondrites that is not present in OC, RC, or CC chondrites (presumably due to their formation in more oxidizing environments) so is not representative of large-scale mixing in the disk (Alexander 2019b). Additionally, the isotopic compositions of RC chondrites are sparsely measured compared to OC and EC chondrites. NC meteorites could have experienced a number of processes (e.g., mixing, chondrule formation, volatile loss, the addition of refractory materials, etc.) that gave these meteorites their specific chemical and isotopic signatures (Alexander 2019b). We do not explore these processes in this study and simply adopt the measured elemental and isotopic compositions of NC chondrites as potential end-members for the compositions of CC meteorites.","Citation Text":["Brennecka et al. 2020"],"Functions Text":["For CC iron meteorites, we examine the isotopic compositions of Mo and Ni, respectively, because these are siderophile elements (so are therefore present in appreciable concentrations in iron meteorites, unlike Ti and Cr) whose compositions have also been relatively well studied in a number of iron meteorites as well as chondrites and their components"],"Functions Label":["Motivation"],"Citation Start End":[[2123,2144]],"Functions Start End":[[1622,1975]]}
{"Identifier":"2015MNRAS.449.4326P__Collaboration_2006_Instance_1","Paragraph":"We now consider upcoming spectroscopic surveys. We consider two cases for the CMB lensing map, including (1) the full Planck CMB lensing map and (2) the Advanced ACTPol5 CMB lensing map. In both cases, we assume that the CMB lensing maps will be estimated using the temperature map and both E and B polarization maps, and we assume the B map only contains noise. We predict the noise in the Planck lensing map assuming the detector sensitivity and beam sizes listed in the Planck Bluebook (Planck Collaboration 2006). Advanced ACTPol will survey 20 000 deg2, and its increased temperature and polarization sensitivity will create a CMB lensing map that is an order of magnitude more sensitive than Planck. The specifications we use for Advanced ACTPol are listed in Table 3. For spectroscopic surveys, we consider the DESI emission line galaxy (ELG), LRG, and quasar surveys, as well as the Euclid H\u03b1 survey and the WFIRST H\u03b1 and O iii combined survey. The properties of the surveys are listed in Table 1. For DESI, we assume the same values as in the DESI Conceptual Design Report:6 bLRGD(z) = 1.7, bELGD(z) = 0.84, bQSOD(z) = 1.2, where D(z) is the growth factor. We also assume a 4\u2009per\u2009cent error in \u03b2 within \u0394z = 0.1 bins. Note that Advanced ACTPol's survey area overlaps with only \u223c75\u2009per\u2009cent of DESI's area; we take this into account in our DESI forecasts. For Euclid and WFIRST ELGs, we assume b(z) = 0.9 + 0.4z, a fit (Takada et al. 2014) to semi-analytic models (Orsi et al. 2010) that compares well with data. We determine the redshift distribution of Euclid H\u03b1 galaxies using the H\u03b1 luminosity function from Colbert et al. (2013) and assume a flux limit of 4\u00d7 10\u221216. This flux limit is in the middle of the range being considered, so the following Euclid forecasts can change accordingly. We also assume a 3\u2009per\u2009cent error in \u03b2 within \u0394z = 0.1 bins for Euclid and WFIRST (Amendola et al. 2013). For all subsequent forecasts, we assume EG measurements over angular scales 100 \u2264 \u2113 \u2264 500.","Citation Text":["Planck Collaboration 2006"],"Functions Text":["We predict the noise in the Planck lensing map assuming the detector sensitivity and beam sizes listed in the Planck Bluebook"],"Functions Label":["Uses"],"Citation Start End":[[490,515]],"Functions Start End":[[363,488]]}
{"Identifier":"2022MNRAS.510.4943S__Murray_&_Dermott_1999_Instance_2","Paragraph":"The gravitational potential of an eccentric companion at the quadrupole order can be decomposed as a sum over circular orbits (e.g. Storch & Lai 2013; Vick, Lai & Fuller 2017):\n(5)$$\\begin{eqnarray*}\r\nU\\left(\\boldsymbol{\\mathbf {r}}, t\\right) = \\sum \\limits _{m=-2}^2 U_{2m} \\left(\\boldsymbol{\\mathbf {r}}, t\\right) ,\r\n\\end{eqnarray*}$$(6)$$\\begin{eqnarray*}\r\nU_{2m}\\left(\\boldsymbol{\\mathbf {r}}, t\\right) &amp;=&amp; -\\frac{GM_2 W_{2m} r^2}{D(t)^3} Y_{2m}(\\theta , \\phi) e^{-imf\\!\\!\\!\\:(t)}, \\\\\r\n&amp;=&amp; -\\frac{GM_2W_{2m} r^2}{a^3}Y_{2m}\\left(\\theta , \\phi \\right) \\sum \\limits _{N = -\\infty }^\\infty \\!\\!F_{Nm}e^{-iN\\Omega t} .\r\n\\end{eqnarray*}$$Here, the coordinate system is centered on the MS star, (r, \u03b8, \u03d5) are the radial, polar, and azimuthal coordinates of $\\boldsymbol{\\mathbf {r}}$ respectively, $W_{2 \\pm 2} = \\sqrt{3\\pi \/10}$, W2 \u00b1 1 = 0, $W_{20} = -\\sqrt{\\pi \/ 5}$, D(t) is the instantaneous distance to the companion, f is the true anomaly, and Ylm denote the spherical harmonics. FNm denote the Hansen coefficients for l = 2 (also denoted $X^N_{2m}$ in Murray & Dermott 1999), which are the Fourier coefficients of the perturbing function, i.e.\n(7)$$\\begin{eqnarray*}\r\n\\frac{a^3}{D(t)^3} e^{-imf\\!\\!\\!\\:(t)} = \\sum \\limits _{N = -\\infty }^\\infty \\!\\!F_{Nm} e^{-iN\\Omega t}.\r\n\\end{eqnarray*}$$The FNm can be written explicitly as an integral over the eccentric anomaly (Murray & Dermott 1999; Storch & Lai 2013):\n(8)$$\\begin{eqnarray*}\r\nF_{Nm} = \\frac{1}{\\pi }\\int \\limits _{0}^{\\pi } \\frac{\\cos \\left[N\\left(E - e\\sin E\\right) - mf(E)\\right]}{\\left(1 - e\\cos E\\right)^2}\\,\\,\\mathrm{d}E.\r\n\\end{eqnarray*}$$By considering the effect of each summand in equation (5), the total torque on the star, energy transfer in the inertial frame, and energy transfer in the star\u2019s corotating frame (which is also the tidal heating rate) can be obtained (Storch & Lai 2013; Vick et al. 2017):\n(9)$$\\begin{eqnarray*}\r\nT = \\sum \\limits _{N = -\\infty }^\\infty F_{N2}^2 T_{\\rm circ}\\left(N\\Omega - 2\\Omega _{\\rm s}\\right),\r\n\\end{eqnarray*}$$(10)$$\\begin{eqnarray*}\r\n\\dot{E}_{\\rm in} &amp;=&amp; \\frac{1}{2}\\sum \\limits _{N = -\\infty }^\\infty \\Bigg [ \\left(\\frac{W_{20}}{W_{22}}\\right)^2 N\\Omega F_{N0}^2 T_{\\rm circ}\\left(N\\Omega \\right) \\\\\r\n&amp;&amp;+\\, N\\Omega F_{N2}^2 T_{\\rm circ}\\left(N\\Omega - 2\\Omega _{\\rm s}\\right) \\Bigg ] ,\r\n\\end{eqnarray*}$$(11)$$\\begin{eqnarray*}\r\n\\dot{E}_{\\rm rot} = \\dot{E}_{\\rm in} - \\Omega _{\\rm s} T .\r\n\\end{eqnarray*}$$Here, dots indicate time derivatives.","Citation Text":["Murray & Dermott 1999"],"Functions Text":["The FNm can be written explicitly as an integral over the eccentric anomaly","F_{Nm} = \\frac{1}{\\pi }\\int \\limits _{0}^{\\pi } \\frac{\\cos \\left[N\\left(E - e\\sin E\\right) - mf(E)\\right]}{\\left(1 - e\\cos E\\right)^2}\\,\\,\\mathrm{d}E."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1390,1411]],"Functions Start End":[[1313,1388],[1457,1607]]}
{"Identifier":"2018MNRAS.478..932H__Anjos_2001_Instance_1","Paragraph":"Although a correlation does exist, it is offset from the one-to-one line that one would expect, overestimating the number of spiral arms by approximately three. This may be due to how mass is assigned to the bulge and disc. We use photometric decompositions of Simard et al. (2011) and Mendel et al. (2014) to assign mass to the bulge and the disc. Such a model fits a classical bulge with n= 4 and an exponential disc. This may cause a systematic for two reasons. First, the photometric decomposition of galaxies may introduce a bias due to image resolution effects. The second issue is the pseudo- versus classical bulge argument \u2013 the model we use assumes an inner classical spherical bulge; bulges instead may be pseudo-bulges, which may not have a spherical shape, and profile well-described by a spherical Hernquist profile (Carollo et al. 1997; Gadotti & dos Anjos 2001; Kormendy et al. 2006; Fisher & Drory 2008; Gadotti 2009). Studying bulges and discs in detail is beyond the scope of this paper. Another possibility is that the assumption that spiral arms are measured at 2Rd may not be valid \u2013 if spiral arms were instead measured closer to the inner regions of galaxies, then this offset is negated. Unfortunately, the binary nature of visual morphological classifications, where arms either are or are not recorded, prevents further investigation of this point. Finally, there may be some spiral arms which are impossible to observe with visual morphology in the way presented in this paper. Of particular note is the case where the model predicts very high spiral arm numbers. In this case, the spiral arms may instead be wakelets which are difficult to observe visually; our observed arm number measurements may therefore be systematically low for these galaxies. Investigating which caveat, or which combination of caveats is responsible requires higher resolution imaging of galaxies than those used in this paper. Any study of this nature would be severely restricted in terms of sample size and completeness compared to the results we present in this paper.","Citation Text":["Gadotti & dos Anjos 2001"],"Functions Text":["The second issue is the pseudo- versus classical bulge argument \u2013 the model we use assumes an inner classical spherical bulge; bulges instead may be pseudo-bulges, which may not have a spherical shape, and profile well-described by a spherical Hernquist profile"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[852,876]],"Functions Start End":[[568,829]]}
{"Identifier":"2015ApJ...806..152S__Ferraro_et_al._2001_Instance_1","Paragraph":"One of the most astonishing characteristics of Liller 1 is the extremely large value of the collision rate parameter. Verbunt & Hut (1987) showed that Liller 1 has the second-highest value of stellar encounter rate (after Terzan 5; see also Lanzoni et al. 2010) among all star clusters in the Galaxy, thus suggesting that it represents an ideal environment where exotic objects, generated by collisions, can form. In fact, it is commonly believed that dynamical interactions in GCs facilitate the formation of close binary systems and exotic objects such as cataclysmic variables (CVs), low-mass X-ray binaries (LMXBs), millisecond pulsars (MSPs), and blue straggler stars (BSSs) (e.g., Bailyn 1992; Paresce et al. 1992; Ferraro et al. 2001, 2009a, 2012; Ransom et al. 2005; Pooley & Hut 2006). Moreover, Hui et al. (2010) found that clusters with large collisional parameters and high metallicity (see also Bellazzini et al. 1995) usually host more MSPs. Indeed, Terzan 5 hosts the largest population of MSPs among all Galactic GCs (Ransom et al. 2005).\n6\n\n\n\n6\n\nNote that Terzan 5 is suspected to not be a genuine GC, because it harbors at least three stellar populations with different iron abundances (Ferraro et al. 2009a; Origlia et al. 2011, 2013; Massari et al. 2014).\n A strong \u03b3-ray emission has been recently detected in the direction of Liller 1 by the Large Area Telescope (LAT) on board Fermi (Tam et al. 2011). This is the most intense emission detected so far from a Galactic GC, again suggesting the presence of a large number of MSPs. However, no direct radio detection of these objects has been obtained so far in this system (Ransom et al. 2005). The only exotic object identified in the cluster is the rapid burster MXB 1730-335, an LMXB observed to emit radio waves and type I and type II X-ray bursts (Hoffman et al. 1978). It seems to be located in the central region of Liller 1, but no optical\/IR counterpart of this object has been found so far (Homer et al. 2001).","Citation Text":["Ferraro et al. 2001"],"Functions Text":["In fact, it is commonly believed that dynamical interactions in GCs facilitate the formation of close binary systems and exotic objects such as cataclysmic variables (CVs), low-mass X-ray binaries (LMXBs), millisecond pulsars (MSPs), and blue straggler stars (BSSs) (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[721,740]],"Functions Start End":[[414,686]]}
{"Identifier":"2022MNRAS.511.1121M__Reig_&_Nespoli_2013_Instance_1","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"],"Functions Text":["The luminosity during the 2020 giant outburst reached a record high, which was significantly higher than the critical luminosity"],"Functions Label":["Background"],"Citation Start End":[[611,630]],"Functions Start End":[[481,609]]}
{"Identifier":"2022MNRAS.515.2188H__Rorai_et_al._2018_Instance_4","Paragraph":"In this work, we follow the method for measuring the IGM thermal state based on Voigt profile decomposition of the Ly\u2009\u03b1 forest (Schaye et al. 1999; Ricotti et al. 2000; McDonald et al. 2001). In this approach, a transmission spectrum is treated as a superposition of multiple discrete Voigt profiles, with each line described by three parameters: redshift zabs, Doppler broadening b, and neutral hydrogen column density $N_{\\rm H\\, {\\small I}}$. By studying the statistical properties of these parameters, i.e. the $b-N_{{{{\\rm H\\, {\\small I}}}}{}}$ distribution, one can recover the thermal information encoded in the absorption profiles. The majority of past applications of this method constrained the IGM thermal state by fitting the low-b-$N_{\\rm H\\, {\\small I}}$ cutoff of the $b-N_{{{{\\rm H\\, {\\small I}}}}{}}$ distribution (Schaye et al. 1999, 2000; Ricotti et al. 2000; McDonald et al. 2001; Rudie, Steidel & Pettini 2012; Boera et al. 2014; Bolton et al. 2014; Garzilli, Theuns & Schaye 2015, 2020; Hiss et al. 2018; Rorai et al. 2018). The motivation for this approach is that the Ly\u2009\u03b1 lines are broadened by both thermal motion and non-thermal broadening resulting from combinations of Hubble flow, peculiar velocities, and turbulence. By isolating the narrow lines in the Ly\u2009\u03b1 forest that constitutes the lower cutoff in $b-N_{{{{\\rm H\\, {\\small I}}}}{}}$ distributions, of which the line-of-sight component of non-thermal broadening is expected to be zero, the broadening should be purely thermal, thus allowing one to constrain the IGM thermal state. However, this method has three crucial drawbacks. First, the IGM thermal state actually impacts all the lines besides just the narrowest lines. Therefore, by restricting attention to data in the distribution outskirts, this approach throws away information and reduces the sensitivity to the IGM thermal state significantly (Rorai et al. 2018; Hiss et al. 2019). Secondly, in practice, determining the location of the cutoff is vulnerable to systematic effects, such as contamination from the narrow metal lines (Hiss et al. 2018; Rorai et al. 2018). Lastly, the results from this approach critically depend on the choice of low-b cutoff fitting techniques, where different techniques might result in inconsistent T0 and \u03b3 measurements (Hiss et al. 2018; Rorai et al. 2018).","Citation Text":["Rorai et al. 2018"],"Functions Text":["Lastly, the results from this approach critically depend on the choice of low-b cutoff fitting techniques, where different techniques might result in inconsistent T0 and \u03b3 measurements"],"Functions Label":["Background"],"Citation Start End":[[2321,2338]],"Functions Start End":[[2117,2301]]}
{"Identifier":"2017ApJ...838...67E__Derekas_et_al._2017_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":["Derekas et al. 2017"],"Functions Text":["Cepheids have also been found to show additional complications that include cycle-to-cycle variations in their light and radial velocity curves (see"],"Functions Label":["Motivation"],"Citation Start End":[[890,909]],"Functions Start End":[[658,806]]}
{"Identifier":"2018MNRAS.474.3280F__Sotomayor-Beltran_et_al._2013_Instance_1","Paragraph":"The ability to retrieve the polarized quantities of a radio source is entirely dependent on the ability to find radio sources within noisy images. It is of importance to planned future surveys to investigate suitable strategies for source-finding in linear polarization with interferometers such as LOFAR, which are very well suited for deep radio surveys (e.g. Hardcastle et al. 2016; Clarke et al. 2017), but are particularly technically challenging due to operation at low radio frequencies (Varenius et al. 2015), with potentially sub-arcsecond angular resolution (Mold\u00f3n et al. 2015), at high sensitivity (Shimwell et al. 2016), and with the ability to make precise Faraday rotation measurements (Sotomayor-Beltran et al. 2013). Nevertheless, finding linearly polarized sources faces many hurdles: (i) at sub-arcminute resolution, the peak in linearly polarized intensity can be offset from the peak in total intensity (see e.g. fig. 1 in O'Sullivan et al. 2015), (ii) the statistics in polarized intensity, $P=\\sqrt{Q^2+U^2}$, are Rician, rather than Gaussian, while all publicly available source-finders are geared towards Gaussian noise statistics, (iii) the full sensitivity is not provided in any single channel of Q, U, or P, and RM Synthesis is therefore required to retrieve the full point-source sensitivity from the data, (iv) sources detected in Q and U can have both positive and negative brightness, and these values oscillate and mix across the observing bandwidth due to Faraday rotation, and (v) in some cases, Q and U images can be more sensitive than I images, which in principle could lead to sources that can be found in P but not in I. Source-finding in circular polarization, Stokes V, is beyond the scope of this paper in which we focus on linear polarization, but also faces similar challenges due to the full sensitivity not being provided in a single channel and the process of Faraday conversion across the observing band. Moreover, the linear feeds used for observations at low radio frequencies with instruments such as LOFAR are more suitable for measuring circular rather than linear polarization, which further increases the difficulty of detecting faint linearly polarized sources.","Citation Text":["Sotomayor-Beltran et al. 2013"],"Functions Text":["It is of importance to planned future surveys to investigate suitable strategies for source-finding in linear polarization with interferometers such as LOFAR,","and with the ability to make precise Faraday rotation measurements"],"Functions Label":["Background","Background"],"Citation Start End":[[702,731]],"Functions Start End":[[147,305],[634,700]]}
{"Identifier":"2017ApJ...837...88B__High_et_al._2010_Instance_1","Paragraph":"With photometry in hand for the majority of our galaxy spectroscopy sample, we can also investigate velocity segregation effects as a function of galaxy luminosity. We use cluster galaxy brightness measurements in units relative to m*, a standard quantity that can be easily incorporated into both observational data and simulated clusters (Cole et al. 2001; Cohen 2002; Rudnick et al. 2006, 2009). Specifically, we use m* values computed from Bruzual & Charlot (2003) models in the same fashion as described in previous SPT publications (High et al. 2010; Song et al. 2012; Bleem et al. 2015). Figure 7 shows the distribution of cluster member galaxies with brightness measurements as described in Section 2.4 plotted relative to the characteristic magnitude, m*, for the full cluster member sample as well as the subsamples of cluster members of different spectral types. We use these data, in combination with the normalized peculiar velocities of each cluster member in the ensemble to plot the expectation value for the absolute peculiar velocity of cluster members as a function of brightness for all galaxies together, as well as for each of the passive, post-starburst, and star-forming galaxy subsets (Figure 8). It is clear that the brightest galaxies\u2014independent of spectral type\u2014universally prefer smaller absolute peculiar velocities, and that when we treat all galaxies together we see a strong drop in the absolute peculiar velocities of galaxies brighter than \n\n\n\n\n\n, while galaxies fainter than this tend to remain approximately flat in absolute peculiar velocity as a function of brightness. There is no statistically significant evidence in our data for an evolution in the presence or of velocity luminosity segregation with redshift. Specifically, if we split our sample into two redshift bins at z = 0.45\u2014which optimally balances the number of bright galaxies in the high and low bins\u2014we see the strong drop in peculiar velocity for bright galaxies in each of the high and low-redshift bins.","Citation Text":["High et al. 2010"],"Functions Text":["Specifically, we use m* values computed from Bruzual & Charlot (2003) models in the same fashion as described in previous SPT publications"],"Functions Label":["Similarities"],"Citation Start End":[[539,555]],"Functions Start End":[[399,537]]}
{"Identifier":"2017AandA...602A.106B__Ehrenreich_et_al._2015_Instance_1","Paragraph":"About a quarter of the known exoplanets orbit at short distances (\u22720.1\u2009au) from their star (from the Exoplanet Encyclopaedia in December 2016; Schneider et al. 2011). Heating by the stellar energy can lead to the expansion of their upper atmospheric layers and their eventual escape. As a result of this expansion, the upper atmosphere produces a deeper absorption than the planetary disk alone when observed in the UV, in particular in the stellar Lyman-\u03b1 (Ly-\u03b1) line of neutral hydrogen (e.g., Vidal-Madjar et al. 2003; Lecavelier des Etangs et al. 2012). Super-Earths and smaller planets display a large diversity in nature and composition (e.g., Seager et al. 2007; Rogers & Seager 2010; Fortney et al. 2013), which can only be investigated through observations of their atmospheres. The reduced scale height of the lower atmospheric layers makes them difficult to probe in the visible and the infrared. In contrast, very deep UV transit signatures can be produced by the upper atmospheres of small planets. The warm Neptune GJ436b, which is the lowest mass planet found evaporating to date (Ehrenreich et al. 2015), shows transit absorption depths up to 60% in the Ly-\u03b1 line. The formation of such an extended exosphere is due in great part to the low mass of GJ436b and the gentle irradiation from its M-dwarf host (Bourrier et al. 2015, 2016b). Few attempts have been made to detect atmospheric escape from Earth-size planets, and Ly-\u03b1 transit observations of the super-Earth 55 Cnc e (Ehrenreich et al. 2012) and HD\u200997658 b (Bourrier et al. 2016a) showed no evidence of hydrogen exospheres. In the case of 55 Cnc e, this non-detection hinted at the presence of a high-weight atmosphere \u2013 or the absence of an atmosphere \u2013 recently supported by the study of its brightness map in the IR (Demory et al. 2016). Understanding the conditions that can lead to the evaporation of Earth-size planets will be necessary to determine the stability of their atmospheres and their possible habitability. For example, large amounts of hydrogen in the upper atmosphere of a close-in terrestrial planet could indicate the presence of a steam envelope being photodissociated, and replenished by evaporating water oceans (Jura 2004; L\u00e9ger et al. 2004). ","Citation Text":["Ehrenreich et al. 2015"],"Functions Text":["The warm Neptune GJ436b, which is the lowest mass planet found evaporating to date","), shows transit absorption depths up to 60% in the Ly-\u03b1 line."],"Functions Label":["Background","Background"],"Citation Start End":[[1096,1118]],"Functions Start End":[[1012,1094],[1118,1180]]}
{"Identifier":"2015MNRAS.451.2663H___2005_Instance_1","Paragraph":"Galaxy formation theory has developed dramatically over the last three decades. \u039b cold dark matter (\u039bCDM) has been established as the standard model for cosmological structure formation, and its parameters have been increasingly tightly constrained by observations. In parallel, simulations of galaxy formation within this standard model have grown in complexity in order to treat more accurately the many baryonic processes that impact the evolution of the galaxy population. Semi-analytic modelling is a particular simulation method which is optimized to connect the observed properties of the galaxy population \u2013 abundances, scaling relations, clustering and their evolution with redshift \u2013 to the astrophysical processes that drive the formation and evolution of individual galaxies (e.g. White 1989; Cole 1991; Lacey & Silk 1991; White & Frenk 1991; Kauffmann, White & Guiderdoni 1993; Cole et al. 1994; Kauffmann et al. 1999; Somerville & Primack 1999; Springel et al. 2001, 2005; Hatton et al. 2003; Kang et al. 2005; Lu et al. 2011; Benson 2012). Simple phenomenological descriptions of the relevant processes are needed, each typically involving uncertain efficiency and scaling parameters. These must be determined by comparison with observation or with more detailed simulations. As the range and quality of observational data have increased, so has the number of processes that must be included to model them adequately, and hence the number of adjustable parameters. In recent years, robust statistical methods have been introduced in order to sample the resulting high-dimensional parameter spaces and to determine the regions that are consistent with specific observational data sets. This development began with the work of Kampakoglou, Trotta & Silk (2008) and Henriques et al. (2009) and has since been extended to a wide range of models and sampling methods (Benson & Bower 2010; Bower et al. 2010; Henriques & Thomas 2010; Lu et al. 2011, 2012; Henriques et al. 2013; Mutch, Poole & Croton 2013; Benson 2014; Ruiz et al. 2015).","Citation Text":["Springel et al.","2005"],"Functions Text":["Semi-analytic modelling is a particular simulation method which is optimized to connect the observed properties of the galaxy population \u2013 abundances, scaling relations, clustering and their evolution with redshift \u2013 to the astrophysical processes that drive the formation and evolution of individual galaxies (e.g."],"Functions Label":["Background"],"Citation Start End":[[959,974],[981,985]],"Functions Start End":[[477,792]]}
{"Identifier":"2020ApJ...897..158S__Sherwood_et_al._2003_Instance_1","Paragraph":"For the study of each of the individual reactions we have characterized the stationary points on their respective potential energy surfaces by means of density functional theory (DFT) calculations. We have employed the mPWB1K exchange and correlation functional of Zhao & Truhlar (2004) in combination with the def2-TZVP basis set (Weigend & Ahlrichs 2005). The main advantage of this functional for this investigation is that it was designed for the easy determination of activation energies. Minima and transition states (TSs) were optimized using the DL-FIND (K\u00e4stner et al. 2009) program of the ChemShell suite (Sherwood et al. 2003; Metz et al. 2014). A search of TS was done by means of potential energy surface scans and posterior optimization in cases where a kinetic barrier was predicted. Assessment of the nature of the stationary points was performed by computing the molecular Hessian of each. Additionally, we have addressed the validity of the DFT method by computing single-point calculations on the relevant DFT geometries at the CCSD(T)-F12\/cc-PVTZ-F12 level of theory using Molpro 2015 (Werner et al. 2012, 2015). We have found excellent agreement between DFT and coupled cluster methods for the reaction barriers, with deviations of less than 1 kcal mol\u22121 in all cases. Furthermore, we have carried out intrinsic reaction coordinate (IRC) calculations to ensure a proper connection between our calculated TSs and their associated minima. For reactions with a barrier, we have computed both classical bimolecular reaction rate coefficients as well as tunneling-corrected ones. The inclusion of quantum tunneling in our calculations has been performed using semi-classical instanton theory (Rommel & K\u00e4stner 2011; Rommel et al. 2011), following a sequential cooling scheme for temperatures below the crossover temperature, and reduced instanton theory for temperatures above (McConnell & K\u00e4stner 2017). Crossover temperatures are defined\u2014with \n\n\n\n\n\n being the frequency of the vibrational imaginary mode in the TS, and kB the Boltzmann constant\u2014as\n6\n\n\n\n\n\n\n","Citation Text":["Sherwood et al. 2003"],"Functions Text":["Minima and transition states (TSs) were optimized using the DL-FIND","program of the ChemShell suite"],"Functions Label":["Uses","Uses"],"Citation Start End":[[616,636]],"Functions Start End":[[494,561],[584,614]]}
{"Identifier":"2022ApJ...925...37W__Lentati_et_al._2015_Instance_1","Paragraph":"The stochastic gravitational-wave background (SGWB)\u2014the primary goal of the search of the PTA collaborations\u2014is expected to be dominant in the nanohertz band, which might originate from supermassive black hole binaries (SMBHBs; Rajagopal & Romani 1995; Sesana 2013), comic strings (Damour & Vilenkin 2005; Blanco-Pillado et al. 2018), the first phase transition (Caprini et al. 2010), and scalar-induced GWs (Yuan et al. 2019). Over the last few decades, the PTA collaborations have not found GW signals, but the increasingly sensitive data sets offer increasingly stringent constraints on the SGWB (van Haasteren et al. 2011; Shannon et al. 2013; Lentati et al. 2015; Shannon et al. 2015; Arzoumanian et al. 2016, 2018; Chen et al. 2020). Recently, the NANOGrav collaboration reported that there is strong evidence in favor of a stochastic common-spectrum process, which is modeled by a power-law spectrum among the pulsars, over the independent red noise processes of each pulsar, in their 12.5 yr data set (Arzoumanian et al. 2020). However, given the lack of statistically significant evidence for quadrupolar spatial correlations, it is inconclusive to claim the detection of an SGWB consistent with general relativity. Note that the tensor transverse (TT) modes giving rise to quadrupolar spatial correlations constitute only two of the six GW polarization modes that are allowed in a general metric theory of gravity, which also includes one scalar transverse (ST) mode, two vectorial longitudinal (VL) modes, and one scalar longitudinal (SL) mode. Later on, Chen et al. (2021) searched for nontensorial SGWBs in the NANOGrav 12.5 yr data set, and found strong Bayesian evidence that the common-spectrum reported by the NANOGrav collaboration had ST spatial correlations. More recently, the PPTA collaboration has also found a common-spectrum process in their second data release (DR2), with PPTA DR2 showing no significant evidence for, or against, TT spatial correlations (Goncharov et al. 2021a).","Citation Text":["Lentati et al. 2015"],"Functions Text":["Over the last few decades, the PTA collaborations have not found GW signals, but the increasingly sensitive data sets offer increasingly stringent constraints on the SGWB"],"Functions Label":["Background"],"Citation Start End":[[648,667]],"Functions Start End":[[428,598]]}
{"Identifier":"2022ApJ...928L..16Y__Yang_&_Zhang_2018_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":["Yang & Zhang 2018"],"Functions Text":["which suggests that at least some FRBs originate from magnetars born from the core collapse of massive stars (e.g.,"],"Functions Label":["Motivation"],"Citation Start End":[[745,762]],"Functions Start End":[[539,654]]}
{"Identifier":"2019MNRAS.482.3550K__Weisz_et_al._2013_Instance_1","Paragraph":"The simplest approach to handling the problem of completeness is to be extremely conservative, and discard all data in regions of parameter space where the observations are not complete or nearly so. However, this invariably requires one to discard much of the available data. A somewhat more sophisticated approach is forward modelling: rather than deriving the mass and age distribution of the population from estimates of mass age for individual clusters, one could instead consider a proposed distribution of masses and ages, predict the resulting photometry distribution including the effects of incompleteness, and adjust parameters of the mass and age distribution until they match the observations. Approaches of this type are widely used in astronomy, for example to infer star formation histories or stellar mass distributions from observed colour\u2013magnitude diagrams (CMDs; e.g. Dolphin 2002; Harris & Zaritsky 2009; Weisz et al. 2013; Conroy & van Dokkum 2016; see Cervi\u00f1o 2013 for a review). However, methods of this type have not previously been applied to deriving the properties of populations of star clusters, at least in part due a unique challenge not present in other applications. In existing applications such as CMD fitting, the forward model is deterministic, i.e. for a given stellar mass, age, and other properties, there is a single predicted colour and magnitude. This is not the case for star clusters with masses \u2272 3000 M\u2299, because such clusters are too small to fully sample the stellar initial mass function (IMF, e.g. Cervi\u00f1o & Luridiana 2004, 2006; da Silva, Fumagalli & Krumholz 2012). As a result, two clusters of the same total mass and age can produce wildly different luminosities and colours. This means that the forward model is not deterministic, but instead depends on an additional random variable that couples non-linearly with the deterministic variables like cluster mass and age. This situation presents computational challenges that are not addressed by existing methods.","Citation Text":["Weisz et al. 2013"],"Functions Text":["Approaches of this type are widely used in astronomy, for example to infer star formation histories or stellar mass distributions from observed colour\u2013magnitude diagrams (CMDs; e.g."],"Functions Label":["Background"],"Citation Start End":[[927,944]],"Functions Start End":[[707,888]]}
{"Identifier":"2019ApJ...871..243Y__Zhou_et_al._1993_Instance_1","Paragraph":"B335 is an isolated Bok globule with an embedded Class 0 protostar at a distance of 100 pc (Keene et al. 1980, 1983; Stutz et al. 2008; Olofsson & Olofsson 2009). The size of the dense core in B335 observed at millimeter wavelengths is \u223c0.1 pc (Saito et al. 1999; Motte & Andr\u00e9 2001; Shirley et al. 2002), and the core is slowly rotating (Saito et al. 1999; Yen et al. 2011; Kurono et al. 2013). Infalling and rotational motions on scales from 100 to 3000 au have been observed in molecular lines with single-dish telescopes and interferometers (Zhou et al. 1993; Choi et al. 1995; Zhou 1995; Saito et al. 1999; Evans et al. 2005, 2015; Yen et al. 2010, 2011, 2015b; Kurono et al. 2013). Nevertheless, no sign of Keplerian rotation was observed with the Atacama Large Millimeter\/submillimeter Array (ALMA) at an angular resolution of 03 (30 au; Yen et al. 2015b), and the envelope rotation on a scale of 100\u20131000 au in B335 is an order of magnitude slower than in other Class 0 and I protostars surrounded by a Keplerian disk with a size of tens of au (Yen et al. 2015a). The presence of a small disk less than 10 au and the slow envelope rotation hints at the effects of the magnetic field on the gas kinematics in B335. In addition, ALMA observations in the C18O and H13CO+ lines show no detectable difference in the infalling velocities of neutral and ionized gas on a 100 au scale with a constraint on the upper limit of the ambipolar drift velocity of 0.3 km s\u22121, suggesting that the magnetic field likely remains well coupled with the matter in the inner envelope in B335 (Yen et al. 2018). The magnetic field structures on a 1000 au scale in B335 also show signs of being dragged toward the center and becoming pinched, as inferred from the ALMA polarimetric observations (Maury et al. 2018). Therefore, B335 is an excellent target to investigate the interplay between the magnetic field and gas motions and the effects of the magnetic field on the dynamics in collapsing dense cores.","Citation Text":["Zhou et al. 1993"],"Functions Text":["Infalling and rotational motions on scales from 100 to 3000 au have been observed in molecular lines with single-dish telescopes and interferometers"],"Functions Label":["Background"],"Citation Start End":[[546,562]],"Functions Start End":[[396,544]]}
{"Identifier":"2016MNRAS.455..449H__Angus_et_al._2012_Instance_1","Paragraph":"The EFE mentioned in the previous section is due to the fact that the MOND equations (3) and (4) are non-linear and involve the total gravitational acceleration with respect to a pre-defined frame (e.g. the CMB frame). Decomposing the total gravitational field \u2207\u03a6 into an internal part $\\boldsymbol g$ and an external field $\\boldsymbol g_{\\rm e}$ and using a similar decomposition for the Newtonian gravitational acceleration ($\\nabla \\boldsymbol \\Phi _{\\rm N}=\\boldsymbol g_{\\rm N}+\\boldsymbol g_{{\\rm Ne}}$) allows us to solve the equations by taking into account the external field. This must typically be done with a numerical Poisson solver (Wu et al. 2008; Angus et al. 2012; L\u00fcghausen, Famaey & Kroupa 2015). Nevertheless, fits to rotation curves in MOND usually neglect the small corrections due to the non-spherical symmetry of the problem, in order to allow for a direct fit of the rotation curve. In the same spirit, and in order to get a first glimpse of the influence of the EFE on rotation curves, we generalize the one-dimensional solution, by using the following formula to fit rotation curves, namely equation (60) from Famaey & McGaugh (2012):\n\n(6)\n\n\\begin{equation}\n\\boldsymbol g=\\nu \\left(\\frac{|\\boldsymbol g_{\\rm N}+\\boldsymbol g_{{\\rm Ne}}|}{a_0}\\right)\\left(\\boldsymbol g_{\\rm N} + \\boldsymbol g_{{\\rm Ne}}\\right)-\\nu \\left(\\frac{g_{{\\rm Ne}}}{a_0}\\right)\\boldsymbol g_{{\\rm Ne}}\\,.\n\\end{equation}\n\nThe 1D version of this formula has been shown to be a good approximation of the true 3D solution from a numerical Poisson solver for a random orientation of the external field, at least for computing the Galactic escape speed (Famaey et al. 2007; Wu et al. 2008). Further work should investigate the range of variation of the actual rotation curve compared to the one obtained in this way, for full numerical solutions of the modified Poisson equation and various orientations of the EFE. As mentioned in Famaey & McGaugh (2012), the EFE is negligible if ge g but can play a significant role when the gravitational field g \u223c ge a0. This condition is always reached at some point in the external part of the galaxies. In this case, the relation (6) shows that the EFE will induce a decrease in the internal gravitational field. In other words, the EFE can lead to a decrease of the external part of the rotation curves. We will study this effect more carefully in Section 3.","Citation Text":["Angus et al. 2012"],"Functions Text":["This must typically be done with a numerical Poisson solver"],"Functions Label":["Uses"],"Citation Start End":[[664,681]],"Functions Start End":[[587,646]]}
{"Identifier":"2021AandA...651A..24H__Simmons_&_Stewart_(1985)_Instance_1","Paragraph":"The two main problems arising when dealing with P are the typically low signal-to-noise ratio of the polarisation signal coming from ERS and the non-Gaussian distribution of its noise statistics. Regarding the former, as mentioned above, the typical polarisation fractions of ERS at frequencies below \u223c10 GHz are at most 10%. This means that only a few ERS are bright enough to be detected in polarisation with present-day technology. A standard procedure to avoid false detections in polarisation is to detect sources in total intensity and then to try to estimate their polarisation properties in a non-blind way3. We follow this approach in this paper. Regarding the latter problem, assuming that the Q and U noises are Gaussian-distributed, P would have a non-Gaussian Rice distribution (Rice 1945). Rician distribution has strictly non-negative support and heavy tails, which firstly biases the estimation of the polarisation of the sources and secondly disrupts the intuitive interpretation of signal-to-noise in terms of \u03c3 thresholds which is used virtually everywhere else in radio astronomy. Simmons & Stewart (1985) discussed four estimators which attempted to correct for biasing in the degree of linear polarisation in the presence of low signal-to-noise ratios. More recently, Arg\u00fceso et al. (2009) studied the problem in the context of CMB astronomy and developed two methods for the detection and estimation of ERS in polarisation data: one that applies the Neyman-Pearson lemma to the Rice distribution, the Neyman-Pearson filter (NPF), and another based on pre-filtering before fusion of Q and U to obtain P, the filtered fusion (FF) method. That work found that under typical CMB-experiment settings, the FF outperforms the NPF both in terms of computational simplicity and accuracy, especially for low fluxes. L\u00f3pez-Caniego et al. (2009) applied the FF to the Wilkinson Microwave Anisotropy Probe (WMAP) five-year data. The same method has been used to study the polarisation of the Planck Second Catalogue of Compact Sources (PCCS2, Planck Collaboration XXVI 2016) and of the QUIJOTE experiment wide survey source catalogue (Herranz et al. 2021). Alternatively, a novel method for the estimation of the polarisation intensity and angle of compact sources in the E and B modes of polarisation based on steerable wavelets has been recently proposed by Diego-Palazuelos et al. (2021).","Citation Text":["Simmons & Stewart (1985)"],"Functions Text":["discussed four estimators which attempted to correct for biasing in the degree of linear polarisation in the presence of low signal-to-noise ratios."],"Functions Label":["Background"],"Citation Start End":[[1101,1125]],"Functions Start End":[[1126,1274]]}
{"Identifier":"2018ApJ...866...20D__Thompson_et_al._2012_Instance_1","Paragraph":"The physical processes involved in the formation of massive O-type stars and their feedback mechanisms are still under debate (Zinnecker & Yorke 2007; Tan et al. 2014). The energetics of O-type stars can affect the origin of new low-mass and massive stars (Deharveng et al. 2010). The massive stars are often surrounded by the bubbles\/rings\/semi-ringlike structures traced at mid-infrared (MIR) 8.0 \u03bcm (Churchwell et al. 2006, 2007) and are also associated with the extended radio continuum emission (e.g., Deharveng et al. 2010). Note that the majority of the studies related to the MIR bubbles are mainly carried out for a single H ii region or several H ii regions on scales of a few parsecs (e.g., Zinnecker & Yorke 2007; Deharveng et al. 2010; Rathborne et al. 2011; Kendrew et al. 2012; Simpson et al. 2012; Tackenberg et al. 2012; Thompson et al. 2012; Tan et al. 2014; Dewangan et al. 2015a, 2015b; Xu et al. 2016a). However, to our knowledge, in the Milky Way, there is still a limited detailed multiwavelength study of large-scale systems (>25 pc) of several MIR bubbles\/H ii regions containing O-type stars, and hence the origin of such extended systems of H ii regions remains unexplored. These systems could be candidates for \u201cmini-starburst\u201d (such as the W43 \u201cmini-starburst\u201d region; Motte et al. 2003). With the availability of the radio recombination line (RRL) and continuum observations (e.g., Lockman 1989; Condon et al. 1998; Anderson & Bania 2009; Jones & Dickey 2012), the MIR survey (e.g., Benjamin et al. 2003), the dust continuum survey at 870 \u03bcm (e.g., Schuller et al. 2009; Urquhart et al. 2018), and the 13CO line survey (e.g., Jackson et al. 2006; Anderson et al. 2009), it appears that the H ii regions located toward the Galactic plane and the inner Galaxy are the promising sites to investigate the extended systems of O-type stars. Such study will enable us to understand the physical conditions in a densely clustered environment linked with the luminous giant H ii regions\/massive star-forming complexes\/mini-starburst candidates in the Galaxy. However, in particular, in the direction of the inner Galaxy, the investigation of an extended system of H ii regions is often restricted by the near\u2013far kinematic distance ambiguity (e.g., Anderson & Bania 2009; Jones & Dickey 2012; Urquhart et al. 2018). In recent years, significant effort has been devoted to resolving the distance ambiguity for H ii regions in the inner Galaxy (see Urquhart et al. 2018 and references therein). In this work, we aim to observationally investigate a large-scale system\/configuration of several H ii regions powered by O-type stars and the origin of such a large system.","Citation Text":["Thompson et al. 2012"],"Functions Text":["Note that the majority of the studies related to the MIR bubbles are mainly carried out for a single H ii region or several H ii regions on scales of a few parsecs (e.g.,"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[838,858]],"Functions Start End":[[531,701]]}
{"Identifier":"2018MNRAS.475.4011B__Burrows_et_al._2011_Instance_1","Paragraph":"ASASSN-14li was discovered by the All Sky Automated Search for Supernova (ASASSN; Shappee et al. 2014) on ut 2014\u201311\u201322.63 (MJD 56983.6) as a 16.5 magnitude source in the V band (Jose et al. 2014; Holoien et al. 2016; Brown et al. 2017a). The position of the source was found to be consistent with the centre of the post-starburst galaxy PGC 043234, with a measured projected separation of 0.04 arcsec. This galaxy is at redshift z = 0.0206 with a luminosity distance of 90.3\u2009Mpc (for cosmological parameters H0 = 73\u2009kms\u22121\u2009Mpc\u22121, \u03a9matter = 0.27, and \u03a9\u039b = 0.73). It was established through archival X-ray observations of PGC 043234 from the ROSAT All-Sky Survey (Voges et al. 1999) that the galaxy does not contain an efficiently accreting AGN, with the count rate implying a luminosity orders of magnitude below standard active nuclei (e.g. Miller et al. 2015). A small number (currently six) of confirmed TDEs, including ASASSN-14li, have also been detected at radio wavelengths and the population may form a bi-modal distribution, consisting of more common non-relativistic \u2018thermal\u2019 events and rarer relativistic jets. Three events (Swift J1644+57; Burrows et al. 2011; Zauderer et al. 2011, Swift J2058+05; Cenko et al. 2012, Swift J1112.2; Brown et al. 2017b) have isotropic \u223c5\u2009GHz luminosities of between 1040 and 1042\u2009erg\u2009s\u22121 whereas the rest (IGR J12580+0134; Irwin et al. 2015, XMMSL1 J0740-85; Alexander et al. 2017, ASASSN-14li; van Velzen et al. 2016; Alexander et al. 2016) have luminosities in the range 1037 to 1039\u2009erg\u2009s\u22121 at similar frequencies. The higher power events are believed to result from observing down the axis of a relativistic jet, resulting in the energy of photons being significantly boosted. Even accounting for boosting, these relativistic events have a higher total energy output than their thermal counterparts. The origin of the radio emission from the thermal events is currently uncertain, with transient jets (van Velzen et al. 2016), non-relativistic winds (Alexander et al. 2016), and shocks driven by unbound material (Krolik et al. 2016) all feasible scenarios. ASASSN-14li is by far the best studied of the \u2018thermal\u2019 TDE category, having been observed extensively at Optical, UV, X-ray (where ASASSN-14li is unusually loud for an optically selected TDE), and radio wavelengths. The high cadence X-ray and radio observations in particular allow for the X-ray\/radio coupling to be probed.","Citation Text":["Burrows et al. 2011"],"Functions Text":["Three events (Swift J1644+57;","have isotropic \u223c5\u2009GHz luminosities of between 1040 and 1042\u2009erg\u2009s\u22121 whereas the rest","have luminosities in the range 1037 to 1039\u2009erg\u2009s\u22121 at similar frequencies. The higher power events are believed to result from observing down the axis of a relativistic jet, resulting in the energy of photons being significantly boosted. Even accounting for boosting, these relativistic events have a higher total energy output than their thermal counterparts."],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1152,1171]],"Functions Start End":[[1122,1151],[1265,1349],[1487,1848]]}
{"Identifier":"2019AandA...627A.135B__Bessell_&_Brett_(1988)_Instance_1","Paragraph":"Optical color-magnitude and near-infrared (NIR) color-color diagrams are used to classify our variable candidates. The 2MASS JHKs data are available for 93 stars, while for the remaining two stars, photometric data are adopted from the UKIDSS Galactic Plane Survey (Lucas et al. 2008). The 2MASS photometry is transformed to the California Institute of Technology (CIT) system using the relations provided on their website3 to compare with the evolutionary models. Figure 6 represents the J\u2005\u2212\u2005H\/H\u2005\u2212\u2005K CCD based on 2MASS data, typically used to classify the YSOs. The YSOs from Rebull et al. (2011) and Ogura et al. (2002) are overplotted in colored symbols. The sequence of dwarf and giants from Bessell & Brett (1988), and the intrinsic locus of CTT stars (Meyer et al. 1997) are also overplotted. The three parallel lines are the reddening vectors drawn from the tip of the giant branch (left), from the base of the MS branch (middle), and from the tip of the intrinsic CTTSs line (right). The extinction ratios to derive these reddening vectors are \n\n\n\n\n\nA\n\nJ\n\/\nH\n\/\nK\n\n\n\nA\nV\n\n\n=\n0.265\n\/\n0.155\n\/\n0.090\n\n\n$ \\frac{A_{J\/H\/K}}{A_V} = 0.265\/0.155\/0.090 $\n\n\n, adopted from Cohen et al. (1981). In general, CTTSs with smaller NIR excess, WTTSs, and field stars (MS and giants) occupy the region between the left and middle reddening vectors. Figure 6 shows that most of the variables that are outliers in the proper motions and lie below the intrinsic CTTSs locus are the MS stars. Two variables (V173 and V177) are members based on their proper motions, but fall below the giant sequence. One of these, V173, does not have kinematic information. The CTTSs with large infrared excess are located in the region between the middle and right reddening vectors, while more moderate CTTSs with smaller infrared excess can also populate the region between left and middle reddening vectors, mixed with reddened WTTSs just above the CTT locus. Some contamination is expected depending on the reddening and IR excess, and also due to variability in single-epoch measurements.","Citation Text":["Bessell & Brett (1988)"],"Functions Text":["The sequence of dwarf and giants from","are also overplotted. The three parallel lines are the reddening vectors drawn from the tip of the giant branch (left), from the base of the MS branch (middle)"],"Functions Label":["Uses","Uses"],"Citation Start End":[[696,718]],"Functions Start End":[[658,695],[777,936]]}
{"Identifier":"2021MNRAS.502.2859N__Evans_&_Howarth_2008_Instance_1","Paragraph":"It is harder to evaluate the behaviour of the young stellar population along the line of sight, since radial velocity measurements for our sample of Cepheids, needed for a thorough study, do not exist. Given this deficit, we provide only a simplified estimate using radial velocities of OBA-type stars from Evans & Howarth (2008). Since they belong to the same young population, we assume that they have a similar distance distribution and kinematics as the Cepheids. Fig. 16 shows the massive star sample in the plane of the sky. Except for the northernmost region (\u03b4 \u2265 \u221272\u25cb), where no data exist, these stars cover a comparable area to the Cepheids (indicated as grey dots in the figure for comparison). The radial velocities show a distinct and well-known gradient across the SMC with higher velocities in the eastern part (see also fig. 5 of Evans & Howarth 2008). Such a gradient in radial velocity is also present in older (few\u2009Gyr) RGB stars (see fig. 9 of Dobbie et al. 2014). This gradient is commonly attributed to rotation of the SMC. Based on our results obtained for the Cepheids, we propose a different interpretation: this line-of-sight velocity gradient may instead be caused by the fact that the nearest parts of the galaxy, in the region of the SMC Wing, move with a higher radial velocity compared with the main body of the galaxy. Given the additional differences in tangential velocities, these outer parts might be in the process of being stripped from the SMC. Diaz & Bekki (2012) show in their simulations that tidal effects can produce a velocity gradient that is similar to that of a rotating disc. We stress again that this interpretation is based on the assumption that the Cepheid sample and the OBA-type stellar sample trace a similar three-dimensional distribution. For any conclusive answer, radial velocities of the Cepheid stars are required. Such measurements will be provided by the One Thousand and One Magellanic Fields (1001MC) survey (Cioni et al. 2019), which is a consortium survey with the forthcoming multi-object spectrograph 4MOST that will be mounted on the VISTA telescope.","Citation Text":["Evans & Howarth (2008)"],"Functions Text":["Given this deficit, we provide only a simplified estimate using radial velocities of OBA-type stars from","Since they belong to the same young population, we assume that they have a similar distance distribution and kinematics as the Cepheids."],"Functions Label":["Uses","Compare\/Contrast"],"Citation Start End":[[307,329]],"Functions Start End":[[202,306],[331,467]]}
{"Identifier":"2019ApJ...887....8H__Marty_et_al._2011_Instance_1","Paragraph":"Silicon carbide is the best-characterized presolar mineral. It was identified more than 30 yr ago (Bernatowicz et al. 1987) because it is tagged with noble gases of anomalous isotopic compositions (Lewis et al. 1994). Subsequently, it was found that the major elements C and Si, and numerous minor elements contained in presolar SiC, have highly anomalous isotopic compositions as well, the fingerprints of nucleosynthetic processes in their parent stars. Based on the isotopic compositions of C, N, and Si, SiC was divided into distinct populations (Zinner 2014). This includes the mainstream grains, which account for about 80%\u201390% of all grains (depending on grain size), and the minor types AB (originally defined as two distinct types A and B), C, X, Y, Z, and (putative) nova grains. The mainstream grains have 12C\/13C ratios between 10 and 100 (solar: 89), and the 14N\/15N ratios of most of them are higher than the solar ratio of 440, the ratio measured for the solar wind (Marty et al. 2011). In a plot of \u03b429Si versus \u03b430Si, the mainstream grains lie along a straight line defined by \u03b429Si = 1.37 \u00d7 \u03b430Si\u201320 (Zinner et al. 2007), where \u03b4xSi = [(xSi\/28Si)grain\/(xSi\/28Si)solar\u22121) \u00d7 1000, and x = 29 or 30, i.e., \u03b4xSi is the per mil deviation from the solar xSi\/28Si ratio. \u03b430Si values of mainstream grains vary between about \u221250\u2030 and +150\u2030. The isotopic compositions of heavy elements show the signatures of slow neutron-capture nucleosynthesis (s-process, K\u00e4ppeler et al. 2011), which points toward low-mass (1.5\u20133 M\u2299) asymptotic giant branch (AGB) stars of about solar or supersolar metallicity as parent stars (e.g., Lugaro et al. 2018, and references therein). The minor type Y and Z grains (a few % of all SiC grains, depending on grain size), which fall to the 30Si-rich side of the Si mainstream line, were also proposed to originate from low-mass AGB stars, but with metallicities lower than solar (Hoppe et al. 1997; Amari et al. 2001b). This low-metallicity scenario, however, was recently questioned (Liu et al. 2019). The type C (\u223c0.1% of all SiC grains) and X grains (\u223c1% of all SiC grains) are believed to originate from core-collapse supernovae (CCSNe; Amari et al. 1992; Hoppe et al. 1996b; Nittler et al. 1996; Gyngard et al. 2010). These grains show strong depletions (X grains) or enrichments (C grains) in the heavy Si isotopes. Their 12C\/13C ratios span a large range from 10 to >10,000; other characteristic features of X and C grains are enrichments in 15N and high initial 26Al\/27Al ratios of typically >0.1. Putative nova grains (\u223c0.1% of all SiC grains) have low 12C\/13C ratios of 10, low 14N\/15N ratios of 40, and high initial 26Al\/27Al ratios of up to 0.2 (Amari et al. 2001a); their origins in the outflows of nova explosions, however, were questioned and SNe were proposed for at least some of the putative nova grains instead (Nittler & Hoppe 2005; Pignatari et al. 2015; Liu et al. 2017a; Hoppe et al. 2018b).","Citation Text":["Marty et al. 2011"],"Functions Text":["The mainstream grains have 12C\/13C ratios between 10 and 100 (solar: 89), and the 14N\/15N ratios of most of them are higher than the solar ratio of 440, the ratio measured for the solar wind"],"Functions Label":["Background"],"Citation Start End":[[982,999]],"Functions Start End":[[790,980]]}
{"Identifier":"2015AandA...579A.102B__Boselli_et_al._2009_Instance_1","Paragraph":"Once corrected for dust attenuation, H\u03b1 luminosities can be transformed into star formation rates (SFR, in\u2009M\u2299 yr-1) using a factor that depends on the assumed IMF and stellar model7: (10)\\begin{equation} {SFR = k({\\rm H}\\alpha) \\times L({\\rm H}\\alpha)_{\\rm cor}} . \\end{equation}SFR=k(H\u03b1)\u00d7L(H\u03b1)cor.We recall that this relation is valid only under the assumption that the mean star formation activity of the emitting galaxies is constant on a timescale of a few Myr, roughly comparable to the typical time spent by the stellar population responsible for the ionisation of the gas on the main sequence (Boselli et al. 2009; Boissier 2013; Boquien et al. 2014). The ionising stars are O and early-B stars, whose typical age is \u2272107 yr. The stationarity condition is generally satisfied in massive, normal, star-forming galaxies undergoing secular evolution. In these objects, the total number of OB associations is significantly larger than the number of HII regions under formation and of OB stars reaching the final stage of their evolution, thus their total H\u03b1 luminosity is fairly constant with time. This might not be the case in strongly perturbed systems or in dwarf galaxies, where the total star formation activity can be dominated by individual giant HII regions (Boselli et al. 2009; Weisz et al. 2012), and the IMF is only stochastically sampled (Lee et al. 2009; Fumagalli et al. 2011; da Silva et al. 2014). The HRS sample is dominated by relatively massive galaxies undergoing secular evolution. For these objects, Eq. (10) can thus be applied. The sample, however, also includes galaxies in the Virgo cluster region, where the perturbation induced by the cluster environment might have affected their star formation rate (e.g. Boselli & Gavazzi 2006, 2014). Models and simulations have shown that in these objects the suppression of star formation occurs on timescales of a few hundred Myr (Boselli et al. 2006, 2008a,b, 2014d). These timescales are relatively long compared to the typical age of O-B stars. The recent work of Boquien et al. (2014) has clearly shown that the Lyman continuum emission tightly follows the rapid variations in the star formation activity of simulated galaxies down to timescales of a few Myrs. We can thus safely consider that the linear relation between the H\u03b1 luminosity and the star formation rate given in Eq. (10) is satisfied in the HRS sample. ","Citation Text":["Boselli et al. 2009"],"Functions Text":["We recall that this relation is valid only under the assumption that the mean star formation activity of the emitting galaxies is constant on a timescale of a few Myr, roughly comparable to the typical time spent by the stellar population responsible for the ionisation of the gas on the main sequence"],"Functions Label":["Uses"],"Citation Start End":[[601,620]],"Functions Start End":[[298,599]]}
{"Identifier":"2021AandA...656A..16C__Bruno_&_Carbone_2013_Instance_3","Paragraph":"Investigations of the turbulent nature of solar wind fluctuations have been ongoing for more than half a century (see, e.g., Bruno & Carbone 2016). Advances have been made consistently thanks to the increasingly accurate measurements of several dedicated space mission as well as to the enormous improvement of numerical calculation, new detailed models and theoretical frameworks, and the development of specific data analysis techniques. Nevertheless, the extremely complex nature of the system and the coexistence of multiple actors, scales, and dynamical regimes have led to a number of questions that remain open (Viall & Borovsky 2020). Among these, the very nature of the turbulent cascade of the solar wind flow and its relationship with the small-scale processes still need to be described in full (Tu & Marsch 1995; Bruno & Carbone 2013; Matthaeus & Velli 2011; Chen 2016). Magnetic field fluctuations have been characterized with great detail at magnetohydrodynamic and kinetic scales, for example, through spectral and high-order moments analysis (Tu & Marsch 1995; Bruno & Carbone 2013). The anisotropic nature of magnetic turbulence has also been addressed, and is still being debated, due to the limited access to three-dimensional measurements in space (see, e.g., Horbury et al. 2008, 2012; Sorriso-Valvo et al. 2010; Yordanova et al. 2015; Verdini et al. 2018; Telloni et al. 2019a; Oughton & Matthaeus 2020). Velocity fluctuations have been studied thoroughly (see, e.g., Sorriso-Valvo et al. 1999; Bruno & Carbone 2013), although the kinetic scales still remain quite unexplored for instrumental limitations, most notably in the sampling time resolution. Both the velocity and magnetic field show highly variable turbulence properties, with well developed spectra, strong intermittency (Sorriso-Valvo et al. 1999), anisotropy, and linear third-order moments scaling (Sorriso-Valvo et al. 2007; Carbone et al. 2011). The level of Alfv\u00e9nic fluctuations (mostly but not exclusively found in fast streams, see e.g., D\u2019Amicis et al. 2011; Bruno et al. 2019) are believed to be associated with the state of the turbulence. In particular, solar wind samples containing more Alfv\u00e9nic fluctuations are typically associated with less developed turbulence, as inferred from both shallower spectra and reduced intermittency (see Bruno & Carbone 2013, and references therein). This is consistent with the expectation that uncorrelated Alfv\u00e9nic fluctuations contribute to reduce the nonlinear cascade by sweeping away the interacting structures (Dobrowolny et al. 1980), as also confirmed by the observed anticorrelation between the turbulent energy cascade rate and the cross-helicity (Smith et al. 2009; Marino et al. 2011a,b).","Citation Text":["Bruno & Carbone 2013"],"Functions Text":["Velocity fluctuations have been studied thoroughly (see, e.g.,","although the kinetic scales still remain quite unexplored for instrumental limitations, most notably in the sampling time resolution."],"Functions Label":["Background","Future Work"],"Citation Start End":[[1518,1538]],"Functions Start End":[[1428,1490],[1541,1674]]}
{"Identifier":"2022MNRAS.514.3894O__Buxton_et_al._2012_Instance_1","Paragraph":"The GBHT outburst light curves could be very complicated, and while the so-called \u2018the main outburst\u2019 could go through the spectral states described above, some GBHTs also show rebrightening episodes during the outburst decay (Kalemci et al. 2013) and\/or an increase in brightness several days after the X-ray flux goes below the detection limits of the most observatories that are sometimes defined as mini-outbursts (Chen et al. 1997). A systematic multiwavelength study of GBHTs in the outburst decay by Kalemci et al. (2013) showed that for most of the systems, a rebrightening (secondary maximum or secondary flare) in OIR occurred \u223c1\u20132 weeks after the soft-to-hard transition. Detection of rebrightening during the outburst decay supports the argument that the formation of a compact jet and its interaction with the accretion environment are imprinted on the multiwavelength behaviour of the GBHTs (Buxton & Bailyn 2004; Kalemci et al. 2005, 2013; Buxton et al. 2012; Din\u00e7er et al. 2012; Corbel et al. 2013). Alternatively, the synchrotron radiation from the hot accretion flow model (Poutanen 1998; Veledina et al. 2013), or the irradiation from the secondary star or outer part of the disc could explain the brightness increase in the OIR bands. In contrast, there are limited number of pointed hard X-ray observations for the mini-outbursts (e.g. XTE J1752 \u2212 223, SWIFT J1745 \u2212 26, and V404 Cyg, Chun et al. 2013; Kalemci et al. 2014; Mu\u00f1oz-Darias et al. 2017) since they have been observed frequently in the soft X-rays and optical (see Chen et al. 1997 for some historical examples, both in black holes and neutron stars). A recent study by Zhang et al. (2019) attempted a classification of the rebrightenings during\/after the main outburst decay based on the the available fluxes and applied this scheme to Swift J1753.5 \u2212 0127, which showed a mini-outburst in radio, optical, and X-rays. It can be seen that different flavours exist depending on whether the source reaches quiescence first. Some sources show multiple mini-outbursts after the initial outburst (e.g. XTE J1650 \u2212 500, MAXI J1535 \u2212 571, Tomsick et al. 2003; C\u00faneo et al. 2020). Although the origin of the mini-outbursts is still debated, an increased mass accretion triggered by the events during the evolution of the primary outburst through heating of the outer parts of the accretion disc (Ertan & Alpar 2002), or the companion star (Augusteijn et al. 1993), is known to be the likely explanations.","Citation Text":["Buxton et al. 2012"],"Functions Text":["Detection of rebrightening during the outburst decay supports the argument that the formation of a compact jet and its interaction with the accretion environment are imprinted on the multiwavelength behaviour of the GBHTs"],"Functions Label":["Similarities"],"Citation Start End":[[955,973]],"Functions Start End":[[683,904]]}
{"Identifier":"2021ApJ...923L..22A__Cordes_&_Jenet_2012_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":["Cordes & Jenet 2012"],"Functions Text":["Orbiting SMBHBs produce a","and transient GW bursts"],"Functions Label":["Background","Background"],"Citation Start End":[[1168,1187]],"Functions Start End":[[484,509],[1115,1138]]}
{"Identifier":"2016MNRAS.461.4176H__Jaffe_&_Kaiser_1995_Instance_1","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"],"Functions Text":["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"],"Functions Label":["Uses"],"Citation Start End":[[1094,1113]],"Functions Start End":[[855,1092]]}
{"Identifier":"2018MNRAS.480.4154C__Maisinger,_Hobson_&_Lasenby_2004_Instance_1","Paragraph":"Classical imaging techniques were developed in the field to solve the RI reconstruction problem, such as clean and its multiscale variants (H\u00f6gbom 1974; Bhatnagar & Corwnell 2004; Cornwell 2008; Stewart, Fenech & Muxlow 2011). In particular, clean builds a model image by iteratively removing point source components from the residuals of the acquired data (at each iteration). clean-based algorithms, however, are typically slow (generally requiring computationally demanding major cycles; cf. Clark clean), requiring fine-tuning and supervision, while providing suboptimal imaging quality (see e.g. Li, Cornwell & de Hoog 2011a; Carrillo, McEwen & Wiaux 2012). Another classical technique is the maximum entropy method (MEM) (Ables 1974; Gull & Daniell 1978), extended to RI imaging by Cornwell & Evans (1985). The MEM approach of Cornwell & Evans (1985) developed for RI imaging considers a regularization problem consisting of a relative entropic prior, a (Gaussian) likelihood term and an additional flux constraint. In principle, MEM requires less fine-tuning and supervision compared to clean and can therefore alleviate part of the shortcomings of clean-based algorithms. However, an optimal metric \u2013 expressed as an entropy functional \u2013 is not known in advance and therefore needs to be chosen individually (Starck et al. 2001; Maisinger, Hobson & Lasenby 2004). Indeed, it is widely known that MEM fails to reconstruct sharp and smooth image features simultaneously. Recently, the theory of compressed sensing (CS) has suggested the use of sparse representation and regularization approaches for the recovery of sparse signals from incomplete linear measurements (Donoho 2006; Candes & Wakin 2008; Candes et al. 2010), which has shown great success. CS techniques based on sparse regularization were ushered into RI imaging for image reconstruction (Suksmono 2009; Wiaux et al. 2009a,b; Wenger et al. 2010; Li et al. 2011a,b; McEwen & Wiaux 2011; Carrillo et al. 2012; Wolz et al. 2013; Carrillo, McEwen & Wiaux 2014; Dabbech et al. 2015; Garsden et al. 2015; Onose et al. 2016; Dabbech et al. 2017; Kartik et al. 2017; Onose, Dabbech & Wiaux 2017; Pratley et al. 2018) and have shown promising results and improvements compared to traditional approaches such as clean-based methods and MEM. In general, such approaches can recover sharp and smooth image features simultaneously (e.g. Carrillo et al. 2012). While sparse approaches have been shown to be highly effective, the best approach to image different sources remains an open question. Algorithms have been developed to scale sparse approaches to big-data (Carrillo et al. 2014; Onose et al. 2016; Cai, Pratley & McEwen 2017a; Kartik et al. 2017; Onose et al. 2017), such as that anticipated from the Square Kilometre Array (SKA1). However, clean-based methods, MEM, and CS-based methods, unfortunately, do not provide any uncertainty quantification about the accuracy of recovered images.","Citation Text":["Maisinger, Hobson & Lasenby 2004"],"Functions Text":["However, an optimal metric \u2013 expressed as an entropy functional \u2013 is not known in advance and therefore needs to be chosen individually"],"Functions Label":["Uses"],"Citation Start End":[[1337,1369]],"Functions Start End":[[1180,1315]]}
{"Identifier":"2020AandA...641A.155V__Spilker_et_al._(2016)_Instance_1","Paragraph":"The right panel of Fig. 7 shows the relation between \u03a3SFR and R52. For each object, we computed \u03a3SFR\u2004=\u2004SFR\/(2\u03c0R2), where R is a representative value of the galaxy radius. The latter is rather arbitrary and it depends on the chosen tracer, the depth, resolution, and wavelength of the observations. Here we adopted the ALMA sizes from circular Gaussian fitting for our sample, assuming R\u2004=\u2004FWHM\/2. As mentioned in Sect. 3.2, this estimate combines all the available lines and continuum measurements, resulting in a size representative of the dust and gas content of each galaxy (Puglisi et al. 2019). We further recomputed the \u03a3SFR for the BzK galaxies in D15, using the Gaussian best-fit results of the rest-frame UV observations to be consistent with our estimates. For the SPT-SMGs, we used the sizes of Spilker et al. (2016), while we employed the 1.4 GHz radio measurements in Liu et al. (2015b) for the local spirals. For reference, we also show the mean values for the BzK galaxies, the local spirals, and ULIRGs as in D15. The best-fit model to the observed points returns a 60% flatter slope than in D15 (Table 3), but the trends are qualitatively similar. We restate that the choice of the tracer, the resolution, and depth of the observations play a major role in setting the exact values of the slope and intercept in our simple linear model, which should be thus taken with a grain of salt. This is particularly true for spatially resolved local objects, where we attempted to replicate the global, galaxy-scale measurements that can be obtained for distant objects. The observed data points in Fig. 7 qualitatively agree with the simulations by Narayanan & Krumholz (2014) and Bournaud et al. (2015), and they support the validity of \u03a3SFR as a good proxy for the gas conditions in galaxies. The total SFR is a worse predictor of the gas excitation conditions (Lu et al. 2014; Kamenetzky et al. 2016), since it does not correlate with the density and temperature probability distribution functions in clouds (Narayanan & Krumholz 2014). Interestingly, this seems to be partially confirmed by the linear regression analysis we applied here (Table 3): when modeling R52 as a function of LIR (\u221dSFR, Fig. 5) and \u03a3SFR, we do find similar slopes, but also a larger correlation coefficient \u03c1 for \u03a3SFR than for LIR. However, LIR alone does correlate with the CO line luminosity ratio.","Citation Text":["Spilker et al. (2016)"],"Functions Text":["For the SPT-SMGs, we used the sizes of"],"Functions Label":["Uses"],"Citation Start End":[[806,827]],"Functions Start End":[[767,805]]}
{"Identifier":"2022MNRAS.512.3137Z__Katz_et_al._1999_Instance_2","Paragraph":"However, it is not straightforward to explain H2 formation in astronomical sources even when the catalytic roles of dust grains are introduced into models. Interstellar species are believed to be formed on cold grain surfaces via the so called Langmuir\u2013Hinshelwood mechanism (Watson & Salpeter 1972; Pickles & Williams 1977; Hasegawa, Herbst & Leung 1992). To form H2, H atoms accrete on dust grains and then bind weakly with surfaces, which is known as physisorption. They can overcome the diffusion barrier and move on the grain surfaces via quantum tunnelling or thermal hopping. However, laboratory and theorectical studies showed that quantum tunnelling does not contribute much to the mobility of H atoms on silicate, carbonaceous or solid amorphous water (ASW) surfaces (Pirronello et al. 1997, 1999; Katz et al. 1999; Nyman 2021). If H atoms encounter other H atoms, then H2 molecules are formed. But H atoms can also desorb and leave grain surfaces. A hydrogen atom must reside on a grain long enough to find a partner H atom to form H2. As the dust temperature increases, the H atom desorption and diffusion rates also increase. So the temperature of grain surfaces must be sufficiently low so that an H atom can encounter another one before it desorbs. On the other hand, the temperature of grain surfaces must be high enough so that H atoms can diffuse on the grain surface. The parameter that measures how strongly species are to bound to grain surfaces is called desorption energy. It was found that if we assume a single H desorption energy in models, the dust temperature range over which efficient H2 formation occurs is narrow (6\u201310 K for olivine grains) (Katz et al. 1999). Moreover, it was found that the highest dust temperature at which H2 can be formed efficiently is less than 17\u2009K (Katz et al. 1999). However, the grain surface temperature in the unshielded diffuse clouds, where hydrogen molecules are believed to be efficiently formed, is around 20 K (Li & Draine 2001).","Citation Text":["Katz et al. 1999"],"Functions Text":["It was found that if we assume a single H desorption energy in models, the dust temperature range over which efficient H2 formation occurs is narrow (6\u201310 K for olivine grains)"],"Functions Label":["Background"],"Citation Start End":[[1674,1690]],"Functions Start End":[[1496,1672]]}
{"Identifier":"2020ApJ...900..100R__White_et_al._2019_Instance_1","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"],"Functions Text":["In the magnetically arrested disk","scenario, the MRI and subsequent turbulence in the inner accretion disk are suppressed due to large-scale magnetic flux (see, e.g.,"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[817,834]],"Functions Start End":[[598,631],[685,816]]}
{"Identifier":"2020ApJ...893..124Z__Tenbarge_&_Howes_2013_Instance_1","Paragraph":"In hydrodynamic and magnetohydrodynamic systems, different fluctuations interact with each other nonlinearly, generating turbulence (Matthaeus et al. 2015). One of the most important characteristics of turbulence is the existence of intermittency among various scales. Spatial intermittency manifests as coherent structures with large gradient at small scales. In ordinary fluids, coherent structures include a tangle of vortex filaments where vorticity is highly concentrated (Frisch 1995). In the turbulent plasma environments, examples of intermittent structures include current sheets, discontinuities, shock waves, and Alfv\u00e9nic vortices (Veltri & Mangeney 1999; Sundkvist et al. 2007; Lion et al. 2016; Wang et al. 2019). The intermittency will influence the measured statistical properties of the fluctuating quantities. For example, the growth of flatness with decreasing scales indicates the existence of intermittency. Plenty of simulation works (Servidio et al. 2011; Karimabadi et al. 2013; Tenbarge & Howes 2013; Wan et al. 2015, 2016; Zhang et al. 2015) have indicated that dissipation, acceleration, and thermalization of turbulent plasmas mainly take place near intermittent structures on kinetic scale, while the actual physical mechanisms behind dissipation remain unclear. There are a variety of diagnostic approaches to measure intermittency. (1) Probability density functions (PDFs) of scale-dependent field increments develop heavy tails because of intermittency, and the tail is more enhanced with increasing intermittency (Marsch & Tu 1994). (2) The scale dependency of the normalized fourth-order moment, known as flatness, \n\n\n\n\n\n, where \u03b4v = v(t + \u03c4)\u2212v(t), is an alternative representation. This quantity increases as the intermittency becomes more significant along with decreasing scales. (3) The p-th order structure function, \n\n\n\n\n\n, will be larger than that for Gaussian PDF without intermittency and the scaling exponent, \u03b6(p), has a nonlinear form when intermittency presents. Based on such diagnostic approaches, the properties of intermittency in plasmas have been widely investigated in the solar atmosphere, magnetosheath, solar wind, termination shock, etc. (Burlaga 1991a, 1991b; Marsch & Liu 1993; Horbury et al. 1995; Macek et al. 2011, 2017; Chasapis et al. 2018). These observational analyses have revealed that intermittency is widely existing in the heliosphere and immensely influencing plasma dynamics.","Citation Text":["Tenbarge & Howes 2013"],"Functions Text":["Plenty of simulation works","have indicated that dissipation, acceleration, and thermalization of turbulent plasmas mainly take place near intermittent structures on kinetic scale, while the actual physical mechanisms behind dissipation remain unclear."],"Functions Label":["Background","Background"],"Citation Start End":[[1002,1023]],"Functions Start End":[[928,954],[1067,1290]]}
{"Identifier":"2018MNRAS.475.4011B__Zauderer_et_al._2011_Instance_1","Paragraph":"ASASSN-14li was discovered by the All Sky Automated Search for Supernova (ASASSN; Shappee et al. 2014) on ut 2014\u201311\u201322.63 (MJD 56983.6) as a 16.5 magnitude source in the V band (Jose et al. 2014; Holoien et al. 2016; Brown et al. 2017a). The position of the source was found to be consistent with the centre of the post-starburst galaxy PGC 043234, with a measured projected separation of 0.04 arcsec. This galaxy is at redshift z = 0.0206 with a luminosity distance of 90.3\u2009Mpc (for cosmological parameters H0 = 73\u2009kms\u22121\u2009Mpc\u22121, \u03a9matter = 0.27, and \u03a9\u039b = 0.73). It was established through archival X-ray observations of PGC 043234 from the ROSAT All-Sky Survey (Voges et al. 1999) that the galaxy does not contain an efficiently accreting AGN, with the count rate implying a luminosity orders of magnitude below standard active nuclei (e.g. Miller et al. 2015). A small number (currently six) of confirmed TDEs, including ASASSN-14li, have also been detected at radio wavelengths and the population may form a bi-modal distribution, consisting of more common non-relativistic \u2018thermal\u2019 events and rarer relativistic jets. Three events (Swift J1644+57; Burrows et al. 2011; Zauderer et al. 2011, Swift J2058+05; Cenko et al. 2012, Swift J1112.2; Brown et al. 2017b) have isotropic \u223c5\u2009GHz luminosities of between 1040 and 1042\u2009erg\u2009s\u22121 whereas the rest (IGR J12580+0134; Irwin et al. 2015, XMMSL1 J0740-85; Alexander et al. 2017, ASASSN-14li; van Velzen et al. 2016; Alexander et al. 2016) have luminosities in the range 1037 to 1039\u2009erg\u2009s\u22121 at similar frequencies. The higher power events are believed to result from observing down the axis of a relativistic jet, resulting in the energy of photons being significantly boosted. Even accounting for boosting, these relativistic events have a higher total energy output than their thermal counterparts. The origin of the radio emission from the thermal events is currently uncertain, with transient jets (van Velzen et al. 2016), non-relativistic winds (Alexander et al. 2016), and shocks driven by unbound material (Krolik et al. 2016) all feasible scenarios. ASASSN-14li is by far the best studied of the \u2018thermal\u2019 TDE category, having been observed extensively at Optical, UV, X-ray (where ASASSN-14li is unusually loud for an optically selected TDE), and radio wavelengths. The high cadence X-ray and radio observations in particular allow for the X-ray\/radio coupling to be probed.","Citation Text":["Zauderer et al. 2011"],"Functions Text":["A small number (currently six) of confirmed TDEs, including ASASSN-14li, have also been detected at radio wavelengths and the population may form a bi-modal distribution, consisting of more common non-relativistic \u2018thermal\u2019 events and rarer relativistic jets. Three events (Swift J1644+57","have isotropic \u223c5\u2009GHz luminosities of between 1040 and 1042\u2009erg\u2009s\u22121","The higher power events are believed to result from observing down the axis of a relativistic jet, resulting in the energy of photons being significantly boosted."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[1173,1193]],"Functions Start End":[[862,1150],[1265,1332],[1563,1725]]}
{"Identifier":"2022ApJ...933..243F__Woosley_&_Bloom_2006b_Instance_1","Paragraph":"Gamma-ray bursts (GRBs) are among the most powerful gamma-ray sources in the universe. They could be generated from the merger of binary compact objects (BCOs; Duncan & Thompson 1992; Usov 1992; Thompson 1994; Metzger et al. 2011) or the death of massive stars (Woosley 1993; Paczy\u0144ski 1998; Woosley & Bloom 2006a). The merger of BCOs; a black hole (BH)\u2013a neutron star (NS) or NS\u2013NS, leading to kilonovae (KNe), is correlated with short-duration gamma-ray bursts (sGRBs; T\n90\n\n10\n\n\n10\n\nT\n90 is defined as the time during which the cumulative number of collected counts above background rises from 5% to 95%. \u2272 2 s; Li & Paczy\u0144ski 1998; Rosswog 2005; Metzger et al. 2010; Kasen et al. 2013; Metzger 2017). On the other hand, long-duration gamma-ray bursts (lGRBs; T\n90 \u2273 2 s; Kouveliotou et al. 1993) are associated with the core collapse (CC) of dying massive stars (Woosley 1993; Galama et al. 1998) leading to supernovae (SNe; Bloom et al. 1999; Woosley & Bloom 2006b). It is believed that in both scenarios large quantities of materials with a wide range of velocities are ejected. In the framework of CC-SNe (depending on the type of SN association), several materials ejected with sub-relativistic velocities less than \u03b2 \u2272 0.4\n11\n\n\n11\nHereafter, we adopt natural units c = \u210f = 1. have been reported (see, e.g., Kulkarni et al. 1998; Bloom et al. 1999; Woosley & Bloom 2006b; Valenti et al. 2008; Gal-Yam 2017; Izzo et al. 2019, 2020; Modjaz et al. 2020; Nicholl et al. 2020). Regarding the merger of two NSs, sub-relativistic materials such as the cocoon, the shock breakout, and the dynamical and wind ejecta are launched with velocities in the range 0.03 \u2272 \u03b2 \u2272 0.8\n12\n\n\n12\nSome authors have considered the shock breakout material in the sub-, trans-, and ultra-relativistic regimes (see, e.g., Kyutoku et al. 2014; Metzger et al. 2015; Fraija et al. 2019c). (see, e.g., Dessart et al. 2009; Metzger & Fern\u00e1ndez 2014; Fern\u00e1ndez et al. 2015; Kyutoku et al. 2014; Metzger et al. 2015; Nagakura et al. 2014; Murguia-Berthier et al. 2014; Lazzati et al. 2017, 2018; Goriely et al. 2011; Hotokezaka et al. 2013; Bauswein et al. 2013; Wanajo et al. 2014). While the mass and velocity inferred for the first GRB\/KN association\n13\n\n\n13\nGRB 170817A\/AT 2017gfo. were M\nej \u2248 (10\u22124\u221210\u22122)M\n\u2299 and \u03b2 \u2248 (0.1\u22120.3), respectively (Coulter et al. 2017; Arcavi et al. 2017; Cowperthwaite et al. 2017; Nicholl et al. 2017; Metzger 2019), the mass and velocity inferred for the first GRB\/SN association\n14\n\n\n14\nGRB 980425\/SN1998bw. was M\nej \u2248 10\u22125\nM\n\u2299 and \u03b2 \u2248 (0.2\u20130.3), respectively (Kulkarni et al. 1998).","Citation Text":["Woosley & Bloom 2006b"],"Functions Text":["On the other hand, long-duration gamma-ray bursts (lGRBs","are associated with the core collapse (CC) of dying massive stars","leading to supernovae (SNe;"],"Functions Label":["Background","Background","Background"],"Citation Start End":[[948,969]],"Functions Start End":[[705,761],[800,865],[901,928]]}
{"Identifier":"2017ApJ...835....2X__Collins_et_al._2012_Instance_1","Paragraph":"On the other hand, a clear physical interpretation of the observed pulse broadening phenomenon requires a good understanding of the interstellar electron density structure. A power-law model of electron density fluctuations is commonly adopted in theoretical constructions on radio wave propagation (Lee & Jokipii 1976; Rickett 1977, 1990) and is compatible with observational indications (e.g., Armstrong et al. 1995). Recent advances in understanding the properties of magnetohydrodynamic (MHD) turbulence (Goldreich & Sridhar 1995; Lithwick & Goldreich 2001; Cho & Lazarian 2002, 2003) stimulate a renewed investigation on density statistics (Beresnyak et al. 2005; Kowal et al. 2007; Lazarian et al. 2008; Burkhart et al. 2009, 2010, 2015; Collins et al. 2012; Federrath & Klessen 2012), which provide important insight into key physical processes such as star formation in the turbulent and magnetized ISM (see reviews by, e.g., McKee & Ostriker 2007; Lazarian et al. 2015). The density spectrum in compressible astrophysical fluids was systematically studied in Kowal et al. (2007) by carrying out an extensive set of MHD numerical simulations with varying compressibility and magnetization. Instead of a single Kolmogorov slope with a power-law index of \n\n\n\n\n\n, significant variations in the spectral slope of density fluctuations are present. For supersonic turbulence, their results are consistent with earlier findings in both magnetized (Beresnyak et al. 2005) and nonmagnetized (Kim & Ryu 2005) fluids. It shows that the density power spectrum becomes shallower as the sonic Mach number (\n\n\n\n\n\n) increases, where VL is the turbulent velocity at the outer scale of turbulence and cs is the sound speed in the medium, and there is a significant excess of density structures at small scales in highly supersonic turbulence. This behavior is naturally expected as the gas is compressed in shocks by supersonic flows and the interacting shocks produce local density enhancements (Mac Low & Klessen 2004; Padoan et al. 2004b).","Citation Text":["Collins et al. 2012"],"Functions Text":["Recent advances in understanding the properties of magnetohydrodynamic (MHD) turbulence","stimulate a renewed investigation on density statistics"],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[744,763]],"Functions Start End":[[420,507],[589,644]]}
{"Identifier":"2017MNRAS.471.2917K__Glenn_et_al._2015_Instance_1","Paragraph":"The only way to access the transitions which trace warmer molecular gas is to get above the atmosphere as did the Herschel Space Observatory The Herschel SPIRE Fourier Transform Spectrometer (FTS) was simultaneously sensitive to the J = 4\u22123 (EJ\u2009=\u20094 = 55 K) up to the J = 13\u221212 (EJ\u2009=\u200913 = 503 K) line of CO for local galaxies. With this access to higher energy CO transitions, Herschel gave us the tools to more finely discern the lower and higher excitation components of the multiphase molecular ISM. To more finely measure the excitation of the molecular ISM is to better understand energy exchange between the star formation and the ISM. Many studies of individual galaxies (Panuzzo et al. 2010; Rangwala et al. 2011; Hailey-Dunsheath et al. 2012; Kamenetzky et al. 2012; Spinoglio et al. 2012; Meijerink et al. 2013; Pellegrini et al. 2013; Rigopoulou et al. 2013; Israel et al. 2014; Papadopoulos et al. 2014; Rosenberg et al. 2014; Schirm et al. 2014; Glenn et al. 2015; Wu et al. 2015) and surveys (Papadopoulos et al. 2010, 2012a,b; Pereira-Santaella et al. 2013, 2014; Greve et al. 2014; Kamenetzky et al. 2014, 2016; Lu et al. 2014; Daddi et al. 2015; Israel, Rosenberg & van der Werf 2015; Liu et al. 2015; Mashian et al. 2015) have shown that the higher J lines of CO arise from warmer, denser gas than the cold gas responsible for the emission of for example CO J = 1\u22120. Though this warm gas is only a small fraction (\u223c10\u2009per\u2009cent) of the total molecular mass, it is responsible for \u223c90\u2009per\u2009cent of the CO luminosity (Kamenetzky et al. 2014). This high CO luminosity cannot be explained by excitation from the ultraviolet (UV) light from young O and B stars in photon-dominated regions (PDRs); mechanical heating is often required. Theorists have attempted to use hydrodynamical galaxy simulations to produce galaxy-integrated spectral line energy distributions (SLEDs) of CO emission (Narayanan et al. 2008; Olsen et al. 2016), but we still lack a comprehensive picture of the mechanisms responsible for this emission. Furthermore, even higher J lines of CO than addressed here (above J = 13\u221212) have been detected, for example with Herschel PACS, indicating a third, even higher temperature component of molecular gas (NGC1068, Hailey-Dunsheath et al. 2012). The CO SLEDs from J = 14\u221213 through J = 30\u221229 show a large range in SLED shape, even among similar galaxies (Mashian et al. 2015).","Citation Text":["Glenn et al. 2015"],"Functions Text":["Many studies of individual galaxies","have shown that the higher J lines of CO arise from warmer, denser gas than the cold gas responsible for the emission of for example CO J = 1\u22120."],"Functions Label":["Background","Background"],"Citation Start End":[[958,975]],"Functions Start End":[[641,676],[1239,1383]]}
{"Identifier":"2022MNRAS.510.1043B__Kukula_et_al._1998_Instance_1","Paragraph":"The origin of the radio emission in radio-loud active galactic nuclei (RL AGN) is clear, luminous relativistic jets of magnetized plasma, which can extend far out, on the host galaxy scale and beyond. Conversely, radio-quiet (RQ) AGN are associated with radio emission which is typically 103 times weaker (as defined by Kellermann et al. 1989), in smaller structures (kpc-pc scale; e.g. Blundell et al. 1996; Nagar et al. 1999; Ulvestad et al. 2005a; Gallimore et al. 2006) with sub-relativistic velocities (e.g. Middelberg et al. 2004; Ulvestad, Antonucci & Barvainis 2005b) compared to RL AGN. The reduced sizes and low brightness of the radio emission of RQ AGN create a major challenge for detailed studies, in sharp contrast with the thoroughly studied RL AGN. The fewer radio studies of RQ AGN (e.g. Barvainis & Antonucci 1989; Kellermann et al. 1994; Barvainis, Lonsdale & Antonucci 1996; Kukula et al. 1998; Ulvestad et al. 2005a; Leipski et al. 2006; Doi et al. 2011; Padovani et al. 2011; Zakamska et al. 2016; Jarvis et al. 2019, 2021; Fawcett et al. 2020; Nyland et al. 2020; Smith et al. 2020b; Baldi et al. 2021a) generally lead to mixed results. This encourages to keep investigating it, as it indicates that the origin of the radio emission in RQ AGN is still an open question. If a number of different processes are indeed involved, then the radio band can be used to probe a range of physical processes, rather than being heavily dominated by a single process, as occurs in RL AGN (see Blandford, Meier & Readhead 2019 Panessa et al. 2019 for reviews). From large scale to small scale: (i) host galaxy star formation (SF) could account for the observed FIR-to-radio emission observed in active and non-active galaxies (Condon et al. 2013; Zakamska et al. 2016); (ii) an AGN-driven wind is expected to shock the interstellar gas, leading to particle acceleration and radio synchrotron emission, which may reach the observed flux level (Jiang et al. 2010); (iii) the intense radiation of the AGN photoionizes large volumes of ambient gas, as supported by the strength of the narrow and broad line emission observed in type-1 AGN, leading to thermal free\u2013free emission in the radio band (Baskin & Laor 2021); (iv) a scaled-down jet, physically similar to the one in RL AGN, but much fainter, less energetic and slower (Barvainis et al. 1996; Gallimore et al. 2006; Talbot, Sijacki & Bourne 2021); (v) a tight radio\/X-ray luminosity relation for AGN (\u223c10\u22125; Laor & Behar 2008) similar to coronally active stars (G\u00fcdel & Benz 1993; G\u00fcdel et al. 2002) suggests that coronal emission from magnetic activity above the accretion disc (Field & Rogers 1993; Gallimore, Baum & O\u2019Dea 1997) may produce the observed radio emission.","Citation Text":["Kukula et al. 199"],"Functions Text":["The fewer radio studies of RQ AGN","generally lead to mixed results. This encourages to keep investigating it, as it indicates that the origin of the radio emission in RQ AGN is still an open question."],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[896,913]],"Functions Start End":[[766,799],[1128,1293]]}
{"Identifier":"2019ApJ...882..144K__Berk_et_al._2001_Instance_1","Paragraph":"The FOCAS and NIRSPEC spectra of PSO J006+39 were obtained at two different epochs separated by 1 yr and 9 months (by slightly less than 3 months in the quasar rest frame). Previously, we found that the PS1 y-band light curve of PSO J006+39 shows brightness variations with a peak-to-peak amplitude of \u223c0.7 mag over \u223c4 yr (Koptelova et al. 2017), which might be due to the flux variations of both continuum and Ly\u03b1 line of PSO J006+39. To infer the brightness state of PSO J006+39 at the epochs of its FOCAS and NIRSPEC observations, we first calculated the spectral slope of the quasar continuum from the NIRSPEC spectrum with a wider wavelength coverage than that of the FOCAS spectrum. Using wavelength intervals of 11100\u201311300, 11400\u201311600, 13085\u201313400, and 14700\u201315200 \u212b we measured a spectral slope of \u03b1\u03bb = \u22121.35 \u00b1 0.26, where the quoted uncertainty is the statistical error of the fit. The fitted power law is shown in Figure 2 with a solid line. The estimated continuum slope is consistent but somewhat flatter than the typical slope of luminous quasars (Zheng & Malkan 1993; Vanden Berk et al. 2001; Selsing et al. 2016). We then fitted the FOCAS data using the power law with a fixed spectral slope of \u03b1\u03bb = \u22121.35 and spectral windows of 9700\u20139850 and 10050\u201310100 \u212b. The spectral windows adopted for the analysis of the FOCAS and NIRSPEC spectra were taken to be similar to the rest-frame wavelength intervals commonly used to fit the continua of quasars (Vanden Berk et al. 2001; Decarli et al. 2010; Lusso et al. 2015) and less affected by the contribution from emission lines on the red side of the Big Blue Bump (BBB; e.g., Malkan 1983). The estimated continuum flux of PSO J006+39 at the epoch of its FOCAS observations is shown in Figure 2 with a dashed line. By comparing the continuum flux at the epochs of the FOCAS and NIRSPEC observations, we find that the brightness state of PSO J006+39 was different at these two epochs. PSO J006+39 was brighter by about 0.8 mag during the FOCAS observations than during the NIRSPEC observations. Thus, the continuum flux of PSO J006+39 might be different at different epochs depending on the brightness state of the quasar. Figure 2 also shows the fluxes of PSO J006+39 in the FOCAS Y, and NIRSPEC N2, N4, and N6 bands at the epochs of the FOCAS and NIRSPEC observations (see also Table 1).","Citation Text":["Vanden Berk et al. 2001"],"Functions Text":["The estimated continuum slope is consistent but somewhat flatter than the typical slope of luminous quasars"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1084,1107]],"Functions Start End":[[954,1061]]}
{"Identifier":"2020AandA...640A.121G__Swain_et_al._2009_Instance_1","Paragraph":"The planetary physical properties that can be derived from measurements with the two most successful exoplanet detection techniques, that is, with the radial velocity method and the transit method, are the planet radius (if the planet transits its star), mass (a lower limit if the planet does not transit its star), and the orbital period and distance, eccentricity, and inclination (if the planet transits its star) (see Perryman 2018, and references therein). The upper atmospheres (above optically thick clouds) of giant exoplanets and those of smaller planets around M stars can be probed using spectroscopy during primary and\/or secondary transits (Bean et al. 2010; Swain et al. 2009, 2008; Tinetti et al. 2007; Nakajima 1983). It appears to be virtually impossible to characterize the (lower) atmospheresand surfaces of small, Earth-like planets in the habitable zones of solar-type stars (B\u00e9tr\u00e9mieux & Kaltenegger 2014; Misra et al. 2014; Kaltenegger & Traub 2009), in particular because the light of the parent star is refracted while traveling through the lower atmosphere of its planet and emerges forever out of reach of terrestrial telescopes (Garc\u00eda Mu\u00f1oz et al. 2012). The (lower) atmosphere and surface of a planet are crucial for determining the habitability of a planet, as they hold information about cloud composition, trace gases in disequilibrium and probably most importantly, liquid surface water (see, e.g., Schwieterman et al. 2018; Kiang et al. 2007a,b, and references therein). For such a characterization of terrestrial-type planets, direct observations of the thermal radiation that they emit or of the light of their parent star that they reflect are required. The numerical results that we present in this paper concern the reflected starlight. Because of the huge distances involved, any measured reflected starlight pertains to the (illuminated and visible part of the) planetary disk. It therefore is a disk-integrated signal.","Citation Text":["Swain et al. 2009"],"Functions Text":["The upper atmospheres (above optically thick clouds) of giant exoplanets and those of smaller planets around M stars can be probed using spectroscopy during primary and\/or secondary transits"],"Functions Label":["Background"],"Citation Start End":[[673,690]],"Functions Start End":[[463,653]]}
{"Identifier":"2020ApJ...900...45L__Helmi_et_al._2018_Instance_1","Paragraph":"Much of the recent work on this topic has focused on gathering observational evidence for hierarchical growth through the identification of stellar streams and tidal interactions, which are thought to be the hallmark of low-mass galaxies having been accreted by more massive galaxies. The Sagittarius dwarf galaxy (Ibata et al. 1994) and the stellar streams associated with it (Belokurov et al. 2006) are some of the most well-studied examples of this type of hierarchical accretion in the Milky Way. Recent data from the Gaia mission (Gaia Collaboration et al. 2016) and other large-scale surveys, as well as analyses that combine kinematic data with information about chemical abundances and\/or specific stellar populations, have led to discoveries of additional substructures in the Galaxy and brought renewed attention to this topic (e.g., Belokurov et al. 2018; Deason et al. 2018; Helmi et al. 2018; Myeong et al. 2018). The Andromeda galaxy (M31) has been studied in detail as well. Structures formed from spatial overdensities of the red giant branch (RGB) population, e.g., the Giant Stellar Stream, have been found in the outer halo of M31 by using resolved RGB star maps (Ibata et al. 2001; Ferguson et al. 2002). Using the surface density of resolved RGB stars, it is possible to infer the surface brightness distribution of the underlying light; structures revealed by this technique, such as faint streams and clumps of RGB stars, suggest a rich history of accretion events and tidal interactions for M31. Deep observations of galaxies outside the Local Group have also revealed streams, shells, and satellites using resolved RGB stars and unresolved low surface brightness (LSB) features (e.g., Janowiecki et al. 2010; Crnojevic et al. 2016; Mihos et al. 2017). However, in galaxies well beyond the Local Group, RGB stars are unresolved and a different tracer of hierarchical assembly processes and tidal interactions is required. Globular clusters (GCs) have a number of properties that make them well suited for this task.","Citation Text":["Helmi et al. 2018"],"Functions Text":["Recent data from the Gaia mission","and other large-scale surveys, as well as analyses that combine kinematic data with information about chemical abundances and\/or specific stellar populations, have led to discoveries of additional substructures in the Galaxy and brought renewed attention to this topic (e.g.,"],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[887,904]],"Functions Start End":[[501,534],[568,843]]}
{"Identifier":"2015ApJ...806..167G__Daughton_&_Karimabadi_2007_Instance_1","Paragraph":"We envision a situation where intense current sheets are developed within a magnetically dominated plasma. Earlier work in nonrelativistic low-\u03b2 plasmas has shown that the gradual evolution of the magnetic field can lead to formation of intense, nearly force-free current layers where magnetic reconnection may be triggered (Galsgaard et al. 2003; Titov et al. 2003). In the present study, the critical parameter is the magnetization parameter defined as \n\n\n\n\n\n, which roughly corresponds to the available magnetic energy per particle. The numerical simulations presented in this paper are initialized from a force-free current layer with \n\n\n\n\n\n (Che et al. 2011; Liu et al. 2013, 2014), which corresponds to a magnetic field with magnitude B0 rotating by 180\u00b0 across the central layer with a half-thickness of \u03bb. No external guide field is included in this study but there is an intrinsic guide field By associated with the central sheet. The plasma consists of electron\u2013positron pairs with mass ratio \n\n\n\n\n\n. The initial distributions are Maxwellian with a spatially uniform density n0 and a thermal temperature (\n\n\n\n\n\n). Particles in the central sheet have a net drift \n\n\n\n\n\n to represent a current density \n\n\n\n\n\n that is consistent with \n\n\n\n\n\n. Since the force-free current sheet does not require a hot plasma component to balance the Lorentz force, this initial setup is more suitable to study reconnection in low-\u03b2 and\/or high-\u03c3 plasmas. The full particle simulations are performed using the VPIC code (Bowers et al. 2009) and NPIC code (Daughton et al. 2006; Daughton & Karimabadi 2007), both of which solve Maxwell equations and push particles using relativistic approaches. The VPIC code directly evolves electric and magnetic fields, whereas in the NPIC code the fields are advanced using the scalar and vector potentials. Although the two codes have very different algorithms, all of the key results are in good agreement for this study, thus providing additional confidence in our conclusions. In addition, we have developed a particle-tracking module to analyze the detailed physics of the particle energization process. In the simulations, we define and adjust \u03c3 by changing the ratio of the electron gyrofrequency \n\n\n\n\n\n to the electron plasma frequency \n\n\n\n\n\n, \n\n\n\n\n\n. For 2D simulations, we have performed simulations with \n\n\n\n\n\n and box sizes \n\n\n\n\n\n, \n\n\n\n\n\n, and \n\n\n\n\n\n, where di is the inertial length \n\n\n\n\n\n. For 3D simulations, the largest case is \n\n\n\n\n\n with \u03c3 = 100. For high-\u03c3 cases (\n\n\n\n\n\n), we choose cell sizes \n\n\n\n\n\n and \n\n\n\n\n\n, so the particle gyromotion scale \n\n\n\n\n\n is resolved. The time step is chosen to correspond to a Courant number \n\n\n\n\n\n, where \n\n\n\n\n\n. The half-thickness of the current sheet is \n\n\n\n\n\n for \n\n\n\n\n\n, \n\n\n\n\n\n for \u03c3 = 400, and \n\n\n\n\n\n for \u03c3 = 1600 in order to satisfy the drift velocity \n\n\n\n\n\n. For both 2D and 3D simulations, we have more than 100 electron\u2013positron pairs in each cell. The boundary conditions for 2D simulations are periodic for both fields and particles in the x direction, while in the z direction the boundaries are conducting for the field and reflecting for the particles. In the 3D simulations, the boundary conditions are periodic for both fields and particles in the y direction, while the boundary conditions in the x and z directions are the same as in the 2D cases. A weak long-wavelength perturbation (Birn et al. 2001) with \n\n\n\n\n\n is included to initiate reconnection. The parameters for different runs are summarized in Table 1, which also lists key results such as maximum energy of particles, spectral index, the fraction of kinetic energy converted from the magnetic energy, and the portion of energy gain arising from the perpendicular electric fields.","Citation Text":["Daughton & Karimabadi 2007"],"Functions Text":["The full particle simulations are performed using","and NPIC code","both of which solve Maxwell equations and push particles using relativistic approaches. The VPIC code directly evolves electric and magnetic fields, whereas in the NPIC code the fields are advanced using the scalar and vector potentials.","Although the two codes have very different algorithms, all of the key results are in good agreement for this study, thus providing additional confidence in our conclusions."],"Functions Label":["Uses","Uses","Background","Similarities"],"Citation Start End":[[1566,1592]],"Functions Start End":[[1444,1493],[1529,1542],[1595,1832],[1833,2005]]}
{"Identifier":"2019ApJ...877...35W__Hu_&_Yang_2014_Instance_1","Paragraph":"Here, for simplicity, computational convenience, and continuity with past works (Kopparapu et al. 2017; Haqq-Misra et al. 2018), we have assumed a slab ocean with zero ocean heat transport. Ocean heat transport may have a significant effect on climate and also the location of clouds (Way et al. 2018). Dynamic ocean heat transport on completely ocean-covered worlds (i.e., no continents) leads to warmer global mean temperatures and a reduction in day\u2013night temperature differences on synchronously rotating planets (Hu & Yang 2014; Del Genio et al. 2019). Yang et al. (2019) argue that ocean heat transport is critical for the treatment of cold, partially ice-covered planets around M dwarf stars; however, it does not have a meaningful effect on the climate and thermal emission phase curves for warm terrestrial planets (Ts \u223c 300 K) residing close to the inner edge of the habitable zone. The presence of continents also presents a significant uncertainty in the net effect of ocean heat transport on climate, as continents can reroute or eliminate day-to-night ocean heat transport entirely, resulting in climate states that are similar to those found in simulations with zero ocean heat transport (Del Genio et al. 2019; Yang et al. 2019). In this work, we have argued that clouds are likely more important in modulating phase curves for habitable planets, compared to surface energy transports. Still, ocean heat transport patterns, combined with variations in continental configurations and topographic prominences, could have significant complex feedback on the surface temperature, sea ice and snow cover, and cloud distributions, respectively, all of which affect thermal and reflected light phase curves. In future works, our group plans to couple our GCM with a dynamic ocean model to study the interactions between oceans, continents, topography, and observable features. However, the long equilibration times for a coupled dynamic atmosphere-ocean model still present a significant computational hinderance for conducting 3D simulations across wider parameter spaces, as we have done here.","Citation Text":["Hu & Yang 2014"],"Functions Text":["Dynamic ocean heat transport on completely ocean-covered worlds (i.e., no continents) leads to warmer global mean temperatures and a reduction in day\u2013night temperature differences on synchronously rotating planets","In future works, our group plans to couple our GCM with a dynamic ocean model to study the interactions between oceans, continents, topography, and observable features."],"Functions Label":["Motivation","Future Work"],"Citation Start End":[[518,532]],"Functions Start End":[[303,516],[1717,1885]]}
{"Identifier":"2021AandA...656A..63M__Shakura_&_Sunyaev_1973_Instance_1","Paragraph":"One possibility could be that the disk, since viewed at such high inclination, might create a \u201cshadow\u201d zone where part of the reflection incoming between Rin,\u20061 and Rin,\u20062 is obscured (see Fig. 10, panel a). As highlighted by the Figure, a jump in the scale height of the disk, located sufficiently far away from its inner edge, could in principle generate such self-shielding effect. A disk flared in the outer region was also put forward by Zdziarski et al. (2021) and Axelsson & Veledina (2021) in order to explain the presence of the outer reflection component and the variability associated to the iron line. However, it is not obvious whether a sudden puffing up of the disk could happen or not in a classical Shakura-Sunyaev disk, like the SAD. A jump could arise from a transition in the mechanism contributing mainly to the opacity in the disk. In the outer regions of a Shakura-Sunyaev disk, the opacity is expected to be mainly due to free-free absorption, while Thomson scattering dominates the opacity in the inner regions of the disk (see Eqs. (2.16) and (2.19) of Shakura & Sunyaev 1973). The geometrical thickness of the disk scales with the radius R following a (slightly) different relation; it is proportional to R\u221221\/20 in the inner Thomson scattering-dominated, regions and to R\u22129\/8 in the outer regions. At the boundary between these two zones (i.e., at a radius Rjump) a jump is expected. We calculated Rjump according to our results for \u1e40in and assuming the same value for \u03b1 in both equations. We found Rjump between 3\u2005\u00d7\u2005104\u2006RG and 7\u2005\u00d7\u2005104\u2006RG in all the epochs considered. In our fits the location of the outer reflector Rin,\u20062 was kept frozen to 300 RG. When the parameter is thawed, it is completely unconstrained, as shown in Fig. 11. Notwithstanding the uncertainty on Rin,\u20062, a boundary at beyond 104\u2006RG is likely located too far away to produce an extra reflection component. However, a crucial role might be played by irradiation from the inner regions of the disk to the outer regions, an ingredient neglected in Shakura-Sunyaev disks. Beyond some boundary radius, we expect that the upper layers of the disk should be heated up by the impinging photons coming from the inner regions, leading to evaporation and to a larger effective scale height.","Citation Text":["Shakura & Sunyaev 1973"],"Functions Text":["In the outer regions of a Shakura-Sunyaev disk, the opacity is expected to be mainly due to free-free absorption, while Thomson scattering dominates the opacity in the inner regions of the disk (see Eqs. (2.16) and (2.19) of","The geometrical thickness of the disk scales with the radius R following a (slightly) different relation; it is proportional to R\u221221\/20 in the inner Thomson scattering-dominated, regions and to R\u22129\/8 in the outer regions. At the boundary between these two zones (i.e., at a radius Rjump) a jump is expected. We calculated Rjump according to our results for \u1e40in and assuming the same value for \u03b1 in both equations. We found Rjump between 3\u2005\u00d7\u2005104\u2006RG and 7\u2005\u00d7\u2005104\u2006RG in all the epochs considered."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1079,1101]],"Functions Start End":[[854,1078],[1104,1596]]}
{"Identifier":"2022AandA...666A..60K__Gustafsson_et_al._2008_Instance_1","Paragraph":"High-precision photometric measurements have shown that the present treatment of limb darkening is not sufficiently accurate and leads to systematic errors in the derived parameters of the exoplanets (e.g., Espinoza & Jord\u00e1n 2016; Morello et al. 2017a; Maxted 2018). Typically, limb darkening is represented by rather simple laws, such as a linear law (Schwarzschild 1906), a quadratic law (Kopal 1950), a square-root law (Diaz-Cordoves & Gimenez 1992), a power-2 law (Hestroffer 1997), or a four-coefficients law (Claret 2000), so that when modeling the transit light curves, limb darkening is parameterized by some set of coefficients. Ideally, these coefficients should be constrained by the stellar modeling. Consequently, many libraries of limb-darkening coefficients covering a wide range of effective temperatures (Teff), surface gravity (log g), and metal-licities (M\/H) have been produced (e.g., Claret 2000; Sing 2010; Claret & Bloemen 2011; Magic et al. 2015) using various radiative transfer codes, such as ATLAS (Kurucz 1993), NextGen (Hauschildt et al. 1999), PHOENIX (Husser et al. 2013), MARCS (Gustafsson et al. 2008), and STAGGER (Magic et al. 2013). However, the limb-darkening parameters diverge between different libraries and often lead to an inadequate quality of fits to the observed transit profiles (Csizmadia et al. 2013). First, this can be due to errors in limb-darkening coefficients introduced by interpolation from the grid of stellar parameters used in these libraries to the actual stellar fundamental parameters. Second, available stellar calculations may not treat mechanisms that affect limb darkening with sufficient accuracy, for instance, convection (Pereira et al. 2013; Chiavassa et al. 2017) or magnetic activity (Csizmadia et al. 2013). Therefore, limb-darkening coefficients are often left as free parameters in a least-squares fit to observed light curves (e.g., Southworth 2008; Claret 2009; Cabrera et al. 2010; Gillon et al. 2010; Csizmadia et al. 2013; Maxted 2018). While this method usually leads to a good quality fit to observed transit profiles, it introduces additional free parameters, resulting in possible biases and degeneracies in the returned planetary radii (Espinoza & Jord\u00e1n 2015; Morello et al. 2017b). The way to reduce these biases and reliably determine planetary radii is to improve theoretical computations of stellar limb darkening.","Citation Text":["Gustafsson et al. 2008"],"Functions Text":["Consequently, many libraries of limb-darkening coefficients covering a wide range of effective temperatures (Teff), surface gravity (log g), and metal-licities (M\/H) have been produced","using various radiative transfer codes, such as","MARCS"],"Functions Label":["Background","Background","Background"],"Citation Start End":[[1111,1133]],"Functions Start End":[[713,897],[971,1018],[1104,1109]]}
{"Identifier":"2022MNRAS.511.2885B__Lu_et_al._2013_Instance_1","Paragraph":"One of our key results is that accounting for a complete stellar mass function can explain the post-starburst preference of TDE hosts. The alternative, currently most widespread explanation is that post-starburst galaxies are initially born with an extremely steep density profile (Stone et al. 2018). Fig. 6 compares the initial conditions adopted in our runs with the ultrasteep initial conditions adopted by Stone et al. (2018) to explain the post-starburst preference, rescaled to our MBH mass according to their footnote 12 (see also the caption of our Fig. 6). The central density in our and their profile differs by several orders of magnitude, and in general it is not straightforward to imagine how star formation could generate such an extremely steep cusp. In fact, Sanders (1998) points out that, too close to the MBH, molecular clouds should get disrupted by its gravitational field preventing star formation; even if the Milky Way nucleus features a population of young stars very close to the MBH (Morris 1993; Ghez et al. 2003; Lu et al. 2013), supporting instead the idea of possible central star formation, it seems unlikely it can develop such steep profiles. For instance, in these systems the time-scale for stellar collisions and mergers can be shorter than the relaxation time-scale at small radii (10\u22122 pc, Stone et al. 2018). Note that if the density profile is initially very steep, relaxation processes render it milder in time, as the equilibrium configuration is the Bahcall & Wolf (1976) solution (\u03c1\u221dr\u22121.75): this means that the ultrasteep configuration should be in place since the starburst stage, i.e. stars should be basically born with \u03c1\u221dr\u2212\u03b3, \u03b3 = 2.5\u20132.75. In order to compare our results with the ones by Stone et al. (2018), we evolve their (monochromatic) profiles shown in Fig. 6 with \u03c1\u221dr\u22122.5, r\u22122.75 and we compare the obtained rates with our complete mass function models in Fig. 7. Even if their rates are overall larger, the early to late event rate is compatible with what we find for much milder profiles featuring a complete mass function; in fact, the ratio between the average TDE rate in the range 250\u2013750 Myr (i.e. roughly the age of post-starburst systems) and 10\u201312 Gyr is equal to 8\u201314 for the Stone et al. (2018) cases, i.e. they are compatible with or smaller than the same ratios computed in our runs (see numbers in parenthesis in Table 1).","Citation Text":["Lu et al. 2013"],"Functions Text":["In fact, Sanders (1998) points out that, too close to the MBH, molecular clouds should get disrupted by its gravitational field preventing star formation; even if the Milky Way nucleus features a population of young stars very close to the MBH","supporting instead the idea of possible central star formation, it seems unlikely it can develop such steep profiles."],"Functions Label":["Background","Similarities"],"Citation Start End":[[1044,1058]],"Functions Start End":[[768,1011],[1061,1178]]}
{"Identifier":"2018AandA...614A..48B__Keselman_&_Nusser_2012_Instance_1","Paragraph":"The driving mechanisms and chronology of the buildup of bulges in late-type galaxies (LTGs) is an issue of key relevance to our understanding of galaxy evolution. According to our current knowledge on bulge demographics in the local universe, a large fraction of LTGs host pseudo-bulges (PBs; e.g., Gadotti 2009; Fisher & Drory 2011; Fern\u00e1ndez Lorenzo et al. 2014) that substantially differ from classical bulges (CBs) in their spectrophotometric and kinematical characteristics. The latter resemble in many respects \u201cold and dead\u201d elliptical galaxies, lacking ongoing star-formation (SF), exhibit a spheroidal shape with inwardly steeply increasing surface brightness profiles (SBPs) being well approximated by the S\u00e9rsic (1963) fitting law with a high (\u22733) exponent \u03b7, show stellar kinematics dominated by velocity dispersion (\u03c3\u22c6) and obey the Kormendy (1977) scaling relations for normal elliptical galaxies (Fisher & Drory 2010). It is observationally established that CBs contain a super-massive black hole (SMBH) with a mass M\u2219 tightly correlating with their stellar mass \n\n${\\cal M}_{\\star,\\textrm{B}}, \\sigma_{*}$M\u22c6,B,\u03c3*\n and optical luminosity (Ho 2008; Kormendy & Ho 2013; see also Ferrarese & Merritt 2000). Traditionally, bulges were thought to invariably form early-on via violent quasi-monolithic gas collapse (Larson 1974) or mergers (Bender et al. 1992; Aguerri et al. 2001; Keselman & Nusser 2012) associated with vigorous nuclear starbursts (Okamoto 2012), with the disk gradually building up around them. Whereas this inside-out galaxy formation scenario appears consistent with important integral characteristics of CBs (e.g., their red colors), it does not offer a plausible explanation for the presence of PBs in present-day LTGs. These generally show ongoing SF, a significant degree of rotational support (Kormendy & Kennicutt 2004, for a review) and flatter\/ellipsoidal shapes with nearly exponential SBPs (\u03b7\u22722; e.g., Drory & Fisher 2007; Fisher & Drory 2010). Even though there is observational evidence that PBs also contain a SMBH (Kormendy et al. 2011; Kormendy & Ho 2013), in some cases revealing itself as an active galactic nucleus (AGN; e.g., Kotilainen et al. 2016; see Kormendy & Ho 2013 for a review), these do not follow the M\u2219 \u2013\u03c3* correlation for CBs, which appears to be consistent with a different formation route. Indeed, the prevailing concept on PB formation is that these entities emerge gradually out of galactic disks through gentle gas inflow spawning quasi-continuous SF and the emergence of a central bulge-like luminosity excess at their centers (e.g., Courteau et al. 1996; Carollo et al. 2001; Kormendy & Kennicutt 2004). Besides bar-driven gas inflow (e.g., Springel & Hernquist 2005), various other mechanisms, such as inward stellar migration, minor mergers with low-mass satellites, or a purely dynamical re-arrangement of the disk (Scannapieco et al. 2010; Guedes et al. 2013; Bird et al. 2012; Roskar et al. 2012; Grand et al. 2014; Halle et al. 2015) have been proposed as further contributors to PB growth along the Gyr-long secular evolution of LTGs.","Citation Text":["Keselman & Nusser 2012"],"Functions Text":["Traditionally, bulges were thought to invariably form early-on via","or mergers","with the disk gradually building up around them.","Whereas this inside-out galaxy formation scenario appears consistent with important integral characteristics of CBs (e.g., their red colors), it does not offer a plausible explanation for the presence of PBs in present-day LTGs."],"Functions Label":["Background","Background","Background","Compare\/Contrast"],"Citation Start End":[[1391,1413]],"Functions Start End":[[1219,1285],[1338,1348],[1475,1523],[1524,1752]]}
{"Identifier":"2016AandA...588A..44Y__Jones_et_al._2014_Instance_3","Paragraph":"The second issue concerns the fact that inside a given region, coreshine is not detected in all the dense clouds observed by Paladini (2014) and Lef\u00e8vre et al. (2014) and that the proportion of clouds exhibiting coreshine varies from one region to another. For instance, 75% of the dense clouds detected in Taurus exhibit coreshine, whereas in most other regions the proportion is closer to 50% (such as Cepheus, Chamaeleon, and Musca)5. On the contrary, there are for instance very few detections in the Orion region. In THEMIS, most of the scattering efficiency originates in the accretion of an a-C:H mantle. This leads to three possible explanations for the absence of detectable coreshine. The first explanation is related to the amount of carbon available in the gas phase. The abundance used by K\u00f6hler et al. (2015) relies on the highest C depletion measurements made by Parvathi et al. (2012) towards regions with \\hbox{$N_{\\rm H} \\geqslant 2 \\times 10^{21}$}NH\u2a7e 2 \u00d7 1021 H\/cm2. Parvathi et al. (2012) highlighted the variability in the carbon depletion in dust depending on the line of sight. Thus, there may be clouds were the amount of carbon available for a-C:H mantle formation is smaller or even close to zero: such regions would be populated with aggregates with a thinner H-rich carbon mantle or no second mantle at all and thus exhibit very little or no coreshine emission. A second explanation is related to the stability of H-rich carbon in the ISM, which depends strongly on the radiation field intensity to local density ratio (Godard et al. 2011; Jones et al. 2014). In low-density regions (according to Jones et al. 2014, \\hbox{$A_{V} \\leqslant 0.7$}AV\u2a7d 0.7 for the standard ISRF), UV photons are responsible for causing the photo-dissociation of CH bonds, a-C:H \u2192 a-C. In transition regions from diffuse ISM to dense clouds (Jones et al. 2014, \\hbox{$0.7 \\leqslant A_{V} \\leqslant 1.2$}0.7 \u2a7d AV\u2a7d 1.2 for the standard ISRF), better shielded from UV photons and where the amount of hydrogen is significantly higher, H-poor carbon can be transformed into H-rich carbon through H atom incorporation, a-C \u2192 a-C:H. Similarly, carbon accreted from the gas phase in these transition regions is likely to be and stay H-rich. Then, in the dense molecular clouds, most of the hydrogen is in molecular form and thus not available to produce a-C:H mantles on the grains. However, this approximately matches the density at which ice mantles start to accrete on the grains, which would partly protect a-C:H layers that had formed earlier (Godard et al. 2011, and references therein). The stability and hydrogenation degree of a-C:H, as well as the exact values of AV thresholds, are both dependent on the timescale and UV field intensity. The resulting a-C \u2194 a-C:H delicate balance could explain why in a quiet region such as Taurus most of the clouds exhibit coreshine, whereas in Orion, where on average the radiation field intensity and hardness are much higher, most clouds do not. A third explanation is related to the age and\/or density of the clouds. In a young cloud, where dust growth is not advanced, or in an intermediate density cloud (\u03c1C ~ a few 103 H\/cm3), the dust population may be dominated by CMM grains instead of AMM(I) dust. Such clouds would be as bright in the IRAC 8 \u03bcm band as in the two IRAC bands at 3.6 and 4.5 \u03bcm, thus not matching the selection criteria defined by Pagani et al. (2010) and Lef\u00e8vre et al. (2014) and would be classified as \u201cno coreshine\" clouds. ","Citation Text":["Jones et al. 2014"],"Functions Text":["In transition regions from diffuse ISM to dense clouds","\\hbox{$0.7 \\leqslant A_{V} \\leqslant 1.2$}0.7 \u2a7d AV\u2a7d 1.2 for the standard ISRF), better shielded from UV photons and where the amount of hydrogen is significantly higher, H-poor carbon can be transformed into H-rich carbon through H atom incorporation, a-C \u2192 a-C:H. Similarly, carbon accreted from the gas phase in these transition regions is likely to be and stay H-rich. Then, in the dense molecular clouds, most of the hydrogen is in molecular form and thus not available to produce a-C:H mantles on the grains."],"Functions Label":["Background","Background"],"Citation Start End":[[1849,1866]],"Functions Start End":[[1793,1847],[1868,2381]]}
{"Identifier":"2021MNRAS.503.6155C__Lovisari_et_al._2017_Instance_2","Paragraph":"Galaxy clusters are the traces of the formation of the largest structures in the Universe and so reliable tools to investigate structures formation and evolution. In principle, this is possible only if and when we have full knowledge of the properties of these objects. The total mass (i.e. the total amount of the dark matter (DM), the intracluster medium (ICM), and the stellar components) is an invaluable quantity when exploring the abundances of clusters along the redshift: a standard way to infer cosmological parameters such as the mean matter density \u03a9m and the amplitude of matter perturbations \u03c38(Planck Collaboration XIII 2016). Furthermore, under the assumption of a simple self-similar model (Kaiser 1986; Voit 2005), we could derive the total mass of the clusters from a few observables in optical, X-ray, or millimetre band (Giodini et al. 2013). This approach results in a few scaling relations valuable when we are interested to obtain averaged results based on some statistics. However, it is prone to the assumed simplified approximations: hydrostatic equilibrium and isothermal and spherical distribution for DM and ICM (Bryan & Norman 1998). It is well known that the hydrostatic equilibrium in haloes is not always satisfied, due to non-thermal pressure contributions from internal motions and turbulence (see e.g. Fang, Humphrey & Buote 2009; Lau, Kravtsov & Nagai 2009; Lagan\u00e1, de Souza & Keller 2010; Rasia et al. 2012; Nelson, Lau & Nagai 2014; Yu, Nelson & Nagai 2015; Biffi et al. 2016; Eckert et al. 2019; Angelinelli et al. 2020; Ansarifard et al. 2020; Gianfagna et al. 2020; Green et al. 2020), pointing out the impact that the dynamical state of those large gravitational bounded objects should have. Several attempts have been made to infer clusters dynamical state, using both observational data and simulations, by analysing the images of the emission in optical (see e.g. Ribeiro, Lopes & Rembold 2013; Wen & Han 2013) and in the X-ray band (see e.g. Rasia, Meneghetti & Ettori 2013; Lovisari et al. 2017; Nurgaliev et al. 2017; Bartalucci et al. 2019; Cao, Barnes & Vogelsberger 2020; Yuan & Han 2020) or of the diffusion of the cosmic microwave background (CMB) photons by thermal Sunyaev\u2013Zel\u2019dovich (tSZ) effect in the millimetre band (Cialone et al. 2018; De Luca et al. 2020, hereafter DL20), or a combination of some of them (see e.g. Mann & Ebeling 2012; Molnar, Ueda & Umetsu 2020; Ricci et al. 2020; The CHEX-MATE Collaboration 2020; Zenteno et al. 2020). Among the possibilities, we have to mention the studies of the clusters morphology in X-ray and tSZ maps. Several indicators are commonly used, such as asymmetry parameter (Schade et al. 1995), light concentration (Santos et al. 2008), third-order power ratio (Buote & Tsai 1995; Wei\u00dfmann et al. 2013), centroid shift (Mohr, Fabricant & Geller 1993; O\u2019Hara et al. 2006), strip parameter, Gaussian fit parameter (Cialone et al. 2018), and so on. They exploit the maps with different apertures and efficiencies and are applied individually or combined together, even with different weights (see e.g. B\u00f6hringer et al. 2010; Nurgaliev et al. 2013; Rasia et al. 2013; Wei\u00dfmann et al. 2013; Mantz et al. 2015; Cui et al. 2016; Lovisari et al. 2017; Cialone et al. 2018; Cao et al. 2020; DL20; Yuan & Han 2020). A complementary approach is by applying thresholds on specific thermodynamic variables. Among the others, the central electron gas density and the core entropy are fairly reliable (Hudson et al. 2010). The azimuthal scatter in radial profiles of gas density, temperature, entropy, or surface brightness (Vazza et al. 2011) is also used as a proxy of the ICM inhomogeneities and correlated to the clusters dynamical state (see e.g. Roncarelli et al. 2013; Ansarifard et al. 2020). Alternatively, the projected sky separations between key positions in the images are resulting in reliable estimators of the dynamical state. Interestingly, the offsets between the bright central galaxy (BCG) and the peaks and\/or the centroids of X-ray or tSZ maps are an indication of how much the relaxation condition is satisfied, with different efficiency (see e.g. Jones & Forman 1984; Katayama et al. 2003; Lin & Mohr 2004; Sanderson, Edge & Smith 2009; Mann & Ebeling 2012; Rossetti et al. 2016; Lopes et al. 2018; DL20; Ricci et al. 2020; Zenteno et al. 2020). To be mentioned also other approaches based on wavelets analysis (Pierre & Starck 1998), on the Minkowski functionals (Beisbart, Valdarnini & Buchert 2001), or on machine learning (see e.g. Cohn & Battaglia 2019; Green et al. 2019; Gupta & Reichardt 2020).","Citation Text":["Lovisari et al. 2017"],"Functions Text":["They exploit the maps with different apertures and efficiencies and are applied individually or combined together, even with different weights (see e.g."],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[3224,3244]],"Functions Start End":[[2948,3100]]}
{"Identifier":"2019ApJ...881...42J__J\u00f8rgensen_1999_Instance_1","Paragraph":"Stellar population evolution studies beyond z \u2248 1 have primarily focused on ages through studies of luminosity changes. Beifiori et al. (2017) used new data for 19 galaxies in z = 1.3\u20131.6 clusters obtained with the Very Large Telescope\/KMOS to extend the redshift coverage of the results regarding the evolution of the mass-to-light (M\/L) ratios of bulge-dominated passive galaxies. The authors used their new results together with the available literature results covering up to z = 1.3 (van Dokkum & Franx 1996; J\u00f8rgensen et al. 1999, 2006, 2014; Kelson et al. 2000; Wuyts et al. 2004; Holden et al. 2005, 2010; Barr et al. 2006; van Dokkum & van der Marel 2007; Saglia et al. 2010; J\u00f8rgensen & Chiboucas 2013) and low-redshift reference data for the Coma cluster (J\u00f8rgensen 1999; J\u00f8rgensen et al. 2006) to further solidify the evidence supporting passive evolution and a formation redshift zform \u2248 2. The formation redshift should be understood as the epoch of the last major star formation episode. At z \u2248 1 the massive (Mass > 1011 M) bulge-dominated galaxies in clusters appear to be in place and mostly passively evolving. Lower mass galaxies may still be added to the red sequence and from then on passively evolve (e.g., S\u00e1nchez-Bl\u00e1zquez et al. 2009; Choi et al. 2014), but see also Cerulo et al. (2016) for results supporting that the red sequence well below L\u22c6 is fully populated in rich clusters already at \n\n\n\n\n\n. Ultimately, the properties of galaxies mapped over a large fraction of the age of the universe, may constrain the models for building the galaxies. It is difficult to understand within the prevailing hierarchical model favored by the \u039bCDM (cold dark matter) cosmology, the existence of such massive passive galaxies with relatively old stellar populations at z \u2248 1, while less massive galaxies appear to harbor younger stellar populations, e.g., J\u00f8rgensen et al. (2017, and references therein), see Kauffmann et al. (2003) for a discussion of this tension between the observational results and the hierarchical models of galaxy formation. However, more recent cosmological simulations like Illustris (Genel et al. 2014; Vogelsberger et al. 2014; Wellons et al. 2015) and UniverseMachine (Behroozi et al. 2019) find that massive quiescent galaxies can be in place by z \u2273 2.","Citation Text":["J\u00f8rgensen 1999"],"Functions Text":["The authors used their new results together with","and low-redshift reference data for the Coma cluster","to further solidify the evidence supporting passive evolution and a formation redshift zform \u2248 2."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[767,781]],"Functions Start End":[[383,431],[713,765],[806,903]]}
{"Identifier":"2020ApJ...899L..10F__Tomida_et_al._2013_Instance_1","Paragraph":"Understanding of the low-mass star formation process has been intensively studied from decades ago, mainly by theoretical work at the beginning (e.g., Larson 1969; Shu 1977; Shu et al. 1987; Inutsuka 2012). They explained that the fragmentation and condensation of molecular clouds result in forming dense cores, which undergo gravitational collapse to form stars. According to the theoretical studies, dense cores eventually harbor the first protostellar cores, the first quasi-hydrostatic object during the star formation process (hereafter the first core; e.g., Larson 1969; Masunaga et al. 1998; Tomida et al. 2013), which provide the initial condition of star formation. Recent magnetohydrodynamic (MHD) simulations demonstrated that the first core with a size of 1\u2013100 au is formed when the central density exceeds \u223c1010 cm\u22123 via the gravitational collapse. The first core is suggested to have a low-velocity (1\u201310 km s\u22121) molecular outflow with a wide opening angle (Machida et al. 2008), which is qualitatively different from the collimated jet driven by a mature protostar. However, it is difficult to identify such an object observationally because the first core has a short lifetime, 103\u2013104 yr, depending on the physical condition of the parental core (Tomida et al. 2010), and does not show bright infrared emission. Although some candidates of the first core were already reported in the past decade (e.g., Chen et al. 2010, 2012; Enoch et al. 2010; Pineda et al. 2011; Pezzuto et al. 2012; Hirano & Liu 2014), the first-core phase is not fully explored observationally and it is supposed to still be the missing link between the isothermal and adiabatic collapse (i.e., prestellar and protostellar core). To search for candidates of the first core, it is essential to perform a survey-type observation toward a large number of starless cores. According to the early dense core survey with an average density of \u2273105 cm\u22123 (Onishi et al. 2002), the lifetime of the starless phase is \u223c4 \u00d7 105 yr (see also Ward-Thompson et al. 2007). The simple calculation tells us that only one out of a few \u00d7 10\u2013100 cores harbors the first core(s).","Citation Text":["Tomida et al. 2013"],"Functions Text":["According to the theoretical studies, dense cores eventually harbor the first protostellar cores, the first quasi-hydrostatic object during the star formation process (hereafter the first core; e.g.,","which provide the initial condition of star formation."],"Functions Label":["Background","Background"],"Citation Start End":[[600,618]],"Functions Start End":[[365,564],[621,675]]}
{"Identifier":"2020AandA...641A.155V__Puglisi_et_al._2019_Instance_3","Paragraph":"The scenario presented above has been formulated in various flavors to individually explain several of the properties reported here. The main addition of this work, namely the excitation of CO in distant main-sequence and starburst galaxies, fits in the general picture that we sketched. The ensemble of properties and correlations that we reported here can be also used to revisit the definition of what a starburst is. A standard operational classification is based on the distance from the observed empirical M\u22c6-SFR correlation, the main sequence. This proved to be a useful distinction and an excellent predictor of several trends (e.g., Sargent et al. 2014), but recent results, including our present and previous analysis (Puglisi et al. 2019), show that the demarcation between starburst and main-sequence galaxies is more blurred that we previously considered. We do detect starburst-like behaviors in galaxies on the main sequence (Elbaz et al. 2018), likely linked to the existence of transitional objects (Popping et al. 2017; Barro et al. 2017b; G\u00f3mez-Guijarro et al. 2019; Puglisi et al. 2019, and in prep. to limit the references to recent works based on submillimeter observations). Such transition might well imply an imminent increase of the SFR, driving the object in the realm of starbursts (e.g., Barro et al. 2017b), or its cessation, bringing the system back onto or even below the main sequence (G\u00f3mez-Guijarro et al. 2019; Puglisi et al. 2019), with the CO properties potentially able to distinguish between these two scenarios. Regardless of these transitional objects, a definition of starburst based on \u03a3SFR, rather than \u0394MS, would naturally better account for the observed molecular gas excitation properties, dust temperatures and opacities, or SFE (see also Elbaz et al. 2011; Rujopakarn et al. 2011; Jim\u00e9nez-Andrade et al. 2018; Tacconi et al. 2020). As an example, in Fig. 8 we show the mean SLED of the subsample of galaxies with both CO (2\u2005\u2212\u20051) and CO (5\u2005\u2212\u20054) coverage, split at its median \u03a3SFR. While only tentative at this stage, this suggests a trend of increasing CO excitation with \u03a3SFR, consistently with Fig. 7 and what mentioned above.","Citation Text":["Puglisi et al. 2019"],"Functions Text":["Such transition might well imply an imminent increase of the SFR, driving the object in the realm of starbursts","or its cessation, bringing the system back onto or even below the main sequence","with the CO properties potentially able to distinguish between these two scenarios."],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Future Work"],"Citation Start End":[[1447,1466]],"Functions Start End":[[1198,1309],[1338,1417],[1469,1552]]}
{"Identifier":"2016ApJ...832..195N__Jin_et_al._2012_Instance_2","Paragraph":"We ignore the density stratification effect in Case I, II, and IIa, because the width of the horizontal current sheet in our simulations is much shorter than the length. The simulation domain extends from x = 0 to x = L0 in the x-direction, and from \n\n\n\n\ny\n=\n\u2212\n0.5\n\n\nL\n\n\n0\n\n\n\n\n to \n\n\n\n\ny\n=\n0.5\n\n\nL\n\n\n0\n\n\n\n\n in the y-direction, in the three cases, with \n\n\n\n\n\n\nL\n\n\n0\n\n\n=\n\n\n10\n\n\n6\n\n\n\n\n m. Outflow boundary conditions are used in the x-direction and inflow boundary conditions in the y-direction. For the inflow boundary conditions, the fluid is allowed to flow into the domain but not to flow out; the gradient of the plasma density vanishes; the total energy is set such that the gradient in the thermal energy density vanishes; a vanishing gradient of parallel components plus divergence-free extrapolation of the magnetic field. For the outflow boundary conditions, the fluid is allowed to flow out of the domain but not to flow in, and the other variables are set by using the same method as the inflow boundary conditions. The horizontal force-free Harris current sheet is used as the initial equilibrium configuration of magnetic fields in Case I,\n13\n\n\n\n\n\n\nB\n\n\nx\n0\n\n\n=\n\u2212\n\n\nb\n\n\n0\n\n\ntanh\n[\ny\n\n\/\n\n(\n0.05\n\n\nL\n\n\n0\n\n\n)\n]\n\n\n\n\n14\n\n\n\n\n\n\nB\n\n\ny\n0\n\n\n=\n0\n\n\n\n\n15\n\n\n\n\n\n\nB\n\n\nz\n0\n\n\n=\n\n\nb\n\n\n0\n\n\n\n\/\n\ncosh\n[\ny\n\n\/\n\n(\n0.05\n\n\nL\n\n\n0\n\n\n)\n]\n.\n\n\nThe magnetic fields in the low solar atmosphere could be very strong (Jin et al. 2009, 2012; Khomenko et al. 2014; Peter et al. 2014; Vissers et al. 2015) and the magnetic field can exceed 0.15 T in both the intranetwork and the network quiet region (e.g., Orozco Su\u00e1rez et al. 2007; Mart\u00ednez Gonz\u00e1lez et al. 2008; Jin et al. 2009, 2012). In the work by Jin et al. 2012, the maximum of the field strength was found to be 0.15 T. The magnetic field could be even stronger in the active region near the sunspot. Therefore, we set b0 = 0.05 T in Case I and Case II, and b0 = 0.15 T in Case IIa. Due to the force-freeness and neglect of gravity, the initial equilibrium thermal pressure is uniform. The initial temperature and plasma density are set as T0 = 4200 K and \u03c10 = 1.66057 \u00d7 10\u22126 kg m\u22123 in Case I, and T0 = 4800 K and \u03c10 = 3.32114 \u00d7 10\u22125 kg m\u22123 in Case II and Case IIa. Therefore, the initial plasma \u03b2 is calculated as \u03b2 \u2243 0.0583 in Case I, \u03b2 \u2243 1.332 in Case II, and \u03b2 \u2243 0.148 in Case IIa. The initial ionization degree is assumed as Yi = 10\u22123 in Case I, and Yi = 1. 2 \u00d7 10\u22124 in Case II and IIa. The magnetic diffusion in this work matches the form computed from the solar atmosphere model in Khomenko & Collados (2012), and we set \n\n\n\n\n\u03b7\n=\n\n[\n\n5\n\u00d7\n\n\n10\n\n\n4\n\n\n\n\n(\n4200\n\n\/\n\nT\n)\n\n\n1.5\n\n\n+\n1.76\n\u00d7\n\n\n10\n\n\n\u2212\n3\n\n\n\n\nT\n\n\n0.5\n\n\n\n\nY\n\n\ni\n\n\n\u2212\n1\n\n\n\n]\n\n\n\n m2 s\u22121 in Case I, and \n\n\n\n\n\u03b7\n=\n[\n5\n\u00d7\n\n\n10\n\n\n4\n\n\n\n\n(\n4800\n\n\/\n\nT\n)\n\n\n1.5\n\n\n\n+\n1.76\n\n\u00d7\n\n\n\n10\n\n\n\u2212\n3\n\n\n\n\nT\n\n\n0.5\n\n\n\n\nY\n\n\ni\n\n\n\u2212\n1\n\n\n]\n\n\n m2 s\u22121 in Case II and IIa. The first part \u223c T\u22121.5 is contributed by collisions between ions and electrons, the second part \n\n\n\n\n\u223c\n\n\nT\n\n\n0.5\n\n\n\n\nY\n\n\ni\n\n\n\u2212\n1\n\n\n\n\n is contributed by collisions between electrons and neutral particles. Small perturbations for both magnetic fields and velocities at t = 0 make the current sheet to evolve and secondary instabilities start to appear later in the three cases. The forms of perturbations are listed below:\n16\n\n\n\n\n\n\nb\n\n\nx\n1\n\n\n=\n\u2212\npert\n\u00b7\n\n\nb\n\n\n0\n\n\n\u00b7\nsin\n\n\n2\n\u03c0\n\n\n\ny\n+\n0.5\n\n\n\nL\n\n\n0\n\n\n\n\n\n\nL\n\n\n0\n\n\n\n\n\n\n\n\u00b7\ncos\n\n\n2\n\u03c0\n\n\n\nx\n+\n0.5\n\n\nL\n\n\n0\n\n\n\n\n\n\nL\n\n\n0\n\n\n\n\n\n\n\n\n\n\n\n17\n\n\n\n\n\n\nb\n\n\ny\n1\n\n\n=\npert\n\u00b7\n\n\nb\n\n\n0\n\n\n\u00b7\ncos\n\n\n2\n\u03c0\n\n\n\ny\n+\n0.5\n\n\nL\n\n\n0\n\n\n\n\n\n\nL\n\n\n0\n\n\n\n\n\n\n\n\u00b7\nsin\n\n\n2\n\u03c0\n\n\n\nx\n+\n0.5\n\n\nL\n\n\n0\n\n\n\n\n\n\nL\n\n\n0\n\n\n\n\n\n\n\n\n\n\n\n18\n\n\n\n\n\n\nv\n\n\ny\n1\n\n\n=\n\u2212\npert\n\u00b7\n\n\nv\n\n\nA\n0\n\n\n\u00b7\nsin\n\n\n\u03c0\n\n\n\ny\n\n\n\n\nL\n\n\n0\n\n\n\n\n\n\n\n\u00b7\n\n\n\n\n\nrandom\n\n\nn\n\n\n\n\nMax\n\n(\n\n\u2223\n\n\nrandom\n\n\nn\n\n\n\u2223\n\n)\n\n\n\n\n,\n\n\nwhere pert = 0.08, vA0 is the initial Alfv\u00e9n velocity, randomn is the random noise function in our code, and \n\n\n\n\nMax\n(\n\u2223\n\n\nrandom\n\n\nn\n\n\n\u2223\n)\n\n\n is the maximum of the absolute value of the random noise function. This random noise function makes the initial perturbations for the velocity in the y-direction to be asymmetric, and such an asymmetry makes the current sheet gradually become more tilted, especially after secondary islands appear. The reconnection process is not really symmetrical in nature (Murphy et al. 2012), this is one of the reasons that we use such a noise function. Another reason is that the asymmetric noise function makes the secondary instabilities develop faster. Figure 1(a) shows the distributions of the current density and magnetic fields at t = 0 in case I.","Citation Text":["Jin et al.","2012"],"Functions Text":["and the magnetic field can exceed 0.15 T in both the intranetwork and the network quiet region"],"Functions Label":["Uses"],"Citation Start End":[[1637,1647],[1654,1658]],"Functions Start End":[[1477,1571]]}
{"Identifier":"2015ApJ...802....8A__Susa_2013_Instance_1","Paragraph":"Early theoretical work on the formation of the first stars, mostly through high-resolution simulations, found that Pop III stars with mass MIII,* \u2273 100 M\u2299 are born in isolation inside minihalos, and thus the \u201cone massive Pop III star per minihalo\u201d paradigm was established (Abel et al. 2002; Bromm et al. 2002; Yoshida et al. 2006). Later, however, several higher-resolution simulations began to observe the formation of binary protostar systems with a smaller mass range, MIII,* \u2243 [10\u201340] M\u2299 (Turk et al. 2009; Stacy et al. 2010). While the universality of the latter finding is in doubt (Greif et al. 2012; Stacy  Bromm 2013; Susa 2013; Hirano et al. 2014; see also Becerra et al. 2015 for the formation of Pop III stars inside more massive halos), it certainly introduces very important subtleties to the old paradigm. One important aspect is that X-ray binary systems may remain after some of these stars die, and can emit X-rays very efficiently (we will quantify their X-ray emissivity in Section 3) through gas accretion rate comparable to the Eddington limit (Mirabel et al. 2011). This then could make the X-ray heating epoch occur earlier than previously thought, or even allow a model where the reionization is dominated by X-ray photons instead of UV photons. Reionization dominated by X-ray photons, due to their long mean free path, will occur much more smoothly in space than reionization by UV photons (e.g., Haiman 2011; Mesinger et al. 2013, and references therein). In addition, X-rays can heat the IGM (see, e.g., recent observational signature reported by Parsons et al. 2014), which impacts the dynamics of the IGM (e.g., Tanaka et al. 2012; Jeon et al. 2014) and \u03b4Tb of IGM (e.g., Mesinger et al. 2013; Fialkov et al. 2014; Jeon et al. 2014). High-redshift X-ray binaries seem to dominate the X-ray background over the active galactic nuclei at the late stage of EoR (z \u223c 6\u20138), if their emissivity and spectral energy distribution (SED) is calibrated by the observed, low-z (0 \u2264 z \u2272 4) X-ray binaries (Fragos et al. 2013).","Citation Text":["Susa 2013"],"Functions Text":["While the universality of the latter finding is in doubt","it certainly introduces very important subtleties to the old paradigm"],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[628,637]],"Functions Start End":[[532,588],[751,820]]}
{"Identifier":"2018AandA...610A..44M__Kr\u00fcger_&_Dreizler_(1992)_Instance_1","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)"],"Functions Text":["The most recent spectroscopic study is from","who reinvestigated the internal rotation measurements and also determined hyperfine coupling parameters due to the nitrogen quadrupole"],"Functions Label":["Background","Background"],"Citation Start End":[[720,744]],"Functions Start End":[[676,719],[745,879]]}
{"Identifier":"2021ApJ...908...15S__Scoville_et_al._2014_Instance_1","Paragraph":"Several methods are commonly used to estimate gas masses. The first method is using CO emission line fluxes (e.g., Daddi et al. 2010; Genzel et al. 2010; Tacconi et al. 2010, 2013). This method has uncertainties on the CO-to-H2 conversion factor, which changes depending on metallicity (Genzel et al. 2012), and on the CO excitation states when higher-J CO lines are used (e.g., Daddi et al. 2015; Riechers et al. 2020). Furthermore, observations of CO lines for galaxies at high redshifts are observationally expensive. The second method is converting a dust mass into a gas mass with an assumed gas-to-dust mass ratio (e.g., Santini et al. 2014; B\u00e9thermin et al. 2015). Because the gas-to-dust mass ratio depends on the metallicity (Leroy et al. 2011; R\u00e9my-Ruyer et al. 2014), metallicity measurements are crucial for estimating the gas mass accurately. Gas masses can also be estimated with an empirically calibrated relation between single-band dust continuum flux at the Rayleigh\u2013Jeans (R-J) tail and gas mass (e.g., Scoville et al. 2014, 2016; Groves et al. 2015). These scaling relations are calibrated with local galaxies or with local galaxies and SMGs up to z \u223c 2. In this method, the gas-to-dust mass ratio is included in the conversion factor, and therefore does not need to be considered. It remains unclear whether the empirical calibration methods are applicable to galaxies at z > 3 or how much scatter there is in the relationships. Given that dust continuum observations take much less time compared to the CO observations, using dust continuum as a tracer of gas has the advantage of increasing the number of galaxies at higher redshifts with individual gas estimates, but these will only be reliable when precise metallicities are available as well. Metallicity measurements based on rest-frame optical emission lines for dustier star-forming galaxies are thought to have larger uncertainties due to strong dust obscuration (e.g., Santini et al. 2010). Herrera-Camus et al. (2018) reported a discrepancy between metallicities derived with rest-frame optical emission lines and far-infrared (FIR) fine-structure lines for local (ultra) luminous infrared galaxies ((U)LIRGs).","Citation Text":["Scoville et al. 2014"],"Functions Text":["Gas masses can also be estimated with an empirically calibrated relation between single-band dust continuum flux at the Rayleigh\u2013Jeans (R-J) tail and gas mass (e.g.,","These scaling relations are calibrated with local galaxies or with local galaxies and SMGs up to z \u223c 2. In this method, the gas-to-dust mass ratio is included in the conversion factor, and therefore does not need to be considered.","It remains unclear whether the empirical calibration methods are applicable to galaxies at z > 3 or how much scatter there is in the relationships."],"Functions Label":["Background","Background","Compare\/Contrast"],"Citation Start End":[[1022,1042]],"Functions Start End":[[856,1021],[1071,1301],[1302,1449]]}
{"Identifier":"2017AandA...597A.114B__Bergin_&_Tafalla_(2007)_Instance_1","Paragraph":"The band-merged Hi-GAL product catalogue (Elia et al. 2016) is built as in Elia et al. (2013) and provides spectral energy distribution (SED) fit parameters to the individual clumps. The average angular size of the clumps is 25\u2032\u2032 at 250 \u03bcm. Using the heliocentric distances provided in the Hi-GAL product catalogue (described in Sect. 3.2.) and the SED fit parameters, the authors of the catalogue are able to provide linear diameters and masses of the clumps. The catalogue explores a wide range of linear diameters and masses, from sub-parsec (\u2264 0.1 pc) to parsec scale (1\u22125 pc) with masses from 1 M\u2299 to 105M\u2299. These wide ranges mean that we probably mix several types of objects, from single star-forming cores to clumps containing multiple cores, even to entire clouds, depending on the distance of object. Most of these sources, however, fulfil the definition of clump, according to the definition of Bergin & Tafalla (2007). Dust temperatures for these clumps have been estimated through a grey-body fit, and searched in the range T = 5\u221240 K. Most of them are found between 10 and 20 K. The appearance of the SED and the parameters obtained through the grey-body fit allow for classifying the evolutionary stage of these objects. Three stages are identified: starless unbound and bound (pre-stellar) objects, and proto-stellar objects. The pre- and proto-stellar stages are distinguished from each other by the presence of a 70 \u03bcm source in a proto-stellar clump (e.g.  Dunham et al. 2008; Ragan et al. 2012; Veneziani et al. 2013). The bound versus unbound identification is obtained by using the mass-radius relation, well known as Larson\u2019s third law, originally formulated as M(r) > 460 M\u2299(r\/ pc)1.9, with r the radius of the source (Larson 1981). Beyond 4\u22125 kpc two effects could lead to misclassifying the pre- and proto-stellar stages. First, different sensitivities of PACS and SPIRE could lead to missing a possible 70 \u03bcm counterpart of a source detected with SPIRE. Second, at large heliocentric distances, two or more pre- and proto-stellar sources could be detected as a single object as a result of lacking resolution, globally and simply labelled as proto-stellar. The first effect was partially mitigated by searching for a possible 70 \u03bcm counterpart that was not originally listed in the single-band catalogue through performing additional source detection at this band using a threshold less demanding than the initial one. Elia et al. (2016) provide statistics and a discussion about the ratio between pre- and proto-stellar clumps. The distribution of the three evolutionary stages is shown in a portion of the Galactic Plane in Fig. 1. Each panel represents an evolutionary stage in the longitude range 26 \u2264l\u2264 31 deg. The pre-stellar clumps are more extended in Galactic latitude than the proto-stellar clumps. The unbound clump distribution is hard to characterise because it is obscured by the proto-stellar clumps in the mid-plane, therefore we only consider clustered over-densities composed of pre- and proto-stellar clumps. These distributions are observed across the entire longitude range of this study. ","Citation Text":["Bergin & Tafalla (2007)"],"Functions Text":["The catalogue explores a wide range of linear diameters and masses, from sub-parsec (\u2264 0.1 pc) to parsec scale (1\u22125 pc) with masses from 1 M\u2299 to 105M\u2299. These wide ranges mean that we probably mix several types of objects, from single star-forming cores to clumps containing multiple cores, even to entire clouds, depending on the distance of object. Most of these sources, however, fulfil the definition of clump, according to the definition of"],"Functions Label":["Uses"],"Citation Start End":[[906,929]],"Functions Start End":[[461,905]]}
{"Identifier":"2022MNRAS.516.2641V__Soleri_&_Fender_2011_Instance_1","Paragraph":"The connection between the accretion flow and compact jets is routinely studied in the X-ray\u2013radio luminosity (LX\u2013LR) diagram. Here, the former traces the accretion luminosity and is a proxy for accretion rate, while the latter traces the jet and its brightness. Compact BH jets show an LX\u2013LR correlation across \u223ceight orders of magnitude in X-ray luminosity (Hannikainen et al. 1998; Corbel et al. 2000, 2003; Gallo et al. 2006); a subset of sources follows a radio-bright correlation with a power-law slope of \u03b2 \u2248 0.6 (where $L_R \\propto L_X^{\\beta }$, while others follow a steeper correlation with \u03b2 \u2273 1 at high X-ray luminosities, before re-joining the other track around LX \u2248 1035 erg\u2009s\u22121 (e.g. Coriat et al. 2011; Soleri & Fender 2011; Carotenuto et al. 2021, although discussion exists regarding the statistical robustness of this separation; e.g. Gallo, Degenaar & van den Eijnden 2018). The situation for NS LMXBs is even more complex. While, as a sample, NS LMXBs are \u223c22 times radio-fainter than BH systems, the individual NS systems do not appear to follow a single correlation; significant scatter exists both between sources and between outbursts of the same source. Furthermore, due to their radio faintness, few sources have been monitored over a large range of X-ray luminosity \u2013 particularly below LX \u2248 1035 erg\u2009s\u22121, few NS LMXBs have been detected in radio (e.g. Tudor et al. 2017; Gusinskaia et al. 2020). For the sample of NS LMXBs, however, a power-law slope of \u03b2 \u2248 0.4\u2013 \u22120.5 has been measured (Gallo et al. 2018), which is similar to BHs. Different radiative efficiencies may be expected for the two types of LMXBs\u2013BHs can advect a fraction of the liberated gravitational energy across the event horizon, while the presence of an NS surface implies that all this liberated energy should, eventually, be radiated. Therefore, a similarity in the LX\u2013LR coupling of the two source classes, which depends on this radiative efficiency, is surprising.","Citation Text":["Soleri & Fender 2011"],"Functions Text":["Compact BH jets show an LX\u2013LR correlation across \u223ceight orders of magnitude in X-ray luminosity","a subset of sources follows a radio-bright correlation with a power-law slope of \u03b2 \u2248 0.6 (where $L_R \\propto L_X^{\\beta }$, while others follow a steeper correlation with \u03b2 \u2273 1 at high X-ray luminosities, before re-joining the other track around LX \u2248 1035 erg\u2009s\u22121 (e.g."],"Functions Label":["Background","Background"],"Citation Start End":[[721,741]],"Functions Start End":[[263,358],[431,700]]}
{"Identifier":"2016MNRAS.462.3441D__Namouni_1999_Instance_5","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"],"Functions Text":["This is not the Kozai\u2013Lidov resonance; in this case, the Kozai\u2013Lidov domain (domain II in","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"],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Uses"],"Citation Start End":[[1851,1863]],"Functions Start End":[[1761,1850],[1865,2012],[2014,2024]]}
{"Identifier":"2020AandA...643A..92S__Huang_et_al._2018a_Instance_1","Paragraph":"Our models show that high signal to noise ratio (S\/N) spiral structures in the gas dynamics, together with a spiral seen in the dust continuum surface density, can put constraints on grain size comparing the observed amplitude of the spiral in the dust with that in the gas. There are multiple candidate sources that are known to have spirals in the continuum or the gas. However, to give an accurate estimate of the surface density and to derive a velocity residual map, we need data with high spatial and velocity resolution. Out of the recent high-resolution Disk Substructures at High Angular Resolution Project (DSHARP) survey of 20 disks, 3 were found to harbor spirals: Elias 27, WaOph 6, and IM Lup (Huang et al. 2018a). Elias 27 harbors the highest contrast spiral arms of the survey, but gravitational instability is a convincing explanation for this source (Pohl et al. 2015; Meru et al. 2017; Forgan et al. 2018, although see also Hall et al. 2018), which makes it unusable for this purpose, as further discussed in Sect. 5.3. Furthermore, Elias 27 and WaOph 6 suffer from cloud and outflow contamination in the CO gas (Andrews et al. 2018), which makes it impossible to use the kinematic data published so far. The IM Lup disk has two high contrast spiral arms in the millimeter continuum, and no cloud contamination (Andrews et al. 2018), so is an ideal candidate to analyze. Unfortunately, no spiral signal is detected in the gas for the IM Lup disk, but in Sect. 4.3 we will show that we can make a prediction on what the spiral looks like in the two velocity components. Other sources are considered by surveying the literature. MWC 758 is a promising source (Dong et al. 2018; Boehler et al. 2018), but the observations are currently too limited to be used. HD 100453 has a spiral detected in scattered light, continuum, and CO gas (Benisty et al. 2017; Rosotti et al. 2020), but the spiral seen in the gas is beyond the size of the continuum disk and the inner region of the disk has a noisy velocity map, potentially due to a warp in the inner disk, which makes the source not useful in our analysis. TW Hya harbors the only spirals in gas dynamics, potentially in the azimuthal velocity (Teague et al. 2019). The spirals are detected in the CO gas emission as well, but no spiral signal is present inthe dust continuum images. In the next subsection, we will analyze the TW Hya disk further and show that the observed spiral signal in the dynamics is interesting for further study, but not useful to test our model.","Citation Text":["Huang et al. 2018a"],"Functions Text":["Out of the recent high-resolution Disk Substructures at High Angular Resolution Project (DSHARP) survey of 20 disks, 3 were found to harbor spirals: Elias 27, WaOph 6, and IM Lup"],"Functions Label":["Uses"],"Citation Start End":[[708,726]],"Functions Start End":[[528,706]]}
{"Identifier":"2016MNRAS.462.1415C__Bolzonella_et_al._2010_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":["Bolzonella et al. 2010"],"Functions Text":["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.","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"],"Functions Label":["Background","Compare\/Contrast"],"Citation Start End":[[1410,1432]],"Functions Start End":[[1119,1409],[1825,2013]]}
{"Identifier":"2018MNRAS.477.3520L__Abolfathi_et_al._2018_Instance_2","Paragraph":"Over time, the data releases have treated the Balmer line regions in different ways. The presence of the artificial curvature was first reported by Busca et al. (2013) in the context of the DR9 data release. To minimize this effect, a different scheme was used in DR12 (Alam et al. 2015, see their table 2) by using a linear function (instead of an iterative b-spline procedure) to interpolate the flux over the masked regions. Surprisingly, we observe that this data reduction change was only applied to the Balmer \u03b2, \u03b3, and \u03b4 lines but not applied to the Balmer \u03b1 line. For the latter one, the original iterative spline interpolation scheme has been used all along. As a result, an absorption-like feature at the location of Balmer \u03b1 is found in SDSS data releases 9 up to now, i.e. the latest data release 14 (Abolfathi et al. 2018). To illustrate this, we show examples of calibration vectors for SDSS BOSS DR9 (Ahn et al. 2012; Dawson et al. 2013), DR12 (Alam et al. 2015), eBOSS DR14 (Dawson et al. 2016; Abolfathi et al. 2018) data release as well as calibration vectors for the MaNGA survey (Bundy et al. 2015) DR14 data release (Abolfathi et al. 2018) in Fig. 6. In the upper panel, we show small-scale features in calibration vectors from the same plate (number 3647) processed by the DR9 (blue), DR12 (green), and DR14 (red) pipelines. The spectra are obtained by normalizing the large-scale features in calibration vectors with a cubic b-spline with break points separated by 50 \u00c5. The three spectra are mostly identical while the smooth curving features at the wavelengths of Balmer series in DR9 disappear in DR12 and DR14. From DR12, the pipeline corrects the features with linear interpolation across wavelength regions with Balmer \u03b2, \u03b3, and \u03b4 lines (Alam et al. 2015). One can also observe the change in the median residual spectra of DR9 and DR12 in Fig. 1. However, the H\u03b1 feature remains uncorrected from DR9 to DR14. We also show a DR7 calibration vector (black) in which most of the wiggles are absent. As pointed previously, this is due to the fact that the DR7 pipeline interpolates the calibration vectors using an effective scale larger than that used in subsequent data releases.","Citation Text":["Abolfathi et al. 2018"],"Functions Text":["To illustrate this, we show examples of calibration vectors for SDSS","eBOSS DR14"],"Functions Label":["Uses","Uses"],"Citation Start End":[[1011,1032]],"Functions Start End":[[837,905],[979,989]]}
{"Identifier":"2018ApJ...854...73I__Schenker_et_al._2013_Instance_1","Paragraph":"The best-fit parameters are shown in Table 3. We find that the uncertainties in M* and \n\n\n\n\n\n are considerably large due to a degeneracy between the two parameters when all parameters are variable. Plotted in Figure 1 is the faint-end slope \u03b1 as a function of redshift. Our results indicate that the best-fit values of \u03b1 are about \u22122 at \n\n\n\n\n\n to 10, which are steeper than those at lower redshift (e.g., \n\n\n\n\n\n at \n\n\n\n\n\n in Bouwens et al. 2015). We show the fitting results at \n\n\n\n\n\n, 8, 9, and 10 in Figures 2, 3, 4, and 5, respectively. The top and bottom panels present the observed number densities and the best-fit luminosity functions in the image plane and the source plane, respectively. We also plot the results of previous blank-field surveys (Ouchi et al. 2009; Bradley et al. 2012; McLure et al. 2013; Oesch et al. 2013; Schenker et al. 2013; Bowler et al. 2014; Bouwens et al. 2015; Finkelstein et al. 2015; Calvi et al. 2016) and recent HFF results in other studies (Atek et al. 2015a; Laporte et al. 2016; McLeod et al. 2016). The best-fit parameters are consistent with those in previous studies. In the top panel of Figure 2, there may be an excess in the observed surface number density at \n\n\n\n\n\n. The reason for this excess is not clear, although using a size\u2013luminosity relation that gives smaller sizes at faint magnitudes may reduce this excess. At \n\n\n\n\n\n, the observed number densities at the bright end are slightly larger than the number densities in the simulation. This is probably due to the existence of an overdense region of \n\n\n\n\n\n dropouts in the Abell 2744 cluster field. We discussed the properties of the overdensity in Ishigaki et al. (2016; see also Zheng et al. 2014 and Atek et al. 2015b). At \n\n\n\n\n\n, although we detect no galaxies, we can place a constraint on the luminosity function from the non-detection. Based on the best-fit parameters where only \n\n\n\n\n\n is variable, \u223c1.4 galaxies are expected to be detected in the HFF fields. The middle panels of Figures 2\u20135 show histograms of the number of the dropouts. It is seen that our samples push the magnitude limits of the luminosity functions significantly by up to \u223c3 magnitude. The correlations between M* and \u03b1 at z \u223c6\u20137 and 8 are presented in Figure 6. \n \n","Citation Text":["Schenker et al. 2013"],"Functions Text":["We also plot the results of previous blank-field surveys"],"Functions Label":["Uses"],"Citation Start End":[[834,854]],"Functions Start End":[[697,753]]}
{"Identifier":"2020AandA...642A..19M__Tempel_et_al._(2014a)_Instance_1","Paragraph":"While this large-scale structure of the Universe (LSS) is also composed of dark matter and gas, it was through the galaxy distribution that it has started to be detected. Galaxy clusters were the first cosmic web features to be identified and studied because they are easily detectable through various techniques. Only with the advent of wide-area spectroscopic redshift surveys have other structures such as filaments begun to be systematically identified. Surveys such as the Two-Degree Field Galaxy Redshift Survey (2dFGRS, Colless et al. 2001), the Sloan Digital Sky Survey (SDSS, York et al. 2000), the Galaxy And Mass Assembly survey (GAMA, Driver et al. 2009), the Vimos Public Extragalactic Redshift Survey (VIPERS, Scodeggio et al. 2018), or the COSMOS survey (Scoville et al. 2007) have allowed us to obtain statistical samples of filaments and other LSS features. For example, Chen et al. (2016) and Tempel et al. (2014a) have produced filament catalogues in the SDSS (but see also the works by Arag\u00f3n Calvo 2007; Sousbie et al. 2011; Rost et al. 2020; Shuntov et al. 2020; Kraljic et al. 2020, some of which also used the same algorithm as we used here). Other works such as Kraljic et al. (2018) and Alpaslan et al. (2014) detected filaments in GAMA, while Malavasi et al. (2017) detected filaments in VIPERS. Additionally, Gott et al. (2005), Iovino et al. (2016), and Kraljic et al. (2018) also identified walls in the SDSS, COSMOS, and GAMA surveys, respectively, while several projects are devoted to the analysis of voids (see e.g. Colberg et al. 2008, for a summary). Recently, not only spectroscopic surveys, but also the increased precision of photometric redshifts (e.g. in COSMOS and in the Canada-France-Hawaii Telescope Legacy Survey, CFHTLS, Laigle et al. 2016; Coupon et al. 2009) allowed for the detection of filaments in volumes of the Universe up to z\u2004\u223c\u20041 (Laigle et al. 2018; Sarron et al. 2019, but also Darvish et al. 2014).","Citation Text":["Tempel et al. (2014a)"],"Functions Text":["For example, Chen et al. (2016) and","have produced filament catalogues in the SDSS"],"Functions Label":["Background","Background"],"Citation Start End":[[911,932]],"Functions Start End":[[875,910],[933,978]]}
{"Identifier":"2020MNRAS.492..444Y__Proga_&_Begelman_2003_Instance_1","Paragraph":"In spite of these similarities, the angular profiles of these properties in two-dimensional simulations reveal a totally different accretion pattern. Within the stagnation radius, winds are mainly accreted in the polar region as a result of the low angular momentum there, whereas material with high angular momentum in the mid-plane outflow is subject to outflow, because the angular momentum of the stellar winds at a given radius decreases towards the pole. This \u2018funnel\u2019 accretion scenario is also found in other simulations with a spherical-like distribution of rotationally accreting material (Proga & Begelman 2003; Ressler et al. 2018). The density and temperature at the equatorial plane are higher than those in the polar region, and thus the highest mass accretion rate is found at the boundary of the polar and disc regions. Therefore the accretion pattern of the two-dimensional simulations is totally different from the inflow\u2013outflow structure in the one-dimensional calculations. Another difference between our results is that the parameters of our best-fitting model, fq and vw, sn, are generally smaller than those in one-dimensional calculations. The smaller fq can be explained in three ways. First, our BH mass is double that in the one-dimensional case, which results in a deeper gravitational potential that allows for a more massive gas reservoir. Thus fq should be smaller to achieve the observed gas density. Second, our outer boundary is set at 17 arcsec, which is larger than the radius of the one-dimensional outer boundary of 12 arcsec. The larger outer boundary causes the mass injection rate by stellar winds to increase by 50 per\u2009cent. Because stellar winds are the only mass source in our models, a smaller fq is required to cancel out this extra mass injection. Third, a smaller vw, sn could also lead to a smaller fq owing to less efficient heating by supernova explosions. The smaller vw, sn is constrained by the temperature at large radii, especially the temperature between 10 and 20 arcsec, which is not fitted in the one-dimensional calculations.","Citation Text":["Proga & Begelman 2003"],"Functions Text":["This \u2018funnel\u2019 accretion scenario is also found in other simulations with a spherical-like distribution of rotationally accreting material"],"Functions Label":["Similarities"],"Citation Start End":[[600,621]],"Functions Start End":[[461,598]]}
{"Identifier":"2017ApJ...850...18H__Murase_et_al._2015_Instance_1","Paragraph":"The heating due to the reprocessing of non-thermal photons produced in the nebula can be efficient even at late times. Here, we treat these processes in a simple way as follows. At early times, electromagnetic cascades proceed in the saturation regime, leading to a flat energy spectrum up to \u223c1 MeV (Metzger et al. 2014). At later times, the spectrum depends on the seed photon spectra, but it can roughly be estimated to be a flat spectrum from \u223c1 eV to \u223c0.1 TeV, while the supernova emission continues, which is expected based on more detailed calculations (e.g., Murase et al. 2015; K. Murase et al. 2017, in preparation). High-energy \u03b3-rays (\u22731 MeV) heat up the ejecta through the Compton scattering and the pair production process. X-ray and UV photons are absorbed and heat up the ejecta through the photoelectric (bound-free) absorption unless the ejecta are fully ionized. Here, we describe the heating rate as\n12\n\n\n\n\n\n\n\n\nQ\n\n\n\u02d9\n\n\n\n\nrad\n\n\n(\nt\n)\n\u2248\n\n\n(\n\n\n\nf\n\n\n\u03b3\n\n\n+\n\n\nf\n\n\n\n\n\nX\n\u2212\nUV\n\n\n,\nbf\n\n\n\n)\n\n\n\n\nL\n\n\nsd\n\n\n,\n\n\nwhere f\u03b3 and \n\n\n\n\n\n\nf\n\n\n\n\n\nX\n\u2212\nUV\n\n\n,\nbf\n\n\n\n\n are the heating efficiencies of \u03b3-rays and X-ray and UV photons to the spin-down luminosity, respectively. We calculate the frequency dependent heating efficiency of \u03b3-rays at each time:\n13\n\n\n\n\n\n\nf\n\n\n\u03b3\n\n\n(\nt\n)\n=\n\n\n\n\n\n\u222b\n\n\n\n\n\u03bd\n\n\nmin\n\n\n\n\n\n\n\u03bd\n\n\nmax\n\n\n\n\n\n\n\nd\n\u03bd\n\n\n\u03bd\n\n\n\n\nmin\n(\n\n\nK\n\n\n\u03b3\n,\n\u03bd\n\n\n\n\n\u03c4\n\n\n\u03b3\n,\n\u03bd\n\n\n,\n1\n)\n\n\n\n\n\u222b\n\n\n1\n\neV\n\n\n1\n\nTeV\n\n\n\n\n\nd\n\u03bd\n\n\n\u03bd\n\n\n\n\n\n\n,\n\n\nwhere the frequency range of \u03b3-rays is \n\n\n\n\n(\nh\n\n\n\u03bd\n\n\nmin\n\n\n,\nh\n\n\n\u03bd\n\n\nmax\n\n\n)\n\n=\n(\n10\n\nkeV\n,\n1\n\nTeV\n)\n\n\n, and h is the Planck constant. Here, \u03c4\u03b3, \u03bd is the optical depth of the ejecta to \u03b3-rays and K\u03b3, \u03bd is the photon inelasticity at a given frequency, where the Klein\u2013Nishina cross section and the cross section for the Bethe\u2013Heitler pair production in the field of a carbon nucleus are taken into account (Chodorowski et al. 1992; Murase et al. 2015). Note that the coefficient of the \u03b3-ray optical depth depends on the density profile of the ejecta. Here, we simply assume a density profile to be constant with the radius. Adopting a realistic density profile may result in different ejecta mass and velocity estimates by a factor of a few.","Citation Text":["Murase et al. 2015"],"Functions Text":["At later times, the spectrum depends on the seed photon spectra, but it can roughly be estimated to be a flat spectrum from \u223c1 eV to \u223c0.1 TeV, while the supernova emission continues, which is expected based on more detailed calculations (e.g.,"],"Functions Label":["Uses"],"Citation Start End":[[567,585]],"Functions Start End":[[323,566]]}
{"Identifier":"2020ApJ...901...41S__Harrington_1973_Instance_1","Paragraph":"Observations have shown that the shape of the Ly\u03b1 line is diverse. It includes broad damped absorption profiles, P-Cygni profiles, double-peak profiles, pure symmetric emission profiles, and combinations thereof (Kunth et al. 1998; Mas-Hesse et al. 2003; Shapley et al. 2003; M\u00f8ller et al. 2004; Noll et al. 2004; Tapken et al. 2004; Venemans et al. 2005; Wilman et al. 2005). This variety can be understood through a detailed radiative transfer calculation, which is analytically solvable only for simple cases (e.g., a static, plane-parallel slab by Harrington 1973 and Neufeld 1990, and a static uniform sphere by Dijkstra et al. 2006). Later, numerical algorithms based on Monte Carlo techniques were developed to solve radiative transfer for more general cases. Now theoretical studies mostly rely on them (e.g., Spaans 1996; Loeb & Rybicki 1999; Ahn et al. 2000, 2002; Zheng & Miralda-Escud\u00e9 2002; Richling 2003; Cantalupo et al. 2005; Dijkstra et al. 2006; Hansen & Oh 2006; Tasitsiomi 2006; Verhamme et al. 2006, 2015; Laursen et al. 2013; Behrens et al. 2014; Duval et al. 2014; Gronke et al. 2015; Smith et al. 2019; Lao & Smith 2020; Michel-Dansac et al. 2020). Meanwhile, a galaxy model needs to be constructed to perform such a radiative transfer calculation. One can adopt a realistic galaxy model from hydrodynamical simulations. Galaxies from such simulations can be useful for performing a statistical study of Ly\u03b1 properties, but they cannot be directly used to model individual galaxies in observations. Therefore it would be better to adopt a simple but manageable toy model for the purpose of reproducing observations. One example for such models is a shell model, in which a central Ly\u03b1 source is surrounded by a constantly expanding, homogeneous, spherical shell of H i medium with dust. Although this shell model has surprisingly well reproduced many observed Ly\u03b1 line profiles (e.g., Ahn 2004; Schaerer & Verhamme 2008; Verhamme et al. 2008; Schaerer et al. 2011; Gronke et al. 2015; Yang et al. 2016; Gronke 2017; Karman et al. 2017), because of its extreme simplicity and contrivance, there is still room for improvement (e.g., see Section 7.2 in Yang et al. 2016; Orlitov\u00e1 et al. 2018).","Citation Text":["Harrington 1973"],"Functions Text":["This variety can be understood through a detailed radiative transfer calculation, which is analytically solvable only for simple cases (e.g., a static, plane-parallel slab by"],"Functions Label":["Background"],"Citation Start End":[[552,567]],"Functions Start End":[[377,551]]}
{"Identifier":"2020MNRAS.496.1051A__Rudolph_et_al._2006_Instance_2","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"],"Functions Text":["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."],"Functions Label":["Differences"],"Citation Start End":[[902,921]],"Functions Start End":[[530,822]]}
{"Identifier":"2021ApJ...917...24Z__Coughlin_et_al._2020b_Instance_2","Paragraph":"Our simulation results show that the median detectable distances of targeted GW events from BH\u2013NS mergers for a single 2nd generation GW detector and a network of such detectors are \u223c300 Mpc and \u223c700 Mpc, respectively (see Table 4). For comparison, Figure 12 shows that the detection rate and detectable distance for HLV (O3) are approximately the same as those for the case when only bKAGRA is running. This is basically consistent with the detection rate and the distance distribution of BH\u2013NS merger candidates detected during LVC O3 (e.g., Anand et al. 2020; Antier et al. 2020b; Coughlin et al. 2020b; Gompertz et al. 2020; Kasliwal et al. 2020). In Section 3, we have shown that the kilonova absolute magnitude at 0.5 days after a BH\u2013NS merger is mainly distributed in the range \u223c\u201310 to \u223c\u201315.5. In view of the fact that the limiting magnitude of the follow-up wide-field survey projects is almost \u227221 mag (e.g., Antier et al. 2020b; Gompertz et al. 2020; Coughlin et al. 2020b; Kasliwal et al. 2020; Wyatt et al. 2020), the maximum detectable distance for BH\u2013NS kilonovae would be \u2272200 Mpc, which can hardly cover the horizon of GW-triggered BH\u2013NS merger events that O3 found (as shown in Figure 10). However, although BH\u2013NS merger kilonovae can hardly be detected for the present search depths, Figures 9 and 11 reveal that there are great opportunities to discover on-axis afterglows associated with sGRBs or orphan afterglows if the BH components have a high-spin distribution. In order to cover the distance range for searching for BH\u2013NS kilonovae for the network of 2nd generation GW detectors as completely as possible, a search limiting magnitude mlimit \u2273 23\u201324 is required as shown in Figure 10. Present survey projects could reach this search limiting magnitude by increasing exposure times and the number of simultaneous exposures. However, the GW candidates during O3 had very large localization areas with an average of thousands of square degrees (Antier et al. 2020b). Increasing exposure times makes it hard for the present survey projects to cover such large localization areas. Therefore, during the HLVK era, we recommend that survey projects may search for jet afterglows after GW triggers with a relatively shallow search limiting magnitude. If BH\u2013NS mergers have a high location precision, a limiting magnitude of mlimit \u2273 23\u201324 can be reached, which gives a higher probability of discovering associated kilonovae.","Citation Text":["Coughlin et al. 2020b"],"Functions Text":["In view of the fact that the limiting magnitude of the follow-up wide-field survey projects is almost \u227221 mag (e.g.,","the maximum detectable distance for BH\u2013NS kilonovae would be \u2272200 Mpc, which can hardly cover the horizon of GW-triggered BH\u2013NS merger events that O3 found (as shown in Figure 10)."],"Functions Label":["Uses","Uses"],"Citation Start End":[[961,982]],"Functions Start End":[[801,917],[1026,1206]]}
{"Identifier":"2022MNRAS.515.2188H__Rorai_et_al._2018_Instance_2","Paragraph":"In this work, we follow the method for measuring the IGM thermal state based on Voigt profile decomposition of the Ly\u2009\u03b1 forest (Schaye et al. 1999; Ricotti et al. 2000; McDonald et al. 2001). In this approach, a transmission spectrum is treated as a superposition of multiple discrete Voigt profiles, with each line described by three parameters: redshift zabs, Doppler broadening b, and neutral hydrogen column density $N_{\\rm H\\, {\\small I}}$. By studying the statistical properties of these parameters, i.e. the $b-N_{{{{\\rm H\\, {\\small I}}}}{}}$ distribution, one can recover the thermal information encoded in the absorption profiles. The majority of past applications of this method constrained the IGM thermal state by fitting the low-b-$N_{\\rm H\\, {\\small I}}$ cutoff of the $b-N_{{{{\\rm H\\, {\\small I}}}}{}}$ distribution (Schaye et al. 1999, 2000; Ricotti et al. 2000; McDonald et al. 2001; Rudie, Steidel & Pettini 2012; Boera et al. 2014; Bolton et al. 2014; Garzilli, Theuns & Schaye 2015, 2020; Hiss et al. 2018; Rorai et al. 2018). The motivation for this approach is that the Ly\u2009\u03b1 lines are broadened by both thermal motion and non-thermal broadening resulting from combinations of Hubble flow, peculiar velocities, and turbulence. By isolating the narrow lines in the Ly\u2009\u03b1 forest that constitutes the lower cutoff in $b-N_{{{{\\rm H\\, {\\small I}}}}{}}$ distributions, of which the line-of-sight component of non-thermal broadening is expected to be zero, the broadening should be purely thermal, thus allowing one to constrain the IGM thermal state. However, this method has three crucial drawbacks. First, the IGM thermal state actually impacts all the lines besides just the narrowest lines. Therefore, by restricting attention to data in the distribution outskirts, this approach throws away information and reduces the sensitivity to the IGM thermal state significantly (Rorai et al. 2018; Hiss et al. 2019). Secondly, in practice, determining the location of the cutoff is vulnerable to systematic effects, such as contamination from the narrow metal lines (Hiss et al. 2018; Rorai et al. 2018). Lastly, the results from this approach critically depend on the choice of low-b cutoff fitting techniques, where different techniques might result in inconsistent T0 and \u03b3 measurements (Hiss et al. 2018; Rorai et al. 2018).","Citation Text":["Rorai et al. 2018"],"Functions Text":["Therefore, by restricting attention to data in the distribution outskirts, this approach throws away information and reduces the sensitivity to the IGM thermal state significantly"],"Functions Label":["Background"],"Citation Start End":[[1891,1908]],"Functions Start End":[[1710,1889]]}
{"Identifier":"2018ApJ...867..101B___2018b_Instance_1","Paragraph":"To fully describe a population of streams, we need a realistic model of the Galaxy. The latest generation of hydrodynamical simulations have produced models that well reproduce a multitude of features observed in galaxies, both as individual objects (e.g., Wetzel et al. 2016) and as a population (e.g., Nelson et al. 2018; Pillepich et al. 2018). Part of their success is in the achieved high resolution, so at the present a Milky-Way-like galaxy is modeled by up to 140 million particles (Wetzel et al. 2016). Ultimately, we would like to have a description of a galaxy that is representative of these models, but at a fraction of the numerical cost. In the classic study, Hernquist & Ostriker (1992) developed a set of basis functions for density and gravitational potential that can reproduce complex morphology of galaxies with a small number of terms (e.g., Lowing et al. 2011; Lilley et al. 2018a, 2018b). These expansions reproduce the force field of an N-body simulation with a precision of a few percent, so representing the gravitational potential in our model with an expansion of basis functions is tempting. However, a truly realistic solution needs to accurately capture not only the current structure of the Galaxy but also its evolution in time. Even though the Milky Way has had a relatively quiet recent merger history, it is currently undergoing a major merger with the Large Magellanic Cloud (e.g., Besla et al. 2007; Pe\u00f1arrubia et al. 2016), which has been a source of gravitational perturbation for at least a billion years\u2014sufficiently long to affect stellar streams. Specifically, Law & Majewski (2010) showed that the only static, ellipsoidal halo that reproduces both the positions and radial velocities along the Sagittarius stream is triaxial. This model correctly predicted proper motions along the stream (Sohn et al. 2015), so it appears to be describing the effective potential well, even though it is cosmologically improbable (Debattista et al. 2013). On the other hand, modeling Sagittarius in a combined system of the Milky Way and the LMC relaxes the requirement for the dark matter halo to be triaxial (Vera-Ciro & Helmi 2013), signaling that having a model of the potential that is correct on average is no guarantee of recovering the true halo shape. To ensure that the complexities added to the model are realistic, the basis function expansion should thus be time dependent, simultaneously describing the interaction between the Milky Way and the LMC, while maintaining a thin and old disk. We relegate the development of such a model and its implementation for mapping the dark matter in the Galaxy to a future study.","Citation Text":["Lilley et al.","2018b"],"Functions Text":["In the classic study, Hernquist & Ostriker (1992) developed a set of basis functions for density and gravitational potential that can reproduce complex morphology of galaxies with a small number of terms (e.g.,","These expansions reproduce the force field of an N-body simulation with a precision of a few percent, so representing the gravitational potential in our model with an expansion of basis functions is tempting. However, a truly realistic solution needs to accurately capture not only the current structure of the Galaxy but also its evolution in time."],"Functions Label":["Background","Motivation"],"Citation Start End":[[884,897],[905,910]],"Functions Start End":[[653,863],[913,1262]]}
{"Identifier":"2021AandA...655A..99D__Carigi_et_al._2005_Instance_2","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)"],"Functions Text":["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","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)."],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[1909,1929]],"Functions Start End":[[1607,1908],[1931,2140]]}
{"Identifier":"2020ApJ...892L..10Y__Macchi_2013_Instance_2","Paragraph":"In this section, we consider the plasma properties under the propagation of strong waves. In strong waves, the motion of electrons in the plasma becomes relativistic. However, different from free electrons that have a relativistic drift velocity in the direction of the incident electromagnetic wave (see Section 3.1), in plasma the space-charge potential is important in preventing the drift of electrons (Waltz & Manley 1978). For nonrelativistic electrons in plasma, if the wave duration \u03c4 is much larger than c\/\u03c9p, where \n\n\n\n\n\n is the plasma frequency, the drift velocity would be close to zero (Waltz & Manley 1978; Sprangle et al. 1990b). In this case, electrons in plasma under a strong wave would have a typical Lorentz factor (\n\n\n\n\n\n) similar to that (\u03b3) in the laboratory frame, so that \n\n\n\n\n\n is satisfied. Due to the relativistic and magnetic force effects, the propagation and dispersion properties of an electromagnetic wave depend on its amplitude. For a circular polarized wave, the dispersion relation in the laboratory frame is given by (e.g., Gibbon 2005; Macchi 2013; Macchi et al. 2013; see the Appendix)\n11\n\n\n\n\n\nThe dispersion relation of strong electromagnetic waves is altered due to the effective electron mass increased by the relativistic effect (e.g., Sarachik & Schappert 1970; Gibbon 2005; Macchi 2013). One can define the effective plasma frequency as\n12\n\n\n\n\n\nso that the wave can propagate in the region where \n\n\n\n\n\n. With respect to the nonrelativistic linear case, this is known as relativistically self-induced transparency. We note that since the dispersion depends on the electromagnetic field amplitude in the nonlinear case, the dispersion relation must be taken with care. The propagation of a pulse will be affected by the complicated effects of nonlinear propagation and dispersion, and finally the spatial and temporal shape of the pulse itself would also be modified. In particular, for linear polarization, the relativistic factor \u03b3 is not a constant (see Section 3.1). The propagation of the linearly polarized wave with a relativistic amplitude would lead to generation of the higher-order harmonics. Sprangle et al. (1990b) proved that the propagation of the first harmonic component, i.e., of the \u201cmain\u201d wave, is still reasonably described by Equation (11) with \n\n\n\n\n\n. Thus, we will directly adopt Equation (11) in the following discussion.","Citation Text":["Macchi 2013"],"Functions Text":["The dispersion relation of strong electromagnetic waves is altered due to the effective electron mass increased by the relativistic effect (e.g.,"],"Functions Label":["Uses"],"Citation Start End":[[1320,1331]],"Functions Start End":[[1134,1279]]}
{"Identifier":"2021ApJ...921...20H__Hayasaki_et_al._2013_Instance_1","Paragraph":"It is still debated if and how all the stellar debris efficiently circularizes via the stream\u2013stream collision. Some hydrodynamical simulations show that the TDE disk retains a significantly elliptical shape because the orbital energy is not dissipated efficiently enough to reduce the eccentricity of the entire disk to zero in a reasonable time (Guillochon et al. 2014; Shiokawa et al. 2015; S\u0105dowski et al. 2016). Lu & Bonnerot (2020) show that a significant fraction of the debris can become unbound causing an outflow from the self-interaction region. Nevertheless, the debris that remains bound eventually contributes to the accretion flow around the SMBH. This part of the debris stream will finally be circularized by energy dissipation, leading to the formation of a small, initially ring-like, accretion disk around the black hole (Hayasaki et al. 2013; Bonnerot et al. 2016; Hayasaki et al. 2016). Note that, in an inefficient debris circularization case, the subsequent fallback material interacts with the outer elliptical debris so that their combined effect on the subsequent evolution of the initial ring is negligible. Angular momentum conservation allows us to estimate the circularization radius of the stellar debris, which is given by\n2\n\n\n\n\n\n\nr\n\n\nc\n\n\n=\n(\n1\n+\n\n\ne\n\n\n*\n\n\n)\n\n\nr\n\n\np\n\n\n,\n=\n\n\n\n\n1\n+\n\n\ne\n\n\n*\n\n\n\n\n\u03b2\n\n\n\n\n\nr\n\n\nt\n\n\n,\n\n\nwhere e* is the orbital eccentricity of the stellar orbit, rp = rt\/\u03b2 is the pericenter distance radius. If debris circularization takes place only through dissipation at the self-interaction shock, the circularization timescale for the non-magnetized, most tightly bound debris can be estimated based on the ballistic approximation (Bonnerot et al. 2017) as\n3\n\n\n\n\n\n\n\n\n\nt\n\n\ncirc\n\n\n\n\n\u2248\n\n\n8.3\n\n\n\n\u03b7\n\n\n\u2212\n1\n\n\n\n\n\n\u03b2\n\n\n\u2212\n3\n\n\n\n\nM\n\n\nbh\n,\n6\n\n\n\u2212\n5\n\n\/\n\n3\n\n\n\n\n\nt\n\n\nmtb\n\n\n\n\n\n\n\n\u223c\n\n\n0.93\n\n\n\n\n\n\n\n\u03b7\n\n\n1.0\n\n\n\n\n\n\n\n\u2212\n1\n\n\n\n\n\u03b2\n\n\n\u2212\n3\n\n\n\n\nM\n\n\nbh\n,\n6\n\n\n\u2212\n7\n\n\/\n\n6\n\n\n\n\n\nm\n\n\n*\n,\n1\n\n\n\u2212\n1\n\n\n\n\n\nr\n\n\n*\n,\n1\n\n\n3\n\n\/\n\n2\n\n\n\nyr\n,\n\n\n\n\n\nwhere the orbital period of the stellar debris on the most tightly bound orbit:\n4\n\n\n\n\n\n\n\n\n\nt\n\n\nmtb\n\n\n\n\n=\n\n\n\n\n\n\u03c0\n\n\n\n\n2\n\n\n\n\n\n\n\n\n1\n\n\n\n\n\u03a9\n\n\n*\n\n\n\n\n\n\n\n\n\n\n\n\n\n\nM\n\n\nbh\n\n\n\n\n\n\nm\n\n\n*\n\n\n\n\n\n\n\n\n\n1\n\n\/\n\n2\n\n\n\n\n\n\n\n\u2248\n\n\n0.11\n\n\n\nM\n\n\nbh\n,\n6\n\n\n1\n\n\/\n\n2\n\n\n\n\n\nm\n\n\n*\n,\n1\n\n\n\u2212\n1\n\n\n\n\n\nr\n\n\n*\n,\n1\n\n\n3\n\n\/\n\n2\n\n\n\nyr\n,\n\n\n\n\n\n\n\n\n\n\n\n\n\n\u03a9\n\n\n*\n\n\n=\n\n\n\n\nGm\n\n\n*\n\n\n\n\/\n\n\n\nr\n\n\n*\n\n\n3\n\n\n\n\n\n\n is the dynamical angular frequency of the star, and we introduce \u03b7(\u22641) as the circularization efficiency which represents how efficiently the kinetic energy at the stream\u2013stream collision is dissipated and the most efficient (\u03b7 = 1) case corresponds to that of Bonnerot et al. (2017). Note that tcirc is not the circularization timescale of all the stellar debris. Our interest here is in the circularization timescale and radius of the most tightly bound debris because the accretion of this debris contributes most to the delayed X-ray peak luminosity in terms of the emitted energy.","Citation Text":["Hayasaki et al. 2013"],"Functions Text":["Nevertheless, the debris that remains bound eventually contributes to the accretion flow around the SMBH. This part of the debris stream will finally be circularized by energy dissipation, leading to the formation of a small, initially ring-like, accretion disk around the black hole"],"Functions Label":["Background"],"Citation Start End":[[842,862]],"Functions Start End":[[557,840]]}
{"Identifier":"2016ApJ...816...41Y__Litvinenko_&_Wheatland_2005_Instance_1","Paragraph":"Some previous observations have shown that surface motions (e.g., shearing motions and converging motions) acting on preexisting coronal fields to form filaments and filament channels always act over a short period of a few days (Gaizauskas et al. 1997; Schmieder et al. 2004; Wang & Muglach 2007; Yan et al. 2015; Yang et al. 2015) to a period of months (Gaizauskas et al. 2001). However, in this event, the filament is formed within about 20 minutes. On this timescale, only the reformation of the filament in the same filament channel after the partial filament eruption has been observed (Tripathi et al. 2009a; Joshi et al. 2014). The rapid formation of the filament is a rare observation that has not been previously reported. Even though our event follows the \u201chead-to-tail\u201d scenario, there is a little difference. It is worth noting that the \u201chead-to-tail\u201d scenario seems to support the idea that the filament is supported by sheared arcades (Litvinenko & Wheatland 2005; Welsch et al. 2005). However, our event may present a picture in which the filament is supported by a twisted flux rope. Therefore, we could not fully exclude any other explanations, although we are inclined toward the \u201chead-to-tail\u201d scenario of Martens & Zwaan. Twisted structures in filaments have been previously observed during their eruptions (Wang et al. 1996; Bi et al. 2012, 2015; Li & Zhang 2013; Yan et al. 2014; Yang et al. 2014; Filippov et al. 2015; Raouafi 2009). In particular, Bi et al. (2012, 2015) found that the erupted filaments appear to be composed of two intertwined twisted flux ropes. However, in the present case, it is found that the flux rope appeared in the course of the formation of the filament, and as reported by Li & Zhang (2013; Yang et al. 2014), it can be observed in all seven AIA EUV lines formed from 0.05 MK to 11 MK. The AIA 304 \u212b observations (see panel (e) in Figure 2) also indicate that the filament may be composed of two sets of intertwined dark threadlike structures. These observations may be direct evidence that the magnetic structure of the formed filament is a flux rope. As mentioned before, the sigmoids and the elongated EUV channel-like structures are evidences for the existence of flux ropes in the higher solar atmosphere. More recently, observational evidence of the detailed structure and evolution of flux ropes in the lower solar atmosphere has been presented by Wang et al. (2015) and Kumar et al. (2015). The flux rope in our case can be observed in all of  the AIA EUV passbands, but it is still in the lower solar atmosphere. Our observation is quite similar to the formation of a chromospheric flux rope as a result of magnetic reconnection between two sheared cool H\u03b1 loops (Wang et al. 2015) in the lower atmosphere, but sheds more light on the formation of the filament. This event may also be a distinct example revealing the evolution of the flux rope in the lower solar atmosphere.","Citation Text":["Litvinenko & Wheatland 2005"],"Functions Text":["It is worth noting that the \u201chead-to-tail\u201d scenario seems to support the idea that the filament is supported by sheared arcades"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[951,978]],"Functions Start End":[[822,949]]}
{"Identifier":"2020MNRAS.494.4382S___2010_Instance_2","Paragraph":"It has been thought that QPOs originate from the innermost part of an accretion disc, which is associated with strong gravity, so that we might detect general relativistic effects. Miller et al. (1998) proposed beat-frequency models and estimated the parameters of NSs using this model. Stella & Vietri (1999) developed the relativistic precession model. In the last 20 years, disc-oscillation and resonance models and wave models have been proposed (e.g. Osherovich & Titarchuk 1999; Abramowicz & Klu\u017aniak 2001; Abramowicz et al. 2003; Zhang 2004; Li & Zhang 2005; Erkut, Psaltis & Alpar 2008; Shi & Li 2009, 2010; Shi 2011; Shi, Zhang & Li 2014, 2018; de Avellar et al. 2018). Shi & Li (2009, 2010) obtained the twin modes of MHD waves in LMXBs (including NS LMXBs and black hole LMXBs), which are considered as the sources of high-frequency QPOs. Shi, Zhang & Li (2014, 2018) also considered the waves produced by the two MHD oscillation modes at the magnetosphere radius as the origin of kHz QPOs. A relationship between the frequencies of the twin-peak kHz QPOs and the accretion rate, in which parallel tracks can be explained, was obtained (Shi, Zhang & Li 2018). Recently, many simulations on the oscillations of accreting tori in the accretion process of NS LMXBs (e.g. Kulkarni & Romanova 2013; Parthasarathy, Klu\u017aniak, \u010cemelji\u0107 2017) have been performed, and almost every model can reproduce some of the observed characteristics of QPOs. However, most models cannot fit the observed data perfectly, the observed data. Belloni, M\u00e9ndez & Homan (2005) suggested that the twin kHz QPOs showed no intrinsically preferred frequency ratio, and this weakened support for the resonance models. Morsink & Stella (1999) were able to fit the overall NS data with different masses and spins of NSs using the relativistic precession model; however, Belloni, M\u00e9ndez & Homan (2007) found that there were deviations between the expected and the observed trends. Recently, T\u00f6r\u00f6k et al. (2016b, 2018) identified the observed QPO frequencies with the frequencies of the epicyclic modes of torus oscillations, and suggested that the relationship between the strong modulation of the X-ray flux and high values of QPO frequencies is connected to the orbital motion in the innermost part of an accretion disc. In addition, there are studies that compared a large set of models with the data of many sources in a complex manner (Lin et al. 2011; T\u00f6r\u00f6k et al. 2012, 2016a).","Citation Text":["Shi & Li","2010"],"Functions Text":["obtained the twin modes of MHD waves in LMXBs (including NS LMXBs and black hole LMXBs), which are considered as the sources of high-frequency QPOs."],"Functions Label":["Background"],"Citation Start End":[[679,687],[695,699]],"Functions Start End":[[701,849]]}
{"Identifier":"2021ApJ...921..179L__Kuznetsov_&_Kolotkov_2021_Instance_1","Paragraph":"Solar flares are powerful eruption events on the Sun associated with a rapid and violent release of magnetic free energy through a reconnection process. A typical flare can radiate at almost all wavelengths constituting the solar spectrum, ranging from radio through optical and ultraviolet (UV) to soft\/hard X-ray (SXR\/HXR) and even \u03b3-rays (e.g., Benz 2017; Tan et al. 2020). Only a small part of the flare radiation is emitted at the shortest wavelengths in the X-ray and extreme-UV (EUV) ranges (Emslie et al. 2012). The quantitative estimation of the radiated flare energy partition suggested about 70% in white light (WL) for solar flares (e.g., Kretzschmar 2011) and 55%\u201380% in WL for stellar flares (e.g., Kuznetsov & Kolotkov 2021). In other words, most of the flare energy is radiated in the longer wavelengths (Kleint et al. 2016). Between those extremes, the solar UV spectrum from 1000 to 3000 \u212b, which can be further split into the far-ultraviolet (FUV), the mid-ultraviolet (MUV), and the near-ultraviolet (NUV), is thought to provide an important contribution to the flare radiation (Woods et al. 2006; Milligan et al. 2014; Dominique et al. 2018). For instance, the Ly\u03b1 spectral line produced by the chromospheric neutral hydrogen, which is centered at 1216 \u212b (in the FUV spectrum), is among the spectral lines in which flares radiate the most (Allred et al. 2005; Curdt et al. 2001; Lu et al. 2021a). The hydrogen Balmer continuum emitted during flares, which is thought to be generated during the recombination of flare-produced free electrons in the chromosphere, is often detected in the MUV and NUV ranges, as well as close to the Balmer recombination edge at 3646 \u212b (Heinzel & Kleint 2014; Kotr\u010d et al. 2016; Dominique et al. 2018). Both the Ly\u03b1 and hydrogen Balmer continuum emissions during solar flares are expected to be nonthermal profiles, i.e., similar to the HXR radiation that is produced by the beam of electrons that are accelerated by the magnetic reconnection during the solar flare (e.g., Avrett et al. 1986; Rubio da Costa et al. 2009; Heinzel & Kleint 2014).","Citation Text":["Kuznetsov & Kolotkov 2021"],"Functions Text":["The quantitative estimation of the radiated flare energy partition suggested about","and 55%\u201380% in WL for stellar flares (e.g.,"],"Functions Label":["Background","Background"],"Citation Start End":[[713,738]],"Functions Start End":[[520,602],[669,712]]}
{"Identifier":"2021AandA...655A..25Z__Garc\u00eda-Burillo_et_al._2014_Instance_2","Paragraph":"Outflows are ubiquitous in both luminous AGN and in local Seyfert galaxies, and occur on a wide range of physical scales, from highly ionised semi-relativistic winds and jets in the nuclear region at subparsec scales to galactic scale outflows seen in mildly ionised, molecular, and neutral gas (Morganti et al. 2016; Fiore et al. 2017; Fluetsch et al. 2019; Lutz et al. 2020; Veilleux et al. 2020, and references therein). In some cases molecular and ionised winds have similar velocities and are nearly co-spatial, suggesting a cooling sequence scenario where molecular gas forms from the cooling of the gas in the ionised wind (Richings & Faucher-Giguere 2017; Menci et al. 2019). Other AGN show ionised winds that are faster than the molecular winds, suggesting a different origin of the two phases (Veilleux et al. 2020, and references therein). The molecular phase is a crucial element of the feeding and feedback cycle of AGN because it constitutes the bulk of the total gas mass and it is the site of star formation activity. On galactic scales massive molecular winds are common in local Seyfert galaxies (e.g. Feruglio et al. 2010; Cicone et al. 2014; Dasyra et al. 2014; Morganti et al. 2015; Garc\u00eda-Burillo et al. 2014, 2017, 2019); these winds likely suppress star formation (i.e. negative feedback) as they reduce the molecular gas reservoir by heating or expelling gas from the host-galaxy ISM. In late-type AGN-host galaxies, the gas kinematics appears complex at all scales, showing several components such as bars, rings, and (warped) discs, with high velocity dispersion regions (e.g. Shimizu et al. 2019; Feruglio et al. 2020; Fern\u00e1ndez-Ontiveros et al. 2020; Alonso-Herrero et al. 2020; Aalto et al. 2020; Audibert et al. 2020). Accurate dynamical modelling of the molecular gas kinematics reveals kinematically decoupled nuclear structures, high velocity dispersion at nuclei, trailing spirals, and evidence of inflows and AGN-driven outflows. (e.g. Combes et al. 2019; Combes 2019, 2021). The outflow driving mechanism (wind shock, radiation pressure, or jet), their multiphase nature, and their relative weights and impact on the galaxy ISM are still open problems (Faucher-Gigu\u00e8re & Quataert 2012; Zubovas & King 2012; Richings & Faucher-Giguere 2017; Menci et al. 2019; Ishibashi et al. 2021). To date, far different outflow phases have been observed only for a handful of sources. Atomic, cold, and warm molecular outflows have been observed in radio galaxies (e.g. Morganti et al. 2007; Dasyra & Combes 2012; Dasyra et al. 2014; Tadhunter et al. 2014; Oosterloo et al. 2017). The nuclear semi-relativistic phase and the galaxy-scale molecular phase have been observed simultaneously in less than a dozen sources, with varied results: in some cases data suggest energy driven flows (Feruglio et al. 2015; Tombesi et al. 2015; Longinotti et al. 2018; Smith et al. 2019), in other cases data suggest momentum driven flows (e.g. Garc\u00eda-Burillo et al. 2014; Feruglio et al. 2017; Fluetsch et al. 2019; Bischetti et al. 2019; Marasco et al. 2020). Fiore et al. (2017), using a compilation of local and high redshift winds, showed that there is a broad distribution of the momentum boost, suggesting that both energy- and momentum-conserving expansion may occur. Enlarging the sample of local AGN-host galaxies with outflows detected in different gas phases is important to understand the nature and driving mechanisms of galaxy-scale outflows.","Citation Text":["Garc\u00eda-Burillo et al. 2014"],"Functions Text":["The nuclear semi-relativistic phase and the galaxy-scale molecular phase have been observed simultaneously in less than a dozen sources, with varied results:","in other cases data suggest momentum driven flows (e.g."],"Functions Label":["Background","Background"],"Citation Start End":[[2953,2979]],"Functions Start End":[[2604,2761],[2897,2952]]}
{"Identifier":"2020MNRAS.499.4666F__Popping_et_al._2017_Instance_1","Paragraph":"An example of these implications is the so-called \u2018dust budget crisis\u2019 introduced in Section 4.4: the dust masses currently estimated at z > 5 are not compatible with standard dust production channels and require an overhaul in our models of the initial mass function for star formation, of supernova production rates, or of dust growth in the ISM. Overall, the dust production rate would need to increase by one to two orders of magnitudes, as shown by Rowlands et al. (2014). The growth of dust grains through accretion in the ISM has been proposed as a solution (e.g. Mancini et al. 2015; Micha\u0142owski 2015; Popping et al. 2017), but there are doubts on the efficiency of accretion at high z, where high dust temperatures due to the CMB (see Section 3.3) keep the desorption time-scale for accreted materials short (Ferrara et al. 2016). The dust budget crisis is not only a problem at high redshift; it is observed, e.g. in the Magellanic Clouds (SMC, LMC). As explained in Srinivasan et al. (2016) using the dust mass fits by Gordon et al. (2014), the dust replenishment time-scale in the SMC from stellar sources alone is expected to be larger than the dust destruction time-scale and, in the worst-case scenario, longer than the lifetime of the Universe. Similarly, the ratio between the best LMC dust mass estimate by Gordon et al. (2014) and the dust injection estimates by Riebel et al. (2012) results in an LMC replenishment time-scale of 34 \u00b1 8 Gyr, exceeding the age of the Universe. Both the high redshift and the local Universe, therefore, show a dust budget crisis that could be alleviated \u2013 and, in the best case scenario, fully resolved \u2013 if the actual dust masses turned out to be lower than currently estimated, as our results suggest. More specifically, Rowlands et al. (2014) mention that dust opacity needs to be increased by just a factor of 7 to solve the high-redshift crisis (provided dust destruction by SNe is not efficient); in the LMC, the aforementioned replenishment time-scale would decrease to less than 2 Gyr if the dust mass were decreased by a factor of 20.","Citation Text":["Popping et al. 2017"],"Functions Text":["The growth of dust grains through accretion in the ISM has been proposed as a solution (e.g.","but there are doubts on the efficiency of accretion at high z, where high dust temperatures due to the CMB (see Section 3.3) keep the desorption time-scale for accreted materials short"],"Functions Label":["Background","Compare\/Contrast"],"Citation Start End":[[610,629]],"Functions Start End":[[478,570],[632,816]]}
{"Identifier":"2018ApJ...853..148C__Shibuya_et_al._2014_Instance_1","Paragraph":"LAE galaxies are defined by a high equivalent width (EW > 20 \u212b) Ly\u03b1 line and are believed to be composed of extremely large regions of active star formation. Many efforts have been made to detect and characterize LAE galaxies (e.g., Conselice et al. 2003; Conselice 2004; Ravindranath et al. 2006; Shimasaku et al. 2006; Bournaud et al. 2007; Ouchi et al. 2008, 2017; Elmegreen et al. 2009a, 2009b; Tacconi et al. 2010; Gronwall et al. 2011; Kashikawa et al. 2011; Mandelker et al. 2014; Moody et al. 2014; Guo et al. 2015). In general, these galaxies appear as clusters of bright clumps, sometimes with a background of continuum emission. Evidence suggests that these clumps are larger and brighter than most star-forming regions in nearby low-redshift galaxies (Elmegreen et al. 2009a). Efforts have been made in quantifying mass, star formation rates, gas composition, and kinematics, as well as other LAE properties (e.g., Nilsson et al. 2009; Ono et al. 2010a, 2010b; Swinbank et al. 2010; Tacconi et al. 2010; Shibuya et al. 2014; Livermore et al. 2015; Nakajima et al. 2016; Hashimoto et al. 2017). These have revealed a wealth of information about the early universe, but they are ultimately limited by LAE surface brightnesses. Most studies rely upon stacks of galaxies and can draw only limited inferences about individual LAEs. Other studies show that LAE dust content, particularly clumpy dust, in the interstellar medium (ISM) can have an impact on most LAE observables (Kobayashi et al. 2007, 2010; Verhamme et al. 2008; Duval et al. 2014). Finkelstein et al. (2009) showed that clumpy dust models can provide a good fit to a set of z \u223c 4.5 LAEs, although they invoked a multiphase ISM that may be unlikely to form in nature (Laursen et al. 2013). Nevertheless, dust in LAE galaxy ISM could cause some of the irregularity in LAE surface-brightness profiles (Buck et al. 2017). With limited resolution, however, it is difficult to make this distinction. A further challenge to morphological studies is that the clump sizes are near the resolution limit of instrumental point spread functions (PSFs) and often cannot be distinguished from point sources (Guo et al. 2015). As a result, direct imaging studies cannot decisively determine whether the clumps are different in nature from star-forming regions in our local universe or if the larger apparent size is merely an artifact of insufficient resolution (Shibuya et al. 2014; Kobayashi et al. 2016; Tamburello et al. 2017; Fisher et al. 2017).","Citation Text":["Shibuya et al. 2014"],"Functions Text":["Efforts have been made in quantifying mass, star formation rates, gas composition, and kinematics, as well as other LAE properties (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[1016,1035]],"Functions Start End":[[789,926]]}
{"Identifier":"2018AandA...618A.128C__Croft_et_al._2005_Instance_1","Paragraph":"As described in Sect. 1, the galaxy protocluster associated with 7C 1756+6520 is characterized by a high fraction of AGN protocluster members: seven AGN, including the central radio galaxy, have been spectroscopically confirmed in close proximity both spatially and in redshift space of the protocluster. This high AGN fraction detected so far, ~23%, makes the overdensity around 7C 1756+6520 similar to the interesting and well-studied cluster around the radio galaxy PKS 1138\u2212262 from this point of view as well (Pentericci et al. 2002; Croft et al. 2005), in addition to the extension (see Sect. 5.2). The source PKS 1138\u2212262 is a massive forming radio galaxy at z ~ 2.16 (Pentericci et al. 1998) that is surrounded by overdensities of Ly\u03b1 emitters (Pentericci et al. 2000), extremely red objects (Kurk et al. 2004a; Koyama et al. 2013), H\u03b1 emitters (Kurk et al. 2004b), X-ray emitters (Pentericci et al. 2002), and an overdensity of dusty starbursts (Dannerbauer et al. 2014; Rigby et al. 2014), several of which are spectroscopically confirmed to be close to the radio galaxy redshift. Five of the 18 X-ray sources (~28%) detected by Pentericci et al. (2002) are AGN, including the central radio galaxy. From the soft X-ray luminosity function of AGN, Pentericci et al. (2002) estimated how many sources are expected in a given region of the cluster PKS 1138\u2212262, finding that it contains about twice the number of expected AGN. More recently, Pentericci et al. (2013) also found high AGN fractions by studying eight galaxy groups from z ~ 0.5 to z ~ 1.1. They found that the fraction of AGN with Log LH > 42 erg s\u22121 in galaxies with MR  \u221220 varies from less than 5% to 22%, with an average value of 6.3%, which is more than double the fraction for massive cluster at similar high redshifts (e.g., Overzier et al. 2005). Martini et al. (2013) estimated that the cluster AGN fraction in a sample of 13 clusters of galaxies at 1  z  1.5 is ~3% for AGN with rest-frame, hard X-ray luminosity greater than LX, H \u2265 1044 erg s\u22121. Based on these findings, the galaxy protocluster around 7C 1756+6520 seems to be a particularly interesting object.","Citation Text":["Croft et al. 2005"],"Functions Text":["This high AGN fraction detected so far, ~23%, makes the overdensity around 7C 1756+6520 similar to the interesting and well-studied cluster around the radio galaxy PKS 1138\u2212262 from this point of view as well"],"Functions Label":["Motivation"],"Citation Start End":[[539,556]],"Functions Start End":[[305,513]]}
{"Identifier":"2021MNRAS.500.2336Y__Lin_et_al._2020_Instance_2","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":["Lin et al. 2020"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[1355,1370]],"Functions Start End":[[931,1171]]}
{"Identifier":"2022AandA...658A.188S__Kreckel_et_al._2019_Instance_1","Paragraph":"We assumed a screen geometry and used PYNEB3 (Luridiana et al. 2015) to correct line fluxes for dust extinction via the H\u03b1\/H\u03b2 ratio, adopting the O\u2019Donnell (1994) reddening law with RV\u2004=\u20043.1 and a theoretical H\u03b1\/H\u03b2\u2004=\u20042.86. The extinction-corrected emission line luminosities of the H\u202fII regions were then computed using the distances reported in Table 1. For every H\u202fII region in our catalog, we also estimated the gas-phase metallicity and the gas ionization parameter by using extinction-corrected emission line fluxes. The gas-phase metallicity O\/H was calculated using the Pilyugin & Grebel (2016) S calibration (Scal hereafter). This calibration relies on three diagnostic line ratios (i.e., [N\u202fII]\/H\u03b2, [S\u202fII]\/H\u03b2, and [O\u202fIII]\/H\u03b2) and provides an empirical calibration against H\u202fII regions that have direct constraints on their nebular temperatures and hence their abundances. This calibration is relatively insensitive to changes in gas pressure or ionization parameter, and we adopted it as fiducial approach in this paper (see Kreckel et al. 2019, for a discussion). In addition, for each galaxy we fit the radial metallicity gradient by using an unweighted least-square linear fitting of the trend between 12\u2005+\u2005log(O\/H) and the deprojected galactocentric radius (see Fig. A.19). The gas ionization parameter represents the ratio between the ionizing photon flux density and the gas hydrogen density. In this paper, we express the ionization parameter as q\u2004=\u2004U\u2005\u00d7\u2005c\u2004=\u2004Q(H0)\/4\u03c0R2n, where c is the speed of light, U is the dimensionless ionization parameter, Q(H0) is the number of hydrogen ionizing photons (E\u2004> \u200413.6 eV) emitted per second, R is the empty (wind-blown) radius of the H\u202fII region, and n its hydrogen density. The ionization parameter is ultimately defined by the structure of an H\u202fII region (e.g., size, gas density, filling factor) and the properties of its ionizing source. Photoionization models show that it can be extracted via different diagnostic lines (Kewley & Dopita 2002; Dors et al. 2011). In this paper we use the calibration proposed by Diaz et al. (1991) based on the [S\u202fIII](9069+9532)\/[S\u202fII](6717+6713) line ratio. As the [S\u202fIII]\u03bb9532 \u00c5 line falls outside the wavelength range covered by MUSE, we assumed that [S\u202fIII]\u03bb9532 \u00c5\u2006= 2.47[S\u2006III]\u03bb9069 \u00c5 according to default atomic data in PYNEB (Luridiana et al. 2015). It should be noted that for about 3000 H\u202fII regions we are not able to estimate the ionization parameter due to lack of detection of the [S\u202fIII]\u03bb9532 \u00c5 emission line.","Citation Text":["Kreckel et al. 2019"],"Functions Text":["This calibration is relatively insensitive to changes in gas pressure or ionization parameter, and we adopted it as fiducial approach in this paper (see",", for a discussion)."],"Functions Label":["Background","Background"],"Citation Start End":[[1034,1053]],"Functions Start End":[[881,1033],[1053,1073]]}
{"Identifier":"2016AandA...587A.159G__Tian_et_al._2014_Instance_2","Paragraph":"One has to be sure to rule out cases where inorganic chemistry can mimic the presence of life (\u201cfalse positives\u201d). Potential abiotic ozone production on Venus- and Mars-like planets has been discussed by Schindler & Kasting (2000, and references therein). While this is based on photolysis of e.g., CO2 and H2O and is thus limited in extent, a sustainable production of abiotic O3 which could build up to a detectable level has been suggested by Domagal-Goldman & Meadows (2010) for a planet within the habitable zone of AD Leonis with a specific atmospheric composition. Indeed, other studies confirm that abiotic buildup of ozone is possible (e.g., Hu et al. 2012; Tian et al. 2014); however, detectable levels are unlikely if liquid water is abundant, as e.g. rainout of oxidized species would keep atmospheric O2 and O3 low (Segura et al. 2007), unless the CO2 concentration is high and both H2 and CH4 emissions are low (Hu et al. 2012). False-positive detection of molecules such as CH4 and O3 is discussed by von Paris et al. (2011). Seager et al. (2013) present a biosignature gas classification. Since abiotic processes cannot be ruled out for individual molecules (e.g. for O3), searches for biosignature molecules should search for multiple biosignature species simultaneously. It has been suggested that the simultaneous presence of O2 and CH4 can be used as an indication for life (Sagan et al. 1993, and references therein). Similarly, Selsis et al. (2002) suggest a so-called \u201ctriple signature\u201d, where the combined detection of O3, CO2 and H2O would indicate biological activity. Domagal-Goldman & Meadows (2010) suggest to simultaneously search for the signature of O2, CH4, and C2H6. Of course, care has to be taken to avoid combinations of biosignature molecules which can be generated abiotically together (see e.g. Tian et al. 2014). The detectability of biosignature molecules is discussed, e.g. by von Paris et al. (2011) and Hedelt et al. (2013). In particular, the simulation of the instrumental response to simulated spectra for currently planned or proposed exoplanet characterization missions has shown that the amount of information the retrieval process can provide on the atmospheric composition may not be sufficient (von Paris et al. 2013). ","Citation Text":["Tian et al. 2014"],"Functions Text":["Of course, care has to be taken to avoid combinations of biosignature molecules which can be generated abiotically together (see e.g."],"Functions Label":["Background"],"Citation Start End":[[1835,1851]],"Functions Start End":[[1701,1834]]}
{"Identifier":"2022ApJ...926...21B__Vida_et_al._2014_Instance_1","Paragraph":"Characterizing the differential rotation (DR) realized at the base and in the convective envelope of solar-type stars is central to the understanding of their magnetic field generation, activity level, and rotation, as it is directly linked to the \u03a9 effect (e.g., stretching of the poloidal magnetic field lines by large-scale shear). Hence, the role of DR in driving the star\u2019s magnetic activity level and field properties should be clarified (Donahue et al. 1996). Doppler imaging (Donati & Collier Cameron 1997; Barnes et al. 2005), asteroseismology (Gizon & Solanki 2004; Reinhold et al. 2013; Garc\u00eda et al. 2014), classical spot models (Lanza et al. 2014), and short-term Fourier transform (Vida et al. 2014) are methods to infer DR. The combination of all these observations on stellar rotation and magnetism helps constrain the trends linking rotation with stellar DR and magnetic activity. Various analyses of stellar DR revealed different dependencies between DR and star\u2019s rotation (\u0394\u03a9 \u221d \u03a9\nn\n), with n varying between 0.15 and 0.7 (Barnes et al. 2005; Reiners 2006; Reinhold et al. 2013). There is no clear consensus in the community for now; some authors are even advocating that such laws should be derived per spectral stellar classes and that the confusion comes from mixing together F and K stars (Balona & Abedigamba 2016). Saar (2011), Brandenburg & Giampapa (2018) also propose that the dependency of the DR with the rotation rate may not be monotonic, with a break near Rossby equals unity. By contrast, a more systematic and stronger dependency is observed with the star\u2019s temperature (\n\n\n\n\u0394\u03a9\u221dTeff8.92\n\n, Barnes et al. 2005, Reinhold et al. 2013; and \n\n\n\n\u0394\u03a9\u221dTeff8.6\n\n, Collier Cameron 2007). Hence, we expect large-scale shear to vary both in amplitude and profile (as a function of latitude and depth) as the global stellar parameters change. Some recent studies have confirmed this is happening in solar-type stars by inverting seismically their profile (Benomar et al. 2018), pointing to a possible antisolar DR state (e.g., slow equator\/fast poles), which was possibly already guessed in F stars (Reiners 2007) and advocated to exist in numerical simulations (Matt et al. 2011; Gastine et al. 2014; Brun et al. 2015, see below).","Citation Text":["Vida et al. 2014"],"Functions Text":["short-term Fourier transform","are methods to infer DR."],"Functions Label":["Background","Background"],"Citation Start End":[[696,712]],"Functions Start End":[[666,694],[714,738]]}
{"Identifier":"2019ApJ...886...15T__Gruendl_&_Chu_2009_Instance_1","Paragraph":"ALMA is capable of resolving internal structures of molecular clouds even in external galaxies. In particular, the Large Magellanic Cloud (LMC) is an ideal laboratory to investigate high-mass star formation thanks to its nearly face-on view (Balbinot et al. 2015) and the close distance, \u223c50 kpc (Schaefer 2008; de Grijs et al. 2014). It is also a great advantage to directly compare the distributions of molecular gas observed by ALMA and positions of massive YSOs identified by Spitzer and Herschel (e.g., Gruendl & Chu 2009; Chen et al. 2010; Seale et al. 2014) without any serious contamination in the line of sight. Earlier studies using the H i gas observations by Fukui et al. (2017) found that there are supergiant shells (Kim et al. 1999, 2003) and kiloparsec-scale gas flows caused by the last tidal interaction between the LMC and the Small Magellanic Cloud (SMC). Therefore, we may be able to examine the relation between such large-scale gas kinematics and the local star formation activities. Our present target in this paper is the N159W-South clump, which was discovered by our previous ALMA Cycle 1 observations (Fukui et al. 2015, hereafter Paper I) with an angular resolution of \u223c1\u2033 (\u223c0.24 pc) toward a GMC in the N159W region (e.g., Johansson et al. 1998; Minamidani et al. 2008, 2011). Paper I revealed that the GMC is composed of many filamentary molecular clouds and discovered the first example of protostellar outflows in the external galaxies. Paper I also found that the protostellar source with a stellar mass of \u223c37 M\u2299 in the N159W-South clump is located toward an intersection of two filaments and suggested that the filament\u2013filament collision triggered the protostar formation. Although the ALMA observations significantly improved our understanding of molecular cloud structures and star formation in this object, much higher-angular-resolution studies are needed to further resolve the filamentary structures down to a width of \u22720.1 pc (see, Arzoumanian et al. 2011, 2019) and investigate the star formation activities therein.","Citation Text":["Gruendl & Chu 2009"],"Functions Text":["It is also a great advantage to directly compare the distributions of molecular gas observed by ALMA and positions of massive YSOs identified by Spitzer and Herschel (e.g.,","without any serious contamination in the line of sight."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[508,526]],"Functions Start End":[[335,507],[565,620]]}
{"Identifier":"2019AandA...627A..53H__Sutherland_&_Bicknell_2007_Instance_1","Paragraph":"A spatial coincidence of the radio jet morphology and velocity dispersion of the ionised gas has already been reported for spatially-resolved spectroscopy of more luminous radio-quiet AGN (e.g. Husemann et al. 2013; Villar-Mart\u00edn et al. 2017) and powerful compact radio sources (e.g. Roche et al. 2016), but it has been correctly proposed that the fast moving plasma itself can lead to radio emission that mimics jet activity (Zakamska & Greene 2014; Hwang et al. 2018). In the case of HE 1353\u22121917 we can rule out that the ionised plasma is creating the radio emission because the high-velocity ionised gas traced by [O\u202fIII] is significantly displaced compared to the observed jet-like radio emission. Hence, we think that the radio jet is transferring its energy and momentum to the ambient medium through an extended shock front, which creates turbulence in a dense clumpy ISM. Such a great impact of the radio jet has been observationally shown in many cases (e.g. Villar-Mart\u00edn et al. 1999, 2014; O\u2019Dea et al. 2002; Nesvadba et al. 2006; Holt et al. 2008; Guillard et al. 2012; Harrison et al. 2015; Santoro et al. 2018; Tremblay et al. 2018; Jarvis et al. 2019) and theoretically supported through detailed hydrodynamic simulations (e.g. Krause & Alexander 2007; Sutherland & Bicknell 2007; Wagner & Bicknell 2011; Wagner et al. 2012; Cielo et al. 2018; Mukherjee et al. 2018). As we discussed in Sect. 3.6, the jet power alone is sufficient to energetically drive the outflow because only a small fraction of the AGN luminosity would impact the thin disc of the galaxy implying conversion efficiencies of more than 10% of Lbol. Hopkins & Elvis (2010) proposed a two-stage process for efficient radiation-driven outflows. They describe a scenario in which an initial weak wind in the hot gas phase, possibly initiated by an accretion disc wind or a radio jet, creates additional turbulence in the surrounding medium so that massive gas clouds will subsequently expand and disperse. This expansion of gas clouds would significantly increase their apparent cross-section with respect to incident radiation field of the AGN. Such a two-stage process may increase the coupling efficiency by an order of magnitude. While we cannot directly confirm this process with our observations, the close alignment of the jet axis and the ionisation cone greatly suggest that the outflow is driven jointly by both mechanical and radiative energy with an unknown ratio of the two. The open question is whether the same powerful outflow could have developed without the fast radio jet impacting the cold gas directly given its unique orientation.","Citation Text":["Sutherland & Bicknell 2007"],"Functions Text":["Hence, we think that the radio jet is transferring its energy and momentum to the ambient medium through an extended shock front, which creates turbulence in a dense clumpy ISM. Such a great impact of the radio jet has been observationally shown in many cases","and theoretically supported through detailed hydrodynamic simulations (e.g."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1269,1295]],"Functions Start End":[[703,962],[1168,1243]]}
{"Identifier":"2018AandA...616A..11G__Quinn_et_al._1993_Instance_1","Paragraph":"In addition to secular evolutionary processes, a disc galaxy like ours is expected to have experienced several accretion events in its recent and early past (Bullock & Johnston 2005; De Lucia & Helmi 2008; Stewart et al. 2008; Cooper et al. 2010; Font et al. 2011; Brook et al. 2012; Martig et al. 2012; Pillepich et al. 2015; Deason et al. 2016; Rodriguez-Gomez et al. 2016). While some of these accretions are currently being caught in the act, like for the Sagittarius dwarf galaxy (Ibata et al. 1994) and the Magellanic Clouds (Mathewson et al. 1974; Nidever et al. 2010; D\u2019Onghia & Fox 2016), we need to find the vestiges of ancient accretion events to understand the evolution of our Galaxy and how its mass growth has proceeded over time. Events that took place in the far past are expected to have induced a thickening of the early Galactic disc, first by increasing the in-plane and vertical velocity dispersion of stars (Quinn et al. 1993; Walker et al. 1996; Villalobos & Helmi 2008, 2009; Zolotov et al. 2009; Purcell et al. 2010; Di Matteo et al. 2011; Qu et al. 2011; Font et al. 2011; McCarthy et al. 2012; Cooper et al. 2015; Welker et al. 2017), and second by agitating the gaseous disc from which new stars are born, generating early stellar populations with higher initial velocity dispersions than those currently being formed (Brook et al. 2004, 2007; Forbes et al. 2012; Bird et al. 2013; Stinson et al. 2013). These complementary modes of formation of the Galactic disc can be imprinted on kinematics-age and kinematics-abundance relations (Str\u00f6mberg 1946; Spitzer & Schwarzschild 1951; Nordstr\u00f6m et al. 2004; Seabroke & Gilmore 2007; Holmberg et al. 2007, 2009; Bovy et al. 2012a, 2016; Haywood et al. 2013; Sharma et al. 2014; Martig et al. 2016; Ness et al. 2016; Mackereth et al. 2017; Robin et al. 2017), and distinguishing between them requires full 3D kinematic information for several million stars, in order to be able to separate the contribution of accreted from in-situ populations, and to constrain impulsive signatures that are typical of accretions (Minchev et al. 2014) versus a more quiescent cooling of the Galactic disc over time. Accretion events that took place in the more recent past of our Galaxy can also generate ripples and rings in a galactic disc (G\u00f3mez et al. 2012b), as well as in the inner stellar halo (Jean-Baptiste et al. 2017). Such vertical perturbations of the disc are further complicated by the effect of spiral arms (D\u2019Onghia et al. 2016; Monari et al. 2016b), which together with the effect of accretion events might explain vertical wave modes as observed in SEGUE andRAVE (Widrow et al. 2012; Williams et al. 2013; Carrillo et al. 2018), as well as in-plane velocity anisotropy (Siebert et al. 2012; Monari et al. 2016b). Mapping the kinematics out to several kiloparsec from the Sun is crucial for understanding whether signs of these recent and ongoing accretion events are visible in the Galactic disc, to ultimately understand to what extent the Galaxy can be represented as a system in dynamical equilibrium (H\u00e4fner et al. 2000; Dehnen & Binney 1998), at least in its inner regions, or to recover the nature of the perturber and the time of its accretion instead from the characteristics and strength of these ringing modes (G\u00f3mez et al. 2012b).","Citation Text":["Quinn et al. 1993"],"Functions Text":["Events that took place in the far past are expected to have induced a thickening of the early Galactic disc, first by increasing the in-plane and vertical velocity dispersion of stars"],"Functions Label":["Background"],"Citation Start End":[[931,948]],"Functions Start End":[[746,929]]}
{"Identifier":"2017ApJ...834...20A__Temi_et_al._2007a_Instance_2","Paragraph":"Lenticular galaxies seems to have a wider range of properties compared to ellipticals that resemble more the old definition of ETGs. However, even in ellipticals, large differences prevail. Recent observations of elliptical galaxies with Spitzer and Herschel (Temi et al. 2005, 2007a, 2007b, 2009; Smith et al. 2012; Agius et al. 2013; Mathews et al. 2013) have revealed that the far-infrared (FIR) luminosity LFIR from these galaxies can vary by \u223c100 among galaxies with similar optical luminosity. The 70 \u03bcm band luminosities (from Temi et al. 2007a, 2009), is a good example of such a huge scatter in the FIR luminosity of elliptical galaxies. Some of the high L70 galaxies are members of a small subset of ellipticals having radio detections of neutral and molecular gas. A few others may be S0 galaxies which, because of their rotationally supported disks, often contain large masses of cold gas and dust. Ellipticals containing large excess masses of dust and cold gas probably result from significant galaxy mergers in the past. However a fraction of elliptical galaxies appear to be completely normal but L70 in these galaxies still ranges over a factor of \u223c30, far larger than can be explained by uncertainties in the estimate of the FIR spectral energy distribution (SED) due to local stellar mass loss. While a significant fraction of the cold gas mass in low- to intermediate-mass ETGs is thought to have an external, merger-related origin (e.g., Davis et al. 2011), in the most massive ETGs the cold gas phases are presumably generated internally (Davis et al. 2011; David et al. 2014; Werner et al. 2014). Mergers with gas- and dust-rich galaxies have often been suggested for the origin of dust in all elliptical galaxies (e.g., Forbes 1991). Although the merger explanation is almost certainly correct in some cases, mergers cannot explain most of the observed scatter in L70. A crucial element in our understanding of the evolution of galaxies toward ETGs is the mutual role played by the major merging of galaxies and the secular star formation quenching. Neither of these scenarios yet accounts for all the observational evidence, and one could assume both contributing to some extent.","Citation Text":["Temi et al. 2007a"],"Functions Text":["The 70 \u03bcm band luminosities (from","is a good example of such a huge scatter in the FIR luminosity of elliptical galaxies."],"Functions Label":["Background","Background"],"Citation Start End":[[534,551]],"Functions Start End":[[500,533],[560,646]]}
{"Identifier":"2021MNRAS.501.4148L__Buchhave_et_al._2012_Instance_1","Paragraph":"We derived the photospheric stellar parameters using three different techniques: the curve-of-growth approach, spectral synthesis match, and empirical calibration. The first method minimizes the trend of iron abundances (obtained from the equivalent width, EW, of each line) with respect to excitation potential and reduced EW respectively, to obtain the effective temperature and the microturbulent velocity, \u03bet. The gravity log\u2009g is obtained by imposing the same average abundance from neutral and ionized iron lines. We obtained the EW measurements using ARESv2\n 5 (Sousa et al. 2015). We used the local thermodynamic equilibrium (LTE) code MOOG\n 6 (Sneden 1973) for the line analysis, together with the ATLAS9 grid of stellar model atmosphere from Castelli & Kurucz (2003). The whole procedure is described in more detail in Sousa (2014). We performed the analysis on a co-added spectrum (SNR > 600), and after applying the gravity correction from Mortier et al. (2014) and adding systematic errors in quadrature (Sousa et al. 2011), we obtained Teff = 5346 \u00b1 69\u2009K, log\u2009g =4.60 \u00b1 0.12, [Fe\/H] =\u22120.40 \u00b1 0.05, and \u03bet =0.78 \u00b1 0.08\u2009km\u2009s\u22121. The spectral synthesis match was performed using the Stellar Parameters Classification tool (SPC, Buchhave et al. 2012, 2014). It determines effective temperature, surface gravity, metallicity, and line broadening by performing a cross-correlation of the observed spectra with a library of synthetic spectra, and interpolating the correlation peaks to determine the best-matching parameters. For technical reasons, we ran the SPC on the 62 GTO spectra only7: the SNR is so high that the spectra are anyway dominated by systematic errors, and including the A40TAC_23 spectra would not change the results. We averaged the values measured for each exposure, and we obtained Teff =5389 \u00b1 50\u2009K, log\u2009g =4.49 \u00b1 0.10, [M\/H] = \u22120.36 \u00b1 0.08, and v\u2009sin\u2009i  2\u2009km\u2009s\u22121. We finally used CCFpams,8 a method based on the empirical calibration of temperature, metallicity, and gravity on several CCFs obtained with subsets of stellar lines with different sensitivity to temperature (Malavolta et al. 2017b). We obtained Teff = 5293 \u00b1 70\u2009K, log\u2009g =4.50 \u00b1 0.15, and [Fe\/H] = \u22120.40 \u00b1 0.05, after applying the same gravity and systematic corrections as for the EW analysis. We list the final spectroscopic adopted values, i.\u2009e. the weighted averages of the three methods, in Table 3. From the co-added HARPS-N spectrum, we also derived the chemical abundances for several refractory elements (Na, Mg, Si, Ca, Ti, Cr, Ni). We used the ARES + MOOG method assuming LTE, as described earlier. The reference for solar values was taken from Asplund et al. (2009), and all values in Table 3 are given relative to the Sun. Details on the method and line lists are described in Adibekyan et al. (2012) and Mortier et al. (2013). This analysis shows that this iron-poor star is alpha-enhanced. Using the average abundances of magnesium, silicon, and titanium to represent the alpha-elements and the iron abundance from the ARES + MOOG method (for consistency), we find that [\u03b1\/Fe] = 0.23.","Citation Text":["Buchhave et al. 2012"],"Functions Text":["The spectral synthesis match was performed using the Stellar Parameters Classification tool (SPC,","It determines effective temperature, surface gravity, metallicity, and line broadening by performing a cross-correlation of the observed spectra with a library of synthetic spectra, and interpolating the correlation peaks to determine the best-matching parameters."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1238,1258]],"Functions Start End":[[1140,1237],[1267,1531]]}
{"Identifier":"2022MNRAS.514.1961R__Prochaska_&_Zheng_2019_Instance_1","Paragraph":"Along with the time-domain detections, we identified J173438.35-504550.4 as a potential host galaxy for FRB 20201123A using robust statistical treatment given the relatively small localization error region. At face value, the low redshift of J173438.35-504550.4 appears at odds with the large dispersion measure for FRB 20201123A (${\\rm DM}_{\\rm FRB}\\approx 434 \\, {\\rm pc \\, cm^{-3}}$). Our Galaxy, however, contributes ${\\rm DM}_{\\rm ISM}\\approx 200 \\, {\\rm pc \\, cm^{-3}}$ (NE2001 gives\u2009229\u2009${\\rm pc \\, cm^{-3}}$ and YMW16 gives 162 ${\\rm pc \\, cm^{-3}}$) from its interstellar medium and a presumed ${\\rm DM}_{\\rm Halo}\\sim 50 \\, {\\rm pc \\, cm^{-3}}$ from its halo (Prochaska & Zheng 2019). This leaves ${\\approx}180 \\, {\\rm pc \\, cm^{-3}}$ for the cosmos (DMcosmic) and the host (DMhost). At z = 0.05, the average cosmic contribution is $\\langle {\\rm DM}_{\\rm cosmic}\\rangle \\sim 42\\, {\\rm pc \\, cm^{-3}}$ (Macquart et al. 2020) but the intrinsic scatter in this quantity is predicted to be large. Adopting the best-fitting model to the Macquart relation by Macquart et al. (2020), the 95 per cent confidence interval is ${\\rm DM}_{\\rm cosmic}= [15,125] \\, {\\rm pc \\, cm^{-3}}$. Allowing for the maximum value of this interval (which would imply a significant foreground galaxy halo), we recover a minimum host contribution of ${\\rm DM}_{\\rm host, min} \\approx 60~\\rm pc~cm^{-3}$. This is consistent with estimates for host galaxy contributions from theoretical and empirical treatments (Prochaska & Zheng 2019; James et al. 2022). For a true DMcosmic value of this sightline closer to (or below) the mean, the host contribution would exceed $100 \\, {\\rm pc \\, cm^{-3}}$. Such values are inferred for other FRB hosts (e.g. FRB\u200920121102A; Tendulkar et al. 2017). In conclusion, we find no strong evidence to rule out the association with J173438.35-504550.4 based on its redshift and DMFRB. The significant host contribution to the DM, combined with the scattering in FRB 20201123A possibly originating in the host, shows that it shares similarities with other highly active, repeating FRBs like FRB 20121102A and FRB 20190520A and potentially resides in a turbulent and dense environment within the host.","Citation Text":["Prochaska & Zheng 2019"],"Functions Text":["Our Galaxy, however, contributes","from its interstellar medium and a presumed ${\\rm DM}_{\\rm Halo}\\sim 50 \\, {\\rm pc \\, cm^{-3}}$ from its halo"],"Functions Label":["Uses","Uses"],"Citation Start End":[[670,692]],"Functions Start End":[[388,420],[559,668]]}
{"Identifier":"2017ApJ...850...20G__Vida\u00f1a_2016_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":["Vida\u00f1a 2016"],"Functions Text":["On the one hand, it is more energetically favorable for the system to populate new degrees of freedom, such as hyperons"],"Functions Label":["Motivation"],"Citation Start End":[[933,944]],"Functions Start End":[[542,661]]}
{"Identifier":"2015ApJ...804..130C__Bertschinger_1985_Instance_1","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"],"Functions Text":["For a cosmic void, embedding requires the underdense interior to be \u201ccompensated\u201d by an overdense bounding ridge, i.e., a compensated void"],"Functions Label":["Background"],"Citation Start End":[[1365,1382]],"Functions Start End":[[1206,1344]]}
{"Identifier":"2019MNRAS.488.2825F__Hamers_et_al._2015_Instance_1","Paragraph":"On the other hand, the distribution of the orbital inclination of the third companion with respect to the inner binary in the triple i3 (bottom panel) is found to peak at \u223c100\u00b0, but with non-negligible tails. For comparison, isolated triples that merge due to the KL mechanism show a very pronounced peak at \u223c90\u00b0, with only a few mergers happening in low-inclination systems (Antonini et al. 2017). There are two possible caveats to this. First, possible resonances between nodal precession and KL oscillations can arise. As shown in Hamers & Lai (2017), this could make even low-inclination systems merge. Secondly, there are three different KL mechanisms competing, thus the eccentricity of the inner binary of the CO triple can be pumped up by the torque either of the SMBH or the third companion in the CO triple, for which the KL time-scale is shorter (e.g. Hamers et al. 2015). We define the following ratios\n(18)\r\n\\begin{eqnarray*}\r\n\\mathcal {R}_{12M}=\\frac{T_{KL}^{123}}{T_{KL}^{12M}}\r\n\\end{eqnarray*}\r\n(19)\r\n\\begin{eqnarray*}\r\n\\mathcal {R}_{123M}=\\frac{T_{KL}^{123}}{T_{KL}^{123M}}\r\n\\end{eqnarray*}\r\n(20)\r\n\\begin{eqnarray*}\r\n\\mathcal {R}_{M}=\\frac{T_{KL}^{123M}}{T_{KL}^{12M}}\\ ,\r\n\\end{eqnarray*}\r\nwhich we plot in Fig. 4 for the BH\u2013BH binaries that merge in our simulations in Models MW, \u03b1 = 2, a3, max = 50 au and different values of \u03b2. Clearly, the shape of the three distributions does not depend on the assumed value of \u03b2, and we find that $\\mathcal {R}_{12M}$, $\\mathcal {R}_{123M}$, and $\\mathcal {R}_{M}$ peak at \u223c10\u22124, \u223c5 \u00d7 10\u22123, and \u223c10\u22121, respectively. We note that, however, KL cycles in a given CO triple may be inactive if the relative inclinations are not in the KL window. As a consequence, the initial dynamical evolution of the CO triple can ultimately be dictated by a KL cycle that takes place on a long time-scale. The latter can in turn excite the relative inclination of one of the other orbits, which could activate KL cycles that take place on a shorter time-scale. Thus, the CO triple may experience rather different and rich dynamical paths, which tend to lead to chaotic behaviour (Grishin et al. 2018).","Citation Text":["Hamers et al. 2015"],"Functions Text":["Secondly, there are three different KL mechanisms competing, thus the eccentricity of the inner binary of the CO triple can be pumped up by the torque either of the SMBH or the third companion in the CO triple, for which the KL time-scale is shorter (e.g."],"Functions Label":["Uses"],"Citation Start End":[[863,881]],"Functions Start End":[[607,862]]}
{"Identifier":"2019AandA...629A..54U__Evans_et_al._2007_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":["Evans et al. 2007"],"Functions Text":["A soft excess below 1.5 keV is also present"],"Functions Label":["Background"],"Citation Start End":[[766,783]],"Functions Start End":[[721,764]]}
{"Identifier":"2018ApJ...860...24P__Warmuth_2015_Instance_3","Paragraph":"Figure 13 shows the temporal evolution of the density, \u03c1, plasma flow velocity, vx, position of the wave crest, PosA, phase speed, vw, and magnetic field component in the z-direction, Bz, for the primary waves in every different case of initial amplitude, \u03c1IA. In Figure 13(a), we observe that the amplitude of the density remains approximately constant at their initial values until the time when the shock is formed and the density amplitude of the primary wave starts decreasing (see Vr\u0161nak & Luli\u0107 2000), i.e., at t \u2248 0.03 (blue), t \u2248 0.04 (red), and t \u2248 0.055 (green). For the case of \u03c1IA = 1.3 (magenta), a decrease of the amplitude of the primary can hardly be observed, as expected for low-amplitude wave (Warmuth 2015). One can see that the larger the initial amplitude, \u03c1IA, the stronger the decrease of the primary wave\u2019s amplitude, which is consistent with observations (Warmuth & Mann 2011; Muhr et al. 2014; Warmuth 2015). The amplitudes decrease to values of \u03c1 \u2248 1.6 (blue), \u03c1 \u2248 1.5 (red), and \u03c1 \u2248 1.4 (green) until the primary wave starts entering the CH. Due to the fact that the waves with larger initial amplitude enter the CH earlier than those with small initial amplitude, we can see in Figure 13(a) that the tracking of the parameters of the faster waves stops at an earlier time than the one for the slower waves. A similar behavior to the one of the density, \u03c1, can be observed for the plasma flow velocity, vx, in Figure 13(b) and the magnetic field component, Bz, in Figure 13(e). Here, the amplitudes decrease from vx = 0.75, Bz = 1.9 (for \u03c1IA = 1.9, blue), vx = 0.6, Bz = 1.7 (for \u03c1IA = 1.7, red), vx = 0.45, Bz = 1.5 (for \u03c1IA = 1.5, green), and vx = 0.27, Bz = 1.3 (for \u03c1IA = 1.3, magenta) to vx = 0.55, Bz = 1.6 (for \u03c1IA = 1.9, blue), vx = 0.46, Bz = 1.5 (for \u03c1IA = 1.7, red), vx = 0.36, Bz = 1.4 (for \u03c1IA = 1.5, green), and vx = 0.25, Bz = 1.25 (for \u03c1IA = 1.3, magenta). Figure 13(c) shows how the primary waves propagate in the positive x-direction. In all four cases of different initial amplitude, \u03c1IA, the phase speed decreases slighty (consistent with observations; see Warmuth et al. 2004 and Warmuth 2015) until the waves enter the CH at different times, i.e., the values for the phase speed start at vw \u2248 2.2 (for \u03c1IA = 1.9, blue), vw \u2248 1.9 (for \u03c1IA = 1.7, red), vw \u2248 1.7 (for \u03c1IA = 1.5, green), and vw \u2248 1.4 (for \u03c1IA = 1.3, magenta) and decrease to vw \u2248 1.5 (for \u03c1IA = 1.9, blue), vw \u2248 1.39 (for \u03c1IA = 1.7, red), vw \u2248 1.2 (for \u03c1IA = 1.5, green), and vw \u2248 1.13 (for \u03c1IA = 1.3, magenta).","Citation Text":["Warmuth 2015"],"Functions Text":["In all four cases of different initial amplitude, \u03c1IA, the phase speed decreases slighty (consistent with observations; see"],"Functions Label":["Similarities"],"Citation Start End":[[2131,2143]],"Functions Start End":[[1983,2106]]}
{"Identifier":"2015AandA...580L...2Z__Fender_et_al._2000_Instance_1","Paragraph":"Synchrotron radiation emitted from one relativistic electron population with density Nrel gyrating along a magnetic field B produces a power-law radio spectrum (S \u221d \u03bd\u03b1) with spectral slope \u03b1thin 0 above a critical frequency \u03bdbreak that depends on B and Nrel. Below this critical frequency self-absorption effects become important and the emission becomes optically thick with a spectral slope \u03b1thick = 2.5 (e.g. Kaiser 2006). For instance, the spectrum of the radio outburst in the pulsar system PSR B1259\u221263 has a spectral index \u03b1 = \u22120.7, which is consistent with optically thin synchrotron emission (Fig. 3 in Connors et al. 2002). In contrast, the jets of microquasars often show flat\/inverted radio spectra as demonstrated by Fender (2001). Plasma and magnetic field variations along the jet in fact may create regions with different B and Nrel values resulting in spectral components with different turn-over frequencies \u03bdbreak. The overlap of the optically thin part of one spectral component with the self-absorbed part of the adjacent one will result in a flat spectrum when observed with a spatial resolution insufficient to resolve the jet (Kaiser 2006). In the microquasar Cygnus X-1, radio emission has been measured with a flat spectrum, i.e. \u03b1 = \u22120.06 \u00b1 0.05 and \u03b1 = 0.07 \u00b1 0.04 (Fender et al. 2000). In Cygnus X-1 the black hole is powered by accretion of the stellar wind of its supergiant companion star. Since the companion is close to filling its Roche lobe, the wind is not symmetric but strongly focused towards the black hole (Mi\u0161kovi\u010dov\u00e1 et al. 2011, and references therein). In contrast, as discussed above in the context of an accretion scenario, accretion in LS I +61\u00b0303  occurs in two particular orbital phases due to the eccentric orbit around the Be star. Is the accretion phase and subsequent ejection in LS I +61\u00b0303  able to develop a microquasar jet with a flat spectrum? As shown by Corbel et al. (2013), in the early phase of jet formation the low density particles in the jet produce an optically thin synchrotron power-law spectrum at GHz frequencies. Only when the density of the jet plasma increases, a transition to higher optical depth occurs resulting in flat\/inverted radio spectra (\u03b1 \u2265 0). In LS I +61\u00b0303 strong evidence for a flat radio spectrum does exist. Early VLA observations at four epochs revealed a flat spectrum between 1.5 and 22 GHz (Gregory et al. 1979). Furthermore, Strickman et al. (1998) carried out multi-frequency VLA observations sparsely covering one orbit and found deviations in the radio spectrum from a simple power law during the outburst. Finally, Massi & Kaufman Bernad\u00f3 (2009) measured even inverted, optically thick spectra between 2.2 GHz and 8.3 GHz. The primary aim of the current work is to extend previous radio observations of LS I +61\u00b0303 and systematically study \u2013 for the first time \u2013 broad-band cm\/mm (2.6\u201332 GHz) radio spectra and their evolution during a complete outburst. The densely sampled observations of LS I +61\u00b0303 were performed over a period of about three weeks (a total of 24\u2009days) using the Effelsberg 100 m telescope at a total of seven frequency bands. The paper is structured as follows. In Sect. 2 we briefly present the observations and data reduction procedures. In Sect. 3 we present the results. Section 4 provides a short discussion and our conclusions. ","Citation Text":["Fender et al. 2000"],"Functions Text":["In the microquasar Cygnus X-1, radio emission has been measured with a flat spectrum, i.e. \u03b1 = \u22120.06 \u00b1 0.05 and \u03b1 = 0.07 \u00b1 0.04"],"Functions Label":["Similarities"],"Citation Start End":[[1294,1312]],"Functions Start End":[[1165,1292]]}
{"Identifier":"2016AandA...591A..30L__Fall_et_al._(2010)_Instance_1","Paragraph":"The catalog reported by Lada & Lada (2003) for embedded clusters is consistent with the mass-size relation R ~ M1\/3 \u2212 M1\/2 for low-mass clusters as pointed out by Murray (2009), where R and M are the gas radius and mass, respectively. While the stellar mass of an embedded cluster is not directly observable and is obtained by assuming an underlying universal initial mass function (IMF), Adams et al. (2006) compiled data from Lada & Lada (2003) and Carpenter (2000) and showed a number-size relation \\hbox{$R_* \\sim N_*^{1\/2}$}R\u2217~N\u22171\/2 between the radius of the cluster and the number of objects it contains. The results of Gutermuth et al. (2009) are also compatible with this relation. The number-size relation is reasonably equivalent to the mass-size relation if we adopt a universal IMF and thus similar averaged mass among clusters. Despite the slight uncertainty in the power-law exponent, varying from 1\/3 to 1\/2, and the scatter of the data, it is clear that embedded clusters follow some mass-size relation. Larger data sets of star-forming clumps, identified as gaseous protoclusters, also exhibit a mass-size relation. Fall et al. (2010) compiled the observations of Shirley et al. (2003), Fa\u00fandez et al. (2004), and Fontani et al. (2005) in their Fig. 1, and found via least-squares regression the relation R \u221d M0.38. A regression fit on the ATLASGAL survey (Urquhart et al. 2014) data gives a dependence of R \u221d M0.50. In their work, they fitted log\u2009(M) to log\u2009(R) and found M \u221d R1.67, of which the power-law exponent is not the inverse of what we inferred. This indicates that the clump properties are differently dispersed in size and mass, and thus there exist some uncertainties in the power-law dependence. Meanwhile, both studies exhibit mass scatter of about 1 dex and are compatible with constant gas surface density, that is, R \u221d M0.5. These observations are performed with molecular lines and dust continuum of the star-forming gas, suggesting that the mass-size relation and probably some other properties of the stellar cluster are established as early as the gas-dominated phase. Pfalzner et al. (2016) pointed out that the mass-size relation for embedded clusters and gaseous protoclusters follow the same trend with different absolute value, which could be explained with star formation efficiency or cluster expansion. One interesting question to ask would be what physical processes actually determine this mass-size relation. Its existence suggests that this primary phase of cluster formation, the gaseous protocluster, is very likely in some equilibrium state governed by the molecular cloud environment in which it resides, and is crucial for understanding the nature of the cluster and, more generally, star formation. An analytical study by Hennebelle (2012) yielded, by linking the gaseous protocluster to properties of the parent cloud, R ~ M1\/3 or R ~ M1\/2 for protoclusters with different accretion schemes. We stress that we are emphasizing a global equilibrium, which does not imply that the structure is not locally collapsing. More precisely, we propose that the large scales are supported by a combination of rotation and turbulence that sets the M-R relation. ","Citation Text":["Fall et al. (2010)"],"Functions Text":["Larger data sets of star-forming clumps, identified as gaseous protoclusters, also exhibit a mass-size relation.","compiled the observations of Shirley et al. (2003), Fa\u00fandez et al. (2004), and Fontani et al. (2005) in their Fig. 1, and found via least-squares regression the relation R \u221d M0.38.","Meanwhile, both studies exhibit mass scatter of about 1 dex and are compatible with constant gas surface density, that is, R \u221d M0.5."],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1133,1151]],"Functions Start End":[[1020,1132],[1152,1332],[1727,1859]]}
{"Identifier":"2021MNRAS.500..291B__Longinotti_et_al._2015_Instance_1","Paragraph":"We have presented the analysis of the current X-ray observations of the disc wind in MCG-03-58-007. Here, multiple and variable wind components with velocities ranging from $\\sim \\! -0.08\\, c$ to $\\sim \\! - 0.2\\, c$ (and potentially up to $0.35\\, c$) are seen at different times. Multi-epoch observations of disc winds, like the one presented here, are crucial for revealing all the possible phases of the disc wind. For example, over a decade worth of observations of PDS\u2009456 revealed that the wind is most likely clumpy and\/or stratified with the ionization ranging from log\u2009($\\xi \/\\rm {erg\\, cm \\, s^{-1})}\\sim 2$\u2009erg cm s\u22121 up to log\u2009($\\xi \/\\rm {erg\\, cm \\, s^{-1})}=6$\u2009erg\u2009cm\u2009s\u22121 and velocities ranging from $\\sim \\! -0.2\\, c$ up to $\\sim \\! -0.46\\, c$ (Reeves et al. 2016, 2018a, 2020). It is possible, as suggested in other examples of ultra fast disc winds, that we are looking at a stratified wind, where multiple components are launched at different disc radii, but not all of them are always detected. This adds MCG-03-58-007 to the small but growing list of multiphase fast X-ray winds. Other examples of AGN with at least two variable phases of the X-ray winds are PG\u20091211+143 (Pounds et al. 2016; Reeves et al. 2018b), IRAS 13224-3809 (Parker et al. 2017; Chartas & Canas 2018; Pinto et al. 2018), 1H\u20090707-495 (Kosec et al. 2018), IRAS 17020+4544 (Longinotti et al. 2015), and PG 1114+445 (Serafinelli et al. 2019). In those cases, multiple phases with a common or different outflowing velocities are detected in the X-ray band. In contrast to most of the cases reported so far, neither of the two phases seen in MCG-03-58-007 requires a different ionization (aside from slice\u2009B) suggesting that we are seeing different streamlines of the same highly ionized flow. The only exception could be the eclipsing event seen in 2015, where a solution is found with a lower ionization for the Fe K intervening absorber. However, what we most likely see during this occultation event is a higher density and lower ionization clump of the wind, which could be faster because its higher opacity makes it easier to accelerate (Waters et al. 2017). Note that this does not imply that the soft X-ray wind components, like the ones seen for example in PDS\u2009456 or PG\u20091211+143, are not present; in contrast to the other examples, MCG-03-58-007 is seen through a relatively high column density (NH \u223c 2 \u00d7 1023\u2009cm\u22122, see Table 2) neutral absorber, therefore these phases may be hidden behind it. MCG-03-58-007 is not the only example where multiple Fe-K zones with the same ionization and outflowing with different velocities had been detected in a single observation. For instance, two simultaneous Fe-K phases were detected at least twice in PDS456 (Reeves et al. 2018a, 2020) and possibly in PG\u20091211+143 (Pounds et al. 2016) and IRAS\u200913349+2438 (Parker et al. 2020).","Citation Text":["Longinotti et al. 2015"],"Functions Text":["Other examples of AGN with at least two variable phases of the X-ray winds are","IRAS 17020+4544"],"Functions Label":["Background","Background"],"Citation Start End":[[1362,1384]],"Functions Start End":[[1099,1177],[1345,1360]]}
{"Identifier":"2016ApJ...827...31F__Das_&_Chakrabarti_2007_Instance_1","Paragraph":"Accretion physics has been extensively studied for decades, particularly in terms of the theoretical aspects including semi-analytic investigations as well as global numerical simulations, in an effort to further understand its physical nature and observational consequences. Many of the works on BH accretion have, in general, revealed an important generic feature of accretion, i.e., the formation of shocks as an accreting plasma is subject to outward forces via a number of decelerating mechanisms (e.g., Abramowicz & Prasanna 1990) and develops a shock front at r = rsh within the radius of the inner edge of a magnetized accretion disk7\n\n7\nArmitage et al. (2001) found an ISCO-like edge in their pseudo-Newtonian MHD accretion simulations.\n, perhaps equivalent to an innermost stable circular orbit (ISCO) for a pure HD Keplerian disk, before crossing an event horizon at r = rH. Previous studies include hydrodynamic shocks (e.g., Nobuta & Hanawa 1994; Lu et al. 1997; Chakrabarti 1990; Fukumura & Tsuruta 2004) and magnetohydrodynamic (MHD) shocks (e.g., Koide et al. 1998, 2000; Das & Chakrabarti 2007; Takahashi et al. 2002, 2006, hereafter T02, T06, respectively; Fukumura & Kazanas 2007b; Fukumura et al. 2007, hereafter F07; Takahashi & Takahashi 2010). In particular, extensive theoretical studies of various types of shocks have been conducted to date in an attempt to understand their dynamical behavior; e.g., shock oscillation in the context of quasi-periodic oscillations and its spectroscopic signatures (e.g., Chakrabarti & Titarchuk 1995; Molteni et al. 1996, 1999; Acharya et al. 2002; Okuda et al. 2004, 2007; Nagakura & Yamada 2008) that may be relevant for X-ray Binaries (XRBs), for example. Independent general relativistic (GR) MHD (GRMHD) simulations of the tilted accretion disk clearly show that the compression of the plunging plasma in the inner region(\n\n\n\n\n\n) leads to the formation of standing shocks (e.g., Fragile et al. 2007; Fragile & Blaes 2008; Generozov et al. 2014) depending on the characteristics of the disk geometry and the BH spin (e.g., Morales et al. 2014). The expected highly magnetized shocked region may perhaps correspond to the magnetically arrested plasma seen in other large-scale simulations (e.g., Tchekhovskoy et al. 2011).","Citation Text":["Das & Chakrabarti 2007"],"Functions Text":["Previous studies include","and magnetohydrodynamic (MHD) shocks"],"Functions Label":["Background","Background"],"Citation Start End":[[1088,1110]],"Functions Start End":[[886,910],[1019,1055]]}
{"Identifier":"2016ApJ...817...12P__Sur_et_al._2007_Instance_2","Paragraph":"Large-scale magnetic fields with strength of the order of 1\u201310 \u03bcG have been observed in disk galaxies (e.g., Beck et al. 1996; Fletcher 2010; Beck 2012; Beck & Wielebinski 2013; Van Eck et al. 2015). The origin of these fields can be explained through mean-field dynamo theory (Ruzmaikin et al. 1988; Beck et al. 1996; Brandenburg & Subramanian 2005a; Kulsrud & Zweibel 2008). The conservation of magnetic helicity is one of the key constraints in these models, and also leads to the suppression of the \u03b1-effect. The operation of the mean-field dynamo automatically leads to the growth of magnetic helicity of opposite signs between the large-scale and small-scale magnetic fields (Pouquet et al. 1976; Gruzinov & Diamond 1994; Blackman & Field 2002). To avoid catastrophic suppression of the dynamo action (\u03b1-quenching), the magnetic helicity due to the small-scale magnetic field should be removed from the system (Blackman & Field 2000, 2001; Kleeorin et al. 2000). Mechanisms suggested to produce these small-scale magnetic helicity fluxes are: advection of magnetic fields by an outflow from the disk through the galactic fountain or wind (Shukurov et al. 2006; Sur et al. 2007; Chamandy et al. 2014), magnetic helicity flux from anisotropy of the turbulence produced by differential rotation (Vishniac & Cho 2001; Subramanian & Brandenburg 2004, 2006; Sur et al. 2007; Vishniac & Shapovalov 2014), and through diffusive flux (Kleeorin et al. 2000, 2002; Brandenburg et al. 2009; Mitra et al. 2010; Chamandy et al. 2014). The outflow of magnetic helicity from the disk through dynamo operation leads to the formation of a corona (Blackman & Field 2000). According to Taylor's hypothesis, an infinitely conducting corona would resistively relax to force-free field configurations under the constraint of global magnetic helicity conservation (Woltjer 1960; Taylor 1974; Finn & Antonsen 1983; Berger & Field 1984; Mangalam & Krishan 2000). In this paper, we include advective and diffusive fluxes in a simple semi-analytic model of a galactic dynamo that transfers magnetic helicity outside the disk and consequently builds up a force-free corona in course of time. We first solve the time-dependent dynamo equations by expressing them as separable in variables r and z. The radial part of the dynamo equation is solved using an eigenvector expansion constructed using the steady-state solutions of the dynamo equation. The eigenvalues of the z part of the solution are obtained by solving a fourth-order algebraic equation, which primarily depends upon the turbulence parameters and the magnetic helicity fluxes. Once the dynamo solutions are written out as parametric functions of these parameters, the evolution of the mean magnetic field is computed numerically by simultaneously solving the dynamical equations for \u03b1-quenching and the growth of large-scale coronal magnetic helicity. Since the large-scale magnetic field lines cross the boundary between the galactic disk and the corona, the magnetic helicity of the large-scale magnetic field in the disk volume is not well defined. Hence we use the concept of gauge-invariant relative helicity (Finn & Antonsen 1983; Berger & Field 1984; Berger 1985) to estimate the large-scale magnetic helicity in the disk and the corona. Here the gauge-invariant relative helicity for the cylindrical geometry is calculated using the prescription given in Low (2006, 2011). We then investigate the dependence of the saturated mean magnetic field strength and its geometry on the magnetic helicity fluxes within the disk and the corresponding evolution of the force-free field in the corona.","Citation Text":["Sur et al. 2007"],"Functions Text":["magnetic helicity flux from anisotropy of the turbulence produced by differential rotation"],"Functions Label":["Background"],"Citation Start End":[[1358,1373]],"Functions Start End":[[1207,1297]]}
{"Identifier":"2022AandA...663A.105P__Brunetti_et_al._2008_Instance_1","Paragraph":"Regardless of the cluster orientation, the spectral index observed for the halo at all available frequencies suggests that it is a USSRH. Despite the number of detected USSRH is still low, radio halos with steep indices are being discovered more and more frequently in the last years thanks to the improved observational capabilities of low-frequency instruments such as GMRT, MWA (Murchison Widefield Array) and LOFAR (Shimwell et al. 2016; Wilber et al. 2018; Bruno et al. 2021; Di Gennaro et al. 2021; Duchesne et al. 2022). An in-depth analysis of all radio halos hosted in Planck clusters and observed in LoTSS, including A1550, has recently been presented in Botteon et al. (2022). USSRH are a prediction of turbulent re-acceleration models (Cassano et al. 2006; Brunetti et al. 2008), in which particles are re-accelerated by turbulence (Brunetti et al. 2001, 2017; Petrosian 2001; Brunetti & Lazarian 2011). On the other hand, the detection of such steep indices is not expected from hadronic (or secondary) models, in which the emission of halos comes from the production of secondary electrons from hadronic collisions between thermal and CR protons (Blasi & Colafrancesco 1999; Dolag & En\u00dflin 2000; Pfrommer et al. 2008). Given that the integrated spectral index observed for the USSRH with LOFAR is \n\n\n\n\n\u03b1\n\n54\n\nMHz\n\n\n144\n\nMHz\n\n\n\u223c\n\u2212\n1.6\n\n\n$ \\alpha_{54\\,\\rm MHz}^{144\\,\\rm MHz} \\sim -1.6 $\n\n\n, we expect an index for the spectral energy distribution8\u03b4\u2004=\u20042\u03b1\u2005\u2212\u20051\u2004=\u2004\u22124.2. If there is no break in the spectrum, the energy budget for these particles would be untenable (Brunetti et al. 2008). Therefore, a break at low energies (\u223cGeV) should exist, suggesting a possible interplay between radiative losses and turbulent re-acceleration during the lifetime of emitting electrons (Brunetti & Jones 2014). Moreover, re-acceleration models predict that a large fraction of halos associated with clusters of masses between 4 and 7\u2005\u00d7\u20051014\u2006M\u2299 should exhibit steep spectra (Cassano et al. 2010, 2012; Brunetti & Jones 2014; Cuciti et al. 2021). The mass of A1550 of \u223c6\u2005\u00d7\u20051014\u2006M\u2299 estimated from Planck Collaboration XXVII (2016) falls in this range9.","Citation Text":["Brunetti et al. 2008"],"Functions Text":["USSRH are a prediction of turbulent re-acceleration models"],"Functions Label":["Background"],"Citation Start End":[[769,789]],"Functions Start End":[[688,746]]}
{"Identifier":"2017MNRAS.470.2517H__D'Angelo_&_Lubow_2010_Instance_1","Paragraph":"We carry out analysis of the orbital properties of our clumps only using the sample as detected by the DDS, as this method is sensitive to most clump masses and semimajor axes. The total semimajor axis evolution of all clumps is shown in the left-hand panel of Fig. 3, which we have already discussed, and refer the reader back to. Circles mark surviving clumps (including clumps that subsume another clump), squares mark destroyed clumps and triangles mark merged clumps. Larger markers correspond to more massive clumps. For destroyed clumps, we take the last measured mass. Roughly half of our most massive clumps migrate radially inwards, which is consistent with migration in locally isothermal discs, as objects exchange angular momentum with the surrounding gas and move inwards. However, about half of our most massive clumps migrate radially outwards. This is known to be possible in radiative discs (Kley & Nelson 2012), but requires either large torques or steep surface density gradients (D'Angelo & Lubow 2010). Large torques can have many sources, but in massive, self-gravitating discs they are likely to be in the form of global spiral arms. We carried out a Fourier analysis on the density structure of our discs, to determine the Fourier amplitude of each m mode (where m is the number of spiral arms). The amplitude, Am, of each mode, m, is calculated by\n(4)\r\n\\begin{equation}\r\nA_{\\mathrm{m}} = \\Bigg | \\sum _{i=1}^{N_{\\mathrm{region}}} \\frac{{\\rm e}^{-im\\phi _i}}{N_{\\mathrm{region}}} \\Bigg |,\r\n\\end{equation}\r\nwhere Nregion is the number of particles in the region we are considering (for our case, R = 20 to R = 100 au), and \u03d5i is the azimuthal angle of the ith particle. Some example amplitudes are shown in Fig. 9, which shows the first 10 Fourier components of the density structure of two discs in their initial state (i.e. when they have just begun to fragment), marked in red, and the same two discs in their final state, marked in black. The discs are from simulation 1 and simulation 5, and their final state can be seen in their column density plots, shown in Fig. 1. These discs were selected because they ran for the same length of time, and they have contrasting final m-modes states, simulation 1 ultimately peaks in the m = 2 mode and simulation 5 ultimately peaks in the m = 6 mode.","Citation Text":["D'Angelo & Lubow 2010"],"Functions Text":["However, about half of our most massive clumps migrate radially outwards.","but requires either large torques or steep surface density gradients"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1001,1022]],"Functions Start End":[[787,860],[931,999]]}
{"Identifier":"2016ApJ...824..138Y__Frail_et_al._2013_Instance_1","Paragraph":"The above qualitative theoretical reasoning raises the question about why would Swift J1834.9\u22120846 be the only magnetar so far powering a wind nebula, given that previous searches around individual magnetars have returned no sign of extended emission attributable to wind nebulae (e.g., Vigan\u00f2 et al. 2014). With only one observed so far, it is difficult to draw any firm conclusions. Nevertheless, Swift J1834.9\u22120846 has some interesting characteristics that are not shared with the entire magnetar population. First, the environment of Swift J1834.9\u22120846 is extremely crowded, with a Fermi GeV source, an H.E.S.S. TeV source, an SNR, a GMC, and an OH maser in its vicinity (Frail et al. 2013; H.E.S.S. Collaboration et al. 2015). The relationship between all these sources is unclear. However, it is tempting to speculate that environmental effects from such a rich field could be playing a role in the production of this wind nebula (e.g., triggering of pair cascade by external gamma rays from a nearby source; Shukre & Radhakrishnan 1982; Istomin and Sob\u2019yanin 2011). Second, the Swift J1834.9\u22120846 X-ray luminosity in quiescence is \n\n\n\n\n\n erg s\u22121. Only five other magnetars (SGR 0418+5729, SGR 1745\u22122900, XTE J1810\u2212197, Swift J1822.3\u22121606, 3XMM J185246.6+003317)20\n\n20\n\nhttp:\/\/www.physics.mcgill.ca\/ pulsar\/magnetar\/main.html\n\n have luminosities \u22721032 erg s\u22121. Among these five, three have the smallest surface B fields measured (SGR 0418+5729, Swift J1822.3\u22121606, 3XMM J185246.6+003317; \n\n\n\n\n\n G), and only one source, SGR 1745\u22122900, has a rotational energy loss rate \n\n\n\n\n\n similar to Swift J1834.9\u22120846, while the rest have \n\n\n\n\n\n at least an order of magnitude lower. Hence, from an observational point of view, it seems that the combination of very weak X-ray luminosity, a magnetar-like B-field strength, and a somewhat large \n\n\n\n\n\n (properties that are only shared by the Galactic center magnetar SGR 1745\u22122900) may favor wind nebula production. Another possibility is that the Swift J1834.9\u22120846 magnetar\/nebula system is an older analog to the Kes 75 system, where the central pulsar evolves into a magnetar while preserving its originial PWN.","Citation Text":["Frail et al. 2013"],"Functions Text":["Nevertheless, Swift J1834.9\u22120846 has some interesting characteristics that are not shared with the entire magnetar population. First, the environment of Swift J1834.9\u22120846 is extremely crowded, with a Fermi GeV source, an H.E.S.S. TeV source, an SNR, a GMC, and an OH maser in its vicinity"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[676,693]],"Functions Start End":[[385,674]]}
{"Identifier":"2021MNRAS.503..594S__Tremaine,_Ostriker_&_Spitzer_1975_Instance_1","Paragraph":"Nuclear star clusters (NSCs) are dense and massive clusters observed with high frequency ($\\gt 80{{\\ \\rm per\\ cent}}$) at the centre of galaxies with stellar masses 109\u20131011 M\u2299 (e.g. B\u00f6ker et al. 2002, 2004; C\u00f4t\u00e9 et al. 2006; Turner et al. 2012; Baldassare et al. 2014; den Brok et al. 2014; Georgiev & B\u00f6ker 2014; S\u00e1nchez-Janssen et al. 2019; Pechetti et al. 2020). These extreme environments often harbour a central supermassive black hole (as in the case of our Galaxy), and represent the most dense stellar systems in the Universe. We refer to the recent and complete review by Neumayer, Seth & B\u00f6ker (2020) for more details about NSCs. There are two main formation channels that are thought to compete in NSC formation: in situ formation via fragmentation of gaseous clouds in the galactic centre (e.g. Loose, Kruegel & Tutukov 1982), or via orbital segregation and merger of massive star clusters that migrate towards the galactic centre via dynamical friction (Tremaine, Ostriker & Spitzer 1975; Capuzzo-Dolcetta 1993). The latter formation channel, named dry-merger scenario, has been widely explored theoretically and numerically (Capuzzo-Dolcetta & Miocchi 2008a, b). For instance, high-resolution N-body models suggested that the Milky Way (MW) NSC might have formed through this mechanism (Antonini et al. 2012; Tsatsi et al. 2017; Arca-Sedda et al. 2020), which might explain both the structure and kinematics of the Galactic NSC. The dry-merger scenario also provides a successful explanation for the potential formation of NSCs in young galaxies and for the seemingly absence of nucleated regions in small dwarf and massive ellipticals (Arca-Sedda & Capuzzo-Dolcetta 2014, 2017). In fact, several semi-analytical models have shown that the dry-merger scenario leads to correlations between the NSC and the host galaxy properties pretty similar to the observed one (Antonini 2013; Arca-Sedda & Capuzzo-Dolcetta 2014; Gnedin, Ostriker & Tremaine 2014; Capuzzo-Dolcetta & Tosta e Melo 2017). None the less, several features of NSCs seem hard to explain as the result of star cluster merging events. For instance, NSCs exhibit a complex star formation history that seems to suggest the occurrence of several episodic star formation events over the entire course of their lifetime (Neumayer et al. 2020). Such a feature can also be easily explained with in situ formation (B\u00f6ker et al. 2004), thus suggesting that the formation and evolution of NSCs is likely the result of both scenarios operating in concert. Here, we present our numerical approach to test this scenario in one of the clearest examples of two massive star clusters caught in the process of merging in the nucleus of the nearby MW-like spiral galaxy NGC\u20094654 shown in Fig. 1 and a zoom in on its nucleus in Fig. 2. Their projected separations, photometric mass, and assumption for the local velocity field of an MW-like galaxy suggest that they should be on a short, few tens of Myr, collision course before they completely merge (Georgiev & B\u00f6ker 2014). However, this pure analytical expectation needs to be tested in order to gain a deeper knowledge on the following: (1) what will be the merging time-scales given the cluster current observational properties; (2) how the physical and observational properties of the final product depend on the merger dynamics and how such properties compare to those of current NSCs in galaxies of similar mass and type as NGC\u20094654; and (3) in the hypothesis that the two clusters contain different stellar populations, what are the expected distributions and fractions of the stars coming from the two progenitors in the merger product.","Citation Text":["Tremaine, Ostriker & Spitzer 1975"],"Functions Text":["There are two main formation channels that are thought to compete in NSC formation:","or via orbital segregation and merger of massive star clusters that migrate towards the galactic centre via dynamical friction"],"Functions Label":["Background","Background"],"Citation Start End":[[968,1001]],"Functions Start End":[[641,724],[840,966]]}
{"Identifier":"2022AandA...659A..41E__Sarbadhicary_et_al._2017_Instance_1","Paragraph":"The age of a neutron star is difficult to measure, as for many other astronomical sources. The most robust way to do it is by identifying the birth supernova of the neutron star. However, this can be done precisely only for a very small number of objects, as 5\u201310 supernovae have historically been observed in our galaxy (Stephenson & Green 2005), and neutron stars are faint sources\u2013practically undetectable at distances beyond the Magellanic Clouds. The explosions, however, leave imprints in the interstellar medium that can remain visible at radio wavelengths for 10\u2005\u2212\u2005100 kyr (Sarbadhicary et al. 2017), thereby allowing the association of pulsars with supernova remnants (SNRs). However, pulsars are rarely found at the centre of SNRs (Frail et al. 1994), as most are expelled like bullets during the explosions possibly due to asymmetries in the process (e.g. Socrates et al. 2005). The transverse velocities of pulsars (based on proper motion and distance estimates) are particularly large, with a mean close to 310 km s\u22121 (Hobbs et al. 2005), which is at least ten times larger than the average velocities for stars in the solar neighbourhood (e.g. Gaia Collaboration 2018). Moreover, some measured velocities range as high as 1000 km s\u22121 (Chatterjee et al. 2005; Deller et al. 2019). Thus, associations between SNRs and pulsars are not always straightforward to make (e.g. see the chapter on young pulsars in Lyne & Graham-Smith 2012). The farther the pulsar is from the explosion site, the higher the possibility that the pulsar and SNR are unrelated. In order to confirm an association, it could be necessary to account for up to 100 kyr of evolution of the SNR (that we assume as the maximum possible age of a SNR), and movement across the Galaxy of the pulsar (e.g. Suzuki et al. 2021). In some situations, proper motion measurements for the pulsars can shed light on the matter. For an association to be secure, the pulsar must be moving away from where the explosion took place (usually adopted as the centre of the SNR), and the time necessary to move the pulsar to its current position must match the age of the system. If such time coincided with an independent age measurement of the SNR or the pulsar, or both, then the association would be concretely confirmed. However, this is rarely possible as SNR and pulsar ages are hard to obtain.","Citation Text":["Sarbadhicary et al. 2017"],"Functions Text":["The explosions, however, leave imprints in the interstellar medium that can remain visible at radio wavelengths for 10\u2005\u2212\u2005100 kyr","thereby allowing the association of pulsars with supernova remnants (SNRs)."],"Functions Label":["Background","Background"],"Citation Start End":[[582,606]],"Functions Start End":[[452,580],[609,684]]}
{"Identifier":"2021MNRAS.505.5427N__Buen-Abad_et_al._2018_Instance_1","Paragraph":"While the status of the S8 discrepancy is perhaps somewhat less clear than that of the H0 tension, it is beyond question that there overall is some disagreement between high- and low-redshift probes of the amplitude of matter fluctuations (see, for instance, Di Valentino et al. 2020c, for a concise review of the problem). It is thus worthwhile to investigate whether new physics might solve or at least alleviate the S8 discrepancy, a possibility that has been investigated in several works. Models that have been contemplated in this sense include for example active and sterile neutrinos (Battye & Moss 2014; MacCrann et al. 2015; Feng, Zhang & Zhang 2017; Vagnozzi et al. 2017; Mccarthy et al. 2018), ultra-light axions (Hlozek et al. 2015), decaying dark matter (DM; Enqvist et al. 2015; Chudaykin, Gorbunov & Tkachev 2018; Di Valentino et al. 2018; Abell\u00e1n et al. 2020; Chen et al. 2021; Pandey, Karwal & Das 2020; Xiao et al. 2020; Abell\u00e1n, Murgia & Poulin2021), extended or exotic DM and\/or dark energy (DE) models and interactions (Kunz, Nesseris & Sawicki 2015; Kumar & Nunes 2016; Pourtsidou & Tram 2016; Gariazzo et al. 2017; Benetti, Graef & Alcaniz 2018; Buen-Abad et al. 2018; Kumar, Nunes & Yadav 2018, 2020a, b; Poulin et al. 2018; Archidiacono et al. 2019; Di Valentino et al. 2019b, 2020b; Lambiase et al. 2019; Vagnozzi et al. 2019; Chamings et al. 2020; Dutta et al. 2020; Heimersheim et al. 2020; Jim\u00e9nez et al. 2020; Choi, Yanagida & Yokozaki 2021) including unified dark sector models (Camera, Martinelli & Bertacca 2019), modified gravity models (Dossett et al. 2015; De Felice & Mukohyama 2017; Nesseris, Pantazis & Perivolaropoulos 2017; Kazantzidis & Perivolaropoulos 2018, 2019; Barros et al. 2020; De Felice, Nakamura & Tsujikawa 2020; Skara & Perivolaropoulos 2020; Zumalacarregui 2020; Marra & Perivolaropoulos 2021), and more generally extended parameter spaces (Di Valentino & Bridle 2018; Di Valentino, Melchiorri & Silk 2020), among the others. It is also worth noting that most of the models invoked to address the S8 discrepancy do so at the expense of worsening the H0 tension, and vice versa (see e.g. Vagnozzi et al. 2018; Poulin et al. 2018; Kumar et al. 2019a; Hill et al. 2020; Alestas & Perivolaropoulos 2021), highlighting the importance of a conjoined analysis of the two tensions (Di Valentino et al. 2020b; Di Valentino, Linder & Melchiorri 2020a).","Citation Text":["Buen-Abad et al. 2018"],"Functions Text":["It is thus worthwhile to investigate whether new physics might solve or at least alleviate the S8 discrepancy, a possibility that has been investigated in several works.","Models that have been contemplated in this sense include for example","extended or exotic DM and\/or dark energy (DE) models and interactions"],"Functions Label":["Motivation","Background","Background"],"Citation Start End":[[1170,1191]],"Functions Start End":[[324,493],[494,562],[971,1040]]}
{"Identifier":"2021ApJ...917L..38K__Mushtukov_et_al._2015a_Instance_1","Paragraph":"With the increase in precision of CRSF measurements, the dependence of CRSFs on luminosity and spin phase has become obvious and is one of the important areas of current X-ray pulsar research, as it potentially allows probing of emission region geometry and properties (Staubert et al. 2019). Of particular interest is the luminosity dependence of the line properties, as it might be used to probe the mechanism of plasma deceleration and the transition between so-called sub- and supercritical accretion regimes associated with the onset of an accretion column first suggested by Basko & Sunyaev (1976). As a result, two main accretion regimes, super- and subcritical accretion, are expected depending on whether the source luminosity is higher or lower than a \u201ccritical luminosity\u201d Lcrit, which strongly depends on the magnetic field of the NS (Basko & Sunyaev 1976; Becker et al. 2012; Mushtukov et al. 2015a). In the supercritical regime, for L > Lcrit, an accretion column forms, and the falling plasma is decelerated by a radiation shock that forms at a certain distance from the NS surface. In this case, the emission height increases with increasing luminosity. In the subcritical regime, for L Lcrit, the infalling matter is presumably decelerated by Coulomb interaction, forming a region whose height decreases with increasing luminosity. At even lower luminosity, the description of the deceleration process of falling material is not very conclusive. A reasonable scenario is that there exists a transition for the regime from Coulomb-stopping to gas-shock dominance beyond which only a small accretion mound forms on the NS surface. Based on this picture, and considering that the local magnetic field strength decreases with height, the correlation between the CRSF line energy and luminosity can be used to trace the accretion regimes. For instance, a transition between the sub- and supercritical regimes was recently reported by Doroshenko et al. (2017) and Vybornov et al. (2018) for V0332+53 during a giant outburst.","Citation Text":["Mushtukov et al. 2015a"],"Functions Text":["As a result, two main accretion regimes, super- and subcritical accretion, are expected depending on whether the source luminosity is higher or lower than a \u201ccritical luminosity\u201d Lcrit, which strongly depends on the magnetic field of the NS"],"Functions Label":["Background"],"Citation Start End":[[889,911]],"Functions Start End":[[605,845]]}
{"Identifier":"2018AandA...609A.131G__Heithausen_2012_Instance_1","Paragraph":"Moreover, there could also be some contribution to the detected temperature asymmetry from high-latitude gas clouds in our Galaxy along the line of sight toward M\u200981. In this respect we note that M\u200981 is at about 40.9\u00b0 north of the Galactic disk, where contamination from the Milky Way is expected to be low. However, interpretation of astronomical observations is often hampered by the lack of direct distance information. Indeed, it is often not easy to judge whether objects on the same line of sight are physically related or not. Since the discovery of the Arp\u2019s Loop (Arp 1965) the nature of the interstellar clouds in this region has been debated; in particular whether they are related to the tidal arms around the galaxy triplet (Sun et al. 2005; de Mello et al. 2008) or to Galactic foreground cirrus (Sollima et al. 2010; Davies et al. 2010). Already Sandage et al. (1976) presented evidence showing that we are observing the M\u200981 triplet through widespread Galactic foreground cirrus clouds and de Vries et al. 1987 built large-scale HI, CO, and dust maps that showed Galactic cirrus emission toward the M\u200981 region with NH \u2243 1\u22122  \u00d7  1020 cm-2. The technique used to distinguish between the emission from extragalactic or Galactic gas and dust relies on spectral measurements and on the identification of the line of sight velocities, which are expected to be different in each case. Unfortunately, in the case of the M\u200981 Group, this technique appears hardly applicable since the radial velocities of extragalactic and Galactic clouds share a similar LSR (local standard of rest) velocity range (Heithausen 2012). Several small-area molecular clouds (SAMS), that is, tiny molecular clouds in a region where the shielding of the interstellar radiation field is too low (so that these clouds cannot survive for a long time), have been detected by Heithausen (2002) toward the M\u200981 Group. More recently, data from the Spectral and Photometric Imaging Receiver (SPIRE) instrument onboard Herschel ESA space observatory and Multiband Imaging Photometer for Spitzer (MIPS) onboard Spitzer allowed the identification of several dust clouds north of the M\u200981 galaxy with a total hydrogen column density in the range 1.5\u20135 \u00d7 1020 cm-2 and dust temperatures between 13 and 17 K (Heithausen 2012). However, since there is no obvious difference among the individual clouds, there was no way to distinguish between Galactic or extragalactic origin although it is likely that some of the IR emission both toward M\u200981 and NGC 3077 is of Galactic origin. Temperature asymmetry studies in Planck data may be indicative of the bulk dynamics in the observed region provided that other Local (Galactic) contamination in the data is identified and subtracted. This is not always possible, as in the case of the M\u200981 Group, and therefore it would be important to identify and study other examples of dust clouds where their origin, either Galactic or extragalactic, is not clear. One such example might be provided by the interacting system toward NGC 4435\/4438 (Cortese et al. 2010) where the SAMS found appear more consistent with Galactic cirrus clouds than with extragalactic molecular complexes. Incidentally, the region A1 within R0.50 has been studied by Barker et al. (2009), who found evidence for the presence of an extended structural component beyond the M\u200981 optical disk, with a much flatter surface brightness profile, which might contain \u224310\u201315% of the M\u200981 total V-band luminosity. However, the lack of both a similar analysis in the other quadrants (and at larger distances from the M\u200981 center) and the study of the gas and dust component associated to this evolved stellar population, hamper our understanding of whether this component may explain the observed temperature asymmetry toward the M\u200981 halo. ","Citation Text":["Heithausen 2012"],"Functions Text":["Unfortunately, in the case of the M\u200981 Group, this technique appears hardly applicable since the radial velocities of extragalactic and Galactic clouds share a similar LSR (local standard of rest) velocity range"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1609,1624]],"Functions Start End":[[1396,1607]]}
{"Identifier":"2022MNRAS.511.1750J__Zirker_1977_Instance_1","Paragraph":"Table 1 depicts the heliolatitudes of the solar offset points on different experiment days and the corresponding heliocentric distances. It is to be noted that owing to the trajectory constraints of the spacecraft along the ecliptic plane, a pure latitudinal dependence (at a constant heliocentric distance) of the spectral index cannot be assessed here. However, variations in the spectral index values with respect to the complementary heliolatitudes of the proximate points of the ray path between 5\u00b0 and 39\u00b0 N can be observed, where the higher-heliolatitude spectra are flatter (lower \u03b1f) than the lower-heliolatitude spectra (see Table 1). Low-heliolatitude coronal regions have streamer structures where the slow solar winds, which have a highly turbulent structure, are generated (Woo & Martin 1997). The fast solar winds emanate from coronal holes at higher heliolatitudes (Zirker 1977). The coronal holes, which generally have a higher magnetic field, lead to a lower sound wave to Alfv\u00e9nic wave ratio and a relatively lower turbulence cascading energy. Consequently, the turbulence requires a longer time (and hence larger heliocentric distances) to achieve a fully developed turbulence spectrum. Thus, at a given heliocentric distance (but varying heliolatitude) of a proximate point of the satellite radio ray path, the fast (high-latitude) solar wind has flatter spectra compared with slow (low-latitude) solar wind spectra, as reported in earlier studies (P\u00e4tzold et al. 1996; Efimov et al. 2008, 2017). Efimov et al. (2008) and Chashei et al. (2007) observed that this change of turbulence regime is prominent during periods of minimum solar activity, as coronal holes (the source of the high-latitude fast solar wind) exist in abundance, and there is relative lack of the active regions that generate slow solar winds. It is, however, intriguing that we notice this aspect in our study, which was conducted during a period of relatively low solar activity.","Citation Text":["Zirker 1977"],"Functions Text":["The fast solar winds emanate from coronal holes at higher heliolatitudes"],"Functions Label":["Uses"],"Citation Start End":[[882,893]],"Functions Start End":[[808,880]]}
{"Identifier":"2021MNRAS.503.3065S__Silverman_1986_Instance_1","Paragraph":"In this work, the difference between G and NG clusters plays a major role. In a first step, we focus our attention on characterizing the structure and distribution of galaxy member properties in each class. In Fig. 1, we show the distributions of: (a) logarithmic virial mass ($\\rm log(\\mathit{ M}_{200}\/M_{\\odot })$); (b) r-band absolute magnitude; (c) a proxy for the concentration of stellar mass in each cluster, defined as $R_{80}\/R_{20}$, where $R_{x}$ is the projected radius within which the stellar mass represents x per\u2009cent of the total stellar mass within $R_{200}$; (d) velocity dispersion along the LOS; and (e) the distribution of cluster mean stellar mass of bright galaxies within $R_{200}$ ($\\langle M_{\\mathrm{ stellar}}^{C} \\rangle$).3 We compare the distributions using two different statistical tests: AD and Wilcoxon Rank Test (Wlx; see Engmann & Cousineau 2011 and Gehan 1965 for a review of both).4 The results are shown in each panel. Aside from the histograms, we add a kernel smoothed curve, shown as shaded area, which is derived directly from the data set using an Epanechnikov Kernel Density estimator (Silverman 1986) with a bandwidth equal to 1.5 times the bin size. We find that G and NG clusters have statistically different distributions. In panel (a), we note that NG clusters tend to have higher values of $M_{200}$ in comparison to G clusters. Namely, we find that 32.4 per\u2009cent (11\/34) of NG clusters have $\\mathrm{ log}(M_{200}\/\\mathrm{ M}_{\\odot })\\gt 14.75$, while this fraction decreases to 6.3 per\u2009cent (9\/143) in G clusters. The majority of G systems ($\\sim 71.3{{\\ \\rm per\\ cent}}$) have $\\mathrm{ log}(M_{200}\/\\mathrm{ M}_{\\odot })\\lt 14.5$. Panel (b) shows an excess of fainter galaxies5 in NG clusters in comparison to G systems. We find that 30.5 per\u2009cent of NG cluster members have $M_{r} \\ge -19.5$, while the equivalent cut in absolute magnitude yields 14.1 per\u2009cent of galaxies in G systems. Panel (c) shows the distribution of concentration in G and NG clusters, revealing that NG cluster galaxies are less concentrated than their G cluster counterparts. Only one NG cluster reaches C > 4.7. In panel (d), we observe an excess of NG clusters with higher velocity dispersion in comparison to G clusters. NG clusters are presumed to be found in a non-virialized state, so that the expected velocity dispersions are higher, and, possibly, the estimated mass may also be an overestimate of the real cluster mass. In any case, this is one more piece of evidence regarding the different state of NG clusters with respect to G systems. Finally, we note in panel (f) an excess of NG clusters with mean stellar mass of $\\rm 3 \\le \\langle \\mathit{ M}_{stellar}^{\\mathit{ C}} \\rangle \\lt 4 \\times 10^{11}\\, M_{\\odot }$. In other words, we find that NG clusters have higher virial mass and radii, an excess of fainter galaxies, are less concentrated, contain more massive B galaxies and have higher velocity dispersion in comparison to G systems.","Citation Text":["Silverman 1986"],"Functions Text":["Aside from the histograms, we add a kernel smoothed curve, shown as shaded area, which is derived directly from the data set using an Epanechnikov Kernel Density estimator","with a bandwidth equal to 1.5 times the bin size."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1134,1148]],"Functions Start End":[[961,1132],[1150,1199]]}
{"Identifier":"2022ApJ...927..106C__Motte_et_al._2007_Instance_1","Paragraph":"To identify the fibers seen in the PPV space, we apply the agglomerative clustering implementation in the scikit-learn package\n7\n\n\n7\n\nhttps:\/\/scikit-learn.org\/stable\/\n (Pedregosa et al. 2011) to the PPV points for each transition. The agglomerative clustering algorithm groups points in N-dimensional space via recursively merging points or clusters (i.e., groups of points) into higher-order clusters such that the pair of points or clusters to be merged minimally increases the linkage distance (Ward 1963). To make the clustering procedure adjustable, we introduce two parameters t\n\nx\n and t\n\nv\n in defining the linkage distance of two PPV points:\n7\n\n\n\nslink,ij=[tx(xi\u2212xj)]2+yi\u2212yj2+vi\u2212vjtv2,\n\nwhere x and y are the spatial coordinates in physical units, and v is the LoS velocity. For each transition, we adjust the parameters (t\n\nx\n, t\n\nv\n) and run the algorithm with the modified coordinates of the PPV points as input. By practice we found that the best values of (t\n\nx\n, t\n\nv\n) for identifying the fiber structures in the H13CO+, N2H+, and NH2D data are (2.8 pc, 1.0 km s\u22121pc\u22121), (0.5 pc, 0.8 km s\u22121pc\u22121), and (1.0 pc, 2.0 km s\u22121pc\u22121), respectively, with which the DR21(OH) ridge is decomposed into 3, 3, and 2 fibers in the PPV space. Figure 4 shows the identified fibers projected on the PoS and their PV plots, with the color coding representing different velocities. In the H13CO+ data, three fibers with roughly north\u2013south orientations are clearly seen; these fibers have distinct velocities, forming a trident with the junction approximately coincident with the DR21(OH) core. We name the three fibers as f1, f2, and f3 from east to west. Among the fibers, f1 is clearly tracing the central-densest part of the dust ridge well known from previous dust-continuum observations (Motte et al. 2007; Hennemann et al. 2012); f2 and f3 are relatively new, not seen in the dust emission, but previous single-dish H13CO+ (1 \u2212 0) observations showed that the ridge slight moves from east to west with the velocity increasing from \u22125 km s\u22121 to 0 km s\u22121 (Schneider et al. 2010), in a manner consistent with the positions and velocities of the three fibers identified here. In particular, f3 is also discernible in the H13CO+(1 \u2212 0) map in Schneider et al. (2010; see their Figure A3), though at a much lower resolution. Given the LoS velocities of the fibers and the existing observations suggesting that the DR21(OH) ridge is in a global collapse (Schneider et al. 2010), one may expect that f1 and f3 are on the far side and the near side along LoS, respectively, and that f2 is probably in the middle. The three fibers in the N2H+ line are less prominent yet are still clearly seen with positions and velocities consistent with those in the H13CO+ line, indicating that they are tracing the same physical entities. While the whole f1 and the southern part of f3 are detected, f2 is fragmented into several parts, probably due to the regional variations of the abundances of the two molecular tracers. On the other hand, in the NH2D line, the whole f1 (though fragmented into parts) and the northern part of f2 are detected, and f3 is completely absent. The appearance of the fibers are also different from that in the H13CO+ and N2H+ lines by their more compact morphologies and smaller widths. In addition, in contrast to what is seen in the H13CO+ and N2H+ lines, the southern end of f1 in the NH2D line is shifted to the west and does not coincide with the DR21(OH) core, indicating that active high-mass star formation has destroyed most of NH2D. The detection information of the fibers in the three transitions is summarized in Table 1.","Citation Text":["Motte et al. 2007"],"Functions Text":["Among the fibers, f1 is clearly tracing the central-densest part of the dust ridge well known from previous dust-continuum observations"],"Functions Label":["Uses"],"Citation Start End":[[1790,1807]],"Functions Start End":[[1653,1788]]}
{"Identifier":"2022MNRAS.509..212D__Farris_et_al._2015_Instance_1","Paragraph":"On the other hand, at separations \u2272 0.01 pc, smaller than those characteristics of spectroscopic binary candidates, many theoretical studies have predicted a significant variability in the observed nuclear light curve due to different physical processes. For example, in studies of the evolution of MBHBs in circumnuclear discs, a modulated gas inflow from the outer gas distribution periodically fuels the accretion discs within the Hill radii of the individual MBHs (that, being smaller than the surrounding circumbinary disc, are commonly referred to as \u2018mini-discs\u2019 in the literature, Artymowicz & Lubow 1994; Ivanov, Papaloizou & Polnarev 1999; Hayasaki, Mineshige & Ho 2008; Cuadra et al. 2009; Roedig et al. 2011, 2012; D\u2019Orazio, Haiman & MacFadyen 2013; Farris et al. 2015; D\u2019Orazio et al. 2016; Miranda, Mu\u00f1oz & Lai 2017; Tang, MacFadyen & Haiman 2017; Bowen et al. 2018; d\u2019Ascoli et al. 2018) as a consequence of the non-axisymmetric and time-dependent potential of the binary. Such modulated inflow could result in a similarly variable luminosity, depending on the properties of the inflowing gaseous streams and of the preexisting mini-discs (see the discussion in Sesana et al. 2012). An alternative cause of observed variability could be the plunging of a very eccentric secondary MBH on to the primary disc, as proposed by Valtonen et al. (2008) for the observed variability of OJ287. Finally, even in the absence of periodic inflows or very eccentric binaries (as expected in the case of a low-mass secondary; D\u2019Orazio et al. 2016; Duffell et al. 2020), variability can be caused by the relativistic Doppler boost of the emitted spectrum during the orbit of the MBHB, resulting in a variable flux observed in fixed observational bands, as proposed for PG 1302-102 in D\u2019Orazio, Haiman & Schiminovich (2015). This last model has the peculiarity of predicting different variability amplitudes at different wavelengths, as demonstrated for the UV versus optical light curves of PG 1302-102 (Xin et al. 2019).","Citation Text":["Farris et al. 2015"],"Functions Text":["For example, in studies of the evolution of MBHBs in circumnuclear discs, a modulated gas inflow from the outer gas distribution periodically fuels the accretion discs within the Hill radii of the individual MBHs (that, being smaller than the surrounding circumbinary disc, are commonly referred to as \u2018mini-discs\u2019 in the literature,","as a consequence of the non-axisymmetric and time-dependent potential of the binary."],"Functions Label":["Background","Background"],"Citation Start End":[[762,780]],"Functions Start End":[[255,588],[903,987]]}
{"Identifier":"2018MNRAS.478..126G___2018b_Instance_1","Paragraph":"It may be appropriate to single out at this point the recent and interesting works by Lin & Ishak (2017a,b) in which the authors run a so-called (dis)cordance test based on using a proposed index of inconsistency (IOI) tailored at finding possible inconsistencies\/tensions between two or more data sets in a systematic and efficient way. For instance, it is well known that there is a persistent discrepancy between the Planck CMB measurements of H0 and the local measurements based on distance ladder (Riess et al. 2016, 2018b). At the same time, if one compares what is inferred from Planck 2015 best-fitting values, the LSS\/RSD measurements generally assign smaller power to the LSS data parametrized in terms of the weighted linear growth rate f(z)\u03c38(z). This feature is of course nothing but the \u03c38-tension we have been addressing in this paper. It is therefore natural to run the IOI test for the different kinds of H0 measurements and also to study the consistency between the H0 and the growth data. For example, upon comparing the constraints on H0 from different methods, Lin & Ishak (2017b) observe a decrease of the IOI when the local H0 measurement is removed. From this fact they conclude that the local measurement of H0 is an outlier compared to the others, what would favour a systematics-based explanation. This situation is compatible with the observed improvement in the statistical quality of the fitting analysis by Sol\u00e0, G\u00f3mez-Valent & de Cruz P\u00e9rez (2017b,c) when the local H0 measurement is removed from the overall fit of the data using the RVM and the \u039bCDM. In this respect, let us mention that a recent model-independent analysis of data on cosmic chronometers and an updated compilation of SNIa seem to favour the lower range of H0 (G\u00f3mez-Valent & Amendola 2018), what would be more along the line of the results found here, which favour a theoretical interpretation of the observed \u03c38 and H0 tensions in terms of vacuum dynamics and in general of DDE (cf. Fig. 10).","Citation Text":["Riess et al","2018b"],"Functions Text":["For instance, it is well known that there is a persistent discrepancy between the Planck CMB measurements of H0 and the local measurements based on distance ladder"],"Functions Label":["Motivation"],"Citation Start End":[[503,514],[522,527]],"Functions Start End":[[338,501]]}
{"Identifier":"2021AandA...655A..12T__Tang_et_al._2017b_Instance_1","Paragraph":"Using the RADEX3 non local thermodynamic equilibrium (LTE) modeling program (van der Tak et al. 2007) with collisional rate coefficients from Wiesenfeld & Faure (2013), we modeled the relation between the gas kinetic temperature and the measured average of para-H2CO 0.5 \u00d7 [(322\u2013221 + 321\u2013220)\/303\u2013202] ratios, adopting a 2.73 K background temperature, an average observational linewidth of 4.0 km s\u22121, and column densities N(para-H2CO) = 2.7 \u00d7 1012 and 3.7 \u00d7 1012 cm\u22122 for N113 and N159W, respectively. The results are shown in Fig. 5. The values of the para-H2CO column density were obtained with APEX data (beam size ~30\u2033; Tang et al. 2017b), which cover similar regions. Different column densities of para-H2CO only weakly affect derived kinetic temperatures (see Fig. 3 in Tang et al. 2017b or Fig. 4 in Tang et al. 2018a; this was also shown in Fig. 13 and discussed in Sect. 4.3.1 of Mangum & Wootten 1993) as long as all lines are optically thin. Considering that the relation between the gas temperature and the para-H2CO line ratio may vary at different spatial densities (see Fig. 2 in Tang et al. 2017b), we modeled it at spatial densities 104, 105, and 106 cm\u22123 in Fig. 5. It appears that Tkin at n(H2) = 105 cm\u22123 is consistently lower than values at 104 and 106 cm\u22123 by \u227223% and \u227234%, respectively, for Tkin \u2272 100 K. Local thermodynamic equilibrium (LTE) is a good approximation for the H2CO level populations under optically thin and high-density conditions (Mangum & Wootten 1993; Tang et al. 2017a,b, 2018b). Following the method applied by Tang et al. (2017b) in their Eq. (2), we plot the relation between the LTE kinetic temperature, TLTE, and the para-H2CO (3\u20132) line ratio in Fig. 5. Apparently, TLTE agrees well with Tnon-LTE at volume densities n(H2) ~ 105 cm\u22123 as long as Tkin \u2272 100 K. Previous observations show that para-H2CO (3\u20132) is sensitive to gas temperature at density 105 cm\u22123 (Ginsburg et al. 2016; Immer et al. 2016; Tang et al. 2017b). The spatial density measured with para-H2CO (303\u2013202) and C18O (2\u20131) in N113 and N159W is n(H2) ~ 105 cm\u22123 on a size of ~30\u2033 (Tang et al. 2017b). Therefore, here we adopt 105 cm\u22123 as an averaged spatial gas density in the N113 and N159W regions.","Citation Text":["Tang et al. 2017b"],"Functions Text":["The values of the para-H2CO column density were obtained with APEX data (beam size ~30\u2033;"],"Functions Label":["Uses"],"Citation Start End":[[626,643]],"Functions Start End":[[537,625]]}
{"Identifier":"2018MNRAS.480.1796S__Mori\u0107_et_al._2010_Instance_1","Paragraph":"To distinguish the emission from AGNs and star formation we use the fact that star-forming galaxies (SFGs) are known to exhibit radio\u2013IR correlation across a wide range of luminosities and redshifts (see Condon 1992; Appleton et al. 2004; Basu et al. 2015). The correlation between 1.4 GHz luminosity and IR luminosity (monochromatic IR luminosity as well as bolometric IR luminosity between 8.0 and 1000 $\\mu$m) in SFGs is attributed to the fact that both radio and IR emission are closely related to star formation (Ivison et al. 2010). AGNs with predominant radio emission from jet deviate from radio\u2013IR correlation by showing radio-excess (Mori\u0107 et al. 2010; Del Moro et al. 2013). Therefore, we examine if our NLS1s show radio-excess in the radio\u2013IR correlation. We note that radio\u2013IR correlation can be represented as the ratio of IR flux to 1.4 GHz radio flux density i.e. q=log(SIR\/$S_{\\rm 1.4 \\, GHz}$) (see Appleton et al. 2004). For our radio-detected NLS1s we estimate q$_{\\rm 22\\, {\\mu }m}$=log($S_{\\rm 22 {\\mu }m}$\/$S_{\\rm 1.4 \\, GHz}$), where IR flux at 22 $\\mu$m is taken from Wide-field Infrared Survey Explorer (WISE),2 and 1.4 GHz flux density is taken from FIRST whenever available, otherwise from NVSS. WISE is an all sky survey carried out at four MIR photometric bands namely W1 [3.6 $\\mu$m], W2 [4.6 $\\mu$m], W3 [12 $\\mu$m], and W4 [22 $\\mu$m], with 5\u03c3 sensitivity of 0.08, 0.11, 1.0, and 6.0 mJy, and angular resolution of 6.1, 6.4, 6.5, and 12.0, respectively (Wright et al. 2010). Using the most recent version of WISE source catalogue (AllWISE3 data release) we obtain MIR counterparts for 481, 480, 440, and 354 of our 498 radio-detected NLS1s with SNR \u2265 5.0 in W1, W2, W3, and W4 bands, respectively. The WISE counterparts of our radio-detected NLS1s are searched within a circle of 2.0\u2009arcsec radius centred at SDSS optical positions. q$_{\\rm 22 {\\mu }m}$ is estimated using k-corrected fluxes, where IR spectral index is derived using W3 and W4 band fluxes, and radio spectral index is derived using 1.4 GHz and 150 MHz flux densities, whenever available, otherwise an average radio spectral index of \u22120.7 is considered.","Citation Text":["Mori\u0107 et al. 2010"],"Functions Text":["AGNs with predominant radio emission from jet deviate from radio\u2013IR correlation by showing radio-excess","Therefore, we examine if our NLS1s show radio-excess in the radio\u2013IR correlation."],"Functions Label":["Uses","Uses"],"Citation Start End":[[644,661]],"Functions Start End":[[539,642],[686,767]]}
{"Identifier":"2018AandA...609A...2M__Titov_&_D\u00e9moulin_(1999)_Instance_1","Paragraph":"The equilibrium between the MFR and the two magnetic charges can lead to an infinite twist on the surface of the MFR because no toroidal field exists outside the MFR. The magnetic field of the line current I0 is included to decrease this to reasonable values, typically less than 4\u03c0, suggested by many observations (Liu & Alexander 2009; Wang et al. 2015; Guo et al. 2012). Combining contributions from line current I0 and the poloidal current inside the MFR, the expression for the toroidal field is written as (7)\\begin{eqnarray} \\label{eq:bt} B_{t}&amp;\\approx&amp; \\frac{\\mu_0 I_0}{2\\pi}\\sqrt{\\frac{2 \\alpha I^2}{a^2 I_0^2}\\left(1-\\frac{r_{\\rm a}^2}{a^2}\\right)+\\frac{1}{R^2}},~r_{\\rm a}\\leq a,\\nonumber \\\\ &amp;=&amp; \\frac{\\mu_0 I_0}{2\\pi r_{\\perp}},~r_{\\rm a}>a. \\end{eqnarray}Bt\u2248\u03bc0I02\u03c02\u03b1I2a2I021\u2212ra2a2+1R2,ra\u2264a,=\u03bc0I02\u03c0r\u22a5,ra>a.For an analytical expression of the poloidal magnetic field, we follow the practice of Titov & D\u00e9moulin (1999), in which poloidal field is given in terms of the vector potential as in (8)\\begin{equation} B_{p}=-\\frac{\\partial A_{t}}{\\partial x}\\frac{\\vec{r}_{\\perp}}{r_{\\perp}}+\\left(\\frac{\\partial A_t}{\\partial r_{\\perp}}+\\frac{A_t}{r_{\\perp}}\\right)\\vec{x}. \\end{equation}Bp=\u2212\u2202At\u2202xr\u22a5r\u22a5+\u2202At\u2202r\u22a5+Atr\u22a5x.The toroidal At outside the MFR (\u03c1 \u2265 a) is (9)\\begin{equation} A_{t}(x,r_{\\perp}) \\approx \\frac{\\mu_{0}I}{2\\pi}\\sqrt{\\frac{R}{r_{\\perp}}}\\mathcal{A}(k), \\end{equation}At(x,r\u22a5)\u2248\u03bc0I2\u03c0Rr\u22a5\ud835\udc9c(k),and inside (\u03c1a) is (10)\\begin{equation} A_{t}(x,r_{\\perp}) \\approx \\frac{\\mu_{0}I}{2\\pi}\\sqrt{\\frac{R}{r_{\\perp}}} (\\mathcal{A}(k_{\\rm a})+\\mathcal{A}'(k_{\\rm a})(k-k_{\\rm a})). \\end{equation}At(x,r\u22a5)\u2248\u03bc0I2\u03c0Rr\u22a5(\ud835\udc9c(ka)+\ud835\udc9c\u2032(ka)(k\u2212ka)).Here, \\begin{eqnarray} &amp;&amp;\\mathcal{A}(k)=k^{-1}[(2-k^2)K(k)-2E(k)], \\\\ &amp;&amp;k=2\\sqrt{\\frac{r_{\\perp}R}{(r_{\\perp}+R)^2+x^2}}, \\\\ &amp;&amp;k_{\\rm a}=2\\sqrt{\\frac{r_{\\perp}R}{4r_{\\perp}R+a^2}}, \\end{eqnarray}\ud835\udc9c(k)=k-1[(2\u2212k2)K(k)\u22122E(k)],k=2r\u22a5R(r\u22a5+R)2+x2,ka=2r\u22a5R4r\u22a5R+a2,and its derivative (14)\\begin{equation} \\mathcal{A}'(k)=\\frac{(2-k^2)E(k)-2(1-k^2)K(k)}{k^2(1-k^2)}\\cdot \\end{equation}\ud835\udc9c\u2032(k)=(2\u2212k2)E(k)\u22122(1\u2212k2)K(k)k2(1\u2212k2)\u00b7These expressions contain the complete elliptic integrals of the first and the second kinds, K(k) and E(k). ","Citation Text":["Titov & D\u00e9moulin (1999)"],"Functions Text":["For an analytical expression of the poloidal magnetic field, we follow the practice of","in which poloidal field is given in terms of the vector potential as in (8)\\begin{equation} B_{p}=-\\frac{\\partial A_{t}}{\\partial x}\\frac{\\vec{r}_{\\perp}}{r_{\\perp}}+\\left(\\frac{\\partial A_t}{\\partial r_{\\perp}}+\\frac{A_t}{r_{\\perp}}\\right)\\vec{x}. \\end{equation}Bp=\u2212\u2202At\u2202xr\u22a5r\u22a5+\u2202At\u2202r\u22a5+Atr\u22a5x."],"Functions Label":["Uses","Uses"],"Citation Start End":[[921,944]],"Functions Start End":[[834,920],[946,1236]]}
{"Identifier":"2020AandA...633A.147B__Miville-Desch\u00eanes_et_al._2017_Instance_1","Paragraph":"Another possible indication of the different nature of the smallest clouds with respect to the largest ones is shown by their behaviour in the velocity dispersion vs. radius relation. This relation, known as the first Larson relation (Larson 1981), is expected to be of the form \u03c3v\u2004\u221d\u2004R\u03b2. The observed non-thermal motions in MCs are always supersonic (thermal motions for gas at T\u2004=\u200410 K is \u22430.1 km s\u22121), and the expected value of \u03b2, if the origin of these motions is purely supersonic turbulence, the so-called Burger turbulence, is \u03b2\u2004=\u20040.5 (McKee & Ostricker 2007). Various values of \u03b2 in MCs have been derived by several authors, ranging from \u03b2\u2004\u223c\u20040.38 (Larson 1981) to \u03b2\u2004=\u20040.6\u2005\u00b1\u20050.3 (Miville-Desch\u00eanes et al. 2017). Restricting the evaluation of the relation to the interstellar clouds of the outer Galaxy, previous estimates give \u03b2\u2004=\u20040.47\u2005\u00b1\u20050.08 (Sodroski 1991), \u03b2\u2004=\u20040.53\u2005\u00b1\u20050.03 (Brand & Wouterloot 1995), \u03b2\u2004=\u20040.45\u2005\u00b1\u20050.04 (May et al. 1997), and \u03b2\u2004=\u20040.53\u2005\u00b1\u20050.06 for MCs in the third quadrant of the Galaxy (Rice et al. 2016), all in agreement with the a Burger-like turbulence. However, the nature of the non-thermal motions in MCs is still debated (Krumholz et al. 2019), and it is not yet clear whether they originate solely from turbulence or from large-scale gravitational collapse (Ballesteros-Paredes et al. 2011; Traficante et al. 2018a,b; Merello et al. 2019). We plot the \u03c3v\u2013R relation for clouds of the FQS catalogue in Fig. 15. Two different regimes appear for clouds with R\u2004\u2265\u20042 pc and clouds with R\u2004 \u20042 pc. For MCs with R\u2004\u2265\u20042 pc the velocity dispersion increases with radius with an exponent \u03b2\u2004=\u20040.59, similar to what is expected from pure supersonic turbulence, and to that previously measured. On the other hand, for MCs with R\u2004 \u20042 pc the relation is almost flat, with \u03b2\u2004=\u20040.08. The non-thermal motions in the smallest clouds seem to be independent of cloud size. These condensations, likely to be only transient structures, are probably not formed by the turbulent ISM, which would lead to a Burger-like spectrum. It may be possible that these structures are the result of a large-scale effect, such as spiral-arm density-wave shock that injects the same amount of kinetic energy to all these objects. This large-scale effect would dominate over the local turbulence or the self-gravity within the smallest clouds, becoming increasingly irrelevant in the larger clouds.","Citation Text":["Miville-Desch\u00eanes et al. 2017"],"Functions Text":["Various values of \u03b2 in MCs have been derived by several authors, ranging from \u03b2\u2004\u223c\u20040.38","to \u03b2\u2004=\u20040.6\u2005\u00b1\u20050.3"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[686,715]],"Functions Start End":[[567,653],[668,684]]}
{"Identifier":"2020MNRAS.496.4468S__Chang_et_al._2018_Instance_1","Paragraph":"More interestingly, there are hints suggesting that the impacts of projection effects extend beyond richness misidentification. First, when Miyatake et al. (2016) reported a possible detection of assembly bias for subsamples of clusters divided based on the concentration of member galaxies, the signal appeared too large compared to theoretical predictions (Wechsler et al. 2006; Dalal et al. 2008) for the \u039bCDM model. Subsequent works found that this large apparent signal might be due to projection effects (Busch & White 2017; Zu et al. 2017; Sunayama & More 2019). Secondly, when More et al. (2016) reported a detection of the so-called \u2018splashback\u2019 radius, a physically motivated boundary of cluster haloes, for redMaPPer clusters, the location of the splashback radius was found to be smaller than theoretical expectations (also see Chang et al. 2018, for a similar detection). Follow-up studies (Shin et al. 2019; Z\u00fcrcher & More 2019) used samples of clusters selected based on the Sunyaev-Zel\u2019dovich effect and found a different location of the splashback radius consistent with theoretical expectations, albeit with larger errors, suggesting that the original location may have been impacted by projection effects. The recent analysis in Murata et al. (2020) further indicated that previous analyses of the splashback radius for optically selected clusters may suffer from projection effects. Third, Murata et al. (2018) developed a forward modelling approach to calibrate the richness\u2013mass relation from a joint measurement of cluster abundances and cluster lensing. However, they found that a population of less massive haloes, down to 1012h\u22121M\u2299, had to be introduced in order for the \u039bCDM model prediction to match the observations. Overall, these studies indicate that projection effects may impact not only cluster richnesses but also other cluster observables such as cluster lensing, which, if true, can render problematic the standard approach for cluster cosmology employed by photometric surveys.","Citation Text":["Chang et al. 2018"],"Functions Text":["More interestingly, there are hints suggesting that the impacts of projection effects extend beyond richness misidentification.","Secondly, when More et al. (2016) reported a detection of the so-called \u2018splashback\u2019 radius, a physically motivated boundary of cluster haloes, for redMaPPer clusters, the location of the splashback radius was found to be smaller than theoretical expectations (also see","for a similar detection)."],"Functions Label":["Motivation","Motivation","Motivation"],"Citation Start End":[[840,857]],"Functions Start End":[[0,127],[570,839],[859,884]]}
{"Identifier":"2022ApJ...927...89W__Stetson_1987_Instance_1","Paragraph":"IRIS always takes short exposure (with the shortest possible exposure time of 2.2 s) and long exposure with specified exposure time (14.5 s) one by one. In order to estimate and subtract the sky background, each field was observed in 10 dithered positions. Raw exposures were sky-subtracted and flat-fielded using standard IRAF (Tody 1986) routines. An astrometric solution was performed with Sextractor (Bertin & Arnouts 1996) and SCAMP (Bertin 2006), and then 10 dithered frames were resampled and combined with SWARP (Bertin 2010) into a single, final image (details of the pipeline used for calibrations can be found in Watermann 2012). Aperture photometry was performed with the DAOPHOT (Stetson 1987) package, and instrumental photometry was tied to the 2MASS system using the constant stars present in a given field as standards (usually more than 3 stars of brightness similar to our target with quality flag AAA in the 2MASS catalog). If there were no comparison stars of similar magnitude in a given field (cases of VY Pyx, SW Tau, and AL Vir), we used long exposures (14.5 s) to measure the brightness of comparison stars while the target was measured in the short exposure (2.2 s). We found a non-negligible color term in J band, and it amounts to \u22120.07 (J \u2212 K\n\ns\n). Internal precision of our photometry is at a level of 0.02 mag. In order to check correctness of our transformation to the 2MASS system, we compared the magnitudes of the constant stars present in the observed fields (transformed to the 2MASS system using an approach identical to our scientific objects) with the corresponding magnitudes from the 2MASS catalog. The observed fields offered very limited numbers of bright constant stars; thus for this test, we mostly used sources that were adopted as comparison stars in the transformation of the photometry of science targets. Each considered star was excluded from the set of comparison stars while transforming its own photometry. This test is presented in Figure 2. The mean difference between IRIS and 2MASS photometry is zero with the error on the mean of 0.002 mag in each band. We adopt this value as our zero-point uncertainty, which contributes to the systematic error of our calibration of PLRs.","Citation Text":["Stetson 1987"],"Functions Text":["Aperture photometry was performed with the DAOPHOT","package"],"Functions Label":["Uses","Uses"],"Citation Start End":[[693,705]],"Functions Start End":[[641,691],[707,714]]}
{"Identifier":"2020ApJ...897...38D__Gough_1990_Instance_1","Paragraph":"The problem of imaging global magnetic fields through helioseismology has been relatively underexplored, as studies have mainly focused on the solar rotational profile and other global and local flow fields (Lavely & Ritzwoller 1992; Basu et al. 1999; Giles 2000; Zhao & Kosovichev 2004; Hanasoge et al. 2012b, 2017). Some studies have analyzed the effects of local fields, such as those contained in a sunspot, on the reduction of wave power (Cally 2000; Schunker & Cally 2006), and changes in wave speed and flow patterns (Gizon et al. 2009; \u0160vanda et al. 2014; Khomenko & Collados 2015; Rabello-Soares et al. 2018; Braun 2019). Nevertheless, early attempts (e.g., Gough 1990) that were made to study the impact of global magnetic fields (Lorentz-stress perturbations) on the eigenfrequencies and eigenfunctions of the standard solar model (model S in Christensen-Dalsgaard et al. 1996) focused on the details of the forward problem of frequency shifts induced by an axisymmetric magnetic field not aligned with the rotation axis. Dziembowski & Goode (2004) considered near-surface, small-scale as well as a large-scale deep toroidal fields at the tachocline. However, lacking a formal relation between Lorentz stresses and the consequent seismic signatures (sensitivity kernels), these efforts suffered from the drawback of restricting the field geometry to make the problem tractable. Several numerical studies have also considered small perturbations around a magnetized background medium (Cameron et al. 2008, 2010; Hanasoge et al. 2012a; Schunker et al. 2013). Attempts at deducing temporal changes in magnetic field configurations from variations in angular velocity have been met with limited success (Antia et al. 2013). In a recent analysis, Cutler (2017) explored the potential of learning about magnetic fields from mode-coupling theory. Hanasoge (2017) derived analytical forms of the Lorentz-stress sensitivity kernels in the context of normal-mode coupling. This has made it possible to consider the treatment of a completely general magnetic field configuration as a perturbation around a hydrostatic background state. This forms the basis of the current study.","Citation Text":["Gough 1990"],"Functions Text":["Nevertheless, early attempts (e.g.,","that were made to study the impact of global magnetic fields (Lorentz-stress perturbations) on the eigenfrequencies and eigenfunctions of the standard solar model (model S in Christensen-Dalsgaard et al. 1996) focused on the details of the forward problem of frequency shifts induced by an axisymmetric magnetic field not aligned with the rotation axis."],"Functions Label":["Background","Background"],"Citation Start End":[[667,677]],"Functions Start End":[[631,666],[679,1032]]}
{"Identifier":"2018MNRAS.474.3162T__Gavazzi_et_al._2004_Instance_1","Paragraph":"We then used the fxcor task to determine the redshifts of selected objects in the field. This task allows to calculate the radial velocity through Fourier cross-correlation between the spectrum of the object under analysis and a reference (template) spectrum (Tonry & Davis 1979). Both spectra are continuum subtracted and Fourier filtered before doing the correlation, while dispersions are equalized by rebinning to the smallest dispersion. For this work, we used two reference spectra: that of NGC\u20094449 (Kennicutt 1992), an irregular galaxy that has well-defined emissions, and the spectrum of NGC\u20094387 (Gavazzi et al. 2004), an elliptical galaxy showing strong absorptions. Both templates were downloaded from NED. For those spectra where it was possible to establish a tentative redshift value, typical emission lines were identified, such as H\u2009\u03b3, [O\u2009ii] (3727\u2009\u00c5), H\u2009\u03b2, [O\u2009iii] (4959\u2009\u00c5) and [O\u2009iii] (5007\u2009\u00c5), and\/or absorptions, like Ca\u2009ii (H+K), Ca+Fe, Na\u2009i and H\u2009\u03b1. Finally, the redshift adopted for each object was computed from Gaussian fits to two or more such features in its spectrum. Three objects (slits #6, #19 and #23) turned out to be Galactic stars5 (we give their radial velocities in Table 1), while no reliable redshift value could be obtained for other five objects (slits #5, #11, #13, #17 and #20), due to the low S\/N ratio of their spectra (we will return to objects #5 and #11 in Section 3.2). Tentative redshifts for two of them (slits #17 and #20) were measured through just one emission line each, which we assumed to be [O\u2009ii] (3727\u2009\u00c5) and H\u2009\u03b4, respectively (based on their colours and spiral morphology). So, their tentative redshift values would be $z_{\\#17} \\simeq 0.476$ and $z_{\\#20} \\simeq 0.479$. No clear emission or absorption lines could be identified in the blazar's spectrum, besides telluric lines and diffuse interstellar bands (DIBs), so no definite redshift value could be established in this way for 3C\u200966A (see, however, Section 3.3 for our analysis on probable foreground absorptions on the blazar's spectrum). Redshifts were thus established for 15 (plus two just tentative) out of the 24 selected objects. Our results are shown in Table 1, where we also include the only galaxy with a previously published redshift within our GMOS field (G2 in Bow97).","Citation Text":["Gavazzi et al. 2004"],"Functions Text":["For this work, we used two reference spectra:","and the spectrum of NGC\u20094387","an elliptical galaxy showing strong absorptions."],"Functions Label":["Uses","Uses","Uses"],"Citation Start End":[[607,626]],"Functions Start End":[[443,488],[577,605],[629,677]]}
{"Identifier":"2018MNRAS.477..392L__Izotov_&_Thuan_1998_Instance_1","Paragraph":"Here, we present new Very Large Telescope (VLT) VIMOS (Le F\u00e9vre et al. 2003) observations of two star-forming dwarf galaxies using the integral field unit (IFU) spectroscopy mode (hereafter VIMOS-IFU). UM\u2009461 (the upper panel in Fig. 1) is a well-studied H\u2009ii\/BCD galaxy (e.g. Taylor et al. 1995; van Zee, Skillman & Salzer 1998; Lagos et al. 2011). This galaxy has been described as formed by two compact and off-centre giant H\u2009ii regions (GH\u2009iiR), some smaller star-forming regions spread across the galaxy disc and an external stellar envelope that is strongly skewed towards the south-west (Lagos et al. 2011). It has been classified as having a cometary-like morphology with an integrated subsolar metallicity of 12 + log(O\/H) = 7.73\u20137.78 (Masegosa, Moles & Campos-Aguilar 1994; Izotov & Thuan 1998; P\u00e9rez-Montero & D\u00edaz 2003). As in most H\u2009ii\/BCD galaxies, UM 461 has an underlying component of old stars (Telles & Terlevich 1997; Lagos et al. 2011) that exhibits an elliptical outer morphology. Deep Near-Infrared observations with the Gemini\/NIRI camera (Lagos et al. 2011) revealed that the star-formation activity in this galaxy is taking place in several star clusters with masses typically between \u223c104\u2009M\u2299 and \u223c106\u2009M\u2299. Fig. 1 shows the Kp band image of UM 461 obtained by Lagos et al. (2011). Using the same notation as Lagos et al. (2011), the main GH\u2009iiR (the brightest one in Fig. 1) in our study is composed of the star clusters nos. 2 and 3, while the faintest one is formed by star clusters nos. 5\u20137. Taylor et al. (1995) proposed that the SE tail in their H\u2009i image of UM\u2009461 was formed as a result of a tidal interaction with UM 462. However, higher resolution H\u2009i maps of UM\u2009461 by van Zee, Skillman & Salzer (1998) did not show the extended SE H\u2009i tail seen in the Taylor H\u2009i map. This discrepancy is attributed to solar interference in the Taylor map (van Zee, Skillman & Salzer 1998). Moreover, the age distribution of the star cluster population in UM 461 indicates that the current starburst has begun within the last few million years (Lagos et al. 2011). This current starburst time-scale is too short to realistically be attributed to a UM\u2009461\/UM\u2009462 interaction.","Citation Text":["Izotov & Thuan 1998"],"Functions Text":["It has been classified as having a cometary-like morphology with an integrated subsolar metallicity of 12 + log(O\/H) = 7.73\u20137.78"],"Functions Label":["Background"],"Citation Start End":[[784,803]],"Functions Start End":[[615,743]]}
{"Identifier":"2016AandA...587A.133G__Wise_et_al._2014_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":["Wise et al. 2014"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[804,820]],"Functions Start End":[[481,782]]}
{"Identifier":"2018AandA...609A..13K__Mucciarelli_et_al._(2017)_Instance_1","Paragraph":"Gaia 1 is a star cluster that was recently discovered by Koposov et al. (2017) in the first Gaia data release (Gaia Collaboration 2016), alongside with another system of lower mass. Its observation and previous detections were seriously hampered by the nearby bright star Sirius, which emphasized the impressive discovery power of the Gaia mission. This object was first characterized as an intermediate-age (6.3 Gyr) and moderately metal-rich (\u22120.7 dex) system, based on isochrone fits to a comprehensive combination of Gaia, 2MASS (Cutri et al. 2003), WISE (Wright et al. 2010), and Pan-STARRS1 (Chambers et al. 2016) photometry. Hence, this object was characterized by Koposov et al. (2017) as a star cluster, most likely of the globular confession. Further investigation of Gaia 1 found a metallicity higher by more than 0.5 dex, which challenged the previous age measurement and rather characterized it as a young (3 Gyr), metal-rich (\u22120.1 dex) object, possibly of extragalactic origin given its orbit that leads it up to ~1.7 kpc above the disk (Simpson et al. 2017). Subsequently, Mucciarelli et al. (2017) measured chemical abundances of six stars in Gaia 1, suggesting an equally high metallicity, but based on their abundance study, the suggestion of an extragalactic origin was revoked. While a more metal-rich nature found by the latter authors conformed with the results by Simpson et al. (2017), the evolutionary diagrams of both studies are very dissimilar and could not be explained by one simple isochrone fit. In particular, it was noted that \u201cthe Simpson et al. (2017) stars do not define a red giant branch in the theoretical plane, suggesting that their parameters are not correct\u201d (Fig. 1 of Mucciarelli et al. 2017). Such an inconsistency clearly emphasizes that a clear-cut chemical abundance scale is inevitable for fully characterising Gaia 1, and to further allow for tailored age determinations, even more so in the light of the seemingly well-determined orbital characteristics, Thus, this work focuses on a detailed chemical abundance analysis of four red giant members of Gaia 1, based on high-resolution spectroscopy, which we complement with an investigation of the orbital properties of this transition object. Combined with the red clump sample of Mucciarelli et al. (2017) and reaching down to the subgiant level (Simpson et al. 2017), stars in different evolutionary states in Gaia 1 are progressively being sampled. ","Citation Text":["Mucciarelli et al. (2017)"],"Functions Text":["Subsequently,","measured chemical abundances of six stars in Gaia 1, suggesting an equally high metallicity, but based on their abundance study, the suggestion of an extragalactic origin was revoked.","While a more metal-rich nature found by the latter authors conformed with the results by Simpson et al. (2017), the evolutionary diagrams of both studies are very dissimilar and could not be explained by one simple isochrone fit."],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1088,1113]],"Functions Start End":[[1074,1087],[1114,1297],[1298,1527]]}
{"Identifier":"2018AandA...615A.148D__Sung_et_al._2013_Instance_1","Paragraph":"We study here the Sco OB1 association (Figs. 1 and 2), using this and other techniques. The general properties of this large OB association, which spans almost 5\u00b0 on the sky, and is surrounded by a ring-shaped HII region called Gum 55, are reviewed by Reipurth (2008). Its central cluster NGC 6231 contains several tens of OB stars, which have been extensively studied. On the other hand, many fewer studies, all recent, were devoted to the full mass spectrum, using optical photometry (Sung et al. 1998, 2013) and X-rays (Sana et al. 2006, 2007; Damiani et al. 2016; Kuhn et al. 2017a,b). The currently accepted distance of NGC 6231 is approximately 1580 pc, and its age is between 2and 8 Myr, with a significant intrinsic spread (Sung et al. 2013; Damiani et al. 2016). No ongoing star formation is known to occur therein, however. Approximately one degree North of the cluster, the loose cluster Trumpler 24 (Tr 24) also belongs to the association. There is little literature on this cluster (Seggewiss 1968; Heske & Wendker 1984, 1985; Fu et al. 2003, 2005) which unlike NGC 6231 lacks a well-defined center and covers about one square degree on the sky. Its age is  10 Myr according toHeske & Wendker (1984, 1985), who find several PMS stars, and its distance is 1570\u20131630 pc according to Seggewiss (1968). Other studies of the entire Sco OB1 association include MacConnell & Perry (1969 \u2013 H\u03b1-emission stars), Schild et al. (1969 \u2013 spectroscopy), Crawford et al. (1971 \u2013 photometry), Laval (Laval 1972a,b \u2013 gas and star kinematics, respectively), van Genderen et al. (1984 \u2013 Walraven photometry), and Perry et al.(1991 \u2013 photometry). At the northern extreme of Sco OB1, the partially obscured HII region G345.45+1.50 and its less obscured neighbor IC4628 were studied by Laval (1972a), Caswell & Haynes (1987), L\u00f3pez et al. (2011), and L\u00f3pez-Calder\u00f3n et al. (2016). They contain massive young stellar objects (YSOs; Mottram et al. 2007), maser sources (Avison et al. 2016), and the IRAS source 16562-3959 with its radio jet (Guzm\u00e1n et al. 2010), outflow (Guzm\u00e1n et al. 2011), and ionized wind (Guzm\u00e1n et al. 2014), and are therefore extremely young (1 Myr or less). The distance of G345.45+1.50 was estimated as 1.9 kpc by Caswell & Haynes (1987), and 1.7 kpc by L\u00f3pez et al. (2011), in fair agreement with distances of Sco OB1 stars. In Fig. 1 of Reipurth (2008) a strip of blue stars is visible, connecting NGC 6231 to the region of IC4628.","Citation Text":["Sung et al.","2013"],"Functions Text":["On the other hand, many fewer studies, all recent, were devoted to the full mass spectrum, using optical photometry"],"Functions Label":["Background"],"Citation Start End":[[487,498],[505,509]],"Functions Start End":[[370,485]]}
{"Identifier":"2016ApJ...832..195N__Jin_et_al._2012_Instance_1","Paragraph":"We ignore the density stratification effect in Case I, II, and IIa, because the width of the horizontal current sheet in our simulations is much shorter than the length. The simulation domain extends from x = 0 to x = L0 in the x-direction, and from \n\n\n\n\ny\n=\n\u2212\n0.5\n\n\nL\n\n\n0\n\n\n\n\n to \n\n\n\n\ny\n=\n0.5\n\n\nL\n\n\n0\n\n\n\n\n in the y-direction, in the three cases, with \n\n\n\n\n\n\nL\n\n\n0\n\n\n=\n\n\n10\n\n\n6\n\n\n\n\n m. Outflow boundary conditions are used in the x-direction and inflow boundary conditions in the y-direction. For the inflow boundary conditions, the fluid is allowed to flow into the domain but not to flow out; the gradient of the plasma density vanishes; the total energy is set such that the gradient in the thermal energy density vanishes; a vanishing gradient of parallel components plus divergence-free extrapolation of the magnetic field. For the outflow boundary conditions, the fluid is allowed to flow out of the domain but not to flow in, and the other variables are set by using the same method as the inflow boundary conditions. The horizontal force-free Harris current sheet is used as the initial equilibrium configuration of magnetic fields in Case I,\n13\n\n\n\n\n\n\nB\n\n\nx\n0\n\n\n=\n\u2212\n\n\nb\n\n\n0\n\n\ntanh\n[\ny\n\n\/\n\n(\n0.05\n\n\nL\n\n\n0\n\n\n)\n]\n\n\n\n\n14\n\n\n\n\n\n\nB\n\n\ny\n0\n\n\n=\n0\n\n\n\n\n15\n\n\n\n\n\n\nB\n\n\nz\n0\n\n\n=\n\n\nb\n\n\n0\n\n\n\n\/\n\ncosh\n[\ny\n\n\/\n\n(\n0.05\n\n\nL\n\n\n0\n\n\n)\n]\n.\n\n\nThe magnetic fields in the low solar atmosphere could be very strong (Jin et al. 2009, 2012; Khomenko et al. 2014; Peter et al. 2014; Vissers et al. 2015) and the magnetic field can exceed 0.15 T in both the intranetwork and the network quiet region (e.g., Orozco Su\u00e1rez et al. 2007; Mart\u00ednez Gonz\u00e1lez et al. 2008; Jin et al. 2009, 2012). In the work by Jin et al. 2012, the maximum of the field strength was found to be 0.15 T. The magnetic field could be even stronger in the active region near the sunspot. Therefore, we set b0 = 0.05 T in Case I and Case II, and b0 = 0.15 T in Case IIa. Due to the force-freeness and neglect of gravity, the initial equilibrium thermal pressure is uniform. The initial temperature and plasma density are set as T0 = 4200 K and \u03c10 = 1.66057 \u00d7 10\u22126 kg m\u22123 in Case I, and T0 = 4800 K and \u03c10 = 3.32114 \u00d7 10\u22125 kg m\u22123 in Case II and Case IIa. Therefore, the initial plasma \u03b2 is calculated as \u03b2 \u2243 0.0583 in Case I, \u03b2 \u2243 1.332 in Case II, and \u03b2 \u2243 0.148 in Case IIa. The initial ionization degree is assumed as Yi = 10\u22123 in Case I, and Yi = 1. 2 \u00d7 10\u22124 in Case II and IIa. The magnetic diffusion in this work matches the form computed from the solar atmosphere model in Khomenko & Collados (2012), and we set \n\n\n\n\n\u03b7\n=\n\n[\n\n5\n\u00d7\n\n\n10\n\n\n4\n\n\n\n\n(\n4200\n\n\/\n\nT\n)\n\n\n1.5\n\n\n+\n1.76\n\u00d7\n\n\n10\n\n\n\u2212\n3\n\n\n\n\nT\n\n\n0.5\n\n\n\n\nY\n\n\ni\n\n\n\u2212\n1\n\n\n\n]\n\n\n\n m2 s\u22121 in Case I, and \n\n\n\n\n\u03b7\n=\n[\n5\n\u00d7\n\n\n10\n\n\n4\n\n\n\n\n(\n4800\n\n\/\n\nT\n)\n\n\n1.5\n\n\n\n+\n1.76\n\n\u00d7\n\n\n\n10\n\n\n\u2212\n3\n\n\n\n\nT\n\n\n0.5\n\n\n\n\nY\n\n\ni\n\n\n\u2212\n1\n\n\n]\n\n\n m2 s\u22121 in Case II and IIa. The first part \u223c T\u22121.5 is contributed by collisions between ions and electrons, the second part \n\n\n\n\n\u223c\n\n\nT\n\n\n0.5\n\n\n\n\nY\n\n\ni\n\n\n\u2212\n1\n\n\n\n\n is contributed by collisions between electrons and neutral particles. Small perturbations for both magnetic fields and velocities at t = 0 make the current sheet to evolve and secondary instabilities start to appear later in the three cases. The forms of perturbations are listed below:\n16\n\n\n\n\n\n\nb\n\n\nx\n1\n\n\n=\n\u2212\npert\n\u00b7\n\n\nb\n\n\n0\n\n\n\u00b7\nsin\n\n\n2\n\u03c0\n\n\n\ny\n+\n0.5\n\n\n\nL\n\n\n0\n\n\n\n\n\n\nL\n\n\n0\n\n\n\n\n\n\n\n\u00b7\ncos\n\n\n2\n\u03c0\n\n\n\nx\n+\n0.5\n\n\nL\n\n\n0\n\n\n\n\n\n\nL\n\n\n0\n\n\n\n\n\n\n\n\n\n\n\n17\n\n\n\n\n\n\nb\n\n\ny\n1\n\n\n=\npert\n\u00b7\n\n\nb\n\n\n0\n\n\n\u00b7\ncos\n\n\n2\n\u03c0\n\n\n\ny\n+\n0.5\n\n\nL\n\n\n0\n\n\n\n\n\n\nL\n\n\n0\n\n\n\n\n\n\n\n\u00b7\nsin\n\n\n2\n\u03c0\n\n\n\nx\n+\n0.5\n\n\nL\n\n\n0\n\n\n\n\n\n\nL\n\n\n0\n\n\n\n\n\n\n\n\n\n\n\n18\n\n\n\n\n\n\nv\n\n\ny\n1\n\n\n=\n\u2212\npert\n\u00b7\n\n\nv\n\n\nA\n0\n\n\n\u00b7\nsin\n\n\n\u03c0\n\n\n\ny\n\n\n\n\nL\n\n\n0\n\n\n\n\n\n\n\n\u00b7\n\n\n\n\n\nrandom\n\n\nn\n\n\n\n\nMax\n\n(\n\n\u2223\n\n\nrandom\n\n\nn\n\n\n\u2223\n\n)\n\n\n\n\n,\n\n\nwhere pert = 0.08, vA0 is the initial Alfv\u00e9n velocity, randomn is the random noise function in our code, and \n\n\n\n\nMax\n(\n\u2223\n\n\nrandom\n\n\nn\n\n\n\u2223\n)\n\n\n is the maximum of the absolute value of the random noise function. This random noise function makes the initial perturbations for the velocity in the y-direction to be asymmetric, and such an asymmetry makes the current sheet gradually become more tilted, especially after secondary islands appear. The reconnection process is not really symmetrical in nature (Murphy et al. 2012), this is one of the reasons that we use such a noise function. Another reason is that the asymmetric noise function makes the secondary instabilities develop faster. Figure 1(a) shows the distributions of the current density and magnetic fields at t = 0 in case I.","Citation Text":["Jin et al.","2012"],"Functions Text":["The magnetic fields in the low solar atmosphere could be very strong"],"Functions Label":["Uses"],"Citation Start End":[[1392,1402],[1409,1413]],"Functions Start End":[[1322,1390]]}
{"Identifier":"2015AandA...580A.135D__Cormier_et_al._2015_Instance_1","Paragraph":"How does the propagation of radiation and the ISM composition affect ISM observables in low-metallicity galaxies? Addressing this question is important to understand the evolution of low-metallicity galaxies, which undergo more bursty star formation than normal galaxies. Nearby star-forming dwarf galaxies present distinct observational signatures compared to well-studied disk galaxies. Dwarfs are usually metal poor, H\u2009i rich, and molecule poor as a result of large-scale photodissociation (e.g., Kunth & \u00d6stlin 2000; Hunter et al. 2012; Schruba et al. 2012). Mid-IR (MIR) and far-IR (FIR) observations have revealed bright atomic lines from H\u2009ii regions ([S\u2009iii], [Ne\u2009iii], [Ne\u2009ii], [O\u2009iii], etc.) and PDRs ([C\u2009ii], [O\u2009i]) (e.g., Hunter et al. 2001; Madden et al. 2006; Wu et al. 2008; Hunt et al. 2010; Cormier et al. 2015). Their spectral energy distributions (SEDs) are also different from spiral and elliptical galaxies and indicative of altered dust properties, with a relatively low abundance of polycyclic aromatic hydrocarbons (PAHs) and perhaps a different dust composition (e.g., Madden et al. 2006; Galliano et al. 2008; R\u00e9my-Ruyer et al. 2013). It is still unknown, however, whether these differences between dwarf and disk galaxies are the direct result of recent star formation activity shaping the ISM or instead a consequence of the low-metallicity ISM that is independent of star formation activity. To answer this, one needs to observe tracers of the interplay between the ISM and various stages of star formation activity. While there are now a number of important studies available on PDR properties modeling FIR lines on large scales in various extragalactic environments (e.g., Kaufman et al. 2006; Vasta et al. 2010; Graci\u00e1-Carpio et al. 2011; Cormier et al. 2012; Parkin et al. 2013) or in our Galaxy under solar-metallicity conditions (e.g., Cubick et al. 2008; Bernard-Salas et al. 2012, 2015), only a few studies are published on individual extragalactic regions (Mookerjea et al. 2011; Lebouteiller et al. 2012). Of particular interest are dwarf galaxies, where the effect due to radiative feedback is expected to be most significant. The goal of this paper is to investigate how the low-metallicity ISM reacts under the effects of star formation in regions that have undergone different histories. The nearby low-metallicity galaxy NGC\u20094214 provides an excellent environment to perform this experiment because it has well-separated star-forming centers, one hosting a super star cluster, which allows us to study the effects of extreme star-forming conditions on the surrounding ISM. ","Citation Text":["Cormier et al. 2015"],"Functions Text":["Mid-IR (MIR) and far-IR (FIR) observations have revealed bright atomic lines from H\u2009ii regions ([S\u2009iii], [Ne\u2009iii], [Ne\u2009ii], [O\u2009iii], etc.) and PDRs ([C\u2009ii], [O\u2009i]) (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[808,827]],"Functions Start End":[[563,733]]}
{"Identifier":"2021MNRAS.501.2112S__Rubin_et_al._2012_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":["Rubin et al. 2012"],"Functions Text":["It is important to do so because different processes, e.g.","pristine accretion","are surely contributing to the same system"],"Functions Label":["Motivation","Motivation","Motivation"],"Citation Start End":[[1139,1156]],"Functions Start End":[[944,1002],[1099,1117],[1294,1336]]}
{"Identifier":"2017ApJ...838...32Y__Somov_et_al._1981_Instance_1","Paragraph":"The question whether electrons are able to penetrate into deep and dense layers of the solar atmosphere is currently under scrutiny. Calculations by Emslie (1978) showed that for electrons and\/or protons to reach the \n\n\n\n\n\n level, their energy has to be of the order of a few MeV, while electron energy in the strongest flares only reaches several hundreds of keV. Considering the continuity equation for electron precipitation, Syrovatskii & Shmeleva (1972) and Dobranskis & Zharkova (2015) calculated that beam electrons with energies above 100 keV can reach the lower chromosphere with a column depth of 2 \u00d7 1021 cm\u22122, while those with energies of the order of 200 keV are even capable of penetrating to the photosphere (1022 cm\u22122; see also Zharkova & Gordovskyy 2006). Works by other authors seem to support the idea that power-law beam electrons may precipitate throughout the entire flaring atmosphere down to the photosphere (Brown 1971; Syrovatskii & Shmeleva 1972), causing its heating via inelastic collisions with the ambient plasma via a hydrodynamic response (Somov et al. 1981; Nagai & Emslie 1984; Fisher et al. 1985; Zharkova & Zharkov 2007). These power-law electrons also cause nonthermal excitation and ionization of hydrogen and other elements, combined with radiative transfer, leading to occurrence of emission in the lines and continua during flaring events (Aboudarham & Henoux 1986; Zharkova & Kobylinskii 1993). Observations indicate that CE may be a combination of radiation that originates (i) in an optically thin chromospheric layer owing to heating caused by beam electrons (thermal model), with subsequent recombination producing enhanced Balmer and Paschen continuum, and (ii) in the photosphere and TMR, where the plasmas are excited and ionized directly by collisions with high-energy electrons, resulting in radiative transfer processes in spectral lines and continua, combined with backwarming radiation supplied by the heated chromosphere and corona (e.g., Hudson 1972; Metcalf et al. 1990; Babin et al. 2016; Kleint et al. 2016). Indeed, since the early 1980s the enhanced Balmer continuum has been observed close to the Balmer cutoff at 3646 \u212b (Neidig 1989; Kerr & Fletcher 2014), and more recently a strong increase of Balmer continuum has been detected over a wide spectral range using IRIS data (Heinzel & Kleint 2014; Kleint et al. 2016).","Citation Text":["Somov et al. 1981"],"Functions Text":["Works by other authors seem to support the idea that power-law beam electrons may precipitate throughout the entire flaring atmosphere down to the photosphere",", causing its heating via inelastic collisions with the ambient plasma via a hydrodynamic response"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1073,1090]],"Functions Start End":[[773,931],[973,1071]]}
{"Identifier":"2015AandA...584A.103S__Potekhin_et_al._2013_Instance_4","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":["Potekhin et al. 2013"],"Functions Text":["and a comparison with the other EoSs of the BSk family","shall be left for future study."],"Functions Label":["Future Work","Future Work"],"Citation Start End":[[2025,2045]],"Functions Start End":[[1907,1961],[2088,2119]]}
{"Identifier":"2021MNRAS.500.3527B__Bernet_et_al._2008_Instance_1","Paragraph":"The origin of kiloparsec scale magnetic fields observed in the nearby galaxies through the polarized radio synchrotron emission (e.g. Fletcher 2010; Beck 2012, etc.) is attributed to the large-scale dynamo operating in the interstellar medium (ISM). This is driven mainly via the helical turbulent motions in the interstellar medium, coupled with the differential shear and vertical density stratification. This mechanism, along with some phenomenological approximations about the properties of background turbulence, in principle explains the growth of magnetic fields from small initial strengths to large-scale equipartition strengths against the diffusive losses (Beck et al. 1996; Shukurov 2005; Beck & Wielebinski 2013), and the characteristic times it takes for the field to reach the equipartition strength turn out to be of the order of $\\sim \\, {\\rm Gyr}\\,$. This is perhaps a much too slow to account for the strong equipartition strength magnetic fields observed in the high-redshift galaxies with z > 1 (e.g. Bernet et al. 2008) or even for that in the slowly rotating nearby galaxies. This discrepancy leads one to invoke some additional mechanism such as cosmic rays (CRs) boosting the typical dynamo action. The idea of CR driven dynamo was initially discussed by Parker (1992), this predicted the possibility of enhanced dynamo action by the virtue of additional CR buoyant instability, that inflates the magnetic field structures (see also Brandenburg 2018). Based on the conventional dynamo formulation Parker further suggested a simple model for the flux loss through the gaseous disc due to buoyancy by substituting the transport terms B\u03d5\/td. These terms are supposed to encapsulate the non advective flux transport associated with the buoyant instability, and leads to the fast dynamo action in characteristic field mixing times. Hanasz & Lesch (2000) indirectly verified such a dynamo action via the numerical simulations of rising magnetic flux tubes and found e-folding times of mean field of the order of $100\\, {\\rm Myr}\\,$. Supplementing this Hanasz, W\u00f3lta\u0144ski & Kowalik (2009), Siejkowski et al. (2010), Kulpa-Dybe\u0142 et al. (2015), Girichidis et al. (2016), etc. also demonstrated the fast amplification of regular magnetic fields via the direct magnetohydrodynamic (MHD) simulation of global galactic ISM including CR driven turbulence, along with the differential shear (but excluding the viscous term). To complement this, we aim here to extend our previous analysis of dynamo mechanism in SN driven ISM turbulence (Bendre, Gressel & Elstner 2015) by including the CR component and investigate the influence of magnetic field dependent propagation of CR on the dynamo. Here we focus on estimating the dynamo coefficients from the direct MHD simulations and effect CR component has on them by comparing with our previous analysis without the CR.","Citation Text":["Bernet et al. 2008"],"Functions Text":["This is perhaps a much too slow to account for the strong equipartition strength magnetic fields observed in the high-redshift galaxies with z > 1 (e.g.","or even for that in the slowly rotating nearby galaxies. This discrepancy leads one to invoke some additional mechanism such as cosmic rays (CRs) boosting the typical dynamo action."],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[1022,1040]],"Functions Start End":[[869,1021],[1042,1223]]}
{"Identifier":"2016AandA...592A..19C__Maraston_et_al._(2009)_Instance_1","Paragraph":"Since the star-formation histories of galaxies (ETGs included, e.g. De Lucia et al. 2006; Maraston et al. 2009) can be stochastic and include multiple bursts, we also verify the full-spectrum fitting capabilities to retrieve more complex SFHs. In particular, we take an 11 Gyr old composite stellar population with an exponentially delayed SF (\u03c4 = 0.3 Gyr) as the main SF episode (this age is compatible with the age of the Universe at z ~ 0.15, which is the median redshift of our sample, see Sect. 3). We then define more complex SFHs by combining this single CSP with a burst of SF at different ages (5, 6, 7 Gyr) and with different mass contributions (3, 5, 10 %). In all cases, we consider a solar metallicity for the main SF episode and, according to the results of Maraston et al. (2009), a subsolar metallicity (Z = 0.004) for the later one. We do not mask any spectral feature of the input spectra, we assume AV  =  0.1 mag for the two components and apply a velocity dispersion of 200 km s-1. We show the results for a S\/N of 80, which matches the typical S\/N of the SDSS median stacked spectra analyzed in the following (see Sect. 3). Useful information can be derived from the comparison between the output SFH obtained from these input simulated spectra and the one provided when the single CSP alone is taken as input SFH. Fig. 5 shows that the single CSP alone is well recovered by the full spectrum fitting. In particular, ~80% of the stellar mass is retrieved within ~1 Gyr from the SFH peak. When a burst is added to this major episode of SF, the full-spectrum fitting is able to recognize the presence of a more complex SFH, as indicated by the tail appearing at smaller ages, and the total mass percentage of the later burst is retrieved within 1 Gyr from the expected age. However, we note that the main episode of SF is spread on a time interval longer than expected, and 50% of the stellar mass is retrieved around ~1 Gyr from the SFR peak. We also find that, in this case, the mean properties of the global stellar population are well retrieved, with a percentage accuracy larger than 10% starting from S\/N ~ 15 for age, ~7 for metallicity, ~20 for AV, ~8 for \u03c3 and that the metallicities of the two SF episodes are separately recovered. These S\/Ns are well below those typical of the stacked spectra analyzed in the following sections. ","Citation Text":["Maraston et al. 2009"],"Functions Text":["Since the star-formation histories of galaxies (ETGs included, e.g.","can be stochastic and include multiple bursts, we also verify the full-spectrum fitting capabilities to retrieve more complex SFHs."],"Functions Label":["Uses","Uses"],"Citation Start End":[[90,110]],"Functions Start End":[[0,67],[112,243]]}
{"Identifier":"2022MNRAS.509.6091H__Tremmel_et_al._2017_Instance_1","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":["Tremmel et al. 2017"],"Functions Text":["Current cosmological hydrodynamic simulations of galaxy formation employ a variety of subgrid models (e.g.","that artificially launch galactic winds, but the results are sensitive to numerical resolution and the exact subgrid model employed"],"Functions Label":["Background","Background"],"Citation Start End":[[669,688]],"Functions Start End":[[389,495],[751,882]]}
{"Identifier":"2022MNRAS.512.4893Z__Toft_et_al._2014_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":["Toft et al. 2014"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[958,974]],"Functions Start End":[[694,899]]}
{"Identifier":"2018ApJ...863..162M__Liu_et_al._2013_Instance_3","Paragraph":"NLFFF extrapolation provides the reconstructed coronal magnetic field for AR 11158 from 2011 February 13 \u2212 2011 February 15 (Figures 1(d)\u2013(f)). The field lines (yellow lines) within the core of the AR have arcade-like structure with a relatively strong twist mainly near the PIL. These figures show that the magnetic field evolved during this period. Although we did not quantitatively compare the field lines with the observation, in general, the reconstructed coronal field morphologies match with the observations in Figures 1(a)\u2013(c). The general morphologies and the locations of the high-twist fields are also in agreement with many previous studies (Jing et al. 2012; Sun et al. 2012; Dalmasse et al. 2013; Inoue et al. 2013, 2014a; Liu et al. 2013; Wang et al. 2013; Aschwanden et al. 2014; Malanushenko et al. 2014; Zhao et al. 2014). Unlike Zhao et al. (2014) who could identify the twisted flux rope from the topology of the reconstructed coronal field, we could not find an obvious topological signature of a flux rope existing in our NLFFF during our analysis time window. It might be due to the fact that there was little magnetic flux with twist higher than one turn in our NLFFF and it is difficult to topologically define it as a flux rope. However, our result is consistent with other NLFFF results (Jing et al. 2012; Sun et al. 2012; Liu et al. 2013; Wang et al. 2013; Inoue et al. 2014a; Malanushenko et al. 2014). The high-twist region in our result is also in agreement with the region with high helicity flux (Dalmasse et al. 2013) and the location of the flare ribbons (Bamba et al. 2013; Liu et al. 2013), as well as the high current density region (Janvier et al. 2014). Figures 1(g)\u2013(i) show the evolution of the twist distribution map, with the magnetic twist of the field lines plotted at the footpoints of field lines according to a color scale. This shows that the high-twist (strongly right-handed twist corresponding to Tw > 0.5) areas are concentrated in only a limited part of the AR. The high-twist area grew and became even more twisted just before the X2.2 flare (Figure 1(i)). Most parts of the AR have twist values less than 0.25, but near the PIL the twist can reach more than 0.5, even up to about a full turn. This is consistent with the results of Sun et al. (2012) and Inoue et al. (2014a). A high-twist (strong negative\/left-handed twist) area also developed in the eastern part of the AR, which did not exist initially on February 13. Both of these high-twist areas produced several flares. However, here we focus on the flares that resulted from the high-twist core region near the center of the AR, where the M6.6 and X2.2 flares occurred.","Citation Text":["Liu et al. 2013"],"Functions Text":["The high-twist region in our result is also in agreement with the","and the location of the flare ribbons"],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[1612,1627]],"Functions Start End":[[1434,1499],[1554,1591]]}
{"Identifier":"2021ApJ...915...86A__Owen_&_Sathyaprakash_1999_Instance_1","Paragraph":"This analysis searches for a GW signal compatible with the inspiral of a BNS or NSBH binary\u2014collectively NS binaries\u2014within 6 s of data associated with an observed short GRB. This stretch of data is the on-source window and runs from \u22125 s to +1 s around the start of the GRB emission (i.e., the GRB trigger time). The surrounding \u223c30\u201390 minutes of data are split into 6 s off-source trials which are also analyzed in order to build a background. Around 30 minutes allows the modeled search to accurately estimate the power spectral density of the available instruments and ensures that it can assess at sub-percent level accuracy the significance of any candidate events found in the on-source window. All the data are processed using PyGRB (Harry & Fairhurst 2011; Williamson et al. 2014), a coherent matched filtering pipeline that is part of the general open-source software PyCBC (Nitz et al. 2020) and has core elements in the LALSuite software library (LIGO Scientific Collaboration 2018). We scan each trial of data and the on-source window in the 30\u20131000 Hz frequency band using a predefined bank of waveform templates (Owen & Sathyaprakash 1999) created with a hybrid geometric\u2013stochastic method (Capano et al. 2016; Dal Canton & Harry 2017) and using a phenomenological inspiral-merger-ringdown waveform model for non-precessing point-particle binaries (IMRPhenomD; Husa et al. 2016; Khan et al. 2016).210\n\n210\nAll waveforms mentioned in this section are generated with the LALSimulation package that is part of the LALSuite software library (LIGO Scientific Collaboration 2018).\n The waveform template bank includes waveforms corresponding to a range of masses ([1.0, 2.8]M for NSs, [1.0, 25.0]M for BHs) and dimensionless spin magnitudes ([0, 0.05] for NSs, [0, 0.998] for BHs) for aligned-spin, zero-eccentricity BNS or NSBH systems that may produce an electromagnetic counterpart via the tidal disruption of the NS (Pannarale & Ohme 2014). Aside from the updated sensitivity of our detectors, the only difference with respect to the second LIGO\u2013Virgo observing run (Abbott et al. 2019b) is that the generation of the bank has been updated to apply more accurate physics to determine whether an NSBH system could produce an accretion disk from this disruption (Foucart et al. 2018). We only search for circularly polarized GWs, which may be emitted by binaries with inclinations of 0\u00b0 or 180\u00b0: such systems have GW amplitudes that are consistent (Williamson et al. 2014) with those of binary progenitors with inclination angles over the full range of viewing angles that we expect for typical brightness GRBs (\u227230\u00b0; Fong et al. 2015), such as those in our sample.","Citation Text":["Owen & Sathyaprakash 1999"],"Functions Text":["We scan each trial of data and the on-source window in the 30\u20131000 Hz frequency band using a predefined bank of waveform templates"],"Functions Label":["Uses"],"Citation Start End":[[1128,1153]],"Functions Start End":[[996,1126]]}
{"Identifier":"2016ApJ...826L..14W__Goudfrooij_et_al._2015_Instance_1","Paragraph":"Figure 1 shows the NGC 419 CMD after field-star decontamination. The cluster has an eMSTO at 20 \u2264 V \u2264 22 mag. We adopted the Padova stellar evolution models (PARSEC CMD 2.7, v. 1.2S; Bressan et al. 2012),8\n\n8\n\nhttp:\/\/stev.oapd.inaf.it\/cgi-bin\/cmd_2.7\n\n for Z = 0.004 (Glatt et al. 2008; Girardi et al. 2009), and fitted the eMSTO extremes with isochrones of log(t yr\u22121) = 9.12 (1.32 Gyr; blue line) and log(t yr\u22121) = 9.31 (2.02 Gyr; red line), adopting a distance modulus (m \u2212 M)0 = 18.90 mag and a visual extinction AV = 0.181 mag (Rubele et al. 2010). The maximum possible age spread implied by the extent of the eMSTO is \u223c700 Myr, similar to the results of Rubele et al. (2010; \u223c700 Myr) and Girardi et al. (2013; \u223c670 Myr). However, the cluster exhibits an SGB that is broad on the blue side and that becomes significantly narrower on the red side. Previous studies of intermediate-age star clusters with eMSTOs usually showed tight SGBs throughout (Mackey & Broby Nielsen 2007; Li et al. 2014b; Bastian & Niederhofer 2015; but see Goudfrooij et al. 2015). This feature is, however, already apparent in the NGC 419 data before field-star decontamination and is, hence, not caused by our data reduction. In addition, the significance level of our field-star decontamination is very high along the SGB, while the good agreement between our iraf\/daophot and dolphot stellar catalogs further confirms the reality of the SGB morphology. Examination of the SGB stars\u2019 spatial distribution reveals that blending is unlikely responsible for the observed narrowing either. The NGC 419 CMD of Goudfrooij et al. (2014) shows a similar trend, although that of Glatt et al. (2008) resembles an SSP more closely. The latter CMD is, however, composed of HST ACS\/High Resolution Channel observations covering a small stellar sample located in the cluster\u2019s central region only. Our data are consistent with the Glatt et al. (2008) CMD for stars drawn from inside the cluster\u2019s core radius. We thus conclude that the observed narrowing of the NGC 419 SGB is real and not caused by artifacts associated with our data reduction (see the Appendix).","Citation Text":["Goudfrooij et al. 2015"],"Functions Text":["Previous studies of intermediate-age star clusters with eMSTOs usually showed tight SGBs throughout","but see"],"Functions Label":["Differences","Differences"],"Citation Start End":[[1036,1058]],"Functions Start End":[[853,952],[1028,1035]]}
{"Identifier":"2022ApJ...926...21B__K\u00e4pyl\u00e4_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":["K\u00e4pyl\u00e4 et al. 2011"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[911,929]],"Functions Start End":[[589,787]]}
{"Identifier":"2022MNRAS.513.5377F__Blum_et_al._2017_Instance_1","Paragraph":"At each heliocentric distance rh, the activity model (Fulle et al. 2020b) is defined by five analytical equations fixing (i) the gas pressure P(s) depending on the depth s from the nucleus surface (Fig. 1 for the CO2 case), (ii) the gas flux Q from the nucleus surface, (iii) the temperature gradient \u2207T at depths of a few cm, (iv) the heat conductivity \u03bbs at depths of a few cm below the nucleus surface, and (v) the temperature Ts of the nucleus surface\n(3)$$\\begin{eqnarray*}\r\nP(s) = P_0 ~f(s) ~\\exp \\left[{- {T_0 \\over {T_s - s ~\\nabla T}}}\\right]\r\n\\end{eqnarray*}$$(4)$$\\begin{eqnarray*}\r\nQ = {{14 ~r ~P(R)} \\over {3 ~R}} \\sqrt{{2 ~m} \\over {\\pi k_B ~(T_s - R ~\\nabla T)}}\r\n\\end{eqnarray*}$$(5)$$\\begin{eqnarray*}\r\n\\nabla T = {\\sqrt{\\Lambda ~Q ~\/ ~\\sigma _B} \\over {8 ~(T_s - R ~\\nabla T) ~R}}\r\n\\end{eqnarray*}$$(6)$$\\begin{eqnarray*}\r\n\\lambda _s = {32 \\over 3} ~(T_s - R ~\\nabla T)^3 ~\\sigma _B ~R\r\n\\end{eqnarray*}$$(7)$$\\begin{eqnarray*}\r\n(1 - A) ~I_\\odot ~\\cos \\theta ~r_h^{-2} = \\epsilon \\sigma _B T_s^4 + \\lambda _s ~\\nabla T + \\Lambda ~Q\r\n,\r\n\\end{eqnarray*}$$where P0, T0, and \u039b values are listed in Table 1, s is the depth from the nucleus surface, $f(s) = 1 - (1 - {s \\over R})^4$ for s \u2264 R, f(s) = 1 elsewhere, r \u2248 50 nm and R \u2248 5 mm are the radii of the grains of which cometary dust consists (Levasseur-Regourd et al. 2018; G\u00fcttler et al. 2019; Mannel et al. 2019) and of the pebbles of which cometary nuclei consist (Blum et al. 2017; Fulle et al. 2020b), m is the mass of the gas molecule, kB is the Boltzmann constant, \u03c3B is the Stefan\u2013Boltzmann constant, A is the nucleus Bond albedo (e.g. A = 1.2 per cent measured at 67P; Fornasier et al. 2015), I\u2299 is the solar flux at the heliocentric distance of Earth, \u03b8 is the solar zenithal angle, and \u03f5 \u2248 0.9 is the nucleus emissivity. Since the gas originates from the superficial pebbles and is assumed to share the temperature Ts \u2212 s \u2207T of refractories and ices, the thermal diffusion due to gas convection is negligible with respect to the sublimation sink \u039b Q. A nucleus is active if the gas pressure overcomes the tensile strength S bonding dust particles to the nucleus surface (Skorov & Blum 2012), thus defining the activity onset for each ice (Table 2), occurring (i) at rh = 85 au for carbon monoxide (Fulle et al. 2020a); (ii) at rh = 60 au for molecular oxygen; (iii) at rh = 52 au for methane; (iv) at rh = 18 au for ethane; (v) at rh = 13 au for carbon dioxide (dotted line in Fig. 1); and (vi) at rh = 3.8 au for water (Fulle et al. 2020b; Ciarniello et al. 2021). The value R \u2248 5 mm has been constrained by several data collected at comet 67P, by laboratory experiments of dust accretion in conditions expected to occur in the solar protoplanetary disc and by observations of other protoplanetary discs (Blum et al. 2017). Other R-values would not provide the best fit of the 67P water-loss time-evolution (Ciarniello et al. 2021).","Citation Text":["Blum et al. 2017"],"Functions Text":["where P0, T0, and \u039b values are listed in Table 1, s is the depth from the nucleus surface, $f(s) = 1 - (1 - {s \\over R})^4$ for s \u2264 R, f(s) = 1 elsewhere, r \u2248 50 nm and R \u2248 5 mm are the radii of the grains of which cometary dust consists","and of the pebbles of which cometary nuclei"],"Functions Label":["Uses","Uses"],"Citation Start End":[[1434,1450]],"Functions Start End":[[1070,1307],[1381,1424]]}
{"Identifier":"2016ApJ...826..168X__Bai_2014_Instance_2","Paragraph":"MRI is considered to be the most promising mechanism driving angular-momentum transport in protoplanetary disks (Balbus & Hawley 1991; Brandenburg et al. 1995; Hawley et al. 1995; Balbus et al. 1996; Balbus & Hawley 1998). However, protoplanetary disks are cold, dense, and, therefore, poorly ionized. The low level of ionization tends to decouple the disk gas from magnetic fields, which generates non-ideal MHD effects: Ohmic dissipation, ambipolar diffusion (AD), and the Hall effect (e.g., Armitage 2011; Turner et al. 2014). These effects quench MRI in different ways: Ohmic dissipation originates from collisions between electrons and neutrals, AD from collisions between ions and neutrals, and the Hall effect from drift between electrons and ions (Fleming et al. 2000; Sano & Stone 2002; Bai & Stone 2011). Ohmic dissipation operates in high-density regions with weak field, AD dominates in highly ionized and low-density regions, and the Hall effect lies in between (Fleming & Stone 2003; Bai & Stone 2011; Bai 2014). So far, the effect of Ohmic dissipation has been best studied. Investigations show the layered accretion in the inner disk, where the midplane region is \u201cdead\u201d due to low ionization while the surface layer is \u201cactive\u201d due to sufficient ionization (Gammie 1996; Jin 1996; Fleming et al. 2000; Fleming & Stone 2003; Turner et al. 2007; Ilgner & Nelson 2008; Oishi & Mac Low 2009; Okuzumi & Hirose 2011). Recent works that take into account both Ohmic dissipation and AD show that AD may render the surface layer and portions of the outer disk inactive (Bai & Stone 2011; Landry et al. 2013; Kalyaan et al. 2015). Bai & Stone (2013) find that MRI is completely suppressed in the inner disk and a strong magnetocentrifugal wind is launched. Three-dimensional simulations that include all three non-ideal MHD effects are also performed (Bai 2014, 2015; Lesur et al. 2014; Simon et al. 2015). In the inner disk, the influence of the Hall effect on midplane angular-momentum transport depends on the orientation of the vertical magnetic field with the disk rotation axis. When the field is aligned with the axis, the enhanced Maxwell stress promotes angular-momentum transport. When the field is anti-aligned with the axis, the midplane remains quiescent. In the outer disk, the Hall effect has little influence on the disk turbulence. Although the inclusion of AD and the Hall effect substantially changes the level of turbulence in the protoplanetary disks, the feature that the viscosity is low in the inner disk and high in the outer disk is still valid. In this study, we assume that gas giant planets form in situ via the core accretion scenario, which implies that their formation locations are always in the low-viscosity region. Since in this study we focus on the relation between photoevaporation and planet formation and gap opening by planets in the disk, we adopt Ohmic dissipation to represent the non-ideal MHD effects on the MRI. We consider that this simplification has little influence on our main calculation results.","Citation Text":["Bai 2014"],"Functions Text":["Three-dimensional simulations that include all three non-ideal MHD effects are also performed"],"Functions Label":["Background"],"Citation Start End":[[1859,1867]],"Functions Start End":[[1764,1857]]}
{"Identifier":"2021AandA...653A..36M__Goulding_&_Alexander_(2009)_Instance_1","Paragraph":"The SFG sample was constructed using the Great Observatories All-Sky LIRG Survey (GOALS sample, Armus et al. 2009), from which we extracted 158 galaxies, with data from Inami et al. (2013), who report the fine-structure lines at high resolution in the 10\u2005\u2212\u200536\u2006\u03bcm interval, and Stierwalt et al. (2014), who include the detections of the H2 molecular lines and the PAH features at low spectral resolution. For those galaxies in the GOALS sample that have a single IRAS counterpart, but more than one source detected in the emission lines, we have added together the line or feature fluxes of all components, to consistently associate the correct line or feature emission to the total IR luminosity computed from the IRAS fluxes. To also cover lower luminosity galaxies, as the GOALS sample only includes luminous IR galaxies (LIRGs) and ultra-luminous IR galaxies (ULIRGs), we included 38 galaxies from Bernard-Salas et al. (2009) and Goulding & Alexander (2009), to reach the total sample of 196 galaxies with IR line fluxes in the 5.5\u2005\u2212\u200535\u2006\u03bcm interval in which an AGN component is not detected. For the Bernard-Salas et al. (2009), Goulding & Alexander (2009), and the GOALS samples, we excluded all the composite starburst-AGN objects identified as those with a detection of [NeV] either at 14.3 or 24.3 \u03bcm. It is worth noting that the original samples from Goulding & Alexander (2009) and Bernard-Salas et al. (2009) have spectra solely covering the central region of the galaxies. To estimate the global SFR, we corrected the published line fluxes of the Spitzer spectra by multiplying them by the ratio of the continuum reported in the IRAS point source catalogue to the continuum measured on the Spitzer spectra extracted from the CASSIS database (Lebouteiller et al. 2015). We assumed here that the line emission scales (at first order) with the IR brightness distribution. In particular, we considered the continuum at 12\u2006\u03bcm for the [NeII]12.8\u2006\u03bcm and [NeIII]15.6\u2006\u03bcm lines, and the continuum at 25 \u03bcm for the [OIV]25.9\u2006\u03bcm, [FeII]26\u2006\u03bcm, [SIII]33.5\u2006\u03bcm, and [SiII]34.8\u2006\u03bcm lines. This correction was not needed for the AGN sample and the GOALS sample because of the greater average redshift of the galaxies in the 12MGS and GOALS samples. In particular, the 12MGS active galaxy sample has a mean redshift of 0.028 (Rush et al. 1993), while the GOALS sample has a mean redshift of 0.026. The galaxies presented by Bernard-Salas et al. (2009) have instead an average redshift of 0.0074, while the sample by Goulding & Alexander (2009) has an average redshift of 0.0044. For the other lines in the 10\u2005\u2212\u200536\u2006\u03bcm interval, Goulding & Alexander (2009) did not report a detection, and we used the data presented in Bernard-Salas et al. (2009) for a total of 15 objects. Both Bernard-Salas et al. (2009) and Goulding & Alexander (2009) reported data from the high-resolution Spitzer-IRS spectra. Data in the 50\u2005\u2212\u2005205\u2006\u03bcm interval were taken from D\u00edaz-Santos et al. (2017). For the GOALS sample, 20 starburst galaxies were taken from Fern\u00e1ndez-Ontiveros et al. (2016), and 23 objects were taken from the ISO-LWS observations of Negishi et al. (2001). As a result, we obtained a total sample of 193 objects. Lastly, the PAH features\u2019 fluxes were measured from the low-resolution Spitzer-IRS spectra by Brandl et al. (2006), including 12 objects from the sample of Bernard-Salas et al. (2009) and 179 objects from Stierwalt et al. (2014).","Citation Text":["Goulding & Alexander (2009)"],"Functions Text":["To also cover lower luminosity galaxies, as the GOALS sample only includes luminous IR galaxies (LIRGs) and ultra-luminous IR galaxies (ULIRGs), we included 38 galaxies from Bernard-Salas et al. (2009) and","to reach the total sample of 196 galaxies with IR line fluxes in the 5.5\u2005\u2212\u200535\u2006\u03bcm interval in which an AGN component is not detected."],"Functions Label":["Uses","Uses"],"Citation Start End":[[933,960]],"Functions Start End":[[727,932],[962,1094]]}
{"Identifier":"2020MNRAS.492..686L__Shiokawa_et_al._2015_Instance_1","Paragraph":"After the disruption phase, the star is tidally stretched into a very long thin stream and the evolution of the stream structure in the transverse and longitudinal directions are decoupled (Kochanek 1994). Thus, the system enters the free-fall phase where each stream segment follows its own geodesic like a test particle (Coughlin et al. 2016). Then, after passing the apocentres of the highly eccentric orbits, the bound debris falls back towards the BH at a rate given by the distribution of specific energy (Evans & Kochanek 1989; Phinney 1989). Due to relativistic apsidal precession, the bound debris, after passing the pericentre, collides violently with the still in-falling stream (see Fig. 1). It has been shown that shocks at the self-intersection point is the main cause of orbital energy dissipation and the subsequent formation of an accretion disc (Rees 1988; Kochanek 1994; Hayasaki, Stone & Loeb 2013; Guillochon, Manukian & Ramirez-Ruiz 2014; Shiokawa et al. 2015; Bonnerot et al. 2016). However, the aftermath of the self-intersection is an extremely complex problem, which depends on the interplay among magnetohydrodynamics, radiation, and general relativity in 3D. No numerical simulations to date have been able to provide a deterministic model for TDEs with realistic star-to-BH mass ratio and high eccentricity (see Stone et al. 2018a, for a review). Many simulations consider either an intermediate-mass BH (e.g. Guillochon et al. 2014, Evans, Laguna & Eracleous 2015; Shiokawa et al. 2015; Sa\u0327dowski et al. 2016) or the disruption of a low-eccentricity (initially bound) star (e.g. Bonnerot et al. 2016; Hayasaki, Stone & Loeb 2016). It is unclear how to extrapolate the simulation results to realistic configurations and provide an answer to the following questions: How long does it take for the bound gas to form a circular accretion disc (if at all)? How much radiative energy is released from the system? What fraction of the radiation is emitted in the optical, UV, or X-ray bands?","Citation Text":["Shiokawa et al. 2015"],"Functions Text":["It has been shown that shocks at the self-intersection point is the main cause of orbital energy dissipation and the subsequent formation of an accretion disc"],"Functions Label":["Background"],"Citation Start End":[[961,981]],"Functions Start End":[[704,862]]}
{"Identifier":"2018ApJ...854..141D__Vegetti_et_al._2014_Instance_1","Paragraph":"Our approach shares its transdimensional nature with that of Brewer et al. (2015), which implemented an independent transdimensional framework to infer substructure in strong lens systems. However, in addition to differences concerning substructure modeling and the sampling method used, our code and analysis were developed independently of theirs, and hence provide an important cross-check of the method and results. It has become customary in cosmological analyses to have at least two publicly available independent codes to check for numerical accuracy and find bugs. While both codes are public, we understand that this approach is not yet ubiquitous in the field of galaxy-scale strong lensing (for instance, it has been difficult thus far to reproduce\/check the results of Vegetti et al. 2014 and Hezaveh et al. 2016), but we hope that our work can help make substructure lensing more reproducible. Given the uniqueness of the constraints it can provide on low-mass substructures and the impact that these detections can have on dark matter physics, we believe it is extremely important to have multiple cross-checks to validate any results. On a more technical level, our approach differs from that of Brewer et al. (2015) in two major ways:\n\n1.\nSubhalo lens modeling. We represent subhalos with NFW profiles with a certain mass, scale, and cutoff radius, whereas Brewer et al. (2015) uses so-called blobs, i.e., deflection as a power law in the subhalo-centric radius. We think that choosing a physically motivated basis function set to represent the deflection field due to subhalos is essential for correctly marginalizing them out, i.e., propagating their (rather large) uncertainties to the parameter of interests such as the subhalo mass function.\n\n\n2.\nThe sampling method and the sampler used to sample from the posterior probability distribution of the (different) lens models given the data. Brewer et al. (2015) uses Diffusive Nested Sampling (DNS) as implemented in RJObject. We use a homebrew reversible-jump MCMC sampler with Metropolis-type within-model proposals.\n\n\n","Citation Text":["Vegetti et al. 2014"],"Functions Text":["While both codes are public, we understand that this approach is not yet ubiquitous in the field of galaxy-scale strong lensing (for instance, it has been difficult thus far to reproduce\/check the results of","and Hezaveh et al. 2016), but we hope that our work can help make substructure lensing more reproducible."],"Functions Label":["Differences","Differences"],"Citation Start End":[[782,801]],"Functions Start End":[[574,781],[802,907]]}
{"Identifier":"2016AandA...586A..80O__Fornasier_et_al._2015_Instance_1","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"],"Functions Text":["In the regions we investigated, the Hapi region displays the lowest spectral slopes 8\u201311%\/100 nm (see also","together with the isolated bright features (IBFs) and the clustered bright features in the Imhotep region."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[848,869]],"Functions Start End":[[741,847],[871,977]]}
{"Identifier":"2021ApJ...906...57S__Mishra-Sharma_et_al._2017_Instance_1","Paragraph":"Complementary to studies using the integrated emission and angular power spectrum of DM annihilation from a population of Galactic subhalos, in this paper we present a novel strategy using one-point photon statistics to search for the annihilation signature. Our technique takes advantage of the information in the entire population of sources, including both those that are resolved and those that are faint and unresolved. The concept of leveraging the one-point photon-count distribution to search for DM has previously been studied in Dodelson et al. (2009) and Feyereisen et al. (2015) in the context of emission from extragalactic sources and in Lee et al. (2009) and Koushiappas et al. (2010) with application to Galactic subhalos. We introduce a method to search for signatures of DM annihilation from a Galactic subhalo population using the non-Poissonian template fitting (NPTF) framework (Malyshev & Hogg 2011; Lee et al. 2015, 2016; Mishra-Sharma et al. 2017), which has previously been applied to characterize unresolved point sources in the inner Galaxy (Lee et al. 2016; Linden et al. 2016; Leane & Slatyer 2019, 2020a, 2020b; Chang et al. 2020; Buschmann et al. 2020) and at high latitudes (Zechlin et al. 2016; Lisanti et al. 2016; Zechlin et al. 2018). Using simulations, we show that the NPTF can constrain DM annihilation from a population of subhalos in the face of astrophysical background emission. We find that using photon statistics to look for collective emission from a subhalo population can be especially promising when a large number of individual subhalo candidates are identified in point-source catalogs. This establishes a method complementary to the established ones based on characterizing individual resolved point sources as subhalo candidates, as well as those based on using the measured 0-point (overall flux) and two-point (angular power spectrum) statistics to characterize a subhalo population. Moreover, note that our methodology is completely independent of assumptions about, e.g., the location of stellar overdensities, so we are less sensitive to certain uncertainties which can bias dwarf galaxy constraints such as those in modeling tracer populations. Thus, this framework provides an important comparison for the dwarf galaxy analyses as well.","Citation Text":["Mishra-Sharma et al. 2017"],"Functions Text":["We introduce a method to search for signatures of DM annihilation from a Galactic subhalo population using the non-Poissonian template fitting (NPTF) framework"],"Functions Label":["Uses"],"Citation Start End":[[945,970]],"Functions Start End":[[739,898]]}
{"Identifier":"2020AandA...635A.121M__Matter_et_al._(2016)_Instance_1","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"],"Functions Text":["Our model consists of four zones; an inner disk (zone 1, as in"],"Functions Label":["Uses"],"Citation Start End":[[63,81]],"Functions Start End":[[0,62]]}
{"Identifier":"2017AandA...601A.134M__Fung_&_Dong_(2015)_Instance_1","Paragraph":"Several predictions for planet(s) shaping the disk of SAO 206462 have been proposed (Muto et al. 2012; Garufi et al. 2013; Fung & Dong 2015; Bae et al. 2016; van der Marel et al. 2016a. Using linear equations from the spiral density wave theory, Muto et al. (2012) suggested two planets with separations beyond ~50 au by fitting independently the two spiral arms seen in Subaru\/HiCIAO data and with masses of ~0.5 MJ by using the amplitude of the spiral wave. Garufi et al. (2013) proposed that one planet of mass 5\u201313 MJ located inside the cavity at a separation of 17.5\u201320 au could be responsible for the different cavity sizes measured for the small and large dust grains. Fung & Dong (2015) presented scaling relations between the azimuthal separation of the primary and secondary arms and the planet-to-star mass ratio for a single companion on a circular orbit with a mass between Neptune mass and 16 MJ around a 1 M\u2299 star. They predicted with 30% accuracy that a single putative planet responsible for both spiral features of SAO 206462 would have a mass of ~6 MJ. Bae et al. (2016) presented dedicated hydrodynamical simulations of the SAO 206462 disk and proposed that both the bright scattered-light feature (Garufi et al. 2013) and the dust emission peak (P\u00e9rez et al. 2014) seen for the southwestern spiral arm result from the interaction of the spiral arm with a vortex, although a vortex alone can account for the S1 brightness peak. They suggested that a 10\u201315 MJ planet may orbit at 100\u2013120 au from the star. However, ALMA observations at two different frequencies seem to contradict a dust particle trapping scenario by a vortex (Pinilla et al. 2015). Stolker et al. (2016) performed new fitting of the spiral arms observed in SPHERE data and found a best-fit solution with two protoplanets located exterior to the spirals: r1\u2009~\u2009168 au, \u03b81\u2009~\u200952\u00b0 and r2\u2009~\u200999 au, \u03b82\u2009~\u2009355\u00b0. van der Marel et al. (2016a) proposed that the features seen in thermal emission in ALMA data and the scattered-light spiral arms are produced by a single massive giant planet located inside the cavity at a separation of ~30 au. Recently, Dong & Fung (2017) used the contrast of the spiral arms to predict a giant planet of ~5\u201310 MJ at ~100 au. ","Citation Text":["Fung & Dong 2015"],"Functions Text":["Several predictions for planet(s) shaping the disk of SAO 206462 have been proposed"],"Functions Label":["Background"],"Citation Start End":[[123,139]],"Functions Start End":[[0,83]]}
{"Identifier":"2021AandA...650A.205V__Jones_et_al._2021_Instance_2","Paragraph":"The question of the evolution of exoplanet systems after the main sequence of their host is generally addressed by studying exoplanets around subgiants, RGB stars, and normal HB (RC) stars (hereafter the \u2019classical\u2019 evolved stars). These classical evolved stars are typically very large stars, with radii ranging from ~ 5\u2212 10 R\u2299 to more than 1000 R\u2299. This is much larger than hot subdwarfs, which have radii in the range ~ 0.1\u22120.3 R\u2299 (Heber 2016). Their mass is typically higher than ~ 1.5 M\u2299, compared to~0.47 M\u2299 for hot subdwarfs. The transit and radial velocity (RV) methods are both challenging for these classical evolved stars because the transit depth is diluted and there are additional noise sources (Van Eylen et al. 2016). Another difficulty forthe question of the fate of exoplanet systems after the RGB phase itself is the difficulty of distinguishing RGB and RC stars based on their spectroscopic parameters alone, which is sometimes hard even with help of asteroseismology (Campante et al. 2019). As a consequence, only large or massive planets are detected around the classical evolved stars (Jones et al. 2021, and references therein). A dearth of close-in giant planets is observed around these evolved stars compared to solar-type main-sequence stars (Sato et al. 2008; D\u00f6llinger et al. 2009). This may be caused by planet engulfment by the host star, but current technologies do not allow us to determine whether smaller planets and remnants (such as the dense cores of former giant planets) are present. The lack of close-in giant planets may also be explained by the intrinsically different planetary formation for these intermediate-mass stars (see the discussion in Jones et al. 2021). Ultimately, the very existence of planet remnants may be linked to the ejection of most of the envelope on the RGB that occurs for hot subdwarfs, while for classical evolved stars, nothing stops the in-spiraling planet inside the host star, and in all cases, the planet finally merges with the star, is fully tidally disrupted, or is totally ablated by heating or by the strong stellar wind. In other words, the ejection of the envelope not only enables the detection of small objects as remnants, but most importantly, it may even be the reason for the existence of these remnants by stopping the spiraling-in in the host star.","Citation Text":["Jones et al. 2021"],"Functions Text":["The lack of close-in giant planets may also be explained by the intrinsically different planetary formation for these intermediate-mass stars (see the discussion in"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1690,1707]],"Functions Start End":[[1525,1689]]}
{"Identifier":"2022ApJ...939L..19I__Petigura_et_al._2017_Instance_1","Paragraph":"Kepler transit observations have shown that planets with sizes between those of Earth (1 R\n\u2295) and Neptune (\u223c4 R\n\u2295) are extremely common (Lissauer et al. 2011; Batalha et al. 2013; Fressin et al. 2013; Howard 2013; Fabrycky et al. 2014; Marcy et al. 2014). Demographics analysis suggests that at least 30%\u201355% of the Sun-like stars host one or more planets within this size range and with orbital periods shorter than 100 days (Mayor et al. 2011; Howard et al. 2012; Fressin et al. 2013; Petigura et al. 2013; Mulders 2018; Mulders et al. 2018; Zhu et al. 2018; He et al. 2019, 2021). Uncertainties in stellar radius estimates from photometric Kepler observations prevented a detailed assessment of the intrinsic planet size distribution (Fulton et al. 2017; Petigura et al. 2017). Specific trends in planet sizes only started to emerge from the data with more precise determination of stellar radii by follow-up surveys (California Kepler Survey, CKS) and the use of Gaia improved parallaxes (Johnson et al. 2017; Van Eylen et al. 2018; Petigura et al. 2022). These studies showed that the size frequency distribution of planets between \u223c1 and \u223c4R\n\u2295 is bimodal with peaks at \u223c1.4 R\n\u2295 and \u223c2.4 R\n\u2295, and a valley at \u223c1.8 R\n\u2295 (Fulton et al. 2017; Fulton & Petigura 2018; Petigura 2020). The best-characterized planets with sizes of about \u223c1.4 R\n\u2295 are consistent with rocky composition, as constrained by their estimated bulk densities (Fortney et al. 2007; Adams et al. 2008; Lopez & Fortney 2014; Weiss & Marcy 2014; Dorn et al. 2015; Wolfgang et al. 2016; Zeng et al. 2016, 2019; Bashi et al. 2017; Chen & Kipping 2017; Otegi et al. 2020). These planets are usually referred to as \u201csuper-Earths.\u201d Planets with sizes of about \u223c2.4R\n\u2295 are consistent with the presence of volatiles\u2014which could reflect either rocky cores with H-He-rich atmospheres or ice\/water-rich planets (Kuchner 2003; Rogers & Seager 2010; Lopez & Fortney 2014; Zeng et al. 2019; Mousis et al. 2020; Otegi et al. 2020). These planets are commonly referred to as \u201cmini-Neptunes.\u201d For a detailed discussion, see reviews by Bean et al. (2021) and Weiss et al. (2022), and references therein.","Citation Text":["Petigura et al. 2017"],"Functions Text":["Uncertainties in stellar radius estimates from photometric Kepler observations prevented a detailed assessment of the intrinsic planet size distribution"],"Functions Label":["Background"],"Citation Start End":[[758,778]],"Functions Start End":[[584,736]]}
{"Identifier":"2015AandA...584A.103S__Lattimer_&_Swesty_1991_Instance_2","Paragraph":"The energy in the inner crust is largely influenced by the properties of the neutron gas and, therefore, the EoS of neutron matter of the different calculations plays an essential role. The NV calculation (Negele & Vautherin 1973) is based on a local energy density functional that closely reproduces the Siemens-Pandharipande EoS of neutron matter (Siemens & Pandharipande 1971) in the low-density regime. The Moskow calculation (Baldo et al. 2007) employs a semi-microscopic energy density functional obtained by combining the phenomenological functional of Fayans et al. (2000) inside the nuclear cluster with a microscopic part calculated in the Brueckner theory with the Argonne v18 potential (Wiringa et al. 1995) to describe the neutron environment in the low-density regime (Baldo et al. 2004). The BBP calculation (Baym et al. 1971a,b) gives the EoS based on the Brueckner calculations for pure neutron matter of Siemens (Siemens & Pandharipande 1971). The LS-Ska (Lattimer & Swesty 1991; Lattimer 2015) and DH-SLy4 (Douchin & Haensel 2001) EoSs were constructed using the Skyrme effective nuclear forces Ska and SLy4, respectively. The SLy4 Skyrme force (Chabanat et al. 1998) was parametrized, among other constraints, to be consistent with the microscopic variational calculation of neutron matter of Wiringa et al. (1988) above the nuclear saturation density. The Shen-TM1 EoS (Shen et al. 1998b,a; Sumiyoshi 2015) was computed using the relativistic mean field parameter set TM1 for the nuclear interaction. The calculations of LS (Lattimer & Swesty 1991; Lattimer 2015) and Shen et al. (Shen et al. 1998b,a; Sumiyoshi 2015) are the two EoS tables in more widespread use for astrophysical simulations. The BSk21 EoS (Pearson et al. 2012; Fantina et al. 2013; Potekhin et al. 2013; Goriely et al. 2010) is based on a Skyrme force with the parameters accurately fitted to the known nuclear masses and constrained, among various physical conditions, to the neutron matter EoS derived within modern many-body approaches which include the contribution of three-body forces. ","Citation Text":["Lattimer & Swesty 1991"],"Functions Text":["The calculations of LS","are the two EoS tables in more widespread use for astrophysical simulations."],"Functions Label":["Background","Background"],"Citation Start End":[[1546,1568]],"Functions Start End":[[1522,1544],[1639,1715]]}
{"Identifier":"2019ApJ...885..168O__Thomas_et_al._2004_Instance_2","Paragraph":"Tidal heating of Io has been shown to be responsible for its widespread volcanism. The tidal heating rate of Jupiter\u2019s tidally locked moon, \n\n\n\n\n\n, driven by forced eccentricities, e, locked by Europa and Ganymede\u2019s Laplace resonance with Io, is the dominant interior heating source. Similarly, the tidal heating of an exomoon will likely dominate the interior energy budget due to the additional stellar tide. Consequently, the tidal heating rate is orders of magnitude higher than at Io, which for an exo-Io of similar rheological properties (\n\n\n\n\n\n, Rs = RIo, \u03c1s = \u03c1Io) can be written as (Cassidy et al. 2009; Equations (19) and (20))\n3\n\n\n\n\n\nwhere \u03c5 = 3 \u00d7 10\u22127 cm3 erg\u22121, and \u03c4s = \u03c4p\/5 based on the tidal stability criterion discussed in Section 2. For utility, we describe the exo-Io\u2019s tidal efficiency as \n\n\n\n\n\n, which can readily be computed for any three-body system as tabulated in Table 4. The enhanced tidal heating described in Equation (3) will also contribute to the surface temperature T0 = Teq + \u0394T0, which is very roughly approximated as\n4\n\n\n\n\n\nwhere \u03c3sb is the Stefan\u2013Boltzmann constant and Teq. At Io, the total neutral volcanic content (SO2, SO, NaCl, KCl, Cl, and dissociation products) ejected to space (Section 4.2.1) by the incident plasma is estimated to be, on average, \u223c1000 kg s\u22121 (e.g., Thomas et al. 2004), varying within an order of magnitude over decades of observations (Burger et al. 2001; Wilson et al. 2002; Thomas et al. 2004). While the source of the dominant gas SO2 is ultimately tidally driven volcanism, the near-surface atmosphere is mostly dominated by the sublimation of SO2 frost (Tsang et al. 2016). By observing the atmospheric evolution of the SO2 column density with heliocentric distance, Tsang et al. (2013) estimated the direct volcanic component to be Nvolc \u223c 6.5 \u00d7 1016 cm\u22122, typically \n\n\n\n\n\n of the total observed SO2 column density. Ingersoll (1989) demonstrated the relative contributions due to both sublimation and volcanic sources in maintaining Io\u2019s atmosphere and established a relationship relating the volcanic source rate to the volcanically supplied atmospheric pressure:\n5\n\n\n\n\n\nThis expression also gives the volcanic column density \n\n\n\n\n\n, where g is the acceleration due to gravity. Adopting an observed atmospheric temperature of Tatm = 170 K by Lellouch et al. (2015) corresponding to an atmospheric scale height of H = 12 km, a thermal velocity \n\n\n\n\n\n equal to 150 m s\u22121, and a sticking coefficient \u03b1 = 0.5 for the SO2 mass of 64 amu yields a volcanic source rate of \n\n\n\n\n\n \u223c 6.9 \u00d7 106 kg s\u22121 of SO2 integrated over Io\u2019s mass MIo. The average volumetric mixing ratio for NaCl to SO2 at Io is observed to be XNaCl \u223c 3 \u00d7 10\u22123 (Lellouch et al. 2003). This leads to a source rate of \n\n\n\n\n\n \u223c 7.4 \u00d7 103 kg s\u22121 of NaCl, somewhat larger than but reasonably consistent with the direct measurement of the NaCl volcanic source rate of (0.8\u20133.1) \u00d7 103 kg s\u22121 (Lellouch et al. 2003). From these estimates, we will adopt \u223c3 \u00d7 103 kg s\u22121 of Na i as the volcanic source rate for Io.","Citation Text":["Thomas et al. 2004"],"Functions Text":["varying within an order of magnitude over decades of observations"],"Functions Label":["Background"],"Citation Start End":[[1443,1461]],"Functions Start End":[[1336,1401]]}
{"Identifier":"2019ApJ...887...40W__Willott_et_al._2017_Instance_1","Paragraph":"In recent years, the Atacama Large Millimeter\/submillimeter Array (ALMA) has carried out comprehensive surveys of the [C II] 158 \u03bcm fine structure line in high-z quasars. For example, Decarli et al. (2018) detected [C II] in 85% of 27 optically selected quasars at z > 5.94. This line is an important coolant that traces the ionized and neutral interstellar medium (ISM) and star-forming activities (Herrera-Camus et al. 2015). The detection of [C II] line and dust continuum emission has revealed a wide range of star formation rates, from a few 10 M\u2299 yr\u22121 to \u22651000 M\u2299 yr\u22121 (Decarli et al. 2018; Izumi et al. 2018). The [C II]\u2013FIR luminosity ratios of these quasar hosts range from 10\u22124 to a few \u00d710\u22123 (Maiolino et al. 2005; Wang et al. 2013; Willott et al. 2017; Decarli et al. 2018), following the trend of decreasing [C II]\u2013FIR ratio with increasing FIR luminosities found with the IR luminous star-forming systems (Malhotra et al. 2001; Luhman et al. 2003; Hailey-Dunsheath et al. 2010; Stacey et al. 2010; Graci\u00e1-Carpio et al. 2011; D\u00edaz-Santos et al. 2013, 2017; Mu\u00f1oz & Oh 2016; Smith et al. 2017; Decarli et al. 2018; Gullberg et al. 2018; Rybak et al. 2019). In particular, the quasar hosts with high FIR luminosities of a few 1012 to 1013 L\u2299 show low [C II]\u2013FIR ratios similar to that found in the ultraluminous infrared galaxies (ULIRGs) and submillimeter galaxies (SMGs; Luhman et al. 2003; D\u00edaz-Santos et al. 2013, 2017; Wang et al. 2013; Rybak et al. 2019). The [C II]\u2013FIR ratio is suggested to be related to the local conditions of the ISM; e.g., a decrease of [C II]\u2013FIR ratio could be due to a high gas temperature (i.e., T > \u0394E\/k \u223c 91 K, \u0394E is the energy separation of the two levels of the [C II]158 \u03bcm transition) where the upper level of the [C II] transition is saturated (Mu\u00f1oz & Oh 2016), or a high gas surface density (e.g., the typical density in compact ULIRGs and SMGs) where more gas is in the molecular phase (Narayanan & Krumholz 2017). It is interesting to see how the resolved distributions of the line and continuum surface brightnesses and emission ratios in these FIR luminous quasar hosts compare to that in the ULIRGs and SMGs (Smith et al. 2017; Gullberg et al. 2018; Rybak et al. 2019).","Citation Text":["Willott et al. 2017"],"Functions Text":["The [C II]\u2013FIR luminosity ratios of these quasar hosts range from 10\u22124 to a few \u00d710\u22123","following the trend of decreasing [C II]\u2013FIR ratio with increasing FIR luminosities found with the IR luminous star-forming systems"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[744,763]],"Functions Start End":[[617,702],[787,918]]}
{"Identifier":"2021AandA...650A.172J__Mangum_&_Shirley_2015_Instance_1","Paragraph":"In the second, complimentary method, the line profiles were fitted using the least-squares method with the SCIPY routine CURVE_FIT. Toward B335, the fit was allowed to include up to four Gaussian line profiles, one for each potential transition in the blended emission feature. Bounds on the fits were estimated from the HDO data, and other non-blended emission lines in the dataset. The best fit does not include a fourth component, corresponding to either the CH3OCH3 line at 316.7904 GHz or the CH3OD transition at 316.7916 GHz (i.e., only three Gaussian profiles were included in the final fit). The exclusion of a fourth component has limited impact on the Gaussian profile for D2O. The resulting fit is shown the right panel of Fig. 4. The column density is estimated from the fitted line profile by assuming local thermodynamic equilibrium (LTE) (Mangum & Shirley 2015). We adopted an excitation temperature of 220 K, which is derived from the synthetic spectrum; this method results in a D2O column density of 4.0 \u00d7 1015 cm\u22122. The HDO transitions presented in Jensen et al. (2019) give a HDO column density of 1.9 \u00d7 1017 cm\u22122 at the sameexcitation temperature, when a beam filling factor is applied for a source size of 0.\u2032\u2032 2. This method then results in a D2O\/HDO ratio of (2.1 \u00b1 0.6) \u00d7 10\u22122. Evidently, the D2O\/HDO ratios toward B335 are somewhat higher when using the Gaussian fits, however both methods yield high D2O\/HDO ratios \u227310\u22122. We note that the HDO column density is different from that presented in Jensen et al. (2019) because a higher excitation temperature of 220 K and a smaller source size are used in this work. Furthermore, the synthetic model includes an optical depth correction and identifies a slight blending of the H\n\n$_{2}^{18}$218\nO line, increasing the HDO\/H2O ratio. Because of this, we consider the column densities presented in this work as more accurate. However, this does not impact the results and discussions of the relative HDO\/H2O abundances in Jensen et al. (2019), in which identical excitation temperatures were used for all sources. The differences in the reported HDO\/H2O ratios in this work and in Jensen et al. (2019) illustrate the dependence on the methodology and assumptions, however, the observed dichotomy between isolated and clustered protostars reported in that paper is robust. To remove the observed dichotomy requires consistently adopting lower excitation temperatures for the isolated sources, independent of the luminosity of the respective sources.","Citation Text":["Mangum & Shirley 2015"],"Functions Text":["The column density is estimated from the fitted line profile by assuming local thermodynamic equilibrium (LTE)"],"Functions Label":["Uses"],"Citation Start End":[[854,875]],"Functions Start End":[[742,852]]}
{"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"],"Functions Text":["In 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"],"Functions Label":["Uses"],"Citation Start End":[[1006,1027]],"Functions Start End":[[700,1004]]}
{"Identifier":"2016AandA...588A..44Y__Jones_et_al._2014_Instance_2","Paragraph":"The second issue concerns the fact that inside a given region, coreshine is not detected in all the dense clouds observed by Paladini (2014) and Lef\u00e8vre et al. (2014) and that the proportion of clouds exhibiting coreshine varies from one region to another. For instance, 75% of the dense clouds detected in Taurus exhibit coreshine, whereas in most other regions the proportion is closer to 50% (such as Cepheus, Chamaeleon, and Musca)5. On the contrary, there are for instance very few detections in the Orion region. In THEMIS, most of the scattering efficiency originates in the accretion of an a-C:H mantle. This leads to three possible explanations for the absence of detectable coreshine. The first explanation is related to the amount of carbon available in the gas phase. The abundance used by K\u00f6hler et al. (2015) relies on the highest C depletion measurements made by Parvathi et al. (2012) towards regions with \\hbox{$N_{\\rm H} \\geqslant 2 \\times 10^{21}$}NH\u2a7e 2 \u00d7 1021 H\/cm2. Parvathi et al. (2012) highlighted the variability in the carbon depletion in dust depending on the line of sight. Thus, there may be clouds were the amount of carbon available for a-C:H mantle formation is smaller or even close to zero: such regions would be populated with aggregates with a thinner H-rich carbon mantle or no second mantle at all and thus exhibit very little or no coreshine emission. A second explanation is related to the stability of H-rich carbon in the ISM, which depends strongly on the radiation field intensity to local density ratio (Godard et al. 2011; Jones et al. 2014). In low-density regions (according to Jones et al. 2014, \\hbox{$A_{V} \\leqslant 0.7$}AV\u2a7d 0.7 for the standard ISRF), UV photons are responsible for causing the photo-dissociation of CH bonds, a-C:H \u2192 a-C. In transition regions from diffuse ISM to dense clouds (Jones et al. 2014, \\hbox{$0.7 \\leqslant A_{V} \\leqslant 1.2$}0.7 \u2a7d AV\u2a7d 1.2 for the standard ISRF), better shielded from UV photons and where the amount of hydrogen is significantly higher, H-poor carbon can be transformed into H-rich carbon through H atom incorporation, a-C \u2192 a-C:H. Similarly, carbon accreted from the gas phase in these transition regions is likely to be and stay H-rich. Then, in the dense molecular clouds, most of the hydrogen is in molecular form and thus not available to produce a-C:H mantles on the grains. However, this approximately matches the density at which ice mantles start to accrete on the grains, which would partly protect a-C:H layers that had formed earlier (Godard et al. 2011, and references therein). The stability and hydrogenation degree of a-C:H, as well as the exact values of AV thresholds, are both dependent on the timescale and UV field intensity. The resulting a-C \u2194 a-C:H delicate balance could explain why in a quiet region such as Taurus most of the clouds exhibit coreshine, whereas in Orion, where on average the radiation field intensity and hardness are much higher, most clouds do not. A third explanation is related to the age and\/or density of the clouds. In a young cloud, where dust growth is not advanced, or in an intermediate density cloud (\u03c1C ~ a few 103 H\/cm3), the dust population may be dominated by CMM grains instead of AMM(I) dust. Such clouds would be as bright in the IRAC 8 \u03bcm band as in the two IRAC bands at 3.6 and 4.5 \u03bcm, thus not matching the selection criteria defined by Pagani et al. (2010) and Lef\u00e8vre et al. (2014) and would be classified as \u201cno coreshine\" clouds. ","Citation Text":["Jones et al. 2014"],"Functions Text":["In low-density regions (according to","\\hbox{$A_{V} \\leqslant 0.7$}AV\u2a7d 0.7 for the standard ISRF), UV photons are responsible for causing the photo-dissociation of CH bonds, a-C:H \u2192 a-C."],"Functions Label":["Background","Background"],"Citation Start End":[[1626,1643]],"Functions Start End":[[1589,1625],[1645,1792]]}
{"Identifier":"2021AandA...654A..80S__Leclercq_et_al._2017_Instance_1","Paragraph":"For this comparison we rely on HST broad-band magnitudes, Ly\u03b1 line fluxes, Ly\u03b1 EW0 estimates, spectral UV slopes \u03b2, continuum magnitudes, and Ly\u03b1 FWHM from the catalog that will be presented by Kerutt et al. (2021). This LAE study is based on the same data and source identification (see Sect. 3) as the current study, but focuses on properties of the emanating Ly\u03b1 emission. The Ly\u03b1 emission fluxes correspond to the measured flux within 3D apertures of three Kron (1980) radii as measured by LSDCat when detecting sources in the MUSE data cubes. We use these Ly\u03b1 fluxes as opposed to obtaining them directly from the TDOSE spectra, as the TDOSE extractions are based on the assumption that the morphological extent of the line emission follows the continuum morphology of the modeled HST images (Sect. 4 and Schmidt et al. 2019). However, Ly\u03b1 emission is known to be extended beyond the continuum (Steidel et al. 2011; Momose et al. 2014; Wisotzki et al. 2016, 2018; Leclercq et al. 2017, 2020) and fluxes based on the TDOSE spectra would therefore be biased. The LSDCat Kron radii fluxes therefore better represents the actual Ly\u03b1 flux emitted by the LAEs. As for the secondary UV emission lines, the EW0(Ly\u03b1) values were calculated by comparing the fluxes to the continuum flux densities estimated from a continuum represented by a power law f\u03bb\u2004\u221d\u2004\u03bb\u03b2. To obtain the spectral slope, Kerutt et al. (2021) first determined the magnitudes from available ancillary broad-band HST photometry, by fitting the rest-frame UV morphology for each of the LAEs using GALFIT (Peng et al. 2010, 2002). This provided morphological parameters for all LAEs including a measure of their effective radius. The estimated absolute UV magnitude at 1500 \u00c5 is also based on these GALFIT models. The spectral slope \u03b2 was then obtained from fitting the continuum power law to the GALFIT-based HST magnitudes. To avoid large scatter in the EW0 measurements presented and analyzed by Kerutt et al. (2021), the EWs are based on the median \u03b2 for the full LAE sample of \u03b2\u2004=\u2004\u22121.97 similar to what was done for the secondary UV emission lines presented here. The FWHM of the Ly\u03b1 emission was measured for each source by fitting a skewed Gaussian profile (Eq. (2) by Shibuya et al. 2014) to the Ly\u03b1 line profiles in 1D spectral extractions weighted by the MUSE PSF to maximize S\/N. These fits also provide a Ly\u03b1 redshift which is more precise than the lead line redshits provided by the LSDCat source identification. We therefore use these redshifts for the analysis of the Ly\u03b1 velocity offsets described in Sect. 8. Finally, Kerutt et al. (2021) provide estimates of the systemic redshifts based on the FWHM and peak separation between any double-peaked LAEs (identified through visual inspection of the 1D spectra) in the sample based on the empirical relations presented by Verhamme et al. (2018). We note that a handful of the z\u2004>\u20042.9 objects studied here are not included in the Kerutt et al. (2021) catalog, as their selection was based only on non-AGN objects with leading Ly\u03b1 emission based on the LSDCat selections (Sect. 3). Hence, for z\u2004>\u20042.9, the objects with IDs 121033078, 601381485, 720470421, 722551008, 722731033, and 723311101 are not included in the LAE parameter comparisons in the following. For further details and for the full value-added catalog of Ly\u03b1-related quantities we refer to Kerutt et al. (2021).","Citation Text":["Leclercq et al. 2017"],"Functions Text":["However, Ly\u03b1 emission is known to be extended beyond the continuum","and fluxes based on the TDOSE spectra would therefore be biased."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[969,989]],"Functions Start End":[[832,898],[997,1061]]}
{"Identifier":"2022AandA...661A.129S__Rodr\u00edguez-Almeida_et_al._2021_Instance_1","Paragraph":"Radio astronomy is recognized as one of the most effective techniques to search for interstellar molecules. By comparing the spectra of candidate molecules in the laboratory with the spectra observed in astronomical surveys, we can determine whether these molecules exist in interstellar space. Therefore, it is necessary to provide rotational spectra of candidates for astronomical detection. Radio astronomy has helped to detect several sulfur-containing molecules in the ISM in recent years: in particular, thiols, the sulfur analogs of alcohols. Methanethiol (or methyl mercaptan, CH3SH) was detected in the Sagittarius B2 (Sgr B2) region of the center of our Galaxy (Linke et al. 1979; Gibb et al. 2000; M\u00fcller et al. 2016; Rodr\u00edguez-Almeida et al. 2021) and in the protostar IRAS 16293-2422 (Majumdar et al. 2016). Two groups reported to have detected several signs of ethanethiol (C2H5SH) in Sgr B2 (M\u00fcller et al. 2016) and Orion (Kolesnikov\u00e1 et al. 2014). Most recently, Rodriguez-Almeida reported the first unambiguous detection of ethanethiol in the ISM, toward the G+0.603-0.027 molecular cloud (Rodr\u00edguez-Almeida et al. 2021). Moreover, several sulfur-containing species have been observed in comets (Altwegg et al. 2017). Some recent efforts, both from spectroscopy and astronomical searches, to detect S-stitutes of other classes of compounds have also been reported. For instance, thioformic acid (HC(O)SH) was very recently detected in G+0.693\u20130.027. Its trans-isomer has an abundance of ~1 \u00d7 10\u201310 (Rodr\u00edguez-Almeida et al. 2021). Conversely, thioformamide (NH2CHS), the counterpart of for-mamide (NH2CHO), was characterized in the laboratory up to 660 GHz, and its transitions were searched for toward the hot cores Sgr B2(N1S) and Sgr B2(N2), but it was not detected (Motiyenko et al. 2020). The rotational spectrum of thioac-etamide was recently analyzed in the 59.6\u2013110.0 GHz frequency region (5.03\u20132.72 mm). Its emission was searched for in regions associated with star formation using the IRAM 30 m ASAI observations toward the prestellar core L1544 and the outflow shock L1157\u2013B1. The molecule was not detected, but the study allowed placing constraints on the thioacetamide abundances (Maris et al. 2019; Remijan et al. 2022).","Citation Text":["Rodr\u00edguez-Almeida et al. 2021"],"Functions Text":["Methanethiol (or methyl mercaptan, CH3SH) was detected in the Sagittarius B2 (Sgr B2) region of the center of our Galaxy"],"Functions Label":["Background"],"Citation Start End":[[729,758]],"Functions Start End":[[550,670]]}
{"Identifier":"2022ApJ...928....3A__Wedemeyer_&_Steiner_2014_Instance_1","Paragraph":"Solar vortex tubes can be spontaneously generated by turbulent convection. In simulations of quiet Sun regions, vortices are found along intergranular lanes (Shelyag et al. 2011a; Kitiashvili et al. 2012; Moll et al. 2012; Silva et al. 2020). These structures have an average lifetime of around 80 s (Silva et al. 2021) and a radius between 40 and 80 km (Shelyag et al. 2013; Silva et al. 2020). Solar kinetic vortex tubes (Silva et al. 2021) act as a sink for magnetic field, creating magnetic flux tubes that expand with height (Kitiashvili et al. 2012; Moll et al. 2012; Silva et al. 2020). The concentration of magnetic flux leads to a high magnetic field tension, which can prevent the magnetic field lines from being twisted by the rotational motion (Shelyag et al. 2011b; Moll et al. 2012; Nelson et al. 2013; Silva et al. 2021). In some cases, twisted magnetic flux tubes appear close enough to flow vortices, leading to magnetic and kinetic vortex structures closely coexisting in regions with high plasma-\u03b2 (Wedemeyer & Steiner 2014; Rappazzo et al. 2019; Silva et al. 2021). The vortical motions can still trigger perturbations along magnetic lines that could lead to wave excitation, e.g., Battaglia et al. (2021). The vorticity evolution in the magnetized solar atmosphere is mainly ruled by the magnetic field, which also influences the general shape of vortices (Shelyag et al. 2011a). Based on the analysis of swirling strength, the part of the vorticity only linked to swirling motion (Shelyag et al. 2011b; Canivete Cuissa & Steiner 2020) showed that the magnetic terms in the swirling equation evolution tend to cancel the hydrodynamic terms close to the solar surface, whereas the magnetic terms dominate alone the production of swirling motion in the chromosphere. The magnetic field also tends to play an important role in the plasma dynamics along the whole vortex tube, as the Lorentz force has a magnitude comparable to the pressure gradient (Silva et al. 2020; Kitiashvili et al. 2013). High-speed flow jets have also been linked to simulated vortex tubes, driven by high-pressure gradients close to the photosphere and by Lorentz force in the weakly magnetized upper solar photosphere (Kitiashvili et al. 2013). In general, the averaged radial profile of magnetic field, angular velocity, pressure gradient inside of the vortex tube at the lower chromosphere and photosphere levels show similar behavior (Silva et al. 2020).","Citation Text":["Wedemeyer & Steiner 2014"],"Functions Text":["In some cases, twisted magnetic flux tubes appear close enough to flow vortices, leading to magnetic and kinetic vortex structures closely coexisting in regions with high plasma-\u03b2"],"Functions Label":["Background"],"Citation Start End":[[1018,1042]],"Functions Start End":[[837,1016]]}
{"Identifier":"2021AandA...649A..58L__Bemporad_et_al._(2018)_Instance_2","Paragraph":"The leading edges of the transients normally leave bright traces in the images of visible light, inspiring many methods that were developed to derive their locations and velocities, such as the icecream cone model (Fisher & Munro 1984), the graduated cylindrical shell (GCS) model (Thernisien 2011), geometric triangulation methods (Liu et al. 2010), mask-fitting methods (Feng et al. 2012), and trace-fitting methods including the point-p, fixed-\u03a6, harmonic mean, and self-similar expansion fitting methods (e.g., Sheeley et al. 1999; Howard et al. 2006; Davies et al. 2012; M\u00f6stl & Davies 2013). To derive the velocity distribution inside one transient rather than only at its leading edge, some other techniques have been proposed. Colaninno & Vourlidas (2006) applied an optical flow tool to extract the velocity vector of a coronal mass ejection (CME) in digital images. Feng et al. (2015) derived the radial velocity profiles of the whole CME from the spatial distribution of its density given by the mass continuum equation. A cross-correlation method was applied to derive continuous 2D speed maps of a CME from coronagraphic images by Bemporad et al. (2018). In their work, the radial shift pixel by pixel is determined by maximizing the cross correlation between the signal in a radial window at one frame and the signal in a radial shifted window at the previous frame, and the radial speed just equals the radial shift over the time interval between the two frames. Ying et al. (2019) improved this cross-correlation method by analyzing data in three steps: forward step (FS), backward step (BS), and average step (AS). In the FS (BS), the 2D velocity map between the current and the previous (next) frame is constructed with almost the same method as Bemporad et al. (2018). In the AS the average, velocity is obtained from the FS and BS. The velocities derived by all these methods are the component of the flow velocity vector projected onto the POS. This may underestimate the velocity especially for transients that do not propagate in the POS. Methods such as the polarizaition ratio technique (Moran & Davila 2004; DeForest et al. 2017) or the local correlation tracking (LCT) method (Mierla et al. 2009) can derive the 3D geometric information of the whole transients, but not the velocity distribution. Bemporad et al. (2018) chose the propagating direction averaged over the whole CME derived by the polarization ratio technique to correct the radial speed in the 2D maps, but the key information along the LOS is still lacking.","Citation Text":["Bemporad et al. (2018)"],"Functions Text":["In the FS (BS), the 2D velocity map between the current and the previous (next) frame is constructed with almost the same method as"],"Functions Label":["Uses"],"Citation Start End":[[1764,1786]],"Functions Start End":[[1632,1763]]}
{"Identifier":"2021ApJ...909..172Z__Read_&_Lebonnois_2018_Instance_1","Paragraph":"Atmospheric superrotation is characterized by eastward wind at the equator, which means the atmosphere there has a higher angular momentum than the solid surface. Atmospheric superrotation is a common phenomenon across the universe. In the solar system, superrotation exists in the atmospheres of Venus, Titan, Saturn, and Jupiter, as well as the stratospheric atmosphere of Earth during the westerly phase of the quasi-biennial oscillation (e.g., Kraucunas & Hartmann 2005; Schneider & Liu 2009; Lutsko 2018; Read & Lebonnois 2018). In order to maintain atmospheric superrotation, there must be momentum transports from higher latitudes to the equator against friction or other processes, according to angular momentum conservation (Hide 1969; Held 1999; Showman et al. 2013). This up-gradient transport into the jet can result from Rossby waves, coupled Rossby\u2013Kelvin waves, mixed Rossby\u2013gravity waves, wave\u2013jet resonance, barotropic instability, or baroclinic instability (Suarez & Duffy 1992; Del Genio & Zhou 1996; Joshi et al. 1997; Lee 1999; Williams 2003; Kraucunas & Hartmann 2005; Schneider & Liu 2009; Caballero & Huber 2010; Mitchell & Vallis 2010; Showman & Polvani 2010, 2011; Showman et al. 2010; Liu & Schneider 2011; Arnold et al. 2012; Pinto & Mitchell 2014; Tsai et al. 2014; Wang & Mitchell 2014; Laraia & Schneider 2015; Lutsko 2018; Read & Lebonnois 2018; Pierrehumbert & Hammond 2019). For example, Kraucunas & Hartmann (2005) suggested that in an Earth-like atmosphere, equatorial superrotation can be generated by equatorward stationary eddy momentum convergence, which is associated with zonal variations in the diabatic heating at low latitudes. Mitchell & Vallis (2010) studied the transition from current Earth-like atmospheric circulation to an equatorial superrotation state. They found that during the spin-up period, superrotation is generated by equatorward momentum convergence associated with both barotropic and baroclinic instabilities.","Citation Text":["Read & Lebonnois 2018"],"Functions Text":["In the solar system, superrotation exists in the atmospheres of Venus, Titan, Saturn, and Jupiter, as well as the stratospheric atmosphere of Earth during the westerly phase of the quasi-biennial oscillation (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[510,531]],"Functions Start End":[[233,447]]}
{"Identifier":"2019ApJ...885...81S__Dekel_et_al._2009b_Instance_1","Paragraph":"This scenario could also explain the reason why star-forming galaxies with a relatively thin disk appeared around z \u223c 1, because the hot-mode accretion is expected not to dominate even in massive halos at z \u2273 2 owing to gas supply through the cold gas stream (e.g., Dekel et al. 2009a; Kere\u0161 et al. 2009). Some of the gas accreting onto a dark matter halo is expected to penetrate surrounding hot gas in a form of filaments of dense and cold infalling gas and directly accrete onto the central galaxy at such high redshift, where the mass accretion rate and matter density tend to be high. Such direct gas supply through the filaments of cold gas could make the gas disk of the central galaxy more turbulent and gravitationally unstable, which leads to thick and clumpy stellar disk and bulge formation\/growth in some cases (e.g., Dekel et al. 2009b; Ceverino et al. 2010; Dekel & Burkert 2014). Thus, it seems to be difficult to form the thin stellar disks even in massive halos at z \u2273 2. After the hot-mode accretion starts to dominate at z \u223c 2 in relatively massive halos, the thin stellar disks may be gradually formed from thinner gas disks and appear around z \u223c 1. If some of these star-forming galaxies with a thin stellar disk quench and evolve into passively evolving galaxies without a violent morphological change as discussed in the previous section, it is understood that the fraction of passively evolving galaxies with a thin shape increases with time at z  1 rather than higher redshifts. Since more massive star-forming galaxies have a sufficient time to form a thin disk through the hot-mode accretion in earlier epochs in this scenario (Noguchi 2019), passively evolving galaxies with a thin shape may also be provided preferentially in more massive galaxy populations at z \u223c 1. This could explain our result that passively evolving galaxies with Mstar = 1010.5\u20131011 M\u2299 already show a thinner shape than those with 1010\u20131010.5 M\u2299 at 0.6  z  1.0. In fact, Bezanson et al. (2018) reported that \u223c64% of quiescent galaxies with Mstar \u223c 1010.5\u20131011 M\u2299 at 0.6  z  1.0 show a significant rotation, while those massive galaxies with Mstar > 1011 M\u2299 show no or little rotation in the LEGA-C survey. Such quiescent galaxies with a significant rotation might be recently quenched galaxies with a relatively thin disk that has grown through the hot-mode gas accretion since z \u2272 2.","Citation Text":["Dekel et al. 2009b"],"Functions Text":["Some of the gas accreting onto a dark matter halo is expected to penetrate surrounding hot gas in a form of filaments of dense and cold infalling gas and directly accrete onto the central galaxy at such high redshift, where the mass accretion rate and matter density tend to be high. Such direct gas supply through the filaments of cold gas could make the gas disk of the central galaxy more turbulent and gravitationally unstable, which leads to thick and clumpy stellar disk and bulge formation\/growth in some cases (e.g.,"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[831,849]],"Functions Start End":[[306,830]]}
{"Identifier":"2022MNRAS.514.2974M__the_2003_Instance_1","Paragraph":"NGC 2992 has been observed at different flux levels (e.g. Fig. 14) and in Marinucci et al. (2018) a detailed analysis of XMM\u2013Newton exposures of this object is presented. In order to provide an holistic view of the spectral properties of the central engine in NGC 2992, we tested our 2019 best-fitting model (model 1 presented in Section 4.3) on the exposures where the source was found in its lowest and highest states. In particular, we considered the XMM\u2013Newton archival observations taken on 2003-05-19 and 2010-11-28 (Obs.IDs. 01479203014 and 0654910901, respectively). We thus reproduced them directly adopting the best-fitting model found for orbit 2 and show it in Fig. 12. We accounted for the different flux states computing the normalization of the primary emission, its associated reflected component, and the one of the apec and Cloudy tables. All the other parameters have been kept fixed to their corresponding best-fitting values already quoted in Table 5. Moreover, to account for the broad emission line found required by the 2003 data, we added a Gaussian component whose width was kept fixed to \u03c3 = 400 eV, in accordance with what literature papers (e.g. Nandra et al. 1997; Shu et al. 2010). This basic procedure led us to the fits shown in Fig. 15, with statistics of \u03c72\/d.o.f. = 218\/170 and \u03c72\/d.o.f. = 181\/140. In the high flux level, (F2\u201310 keV = (9.5 \u00b1 0.1) \u00d7 10\u221211 erg\u2009cm\u22122 s\u22121), the normalization of the power law is Normpo = (3.00 \u00b1 0.01) \u00d7 10\u22122 ph.\u2009keV\u22121\u2009cm2 s\u22121, about twice of what found in 2019 while the amount of reflected flux is fully consistent with what found in 2019 as we obtained NormMyTorus = (9.7 \u00b1 3.7) \u00d7 10\u22122 ph.\u2009keV\u22121\u2009cm2 s\u22121. On the other hand, in the 2010 low flux level exposure (F2\u201310 keV  = (2.9 \u00b1 0.2) \u00d7 10\u221212 erg\u2009cm\u22122 s\u22121), we found the power-law normalization to be Normpo = (4.7 \u00b1 0.2) \u00d7 10\u22124 ph.\u2009keV\u22121\u2009cm2 s\u22121 about 20\u201330 times lower than in 2019. The normalization of the reflected component modelled using MyTorus is NormMyTorus = (9.5 \u00b1 1.2) \u00d7 10\u22123 ph.\u2009keV\u22121\u2009$\\rm cm^{2}$\u2009s\u22121, a factor of \u223c10 less than in 2019. Therefore, on very long time-scales and during a prolonged low state of the source in 2010, the strength of the reflector appears to respond to the continuum. However, the smaller value of the reflected component found in 2010 can be explained by the torus reflecting the primary continuum of NGC 2992 during a low flux state. Observing the light curves in Fig. 1, before 2010 NGC 2992 was observed in a very low flux state, even lower than the one in 2021. Such a long-term adjustment suggests that reflected spectrum emerges far from the central engine. Finally, in accordance with previous studies, the Fe K\u03b1 emission line of NGC 2992 has an unresolved component correlating with the primary flux and emerging from the broad line region. However, the current MyTorus-based model accounts for whole Fe K\u03b1 flux (see residuals in Fig. 15) and data do not require any additional Gaussian component. Below 1 keV, the non-variable behaviour of the distant scattering off the NLR can be witnessed in the top panels of Fig. 7 or in the first bin of the excess variance spectra in both Figs 6 and  8.","Citation Text":["Nandra et al. 1997"],"Functions Text":["Moreover, to account for the broad emission line found required by the 2003 data, we added a Gaussian component whose width was kept fixed to \u03c3 = 400 eV, in accordance with what literature papers (e.g."],"Functions Label":["Uses"],"Citation Start End":[[1175,1193]],"Functions Start End":[[973,1174]]}
{"Identifier":"2019MNRAS.484.3356G__H\u00e4ring-Neumayer_et_al._2006_Instance_1","Paragraph":"Optical-NIR Spectral Energy Distribution (SED) modelling of diffuse light (e.g. Carson et al. 2015; Dale et al. 2016) at high spatial resolution can provide a map of the spatial variation and composition of the main stellar population components of the NSC, as well as their host galaxy. Such analysis can also unveil (an additional) contribution to the optical-NIR light from accretion on to an obscured nuclear MBH and\/or nuclear star formation activity (e.g. Noll et al. 2009; Drouart et al. 2016). To understand NSC formation, it is therefore critical to be able to disentangle such degeneracies. For example, follow-up spectroscopic observations aiming at decomposing the main stellar populations and the dynamical imprint of a MBH rely on a good mass model to predict the stellar population velocity profile (e.g. H\u00e4ring-Neumayer et al. 2006; Seth et al. 2010; Neumayer et al. 2011; Neumayer & Walcher 2012; Nguyen et al. 2018). Observations of NSCs have shown that although they contain young stellar populations and extended star formation histories (SFHs), the most dominant one by mass is the oldest (\u22733\u2009Gyr), where more than 50 per cent of the mass of the cluster has formed (e.g. Walcher et al. 2006; Kacharov et al. 2018, from spectral modelling in the optical). Therefore, characterizing the spatial structure of NSCs in the NIR is of particular importance, because that is where most of the stellar light of the old stellar population is emitted that allows us to trace most of the gravitating mass. Therefore, characterizing NSCs in the NIR can provide additional constraints as to which of the leading NSC formation scenarios had a major role in the formation of particular NSC. However, to be able to achieve this for a larger sample and of more distant galaxies, efficient high spatial resolution observations are required, such as with the presented here wide-field ground-layer AO NIR observations in the NIR with ARGOS at the LBT.","Citation Text":["H\u00e4ring-Neumayer et al. 2006"],"Functions Text":["To understand NSC formation, it is therefore critical to be able to disentangle such degeneracies. For example, follow-up spectroscopic observations aiming at decomposing the main stellar populations and the dynamical imprint of a MBH rely on a good mass model to predict the stellar population velocity profile (e.g."],"Functions Label":["Motivation"],"Citation Start End":[[820,847]],"Functions Start End":[[502,819]]}
{"Identifier":"2021MNRAS.508.3111M__Garton,_Gallagher_&_Murray_2018_Instance_1","Paragraph":"The initially used automated methods for detection of solar activity in the solar images are based on the predefined rules inferred from the appearance and usual characteristics of the structures in the solar atmosphere (Henney & Harvey 2005; Krista & Gallagher 2009; P\u00e9rez-Su\u00e1rez et al. 2011). However, it is not possible to capture all the unique solar structures using generally defined rules. Therefore, methods with more advanced mathematical algorithms were introduced. The Spatial Possibilistic Clustering Algorithm (SPoCA) (Barra et al. 2009; Verbeeck et al. 2014) or Coronal Hole Identification via Multi-thermal Emission Recognition Algorithm (CHIMERA) (Garton, Gallagher & Murray 2018) have been found to be very effective and are widely used in online solar data visualization tools1$^,$\u2009.2 The SPoCA also provides entries for catalogues of coronal holes and active regions within the Heliophysics Events Knowledgebase (HEK) (Hurlburt et al. 2012) that is commonly used in the SolarSoft (Freeland & Handy 1998) and SunPy (The SunPy Community et al. 2020) frameworks. As will be presented later, these algorithms still have limitations for the precise segmentation of structures in the solar corona. Due to the advances in computer vision in recent years, approaches based on machine-learning techniques are able to extend the standard methods (Aschwanden 2010). Conventional machine-learning techniques as support vector machine (SVM), Decision Tree, or Random Forest could improve the detection of coronal holes as they provide automated distinguishing from filaments in EUV solar images (Reiss et al. 2015; Delouille et al. 2018). On the assumption that the techniques based on convolutional neural networks (CNN) are the most convenient techniques for image segmentation tasks (Lecun, Bengio & Hinton 2015), a new era in automated processing of solar images has begun. Illarionov & Tlatov (2018) demonstrated that for the segmentation of coronal holes, CNN provides quantitatively comparable results as the SPoCA and CHIMERA algorithms. Jarolim et al. (2021) provided a method called CHRONOS based on CNN with progressively growing network approach for robust segmentation of coronal holes.","Citation Text":["Garton, Gallagher & Murray 2018"],"Functions Text":["Therefore, methods with more advanced mathematical algorithms were introduced.","Coronal Hole Identification via Multi-thermal Emission Recognition Algorithm (CHIMERA)","have been found to be very effective and are widely used in online solar data visualization tools1$^,$\u2009.2"],"Functions Label":["Background","Background","Background"],"Citation Start End":[[664,695]],"Functions Start End":[[397,475],[576,662],[697,802]]}
{"Identifier":"2021MNRAS.506.5129B__Momose_et_al._2014_Instance_1","Paragraph":"More recently, Lyman \u03b1 haloes (LAHs) have been discovered around star-forming galaxies that show Ly\u2009\u03b1 emission far beyond the galaxies\u2019 optical bodies, tracing the circumgalactic rather than interstellar gas (e.g. Hayes et al. 2013). While LAHs are fainter and smaller than LABs in their Ly\u2009\u03b1 extent, they might be a generic feature around Ly\u2009\u03b1 emitting galaxies. Narrow-band imaging can efficiently detect LAHs at targeted redshifts through stacking (Hayashino et al. 2004; Steidel et al. 2011; Matsuda et al. 2012; Feldmeier et al. 2013), and narrow-band surveys enable ultradeep, blind samples of LAHs around distant galaxies (Momose et al. 2014, 2016; Kakuma et al. 2019). In the last years modern surveys performed with integral field unit (IFU) spectrographs on 10\u2009m-class telescopes, such as the Multi Unit Spectroscopic Explorer (MUSE) and the Keck Cosmic Web Imager (KCWI), take place. These new instruments allow the study of individual, faint LAHs opposed to previous narrow-band stacks. Along with the IFUs\u2019 spectral resolution these recent surveys largely increase the information available from LAH observations. Hundreds of individually extended Lyman \u03b1 haloes at z \u2273 2 have been revealed since (Wisotzki et al. 2016). Many of these are specifically targeted samples which focus on bright quasars (Borisova et al. 2016; Cai et al. 2019; Guo et al. 2020; O\u2019Sullivan et al. 2020), based on strong earlier evidence of enhanced Ly\u2009\u03b1 emission around active galactic nuclei (AGNs Cantalupo et al. 2014; Arrigoni Battaia et al. 2016; Arrigoni Battaia et al. 2019; Farina et al. 2019). Others exploit the large field of view of MUSE, in particular, to conduct blind surveys for LAHs around more typical, generally star-forming galaxies (Leclercq et al. 2017; Wisotzki et al. 2018; Leclercq et al. 2020). At the same time, follow-up with other instruments such as ALMA reveals complementary views on other gas phases within LAHs including CO (Emonts et al. 2019).","Citation Text":["Momose et al. 2014"],"Functions Text":["Narrow-band imaging can efficiently detect LAHs at targeted redshifts through stacking","and narrow-band surveys enable ultradeep, blind samples of LAHs around distant galaxies"],"Functions Label":["Background","Background"],"Citation Start End":[[630,648]],"Functions Start End":[[364,450],[541,628]]}
{"Identifier":"2021MNRAS.508.3499V__Takahashi,_Witti_&_Janka_1994_Instance_1","Paragraph":"The astrophysical environments of r-process events in the cosmos usually involve explosive physical conditions, because very high neutron densities are required. The neutron-rich isotopes produced along the r-process path have a very short half-life, causing the neutron-rich nuclei to \u03b2-decay over time-scales of the order of milliseconds, if they did not undergo further rapid neutron capture. Examples of r-process sites that have been proposed and investigated by theoretical studies include (i) neutrino-driven winds from proto-neutron stars (NSs; Takahashi, Witti & Janka 1994; Woosley et al. 1994; Qian & Woosley 1996; Hoffman, Woosley & Qian 1997; Otsuki et al. 2000; Thompson, Burrows & Meyer 2001; Wanajo et al. 2009; Hansen et al. 2013; Wanajo 2013); (ii) neutrino-driven winds from highly magnetic and potentially rapidly rotating proto-magnetars (Thompson 2003; Thompson, Chang & Quataert 2004; Metzger, Thompson & Quataert 2007, 2008; Vlasov, Metzger & Thompson 2014; Vlasov et al. 2017; Thompson & ud-Doula 2018); (iii) neutrino-driven winds around the accretion disc of a black hole (Pruet, Thompson & Hoffman 2004; Metzger et al. 2008; Wanajo & Janka 2012; Siegel, Barnes & Metzger 2019); (iv) electron-capture SNe (see e.g. Wanajo et al. 2009; Cescutti et al. 2013; Kobayashi, Karakas & Lugaro 2020); (v) magneto-rotationally driven SNe (Burrows et al. 2007; Winteler et al. 2012; Cescutti & Chiappini 2014; M\u00f6sta et al. 2014, 2015, 2018; Nishimura, Takiwaki & Thielemann 2015; Nishimura et al. 2017; Halevi & M\u00f6sta 2018; Reichert et al. 2021); and (vi) neutron-star mergers (Lattimer et al. 1977; Freiburghaus, Rosswog & Thielemann 1999; Argast et al. 2004; Goriely, Bauswein & Janka 2011; Rosswog 2013; Matteucci et al. 2014; Cescutti et al. 2015; Vincenzo et al. 2015; Kobayashi et al. 2020). Since it is likely that all these mechanisms have contributed to r-process nucleosynthesis at some level, the theoretical studies have focused on exploring the frequency of each event and the predicted template of the corresponding r-process ejecta.","Citation Text":["Takahashi, Witti & Janka 1994"],"Functions Text":["Examples of r-process sites that have been proposed and investigated by theoretical studies include (i) neutrino-driven winds from proto-neutron stars (NSs;"],"Functions Label":["Background"],"Citation Start End":[[553,582]],"Functions Start End":[[396,552]]}
{"Identifier":"2018ApJ...861...77M__Pillai_2017_Instance_1","Paragraph":"Previous Herschel studies confirmed numerical calculations (e.g., Nagai et al. 1998) showing that a parallel orientation between filaments and the local magnetic field is to be expected for filaments having low column densities (e.g., Peretto et al. 2012; Busquet et al. 2013; Palmeirim et al. 2013; Andr\u00e9 et al. 2014; Cox et al. 2016; Panopoulou et al. 2016). In fact, a variety of orientations have previously been observed, and the alignment of filaments and apparent field direction appear to vary between regions and for different gas densities (cf., Pillai 2017, for a review). For example, there is observational evidence that low-column density filaments, or \u201cstriations,\u201d are aligned parallel to the magnetic field lines in Taurus, as they are more susceptible to the magnetic influence than higher column density structures (Palmeirim et al. 2013). These striations may therefore promote mass accretion onto larger filaments, in which stars are finally formed. Such studies, however, were mostly performed using lower-resolution dust column density maps obtained with Herschel (e.g., Li et al. 2013; Palmeirim et al. 2013; Soler et al. 2017) or polarization data tracing relatively low AV environments (e.g., Busquet et al. 2013; Fissel et al. 2016; Planck Collaboration Int. XXXV 2016; Soler et al. 2016; Jow et al. 2018). Nevertheless, as discussed in Section 4.3, the filament F2 studied in this work has molecular hydrogen column densities roughly 2 to 3 orders of magnitude larger than the Herschel filaments studied in Cygnus or Taurus, and may therefore possibly be self-gravitating. At face value, these results are at odds with the widely accepted paradigm of magnetically regulated star formation in filaments, in which a perpendicular orientation of dense filaments with respect to the magnetic field is expected (e.g., Peretto et al. 2012; Busquet et al. 2013; Palmeirim et al. 2013; Andr\u00e9 et al. 2014; Cox et al. 2016; Panopoulou et al. 2016).","Citation Text":["Pillai 2017"],"Functions Text":["In fact, a variety of orientations have previously been observed, and the alignment of filaments and apparent field direction appear to vary between regions and for different gas densities (cf.,","for a review)"],"Functions Label":["Background","Background"],"Citation Start End":[[556,567]],"Functions Start End":[[361,555],[569,582]]}
{"Identifier":"2022ApJ...927...61K__SN_2011f_Instance_1","Paragraph":"Statistical investigations of a large sample of photometric data are useful in exploring the bulk properties of various types of events. Nevertheless, such studies are generally limited to LCs in selected passbands, and also the follow-up covers a short duration of the SN evolution. Further, the spectroscopic follow-up of most objects is restricted to early phases. Nonuniformity in the data sample may also be present, as these are collected at different observing facilities and detectors. Studying individual events with proper monitoring at different phases (both photospheric and nebular) is hence extremely important. The very early phase observations (hours to days after explosion) of these events are useful to constrain the progenitor radius at its end stage. This needs a very early detection and quick follow-up, which is not always possible considering their random occurrence in the sky. The large-area surveys with high cadence have contributed significantly in this regard. SN 1993J (Type IIb) is the first event among SE-SNe that shows evidence of a prominent cooling tail just after explosion (Richmond et al. 1994; Barbon et al. 1995). During the past two decades, several Type IIb SNe with such interesting features have been monitored and studied well, e.g., SN 2008ax (Pastorello et al. 2008; Roming et al. 2009), SN 2011dh (Arcavi et al. 2011), SN 2011fu (Kumar et al. 2013; Morales-Garoffolo et al. 2015), SN 2013df (Fremling et al. 2014), and SN 2016gkg (Arcavi et al. 2017; Tartaglia et al. 2017). A handful of Type Ib events have also been discovered at very early phases, such as SN 1999ex (Hamuy et al. 2002; Stritzinger et al. 2002), SN 2008D (Mazzali et al. 2008; Soderberg et al. 2008; Malesani et al. 2009; Modjaz et al. 2009; Bersten et al. 2013), iPTF13bvn (Bersten et al. 2014; Fremling et al. 2016), and LSQ13abf (Stritzinger et al. 2020), and their observational properties are studied in detail. Recently, Prentice et al. (2019) analyzed the properties of 18 SE-SNe and discussed the implications for their progenitors. Direct detection of SE-SNe progenitor candidates has been possible only for a few objects, e.g., SN 1993J (Maund et al. 2004), SN 2011dh (Maund et al. 2011; Van Dyk et al. 2013), SN 2001ig (Ryder et al. 2018), iPTF13bvn (Cao et al. 2013; Eldridge et al. 2015; Kuncarayakti et al. 2015; Folatelli et al. 2016), SN 2016gkg (Kilpatrick et al. 2017; Tartaglia et al. 2017), SN 2017ein (Kilpatrick et al. 2018; Van Dyk et al. 2018; Xiang et al. 2019), and SN 2019yvr (Kilpatrick et al. 2021). Along with very early phases, the temporal observations during maximum to nebular phases are equally important to estimate various explosion parameters and progenitor properties. Detailed investigation of more events can provide an alternative way to understand various progenitor channels.","Citation Text":["Kumar et al. 2013"],"Functions Text":["During the past two decades, several Type IIb SNe with such interesting features have been monitored and studied well, e.g.,","SN 2011fu"],"Functions Label":["Background","Background"],"Citation Start End":[[1381,1398]],"Functions Start End":[[1157,1281],[1370,1379]]}
{"Identifier":"2017MNRAS.465..213B__David_et_al._1993_Instance_1","Paragraph":"As a baseline for understanding how the scaling relations evolve as a function of mass and redshift, we adopt the following self-similar scalings:\n\n(2)\n\n\\begin{equation}\nM_{\\rm {gas},\\Delta }\\propto M_{\\Delta },\n\\end{equation}\n\n\n(3)\n\n\\begin{equation}\nT_{\\Delta }\\propto M^{2\/3}_{\\Delta }E^{2\/3}(z),\n\\end{equation}\n\n\n(4)\n\n\\begin{equation}\nY_{\\rm {X},\\Delta }\\propto M_{\\Delta }^{5\/3}E^{2\/3}(z),\n\\end{equation}\n\n\n(5)\n\n\\begin{equation}\nY_{\\rm {SZ},\\Delta }\\propto M_{\\Delta }^{5\/3}E^{2\/3}(z),\n\\end{equation}\n\n\n(6)\n\n\\begin{equation}\nL_{\\Delta }^{\\rm {X,bol}}\\propto M^{4\/3}_{\\Delta }E^{7\/3}(z),\n\\end{equation}\n\n\n(7)\n\n\\begin{equation}\nL_{\\Delta }^{\\rm {X,bol}}\\propto T^{2}E(z),\n\\end{equation}\n\nwhere $E(z)\\equiv H(z)\/H_0=\\sqrt{\\Omega _{\\text{m}}(1+z)^3+\\Omega _{\\Lambda }}$, \u0394 is the chosen overdensity relative to the critical density and YX is the X-ray analogue of the integrated SZ effect. These are derived in Appendix B. Although shown to be too simplistic by the first X-ray studies of clusters (Mushotzky 1984; Edge & Stewart 1991; David et al. 1993), the self-similar relations allow us to investigate if astrophysical processes are less significant in more massive clusters or at higher redshift. To enable a comparison with the self-similar predictions, and previous work, we fit the scaling relations of our samples at each redshift. We derive a median relation by first binning the clusters into bins of log mass (width: 0.1 dex) or log temperature (width: 0.07 dex) and then computing the median in each bin with more than 10 clusters. We also remove the evolution in normalization predicted by self-similar relations. The medians of the bins are then fitted with a power law of the form\n\n(8)\n\n\\begin{equation}\nE^{\\beta }(z)Y=10^A\\left(\\frac{X}{X_0}\\right)^{\\alpha }\\!\\!,\n\\end{equation}\n\nwhere A and \u03b1 describe the normalization and slope of the best fit, respectively, \u03b2 removes the expected self-similar evolution with redshift, X is either the total mass or temperature and Y is the observable quantity (Mgas, LX, bol, etc.). X0 is the pivot point, which we set to 4 \u00d7 1014\u2009M\u2299 for observable\u2013mass relations and to 6\u2009keV for observable\u2013temperature relations. We note that we fix the pivot for all samples and all redshifts. Fitting to the medians of bins, rather than individual clusters, prevents the fit from being dominated by low-mass objects, which are significantly more abundant due to the shape of the mass function. For the hot sample and its relaxed subset, there are too few bins with 10 or more clusters to reliably derive a best-fitting relation at z \u2265 1. By limiting our sample to systems with M500 \u2265 1014\u2009M\u2299, we avoid any breaks in the power-law relations that have been seen both observationally and in previous simulation work (Le Brun et al. 2016).","Citation Text":["David et al. 1993"],"Functions Text":["Although shown to be too simplistic by the first X-ray studies of clusters","the self-similar relations allow us to investigate if astrophysical processes are less significant in more massive clusters or at higher redshift."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1036,1053]],"Functions Start End":[[923,997],[1056,1202]]}
{"Identifier":"2019MNRAS.487...24G__Rogers_2015_Instance_1","Paragraph":"NASA\u2019s Kepler mission has unveiled a wealth of new planetary systems (e.g. Borucki et al. 2010). These systems offer new insights into the process of planet formation and evolution. One of Kepler\u2019s key findings is that the most common planets in our Galaxy, observed to date, are between 1 and 4 R\u2295, i.e. larger than Earth but smaller than Neptune (Fressin et al. 2013; Petigura, Marcy & Howard 2013). Further observations revealed a transition in average densities at planet sizes \u223c1.5 R\u2295 (Marcy et al. 2014; Rogers 2015), with smaller planets having densities consistent with rocky compositions while larger planets having lower densities indicating significant H\/He envelopes. In addition, Owen & Wu (2013) noticed a bimodal distribution of observed planet radii. Since then, refined measurements have provided strong observational evidence for the sparseness of short-period planets in the size range of \u223c1.5\u20132.0 R\u2295 relative to the smaller and larger planets, yielding a valley in the small exoplanet radius distribution (e.g. Fulton et al. 2017; Fulton & Petigura 2018). For example, the California-Kepler Survey reported measurements from a large sample of 2025 planets, detecting a factor of \u223c2 deficit in the relative occurrence of planets with sizes \u223c1.5\u20132.0 R\u2295 (Fulton et al. 2017). Studies suggest that this valley likely marks the transition from the smaller rocky planets: \u2018super-Earths\u2019, to planets with significant H\/He envelopes typically containing a few per cent of the planet\u2019s total mass: \u2018sub-Neptunes\u2019 (e.g. Lopez & Fortney 2013, 2014; Owen & Wu 2013; Rogers 2015; Ginzburg, Schlichting & Sari 2016). Furthermore, the location of this valley is observed to decrease to smaller planet radii, Rp, with increasing orbital period, P. In a recent study involving asteroseismology-based high precision stellar parameter measurements for a sample of 117 planets, a slope $\\text{d log} R_\\mathrm{ p}\/ \\text{d log} P = -0.09^{+0.02}_{-0.04}$ was reported for the radius valley by Van Eylen et al. (2018). A similar value for the slope of $-0.11^{+0.03}_{-0.03}$ was reported by Martinez et al. (2019).","Citation Text":["Rogers 2015"],"Functions Text":["Further observations revealed a transition in average densities at planet sizes \u223c1.5 R\u2295","with smaller planets having densities consistent with rocky compositions while larger planets having lower densities indicating significant H\/He envelopes."],"Functions Label":["Background","Background"],"Citation Start End":[[510,521]],"Functions Start End":[[402,489],[524,679]]}
{"Identifier":"2022MNRAS.515.5121B__Spiniello_et_al._2018_Instance_1","Paragraph":"There are multiple sources of spurious detections in strong lensing analysis, and some of them are just image artefacts. None the less, others are images that might look like a lens, as in the case of edge-on galaxies and some spiral galaxies. One popular approach to reduce this issue is to use colours and look for red galaxies with a blue object close. The colour information is considered an important feature to find lenses (see e.g. Ostrovski et al. 2018; Spiniello et al. 2018). The use of colour queries is considered at least an interesting way to make a pre-selection and improve the final deep learning result purity. In fact, even in the I SGLC results the single-band scenario found lower performance in terms of the considered metric, the AUC of ROC. The II SGLC in a Euclid-like scenario allowed to make deep learning classifications using images with colour information in multiple bands but also higher resolution images. Interestingly the networks using only HJY bands did not find a competitive fit. On the other hand, using the VIS band only with a higher resolution found better results than HJY, which is close to a random guess. Still, the VIS band only has inferior results if compared to the runs using all information. This result suggests that the use of high-resolution images might play an important role and was in fact, more relevant than the multiple bands in the cases tested in this contribution. It is worth mentioning that the threshold defined by the maximum F\u03b2, with \u03b2 = 0.001, privileged a pure sample instead of a complete one, obtaining in our best network, HJY + VIS with alternative pre-processing, around ${\\sim} 99{{\\ \\rm per\\ cent}}$ purity with completeness of around $45{{\\ \\rm per\\ cent}}$. This choice is justified for the same reason a pre-selection of targets is implemented by colour queries and other methods to reduce the number of non-lenses in a survey where we have billions of non-lenses for thousands of lenses.","Citation Text":["Spiniello et al. 2018"],"Functions Text":["The colour information is considered an important feature to find lenses (see e.g.","The use of colour queries is considered at least an interesting way to make a pre-selection and improve the final deep learning result purity."],"Functions Label":["Uses","Uses"],"Citation Start End":[[462,483]],"Functions Start End":[[356,438],[486,628]]}
{"Identifier":"2015ApJ...799..138S__Yuan_&_Kewley_2009_Instance_1","Paragraph":"We present these results with one very important caveat. Accurately determining metallicities at different redshifts is of key importance to studying the evolution of the MZR. In the local universe, relationships between strong emission line ratios and metallicity can be calibrated to \u00e2\u0080\u009cdirect\u00e2\u0080\u009d electron temperature-determined metallicities from measuring auroral lines such as [O\u00e2\u0080\u0089iii]\u00c2 \u00ce\u00bb4363 (Pettini & Pagel 2004; Pilyugin & Thuan 2005) or photoionization models of star-forming regions (Zaritsky et al. 1994; Kewley & Dopita 2002; Kobulnicky & Kewley 2004; Tremonti et al. 2004). At redshifts above z \u00e2\u0088\u00bc 1, it is nearly impossible to detect weak auroral lines for directly determining metallicity (but see Yuan & Kewley 2009; Rigby et al. 2011; Brammer et al. 2012a; Christensen et al. 2012; Maseda et al. 2014). Creating photoionization models that suitably represent high-redshift star-forming regions requires knowledge of physical parameters which have been poorly constrained up to this point. Thus, it is unknown if local metallicity calibrations hold at high redshifts. Figure 6 shows a comparison between metallicities determined using the O3N2 indicator and the N2 indicator for both local SDSS galaxies (grey points) and MOSDEF z \u00e2\u0088\u00bc 2.3 galaxies (black points). The black dashed line indicates a one-to-one relationship. If local calibrations do indeed hold at high redshifts, then the relationship between metallicities determined from different indicators should not evolve with redshift. It is clear that the z \u00e2\u0088\u00bc 2.3 galaxies are offset below the local galaxies. The dotted line is the best-fit line of slope unity to the individual z \u00e2\u0088\u00bc 2.3 galaxies, yielding an offset of \u00e2\u0088\u00920.1 dex from a one-to-one correspondence, over twice that displayed by the SDSS sample. Steidel et al. (2014) found an offset slightly larger than this at z \u00e2\u0088\u00bc 2.3. This offset demonstrates that the two metallicity indicators are not evolving in the same way with redshift, and shows the need of metallicity calibrations appropriate for high-redshift galaxies.","Citation Text":["Yuan & Kewley 2009"],"Functions Text":["At redshifts above z \u00e2\u0088\u00bc 1, it is nearly impossible to detect weak auroral lines for directly determining metallicity (but see"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[717,735]],"Functions Start End":[[590,716]]}
{"Identifier":"2016ApJ...817..152X__Schady_et_al._2007_Instance_1","Paragraph":"The connection between long-duration GRBs (LGRBs) and SNe was predicted theoretically (Colgate 1974; Woosley 1993) and has been verified observationally (e.g., Galama et al. 1998; Hjorth et al. 2003; see a review in Woosley & Bloom 2006). They usually happen in the star formation regions of the galaxies (e.g., Paczy\u0144ski 1998; see the reviews in Woosley & Bloom 2006 and Kumar & Zhang 2015). The immensely bright afterglows illuminate the gas and dust within the star-forming regions of the host galaxy and intervening intergalactic medium along the GRB line of sight. Their spectra are usually featureless power-laws or broken power-laws, which can be well described by the synchrotron radiations of relativistic electrons. Therefore, GRB afterglows are good probes of burst environment and the interstellar dust and gas in distant, star-forming galaxies (Metzger et al. 1997; Jensen et al. 2001; Savaglio et al. 2003; Vreeswijk et al. 2004; Chen et al. 2005; Prochaska et al. 2007a; Schady et al. 2007, 2010; Watson et al. 2007; Fox et al. 2008; Starling et al. 2008; Jang et al. 2011; Xin et al. 2011). GRB afterglow spectra with Ly\u03b1 absorption features indicate the presence of large column densities of cold neutral gas within GRB host galaxies, and their hydrogen column densities (\n\n\n\n\n\n) are usually larger than \n\n\n\n\n\n cm\u22122 (e.g., Prochaska et al. 2007b; Schady 2012). The damped Ly\u03b1 systems (DLA) may represent the ISM near the GRBs in a few kiloparsecs (Kpc), but not gas directly local to the GRB (Prochaska et al. 2007b). Thus GRB optical afterglows may be used as probes of the ISM in their host galaxies, as the ISM observed is less affected by the GRB or its progenitor (Watson et al. 2007). The visual dust extinctions (\n\n\n\n\n\n) along the GRB lines of sight of many GRBs are low. As shown in Greiner et al. (2011), about 50% of GRBs observed with GROND after the launch of Swift mission have \n\n\n\n\n\n. In addition, the early optical light curves of about one-third of GRBs show a smooth onset bump (Li et al. 2012). It may be due to the deceleration of the GRB fireball by the ambient medium (Sari & Piran 1999; Kobayashi & Zhang 2007). In this scenario, the rising slope of the bump is determined by the medium density profile (\n\n\n\n\n\n) and the spectrum index of the accelerated electrons \n\n\n\n\n\n], says, \n\n\n\n\n\n (Liang et al. 2013). Hence, The afterglow onset bumps would be also an ideal probe to study the properties of the fireball and the profile of the circumburst medium. Liang et al. (2013) found that \n\n\n\n\n\n (see also Watson et al. 2007; Jin et al. 2012).","Citation Text":["Schady et al. 2007"],"Functions Text":["Therefore, GRB afterglows are good probes of burst environment and the interstellar dust and gas in distant, star-forming galaxies"],"Functions Label":["Motivation"],"Citation Start End":[[986,1004]],"Functions Start End":[[726,856]]}
{"Identifier":"2020ApJ...902...98G__Tacconi_et_al._2020_Instance_1","Paragraph":"On balance, a large abundance of baryon-dominated, dark matter cored galaxies at z \u223c 2, most strongly correlated with baryonic surface density, angular momentum, and central bulge mass, may be most naturally accounted for by the interaction of baryons and dark matter during the formation epoch of massive halos. Massive halos (log(Mhalo\/M\u2299) > 12) formed for the first time in large abundances in the redshift range z \u223c 1\u20133 (Press & Schechter 1974; Sheth & Tormen 1999; Mo & White 2002; Springel et al. 2005). At the same time, gas accretion rates were maximal (Tacconi et al. 2020). This resulted in high merger rates (Genel et al. 2008, 2009; Fakhouri & Ma 2009), very efficient baryonic angular momentum transport (Dekel et al. 2009; Zolotov et al. 2015), formation of globally unstable disks, and radial gas transport by dynamical friction (Noguchi 1999; Immeli et al. 2004; Genzel et al. 2008; Bournaud & Elmegreen 2009; Bournaud et al. 2014; Dekel & Burkert 2014). These processes enabled galaxy mass doubling on a timescale  0.4 Gyr at z \u223c 2\u20133, and massive bulge formation by disk instabilities and compaction events on 1 Gyr timescales. However, central baryonic concentrations would naturally also increase central dark matter densities through adiabatic contraction (Barnes & White 1984; Blumenthal et al. 1986; Jesseit et al. 2002). For adiabatic contraction to be ineffective requires the combination of kinetic heating of the central dark matter cusp by dynamical friction from in-streaming baryonic clumps (El-Zant et al. 2001; Goerdt et al. 2010; Cole et al. 2011), with feedback from winds, supernovae, and AGNs driving baryons and dark matter out again (Dekel & Silk 1986; Pontzen & Governato 2012, 2014; Martizzi et al. 2013; Freundlich et al. 2020; K. Dolag et al. 2020, in preparation). Using idealized Monte Carlo simulations, El-Zant et al. (2001) demonstrated that dynamical friction acting on in-spiraling gas clumps can provide enough energy to heat up the central dark matter component and create a finite dark matter core (see also A. Burkert et al. 2020, in preparation). They argue that dark matter core formation in massive galaxies would require that clumps be compact, such that they avoid tidal and ram-pressure disruption, and have masses of >108 M\u2299. Other idealized simulations (e.g., Tonini et al. 2006) confirm these results.","Citation Text":["Tacconi et al. 2020"],"Functions Text":["At the same time, gas accretion rates were maximal"],"Functions Label":["Background"],"Citation Start End":[[562,581]],"Functions Start End":[[510,560]]}
{"Identifier":"2017ApJ...850..197P__Popham_&_Narayan_1991_Instance_1","Paragraph":"It is worth noting that more ECSNe are predicted for systems with a mass ratio close to unity, as the development of contact happens at longer periods for higher-q systems (see Figure 12). As the primary star starts transferring mass to the secondary, the orbit shrinks, until the mass ratio is reversed. This reversal happens earlier for mass ratios close to 1 and later for lower mass ratios, increasing the chance for contact (de Mink et al. 2013). This primarily affects the numbers of ECSNe, not so much the initial mass range (except for Case A systems). This is also true for the value of \u03b2, which controls the amount of matter that is lost from the system (i.e., \n\n\n\n\n\n is conservative mass transfer, no mass that is transferred from the primary to the secondary is lost from the system; \n\n\n\n\n\n is completely nonconservative mass transfer, all mass that is transferred from the primary is lost from the system, no accretion onto the secondary). For the sake of our parameter study we chose various fixed values of \u03b2, while the mass transfer efficiency in real systems varies in time and will depend on the evolutionary phase of both stars, the amount of matter already accreted onto the secondary, and how that has affected its spin rate. Several mechanisms have been suggested that control the efficiency of mass transfer, mass accretion, and mass loss, including the existence of an accretion disk that regulates the amount of mass and angular momentum that can be accreted (Paczy\u0144ski 1991; Popham & Narayan 1991; Deschamps et al. 2013), the necessity of the secondary to stay below critical rotation (Packet 1981), and the effects of tides on the stellar spins and the stellar orbit (Zahn 1977; Hurley et al. 2002). Work by Deschamps et al. (2013) and van Rensbergen et al. (2008) for systems with slightly lower masses suggests periods with values of \u03b2 close to 1 (i.e., very inefficient mass transfer), while simulations with a strong tidal interaction (i.e., a short spin\u2013orbit synchronization timescale) suggest shorter periods of moderately inefficient mass transfer (Paxton et al. 2015). Although our models suggest that the efficiency of mass transfer does not really affect the mass range for ECSNe, it does, however, strongly affect the range of initial periods that can lead to an ECSN. In addition, the evolution of the secondary will be affected. It will most likely rapidly spin up after the onset of mass transfer and maintain near-critical rotation for possibly extended periods of time. This will induce strong rotational mixing (de Mink et al. 2008a, 2013; Langer 2012), causing possibly quasi-chemically homogeneous evolution (Maeder 1987; Langer 2012), and alter the evolution of the star beyond just the simple fact of mass accretion (Hirschi et al. 2004; de Mink & Mandel 2016; Marchant et al. 2016). Although it is not clear to what extent this will affect the incidence of ECSNe, the effects of tides, mass and angular momentum transfer and loss, and near-critical rotation of the secondary are possibly important and will be discussed in a forthcoming paper, in addition to the effects of additional mixing and convection criteria.","Citation Text":["Popham & Narayan 1991"],"Functions Text":["Several mechanisms have been suggested that control the efficiency of mass transfer, mass accretion, and mass loss, including the existence of an accretion disk that regulates the amount of mass and angular momentum that can be accreted"],"Functions Label":["Motivation"],"Citation Start End":[[1501,1522]],"Functions Start End":[[1247,1483]]}
{"Identifier":"2022MNRAS.509.3488I__Siana_et_al._2008_Instance_1","Paragraph":"In Fig. 11, we present the evolution of the quasar bolometric luminosity function (LF) from $z \\sim 3$ down $z \\sim 0$. Even though these functions give the number density of accreting black holes in different luminosity bins, they have been a powerful tool to extract information on how MBHs grow with cosmic time, on the geometry of the accretion discs and other fundamental quantities such as the black hole spins and radiative efficiencies. In this work, we only focus on the very bright objects, i.e. ${\\gt}10^{45}\\, \\rm erg\\,s^{-1}$, avoiding the comparison with lower luminosity given the current limitations on observational and theoretical models. In particular, from an observational standpoint, the covered area and depth of current surveys pose serious challenges when extracting statistical properties of the LF at the faint end (Siana et al. 2008; Masters et al. 2012; McGreer et al. 2013; Niida et al. 2016; Akiyama et al. 2018). Even more, dust attenuation effects might play an important role in shaping current measurements. On the other hand, current theoretical works show a large excesses at luminosity ${\\lt}10^{45}\\, \\rm erg\\,s^{-1}$. In order to reconcile observations with predictions, these works have played with empirical relations for obscuring accreting black holes or with the efficiency of the seeding process (see e.g Degraf, Di Matteo & Springel 2010; Fanidakis et al. 2012; DeGraf & Sijacki 2020). Even though these works provide interesting results shedding light on the nature of low-luminous quasars, the treatment of seeding or dust obscuration is beyond the scope of this paper. As shown in Fig. 11, the fiducial model is compatible with current observations of the quasar LF, showing a sharp cut-off at larger luminosity (Shen et al. 2020). On the other hand, the models with higher gas accretion display a completely different behaviour. Boosting the gas accretion during DI leads to a larger excess of bright quasars at $z \\gt 1.0$. For instance, at $z \\sim 2$ and for luminosities ${\\gt}10^{46}\\, \\rm erg\\,s^{-1}$, the models with $A_{\\rm yr^{-1}} \\sim 1.92 \\times 10^{-15}$ and $A_{\\rm yr^{-1}} \\sim 2.67 \\times 10^{-15}$ are systematically overpredicting the number density by a factor of ${\\sim }1$ and ${\\sim }2\\, \\rm dex$, respectively. A similar behaviour is seen at $z \\sim 3$. At lower redshifts ($z \\lt 1.0$), the model follows both the fiducial results and the observed trends. This is principally caused by the decrease of important DIs events at these redshifts. Regarding the IM models, we can see similar trends at $z \\gt 2$, where the bright end of the LF is systematically larger than the observed one. We highlight that the difference is larger with $A_{\\rm yr^{-1}} \\sim 2.67 \\times 10^{-15}$. Interestingly, the excess with respect to the observations is smaller than with the IDI model. This is principally caused by the fact that DI events are more important than mergers at these redshifts (Izquierdo-Villalba et al. 2020). At lower redshifts, we can see larger differences with respect to the fiducial and the IDI models: IM model is systematically overprotecting the bright end of the LF (${\\gt}10^{46}\\, \\rm erg\\,s^{-1}$). Such differences can be a factor of 3 (1.5) by $z \\sim 0$ up to a factor of 5 (2) at $z \\sim 0.5$ for $A_{\\rm yr^{-1}} \\sim 1.92 \\times 10^{-15}$ ($A_{\\rm yr^{-1}} \\sim 2.67 \\times 10^{-15}$).","Citation Text":["Siana et al. 2008"],"Functions Text":["In this work, we only focus on the very bright objects, i.e. ${\\gt}10^{45}\\, \\rm erg\\,s^{-1}$, avoiding the comparison with lower luminosity given the current limitations on observational and theoretical models. In particular, from an observational standpoint, the covered area and depth of current surveys pose serious challenges when extracting statistical properties of the LF at the faint end"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[843,860]],"Functions Start End":[[445,841]]}
{"Identifier":"2021MNRAS.508.2743A__Shakura_&_Sunyaev_1973_Instance_1","Paragraph":"The gas particles are distributed such that the initial surface density profile has the power law:\n(7)$$\\begin{eqnarray*}\r\n\\Sigma _\\mathrm{g}=\\Sigma _{\\mathrm{g},0}\\left(\\frac{R}{R_{\\rm in}}\\right)^{-p},\r\n\\end{eqnarray*}$$where \u03a3g, 0 is a normalization constant. We choose a locally isothermal equation of state with a sound speed cs such that\n(8)$$\\begin{eqnarray*}\r\nc_\\mathrm{s} = c_\\mathrm{s,in} \\left(\\frac{R}{R_{\\rm in}} \\right)^{-q}.\r\n\\end{eqnarray*}$$where cs, in is the sound speed at the inner radius. We choose power law indices p = 1 and q = 0.5. We use a constant SPH artificial viscosity coefficient \u03b1SPH \u2248 0.55 such that the corresponding disc viscosity coefficient (Shakura & Sunyaev 1973) is \u03b1d = 0.01 at the inner radius. We note, however, that our choice of the indices p and q impact the radial dependence of the viscosity coefficient \u03b1d. The mapping between the SPH and disc viscosity coefficient is (e.g. Artymowicz & Lubow 1994, Murray 1996, Lodato & Pringle 2007)\n(9)$$\\begin{eqnarray*}\r\n\\alpha _\\mathrm{d} \\propto \\alpha _{\\rm SPH} \\frac{\\langle h \\rangle }{H}\r\n,\r\n\\end{eqnarray*}$$where \u3008h\u3009 is the vertically averaged smoothing length and H is the disc thickness, which depends on disc radius:\n(10)$$\\begin{eqnarray*}\r\nH = \\frac{c_\\mathrm{s}}{\\Omega _\\mathrm{k}} \\propto R^{3\/2-q}\r\n.\r\n\\end{eqnarray*}$$The average smoothing length in 3D has the radial dependence:\n(11)$$\\begin{eqnarray*}\r\n\\langle h \\rangle \\propto \\rho ^{-1\/3} \\propto \\left(\\frac{\\Sigma }{H}\\right)^{-1\/3} \\propto R^{1\/2+(p-q)\/3}\r\n.\r\n\\end{eqnarray*}$$Hence, the choice of p and q determines the radial dependence of the disc viscosity coefficient. A common choice is to use p = 3\/2 and q = 3\/4, which results in a radially constant \u03b1d (e.g. Lodato & Price 2010). In this paper, we decide to use a physically motivated choice of indices (e.g. Aly et al. 2018), even though it results in a non-constant disc viscosity coefficient. For our choice of p = 1 and q = 0.5, the resulting radial dependence of the disc viscosity coefficient is \u03b1d \u221d R\u22121\/3. To prevent particle interpenetration, we employ the recommended value \u03b2SPH = 2 for the quadratic viscosity coefficient (Price & Federrath 2010; Meru & Bate 2012).","Citation Text":["Shakura & Sunyaev 1973"],"Functions Text":["We use a constant SPH artificial viscosity coefficient \u03b1SPH \u2248 0.55 such that the corresponding disc viscosity coefficient","is \u03b1d = 0.01 at the inner radius."],"Functions Label":["Uses","Uses"],"Citation Start End":[[681,703]],"Functions Start End":[[558,679],[705,738]]}
{"Identifier":"2022ApJ...926...85S__Brogi_et_al._2016_Instance_1","Paragraph":"A prime example of the unique benefits of the intersection of the 3D nature of (ultra)hot exoplanet atmospheres and high-resolution spectroscopy lies in the ultrahot Jupiter WASP-76b, a gas giant orbiting an F7 star (West et al. 2016) that is well studied at lower resolution (Fu et al. 2017, 2021; Fisher & Heng 2018; Tsiaras et al. 2018; Edwards et al. 2020; von Essen et al. 2020). Using the high-resolution (R \u2248 138,000) ESPRESSO spectrograph on the Very Large Telescope (Pepe et al. 2010, 2013), the Ehrenreich et al. (2020) team was able to produce novel, high signal-to-noise ratio (S\/N), phase-resolved transmission spectra of this target across two separate transits. Curiously, an anomalous Doppler signature in the planet\u2019s transmission spectrum was detected: between 0 and \u22125 km s\u22121 at ingress, but roughly \u221211 km s\u22121 by egress. These speeds far exceed the few km s\u22121 planet-frame velocities detected on other planets (Snellen et al. 2010; Brogi et al. 2016). The detection team attributes this variable and strong blueshift to an asymmetric distribution of atomic iron in the planet\u2019s atmosphere, with a considerable amount of iron existing in the gas phase east of the substellar point, but cooling and condensing out as it makes its way to the much colder nightside. This asymmetry would cause a progressively blueshifted signal as the gas-phase iron region rotates into view over the course of the planet\u2019s transit. Ehrenreich et al. (2020) posit that their signal is composed of two independent Doppler components: solid-body rotation of \u00b15.3 km s\u22121  (the tidally locked equatorial velocity of the planet), and a uniform day\u2013night wind contributing an additional \u22125.3 km s\u22121 across both limbs, with gas-phase iron only present on the evening limb of the planet (Figure 1). Hence, the approaching limb would exhibit a blueshift from both rotation and winds totaling \u201310.6 km s\u22121, and the receding limb would produce no Doppler signal, as it would contain no gas-phase iron to absorb starlight.","Citation Text":["Brogi et al. 2016"],"Functions Text":["These speeds far exceed the few km s\u22121 planet-frame velocities detected on other planets"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[952,969]],"Functions Start End":[[841,929]]}
{"Identifier":"2018AandA...615A.148D__Sung_et_al._2013_Instance_2","Paragraph":"We study here the Sco OB1 association (Figs. 1 and 2), using this and other techniques. The general properties of this large OB association, which spans almost 5\u00b0 on the sky, and is surrounded by a ring-shaped HII region called Gum 55, are reviewed by Reipurth (2008). Its central cluster NGC 6231 contains several tens of OB stars, which have been extensively studied. On the other hand, many fewer studies, all recent, were devoted to the full mass spectrum, using optical photometry (Sung et al. 1998, 2013) and X-rays (Sana et al. 2006, 2007; Damiani et al. 2016; Kuhn et al. 2017a,b). The currently accepted distance of NGC 6231 is approximately 1580 pc, and its age is between 2and 8 Myr, with a significant intrinsic spread (Sung et al. 2013; Damiani et al. 2016). No ongoing star formation is known to occur therein, however. Approximately one degree North of the cluster, the loose cluster Trumpler 24 (Tr 24) also belongs to the association. There is little literature on this cluster (Seggewiss 1968; Heske & Wendker 1984, 1985; Fu et al. 2003, 2005) which unlike NGC 6231 lacks a well-defined center and covers about one square degree on the sky. Its age is  10 Myr according toHeske & Wendker (1984, 1985), who find several PMS stars, and its distance is 1570\u20131630 pc according to Seggewiss (1968). Other studies of the entire Sco OB1 association include MacConnell & Perry (1969 \u2013 H\u03b1-emission stars), Schild et al. (1969 \u2013 spectroscopy), Crawford et al. (1971 \u2013 photometry), Laval (Laval 1972a,b \u2013 gas and star kinematics, respectively), van Genderen et al. (1984 \u2013 Walraven photometry), and Perry et al.(1991 \u2013 photometry). At the northern extreme of Sco OB1, the partially obscured HII region G345.45+1.50 and its less obscured neighbor IC4628 were studied by Laval (1972a), Caswell & Haynes (1987), L\u00f3pez et al. (2011), and L\u00f3pez-Calder\u00f3n et al. (2016). They contain massive young stellar objects (YSOs; Mottram et al. 2007), maser sources (Avison et al. 2016), and the IRAS source 16562-3959 with its radio jet (Guzm\u00e1n et al. 2010), outflow (Guzm\u00e1n et al. 2011), and ionized wind (Guzm\u00e1n et al. 2014), and are therefore extremely young (1 Myr or less). The distance of G345.45+1.50 was estimated as 1.9 kpc by Caswell & Haynes (1987), and 1.7 kpc by L\u00f3pez et al. (2011), in fair agreement with distances of Sco OB1 stars. In Fig. 1 of Reipurth (2008) a strip of blue stars is visible, connecting NGC 6231 to the region of IC4628.","Citation Text":["Sung et al. 2013"],"Functions Text":["The currently accepted distance of NGC 6231 is approximately 1580 pc, and its age is between 2and 8 Myr, with a significant intrinsic spread"],"Functions Label":["Background"],"Citation Start End":[[732,748]],"Functions Start End":[[590,730]]}
{"Identifier":"2018MNRAS.473.1512A__Eichler_et_al._1989_Instance_1","Paragraph":"In an attempt to understand the radio properties of GRBs, Chandra & Frail (2012) conducted a complete investigation of all historical events observed in the radio domain. These included both of the main GRB populations (Kouveliotou et al. 1993): long-duration GRBs (likely produced by massive stellar collapse where the gamma-ray emission lasts for more than 2\u2009s; Woosley 1993; Kulkarni et al. 1998; Woosley & Bloom 2006) and short-duration GRBs (likely caused by the coalescence of two neutron stars or a neutron star and black hole, which lasts for less than 2\u2009s; Lattimer & Schramm 1976; Eichler et al. 1989; Narayan, Paczynski & Piran 1992). Only 30\u2009per\u2009cent of their sample had a detectable radio afterglow, with the radio emission peaking within a very narrow flux range. This led them to conclude that the low percentage of detections was likely due to the sensitivity of radio telescopes rather than there being two distinct GRB populations: radio-bright and radio-faint. Ghirlanda et al. (2013) and Burlon et al. (2015) then conducted simulations to demonstrate that potentially all Swift GRBs will be detectable at radio frequencies with phase 1 of the Square Kilometre Array (SKA), specifically SKA1-MID in Band 5 (\u223c9\u2009GHz)1 between 2 and 10 \u2009d post-burst, as well as with the recently upgraded Karl G. Jansky Very Large Array (VLA)2 and MeerKAT (the South African SKA precursor telescope; Jonas 2009). In fact, SKA1-MID will be so sensitive it could detect the radio counterparts from GRBs with gamma-ray emission up to five times fainter than those currently detected with Swift-BAT (note that these simulations do not account for radio emission produced by the reverse-shock, only considering contributions from the forward-shock; Burlon et al. 2015). However, a study conducted by Hancock, Gaensler & Murphy (2013), which involved visibility stacking of VLA GRB radio observations, suggested the low radio detection rate may be due to there being separate radio-bright and radio-faint GRB populations, and that \u226470\u2009per\u2009cent are likely to be truly radio bright.","Citation Text":["Eichler et al. 1989"],"Functions Text":["These included","and short-duration GRBs (likely caused by the coalescence of two neutron stars or a neutron star and black hole, which lasts for less than 2\u2009s;"],"Functions Label":["Background","Background"],"Citation Start End":[[591,610]],"Functions Start End":[[171,185],[422,565]]}
{"Identifier":"2018AandA...609A.131G__Heithausen_2012_Instance_2","Paragraph":"Moreover, there could also be some contribution to the detected temperature asymmetry from high-latitude gas clouds in our Galaxy along the line of sight toward M\u200981. In this respect we note that M\u200981 is at about 40.9\u00b0 north of the Galactic disk, where contamination from the Milky Way is expected to be low. However, interpretation of astronomical observations is often hampered by the lack of direct distance information. Indeed, it is often not easy to judge whether objects on the same line of sight are physically related or not. Since the discovery of the Arp\u2019s Loop (Arp 1965) the nature of the interstellar clouds in this region has been debated; in particular whether they are related to the tidal arms around the galaxy triplet (Sun et al. 2005; de Mello et al. 2008) or to Galactic foreground cirrus (Sollima et al. 2010; Davies et al. 2010). Already Sandage et al. (1976) presented evidence showing that we are observing the M\u200981 triplet through widespread Galactic foreground cirrus clouds and de Vries et al. 1987 built large-scale HI, CO, and dust maps that showed Galactic cirrus emission toward the M\u200981 region with NH \u2243 1\u22122  \u00d7  1020 cm-2. The technique used to distinguish between the emission from extragalactic or Galactic gas and dust relies on spectral measurements and on the identification of the line of sight velocities, which are expected to be different in each case. Unfortunately, in the case of the M\u200981 Group, this technique appears hardly applicable since the radial velocities of extragalactic and Galactic clouds share a similar LSR (local standard of rest) velocity range (Heithausen 2012). Several small-area molecular clouds (SAMS), that is, tiny molecular clouds in a region where the shielding of the interstellar radiation field is too low (so that these clouds cannot survive for a long time), have been detected by Heithausen (2002) toward the M\u200981 Group. More recently, data from the Spectral and Photometric Imaging Receiver (SPIRE) instrument onboard Herschel ESA space observatory and Multiband Imaging Photometer for Spitzer (MIPS) onboard Spitzer allowed the identification of several dust clouds north of the M\u200981 galaxy with a total hydrogen column density in the range 1.5\u20135 \u00d7 1020 cm-2 and dust temperatures between 13 and 17 K (Heithausen 2012). However, since there is no obvious difference among the individual clouds, there was no way to distinguish between Galactic or extragalactic origin although it is likely that some of the IR emission both toward M\u200981 and NGC 3077 is of Galactic origin. Temperature asymmetry studies in Planck data may be indicative of the bulk dynamics in the observed region provided that other Local (Galactic) contamination in the data is identified and subtracted. This is not always possible, as in the case of the M\u200981 Group, and therefore it would be important to identify and study other examples of dust clouds where their origin, either Galactic or extragalactic, is not clear. One such example might be provided by the interacting system toward NGC 4435\/4438 (Cortese et al. 2010) where the SAMS found appear more consistent with Galactic cirrus clouds than with extragalactic molecular complexes. Incidentally, the region A1 within R0.50 has been studied by Barker et al. (2009), who found evidence for the presence of an extended structural component beyond the M\u200981 optical disk, with a much flatter surface brightness profile, which might contain \u224310\u201315% of the M\u200981 total V-band luminosity. However, the lack of both a similar analysis in the other quadrants (and at larger distances from the M\u200981 center) and the study of the gas and dust component associated to this evolved stellar population, hamper our understanding of whether this component may explain the observed temperature asymmetry toward the M\u200981 halo. ","Citation Text":["Heithausen 2012"],"Functions Text":["More recently, data from the Spectral and Photometric Imaging Receiver (SPIRE) instrument onboard Herschel ESA space observatory and Multiband Imaging Photometer for Spitzer (MIPS) onboard Spitzer allowed the identification of several dust clouds north of the M\u200981 galaxy with a total hydrogen column density in the range 1.5\u20135 \u00d7 1020 cm-2 and dust temperatures between 13 and 17 K","However, since there is no obvious difference among the individual clouds, there was no way to distinguish between Galactic or extragalactic origin although it is likely that some of the IR emission both toward M\u200981 and NGC 3077 is of Galactic origin."],"Functions Label":["Background","Compare\/Contrast"],"Citation Start End":[[2282,2297]],"Functions Start End":[[1899,2280],[2300,2551]]}
{"Identifier":"2021ApJ...919L..23F__Tendulkar_et_al._2021_Instance_1","Paragraph":"Alongside studies of their emission properties, examining the environments of FRBs on subparcsec to kiloparsec scales can be equally informative. Thus far, only a fraction of known extragalactic repeating FRBs have been localized to host galaxies (Chatterjee et al. 2017; Tendulkar et al. 2017; Marcote et al. 2020; Macquart et al. 2020; Heintz et al. 2020; Bhandari et al. 2021; Bhardwaj et al. 2021a, 2021b; Li et al. 2021).17\n\n17\nFRB 20201124A is the fifth announced repeating FRB with a host galaxy. In total, there are eight such FRB sources known as of 2021 August.\n Much closer by, FRB-like emission has been detected from a Milky Way magnetar, SGR 1935 + 2154 (CHIME\/FRB Collaboration et al. 2020; Bochenek et al. 2020). All identified repeating FRB hosts have evidence for low to modest ongoing star formation rates of \u223c0.06\u22122 M\u2299 yr\u22121 (Gordon et al. 2004; Bhandari et al. 2020a; Heintz et al. 2020), several exhibit spiral arm morphologies (Bhardwaj et al. 2021a; Mannings et al. 2021; Tendulkar et al. 2021), and their stellar populations span a range of stellar masses, \u223c108\u22121010.5 M\u2299 (de Blok et al. 2008; Bhandari et al. 2020a; Heintz et al. 2020; Mannings et al. 2021). At face value, these characteristics, coupled with the absence of any quiescent host galaxy identifications for repeating FRBs, may indicate that FRBs are connected to host galaxies with ongoing star formation. However, studies of their more local environments reveal a rich diversity. For instance, the discovery of the repeating FRB 20200120E in an old (9.1 Gyr) globular cluster on the outskirts of the grand design spiral galaxy M81 (Bhardwaj et al. 2021a; Kirsten et al. 2021a) and the detection of the repeating FRB 20121102A embedded in a star-forming knot in its dwarf host galaxy (Bassa et al. 2017) seemingly represent polar opposite local environments. If all repeaters discovered to date originate from the same type of progenitor, then the progenitor model must accommodate the observed diversity of both local and galactic environments. Additionally, any connection in progenitors to the population of apparent nonrepeaters remains opaque.","Citation Text":["Tendulkar et al. 2021"],"Functions Text":["several exhibit spiral arm morphologies"],"Functions Label":["Background"],"Citation Start End":[[995,1016]],"Functions Start End":[[909,948]]}
{"Identifier":"2020AandA...634A..81B__Bleem_et_al._2015_Instance_1","Paragraph":"After checking the results of the U-net in the test area, we applied the same detection method to the full-sky SZ prediction map, with a detection threshold of pmax\u2004=\u20040.1 in order to recover the maximum number of Planck_z clusters. We detected 20 204 sources in the full-sky map with the U-net with pmax\u2004=\u20040.1. We compared the detections with the three catalogues of known galaxy clusters, Planck_z, Planck_no-z, and MCXCwP. Among the 20 204 detected sources, 98.5% of the Planck_z clusters are recovered, together with 76.4% of Planck_no-z clusters, and 20.8% of MCXCwP clusters. Moreover, 11 cluster are identified by ACT (Hasselfield et al. 2013) and 98 clusters are identified by the South Pole Telescope (SPT; Bleem et al. 2015), but are not included in the Planck PSZ2 catalogue. This means that 18 415 sources do not belong to any of the catalogues. We investigated the nature of the sources detected with the U-net. First, we cross-matched the sample of 18 415 sources with Planck point sources. Only 6.1% are matched within a cross-match radius of 5 arcmin with the positions of the Planck catalogue of galactic cold cores, and only 0.2% are matched with the positions of the Planck sources identified at 353 GHz. Second, we stacked at their positions 16 maps in different wavelengths, each of them potentially probing different galaxy cluster counterparts. Some of the 16 maps are also based on Planck data and thus are not independent, but some of the maps are independent and may show indications of galaxy cluster counterparts in other wavelengths, that is, in near-infra-red (where galaxies emit) and in X-rays (where the same gas emits as is detected with the SZ). The 16 maps are the Planck SZ MILCA map, the 6 Planck HFI frequency maps, the IRIS map at 100 \u03bcm (Miville-Desch\u00eanes & Lagache 2005), the CMB lensing map (based on Planck; Planck Collaboration VIII 2020), 4 galaxy over-density maps of all galaxies (called GAL ALL), passive galaxies (called GAL P), transitioning galaxies (called GAL T), and active galaxies (called GAL A) following Bonjean et al. (2019), star formation rate density maps (called SFR), a stellar mass density maps (called Mstar) constructed with the method from Bonjean et al. (2019), and finally, the ROSAT X-ray map (ByoPiC4 product). The Planck maps were masked from the Planck Catalogue of Compact Sources (PCCS, Planck Collaboration XXVI 2016), and the ROSAT map was masked from the point sources detected in ROSAT, Chandra, and XMM-Newton (Boller et al. 2016; Evans et al. 2010; Rosen et al. 2016, respectively). The result of the 16 stacks are shown in Fig. 5.","Citation Text":["Bleem et al. 2015"],"Functions Text":["98 clusters are identified by the South Pole Telescope (SPT;","but are not included in the Planck PSZ2 catalogue."],"Functions Label":["Uses","Differences"],"Citation Start End":[[715,732]],"Functions Start End":[[654,714],[735,785]]}
{"Identifier":"2022MNRAS.513.3458B__Robertson,_Massey_&_Eke_2017_Instance_1","Paragraph":"Among the most viable mechanisms of cusp-core transformation that require changes to the assumed cosmogony is one that was proposed specifically as a possible solution to the cusp-core problem. It proposes that the DM is in fact not collisionless but self-interacting (SIDM; Spergel & Steinhardt 2000; Yoshida et al. 2000; Dav\u00e9 et al. 2001; Col\u00edn et al. 2002; Vogelsberger, Zavala & Loeb 2012; Rocha et al. 2013; see Tulin & Yu 2018 for a review). In SIDM, particles can exchange energy and momentum through elastic scattering, causing an outside-in energy redistribution within the centre of DM haloes, resulting in the formation of an isothermal core. The time-scale on which an initially cuspy SIDM halo forms a flat and isothermal core is roughly given by the time it takes for each DM particle in the inner halo to scatter at least once (Vogelsberger et al. 2012; Rocha et al. 2013). The strength of the self-interaction in SIDM models is parametrized in terms of the momentum transfer cross-section per unit mass, \u03c3T\/m\u03c7. Depending on the specific SIDM model, \u03c3T\/m\u03c7 can either be constant or dependent on the relative velocity between the two scattering DM particles. SIDM is an efficient mechanism of cusp-core transformation in dwarf-size haloes for $\\sigma _T\/m_\\chi \\gtrsim 1\\, {\\rm cm^2g^{-1}}$, whereas SIDM haloes are virtually indistinguishable from CDM haloes if $\\sigma _T\/m_\\chi \\lesssim 0.1\\, {\\rm cm^2g^{-1}}$ (Zavala, Vogelsberger & Walker 2013). The most stringent and precise constraints on the self-interaction cross-section have been put on the scales of galaxy clusters (e.g. Robertson, Massey & Eke 2017; Robertson et al. 2019) and large elliptical galaxies (Peter et al. 2013), where observations require that $\\sigma _T\/m_chi \\lesssim 1\\, {\\rm cm^2g^{-1}}$. On smaller scales, Read, Walker & Steger (2018) concluded that $\\sigma _T\/m_\\chi \\lesssim 0.6\\, {\\rm cm^2g^{-1}}$, based on their findings that the central density profile of the MW dwarf spheroidal galaxy Draco is cuspy (see also the SIDM results of Valli & Yu 2018). Moreover, based on a DM only analysis of the updated too-big-to-fail problem, Zavala et al. (2019) concluded that SIDM models with a constant cross-section of $\\sigma _T\/m_\\chi \\sim 1\\, {\\rm cm^2g^{-1}}$ fail to explain the apparently large central densities of the host haloes of the ultra-faint satellites of the MW (Errani, Pe\u00f1arrubia & Walker 2018). It should be pointed out that the constraints on \u03c3T\/m\u03c7 on the scale of dwarf galaxies are affected by significantly larger systematic uncertainties than on the scales of galaxy clusters or elliptical galaxies. Moreover, Zavala et al. (2019) demonstrate that SIDM with a strongly velocity-dependent self-interaction cross-section may provide a natural explanation for the observed diversity in the rotation curves of the MW dwarf spheroidals (see also Correa 2021). The strong dependence of the self-interaction cross-section on the typical DM velocities would create a bimodal distribution of rotation curves in the MW satellites in which the heavier haloes have constant density cores while the lighter haloes have undergone gravothermal collapse and have very steep central cusps as a consequence. The same mechanism of gravothermal collapse might be accelerated by tidal interactions in the environment of the MW leading to an agreement between constant cross-section SIDM models with $\\sigma _T\/m_\\chi \\sim 3\\, {\\rm cm^2g^{-1}}$ and the internal kinematics of MW satellites (e.g. Kahlhoefer et al. 2019; Sameie et al. 2020).","Citation Text":["Robertson, Massey & Eke 2017"],"Functions Text":["The most stringent and precise constraints on the self-interaction cross-section have been put on the scales of galaxy clusters (e.g.","where observations require that $\\sigma _T\/m_chi \\lesssim 1\\, {\\rm cm^2g^{-1}}$."],"Functions Label":["Background","Background"],"Citation Start End":[[1600,1628]],"Functions Start End":[[1466,1599],[1704,1784]]}
{"Identifier":"2017MNRAS.465..383R__Archibald_et_al._2016_Instance_1","Paragraph":"In previous work (Rogers & Safi-Harb 2016, hereafter RSH16), we studied a parametrized phenomenological model for describing magnetic field growth in neutron star (NS) evolution. Assuming that the external dipole field is buried by an intense process of fall-back accretion after the formation of the NS (Muslimov & Page 1995, 1996), the slow diffusion from the stellar surface exerts a time-dependent torque on the NS (Geppert, Page & Zannias 1999). This process may provide an explanation for the observed braking indices of young pulsars with n \u2260 3 (Espinoza 2012; Archibald et al. 2016), in contrast to the prediction of a rotating magnetic dipole in vacuum (n = 3; Ostriker & Gunn 1969). Magnetic field growth provides an evolutionary link between apparently disparate classes of NS, such as the extremely weak field1 subset of NSs known as central compact objects (CCOs), high magnetic field pulsars (HBPs) and the X-ray dim isolated NSs (XDINSs). On the other hand, magnetic field decay has been invoked to describe the evolution of the anomalous X-ray pulsars (AXPs) and soft gamma repeaters (SGRs). These objects are conventionally described by the magnetar model that posits the development of large magnetic fields via dynamo action from rapid rotation early in the life of the NS, necessary to support the core of the massive progenitor from collapse (Duncan & Thompson 1992; Akiyama et al. 2003; Thompson, Quataert & Burrows 2005). However, the distinction between the apparently rotation-powered HBPs, with dipole fields just above the quantum critical limit, BQED = 4.4 \u00d7 1013 G, and the magnetically powered AXPs and SGRs (B \u2265 BQED) is significantly blurred after the HBPs were observed displaying magnetar-like activity (Gavriil et al. 2008; Kumar & Safi-Harb 2008; G\u00f6\u011f\u00fcs et al. 2016), and radio emission was observed from magnetars (Camilo et al. 2007). Moreover, the SNRs associated with the AXPs and SGRs show evidence of \u2018typical\u2019 explosion energies (\u22641051 erg) and not superenergetic as one would expect from a rapidly rotating proto-NS (Vink & Kuiper 2006; Kumar, Safi-Harb & Gonzalez 2012; Safi-Harb & Kumar 2013; Kumar et al. 2014). These complications are difficult to reconcile with the standard magnetar picture.","Citation Text":["Archibald et al. 2016"],"Functions Text":["This process may provide an explanation for the observed braking indices of young pulsars with n \u2260 3","in contrast to the prediction of a rotating magnetic dipole in vacuum"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[568,589]],"Functions Start End":[[451,551],[592,661]]}
{"Identifier":"2018ApJ...869...12S__Duc_et_al._2015_Instance_1","Paragraph":"To verify the predictions of CDM, there have been recent attempts to compare the amount of stellar mass observed in the outskirts of galaxies with the mass fraction in accreted stars predicted by simulations (e.g., Font et al. 2008; Pillepich et al. 2014; Merritt et al. 2016; D\u2019Souza & Bell 2018; Elias et al. 2018; Huang et al. 2018). For the purposes of this paper, we will use the term \u201cstellar halo\u201d in an observational sense, referring to the faint structure in the outskirts of galaxies beyond the central concentration of stellar mass. In practice, there are many observational definitions of this term that can include radial profile, surface brightness, or metallicity characteristics. Most recent observational attempts to characterize the stellar halos of galaxies, including the MW (e.g., Carollo et al. 2010) and the Andromeda galaxy (M31; e.g., Courteau et al. 2011), have used analysis of resolved stellar populations to identify the accreted component, usually by searching for an old, metal-poor population extending far from the central galaxy (e.g., Seth et al. 2007; Cockcroft et al. 2013). However, this method requires extremely deep images, mainly obtained using the Hubble Space Telescope (with the exception of Greggio et al. 2014), and has therefore been limited to a small handful of galaxies so far. Observing stellar halos in integrated light is, in principle, more easily scalable to the sample sizes needed to explore the wide variation in the accreted component that is predicted by simulations (e.g., Bakos & Trujillo 2012; D\u2019Souza et al. 2014; Duc et al. 2015; Merritt et al. 2016; Huang et al. 2018), presuming that it is possible to account for the contribution of scattered light (de Jong 2008; Slater et al. 2009; Sandin 2014). However, the lack of resolved stellar population information makes it far more challenging to identify the regions of an image dominated by accreted material. So far, no work has attempted to account for how the method used to select the stellar halo from a galaxy observed in integrated light may bias the comparison to simulations, where the provenance of material is perfectly known and a variety of definitions of \u201cstellar halo\u201d are imposed. Despite efforts such as that in Rodriguez-Gomez et al. (2016) to understand whether the mass in the stellar halo comes primarily from accreted material or from stars formed in the central galaxy (which in this work we call \u201cformed in situ\u201d) and expelled to the halo, and its dependence with separation from the central galaxy, it is not straightforward to apply these results to the spatial selections in projection that are commonly used in integrated-light images. In fact, most prior work has focused on comparisons between observed galaxies and predictions for the stellar halo based on dark matter\u2013only (DM-only) simulations tagged with stars, where the stellar halo is by definition 100% accreted. However, in simulations that include baryonic physics, both of these distinct channels are observed to contribute to the stellar halos of galaxies (Font et al. 2011; Cooper et al. 2013; Tissera et al. 2013; Pillepich et al. 2015; Angl\u00e9s-Alc\u00e1zar et al. 2017; G\u00f3mez et al. 2017a), and both are interesting for what they tell us about the process of galaxy formation, as well as the cosmology in which galaxies are formed.","Citation Text":["Duc et al. 2015"],"Functions Text":["Observing stellar halos in integrated light is, in principle, more easily scalable to the sample sizes needed to explore the wide variation in the accreted component that is predicted by simulations (e.g.,"],"Functions Label":["Motivation"],"Citation Start End":[[1579,1594]],"Functions Start End":[[1329,1534]]}
{"Identifier":"2019MNRAS.484..892R__Wolf_et_al._2009_Instance_1","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"],"Functions Text":["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,"],"Functions Label":["Uses"],"Citation Start End":[[907,923]],"Functions Start End":[[690,906]]}
{"Identifier":"2021AandA...655A.111K__Rojas-Arriagada_et_al._(2019)_Instance_2","Paragraph":"Over the last decade, the radial and vertical dependences of the metallicity-alpha-element distribution have been studied in more and more detail with increasingly larger samples (e.g., Bensby et al. 2011; Anders et al. 2014; Nidever et al. 2014; Hayden et al. 2015; Queiroz et al. 2020). Figure 6 is mostly consistent with similar plots shown in the above papers. In the inner 10 kpc, it displays two over-densities, a high alpha-element (here [Mg\/Fe]), and a low one. Between Rg\u2004=\u20046 and 10 kpc, the two over-densities define two different sequences. In Appendix E, we note that when the sample is restricted to a \u00b1500 pc layer around the Galactic plane, two close but separated sequences are observed in the Rg\u2004\u2208\u2004[4,\u20066] kpc interval. Because of their scale height (Bovy et al. 2012), kinematics (Bensby et al. 2003), and age properties (Haywood et al. 2013), these two sequences are associated with the thick disc (high-alpha) and thin disc (low-alpha), respectively. Moving inward of Rg\u2004=\u20044\u2005\u2212\u20056 kpc, Fig. 6 shows that the two over-densities connect through a zone of lower density to form a single sequence. This is in agreement with the observations of Hayden et al. (2015), Bensby et al. (2017), Zasowski et al. (2019), Bovy et al. (2019), and Lian et al. (2020a,b), who also report a single sequence in the inner disc and\/or in the bulge\/bar area. Conversely, Rojas-Arriagada et al. (2019) and Queiroz et al. (2020) observe two sequences in the inner regions. In Appendix F, we compare the distributions of different APOGEE DR16 alpha elements in the ([Fe\/H], [\u03b1\/Fe]) plane (restricting the sample to the stars contained in the Rg\u2004\u2208\u2004[0,\u20062] kpc interval). The different elements produce different patterns: the global alpha-element abundance5 and oxygen show a double sequence, while magnesium, silicon, and calcium present a single sequence. This could explain, at least partly, why Queiroz et al. (2020), who use a combined \u03b1-element abundance, observe a double sequence, while we see a single one with magnesium. However, this does not explain the discrepancy with Rojas-Arriagada et al. (2019), who also used magnesium. Beyond Rg\u2004=\u200410 kpc, the high-alpha sequence gradually vanishes. This is in agreement with the finding that the thick disc has a shorter scale length than the thin disc (Bensby et al. 2011; Cheng et al. 2012; Bovy et al. 2012). It should be emphasised that in this paragraph the term \u2018sequence\u2019 is used in the geometrical sense. It does not presuppose the number of chemical tracks that form the sequence or sequences. In particular, based on Fig. 6, it can not be excluded that the single geometrical sequence observed in the inner disc be made of two chemical tracks, with the low-alpha one restricted to a narrow metallicity range. We discuss and propose an interpretation of the inner disc sequence in Sect. 5.","Citation Text":["Rojas-Arriagada et al. (2019)"],"Functions Text":["However, this does not explain the discrepancy with","who also used magnesium."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[2073,2102]],"Functions Start End":[[2021,2072],[2104,2128]]}
{"Identifier":"2020ApJ...898...92C__Occhiogrosso_et_al._2011_Instance_1","Paragraph":"The first interstellar discovery of MF took place in 1975 from the microwave emission spectrum of Sgr B2, which agrees with the rotational constants of the syn isomer (Curl 1959; Brown et al. 1975; Nummelin et al. 2000). Subsequently, it was detected in several other sources such as comet Hale\u2013Bopp, and protostars and hot corinos (e.g., in Orion A and the protoplanetary nebula CRL 618) (Ellder et al. 1980; Ikeda et al. 2001; Cazaux et al. 2003; Bottinelli et al. 2004a, 2004b; Remijan et al. 2005, 2006; Remijan & Hollis 2006). Later, the less stable trans rotamer was detected in both the laboratory and the interstellar medium (ISM; Ruschin & Bauer 1980; Blom & G\u00fcnthard 1981; M\u00fcller et al. 1983; Neill et al. 2011, 2012). The detailed mechanistic route for the interstellar synthesis of MF is a subject of considerable debate (Horn et al. 2004; Herbst 2005; Garrod & Herbst 2006; Snyder 2006; Bennett & Kaiser 2007; Occhiogrosso et al. 2011; Lawson et al. 2012). The initially assumed gas-phase model involving an ion\u2013molecule reaction of protonated methanol (\n\n\n\n\n\n) with formaldehyde (H2CO) fails to reproduce the large MF column density, because this pathway demands a significant activation barrier (\u223c128 kJ mol\u22121) for one of the intermediate steps (Horn et al. 2004). Instead, gas-phase radiative association between the methyl cation (\n\n\n\n\n\n) and formic acid (HCOOH) seems energetically promising to produce H+MF and thereafter neutral MF via dissociative recombination with electrons:\n1\n\n\n\n\n\n\n\n2\n\n\n\n\n\nHowever, this reaction scheme fails to reproduce the huge MF concentration observed in hot cores (Horn et al. 2004). Alternative routes suggested more recently involve ion\u2013molecule reactions with a low barrier or no barrier between protonated methanol and formic acid (acid-catalyzed Fischer esterification) or between protonated formic acid and methanol (\n\n\n\n\n\n transfer):\n3\n\n\n\n\n\n\n\n4\n\n\n\n\n\nfollowed by reaction (2) or exothermic proton transfer of \n\n\n\n\n\n (H+MF) to a base with a higher proton affinity than MF (Ehrenfreund & Charnley 2000; Neill et al. 2011, 2012).","Citation Text":["Occhiogrosso et al. 2011"],"Functions Text":["The detailed mechanistic route for the interstellar synthesis of MF is a subject of considerable debate"],"Functions Label":["Background"],"Citation Start End":[[923,947]],"Functions Start End":[[729,832]]}
{"Identifier":"2021ApJ...923...59V__Feruglio_et_al._2015_Instance_1","Paragraph":"We compute the emission line luminosity using the following equation:\n1\n\n\n\nLCO\u2032=3.25\u00d7107SCO\u0394vDL2(1+z)3\u03bdobs2Kkms\u22121pc2,\n\nwhere \u03bd\nobs is the observed CO transition frequency, D\n\nL\n is the luminosity distance, and S\nCO\u0394v is the line-integrated flux in units of Jy km s\u22121. We convert the observed CO transition luminosity into CO (1\u20130) luminosity (L\n\n\n\n\nCO(1\u22120)\u2032\n\n) by assuming that the low-J CO transitions are thermalized and optically thick, so \n\n\n\nLCO4\u20133\u2032=LCO3\u20132\u2032=LCO1\u20130\u2032\n\n. Using the ratios (\n\n\n\nrJ1=LCOJ\u2192J\u22121\u2032\/LCO1\u20130\u2032\n\n) from Carilli & Walter (2013) with r\n31=0.97 and r\n41 = 0.87 did not significantly change our results. Furthermore, in 3C 298, we found that the molecular gas is consistent with being thermalized and optically thick (Vayner et al. 2017). These physical conditions are consistent with what is found for Mrk 231 (Feruglio et al. 2015). Finally, we convert the \n\n\n\nLCO1\u20130\u2032\n\n line luminosity into molecular gas mass using the CO-to-H2 conversion factor: \u03b1\nCO with units of (K km s\u22121pc2)\u22121. For sources where we do not detect any narrow CO emission at the systemic redshift, we place a limit on the molecular gas mass over an aperture equal to the beam size for an emission line with a velocity FWHM of 250 km s\u22121. The molecular gas mass limits can be linearly scaled with a different \u03b1\nCO value. For sources with detected molecular gas at the quasar\u2019s systemic redshift, we compute the radius of the molecular gas region, which allows us to measure the gas surface density. All radii are computed using a curve-of-growth method. Effective radii refer to a region that encloses 50% of the flux, while \u201cmaximum\u201d extent refers to a size scale that encloses 90% of the flux. In all cases, narrow emission at the systemic redshift of the quasar is spatially resolved by our observations. We deconvolve the size of the beam from all radius measurements. For sources with no detected CO emission, we use the molecular gas mass limit and the beam of the observations as a proxy for the radius. Values associated with the molecular gas at the systemic redshift are summarized in Table 3.","Citation Text":["Feruglio et al. 2015"],"Functions Text":["These physical conditions are consistent with what is found for Mrk 231"],"Functions Label":["Similarities"],"Citation Start End":[[831,851]],"Functions Start End":[[758,829]]}
{"Identifier":"2016ApJ...829...53V__Piconcelli_et_al._2005_Instance_1","Paragraph":"Two spectroscopically confirmed X-ray AGNs in the cluster core (#607, 661 in G13) are suitable candidates for ionizing the nebula. The depth of the new Chandra observation, coupled with an optimal on-axis alignment, allowed us to perform a basic X-ray spectral analysis despite the limited photon statistics (34 and 20 net counts in the observed 0.5\u20137 keV band for sources #607 and 661, respectively). Source #607 is characterized by a power-law spectrum with photon index \u0393 = 2.0 \u00b1 0.6; the observed 2\u201310 keV flux is 1.7+1.1\u22120.6 \u00d7 10\u221215 erg cm\u22122 s\u22121, corresponding to a rest-frame 2\u201310 keV luminosity of \n\n\n\n\n\n erg s\u22121, typical of a luminous Seyfert galaxy. The X-ray spectrum of source #661, the point-like Ly\u03b1 emitter (Figure 2), is flat: fitting the data with a power-law model provides \n\n\n\n\n\n, highly indicative of strong obscuration. We then included an absorption component and fixed the photon index to 1.8, as expected for the intrinsic AGN emission (e.g., Piconcelli et al. 2005). This model results in a column density of \n\n\n\n\n\n cm\u22122, i.e., consistent with marginal Compton-thick absorption (1.5 \u00d7 1024 cm\u22122). The tentative detection of an iron K\u03b1 emission line at 6.4 keV (with equivalent width of \u22482.4 keV rest frame), if confirmed, would further support the heavily obscured nature of source #661. The derived 2\u201310 keV flux is (7.4 \u00b1 2.2) \u00d7 10\u221215 erg cm\u22122 s\u22121, corresponding to a rest-frame luminosity of \n\n\n\n\n\n erg s\u22121, placing source #661 in the quasar regime. We do not detect any bright counterpart in deep Jansky Very Large Array observations at 3 GHz down to 2.7 \u03bcJy (rms), and we thus classify source #661 as radio-quiet. From aperture photometry, we estimated a Ly\u03b1 flux of \n\n\n\n\n\n erg cm\u22122 s\u22121, corresponding to a luminosity of (1.9 \u00b1 0.2) \u00d7 1042 erg s\u22121. The spectral energy distribution (SED) of #661 is shown in Figure 7. From SED modeling, which benefits from near-, mid-, and far-IR observations from Spitzer and Herschel, we estimated a bolometric luminosity for the AGN of (2.7 \u00b1 1.5) \u00d7 1045 erg s\u22121. A similar value (3.2 \u00b1 0.6 \u00d7 1045 erg s\u22121) is derived using the observed [O iii]\u03bb5007 \u00c5 luminosity obtained from recent Subaru\/MOIRCS spectroscopy of the galaxy (Valentino et al. 2015), converted into a bolometric luminosity as \n\n\n\n\n\n (Heckman et al. 2004). Assuming the luminosity-dependent bolometric correction as in Lusso et al. (2012), we predict an intrinsic 2\u201310 keV luminosity for source #661 of \n\n\n\n\n\n erg s\u22121. This value is consistent, within the uncertainties due to the adopted relations and measurements, with that derived from the X-ray spectral analysis reported above.","Citation Text":["Piconcelli et al. 2005"],"Functions Text":["We then included an absorption component and fixed the photon index to 1.8, as expected for the intrinsic AGN emission (e.g.,"],"Functions Label":["Uses"],"Citation Start End":[[966,988]],"Functions Start End":[[840,965]]}
{"Identifier":"2018AandA...619A..13V__Saviane_et_al._2012_Instance_4","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)"],"Functions Text":["For internal consistency, all EWs in this work have been adjusted to the measurement scale of"],"Functions Label":["Uses"],"Citation Start End":[[2136,2157]],"Functions Start End":[[2042,2135]]}
{"Identifier":"2018ApJ...856...19N__Fischer_et_al._2010_Instance_1","Paragraph":"Models of the ECSN progenitor cores suggest the onset of the electron-capture instability occurs at a unique ONeMg core mass in the mass range of 1.366\u20131.377 M\u2299. (Miyaji et al. 1980; Nomoto 1984, 1987; Podsiadlowski et al. 2005; Takahashi et al. 2013). Electron captures cause the core to contract, and O and Ne burning is ignited in the central regions and propagates outwards in a deflagration front (Schwab et al. 2015), processing material to nuclear statistical equilibrium, where further electron captures and photdissociation accelerates the collapse (Miyaji et al. 1980; Nomoto 1987; Takahashi et al. 2013). Whether the core collapses or the deflagration disrupts the core depends sensitively on the ignition density (Isern et al. 1991; Jones et al. 2016). If the core does collapse, the explosion proceeds via delayed explosion on short timescales (Mayle & Wilson 1988; Kitaura et al. 2006; Fischer et al. 2010), and 2D simulations suggest the explosion occurs before significant convection has had time to develop (Wanajo et al. 2011) and hence a symmetric explosion results. This, coupled with the steep density gradient at the core surface, leads to very little mass loss from the core; estimates of mass loss include of order 10\u22123 M\u2299 (Podsiadlowski et al. 2005), 10\u22122 M\u2299 (Kitaura et al. 2006), and 1.39 \u00d7 10\u22122 M\u2299 (1.14 \u00d7 10\u22122 M\u2299) for the 1D (2D) models of Wanajo et al. (2009, 2011). Therefore the ONeMg progenitor core mass is a good estimate of the baryon mass MB of the resulting NS (Podsiadlowski et al. 2005). Indeed, PSR J0737-3039A and the companion to PSR J1756-2251 have gravitational masses consistent with baryon masses \u223c1.37 M\u2299 when their gravitational binding energies are taken into account (Lattimer & Yahil 1989). Population synthesis calculations incorporating the various binary evolution channels that might lead to production of NSs via ECSNe show that J0737-3039B most likely formed in an ECSN, and the companion to PSR J1756-2251 is consistent with such a formation scenario (Andrews et al. 2015). Other systems with candidates for ECSNe formation also exist (Keith et al. 2009; Chen et al. 2011).","Citation Text":["Fischer et al. 2010"],"Functions Text":["If the core does collapse, the explosion proceeds via delayed explosion on short timescales","and 2D simulations suggest the explosion occurs before significant convection has had time to develop","and hence a symmetric explosion results."],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[900,919]],"Functions Start End":[[765,856],[922,1023],[1045,1085]]}
{"Identifier":"2018AandA...615A.155H__Chou_et_al._2007_Instance_1","Paragraph":"While the velocity field is very regular, the velocity dispersion looks more complex. Instead of one peak extended along a southeast-northwest axis (Figs. 10, and 11), we find three (Fig. 13). The central one, elongated along the main axis of the galaxy and not perpendicular to it, is accompanied by two additional regions with wider lines to the northeast and southwest, separated by a total of 5.\u2032\u2032 9 \u00b1 0.\u2032\u20322. At these locations we observe the transition between rigid body rotation and steep rotation curve to a flat one (see also Chou et al. 2007). At a radius of 2.\u2032\u203245 (\u224845 pc) and with a rotation velocity of 140 km s\u22121 we obtain with Eq. (1) of Mauersberger et al. (1996; \u03b7 = 1) an enclosed mass of M2.45 = 2.1 \u00d7 108 M\u2299 with an estimated error of 10%, in good agreement with Cunningham & Whiteoak (2005). For comparison, Roy et al. (2010) obtain with the H92\u03b1 line 3 \u00d7 107 M\u2299 for the mass inside a radius of 1\u2033 (\u224819 pc). Furthermore, we note that the nuclear region is slightly lopsided: The center of the line connecting the two outer peaks of line width is located \u22481.\u2032\u20324 southwest of the continuum peak (see Table 3). Systemic velocities (Vbarycentric \u2248 571 km s\u22121) are found about 0.\u2032\u20329 (\u224815\u201320 pc) southwest of the continuum peak and may be closer to the position of the maser disk (Greenhill et al. 1997) given in Sect. 3.1. While uncertainties with respect to the value of the systemic velocity (Sect. 2) and the relative positions of the continuum peak and the maser disk are significant, most of the star formation represented by the continuum emission (see Bendo et al. 2016) appears to originate slightly northeast of the dynamical center. The lopsidedness of the central region of NGC 4945 also explains why molecular lines with significant absorption show integrated intensity peaks slightly shifted to the southwest: With the bulk of the continuum arising from the northeast, line emission is likely more quenched by absorption at this side of the center than in the southwestern part of the inclined nuclear disk (see Sect. 4.1.2 and Fig. 7).","Citation Text":["Chou et al. 2007"],"Functions Text":["While the velocity field is very regular, the velocity dispersion looks more complex. Instead of one peak extended along a southeast-northwest axis (Figs. 10, and 11), we find three (Fig. 13). The central one, elongated along the main axis of the galaxy and not perpendicular to it, is accompanied by two additional regions with wider lines to the northeast and southwest, separated by a total of 5.\u2032\u2032 9 \u00b1 0.\u2032\u20322. At these locations we observe the transition between rigid body rotation and steep rotation curve to a flat one (see also"],"Functions Label":["Similarities"],"Citation Start End":[[535,551]],"Functions Start End":[[0,534]]}
{"Identifier":"2021MNRAS.507.4389G__Masters_et_al._2011_Instance_2","Paragraph":"Erwin (2018) showed that, in a sample drawn from the Spitzer Survey of Stellar Structure in Galaxies (S4G), the bar fraction is constant over a range of (g \u2212r) colours and gas fractions. Their bar fraction does not increase, but rather decreases for stellar masses higher than \u223c 109.7M\u2299. These results are in contrast to many SDSS-based studies cited above. Erwin (2018) argues that this apparent contradiction can be explained if SDSS-based studies miss bars in low-mass blue galaxies. In Figs 5 and 6, we showed that the newly detected bars in GZD (compared to GZ2) are weak bars in low-mass blue galaxies. Nevertheless, the \u2018combined\u2019 bar fraction in Fig. 6 is not constant over (g \u2212r) colour and agrees well with Masters et al. (2011) for redder colours [(g \u2212r) colour > 0.5]. Additionally, our \u2018combined\u2019 bar fraction remains roughly constant over stellar mass. As mentioned before, we conclude that strong bars drive the trends of bar fraction with (g \u2212r) colour, stellar mass, and SFR observed in other studies (Nair & Abraham 2010b; Masters et al. 2011, 2012; Vera et al. 2016; Cervantes Sodi 2017). However, the addition of weak bars in low-mass blue galaxies is insufficient to resolve the apparent disagreement between Erwin (2018) and many SDSS-based studies (Masters et al. 2011, 2012; Vera et al. 2016; Cervantes Sodi 2017; Kruk et al. 2018), which instead seems likely to be due to the very different sample selection of the S4G and SDSS galaxy samples. For example, the median stellar mass of the sample used in Erwin (2018) is \u223c109.6M\u2299 (based on their Fig. 4 and the bins in the top left-hand panel of their Fig. 5). However, the median stellar mass of our sample is 1010.6M\u2299. As stellar mass correlates with many parameters (including bar length), this can have major consequences. Additionally, as Erwin (2018) notes, there is also the issue of resolution to consider. With an r-band FWHM of 1.18 arcsec from DECaLS (Dey et al. 2019) and a mean redshift of 0.036, the mean linear resolution of our sample is approximately 834 pc, which is higher than the 165 pc of Erwin (2018). This explains why they observe many sub-kpc bars, while we do not. These differences in stellar mass and resolution will manifest themselves in the conclusions, so a more detailed analysis is needed for a proper comparison with Erwin (2018).","Citation Text":["Masters et al. 2011"],"Functions Text":["As mentioned before, we conclude that strong bars drive the trends of bar fraction with (g \u2212r) colour, stellar mass, and SFR observed in other studies"],"Functions Label":["Extends"],"Citation Start End":[[1041,1060]],"Functions Start End":[[867,1017]]}
{"Identifier":"2019ApJ...883..149A__Rodriguez_et_al._2015_Instance_1","Paragraph":"A key question that remains unanswered is how BBHs are formed. Viable formation channels include isolated binary evolution (e.g., Bethe & Brown 1998; Belczynski et al. 2002, 2014, 2016; Dominik et al. 2013; Mennekens & Vanbeveren 2014; Spera et al. 2015; Eldridge & Stanway 2016; Mandel & de Mink 2016; Marchant et al. 2016; Mapelli et al. 2017; Stevenson et al. 2017; Barrett et al. 2018; Giacobbo & Mapelli 2018; Kruckow et al. 2018; Mapelli & Giacobbo 2018) and dynamical encounters in dense stellar environments, such as globular clusters (e.g., Portegies Zwart & McMillan 2000; O\u2019Leary et al. 2006; Sadowski et al. 2008; Downing et al. 2010, 2011; Rodriguez et al. 2015, 2016a, 2016b; Askar et al. 2017; Fragione & Kocsis 2018; Rodriguez & Loeb 2018; Samsing 2018; Samsing et al. 2018; Zevin et al. 2019), young star clusters (e.g., Banerjee et al. 2010; Ziosi et al. 2014; Mapelli 2016; Banerjee 2017, 2018; Di Carlo et al. 2019; Kumamoto et al. 2019), and galactic nuclei (e.g., O\u2019Leary et al. 2009; Antonini & Perets 2012; Antonini & Rasio 2016; Petrovich & Antonini 2017; Stone et al. 2017a, 2017b; Rasskazov & Kocsis 2019). Moreover, the dynamical process known as Kozai\u2013Lidov (KL) resonance (Kozai 1962; Lidov 1962) can trigger the merger of a BBH, even if the BBH has not been formed in a dense star cluster. In fact, if the BBH is orbited by a tertiary body (i.e., the BBH is the inner binary of a stable hierarchical triple system), the KL mechanism triggers oscillations of the BBH\u2019s eccentricity, which might speed up the merger by gravitational-wave emission. Each channel is expected to produce black hole mergers with different mass and spin distributions (Mandel & O\u2019Shaughnessy 2010; Abbott et al. 2016a; Rodriguez et al. 2016c; Farr et al. 2017; Abbott et al. 2019d). The limited statistics from the low number of systems detected through gravitational waves and model uncertainties so far do not allow strong constraints on the formation channels.","Citation Text":["Rodriguez et al. 2015"],"Functions Text":["A key question that remains unanswered is how BBHs are formed. Viable formation channels include","and dynamical encounters in dense stellar environments, such as globular clusters (e.g.,"],"Functions Label":["Background","Background"],"Citation Start End":[[653,674]],"Functions Start End":[[0,96],[461,549]]}
{"Identifier":"2015MNRAS.450.2749G__Brinchmann_et_al._2004_Instance_1","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"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[824,846]],"Functions Start End":[[643,822]]}
{"Identifier":"2015AandA...584A..75V__Essen_et_al._(2014)_Instance_8","Paragraph":"The data presented here comprise quasi-simultaneous observations during secondary eclipse of WASP-33 b around the V and Y bands. The predicted planet-star flux ratio in the V-band is 0.2 ppt, four times lower than the accuracy of our measurements. Therefore, we can neglect the planet imprint and use this band to measure the stellar pulsations, and most specifically to tune their current phases (see phase shifts in von Essen et al. 2014). Particularly, our model for the light contribution of the stellar pulsations consists of eight sinusoidal pulsation frequencies with corresponding amplitudes and phases. Hence, to reduce the number of 24 free parameters and the values they can take, we use prior knowledge about the pulsation spectrum of the star that was acquired during von Essen et al. (2014). As the frequency resolution is 1\/\u0394T (Kurtz 1983), 3.5 h of data are not sufficient to determine the pulsations frequencies. Therefore, during our fitting procedure we use the frequencies determined in von Essen et al. (2014) as starting values plus their derived errors as Gaussian priors. As pointed out in von Essen et al. (2014), we found clear evidences of pulsation phase variability with a maximum change of 2 \u00d7 10-3 c\/d. In other words, as an example after one year time a phase-constant model would appear to have the correct shape with respect to the pulsation pattern of the star, but shifted several minutes in time. To account for this, the eight phases were considered as fitting parameters. The von Essen et al. (2014) photometric follow-up started in August, 2010, and ended in October, 2012, coinciding with these LBT data. We then used the phases determined in von Essen et al. (2014) during our last observing season as starting values, and we restricted them to the limiting cases derived in Sect. 3.5 of von Essen et al. (2014), rather than allowing them to take arbitrary values. The pulsation amplitudes in \u03b4 Scuti stars are expected to be wavelength-dependent (see e.g. Daszy\u0144ska-Daszkiewicz 2008). Our follow-up campaign and these data were acquired in the blue wavelength range. Therefore the amplitudes estimated in von Essen et al. (2014), listed in Table 1, are used in this work as fixed values. This approach would be incorrect if the pulsation amplitudes would be significantly variable (see e.g., Breger et al. 2005). Nonetheless, the short time span of LBT data, and the achieved photometric precision compared to the intrinsically low values of WASP-33\u2019s amplitudes, make the detection of any amplitude variability impossible. ","Citation Text":["von Essen et al. (2014)"],"Functions Text":["Therefore the amplitudes estimated in","listed in Table 1, are used in this work as fixed values."],"Functions Label":["Uses","Uses"],"Citation Start End":[[2148,2171]],"Functions Start End":[[2110,2147],[2173,2230]]}
{"Identifier":"2016AandA...588A..44Y__Jones_et_al._2014_Instance_1","Paragraph":"The second issue concerns the fact that inside a given region, coreshine is not detected in all the dense clouds observed by Paladini (2014) and Lef\u00e8vre et al. (2014) and that the proportion of clouds exhibiting coreshine varies from one region to another. For instance, 75% of the dense clouds detected in Taurus exhibit coreshine, whereas in most other regions the proportion is closer to 50% (such as Cepheus, Chamaeleon, and Musca)5. On the contrary, there are for instance very few detections in the Orion region. In THEMIS, most of the scattering efficiency originates in the accretion of an a-C:H mantle. This leads to three possible explanations for the absence of detectable coreshine. The first explanation is related to the amount of carbon available in the gas phase. The abundance used by K\u00f6hler et al. (2015) relies on the highest C depletion measurements made by Parvathi et al. (2012) towards regions with \\hbox{$N_{\\rm H} \\geqslant 2 \\times 10^{21}$}NH\u2a7e 2 \u00d7 1021 H\/cm2. Parvathi et al. (2012) highlighted the variability in the carbon depletion in dust depending on the line of sight. Thus, there may be clouds were the amount of carbon available for a-C:H mantle formation is smaller or even close to zero: such regions would be populated with aggregates with a thinner H-rich carbon mantle or no second mantle at all and thus exhibit very little or no coreshine emission. A second explanation is related to the stability of H-rich carbon in the ISM, which depends strongly on the radiation field intensity to local density ratio (Godard et al. 2011; Jones et al. 2014). In low-density regions (according to Jones et al. 2014, \\hbox{$A_{V} \\leqslant 0.7$}AV\u2a7d 0.7 for the standard ISRF), UV photons are responsible for causing the photo-dissociation of CH bonds, a-C:H \u2192 a-C. In transition regions from diffuse ISM to dense clouds (Jones et al. 2014, \\hbox{$0.7 \\leqslant A_{V} \\leqslant 1.2$}0.7 \u2a7d AV\u2a7d 1.2 for the standard ISRF), better shielded from UV photons and where the amount of hydrogen is significantly higher, H-poor carbon can be transformed into H-rich carbon through H atom incorporation, a-C \u2192 a-C:H. Similarly, carbon accreted from the gas phase in these transition regions is likely to be and stay H-rich. Then, in the dense molecular clouds, most of the hydrogen is in molecular form and thus not available to produce a-C:H mantles on the grains. However, this approximately matches the density at which ice mantles start to accrete on the grains, which would partly protect a-C:H layers that had formed earlier (Godard et al. 2011, and references therein). The stability and hydrogenation degree of a-C:H, as well as the exact values of AV thresholds, are both dependent on the timescale and UV field intensity. The resulting a-C \u2194 a-C:H delicate balance could explain why in a quiet region such as Taurus most of the clouds exhibit coreshine, whereas in Orion, where on average the radiation field intensity and hardness are much higher, most clouds do not. A third explanation is related to the age and\/or density of the clouds. In a young cloud, where dust growth is not advanced, or in an intermediate density cloud (\u03c1C ~ a few 103 H\/cm3), the dust population may be dominated by CMM grains instead of AMM(I) dust. Such clouds would be as bright in the IRAC 8 \u03bcm band as in the two IRAC bands at 3.6 and 4.5 \u03bcm, thus not matching the selection criteria defined by Pagani et al. (2010) and Lef\u00e8vre et al. (2014) and would be classified as \u201cno coreshine\" clouds. ","Citation Text":["Jones et al. 2014"],"Functions Text":["A second explanation is related to the stability of H-rich carbon in the ISM, which depends strongly on the radiation field intensity to local density ratio"],"Functions Label":["Background"],"Citation Start End":[[1569,1586]],"Functions Start End":[[1391,1547]]}
{"Identifier":"2017AandA...600A.138C__Wakeford_et_al._(2016)_Instance_1","Paragraph":"Gibson (2014) proposed that marginalization over many systematics is more robust than simple model selection, and that the BIC-based model selection could be the worst criterion in their experiments. To assess the impact of BIC-based model selection choices (hereafter Method 1) on our derived transmission spectrum, we also performed a separate analysis on the spectroscopic light curves employing the systematics marginalization approach (hereafter Method 2). We followed the implication of this approach described in Wakeford et al. (2016), and refer the reader to that work for more details. Instead of using the wavelet-based MCMC to account for the correlated noise (see Method 1), for simplicity, Method 2 employed the MPFIT package to fit the data and accounted for the correlated noise using the time-averaging \u03b2 approach. We calculated the marginal likelihood for all the systematics models using the Akaike information criterion (AIC; Akaike 1973) as the approximation, that is, \\hbox{$\\ln\\mathcal{P}(D|S_q)\\approx-\\mathrm{AIC}\/2$}ln\ud835\udcab(D | Sq) \u2248 \u2212AIC\/2, which provides more adequate fits and performs better than BIC as suggested by Gibson (2014). The resulting Rp\/R\u22c6 in each spectroscopic channel was then calculated as the marginal-likelihood-weighted average of the best-fitting values from all the systematics models, whose uncertainty was propagated from both the deviation from the weighted average and the best-fitting error bar for each systematics model. The derived Rp\/R\u22c6 values are also listed in Table 3. The middle and bottom panels of Fig. 5 show the transmission spectra derived by Method 2 and the comparison between these two methods, respectively. The great consistency confirms that the BIC-based model selection in this work does not bias the derived transmission spectrum. Since the two methods have almost the same transit-depth values in any given spectral channel, we decide to present the results from Method 2 in the following discussion, as they have smaller error bars. ","Citation Text":["Wakeford et al. (2016)"],"Functions Text":["We followed the implication of this approach described in","and refer the reader to that work for more details."],"Functions Label":["Uses","Uses"],"Citation Start End":[[520,542]],"Functions Start End":[[462,519],[544,595]]}
{"Identifier":"2019ApJ...872...52C__Ranjan_et_al._2017_Instance_1","Paragraph":"Studies aiming at measuring and modeling solar radiation and its variability are strongly motivated by the impact that solar irradiance (that is, the electromagnetic energy emitted by the Sun received at the top of Earth\u2019s atmosphere in units of area and time), especially in the UV, has on the chemistry and physical properties of Earth\u2019s atmosphere and climate (e.g., Gray et al. 2010; Matthes et al. 2017). Studies of solar variability have been recently also driven by the necessity of improving our understanding of stellar variability (see Fabbian et al. 2017, for a recent review), which, in turn, is essential to characterize the habitable zones of stars and the atmospheres of their exoplanets. As for Earth, modeling of exoplanet atmospheres requires as fundamental input the spectral energy distribution of the hosting star, especially UV and shorter wavelengths (e.g., Tian et al. 2014; Ranjan et al. 2017; Rugheimer & Kaltenegger 2018). Unfortunately, measurements of UV radiation are strongly hampered by the interstellar medium absorption (up to 70%\u201390%), which is significant even for relatively close stars, so that estimates of stellar UV radiation strongly rely on modeling (see Linsky 2017, for a recent review). Moreover, because there is no mission scheduled in the near future to observe stellar spectra in the UV, after the Hubble Space Telescope ceases operations, the characterization of UV spectra of stars hosting exoplanets that will be discovered by current and future missions (e.g., TESS or James Webb Space Telescope) will necessarily rely on indirect estimates, performed, for instance, through the use of semi-empirical models (e.g., Mauas et al. 1997; Fontenla et al. 2016; Bus\u00e1 et al. 2017) or proxies (e.g., Stelzer et al. 2013; Shkolnik et al. 2014). Stellar irradiance variability also affects the detectability of exoplanets (see, e.g., the recent review by Oshagh 2018). The passage over the disk of spots and faculae may induce photometric variations of amplitude similar to or larger than photometric variations induced by planetary transits. Moreover, the presence of active regions may alter spectral line profiles, thus hindering exoplanet detections performed through radial velocity measurements. Similarly, spectroscopic techniques that allow us to estimate the physical properties of exoplanet atmospheres (see, e.g., Kreidberg 2017, for a recent view) require as fundamental input the spectra synthetized through models representing quiet and active regions (faculae and sunspots). Finally, stellar irradiance variability observed at different spectral ranges, especially in the UV, is a fundamental observable for the characterization of the magnetic activity of a star, and therefore for the understanding of dynamo processes in stellar objects (e.g., Reinhold et al. 2013; Basri 2016; Salabert et al. 2016). Because stellar photometric and spectral variability can be modeled using the semi-empirical approaches developed for the Sun described above (see, e.g., Shapiro et al. 2016; Witzke et al. 2018), understanding the limitations of current irradiance models is fundamental to improving our capability of modeling stellar variability.","Citation Text":["Ranjan et al. 2017"],"Functions Text":["As for Earth, modeling of exoplanet atmospheres requires as fundamental input the spectral energy distribution of the hosting star, especially UV and shorter wavelengths (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[899,917]],"Functions Start End":[[704,880]]}
{"Identifier":"2019ApJ...876L..28D__Lamb_et_al._2018_Instance_1","Paragraph":"In Figures 1 and 2 we show that the X-ray (1.7 keV5\n\n5\nThis value corresponds to the geometric mean of the XRT energy band, at which the error of the estimated flux can be reasonably suppressed.\n), optical (R-band), and radio (6 GHz) fluxes varied with the time of observation applied to the proper corrections if observed at a distance of 200 Mpc, motivated by the fact that the averaged sensitive range of the Advanced LIGO\/Virgo detectors in their full-sensitivity run is about 210 Mpc, for the current samples. Due to the faintness of the SGRB afterglow emission, there are gaps of the data between the previous more distant events and GW170817\/GRB 170817A (please note that for the latter we only consider the quick decline phase as the early part is significantly influenced by the beam effect of the off-axis outflow). Therefore we extrapolate the very late (t > 200 day) X-ray and optical afterglow data of GRB 170817A to t \u223c 2 day after the burst and then compare them to other events. The radio to X-ray spectrum of the forward shock afterglow emission of GW170817\/GRB 170817A is f\u03bd \u221d \u03bd\u22120.6, which yields a p = 2.2 in the slow-cooling synchrotron radiation scenario (Lamb et al. 2018; Troja et al. 2018). In the jet model, such a p can also reasonably account for the very late flux decline of \n\n\n\n\n\n (Lamb et al. 2018). The extrapolation function of the forward shock emission of GRB 170817A to early times is thus taken as f \u221d t\u22122.2. Surprisingly, the forward shock afterglow emission of GW170817\/GRB 170817A, the first neutron star merger event detected by Advanced LIGO\/Virgo, are among the brightest ones for all SGRBs detected so far. Just a few events have X-ray afterglow emission brighter than that of GRB 170817A, as demonstrated in the right panels of Figure 1. The same conclusion holds for the optical and radio afterglow data, as shown in Figure 2, though these two samples are rather limited. We have also compared the distribution of the isotropic gamma-ray energy Eiso, calculated in the rest-frame energy band of 1\u2013104 keV, for the SGRBs with well-measured spectra, and found no significant difference for the SGRBs with and without \u201clong-lasting\u201d afterglow emission (see Figure 3; where the number of events for the X-ray sample are smaller than that presented in Figure 1 because some bursts lack reliable spectral measurements). GRB170817A and GRB 150101B (Troja et al. 2018), two short events with the weakest detected prompt emission, have \u201cbright\u201d late-time afterglow emission because of their off-axis nature.","Citation Text":["Lamb et al. 2018"],"Functions Text":["The radio to X-ray spectrum of the forward shock afterglow emission of GW170817\/GRB 170817A is f\u03bd \u221d \u03bd\u22120.6, which yields a p = 2.2 in the slow-cooling synchrotron radiation scenario"],"Functions Label":["Uses"],"Citation Start End":[[1177,1193]],"Functions Start End":[[995,1175]]}
{"Identifier":"2021MNRAS.504.2168G__Steiner_et_al._2011_Instance_2","Paragraph":"Finally, we attempt to characterize the reflection component using the full 2\u201335\u2009keV spectra with a sophisticated model [M4: ${\\tt{\\rm constant}}$*${\\tt{\\rm TBabs}}$*(${\\tt{\\rm simplr}}$*${\\tt {\\rm kerrbb2}}$+${\\tt{\\rm kerrconv}}$*(${\\tt{\\rm ireflect}}$*${\\tt{\\rm simplc}}$)), to evaluate the impact on the spin measurement and understand the changes of the accretion flow and the interaction between the disc and the corona. This model features a self-consistent treatment of the thermal, Compton scattering and the reflection component: ${\\tt {\\rm kerrbb2}}$ describes the thermal component and supplies the seed photons for ${\\tt{\\rm simplr}}$ (a modified version of ${\\tt{\\rm simpl}}$, Steiner et al. 2011) to generate the Compton component; while a portion of the Compton component will escape to reach an observer, the remains (refer as ${\\tt{\\rm simplc}}$, Steiner et al. 2011) will strike back to the disc to generate the reflected component. The reflection fraction Rref in ${\\tt{\\rm ireflect}}$ (Magdziarz & Zdziarski 1995), defined as the ratio of the Compton photons striking back to the disc to that escaping to infinity, is restricted to negative value thereby only the reflected component is returned by ${\\tt{\\rm ireflect}}$. It is linked to the reflection constant parameter x in ${\\tt{\\rm simplr}}$ via the relation x = 1 + |Rref| (Gou et al. 2011). We set the elemental abundance to unity and the iron abundance AFe to five times the solar abundance (Bharali et al. 2019; Buisson et al. 2019; Xu et al. 2020). The disc temperature Tin is fixed at the value returned by ${\\tt{\\rm diskbb}}$ (M1, refer to Gou et al. 2011). The ionization parameter \u03be is fixed at 1000 (i.e. log(\u03be) = 3, Xu et al. 2020; Buisson et al. 2019), as it is difficult to be constrained. Finally we use ${\\tt{\\rm kerrconv}}$ (Brenneman & Reynolds 2006) to apply relativistic effects assuming an unbroken emissivity profile with index q = 3. The key parameters in ${\\tt {\\rm kerrbb2}}$ and ${\\tt{\\rm kerrconv}}$ are linked together.","Citation Text":["Steiner et al. 2011"],"Functions Text":["while a portion of the Compton component will escape to reach an observer, the remains (refer as ${\\tt{\\rm simplc}}$,","will strike back to the disc to generate the reflected component."],"Functions Label":["Uses","Uses"],"Citation Start End":[[864,883]],"Functions Start End":[[746,863],[885,950]]}
{"Identifier":"2015ApJ...807..148G__Rees_&_M\u00e9sz\u00e1ros_1994_Instance_2","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"],"Functions Text":["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;"],"Functions Label":["Background"],"Citation Start End":[[1038,1058]],"Functions Start End":[[858,1037]]}
{"Identifier":"2021AandA...655A..12T__Tang_et_al._2017b_Instance_6","Paragraph":"Using the RADEX3 non local thermodynamic equilibrium (LTE) modeling program (van der Tak et al. 2007) with collisional rate coefficients from Wiesenfeld & Faure (2013), we modeled the relation between the gas kinetic temperature and the measured average of para-H2CO 0.5 \u00d7 [(322\u2013221 + 321\u2013220)\/303\u2013202] ratios, adopting a 2.73 K background temperature, an average observational linewidth of 4.0 km s\u22121, and column densities N(para-H2CO) = 2.7 \u00d7 1012 and 3.7 \u00d7 1012 cm\u22122 for N113 and N159W, respectively. The results are shown in Fig. 5. The values of the para-H2CO column density were obtained with APEX data (beam size ~30\u2033; Tang et al. 2017b), which cover similar regions. Different column densities of para-H2CO only weakly affect derived kinetic temperatures (see Fig. 3 in Tang et al. 2017b or Fig. 4 in Tang et al. 2018a; this was also shown in Fig. 13 and discussed in Sect. 4.3.1 of Mangum & Wootten 1993) as long as all lines are optically thin. Considering that the relation between the gas temperature and the para-H2CO line ratio may vary at different spatial densities (see Fig. 2 in Tang et al. 2017b), we modeled it at spatial densities 104, 105, and 106 cm\u22123 in Fig. 5. It appears that Tkin at n(H2) = 105 cm\u22123 is consistently lower than values at 104 and 106 cm\u22123 by \u227223% and \u227234%, respectively, for Tkin \u2272 100 K. Local thermodynamic equilibrium (LTE) is a good approximation for the H2CO level populations under optically thin and high-density conditions (Mangum & Wootten 1993; Tang et al. 2017a,b, 2018b). Following the method applied by Tang et al. (2017b) in their Eq. (2), we plot the relation between the LTE kinetic temperature, TLTE, and the para-H2CO (3\u20132) line ratio in Fig. 5. Apparently, TLTE agrees well with Tnon-LTE at volume densities n(H2) ~ 105 cm\u22123 as long as Tkin \u2272 100 K. Previous observations show that para-H2CO (3\u20132) is sensitive to gas temperature at density 105 cm\u22123 (Ginsburg et al. 2016; Immer et al. 2016; Tang et al. 2017b). The spatial density measured with para-H2CO (303\u2013202) and C18O (2\u20131) in N113 and N159W is n(H2) ~ 105 cm\u22123 on a size of ~30\u2033 (Tang et al. 2017b). Therefore, here we adopt 105 cm\u22123 as an averaged spatial gas density in the N113 and N159W regions.","Citation Text":["Tang et al. 2017b"],"Functions Text":["Previous observations show that para-H2CO (3\u20132) is sensitive to gas temperature at density 105 cm\u22123"],"Functions Label":["Uses"],"Citation Start End":[[1953,1970]],"Functions Start End":[[1811,1910]]}
{"Identifier":"2021MNRAS.500.3083C__Lupi_&_Bovino_2020_Instance_2","Paragraph":"Previous theorethical studies have outlined that the [C\u2009ii] emission originates from the cold (with temperatures of a few 100 K) neutral medium and from photo-dissociation regions (PDR, Vallini et al. 2013). This seems to suggest that its presence closely traces star formation sites, resulting in a linear relation, as found by De Looze et al. (2014) and Herrera-Camus et al. (2015). While at low-redshift and close to solar metallicity such a relation is well established, as shown by several observations (De Looze et al. 2014; Herrera-Camus et al. 2018) and also numerical simulations (see, e.g. Lupi & Bovino 2020), significant deviations can arise in different ISM conditions, like at lower metallicity or in presence of a strong ionizations field, that are more typically found in the high-redshift Universe. To address the impact of such conditions, several studies have analysed the [C\u2009ii] emission from typical high-redshift galaxies, by post-processing hydrodynamic zoom-in cosmological simulations with cloudy (Ferland et al. 2017; see, e.g. Olsen et al. 2017; Pallottini et al. 2017, 2019; Katz et al. 2019), or via ad-hoc methods, as in Arata et al. (2020), or also via on-the-fly non-equilibrium chemistry (Lupi et al. 2020). The main conclusion in all these studies is that a [C\u2009ii] deficit exists at high redshift, most likely due to the starbursting nature of these galaxies rather than their metallicity, since most of these systems are close to solar (see, e.g. Vallini et al. 2015; Lupi & Bovino 2020). Other studies have evidenced a weak dependence of [C\u2009ii] on metallicity. Harikane et al. (2020) showed that the L[C\u2009ii]\/SFR ratio does not show a strong dependence on metallicity, which to first approximation was interpreted as the result of the proportionality between C abundance and metallicity Z and an inverse proportionality between PDR column density and Z in a dust-dominated shielding regime (Kaufman, Wolfir & Hollenbach 2006). In this framework, if PDR give a large conribution to the [C\u2009ii] emissivity, the [C\u2009ii] luminosity is not expected to strongly depend on Z (see also Ferrara et al. 2019; Pallottini et al. 2019).","Citation Text":["Lupi & Bovino 2020"],"Functions Text":["The main conclusion in all these studies is that a [C\u2009ii] deficit exists at high redshift, most likely due to the starbursting nature of these galaxies rather than their metallicity, since most of these systems are close to solar (see, e.g."],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1503,1521]],"Functions Start End":[[1241,1481]]}
{"Identifier":"2020MNRAS.494.6012W__Middleton_et_al._2015a_Instance_1","Paragraph":"Walton et al. (2017) suggested that the broad-band spectral variability seen in Holmberg\u2009IX X-1, similar to that reported here for NGC\u20091313 X-1, could potentially be related to the presence of the expected funnel-like geometry for the inner accretion flow. In such a scenario, the funnel is expected to geometrically collimate the emission from the innermost regions within the funnel (discussed further in Section 4.2). Regardless of the nature of the accretor (black hole or neutron star), the highest energy emission probed by NuSTAR is usually expected to arise from these regions, either powered by a centrally located Compton-scattering corona (e.g. Reis & Miller 2013), or a centrally located accretion column. The stability of this emission would therefore imply that any geometrical collimation it experiences remains roughly constant, despite the change in observed broad-band X-ray flux (which would suggest a change in accretion rate, $\\dot{M}$). In principle, an increase in accretion rate would be expected to result in an increase in the scale height of the funnel (e.g. King 2008; Middleton et al. 2015a). However, while this must happen over some range of $\\dot{M}$ in order for the disc structure to transition from the thin disc expected for standard sub-Eddington accretion to the funnel-like geometry expected for super-Eddington accretion, as discussed by Lasota et al. (2016), once the disc reaches the point of being fully advection-dominated the opening angle of the disc should tend to a constant (H\/R \u223c 1, where H is the scale height of the disc at radius R). Walton et al. (2017) speculated that once this occurs, rather than closing the funnel further, an increase in $\\dot{M}$ instead simply increases the characteristic radius within which geometric beaming occurs, such that emission that is already within this region (the highest energies probed) experiences no further collimation with an increase in $\\dot{M}$, while emission from larger radii (i.e. from more intermediate energies) does still become progressively more focused, and would exhibit stronger variability. In essence, this idea invokes a radially dependent beaming factor in which the beaming of the innermost regions has saturated to explain (in only a qualitative sense) the unusual, energy-dependent broad-band spectral evolution seen from Holmberg\u2009IX X-I (and now NGC\u20091313 X-1).","Citation Text":["Middleton et al. 2015a"],"Functions Text":["In principle, an increase in accretion rate would be expected to result in an increase in the scale height of the funnel (e.g.","However, while this must happen over some range of $\\dot{M}$ in order for the disc structure to transition from the thin disc expected for standard sub-Eddington accretion to the funnel-like geometry expected for super-Eddington accretion, as discussed by Lasota et al. (2016), once the disc reaches the point of being fully advection-dominated the opening angle of the disc should tend to a constant (H\/R \u223c 1, where H is the scale height of the disc at radius R)."],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[1097,1119]],"Functions Start End":[[959,1085],[1122,1586]]}
{"Identifier":"2021MNRAS.507.2115M__Hoyle_2016_Instance_1","Paragraph":"In astrophysics, the number of studies that apply ML techniques has risen substantially in the last years. Unsupervised learning algorithms have been used to identify different kinematic components of simulated galaxies (Obreja et al. 2018, 2019), to compare stellar spectra (Traven et al. 2017), to classify pulsars (Lee et al. 2012), and to find high-redshift quasars (Polsterer, Zinn & Gieseke 2013). Supervised learning has been used to classify variable stars (Richards et al. 2011), to classify galaxies morphologically (Huertas-Company et al. 2008), and to determine the redshift of galaxies (Hoyle et al. 2015; Hoyle 2016; D\u2019Isanto & Polsterer 2018). Recently, ML has also been used to connect the properties of galaxies and dark matter haloes using supervised learning techniques. Kamdar, Turk & Brunner (2016a), Kamdar, Turk & Brunner (2016b) use tree-based methods to predict several galaxy properties from a set of halo properties and train the models on galaxy catalogues obtained from semi-analytic models and the Illustris hydrodynamic simulation (Vogelsberger et al. 2014). Sullivan, Iliev & Dixon (2018) train a simple neural network with one hidden layer to predict the baryon fraction within a dark matter halo at high redshift, given several halo properties (features). As training data they use the results of a cosmological hydrodynamic simulation with Ramses-RT. Similarly, Agarwal, Dav\u00e9 & Bassett (2018) use several ML methods to link input halo properties to galaxy properties, training on the Mufasa cosmological hydrodynamical simulation. Taking a reverse approach, Calderon & Berlind (2019) train tree-based methods and a neural network to derive halo mass from galaxy properties, training on an SDSS group catalogue. The limitation of all these studies is the supervised training and the training data. As labelled galaxy-halo data is not available for observed systems, the data for supervised learning has to be taken from a model. Even if the ML algorithms learn to reproduce the training data perfectly, the connection between galaxy and halo properties is the same as in the simulations. If the simulations predict the true relations poorly, so will the ML method. Therefore ML algorithms should not be trained on simulated data, but on observed data directly.","Citation Text":["Hoyle 2016"],"Functions Text":["Supervised learning has been used to","and to determine the redshift of galaxies"],"Functions Label":["Background","Background"],"Citation Start End":[[619,629]],"Functions Start End":[[404,440],[557,598]]}
{"Identifier":"2020AandA...641A.126B__Lyutikov_et_al._2005_Instance_1","Paragraph":"Many low-luminosity active galactic nuclei (LLAGN) display prominent jets and compact cores that are sources of highly nonthermal continuum radio emission (see, e.g., Heeschen 1970; Wrobel & Heeschen 1991). The observational signatures of the compact cores have been reproduced using models that produce self-absorbed synchrotron emission in the jet (Falcke & Biermann 1995; Falcke et al. 2004) or in a magnetized accretion flow (Narayan et al. 1998; Yuan et al. 2003; Broderick & Loeb 2006; Moscibrodzka et al. 2009; Dexter et al. 2009; see also Falcke et al. 2001). This radiation is emitted by relativistic electrons gyrating around magnetic field lines. In the optically thin limit, the emission is significantly polarized (Jones & Hardee 1979), an effect that has been observed in higher-luminosity AGN sources (Gabuzda et al. 1996; Gabuzda & Cawthorne 2000; Lyutikov et al. 2005). The polarized emission from an accreting AGN can therefore yield information about the magnetic-field morphology of the source, which may be crucial to the evolution of the accretion flow of the AGN. The Event Horizon Telescope (EHT) is a worldwide millimeter-wavelength array capable of resolving the black-hole shadow (Goddi et al. 2017; Event Horizon Telescope Collaboration 2019); this is a characteristic feature of the radio-frequency emission from optically thin AGN at the scale of the event horizon (Falcke et al. 2000; Broderick & Narayan 2006), although the black-hole shadow may be obscured or exaggerated in certain accretion scenarios (see Gralla et al. 2019 and Narayan et al. 2019). The EHT can also determine the polarization state of such emission: Johnson et al. (2015) report 1.3 mm observations (230 GHz) that indicate partially ordered magnetic fields within a region of about six Schwarzschild radii around the event horizon of Sagittarius A* (Sgr A*), the supermassive black hole in the center of the Milky Way. Bower et al. (2003) reported stable long-term behavior and short-term variability in Sgr A* rotation measure, implying a complex inner region (within 10 Schwarzschild radii) in which both emission and propagation effects are important to the observed polarization. Hada et al. (2016) studied the central black hole in the galaxy M 87, and observed a bright feature with (linear) polarization degree of 0.2 at 86 GHz at the jet base. Observations in infrared by Gravity Collaboration (2018) were consistent with a model in which a relativistic \u201chot spot\u201d of material, orbiting near the innermost stable circular orbit (ISCO) of Sgr A* in a poloidal magnetic field, emits polarized synchrotron radiation.","Citation Text":["Lyutikov et al. 2005"],"Functions Text":["In the optically thin limit, the emission is significantly polarized",", an effect that has been observed in higher-luminosity AGN sources"],"Functions Label":["Background","Background"],"Citation Start End":[[864,884]],"Functions Start End":[[658,726],[748,815]]}
{"Identifier":"2019ApJ...873...32M__Deming_et_al._2007_Instance_1","Paragraph":"Inspired by these works, we investigated the possibility of introducing a color diagram for the characterization of irradiated planets using their effective temperature instead of their absolute magnitude. The effective temperature can be used as a proxy for the luminosity\/absolute magnitude because the reference radius is assumed to be constant; see Section 3. To take an even more practical approach, we chose a normalized color parameter based on the Spitzer Infrared Array Camera (IRAC; Allen et al. 2004; Fazio et al. 2004) as a commonly used photometer for the observation of exoplanets (see, e.g., Charbonneau et al. 2005; Deming et al. 2007; Todorov et al. 2009; Sing et al. 2016). The IRAC photometric channels 1 and 2 are centered at 3.6 and 4.5 \u03bcm, respectively. Channel 1 (3.6 \u03bcm) is more suited to studying CH4\/H2O spectral features, while channel 2 (4.5 \u03bcm) is more sensitive to CO\/CO2 features (see, e.g., D\u00e9sert et al. 2009; Swain et al. 2009a, 2009b). Depending on the type of spectroscopy, i.e., transmission or emission, the ratio of the transit depth or the ratio of the secondary eclipse depth at these channels could potentially provide information regarding the relative presence of these molecules in the atmosphere of a planet. For transmission spectroscopy, we define this ratio as\n15\n\n\n\n\n\nwhere IRAC\u03bb is the transition depth observed at the wavelength \u03bb (\u03bcm) channel. In transmission spectroscopy, absorption features appear as positive signals in the unit of transit depth. This is not the case for emission spectroscopy, where absorption features are negative signals with respect to a blackbody curve. As a result, we rearrange the terms and define this ratio of channels for emission spectroscopy as\n16\n\n\n\n\n\nwhere IRAC\u03bb is the secondary eclipse depth observed at the wavelength \u03bb (\u03bcm) channel. By applying IRAC\u2019s spectral response curves to our 56,448 synthetic spectra, we estimated these ratios for the two spectroscopy methods. Figure 10 shows the IRAC synthetic color\u2013temperature diagrams for cloud-free atmospheres under equilibrium chemistry conditions.","Citation Text":["Deming et al. 2007"],"Functions Text":["To take an even more practical approach, we chose a normalized color parameter based on the Spitzer Infrared Array Camera","as a commonly used photometer for the observation of exoplanets (see, e.g.,"],"Functions Label":["Uses","Uses"],"Citation Start End":[[632,650]],"Functions Start End":[[364,485],[531,606]]}
{"Identifier":"2022ApJ...938...92B__Favier_et_al._2010_Instance_1","Paragraph":"Flows with \n\n\n\nRem\u226a1\n\n and \n\n\n\nN\u223c\ue23b(1)\n\n have the distinct property that the induced magnetic field is quickly diffused away, yet the Lorentz force is not negligible. This limit is referred to as the quasi-static approximation to MHD (which we call \u201cQMHD\u201d henceforth; Moffatt 1967; Sommeria & Moreau 1982; Davidson 1995; Knaepen & Moreau 2008; Davidson 2013), and has been studied mainly in metallurgy and in MHD experiments due to the typically low conductivity of liquid metals (Alemany et al. 1979; Sommeria 1988; Gallet et al. 2009; Klein & Poth\u00e9rat 2010; Poth\u00e9rat & Klein 2014; Baker et al. 2018), although recent numerical studies on its turbulent properties and anisotropy have been done as well (Zikanov & Thess 1998; Burattini et al. 2008; Favier et al. 2010, 2011; Reddy & Verma 2014; Verma 2017). After nondimensionalizing the equations of MHD using the uniform density \u03c1, \u2113, and u, and taking the limits above, one is left with a single dynamical equation for the velocity\n2\n\n\n\n\u2202v\u2202t+v\u00b7\u2207v=\u2212\u2207p*\u2212Ro\u22121x\u02c6\u2225\u03a9\u00d7v\u2212N\u2207\u22122(x\u02c6\u2225B0\u00b7\u2207)2v+F,\n\nwhere p\n* is the total pressure modified by rotation and magnetic pressure, Ro\n\u22121 \u2261 2\u03a9\u2113\/u is the inverse Rossby number (quantifying the relative strength of the Coriolis force), \n\n\n\nx\u02c6\u2225\u03a9\n\n and \n\n\n\nx\u02c6\u2225B0\n\n are unit vectors in the direction of rotation and the background magnetic field, respectively, and \nF\n is a generic forcing term that can include dissipation such as viscosity and a body force (to be specified in Section 3). The background magnetic field is fixed in time and is uniform in space, such that \n\n\n\n\u2207\u00d7B0=B0\u2207\u00d7x\u02c6\u2225B0=0\n\n. Care must be taken if considering a spatially dependent background magnetic field \nB\n\n\n0\n(\nx\n), as the resulting equation will not be the same. See the discussion in Section 5. This equation is accompanied with the incompressibility condition \u2207 \u00b7 \nv\n = 0. The induced magnetic field can be found using a diagnostic relation\n3\n\n\n\nb=\u2212\u2207\u22122x\u02c6\u2225B0\u00b7\u2207v,\n\nwhich would be \n\n\n\nb=\u2212\u2207\u22122B0\u00b7\u2207v\/\u03b7\n\n in dimensional variables.","Citation Text":["Favier et al. 2010"],"Functions Text":["This limit is referred to as the quasi-static approximation to MHD","although recent numerical studies on its turbulent properties and anisotropy have been done as well"],"Functions Label":["Background","Background"],"Citation Start End":[[748,766]],"Functions Start End":[[166,232],[602,701]]}
{"Identifier":"2019MNRAS.484.4083H__Pytte,_Mcpherron_&_Kokubun_1976_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":["Pytte, Mcpherron & Kokubun 1976"],"Functions Text":["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"],"Functions Label":["Uses"],"Citation Start End":[[1085,1116]],"Functions Start End":[[847,1083]]}
{"Identifier":"2018AandA...609A..13K__Mucciarelli_et_al._(2017)_Instance_3","Paragraph":"Gaia 1 is a star cluster that was recently discovered by Koposov et al. (2017) in the first Gaia data release (Gaia Collaboration 2016), alongside with another system of lower mass. Its observation and previous detections were seriously hampered by the nearby bright star Sirius, which emphasized the impressive discovery power of the Gaia mission. This object was first characterized as an intermediate-age (6.3 Gyr) and moderately metal-rich (\u22120.7 dex) system, based on isochrone fits to a comprehensive combination of Gaia, 2MASS (Cutri et al. 2003), WISE (Wright et al. 2010), and Pan-STARRS1 (Chambers et al. 2016) photometry. Hence, this object was characterized by Koposov et al. (2017) as a star cluster, most likely of the globular confession. Further investigation of Gaia 1 found a metallicity higher by more than 0.5 dex, which challenged the previous age measurement and rather characterized it as a young (3 Gyr), metal-rich (\u22120.1 dex) object, possibly of extragalactic origin given its orbit that leads it up to ~1.7 kpc above the disk (Simpson et al. 2017). Subsequently, Mucciarelli et al. (2017) measured chemical abundances of six stars in Gaia 1, suggesting an equally high metallicity, but based on their abundance study, the suggestion of an extragalactic origin was revoked. While a more metal-rich nature found by the latter authors conformed with the results by Simpson et al. (2017), the evolutionary diagrams of both studies are very dissimilar and could not be explained by one simple isochrone fit. In particular, it was noted that \u201cthe Simpson et al. (2017) stars do not define a red giant branch in the theoretical plane, suggesting that their parameters are not correct\u201d (Fig. 1 of Mucciarelli et al. 2017). Such an inconsistency clearly emphasizes that a clear-cut chemical abundance scale is inevitable for fully characterising Gaia 1, and to further allow for tailored age determinations, even more so in the light of the seemingly well-determined orbital characteristics, Thus, this work focuses on a detailed chemical abundance analysis of four red giant members of Gaia 1, based on high-resolution spectroscopy, which we complement with an investigation of the orbital properties of this transition object. Combined with the red clump sample of Mucciarelli et al. (2017) and reaching down to the subgiant level (Simpson et al. 2017), stars in different evolutionary states in Gaia 1 are progressively being sampled. ","Citation Text":["Mucciarelli et al. (2017)"],"Functions Text":["Combined with the red clump sample of","stars in different evolutionary states in Gaia 1 are progressively being sampled."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[2283,2308]],"Functions Start End":[[2245,2282],[2372,2453]]}
{"Identifier":"2016AandA...595A..72M__Vergani_et_al._2015_Instance_2","Paragraph":"On the other hand, the Australia Telescope Compact Array (ATCA) 21\u2009cm line survey of GRB host galaxies revealed high levels of atomic hydrogen (H\u2009i), suggesting that the connection between atomic gas and star formation is stronger than previously thought (Micha\u0142owski et al. 2015). Star formation may be directly fuelled by atomic gas, as has been theoretically shown to be possible (Glover & Clark 2012; Krumholz 2012; Hu et al. 2016), and this is supported by the existence of H\u2009i-dominated, star-forming regions in other galaxies (Bigiel et al. 2008, 2010; Fumagalli & Gavazzi 2008; Elmegreen et al. 2016). This can happen in a low metallicity gas that is recently acquired, even if the metallicity in other parts of a galaxy is higher, near the onset of star formation because cooling of gas (necessary for star formation) is faster than the H\u2009i-to-H2 conversion (Krumholz 2012). Indeed, large atomic gas reservoirs, together with low molecular gas masses (Hatsukade et al. 2014; Stanway et al. 2015b) and stellar masses (Perley et al. 2013, 2015; Vergani et al. 2015), indicate that GRB hosts are preferentially galaxies that have very recently started a star formation episode. This provides a natural route for forming GRBs in low-metallicity environments, as found for most GRB hosts (Fruchter et al. 2006; Modjaz et al. 2008; Levesque et al. 2010a; Han et al. 2010; Boissier et al. 2013; Schulze et al. 2015; Vergani et al. 2015; Japelj et al. 2016; Perley et al. 2016), except of a few examples of hosts with solar or super-solar metallicities (Prochaska et al. 2009; Levesque et al. 2010b; Kr\u00fchler et al. 2012; Savaglio et al. 2012; Elliott et al. 2013; Schulze et al. 2014; Hashimoto et al. 2015; Schady et al. 2015; Stanway et al. 2015a). Indeed, the GRB collapsar model requires that most of the GRB progenitors have low metallicity (below solar) in order to reduce the loss of mass and angular momentum that is required for launching the jet (Yoon & Langer 2005; Yoon et al. 2006; Woosley & Heger 2006). We note however that other models, while still predicting the metallicity preference (e.g. Izzard et al. 2004; Podsiadlowski et al. 2004; Detmers et al. 2008), allow higher metallicities owing to differential rotation (Georgy et al. 2012), binary evolution (Podsiadlowski et al. 2010; van den Heuvel & Portegies Zwart 2013), or weaker magnetic fields (Petrovic et al. 2005). ","Citation Text":["Vergani et al. 2015"],"Functions Text":["This provides a natural route for forming GRBs in low-metallicity environments, as found for most GRB hosts"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1418,1437]],"Functions Start End":[[1184,1291]]}
{"Identifier":"2019ApJ...872..143B__Seckel_et_al._1991_Instance_1","Paragraph":"The gamma-ray emission from the solar disk due to CR cascades in the solar atmosphere is denoted as a disk component. This secondary gamma-ray produced by the hadronic interaction of cosmic ray with the solar surface was first proposed by Dolan & Fazio (1965). While only upper limits were obtained by early measurements over the range 20 keV\u201310 MeV (Peterson et al. 1966). A detailed theoretical model for gamma-rays from the collision of cosmic ray with the solar atmosphere was presented by Seckel et al. (1991). The predicted gamma-ray flux at energies from 10 MeV to 10 GeV has a large uncertainty, being sensitive to the assumptions about the cosmic-ray transport in the magnetic field near the Sun. Gamma rays from the Sun were first detected by the Energetic Gamma-ray Experiment Telescope (Orlando & Strong 2008). The measured flux from 100 MeV to 2 GeV was within the range of the theoretical predictions. The Fermi collaboration (Abdo et al. 2011) reported the detection of high energy gamma-rays at 0.1\u201310 GeV from the quiescent Sun using the first 1.5 yr data. However, the measured solar disk emission flux was about a factor of seven higher than that predicted about this disk component by a \u201cnominal\u201d model (Seckel et al. 1991). This mismatch motivated Ng et al. (2016) to analyze 6 yr of public Fermi-LAT data. The obtained gamma-ray spectrum follows a simple power-law shape (\u03b1 = \u22122.3) in 1\u2013100 GeV without any evident high energy cutoff. For the flux in 1\u201310 GeV, a significant time variation of the solar disk gamma-ray flux that anticorrelates with solar activity was discovered, suggesting that the solar magnetic field would play an important role. An updated analysis with 9 yr of Fermi-LAT data, from 2008 August 7 to 2017 July 27, was performed, and Tang et al. (2018) confirmed these results and extended the gamma-ray spectrum up to >200 GeV. Notably, the bright gamma-ray flux above 100 GeV is dominant only during solar minimum at the end of Cycle 23 (Linden et al. 2018). The HAWC measurements in periods of high solar activity may support these findings (Albert et al. 2018a). Data collected from 2014 November to 2017 December, the second half of solar cycle 24, have been used to set strong upper limits on the flux of 1\u2013100 TeV gamma-rays from the solar disk, about 10% of the maximum gamma-ray flux estimated by Linden et al. (2018). The HAWC 95% upper limit at 1 TeV is about 13% of the flux extrapolated from the solar minimum Fermi-LAT gamma-ray spectrum.","Citation Text":["Seckel et al. (1991)"],"Functions Text":["A detailed theoretical model for gamma-rays from the collision of cosmic ray with the solar atmosphere was presented by","The predicted gamma-ray flux at energies from 10 MeV to 10 GeV has a large uncertainty, being sensitive to the assumptions about the cosmic-ray transport in the magnetic field near the Sun."],"Functions Label":["Background","Background"],"Citation Start End":[[494,514]],"Functions Start End":[[374,493],[516,705]]}
{"Identifier":"2017ApJ...850...18H__Murase_et_al._2015_Instance_2","Paragraph":"The heating due to the reprocessing of non-thermal photons produced in the nebula can be efficient even at late times. Here, we treat these processes in a simple way as follows. At early times, electromagnetic cascades proceed in the saturation regime, leading to a flat energy spectrum up to \u223c1 MeV (Metzger et al. 2014). At later times, the spectrum depends on the seed photon spectra, but it can roughly be estimated to be a flat spectrum from \u223c1 eV to \u223c0.1 TeV, while the supernova emission continues, which is expected based on more detailed calculations (e.g., Murase et al. 2015; K. Murase et al. 2017, in preparation). High-energy \u03b3-rays (\u22731 MeV) heat up the ejecta through the Compton scattering and the pair production process. X-ray and UV photons are absorbed and heat up the ejecta through the photoelectric (bound-free) absorption unless the ejecta are fully ionized. Here, we describe the heating rate as\n12\n\n\n\n\n\n\n\n\nQ\n\n\n\u02d9\n\n\n\n\nrad\n\n\n(\nt\n)\n\u2248\n\n\n(\n\n\n\nf\n\n\n\u03b3\n\n\n+\n\n\nf\n\n\n\n\n\nX\n\u2212\nUV\n\n\n,\nbf\n\n\n\n)\n\n\n\n\nL\n\n\nsd\n\n\n,\n\n\nwhere f\u03b3 and \n\n\n\n\n\n\nf\n\n\n\n\n\nX\n\u2212\nUV\n\n\n,\nbf\n\n\n\n\n are the heating efficiencies of \u03b3-rays and X-ray and UV photons to the spin-down luminosity, respectively. We calculate the frequency dependent heating efficiency of \u03b3-rays at each time:\n13\n\n\n\n\n\n\nf\n\n\n\u03b3\n\n\n(\nt\n)\n=\n\n\n\n\n\n\u222b\n\n\n\n\n\u03bd\n\n\nmin\n\n\n\n\n\n\n\u03bd\n\n\nmax\n\n\n\n\n\n\n\nd\n\u03bd\n\n\n\u03bd\n\n\n\n\nmin\n(\n\n\nK\n\n\n\u03b3\n,\n\u03bd\n\n\n\n\n\u03c4\n\n\n\u03b3\n,\n\u03bd\n\n\n,\n1\n)\n\n\n\n\n\u222b\n\n\n1\n\neV\n\n\n1\n\nTeV\n\n\n\n\n\nd\n\u03bd\n\n\n\u03bd\n\n\n\n\n\n\n,\n\n\nwhere the frequency range of \u03b3-rays is \n\n\n\n\n(\nh\n\n\n\u03bd\n\n\nmin\n\n\n,\nh\n\n\n\u03bd\n\n\nmax\n\n\n)\n\n=\n(\n10\n\nkeV\n,\n1\n\nTeV\n)\n\n\n, and h is the Planck constant. Here, \u03c4\u03b3, \u03bd is the optical depth of the ejecta to \u03b3-rays and K\u03b3, \u03bd is the photon inelasticity at a given frequency, where the Klein\u2013Nishina cross section and the cross section for the Bethe\u2013Heitler pair production in the field of a carbon nucleus are taken into account (Chodorowski et al. 1992; Murase et al. 2015). Note that the coefficient of the \u03b3-ray optical depth depends on the density profile of the ejecta. Here, we simply assume a density profile to be constant with the radius. Adopting a realistic density profile may result in different ejecta mass and velocity estimates by a factor of a few.","Citation Text":["Murase et al. 2015"],"Functions Text":["Here, \u03c4\u03b3, \u03bd is the optical depth of the ejecta to \u03b3-rays and K\u03b3, \u03bd is the photon inelasticity at a given frequency, where the Klein\u2013Nishina cross section and the cross section for the Bethe\u2013Heitler pair production in the field of a carbon nucleus are taken into account"],"Functions Label":["Uses"],"Citation Start End":[[1845,1863]],"Functions Start End":[[1549,1818]]}
{"Identifier":"2020ApJ...898...92C__Harrison_&_Tsang_1976_Instance_1","Paragraph":"We recorded the IRPD spectra of mass-selected H+MF and H+MF-Ln\u22642 clusters (L = Ar\/N2) in the 2950\u20133650 cm\u22121 spectral range in a tandem quadrupole mass spectrometer interfaced with an electron ionization source (Dopfer 2003, 2005), a setup used previously to record IR spectra of a variety of hydrocarbon cations and their clusters (Dopfer et al. 2002; Solc\u00e0 & Dopfer 2002; Andrei et al. 2008; Patzer et al. 2010; Chatterjee & Dopfer 2017, 2018a, 2018b, 2020; Chatterjee et al. 2019). Briefly, the parent ions are produced in an ion source, which combines a pulsed supersonic expansion with electron (chemical) ionization close to the nozzle orifice. The expanding gas mixture is generated by passing a carrier gas mixture of Ar (N2) and 5% H2 in He in a ratio 5:1 at 10 bar through a reservoir containing liquid MF (Sigma-Aldrich, 99%, heated to 323 K). One possible reaction pathway begins with electron ionization of H2, followed by exothermic proton transfer reactions to form H+MF and subsequent three-body association to generate H+MF-Ln clusters (Harrison & Tsang 1976; Hopkinson et al. 1979; Dopfer et al. 1999; Lawson et al. 2012; Chatterjee & Dopfer 2018a), according to:\n5\n\n\n\n\n\n\n\n6\n\n\n\n\n\n\n\n7\n\n\n\n\n\n\n\n8\n\n\n\n\n\nThis route is reminiscent of the synthesis of interstellar protonated molecules (Oka 2006; Etim et al. 2017). An alternative pathway involves charge transfer from Ar+ (or \n\n\n\n\n\n or \n\n\n\n\n\n) and subsequent self-protonation (Lawson et al. 2012; Chatterjee & Dopfer 2018a):\n5a\n\n\n\n\n\n\n\n6a\n\n\n\n\n\n\n\n7a\n\n\n\n\n\nThe desired ions are mass-selected in the first quadrupole and irradiated in an adjacent octupole with a tunable IR laser pulse (\u03bdIR, 10 Hz, 2\u20135 mJ pulse\u22121, 1 cm\u22121 bandwidth) of a Nd:YAG-pumped optical parametric oscillator. Calibration of \u03bdIR to better than 1 cm\u22121 is achieved by a wavemeter. Resonant vibrational excitation upon single-photon absorption leads to the fragmentation of neutral ligands or molecules:\n9a\n\n\n\n\n\n\n\n9b\n\n\n\n\n\n\n\n9c\n\n\n\n\n\nResulting fragment ions are mass-selected by the second quadrupole and monitored with a Daly detector as a function of \u03bdIR to derive the IRPD spectrum of the parent ions, which is linearly normalized for the energy of the laser pulse. The contribution from metastable decay is subtracted from the laser-induced dissociation signal by triggering the ion source at twice the laser repetition rate. The observed widths of the vibrational bands mainly originate from unresolved rotational structure, sequence hot bands involving inter- and intramolecular modes, lifetime broadening, and congestion produced from overlap of various structural isomers. Low-energy collision-induced dissociation (CID) mass spectra at \u223c10 eV collision energy in the laboratory frame are recorded by introducing 10\u22125 mbar N2 into the octople to confirm the composition of the ions and their clusters and establish their fragmentation channels.","Citation Text":["Harrison & Tsang 1976"],"Functions Text":["One possible reaction pathway begins with electron ionization of H2, followed by exothermic proton transfer reactions to form H+MF and subsequent three-body association to generate H+MF-Ln clusters"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1053,1074]],"Functions Start End":[[854,1051]]}
{"Identifier":"2022AandA...658A.194P__Khata_et_al._2020_Instance_1","Paragraph":"The stellar photospheric parameters we collected from literature for the benchmark stars are summarized in Table A.1. Although most benchmark stars have v sini  2 km s\u22121 (Reiners et al. 2018), there are two stars with larger values: J07558+833 (12.1 km s\u22121) and J13005+056 (16.4 km s\u22121). These stars are useful to investigate the performance of the algorithms when dealing with higher rotational velocities. The literature values were derived with different methods. These methods include: interferometry to estimate the stellar radius and Teff (Boyajian et al. 2012; S\u00e9gransan et al. 2003; von Braun et al. 2014; Berger et al. 2006; Newton et al. 2015), synthetic model fitting using BT-Settl models to determine Teff (Gaidos et al. 2014; L\u00e9pine et al. 2013; Gaidos & Mann 2014; Mann et al. 2015) and log g (L\u00e9pine et al. 2013), empirical relations to derive stellar mass in the form of mass-luminosity relations (Mann et al. 2015; Khata et al. 2020; Boyajian et al. 2012; Berger et al. 2006; S\u00e9gransan et al. 2003), along with the mass-magnitude relations (Maldonado et al. 2015), mass-radius relations (von Braun et al. 2014), mass\u2013Teff relations (Gaidos & Mann 2014; Gaidos et al. 2014), empirical relations to derive the stellar radius in the form of mass-radius relations (Maldonado et al. 2015) and Teff\u2013radius relations (Gaidos & Mann 2014; Gaidos et al. 2014; Houdebine et al. 2019), pEW measurements to determine Teff (Maldonado et al. 2015; Neves et al. 2014; Newton et al. 2015) and [Fe\/H] (Maldonado et al. 2015; Neves et al. 2014; Gaidos et al. 2014; Mann et al. 2015), the definition of spectral indices such as the H2O-K2 index to estimate Teff (Rojas-Ayala et al. 2012), as well as the combination of the H2O-K2 index with pEWs to derive [Fe\/H] (Rojas-Ayala et al. 2012; Khata et al. 2020), the stellar radius and Teff (Khata et al. 2020), and spectral curvature indices for the determination of Teff (Gaidos & Mann 2014). Additionally, [Fe\/H] was derived by using color-magnitude metallicity relations (Dittmann et al. 2016), atomic line strength relations (Gaidos & Mann 2014), and spectral feature relations (Terrien et al. 2015). Terrien et al. (2015) used K-band magnitudes and the Dartmouth Stellar Evolution Program (Dotter et al. 2008) to derive the stellar radius, whereas Mann et al. (2015) employed the Boltzmann equation with Teff determined from synthetic model fits. Last, but not least, Houdebine et al. (2019) derived Teff from photometric colors. For more details on the individual methods, we refer to the descriptions in the corresponding works.","Citation Text":["Khata et al. 2020"],"Functions Text":["These methods include:","empirical relations to derive stellar mass in the form of mass-luminosity relations"],"Functions Label":["Background","Background"],"Citation Start End":[[933,950]],"Functions Start End":[[467,489],[830,913]]}
{"Identifier":"2015ApJ...813...47M__Morley_et_al._2012_Instance_1","Paragraph":"The consequences of a potential rainout for a planetary atmosphere can be manyfold. First of all the rainout removes metals from the atmosphere, relocating them to deeper layers. Hence the corresponding grain or droplet opacity will be missing from higher atmospheric layers. Because we do not include cloud opacities in our calculations we make the implicit assumption of a rainout of the condensed particles, although we do not model it, the net effect being the removal of metals from the higher layers. It has to be kept in mind, however, that the chemical equilibrium solution of the gas abundances in chemical contact with the condensed species is not necessarily the same as it would be when assuming a rainout. Our implicit assumption of a rainout is also applicable when considering the gaseous Na and K alkali abundances. In our models MgSiO3 condenses at temperatures below \u223c1600 K. In principle this silicate material can further react with the alkali atoms to form alkali feldspars (such as albite and orthoclase), removing the gaseous alkalis from the gas for T \u2272  1600 K (see, e.g., Lodders 2010). We do not consider these feldspars in our condensation model, such that the alkali atoms stay in the gas, as they would in a silicate rainout scenario. It has been found that alkali atoms are present in cool brown dwarf atmospheres, indicating that silicate rainout may occur in these objects (Marley et al. 2002; Morley et al. 2012). Another consequence of condensed material can be the formation of a cloud deck, close to and above the layers of the atmosphere hot enough the evaporate the in-falling cloud particles again. Such cloud decks can heat the atmosphere locally and in the layers below, by making the atmosphere more opaque to the planet\u2019s intrinsic flux, effectively acting like a blanket covering the lower layers of the atmosphere (see, e.g., Helling & Casewell 2014; Morley et al. 2014). If the cloud layer is optically thick close to the planet\u2019s photosphere it will leave an imprint on the planet\u2019s spectral appearance and and may reduce the contrast of absorption features. The height of the cloud deck depends critically on the planets effective temperature and also on its surface gravity since the condensation temperature is pressure-dependent. The cooler an object is, the deeper in its interior the clouds will reside. Therefore the spectral imprint of clouds will vary with temperature, similar to the behavior in brown dwarf atmospheres. Silicate clouds with a high optical depth are thought to reside in the photospheres L4\u2013L6 type brown dwarfs (\n\n\n\n\n\n \u223c 1500\u20131700 K) where they affect the spectra. For cooler objects the cloud deck lies below the photosphere and the clouds are no longer seen (see, e.g., Lodders & Fegley 2006). In our atmospheres we checked the possible locations of the cloud decks (i.e., the layers below which the condensates evaporate). We found that the silicate evaporation layer of planets with \n\n\n\n\n\n = 1000 K  and \n\n\n\n\n\n K is always located at pressures far higher than that of the photosphere, such that we do not expect any spectral impact of a cloud layer. For effective temperatures between 1500 and 1750 K the evaporation layer lies close to and above the photosphere (in altitude), such that a cloud deck could potentially affect the spectrum. For increasing log(g) the photosphere shifts to layers of deeper pressure, but so does the evaporation layer, as condensation is pressure dependent. Note that this temperature range is close to the effective temperature where L4\u2013L6 dwarfs are thought to be most strongly affected by silicate clouds. For higher temperatures the evaporation layer is far above the photosphere such that we do not expect clouds to be of importance.","Citation Text":["Morley et al. 2012"],"Functions Text":["It has been found that alkali atoms are present in cool brown dwarf atmospheres, indicating that silicate rainout may occur in these objects"],"Functions Label":["Motivation"],"Citation Start End":[[1427,1445]],"Functions Start End":[[1265,1405]]}
{"Identifier":"2019MNRAS.485.3185C__Orienti_et_al._2015_Instance_1","Paragraph":"We detected radio emission for six sources, but while the morphology of radio emission in four over six cases is compact on $\\le \\,$arcseconds scales, predominantly unresolved, for NGC 3185 and NGC 3941 we find a more complex morphology. The estimated radio luminosities at this frequency are L$\\, \\sim \\, 5\\, \\times \\, 10^{20}$ and ${\\sim }\\, 6\\, \\times \\, 10^{18}$\u2009W\u2009Hz\u22121 for NGC 3185 and NGC 3941, respectively, well below the radio-loud\/radio-quiet threshold of L$\\, \\sim \\, 10^{23}$\u2009W\u2009Hz\u22121 defined by Condon (1992). Considering NGC 3185, the structure observed in L band has a size of $5\\, \\times \\,$3.8 arcsec, which translates into a linear scale of $0.5\\, \\times \\,$0.4 kpc.6 Analogous considerations can be made for the C and X Bands, leading to linear scales of $0.38\\, \\times \\,$0.36 kpc. There is evidence for radio emission spread over the host galaxy scale for a number of Seyferts when observed with adequate angular resolution and sensitivity (Orienti et al. 2015). In particular, circumnuclear starburst rings with knots of star formation are observed. Orienti & Prieto (2010) observed diffuse radio emission with knots of star formation for a number of Seyferts, on scales smaller than 1 kpc (NGC 5506, NGC 7469, and NGC 7582). Analogously, the morphology of the radio emission in the case of NGC 3185 as observed in our radio maps may be consistent with emission from circumnuclear rings of star formation. In order to check this hypothesis, we overlapped L band and C band radio contours to an archivial HST image7 for NGC 3185, as shown in Fig. 4 . The radio contours overlap with optical emission, from which a nearly circular, ring-like structure emerges. However, further work is needed to understand the link between radio emission and star formation in the form of circumnuclear rings in this source. In particular, intermediate angular resolution scales radio observations (such as those of e-MERLIN) combined with H \u03b1 maps would allow us to confirm the star formation hypothesis. If we consider the parent sample with our new data, then we find that the average spectral slope is compatible with flat (0.30 \u00b1 0.10), with a nearly equal number of flat and steep slopes. If we consider the two sub-populations of type-1 and type-2 Seyferts, then no significant differences are found with respect to the average spectral index and the average radio-loudness parameter, the only difference being that type-1 sources are slightly more luminous. Considering the X-ray radio loudness parameter, the black hole mass and the Eddington ratio, we do not find significant differences between the type-1 and type-2 sub-classes. We note that, even though the sources in the Reference sample have an Eddington ratio nearly an order of magnitude smaller than the average for the Parent sample, the X-ray radio-loudness parameter does not exhibit anomalous values.","Citation Text":["Orienti et al. 2015"],"Functions Text":["There is evidence for radio emission spread over the host galaxy scale for a number of Seyferts when observed with adequate angular resolution and sensitivity"],"Functions Label":["Background"],"Citation Start End":[[960,979]],"Functions Start End":[[800,958]]}
{"Identifier":"2017ApJ...845...92R__Mosqueira_&_Estrada_2003a_Instance_1","Paragraph":"Concerning case (1), the first explanation that has been proposed is an increasing relative velocity among the building blocks with decreasing distance from the planet leading to substantial water loss in the case of the most energetic impacts (Estrada & Mosqueira 2006), which occurred closer to Jupiter. Nonetheless, this scenario has been discarded by a detailed study showing that Io and Europa analogs exhibit an overabundance of water when they are formed via an N-body code simulating imperfect accretion and water loss during collisions (Dwyer et al. 2013). A second explanation is that the observed water gradient among the satellites results from an outwardly decreasing temperature of the CPD, leading to the existence of a snowline at a given radial distance from Jupiter (see, e.g., Lunine & Stevenson 1982). In this case, bodies that formed inward of the snowline (Io) accreted from essentially water-poor building blocks, whereas bodies that formed outward of the snowline (Ganymede, Callisto) formed from a primordial mixture of water ice and silicates (e.g., Canup & Ward 2002; Mosqueira & Estrada 2003a, 2003b; Mousis & Gautier 2004). Within this scenario, the low water content of Europa is puzzling. So far, Europa\u2019s water content has been mostly attributed to its formation both outward and inward of the snowline due to either (i) its migration inward of the snowline during formation (i.e., growth), (ii) the progressive cooling of the disk and thus inward migration of the snowline during its formation, or (iii) an interplay between the two mechanisms (Alibert et al. 2005; Canup & Ward 2009). However, the evolution of the CPD has been systematically modeled using an ad hoc parametrization of the turbulent viscous disk (the so-called \u03b1-viscosity, Shakura & Sunyaev 1973) which governs the temperature evolution and lifetime of the disk. While providing a good starting point for evolutionary disk models, this kind of parametrization has been highly questioned in recent years (Bai & Stone 2013; Simon et al. 2013; Gressel et al. 2015). Hence, using a predefined \u03b1-viscosity prescription to describe the CPD\u2019s evolution and provide hints on Europa\u2019s formation remains questionable. The same remark holds for planet (or satellite) migration, which has also been extensively studied within recent years (see, e.g., Paardekooper et al. 2010; Bitsch et al. 2014). These studies have shown that in realistic disk conditions, migrating planets can behave significantly differently from what was previously thought, i.e., a persistent inward motion (e.g., Tanaka et al. 2002), due to the existence of regions where the migration is halted and even reversed. Because the studies of satellite formation have been based so far on the migration formulation of Tanaka et al. (2002; e.g., Canup & Ward 2002, 2006; Alibert et al. 2005; Sasaki et al. 2010) their proposed growth\/migration scenario is questionable.","Citation Text":["Mosqueira & Estrada 2003a"],"Functions Text":["In this case, bodies that formed inward of the snowline (Io) accreted from essentially water-poor building blocks, whereas bodies that formed outward of the snowline (Ganymede, Callisto) formed from a primordial mixture of water ice and silicates (e.g.,","Within this scenario, the low water content of Europa is puzzling."],"Functions Label":["Background","Differences"],"Citation Start End":[[1095,1120]],"Functions Start End":[[822,1075],[1153,1219]]}
{"Identifier":"2021AandA...646A..21C__Polyansky_et_al._(2018)_Instance_1","Paragraph":"(1a) Work is underway for an updated ExoMol line list for SiO which will extend into the ultraviolet. The current line list only considers vibration-rotation transitions and so the current maximum wavenumber was set at 6049 cm\u22121. (1b) The HITEMP line list for NO includes data from the ExoMol NOname line list (Wong et al. 2017). (1c) The Ames line list (Huang et al. 2017) and the CDSD-4000 databank (Tashkun & Perevalov 2011) are also available for CO2, as well as the HITEMP compilation (Rothman et al. 2010). (1d) The previous ExoMol line list for H2O, BT2 (Barber et al. 2006), is only complete up to temperatures of 3000 K, whereas the more accurate ExoMol POKAZATEL line list Polyansky et al. (2018) is complete up to 4000 K. (1e) There is also a line list for MgH from ExoMol Yadin (Yadin et al. 2012). However, since it only covers the ground electronic X2\u03a3+ state, and so is less complete than the more recent MoLLIST line list of Gharib-Nezhad et al. (2013), we use the latter. (1f) Previous to the ExoMol aCeTY line list of Chubb et al. (2020b), the main sources of data for acetylene were from HITRAN (Gordon et al. 2017) and ASD-1000 (Lyulin & Perevalov 2017). The data from HITRAN are only applicable for studies performed at room temperature, and were shown in Chubb et al. (2020b) to be inadequate for high-temperature applications. ASD-1000 was a vast improvement, although there does seem to be opacity missing from some of the hot bands when compared to ExoMol aCeTY in Chubb et al. (2020b). (1g) The previous ExoMol line list for CH4, called 10\u201310 (Yurchenko & Tennyson 2014), is only complete up to 1500 K. The updated 34\u201310 line list is therefore recommended instead. Future updates of the database will investigate using data for methane based on recent line lists from either TheoReTs (Rey et al. 2017) or HITEMP (Hargreaves et al. 2020); these are currently expected to be more accurate when considering high-resolution applications. For low-resolution applications, we expect the quality of the ExoMol line list used here to be sufficient, particularly because completeness is more important than accuracy at lower resolutions (Yurchenko et al. 2014). (1h) The energy states from Coppola et al. (2011) are used in the Amaral et al. (2019) line list for HD+. (1i) The energy states from Engel et al. (2005) are used in the Amaral et al. (2019) line list for HeH+. (1j) The pressure and temperature broadened profiles for the resonance doublets of Na and K are computed using Allard et al. (2016, 2019). See Sect. 3.2.3 for a discussion on the broadening profiles of these atoms.","Citation Text":["Polyansky et al. (2018)"],"Functions Text":["whereas the more accurate ExoMol POKAZATEL line list","is complete up to 4000 K."],"Functions Label":["Background","Background"],"Citation Start End":[[683,706]],"Functions Start End":[[630,682],[707,732]]}
{"Identifier":"2022ApJ...935..148M__Shamasundar_et_al._2011_Instance_1","Paragraph":"Using the high-level ab initio calculations implemented in the MOLPRO 2015 software package (Werner et al. 2015, 2020), the electronic structures of CO have been determined. In our calculations, the molecular orbitals (MOs) and ground-state energies are computed by the HF method. Then, the CASSCF method (Knowles & Werner 1985; Werner & Knowles 1985) is applied to perform the state-averaged calculation to generate multi-configuration wave functions by utilizing the HF MOs as the starting orbitals. Finally, based on the CASSCF wave functions, the internally contracted multi-reference configuration interaction method with the Davidson correction (icMRCI+Q) (Knowles & Werner 1988; Werner et al. 1988; Knowles & Werner 1992; Shamasundar et al. 2011) is performed to consider the dynamic correlation and size-consistency error. The augmented correlation consistent polarized weighted core valence quintuplet aug-cc-pwCV5Z-DK Gaussian basis set is selected to describe the carbon and oxygen atoms, which is found to give an excellent production of the potential energy and dipole moments for electronic states as the CO molecule dissociates, as mentioned in previous publications (Dunning 1989; De Jong et al. 2001; Peterson & Dunning 2002). The scalar relativistic effect is considered by the third-order Douglas\u2013Kroll Hamiltonian approximation (Reiher & Wolf 2004, 2004). The PECs, PDMs, and TDMs for the singlet and triplet states were computed at the internuclear distances from 0.8\u20137.5 \u00c5 with step sizes of 0.05 \u00c5 for 0.8\u20130.9 \u00c5, 0.02 \u00c5 for 0.9\u20131.5 \u00c5, 0.05 \u00c5 for 1.5\u20132.6 \u00c5, 0.1 \u00c5 for 2.6\u20136 \u00c5, and 0.5 \u00c5 for 6\u20137.5 \u00c5, and those for the quintet states were computed at the internuclear distances from 0.8\u201315 \u00c5 with step sizes of 0.05 \u00c5 for 0.8\u20130.9 \u00c5, 0.02 \u00c5 for 0.9\u20131.5 \u00c5, 0.05 \u00c5 for 1.5\u20132.6 \u00c5, 0.1 \u00c5 for 2.6\u20136 \u00c5, and 0.5 \u00c5 for 6\u201315 \u00c5. The obtained PECs can be also used to determine the dissociation energy D\n\ne\n, the electronic excitation energy relative to the ground state T\n\ne\n, the internuclear separation R\n\ne\n, the harmonic frequency \u03c9\n\ne\n, the first-order anharmonic constant \u03c9\n\ne\n\n\u03c7\n\ne\n, the rotational constant B\n\ne\n, and the vibrational coupling constant \u03b1\n\ne\n.","Citation Text":["Shamasundar et al. 2011"],"Functions Text":["Finally, based on the CASSCF wave functions, the internally contracted multi-reference configuration interaction method with the Davidson correction (icMRCI+Q)","is performed to consider the dynamic correlation and size-consistency error."],"Functions Label":["Uses","Uses"],"Citation Start End":[[729,752]],"Functions Start End":[[502,661],[754,830]]}
{"Identifier":"2019MNRAS.487.3702O__Owen_et_al._2011b_Instance_1","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"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[1764,1781]],"Functions Start End":[[1598,1762]]}
{"Identifier":"2020ApJ...898...49N__Sowmya_et_al._2014_Instance_1","Paragraph":"In the line cores of Q\/I and U\/I profiles, we see depolarization and rotation for fields in the range \n\n\n\n\n\n G. These are due to the Hanle effect. For \n\n\n\n\n\n G, we see the signatures of level-crossings in the line cores of (Q\/I, U\/I) profiles, namely they tend toward the non-magnetic value (see Figures 1(b) and 2(b)). We recall that, traditionally the loops in the polarization diagram (namely, a plot of Q\/I versus U\/I for a given wavelength and for a range of magnetic field strength or orientation values) are identified to be due to the level-crossings in the incomplete PBE regime (see, e.g., Bommier 1980; Landi Degl\u2019Innocenti & Landolfi 2004, see also Sowmya et al. 2014). When a given curve in the polarization diagram forms a loop the Q\/I and U\/I values tend toward the non-magnetic value. Based on this we identify the above noted behavior in the line cores of (Q\/I, U\/I) profiles for the mentioned field-strength regime as to be the signatures of level-crossings in the incomplete PBE regime. Polarization diagrams require the use of very fine grids of magnetic field strength or orientation. With the radiative transfer calculations presented in this paper, it is computationally difficult to produce such diagrams. For \n\n\n\n\n\n G, transverse Zeeman effect like signatures are seen in the line core of (Q\/I, U\/I) profiles (see Figures 1(c) and 2(c)). The Faraday rotation (del Pino Alem\u00e1n et al. 2016; Alsina Ballester et al. 2017; Sampoorna et al. 2017), which results in depolarization in the wings of Q\/I and generation of U\/I in the wings, strongly influences the wings of U\/I profiles for the entire field-strength regime considered here, while it shows up in Q\/I for \n\n\n\n\n\n G. For the cases of theoretical model line and the isothermal model atmosphere considered in this section, the Voigt effect starts to show up in U\/I for \n\n\n\n\n\n G and in Q\/I for \n\n\n\n\n\n G, and its signatures are similar to those discussed in Sampoorna et al. (2019a). Also we see the signatures of incomplete PBE in the V\/I profiles, which are now asymmetric about the line center for fields up to 30 G.","Citation Text":["Sowmya et al. 2014"],"Functions Text":["We recall that, traditionally the loops in the polarization diagram (namely, a plot of Q\/I versus U\/I for a given wavelength and for a range of magnetic field strength or orientation values) are identified to be due to the level-crossings in the incomplete PBE regime","see also"],"Functions Label":["Background","Background"],"Citation Start End":[[661,679]],"Functions Start End":[[320,587],[652,660]]}
{"Identifier":"2020ApJ...888..126Z__Warren_et_al._2006_Instance_1","Paragraph":"Achondritic meteorites are fragments of differentiated asteroids or planetary bodies of the solar system. Ureilites are coarse-grained ultramafic (olivine\u2013pyroxene) achondrites. They include accessory minerals of metal and sulfide associated with high abundances of carbon phases (on average 3 vol% and up to \u223c7 vol%; Goodrich et al. 2015), including graphite and high-pressure diamond (Goodrich 1992; Mittlefehldt et al. 1998; Nabiei et al. 2018). According to their petrology, ureilites are further divided into main group ureilites (formerly referred to as monomict or unbrecciated ureilites; accounting for 95%) and rare polymict ureilites (or brecciated ureilites). The main group ureilites represent the mantle residues after the extraction of feldspar-rich magmas (Cohen et al. 2004; Bischoff et al. 2014; Barrat et al. 2016) and a sulfur-rich iron melt (Goodrich et al. 2004; Warren et al. 2006; Goodrich et al. 2007; Rankenburg et al. 2008; Barrat et al. 2015), whereas polymict ureilites are breccias containing fragments of main group ureilites as well as non-ureilite clasts (Goodrich et al. 2004, 2015). Unlike samples from other differentiated planetary bodies (e.g., terrestrial and lunar samples, Martian meteorites, angrites), some primitive isotopic geochemical features (e.g., O isotope heterogeneities; Greenwood et al. 2005, 2006, 2017) have been preserved in ureilites, which are considered incompatible with a magma ocean event (e.g., Goodrich 1992; Goodrich et al. 2004, 2015). Therefore, the geochemical signatures of ureilite meteorites can provide valuable information about the origin of the ureilite parent body (UPB) and the early evolution of the solar system. For instance, olivine cores from individual ureilite meteorites cover a range of Fe\/Mn ratios from 3 to 57 (Goodrich 1992; Goodrich et al. 2004, 2007; Downes et al. 2008) and \u039417O values (mass-independent O isotopic variations) from \u22122.5\u2030 to \u22120.2\u2030 (Clayton & Mayeda 1988, 1996; Greenwood et al. 2017). The whole rock \u039417O values further correlate with both the Fe\/Mn ratios and magnesium number (Mg#: molar ratio of Mg\/Mg+Fe) of their olivine cores (Goodrich et al. 2004, 2015). The origin of these covariations is debated and multiple scenarios have been proposed to explain these features, including (1) a \u201csmelting\u201d process which causes redox reactions between C and FeO during UPB differentiation (Singletary & Grove 2003; Goodrich et al. 2007); (2) oxidation of metal due to the presence of water prior to igneous differentiation (Sanders et al. 2017); (3) inherited nebular redox variability from the precursor materials that were not homogenized during partial melting (Warren & Huber 2006; Warren 2012); (4) UPB accretion from the mixing of two different chemical and isotopic reservoirs (Barrat et al. 2017). As such, ureilite genesis still remains unclear. Understanding the origin of these geochemical correlations in ureilites is further complicated by the fact that the timing of formation and early evolution of the UPB is poorly constrained.","Citation Text":["Warren et al. 2006"],"Functions Text":["The main group ureilites represent the mantle residues after the extraction of","and a sulfur-rich iron melt"],"Functions Label":["Background","Background"],"Citation Start End":[[884,902]],"Functions Start End":[[671,749],[833,860]]}
{"Identifier":"2018AandA...618A..38K__Kirchschlager_&_Wolf_2013_Instance_1","Paragraph":"In Fig. 1, we investigate the dust grain size distributions at the periastron and apastron of debris disks a belt eccentricity of eb = 0.4 (with dynamical excitation \u0394eb = 0.1) for different material strengths. Since the radiation pressure has a strong influence on the cutoff in the grain size distribution at smallest sizes, the wavy patterns of the grain size distributions start around the smallest bound grains. We note that the blowout size can be increased with particle porosity and increasing stellar temperature (Burns et al. 1979; Kirchschlager & Wolf 2013; Brunngr\u00e4ber et al. 2017). We find different characteristic wavy patterns in the grain size distributions starting with the depletion of grains with radii below the blowout radius abo, which are caused by the lack of \u03b2 > 0.5 dust grains (Th\u00e9bault et al. 2003), depending on different collisional evolutions (Figs. 1 and 2). For grains with radii ~3 \u03bcm at the apastron side and grains with radii ~6 \u03bcm at the periastron side, a strong wavy pattern develops. This depletion leads to an over-density of slightly larger grains (\u201cfirst peak\u201d) because grains with radii a abo are depleted and thus can no longer contribute efficiently to the destruction and erosion processes anymore. The overabundance of grains with radii around the first peak (~1.5 abo) in turn induces another depletion of grains with radii around the \u201csecond depletion\u201d (~4\u201340 abo) that is caused by small high-\u03b2 grains originating inside the disks. Thus, an efficient destruction is responsible for the deep depletion of objects with radii of up to ~4\u221240 abo. Qualitatively, this first wavy pattern is less pronounced in the size distribution of higher material strengths, where the impact velocities and thus their rate of destructive collisions is significantly lower. A strong depletion for lower material strengths is found for grains with radii ~60 \u03bcm, while the depletion in the case of higher material strengths is valid for radii \u226410 \u03bcm (Fig. 1). This depletion eventually leads to an over-density of grains with radii around the \u201csecond peak\u201d. The overabundance of these grains is shifted from ~100 \u03bcm (high material strength) to ~1000 \u03bcm (very low material strength).","Citation Text":["Kirchschlager & Wolf 2013"],"Functions Text":["We note that the blowout size can be increased with particle porosity and increasing stellar temperature"],"Functions Label":["Uses"],"Citation Start End":[[542,567]],"Functions Start End":[[417,521]]}
{"Identifier":"2021ApJ...910...84E__Rappazzo_&_Velli_2011_Instance_1","Paragraph":"Based on these results, a series of 3D numerical simulations solving the simplified reduced magnetohydrodynamic (RMHD) equations, introduced by Strauss (1976), in Cartesian geometry were performed. The goal of these simulations was twofold: (1) to determine how a coronal loop responds to different photospheric velocity patterns, and (2) to investigate how the electric current sheets are formed, as well as the details of the heating that occurs through the dissipation of magnetic energy (Rappazzo et al. 2007, 2008, 2010; Rappazzo & Velli 2011; Rappazzo & Parker 2013; Rappazzo 2015). The results detailed in these papers confirmed the 2D results that the energy flux entering the corona because of photospheric motions causes a turbulent cascade of energy toward small scales, with electric current sheets continuously being created and destroyed throughout the coronal loop. The turbulence that develops is highly intermittent. The dissipation of energy occurs in current sheets, which are localized structures. Hence, the heating also occurs on small spatial scales. (Note that there is also particle acceleration, although that is beyond the scope of the present model.) When the loop is fully turbulent, it achieves a statistically steady state in which, on average, the Poynting flux induced at the boundaries by footpoint convection is balanced by the dissipation of energy in the electric current sheets. In this steady state, saturation occurs for the fluctuating magnetic energy, kinetic energy, mean square electric current, enstrophy, and other quantities that are then seen to fluctuate around their mean values. When the system becomes nonlinear, its behavior is highly dynamical and chaotic. This state occurs independently of the detailed form of the boundary velocity. A statistically steady state is achieved in which, although the magnetic field footpoints are convected continuously by the boundary flows, the resulting topology of the total magnetic field is not simply a mapping of those flows. By the same token, the turbulent dynamics that occur are due to the inherent nonlinear nature of the system, rather than being a consequence of the complexity of the magnetic field footpoint configuration.","Citation Text":["Rappazzo & Velli 2011"],"Functions Text":["The goal of these simulations was twofold: (1) to determine how a coronal loop responds to different photospheric velocity patterns, and (2) to investigate how the electric current sheets are formed, as well as the details of the heating that occurs through the dissipation of magnetic energy"],"Functions Label":["Motivation"],"Citation Start End":[[526,547]],"Functions Start End":[[198,490]]}
{"Identifier":"2022AandA...665A.115C__Clark_&_Steele_(2000)_Instance_2","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)"],"Functions Text":["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."],"Functions Label":["Background"],"Citation Start End":[[1749,1770]],"Functions Start End":[[1771,2083]]}
{"Identifier":"2022MNRAS.511.1819S__Bengaly_et_al._2019_Instance_1","Paragraph":"A sky survey at multiple frequencies, carried out with the Square Kilometre Array (SKA) could be used for an independent investigation of the controversial dipole anisotropy with a much superior sensitivity and thereby settle the question of the CP, hopefully, in a more decisive manner. In fact it was argued by Crawford (2009) that it is not possible to detect a radio dipole at more than 1\u03c3 level (for a presumed radio dipole amplitude equal to the CMBR dipole, i.e. p = 1), in a survey like the NVSS as the latter does not have sufficient number of sources. The conclusion drawn instead was that in order to make a positive detection of a radio dipole we may have to wait for the SKA data, going to sub-\u00b5Jy flux-density levels, and thus having much larger number of sources, to get a better signal to noise. In spite of this deterrent prediction, a radio dipole at statistically significant level was detected from the NVSS data itself (Singal 2011) which became possible only because the radio dipole amplitude turned out to be actually a factor of \u223c4 larger than the CMBR dipole. Although SKA might provide number counts at sub-\u00b5Jy levels (Schwarz et al. 2015; Bengaly et al. 2019), however, one may need to be wary of possible caveats in using radio source number counts at these flux-density levels as one might start seeing at sub-mJy levels, in addition to the powerful distant radio sources, a substantially increasing fraction of very different populations of radio sources, e.g. nearby normal galaxies, starburst galaxies, and even galactic sources (Windhorst 2003; Padovani 2011; Luchsinger et al. 2015). Instead what would be important is to get surveys at different frequencies from SKA, at a few mJy levels or above where the radio source population may comprise mostly powerful radio galaxies and quasars, and then investigate the dipoles to see how genuine is the difference in dipoles from surveys at widely separated frequency bands like that seen in the NVSS and TGSS dipoles. The multiple frequency flux-density measurements should also allow for the chosen sample a direct estimate of an optimum value of the spectral index, that enters in the expression for the dipole amplitude. The question of CP could be resolved in a positive manner if it is found that multiple-frequency radio source surveys yield dipoles consistent with the CMBR dipole, though it would still remain to be explained why the presently determined dipoles are not consistent with that. However, in case it does turn out that the SKA surveys with better statistical accuracies also yield dipole magnitudes which are significantly different from the CMBR dipole, even though might be pointing along the CMBR dipole direction, as is seen in the presently determined dipoles, then it will certainly be a big problem for the CP, and consequently all conventional cosmological models, including the ones dealing with the question of dark energy, could be in serious jeopardy.","Citation Text":["Bengaly et al. 2019"],"Functions Text":["Although SKA might provide number counts at sub-\u00b5Jy levels","however, one may need to be wary of possible caveats in using radio source number counts at these flux-density levels as one might start seeing at sub-mJy levels, in addition to the powerful distant radio sources, a substantially increasing fraction of very different populations of radio sources, e.g. nearby normal galaxies, starburst galaxies, and even galactic sources"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1167,1186]],"Functions Start End":[[1086,1144],[1189,1561]]}
{"Identifier":"2021ApJ...909..175Y__Buzzicotti_et_al._2018_Instance_3","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"],"Functions Text":["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"],"Functions Label":["Uses"],"Citation Start End":[[1468,1490]],"Functions Start End":[[1291,1466]]}
{"Identifier":"2015AandA...580A..71L__Sutton_et_al._(2013)_Instance_2","Paragraph":"The simplest two component model (power law + disk) is a phenomenological model often used to describe the spectra of ULXs as an empirical description of a disk plus corona geometry. In the presence of a cool (kT ~ 0.1\u22120.4 keV) and luminous (L ~ 1039\u22121040 erg\/s) disk, it allows inferring the presence of intermediate-mass black holes (e.g., Makishima et al. 2000). This is not the case for M33 X-8, where the disk component describes the high-energy part of the spectrum well and appears to be hot (kT ~ 1.15 keV), leaving a soft excess that is accounted for by the power law. The overall disk parameters are then inconsistent with a massive black hole, but instead are more typical of an ordinary stellar mass black hole: using the relationship between mass, temperature, and luminosity in a standard disk (see, e.g., Makishima et al. 2000), we derive a mass of ~10 M\u2299 for a nonrotating black hole, consistent with the estimation obtained by data from other satellites (e.g., Foschini et al. 2006; Weng et al. 2009; Isobe et al. 2012). Sutton et al. (2013) developed a classification scheme based on a disk+power law fit, to be applied to ULX spectra, according to which the spectral state of an ULX source can be defined by the disk temperature, the power-law slope, and the ratio between the flux contribution of the two spectral components in the 0.3\u22121 keV band. Our result is consistent with that found by Sutton et al. (2013) using XMM-Newton data, and, according to their classification, it identifies M33 X-8 as a broadened disk source, or in other words, as a source whose spectrum is dominated by emission from a hot disk (see Table 2) and where the additional soft component may be the effect of an unrealistic description of the disk spectrum by the diskbb model. In fact, such hot-disk\/soft power-law spectra are difficult to explain in the context of the analogy of ULXs with GBHs: the thermal state of GBHs is indeed characterized by a hot disk, but the presence of a soft power-law-like component in addition to the disk is unusual, and its physical interpretation is not simple: if this component is due to the presence of a Comptonized corona, we do not expect it to be dominant at energies lower than the temperature of the seed photons that come from the disk. ","Citation Text":["Sutton et al. (2013)"],"Functions Text":["Our result is consistent with that found by","using XMM-Newton data,","and, according to their classification, it identifies M33 X-8 as a broadened disk source, or in other words, as a source whose spectrum is dominated by emission from a hot disk (see Table 2) and where the additional soft component may be the effect of an unrealistic description of the disk spectrum by the diskbb model."],"Functions Label":["Similarities","Similarities","Uses"],"Citation Start End":[[1412,1432]],"Functions Start End":[[1368,1411],[1433,1455],[1456,1776]]}
{"Identifier":"2021AandA...648A..73B__Marois_et_al._2008_Instance_1","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"],"Functions Text":["In color-magnitude space, YSES 2b is very close to the innermost three planets of the HR 8799 multi-planetary system"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1133,1151]],"Functions Start End":[[1015,1131]]}
{"Identifier":"2019AandA...622A.180L__Cappellari_et_al._2011_Instance_1","Paragraph":"Large spectroscopic and photometric surveys, such as the Sloan Digital Sky Survey (SDSS, York et al. 2000), have revolutionized astrophysics in many fields, but can only deliver a limited view of the star formation activity in the local universe, given that instruments have a small field of view (FoV) in the case of spectroscopic surveys, and a low spectral resolution in the case of photometric ones. In recent years, integral field spectroscopic (IFS) surveys such as the Calar Alto Legacy Integral Field spectroscopy Area (CALIFA, S\u00e1nchez et al. 2012a) survey used in this work, have overcome these problems with the use of instruments with larger FoVs and a fully spectral coverage (e.g SAURON, Bacon et al. 2001; ATLAS3D, Cappellari et al. 2011; SAMI, Croom et al. 2012; VENGA, Blanc et al. 2013; and MaNGA, Bundy et al. 2015). However, these surveys still have limitations, such as the lack of a large contiguous observed area to trace the environment of nearby galaxies. They also suffer from selection biases due to the exclusion of galaxies with angular sizes that do not fit in the FoV of the integral field units (IFUs) that need to be considered (Walcher et al. 2014). These problems can be circumvented with multi-filter photometric surveys, which use a set of intermediate and narrow-band filters designed to provide the required spectral information while still covering a large contiguous area. The Javalambre Photometric Local Universe Survey (J-PLUS1, Cenarro et al. 2019) is currently operating to observe thousands of square degrees of the northern sky from the Observatorio Astrof\u00edsico de Javalambre (OAJ2) in Teruel, Spain. The survey is being carried out with the 0.83 m JAST\/T80 telescope and the panoramic camera T80Cam (Marin-Franch et al. 2015), with a 2 deg2 FoV. A set of twelve broad, intermediate, and narrow-band optical filters is used (Fig. 1 and Table 1), optimized to provide an adequate sampling of the spectral energy distribution (SED) of millions of stars in our galaxy. These SEDs will be required for the photometric calibration of the Javalambre Physics of the accelerating universe Astrophysical Survey (J-PAS3, Benitez et al. 2014). In addition, the position of the filters, the exposure times, and the survey strategy, are suitable to perform science that will expand our knowledge in many fields of astrophysics. Further details on the OAJ, instrumentation, filter set, J-PLUS photometric calibration process, strategy, and several science applications can be found in the J-PLUS presentation paper (Cenarro et al. 2019).","Citation Text":["Cappellari et al. 2011"],"Functions Text":["In recent years, integral field spectroscopic (IFS) surveys","have overcome these problems with the use of instruments with larger FoVs and a fully spectral coverage (e.g","ATLAS3D,"],"Functions Label":["Background","Background","Background"],"Citation Start End":[[729,751]],"Functions Start End":[[404,463],[584,692],[720,728]]}
{"Identifier":"2022ApJ...934...73A__Ryan_et_al._2011_Instance_1","Paragraph":"In contrast, deep narrow-field surveys can reach more distant UCD populations, enabling measurement of disk structure and a greater proportion of halo and thick disk sources. The majority of deep surveys for UCDs (Table 1) have been undertaken with the Hubble Space Telescope (HST), as these objects often comprise a foreground to extragalactic surveys. Early work in this area includes measurement of M dwarf number counts in the HST Deep Field and Large Area Multi-Color Survey Groth Strip (Gould et al. 1997; Kerins 1997; Chabrier & Mera 1997). Analysis of these samples determined M dwarf thin and thick disk vertical scaleheights of \u223c325 pc and \u223c650 pc, respectively, and ruled out very low-mass stars as being an appreciable component (1%) of Galactic halo dark matter. Ryan et al. (2005) performed one of the first deep photometric surveys of distant UCDs, identifying 28 candidate L and T dwarfs in 135 arcmin2 of deep imaging data obtained with the HST Advanced Camera for Surveys (ACS) instrument, selected by their i \u2212 z colors to a limiting magnitude of z  25. They determined a thin disk vertical scaleheight of \u223c350 pc, similar to prior measurements of deep M dwarf star counts. (Ryan et al. 2011) subsequently identified 17 candidate late-M, L, and T dwarfs in 232 arcmin2 of HST\/Wide Field Camera 3 (WFC3) imaging of the Great Observatories Origins Deep Survey (Giavalisco et al. 2004) using optical and near-infrared color selection, and determined a thin disk vertical scaleheight for these sources of 290 \u00b1 40 pc. Deep ground-based surveys have also identified samples of distant UCDs. Kakazu et al. (2010) identified seven late-L and T dwarfs in 9.3 deg2 of optical and infrared imaging data from the Subaru Suprime-Cam Hawaii Quasar and T dwarf survey to a limiting magnitude of z  23.3, spectroscopically confirming several of the targets. From this small sample, Kakazu et al. (2010) inferred a thin disk vertical scaleheight of \u223c400 pc for brown dwarfs. Sorahana et al. (2019) used the larger (130 deg2) and deeper (z  24) Hyper Suprime-Cam Subaru Strategic Program survey (Aihara et al. 2018) to photometrically identify 3,665 L dwarfs, and inferred an average thin disk vertical scaleheight of 340\u2013420 pc. Carnero Rosell et al. (2019) used multi-band imaging data from the Dark Energy Survey (The Dark Energy Survey Collaboration 2005), combined with photometry from wide-field imaging surveys, to photometrically identify and classify 11,745 L0\u2013T9 dwarfs to a limiting magnitude of z \u2264 22, and estimated a thin disk vertical scaleheight of \u223c450 pc. Recently, Warren et al. (2021) compiled a sample of 34,000 M7-L3 UCDs by searching over a large area of 3,070 deg2 in the Sloan Digital Sky Survey (SDSS; York et al. 2000) and UKIRT Infrared Deep Sky Survey (UKIDSS) down to J = 17.5, and measured a scaleheight of \u223c270 pc. These last three studies, which comprise the largest compilations of UCDs to date, use multiple colors to segregate UCDs from other background sources (Skrzypek et al. 2016).","Citation Text":["Ryan et al. 2011"],"Functions Text":["subsequently identified 17 candidate late-M, L, and T dwarfs in 232 arcmin2 of HST\/Wide Field Camera 3 (WFC3) imaging of the Great Observatories Origins Deep Survey","using optical and near-infrared color selection, and determined a thin disk vertical scaleheight for these sources of 290 \u00b1 40 pc."],"Functions Label":["Background","Background"],"Citation Start End":[[1194,1210]],"Functions Start End":[[1212,1376],[1402,1532]]}
{"Identifier":"2018AandA...613A...3Q__Kelly_et_al._2017_Instance_2","Paragraph":"As a prototypical Seyfert 2 galaxy with starburst at a distance of 14.4 Mpc (1\u2033 = 72 pc, Bland-Hawthorn et al. 1997), NGC 1068 was observed at radio (Greenhill et al. 1996), millimeter (Schinnerer et al. 2000), infrared (Jaffe et al. 2004), optical (Antonucci & Miller 1985), UV (Antonucci et al. 1994), and X-ray (Kinkhabwala et al. 2002). High spatial resolution CO (1\u20130) observations show two molecular spiral arms with a diameter of ~40\u2033 and a northern half-bar, while a CO (2\u20131) map reveals a nuclear ring with two bright knots in the CND region (Schinnerer et al. 2000). The dense gas fraction as traced by HCN (1\u20130) (Tacconi et al. 1994; Helfer & Blitz 1995) and CS (2\u20131) (Tacconi et al. 1997; Takano et al. 2014) in the nuclear region is higher than the two arms. Observations of CO (3\u20132) (Krips et al. 2011; Tsai et al. 2012; Garc\u00eda-Burillo et al. 2014) showed that the difference of molecular gas temperatures between the nuclear region and the two arms was not as large as that of densities. Dozens of molecular lines at millimeter wavelength were detected at CND with single-dish observations (Usero et al. 2004; Nakajima et al. 2011, 2013; Aladro et al. 2013). Moreover, several molecules were detected and resolved toward NGC 1068 with interferometers in the past few years (Tosaki et al. 2017; Kelly et al. 2017; Furuya & Taniguchi 2016; Izumi et al. 2016; Imanishi et al. 2016; Nakajima et al. 2015; Viti et al. 2014; Takano et al. 2014; Garc\u00eda-Burillo et al. 2014, 2016). The molecular gas in the CND region was denser and hotter than that in the starburst ring, while chemical properties in the two regions were also different (Viti et al. 2014). The highest molecular gas temperature was higher than 150 K, and the gas density was above 105 cm\u22123 in the CND region (Viti et al. 2014). The distribution of different species of molecules were also different: CO isotopic species, for instance, were enhanced in the starburst ring, while the shock\/dust related molecules were enhanced in the CND region (Nakajima et al. 2015). The spatially resolved observations showed that the CND region was a complex dynamical system. For instance, the east and west dots were dominated by a fast shock and a slower shock (Kelly et al. 2017), while the dust torus also showed complex kinematics (Garc\u00eda-Burillo et al. 2016). Gas inflow was driven by a past minor merger (Furuya & Taniguchi 2016), while the outflow was AGN driven (Garc\u00eda-Burillo et al. 2014). We conducted adeeper survey of millimeter lines toward the CND region of NGC 1068 with the IRAM 30 m telescope, with the goal to quantify the gas properties in the CND. Compared to previous single-dish observations, our data probe weaker transition lines, which could place more constraints on the physical and chemistry properties of the CND.","Citation Text":["Kelly et al. 2017"],"Functions Text":["The spatially resolved observations showed that the CND region was a complex dynamical system. For instance, the east and west dots were dominated by a fast shock and a slower shock"],"Functions Label":["Background"],"Citation Start End":[[2225,2242]],"Functions Start End":[[2042,2223]]}
{"Identifier":"2022MNRAS.512..439C__Johnson,_Sangwan_&_Shankaranarayanan_2022_Instance_1","Paragraph":"It is still unclear whether this incompatibility is evidence against the spatially flat \u039bCDM model or is caused by unidentified systematic errors in one of the established cosmological probes or by evolution of the parameters themselves with the redshift (Dainotti et al. 2021b, 2022). Newer, alternate cosmological probes could help alleviate this issue. Recent examples of such probes include reverberation-mapped quasar (QSO) measurements that reach to redshift z \u223c 1.9 (Czerny et al. 2021; Khadka et al. 2021a,b; Yu et al. 2021; Zaja\u010dek et al. 2021), H\u2009ii starburst galaxy measurements that reach to z \u223c 2.4 (Mania & Ratra 2012; Ch\u00e1vez et al. 2014; Gonz\u00e1lez-Mor\u00e1n et al. 2019, 2021; Cao, Ryan & Ratra 2020, 2022a; Cao et al. 2021a; Johnson, Sangwan & Shankaranarayanan 2022; Mehrabi et al. 2022), QSO angular size measurements that reach to z \u223c 2.7 (Cao et al. 2017, 2020, 2021a; Ryan, Chen & Ratra 2019; Lian et al. 2021; Zheng et al. 2021), QSO flux measurements that reach to z \u223c 7.5 (Risaliti & Lusso 2015, 2019; Khadka & Ratra 2020a,b, 2021, 2022; Lusso et al. 2020; Yang, Banerjee & \u00d3 Colg\u00e1in 2020; Li et al. 2021; Lian et al. 2021; Luongo et al. 2021; Rezaei, Sol\u00e0 Peracaula & Malekjani 2021; Zhao & Xia 2021),1 and the main subject of this paper, gamma-ray burst (GRB) measurements that reach to z \u223c 8.2 (Amati et al. 2008, 2019; Cardone, Capozziello & Dainotti 2009; Cardone et al. 2010; Samushia & Ratra 2010; Dainotti et al. 2011, 2013a,b; Postnikov et al. 2014; Wang, Dai & Liang 2015; Wang et al. 2016, 2022; Fana Dirirsa et al. 2019; Khadka & Ratra 2020c; Hu, Wang & Dai 2021; Dai et al. 2021; Demianski et al. 2021; Khadka et al. 2021c; Luongo et al. 2021; Luongo & Muccino 2021; Cao et al. 2021a). Some of these probes might eventually allow for a reliable extension of the Hubble diagram to z \u223c 3\u20134, well beyond the reach of Type Ia supernovae. GRBs have been detected to z \u223c 9.4 (Cucchiara et al. 2011), and might be detectable to z = 20 (Lamb & Reichart 2000), so in principle GRBs could act as a cosmological probe to higher redshifts than 8.2.","Citation Text":["Johnson, Sangwan & Shankaranarayanan 2022"],"Functions Text":["Newer, alternate cosmological probes could help alleviate this issue.","Recent examples of such probes include","H\u2009ii starburst galaxy measurements that reach to z \u223c 2.4"],"Functions Label":["Motivation","Background","Background"],"Citation Start End":[[736,777]],"Functions Start End":[[286,355],[356,394],[555,611]]}
{"Identifier":"2015AandA...573A.102B__Fern\u00e1ndez_et_al._1999_Instance_1","Paragraph":"The cumulative total absolute magnitude distribution of the comets obeys N(  HT) \u221d 10\u2212 \u03b1THT. Since we have imposed that the total brightness of the comets scales as BT \u221d fD2, it is easy to show that the slope of the total absolute magnitude distribution, \u03b1T, is equal to the slope of the nuclear absolute magnitude distribution, \u03b1. The latter is related to the cumulative size-frequency distribution, N( >D) \u221d D\u2212 \u03b3, where \u03b3 = 5\u03b1. Even though there is a lot of scatter in the D \u2212 HT diagram caused by variation in f from one comet to the next (Fern\u00e1ndez et al. 1999), there is a clear correlation between D and HT in Fig. 1 which must be caused by the underlying size distribution. For JFCs with diameters between approximately 2 km and 10 km the slope \u03b3 ~ 2 (Meech et al. 2004; Snodgrass et al. 2011), corresponding to \u03b1 = 0.4. The number of JFCs with q  2.5 AU and HT  10.8 is then \\hbox{$N_{\\rm vJFC}=294_{-235}^{+556}$}NvJFC=294-235+556, about three times higher than the number reported in Levison & Duncan (1997) and Di Sisto et al. (2009). This is likely to be a lower limit because of the aforementioned incompleteness. However, the fading alters the mean active lifetime, \u03c4vFJC, as well. From the output of our simulations and using the delayed power law we computed a weighted mean period of JFCs with q  2.5 AU of 7.94 yr and a corresponding active lifetime \\hbox{$\\tau_{\\rm vJFC} =1969^{+7479}_{-1540}$}\u03c4vJFC=1969-1540+7479 yr, lower than previous estimates (Di Sisto et al. 2009; Duncan & Levison 1997; Fern\u00e1ndez et al. 2002). Following Brasser & Morbidelli (2013) the corresponding number of objects in the scattered disc with this updated active lifetime, taking into account the uncertainties in all relevant quantities, is then \\hbox{$N_{\\rm SD}= 5.9^{+2.2}_{-5.1} \\times 10^9$}NSD=5.9-5.1+2.2\u00d7109. That same work computed an Oort cloud population of NOC = (7.6 \u00b1 3.3) \u00d7 1010 for objects with D> 2.3 km. From our new analysis the Oort cloud to scattered disc population ratio turns out to be 13\\hbox{$_{-11}^{+77}$}+77-11, which is consistent with the ratio of 12 \u00b1 1 from simulations (Brasser & Morbidelli 2013). Thus, it is likely that the Oort cloud and scattered disc formed at the same time from the same source, and thus is consistent with a formation during the giant planet instability. ","Citation Text":["Fern\u00e1ndez et al. 1999"],"Functions Text":["Even though there is a lot of scatter in the D \u2212 HT diagram caused by variation in f from one comet to the next"],"Functions Label":["Uses"],"Citation Start End":[[543,564]],"Functions Start End":[[430,541]]}
{"Identifier":"2017MNRAS.471.4286F__Colpi_2014_Instance_1","Paragraph":"Following the first detection by ROSAT (Komossa & Bade 1999; Bade, Komossa & Dahlem 2016), about 50 TDEs have been observed (Komossa 2015) in hard X-ray (Bloom et al. 2011; Burrows et al. 2011; Cenko et al. 2012; Pasham et al. 2015), soft X-ray (Komossa & Bade 1999; Donley et al. 2002; Esquej et al. 2008; Maksym et al. 2010; Saxton et al. 2012, 2017; Bade et al. 2016), UV (Stern et al. 2004; Gezari et al. 2006, 2008, 2009) and optical (van Velzen et al. 2011; Gezari et al. 2012; Arcavi et al. 2014; Chornock et al. 2014; Holoien et al. 2014; Vinko et al. 2015) wavelengths. Some of the observed TDEs exhibit unusual properties. In particular, one of the detected TDEs shows an excess of variability in its light curve (Saxton et al. 2012) which can be explained if the black hole is actually a binary with a mass of 106\u2009M\u2299, mass ratio of 0.1 and semimajor axis of 0.6 milliparsecs (Liu, Li & Komossa 2014). This candidate appears to have one of the most compact orbits among the known SMBH binaries and has overcome the \u2018final parsec problem\u2019 (Colpi 2014). Upon coalescence, it will be a strong source of gravitational wave emission in the sensitivity range of the evolved Laser Interferometer Space Antenna (eLISA). Three other TDEs appear to be very bright in X-rays with peak soft X-ray isotropic luminosity being highly super-Eddington (Burrows et al. 2011; Cenko et al. 2012; Brown et al. 2015), while follow-up observations showed that these events were also associated with bright, compact, variable radio synchrotron emission (Zauderer et al. 2011; Cenko et al. 2012). The observed high X-ray luminosity can be explained if the tidal disruption of stars in these cases powered a highly beamed relativistic jet pointed at the observer (Tchekhovskoy et al. 2014). Based on these three observations, Kawamuro et al. (2016) concluded that 0.0007\u2009per\u2009cent\u201334\u2009per\u2009cent of all TDEs source relativistic jets, while Bower et al. (2013) and van Velzen et al. (2013) estimated that \u227210\u2009per\u2009cent of TDEs produce jetted emission at the observed level. Formation of jets in TDEs is a topic of active research, e.g. works by Metzger, Giannios & Mimica (2012), Mimica et al. (2015) and Generozov et al. (2017).","Citation Text":["Colpi 2014"],"Functions Text":["This candidate appears to have one of the most compact orbits among the known SMBH binaries and has overcome the \u2018final parsec problem\u2019"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1049,1059]],"Functions Start End":[[912,1047]]}
{"Identifier":"2022MNRAS.517.4529B__Zaroubi,_Hoffman_&_Dekel_1999_Instance_1","Paragraph":"The other criteria that we can use to categorize the reconstruction methods is whether the reconstruction is performed using forward-modelling or uses a direct inversion from the data. Inverting non-linear problems from partial, noisy, observations is an ill-posed inverse problem, which makes forward-modelled Bayesian methods particularly suitable for the task of reconstruction of high-dimensional fields. Bayesian reconstruction methods have become increasingly popular in cosmology and have been applied in a range of different applications such as initial conditions reconstruction (Jasche & Wandelt 2013; Modi, Feng & Seljak 2018; Jasche & Lavaux 2019), weak lensing (Fiedorowicz et al. 2022; Porqueres et al. 2021, 2022; Boruah, Rozo & Fiedorowicz 2022), and CMB lensing (Millea et al. 2021; Millea, Anderes & Wandelt 2020). Such methods have also been used for the local velocity field reconstruction. The simplest of such methods uses a Wiener filtering technique (Zaroubi, Hoffman & Dekel 1999). This approach assumes that the density\/velocity field is described as a Gaussian random field and the Wiener filtered reconstruction is the maximum-a-posteriori (MAP) solution for the problem. The Wiener filtering approach has been extended to account for uncertainties and biases in the reconstruction using a constrained realization approach (Hoffman & Ribak 1991; Hoffman, Courtois & Tully 2015; Hoffman et al. 2018; Lilow & Nusser 2021) An alternative way to account for the biases in the reconstruction in Wiener filtering is using the unbiased minimal variance approach (Zaroubi 2002). Another similar approach is the Bayesian hierarchical method, virbius (Lavaux 2016), which is based on the constrained realization approach but accounts for many different systematic effects in its analysis. This approach has been been applied to the Cosmicflows-3 (Tully, Courtois & Sorce 2016) data set by Graziani et al. (2019). A similar reconstruction code, hamlet, was introduced in Valade et al. (2022). However, these methods fail to account for the inhomogeneous Malmquist (IHM) bias which is an important source of systematic error in peculiar velocity analysis. The IHM bias arises from an incorrect assumption on the distribution of peculiar velocity tracers due to neglecting the line-of-sight inhomogeneities.","Citation Text":["Zaroubi, Hoffman & Dekel 1999"],"Functions Text":["The simplest of such methods uses a Wiener filtering technique","This approach assumes that the density\/velocity field is described as a Gaussian random field and the Wiener filtered reconstruction is the maximum-a-posteriori (MAP) solution for the problem."],"Functions Label":["Background","Background"],"Citation Start End":[[975,1004]],"Functions Start End":[[911,973],[1007,1199]]}
{"Identifier":"2020AandA...641A.155V__Elmegreen_et_al._2007_Instance_1","Paragraph":"It has also become evident that the normalization of the MS rapidly increases with redshift: distant galaxies form stars at higher paces than in the local Universe, at fixed stellar mass (e.g., Daddi et al. 2007; Elbaz et al. 2007; Whitaker et al. 2012; Speagle et al. 2014; Schreiber et al. 2015). This trend could be explained by the availability of copious molecular gas at high redshift (Daddi et al. 2010a; Tacconi et al. 2010, 2018; Scoville et al. 2017a; Riechers et al. 2019; Decarli et al. 2019; Liu et al. 2019a), ultimately regulated by the larger accretion rates from the cosmic web (Kere\u0161 et al. 2005; Dekel et al. 2009a). Moreover, higher SFRs could be induced by an increased efficiency of star formation due to the enhanced fragmentation in gas-rich, turbulent, and gravitationally unstable high-redshift disks (Bournaud et al. 2007, 2010; Dekel et al. 2009b; Ceverino et al. 2010; Dekel & Burkert 2014), reflected on their clumpy morphologies (Elmegreen et al. 2007; F\u00f6rster Schreiber et al. 2011; Genzel et al. 2011; Guo et al. 2012, 2015; Zanella et al. 2019). IR-bright galaxies with prodigious SFRs well above the level of the MS are observed also in the distant Universe, but their main physical driver is a matter of debate. While a star formation efficiency (SFE\u2004=\u2004SFR\/Mgas) monotonically increasing with the distance from the main sequence (\u0394MS\u2004=\u2004SFR\/SFRMS, Genzel et al. 2010, 2015; Magdis et al. 2012; Tacconi et al. 2018, 2020) could naturally explain the existence of these outliers, recent works suggest the concomitant increase of gas masses as the main driver of the starbursting events (Scoville et al. 2016; Elbaz et al. 2018). In addition, if many bright starbursting (sub)millimeter galaxies (SMGs, Smail et al. 1997) are indeed merging systems as in the local Universe (G\u00f3mez-Guijarro et al. 2018, and references therein), there are several well documented cases of SMGs hosting orderly rotating disks at high redshift (e.g., Hodge et al. 2016, 2019; Drew et al. 2020), disputing the pure merger scenario. The same definition of starbursts as galaxies deviating from the main sequence has been recently questioned with the advent of high spatial resolution measurements of their dust and gas emission. Compact galaxies with short depletion timescales typical of SBs are now routinely found on the MS, being possibly on their way to leave the sequence (Barro et al. 2017a; Popping et al. 2017; Elbaz et al. 2018; G\u00f3mez-Guijarro et al. 2019; Puglisi et al. 2019; Jim\u00e9nez-Andrade et al. 2019); or galaxies moving within the MS scatter, due to mergers unable to efficiently boost the star formation (Fensch et al. 2017) or owing to gravitational instabilities and gas radial redistribution (Tacchella et al. 2016).","Citation Text":["Elmegreen et al. 2007"],"Functions Text":["Moreover, higher SFRs could be induced by an increased efficiency of star formation due to the enhanced fragmentation in gas-rich, turbulent, and gravitationally unstable high-redshift disks","reflected on their clumpy morphologies"],"Functions Label":["Background","Background"],"Citation Start End":[[961,982]],"Functions Start End":[[636,826],[921,959]]}
{"Identifier":"2016ApJ...820..113J__Malanushenko_et_al._2009_Instance_1","Paragraph":"There have been previous studies that utilized coronal extreme ultraviolet (EUV) images in combination with models to try to more accurately determine the magnetic field. Conlon and Gallagher (Conlon & Gallagher 2010) used Extreme Ultraviolet Imager (EUVI) images to choose the best value of the parameter \u03b1 in an LFFF model of a coronal active region. Aschwanden and co-authors (Aschwanden & Sandman 2010; Aschwanden 2013; Aschwanden & Malanushenko 2013) created NLFFF models of active regions by forward-fitting the underlying magnetograms with sets of buried magnetic monopoles and selecting the corresponding \u03b1 values for each monopole that produced the greatest agreement with coronal loops traced in EUVI images. Malanushenko and co-authors (Malanushenko et al. 2009, 2012) have presented a method for deriving an NLFFF model from sparsely distributed EUV loop observations. These studies have focused on modeling the field in active regions, where the high density results in bright EUV emission and the limited size of the region allows for a reasonable number of free parameters. Additionally, the complexity of these methods suggests that their primary application would be the detailed study of a specific region of interest, rather than casual production of models for programmatic use. In contrast, the purpose of our study has been to develop a method for the fast production of global coronal magnetic field models based on widely available synoptic magnetograms and coronal images. Global models are necessary for studies on a wide array of topics, including the connectivity of active regions (Tadesse et al. 2012; Schrijver et al. 2013), the topology of the corona through the solar cycle (Wang & Sheeley 2003; Platten et al. 2014), and as a larger context for studies of coronal activity in localized regions (Conlon & Gallagher 2010; Schrijver et al. 2013). They are also instrumental for global heliospheric simulation and interpreting in situ measurements by upcoming near-Sun missions, Solar Orbiter (SO) and Solar Probe Plus (SPP).","Citation Text":["Malanushenko et al. 2009"],"Functions Text":["Malanushenko and co-authors","have presented a method for deriving an NLFFF model from sparsely distributed EUV loop observations.","These studies have focused on modeling the field in active regions, where the high density results in bright EUV emission and the limited size of the region allows for a reasonable number of free parameters. Additionally, the complexity of these methods suggests that their primary application would be the detailed study of a specific region of interest, rather than casual production of models for programmatic use. In contrast, the purpose of our study has been to develop a method for the fast production of global coronal magnetic field models based on widely available synoptic magnetograms and coronal images."],"Functions Label":["Background","Background","Compare\/Contrast"],"Citation Start End":[[748,772]],"Functions Start End":[[719,746],[780,880],[881,1497]]}
{"Identifier":"2015AandA...584A..75V__Essen_et_al._(2014)_Instance_7","Paragraph":"The data presented here comprise quasi-simultaneous observations during secondary eclipse of WASP-33 b around the V and Y bands. The predicted planet-star flux ratio in the V-band is 0.2 ppt, four times lower than the accuracy of our measurements. Therefore, we can neglect the planet imprint and use this band to measure the stellar pulsations, and most specifically to tune their current phases (see phase shifts in von Essen et al. 2014). Particularly, our model for the light contribution of the stellar pulsations consists of eight sinusoidal pulsation frequencies with corresponding amplitudes and phases. Hence, to reduce the number of 24 free parameters and the values they can take, we use prior knowledge about the pulsation spectrum of the star that was acquired during von Essen et al. (2014). As the frequency resolution is 1\/\u0394T (Kurtz 1983), 3.5 h of data are not sufficient to determine the pulsations frequencies. Therefore, during our fitting procedure we use the frequencies determined in von Essen et al. (2014) as starting values plus their derived errors as Gaussian priors. As pointed out in von Essen et al. (2014), we found clear evidences of pulsation phase variability with a maximum change of 2 \u00d7 10-3 c\/d. In other words, as an example after one year time a phase-constant model would appear to have the correct shape with respect to the pulsation pattern of the star, but shifted several minutes in time. To account for this, the eight phases were considered as fitting parameters. The von Essen et al. (2014) photometric follow-up started in August, 2010, and ended in October, 2012, coinciding with these LBT data. We then used the phases determined in von Essen et al. (2014) during our last observing season as starting values, and we restricted them to the limiting cases derived in Sect. 3.5 of von Essen et al. (2014), rather than allowing them to take arbitrary values. The pulsation amplitudes in \u03b4 Scuti stars are expected to be wavelength-dependent (see e.g. Daszy\u0144ska-Daszkiewicz 2008). Our follow-up campaign and these data were acquired in the blue wavelength range. Therefore the amplitudes estimated in von Essen et al. (2014), listed in Table 1, are used in this work as fixed values. This approach would be incorrect if the pulsation amplitudes would be significantly variable (see e.g., Breger et al. 2005). Nonetheless, the short time span of LBT data, and the achieved photometric precision compared to the intrinsically low values of WASP-33\u2019s amplitudes, make the detection of any amplitude variability impossible. ","Citation Text":["von Essen et al. (2014)"],"Functions Text":["and we restricted them to the limiting cases derived in Sect. 3.5 of","rather than allowing them to take arbitrary values."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1830,1853]],"Functions Start End":[[1761,1829],[1855,1906]]}
{"Identifier":"2021MNRAS.503..354G__Hou_&_Han_2014_Instance_2","Paragraph":"The spatial distribution of OB stars and associations, young long-period Cepheids and open clusters, star-forming regions, H\u2009ii regions, interstellar dust, and giant molecular and neutral gas clouds in the solar vicinity that have been in existence generally \u03c4 \u2272 108 yr is known to correlate with the location of the inner Sagittarius, the closest Orion, and outer Perseus spiral arm segments. (The distances for the vast majority of these spiral tracers have been determined in the literature with trigonometric or photometric methods.) The Sun is situated at the inner edge of the Orion arm (Levine et al. 2006; Hou & Han 2014; Nakanishi & Sofue 2016; Xu et al. 2018, 2021; Lallement et al. 2019; Reid et al. 2019; Skowron et al. 2019; Cantat-Gaudin et al. 2020; Fig. 2 above).3 These three spatial features nearby to the Sun appear to form part of the global spiral structure in the Galaxy. Contrary, the objects of older population with larger random velocities, for instance, main-sequence A\u2013K stars or the oldest Cepheids and open clusters, do not currently follow the exact location of those arms (e.g. Cantat-Gaudin et al. 2020, fig. 8 therein; Griv et al. 2020, fig. 7 therein). The latter can be explained by the difference in rotation velocity between the spiral density waves and the objects. Investigating the velocity field of Xu et al.\u2019s (2018) O and early B-type stars in the framework of the Lin\u2013Shu density-wave proposal, we also found that the Sun lies within the Orion arm, at the inner edge of this spiral. The radial distance from the Sun to the centre of the Orion arm is \u22480.2 kpc in the direction of the Galactic anticentre, the centre of the Sagittarius arm is \u22481.8 kpc from the Sun in the direction of the GC, and the width of the arms is \u22480.5 kpc. The radial distance between the centres of the Orion and Sagittarius arms near the Sun is \u03bbrad \u2248 2 kpc (cf. Hou & Han 2014; Wu et al. 2014; Bovy et al. 2015). As for us, the nearest Orion spiral arm forms part of the dominant density-wave structure of the system.","Citation Text":["Hou & Han 2014"],"Functions Text":["The radial distance between the centres of the Orion and Sagittarius arms near the Sun is \u03bbrad \u2248 2 kpc (cf."],"Functions Label":["Background"],"Citation Start End":[[1883,1897]],"Functions Start End":[[1775,1882]]}
{"Identifier":"2019ApJ...871...82G__Tarr_et_al._2014_Instance_1","Paragraph":"By considering the UV emission at several wavelengths, we have been able to reconstruct the evolution of the EFR at different layers using radiance maps. The optically thin view provided by the O i \u03bb1355.6 line (Lin & Carlsson 2015) images the AFS that formed above the EFR. AFSs are typically observed in absorption in the chromospheric layers during flux emergence and reflect the serpentine nature of the emerging fields (Bruzek 1980; Spadaro et al. 2004; Zuccarello et al. 2005; Murabito et al. 2017). They also can reconnect with the ambient field (e.g., Zuccarello et al. 2008; Tarr et al. 2014; Su et al. 2018). Here the AFS appears in emission when observed in the O i \u03bb1355.6 line, being brighter than the background. Some brightness enhancements are found along the AFS, suggesting the occurrence of small-scale energy release events related to the reconnection of the emerging field lines with the ambient field (see, e.g., Huang et al. 2018). The view of optically thick Mg ii k and C ii \u03bb1335 lines (Leenaarts et al. 2013; Rathore et al. 2015) offers the possibility to see the counterpart of the AFS in the upper chromosphere. Threads observed in absorption, departing from the AFS, cover the whole structure. It should be stressed that the reconstructed radiance maps image the EFR along the entire IRIS scan: this means that the individual threads at every single slit position along the x-direction were only transiently observed. From the comparison with IRIS SJIs in the 2796 \u212b passband, it can be deduced that threads correspond to the surges, being viewed as individual ejections in slice imaging. The westernmost part of the UV burst, being not covered by threads, is seen as a compact brightening in the Mg ii k and C ii \u03bb1335 radiance maps, in the latter having a higher contrast with respect to the background. Conversely, the optically thin emission in the Si iv \u03bb\u03bb1402 and 1394 lines clearly illustrates the UV burst and other bright knots in the EFR area. Apparently, threads have no counterpart at Si iv \u03bb\u03bb1402 and 1394 formation heights, although fluffy elongated structures appear in the northern part of the UV burst. Finally, dark elongated absorption structures are distinctly visible in the SDO\/AIA 193 filtergrams, which are highly reminiscent of the threads seen by IRIS. To their west, a compact brightening cospatial to the UV burst is also detected, being not obscured by such structures.","Citation Text":["Tarr et al. 2014"],"Functions Text":["They also can reconnect with the ambient field (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[584,600]],"Functions Start End":[[506,559]]}
{"Identifier":"2017AandA...603A.107A__Leconte_et_al._(2015)_Instance_1","Paragraph":"The second equation of our system is the conservation of mass, (12)\\begin{equation} \\dfrac{\\partial \\rho}{\\partial t} + \\boldsymbol{\\nabla}. \\left( \\rho {\\vec V} \\right) = 0, \\end{equation}\u2202\u03c1\u2202t+\u2207.\u03c1V=0,which, in spherical coordinates, writes (13)\\begin{equation} \\dfrac{\\partial \\delta \\rho}{\\partial t } + \\frac{1}{r^2} \\dfrac{\\partial }{\\partial r} \\left( r^2 \\rho_0 V_r \\right) + \\frac{\\rho_0}{r \\sin \\theta} \\left[ \\dfrac{\\partial }{\\partial \\theta} \\left( \\sin \\theta V_\\theta \\right) + \\dfrac{\\partial V_\\varphi}{\\partial \\varphi} \\right] = 0. \\label{conservation_masse} \\end{equation}\u2202\u03b4\u03c1\u2202t+1r2\u2202\u2202r(r2\u03c10Vr)+\u03c10rsin\u03b8\u2202\u2202\u03b8sin\u03b8V\u03b8+\u2202V\u03d5\u2202\u03d5=0.The thermal forcing (J) appears on the right-hand side of the linearized heat transport equation (see CL70, Gerkema & Zimmerman 2008) given by (14)\\begin{equation} \\frac{1}{\\Gamma_1 p_0} \\dfrac{\\partial \\delta p}{\\partial t} - \\frac{1}{\\rho_0} \\dfrac{\\partial \\delta \\rho}{\\delta t} + \\frac{N^2}{g} V_r = \\kappa \\frac{\\rho_0}{p_0} \\left[ J - J_{\\rm rad} \\right], \\label{transport_chaleur_1} \\end{equation}1\u03931p0\u2202\u03b4p\u2202t\u22121\u03c10\u2202\u03b4\u03c1\u03b4t+N2gVr=\u03ba\u03c10p0J\u2212Jrad,where \\hbox{$ \\kappa = \\left( \\Gamma_1 - 1 \\right)\/\\Gamma_1 $}\u03ba=\u03931(\u22121)\/\u03931 and Jrad is the power per mass unit radiated by the atmosphere, supposed to behave as a grey body. We consider that Jrad\u221d\u03b4T. This hypothesis is known as \u201cNewtonian cooling\u201d and was used by Lindzen & McKenzie (1967) to introduce radiation analytically in the classical theory of atmospheric tides (see also Dickinson & Geller 1968). Physically, it corresponds to the case of an optically thin atmosphere in which the flux emitted by a layer propagates upwards or downwards without being absorbed by the other layers. In optically thick atmospheres, such as on Venus (Lacis 1975), this physical condition is not verified. Indeed, because of a stronger absorption, the power emitted by a layer is almost totally transmitted to the neighbourhood. Therefore, this significant thermal coupling within the atmospheric shell should ideally be taken into account in a rigorous way, which would lead to great mathematical difficulties (e.g. complex radiative transfers, Laplacian operators) in our analytical approach. However, recent numerical simulations of thermal tides in optically thick atmospheres by Leconte et al. (2015) show behaviour of the flow that is in good agreement with a model that uses radiative cooling. Therefore, in this work we assumeNewtonian cooling as a first model of the action of radiation on atmospheric tides. Newtonian cooling brings a new characteristic frequency, denoted \u03c30, which we call radiative frequency and which depends on the thermal capacity of the atmosphere. The radiative power per unit mass is thus written (15)\\begin{equation} J_{\\rm rad} = \\frac{p_0 \\sigma_0}{\\kappa \\rho_0 T_0} \\delta T. \\label{phirad} \\end{equation}Jrad=p0\u03c30\u03ba\u03c10T0\u03b4T.Like the basic fields p0, \u03c10, and T0, the radiative frequency varies with r and defines the transition between the dynamical regime, where the radiative losses can be ignored, and the radiative regime, where they predominate in the heat transport equation. Assuming that the radiative emission of the gas is proportional to the local molar concentration C0 = \u03c10\/M, it is possible to express Jrad and \u03c30 as functions of the physical parameters of the fluid (Appendix D), (16)\\begin{equation} J_{\\rm rad} = \\frac{8 \\epsilon_a}{M} \\mathscr{S} T_0^3 \\delta T \\label{phirad2} \\end{equation}Jrad=8\u03f5aMST03\u03b4Tand (17)\\begin{equation} \\sigma_0 \\left( r \\right) = \\frac{8 \\kappa \\epsilon_a \\mathscr{S} }{\\mathscr{R}_{\\rm GP} } T_0^3, \\label{sigma0} \\end{equation}\u03c30(r)=8\u03ba\u03f5aSRGPT03,the parameter \u03f5a being an effective molar emissivity coefficient of the gas and S= 5.670373 \u00d7 10-8 W m-2 K-4 the Stefan-Boltzmann constant (Mohr et al. 2012). The substitution of Eq. (15) in Eq. (14) yields (18)\\begin{equation} \\frac{1}{\\Gamma_1 P_0} \\dfrac{\\partial \\delta p}{\\partial t} - \\frac{1}{\\rho_0} \\dfrac{\\partial \\delta \\rho}{\\delta t} + \\frac{N^2}{g} V_r = \\kappa \\frac{\\rho_0}{p_0} J - \\sigma_0 \\frac{\\delta T}{T_0}\\cdot  \\label{transport_chaleur_2} \\end{equation}1\u03931P0\u2202\u03b4p\u2202t\u22121\u03c10\u2202\u03b4\u03c1\u03b4t+N2gVr=\u03ba\u03c10p0J\u2212\u03c30\u03b4TT0\u00b7Finally, the system is closed by the perfect gas law (19)\\begin{equation} \\frac{\\delta p}{p_0} = \\frac{\\delta T}{T_0} + \\frac{\\delta \\rho}{\\rho_0}\\cdot \\label{GPlaw} \\end{equation}\u03b4pp0=\u03b4TT0+\u03b4\u03c1\u03c10\u00b7Substituting Eq. (19) in Eq. (18), we eliminate the unknown \u03b4T, and obtain (20)\\begin{equation} \\frac{1}{\\Gamma_1 p_0} \\left( \\dfrac{\\partial \\delta p}{\\partial t} \\! +\\! \\Gamma_1 \\sigma_0 \\delta p \\right) + \\frac{N^2}{g} \\dfrac{\\partial \\xi_r}{\\partial t} = \\frac{ \\kappa \\rho_0}{p_0} J + \\frac{1}{\\rho_0} \\left( \\dfrac{\\partial \\delta \\rho}{\\partial t}\\! +\\! \\sigma_0 \\delta \\rho \\right). \\label{transport_chaleur_3} \\end{equation}1\u03931p0\u2202\u03b4p\u2202t+\u03931\u03c30\u03b4p+N2g\u2202\u03ber\u2202t=\u03ba\u03c10p0J+1\u03c10\u2202\u03b4\u03c1\u2202t+\u03c30\u03b4\u03c1.Because of the rotating motion of the perturber in the equatorial frame (RE;T), a tidal perturbation is supposed to be periodic in time (t) and longitude (\u03d5). So, any perturbed quantity f of our model can be expanded in Fourier series of t and \u03d5(21)\\begin{equation} f = \\sum_{m,\\sigma} f^{m,\\sigma} \\left( \\theta , r \\right) e^{i \\left( \\sigma t + m \\varphi \\right)}, \\label{perturbation} \\end{equation}f=\u2211m,\u03c3fm,\u03c3(\u03b8,r)ei\u03c3t+m\u03d5,the parameter \u03c3 being the tidal frequency of a Fourier component and m its longitudinal degree1. We also introduce the spin parameter (22)\\begin{equation} \\nu \\left( \\sigma \\right) = \\frac{2 \\Omega}{\\sigma}, \\label{nu} \\end{equation}\u03bd(\u03c3)=2\u03a9\u03c3,which defines the possible regimes of tidal gravito-inertial waves:","Citation Text":["Leconte et al. (2015)"],"Functions Text":["However, recent numerical simulations of thermal tides in optically thick atmospheres by","show behaviour of the flow that is in good agreement with a model that uses radiative cooling.","Therefore, in this work we assumeNewtonian cooling as a first model of the action of radiation on atmospheric tides."],"Functions Label":["Similarities","Similarities","Uses"],"Citation Start End":[[2251,2272]],"Functions Start End":[[2162,2250],[2273,2367],[2368,2484]]}
{"Identifier":"2019AandA...629A..63J__XIX_2015_Instance_1","Paragraph":"Our knowledge of magnetic fields in molecular clouds is based mainly on light polarisation, the optical and near-infrared (NIR) observations of background stars (Goodman et al. 1995; Whittet et al. 2001; Pereyra & Magalh\u00e3es 2004; Alves et al. 2008; Chapman et al. 2011; Cox et al. 2016; Neha et al. 2018; Kandori et al. 2018), and the polarised dust emission at far-infrared (FIR), sub-millimetre, and radio wavelengths (Ward-Thompson et al. 2000; Koch et al. 2014; Matthews et al. 2014; Fissel et al. 2016; Pattle et al. 2017). The Planck survey provides a large amount of data for polarisation studies at cloud scales (Planck Collaboration Int. XX 2015; Planck Collaboration Int. XIX 2015; Planck Collaboration Int. XXXIII 2016). The Planck data have been used especially to study the polarisation fraction and the correlations in the relative morphology of column density and magnetic field structures (Planck Collaboration Int. XX 2015; Planck Collaboration Int. XXXIII 2016; Malinen et al. 2016; Soler et al. 2016; Alina et al. 2019). Particularly, the drop of polarisation fraction p towards PGCC clumps has been observed with high significance in the Planck 353 GHz data (Alina et al. 2019; Ristorcelli et al, in prep.). The variations ofp are interesting because they are related to the configuration of the magnetic fields in clumps and cores during the star formation process. However, p is also affected by variations in the efficiency of the grain alignment, as predicted, for example, by the theory of radiative torque alignment (RAT; Lazarian et al. 1997; Cho & Lazarian 2005; Hoang & Lazarian 2014) and demonstrated by numerical simulations (Pelkonen et al. 2009; Brauer et al. 2016; Reissl et al. 2018). These suggest that high optical depths and more frequent gascollisions should significantly reduce the grain alignment and thus the polarised emission observable from within the clumps. ThePGCC provides a statistically significant sample to study these questions observationally, although the Planck resolution limits the investigations to structures that are typically much larger than an individual cloud core. However, polarisation of selected PGCCs has already been studied at higher resolution with the SCUBA-2 POL-2 instrument at JCMT (Liu et al. 2018b,c; Juvela et al. 2018c), and many more will be covered by ongoing surveys (Ward-Thompson et al. 2017).","Citation Text":["Planck Collaboration Int. XIX 2015"],"Functions Text":["The Planck survey provides a large amount of data for polarisation studies at cloud scales"],"Functions Label":["Background"],"Citation Start End":[[656,690]],"Functions Start End":[[529,619]]}
{"Identifier":"2017AandA...604A.118T__Pilkington_et_al._(2012)_Instance_1","Paragraph":"A large effort has been invested in understanding the chemical patterns of galaxies using analytical and numerical chemical modelling (e.g. Brook et al. 2007; Calura et al. 2012; Moll\u00e1 et al. 2015). In particular, hydrodynamical simulations provide the chemical enrichment of baryons as galaxies are assembled in a cosmological framework, opening the possibility of understanding the interplay of different physical processes in the non-linear regime of evolution (e.g. Mosconi et al. 2001; Lia et al. 2002; Wiersma et al. 2009). Pilkington et al. (2012) carried out a comparison of the metallicity gradients in the ISM (traced by young SPs) obtained by different numerical and analytical models. Gibson et al. (2013) analysed different models of SN feedback schemes reporting different evolution for the gas-phase metallicity gradients. No evolution was reported when an enhanced SN model was used. Tissera et al. (2016b) studied the gas-phase metallicity gradients in discs and the specific star formation of the galaxies. For gas-phase metallicities, the simulated galaxies showed a correlation with stellar mass, which was erased when the metallicity gradients were renormalised by the effective radius, reproducing observations by Ho et al. (2015, and references there in). They also found indications of a correlation between the abundance slopes and the specific star formation rate (sSFR) in agreement with observational findings (Stott et al. 2014). As a function of redshift, the metallicity gradients of the gas-phase disc components were found to be more negative at higher redshift, principally for lower stellar-mass galaxies. The fraction of galaxies with positive metallicity gradients increased with increasing redshift and were found to be associated to mergers and interactions. The trends found for the gas-phase metallicity gradients by Tissera et al. (2016b) are relevant to the discussion of this paper because we are using the same set of simulated galaxies. ","Citation Text":["Pilkington et al. (2012)"],"Functions Text":["carried out a comparison of the metallicity gradients in the ISM (traced by young SPs) obtained by different numerical and analytical models"],"Functions Label":["Background"],"Citation Start End":[[530,554]],"Functions Start End":[[555,695]]}
{"Identifier":"2017AandA...600A..67P__Nakariakov_et_al._2003_Instance_1","Paragraph":"Many theoretical models have been proposed to explain the generation of QPPs. The most elaborated model of QPPs considers MHD oscillations, which affect almost all aspects of the flare emission generation. Indeed, QPPs are involved in triggering the magnetic reconnection, modulating the reconnection rate, accelerating and transporting non-thermal electrons, and changing the physical conditions in emitters (Nakariakov & Melnikov 2009). Other models are based on a sandpile system with self-organized critical states (Lu & Hamilton 1991; Baiesi et al. 2008), the quasi-stabilized system of non-linear plasmas governed by an oscillatory phase of wave-wave or wave-particle interactions (Aschwanden 1987). Different MHD oscillation modes have been identified to be responsible for QPPs in a single flaring loop (Nakariakov et al. 2003; Melnikov et al. 2005; Warmuth et al. 2005; Inglis et al. 2008; Kupriyanova et al. 2010, 2013; Kim et al. 2012). Kolotkov et al. (2015) studied the QPPs of the microwave emission generated in a X3.2-class solar flare. They found three well-defined intrinsic modes with mean periods of 15, 45, and 100s. These authors proposed that the 100 s and 15s modes are likely to be associated with fundamental kink and sausage modes of the flaring loop, respectively. The 100 s oscillations could also be caused by the fundamental longitudinal mode. The 45 s mode, on the other hand, could be the second standing harmonic of the kink mode. Inglis & Nakariakov (2009) reported a multi-periodic oscillatory event with three distinct periods, namely 28s, 18s, and 12s. They argued that the cause of this multi-periodic event is likely to be a kink mode that periodically triggers magnetic reconnection. Similar QPPs could be generated by different mechanisms. To discover, understand, and distinguish these mechanisms correctly, detailed and multi-wavelength observations are required. Within the framework of the EU FP7-project SOLSPANET1, we are developing a solar and space-weather knowledge base that will allow such extensive and detailed studies. ","Citation Text":["Nakariakov et al. 2003"],"Functions Text":["Different MHD oscillation modes have been identified to be responsible for QPPs in a single flaring loop"],"Functions Label":["Background"],"Citation Start End":[[812,834]],"Functions Start End":[[706,810]]}
{"Identifier":"2017MNRAS.472..205S__Becerra_et_al._2015_Instance_1","Paragraph":"Throughout the initial collapse the halo structure is well approximated by ellipsoidal collapse models. We therefore explore radial profiles of various physical quantities to extract information about the galactic environment. The density is illustrated in Fig. 2, which is reasonably approximated by a broken power law over the range r \u2208 (10\u22123, 103)\u2009pc. The break radius is due to the choice of the density resolution threshold, or equivalently, the limited resolution implies the pre-formation profile corresponds to roughly a dynamical time, as evaluated at the maximum resolved density, of \u223c10\u2009kyr prior to the formation of the protostar. If the evolution proceeds under self-similar, isothermal collapse then the break in the profile will shift to smaller scales, eventually reaching the radius of the protostar (Abel, Bryan & Norman 2002; Becerra et al. 2015). We note that secondary infall and accretion results in a density distribution that is steeper than the \u03c1 \u221d r\u22122 profile produced by violent relaxation. We find a power-law scaling of \u03c1 \u2248 r\u2212 7\/3 in good agreement with Bertschinger (1985) where \u03c1 \u221d r\u22129\/4 and Wise, Turk & Abel (2008) where \u03c1 \u221d r\u221212\/5, except in the core where the slope flattens off to \u03c1 \u221d r\u22120.4 around r \u2248 0.3\u2009pc. During the initial collapse and formation, the gas remains neutral with only a small abundance of free protons. Specifically, the ionization fraction, $x_{\\rm {H \\small {II}}} \\equiv n_{\\rm {H \\small {II}}} \/ n_{\\rm H}$, is typically of order 10\u22125 to 10\u22123 depending on the density, such that $n_{\\rm {H \\small {II}}} \\sim 0.1\\,{\\rm cm}^{-3}\\,(r\/10\\,{\\rm pc})^{-5\/3}$. As expected for direct collapse, the central region has access to at least 105\u2009M\u2299 of gas within a radius of \u22481.3\u2009pc. More broadly, the enclosed baryonic mass is $M_{<r} \\approx 4 \\pi \\int _0^r \\rho (r) r^2 {\\rm d}r$ (for 1D profiles), calculated explicitly as Mr \u2261 \u2211r\u03c1iVi where the sum is over all Voronoi cells within a radius r, and the subscript denotes cell quantities for density \u03c1i and volume Vi. For convenience and completeness, in Table 1 we provide radial scaling relations for several relevant quantities, calculated as mass-weighted averages within shell volumes, i.e.\n(1)\r\n\\begin{eqnarray}\r\n\\langle q \\rangle _{\\rm shell} \\equiv \\frac{\\int q \\rho \\, {\\rm d}V}{\\int \\rho \\, {\\rm d}V} \\approx \\frac{\\sum q_i \\rho _i V_i}{\\sum \\rho _i V_i} \\,,\r\n\\end{eqnarray}\r\nwhere the discretized version sums over all cells within each shell. Similarly, we obtain mass-weighted line of sight averages with\n(2)\r\n\\begin{eqnarray}\r\n\\langle q \\rangle _{\\rm LOS} \\equiv \\frac{\\int q \\rho \\, {\\rm d}\\ell }{\\int \\rho \\, {\\rm d}\\ell } \\approx \\frac{\\sum q_i \\rho _i \\Delta \\ell _i}{\\sum \\rho _i \\Delta \\ell _i} \\,,\r\n\\end{eqnarray}\r\nwhere the integral is along radial rays within each shell and the summation is the discretized version calculated by ray tracing.","Citation Text":["Becerra et al. 2015"],"Functions Text":["If the evolution proceeds under self-similar, isothermal collapse then the break in the profile will shift to smaller scales, eventually reaching the radius of the protostar"],"Functions Label":["Uses"],"Citation Start End":[[845,864]],"Functions Start End":[[643,816]]}
{"Identifier":"2021ApJ...910..120K__Schlafly_et_al._2018_Instance_1","Paragraph":"Many of these objects have been confirmed as DNe through an identification spectrum or a measurement of the superhump period, but we find that 27 sources were not confirmed through any method. Quiescent multiband photometry of these remaining candidates can provide insights into the distance and ultimately constrain the luminosity of the transient. In the Galactic plane, a candidate that is relatively blue is likely close by and therefore will have a lower luminosity at peak brightness, suggesting a DN outburst. Conversely, a candidate that is brighter in redder filters is likely reddened by dust, implying a higher luminosity and a CN outburst. This strategy\u2014identifying highly reddened quiescent counterparts\u2014is only possible for candidates within a few degrees of the Galactic plane and that can be securely matched to multiband optical catalogs. As described in Section 2.4, we find quiescent counterparts in Pan-STARRS and also add in coverage of southerly declinations by using the DeCAPS catalog (Schlafly et al. 2018; cross-matching for both catalogs is as explained in Appendix A.3). We use the griz photometry to estimate reddening by fitting the observed spectral energy distributions of the CVs in question to dereddened SDSS CV colors from Kato et al. (2012) by varying the amount of reddening according to extinction laws (Cardelli et al. 1989; Mathis 1990). For each of our candidates, this results in a distribution of extinction values that are plugged into the three-dimensional all-sky extinction map stitched together in Bovy et al. (2016) to find a range of plausible distances. This strategy only worked for 19 of the candidates: those that we were able to cross-match and in fields close to the plane, where a large amount of dust can constrain the distance. We find that all 19 are consistent with the peak luminosity of a dwarf nova, and none are consistent with the peak luminosity of a classical nova. This analysis will also be useful to shed light on the nature of CV candidates that are discovered in the future, so it is made publicly available as an iPython notebook at https:\/\/github.com\/amkawash\/CV_colors_luminosity.\n","Citation Text":["Schlafly et al. 2018"],"Functions Text":["As described in Section 2.4, we find quiescent counterparts in Pan-STARRS and also add in coverage of southerly declinations by using the DeCAPS catalog (","; cross-matching for both catalogs is as explained in Appendix A.3)."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1011,1031]],"Functions Start End":[[857,1011],[1031,1099]]}
{"Identifier":"2022AandA...666A..95H__Hartman_et_al._2022_Instance_1","Paragraph":"Scaling relations for the core radii rc, core densities \u03b4c, and core masses Mc as functions of the total halo mass M200 were fitted to the simulated halo populations, which largely agree with hydrostatic considerations of the halo cores where rc is nearly constant, as well as velocity dispersion tracing in the halo envelope, \n\n\n\n\nv\nc\n2\n\n\u2248\nv\n\n200\n2\n\n\n\n$ v^2_{\\rm c} \\approx v^2_{200} $\n\n\n\n. However, these trends do not agree with those obtained by fitting the Burkert profile to nearby galaxies in the SPARC dataset and the classical Milky Way dSphs. This poses an issue for SIBEC-DM with Rc\u2004\u2273\u20041 kpc and a largely CDM-like matter power spectrum at late times (Harko 2011; Harko & Mocanu 2012; Velten & Wamba 2012; Freitas & Gon\u00e7alves 2013; Bettoni et al. 2014; de Freitas & Velten 2015; Hartman et al. 2022), as was used in our simulations, although these scenarios are not well-motivated. For SIBEC-DM given by the field Lagrangian in Eq. (1), the self-interaction is constrained to Rc\u2004\u20041 kpc, otherwise an early radiation-like period and a large comoving Jeans\u2019 length washes out too much structure to be consistent with observations (Shapiro et al. 2021; Hartman et al. 2022). In fact, Shapiro et al. (2021) found by using constraints on FDM as a proxy for SIBEC-DM, and matching their transfer function cut-offs and HMFs, that the SIBEC-DM self-interaction should be as low as Rc\u2004\u223c\u200410 pc to not be in conflict with observations. We were unable to probe SIBEC-DM with initial conditions and parameters consistent with the Lagrangian in Eq. (1), since the large gap between the halo cores and the cut-off scale requires both a large simulation box and very high spatial resolution. It should be noted that our SIBEC-DM-only simulations do provide a better agreement with the slopes in observed scaling relations than FDM. In particular, FDM simulations generally find \n\n\n\n\nM\nc\n\n\u223c\n\nM\n\n200\n\n\u03b3\n\n\n\n$ M_{\\mathrm{c}}\\sim M_{200}^{\\gamma} $\n\n\n with 1\/3\u2004\u2004\u03b3\u2004\u20040.6, while we find \u03b3\u2004\u2248\u20040.75, which is closer to the observed \u03b3\u2004\u2248\u20041.1. Additionally, FDM halos have core radii that generally decrease with the halo mass, while we find a slightly increasing trend due to larger halos experiencing more thermal heating, although not as steep as in the SPARC dataset and the Milky Way dSphs.","Citation Text":["Hartman et al. 2022"],"Functions Text":["However, these trends do not agree with those obtained by fitting the Burkert profile to nearby galaxies in the SPARC dataset and the classical Milky Way dSphs. This poses an issue for SIBEC-DM with Rc\u2004\u2273\u20041 kpc and a largely CDM-like matter power spectrum at late times","as was used in our simulations, although these scenarios are not well-motivated."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[789,808]],"Functions Start End":[[392,660],[811,891]]}
{"Identifier":"2021ApJ...919..140S__Bartos_et_al._2017_Instance_1","Paragraph":"Resonant dynamical friction may have applications beyond the relaxation of IMBHs examined in this paper. It may affect all objects in stellar clusters much more massive than the individual constituents of the disk, if present, including massive stars, stellar mass black holes (BHs), or the center of mass of massive binaries. Furthermore, it is also expected to operate in any type of disk with a high number of particles, including active galactic nucleus (AGN) accretion disks. Previously, it has been argued that stars and BHs crossing the disk on low-inclination orbits get captured by Chandrasekhar dynamical friction into the disk (Bartos et al. 2017; Panamarev et al. 2018; Tagawa et al. 2020). An interesting implication is that, if BHs settle into the disk, they interact dynamically and form BH\u2013BH binaries efficiently, and frequent dynamical interactions and gas effects drive the BHs to merger, producing gravitational waves (GWs) detectable by LIGO, VIRGO, and KAGRA (McKernan et al. 2014, 2018; Bartos et al. 2017; Leigh et al. 2018; Yang et al. 2019; Tagawa et al. 2020, 2021; Samsing et al. 2020). Mergers are also facilitated by Lidov\u2013Kozai oscillations in anisotropic systems (Heisler & Tremaine 1986; Petrovich & Antonini 2017; Hamilton & Rafikov 2019). The results in this paper show that resonant dynamical friction may accelerate the capture of objects in the accretion disks by a factor proportional to the SMBH mass over the local disk mass for large orbital inclinations. Pressure and viscosity in a gaseous disk do not inhibit the orbit-averaged torque from the IMBH, which leads to realignment and the warping of the disk (Bregman & Alexander 2012). Thus, RDF may efficiently catalyze the alignment of the orbital planes of BHs even in low-luminosity AGN or Seyfert galaxies with relatively small disk masses, which may not be possible for Chandrasekhar dynamical friction. In fact, this mechanism extends the scope of the \u201cAGN merger channel\u201d for GW source populations even beyond low-luminosity AGN and Seyfert galaxies, as it may organize BHs into disks also in nonactive galaxies with nuclear stellar disks.","Citation Text":["Bartos et al. 2017"],"Functions Text":["Previously, it has been argued that stars and BHs crossing the disk on low-inclination orbits get captured by Chandrasekhar dynamical friction into the disk"],"Functions Label":["Background"],"Citation Start End":[[639,657]],"Functions Start End":[[481,637]]}
{"Identifier":"2022MNRAS.516.4833J__Munday_et_al._2020_Instance_1","Paragraph":"Bond & Ciardullo (2018) identify a periodicity of 2.06 d in their data, with a similar period clearly evident in the newly acquired data. However, no single period and time of first minima could be found that fits all the data (perhaps not unexpected given that the variability is clearly not constant, see Section 2.3 for further discussion). In order to better constrain the period, as well as the changes between observing epochs, we focus on the periodicitiy of all the data taken between 2011 and 2015 (shown in the lower left panel of Fig. 1). Given the nature of the variability, we measure the period via the reduced \u03c72 of a sinusoidal fit as a function of frequency (ideally suited to sinusoidal variability, e.g. Horne, Wade & Szkody 1986; Munday et al. 2020). The resulting periodogram is shown in Fig. 2. Two strong minima in the reduced \u03c72 are found, the strongest at 2.061 \u00b1 0.005 d (consistent with the period of Bond & Ciardullo 2018) and a second slightly shallower at 1.938 \u00b1 0.005 d. The data from 2010, combining that of Bond & Ciardullo (2018) and our own from SAAO, phases well on the first period but not the second, lower period, likely indicating that this is an alias. As such, we conclude that the rotation period is indeed \u22482.06 d and that there is no strong evidence for a change in period before and after the phase of negligible variability in early 2011. There is, however, a clear shift in phase between the 2010 data and the later data used to derive the period (see Fig. 2). There is no statistically significant evidence for a shift between the 2011\u20132015 data and the later data from 2022, although the lack of data in the intervening seven years as well as the uncertainty on the period mean that we cannot exclude a phase shift between these two data sets. The 2022 data does, however, appear to have a larger photometric amplitude, more consistent with the 2010 data. None the less, if the variability did remain constant throughout this time, we can refine the period to P = 2.060561 \u00b1 0.000002 d.","Citation Text":["Munday et al. 2020"],"Functions Text":["Given the nature of the variability, we measure the period via the reduced \u03c72 of a sinusoidal fit as a function of frequency (ideally suited to sinusoidal variability, e.g."],"Functions Label":["Uses"],"Citation Start End":[[750,768]],"Functions Start End":[[550,722]]}
{"Identifier":"2022MNRAS.516.3175B__Law_et_al._2015_Instance_1","Paragraph":"We use the final version of the MaNGA data (Abdurro\u2019uf et al. 2022). MaNGA is the survey project in SDSS-IV (Sloan Digital Sky Survey) using IFU observations of galaxies to produce spatially resolved spectroscopic data (Gunn et al. 2006; Bundy et al. 2015; Drory et al. 2015; Blanton et al. 2017; Aguado et al. 2019; Abdurro\u2019uf et al. 2022). MaNGA observed approximately 10\u2009000 galaxies, selected as an unbiased sample in terms of stellar mass (${\\it M}_{\\star } \\gt 10^{9}\\, {\\rm M}_{\\odot }$) and environments (Law et al. 2015; Yan et al. 2016a; Wake et al. 2017). The MaNGA targets were selected as two main subgroups and two minor subgroups. Among the main subgroups, the primary sample of about 5000 galaxies was selected to observe the central part out to $\\rm 1.5{\\it R}_{eff}$. The secondary sample of the main subgroups includes about 3300 galaxies, observed by the IFU bundle out to $\\rm 2.5{\\it R}_{eff}$. These samples were selected without bias and covers a wide range in galaxy masses. The first of the two minor subgroups is a \u2018colour-enhanced sample\u2019 of about 1700 galaxies. It is selected to include blue massive galaxies and green valley galaxies that are important to study the quenching process. These targets are observed with an IFU coverage of $\\rm 1.5{\\it R}_{eff}$. A second minor sample includes about 1000 ancillary targets, selected for various observations using the unique capability of the MaNGA instrument (Yan et al. 2016a). The IFU bundles of the MaNGA consist of 19\u2013127 fibres with hexagonal shape covering 12\u2013$\\rm 32~arcsec$ in diameter on sky. The spatial resolution of MaNGA is $\\rm 2.5~arcsec$, which corresponds to 1.3\u2013$\\rm 4.5~kpc$ for the primary sample and 2.2\u2013$\\rm 5.1~kpc$ for the secondary sample. The BOSS spectrographs used by MaNGA provide spectra from 3600 to 10\u2009300 \u00c5  with a spectral resolution of about 2100\u20136000 \u00c5  (Smee et al. 2013; Yan et al. 2016b). The observations were conducted to reach an signal-to-noise ratio (S\/N) in the stellar continuum at $\\rm 1{\\it R}_{eff}$ of $\\rm 14$\u2013$\\rm 35$ per spatial sample, which required about 3 h net integration for each target.","Citation Text":["Law et al. 2015"],"Functions Text":["MaNGA observed approximately 10\u2009000 galaxies, selected as an unbiased sample in terms of stellar mass (${\\it M}_{\\star } \\gt 10^{9}\\, {\\rm M}_{\\odot }$) and environments"],"Functions Label":["Background"],"Citation Start End":[[513,528]],"Functions Start End":[[342,511]]}
{"Identifier":"2020MNRAS.492.5247S__Sasaki_et_al._2018_Instance_1","Paragraph":"A lognormal mass function:\n(7)$$\\begin{eqnarray}\r\n\\psi (M;\\mu ,\\sigma)=\\frac{1}{M\\sqrt{2\\pi \\sigma ^2}}\\exp {\\left(- \\frac{\\ln ^2(M\/\\mu)}{2\\sigma ^2}\\right)},\r\n\\end{eqnarray}$$where \u03bc > 0 and \u03c3 > 0. The mean PBH mass in the assumption of a lognormal mass function is $\\bar{M}=\\mu \\exp \\left(-\\sigma ^2\/2\\right)$. Such mass function is a good approximation for a large class of PBH formation scenarios, e.g. axion-curvaton, running-mass, and single field double inflation (Green 2016; Kannike et al. 2017). The (\u03bc, \u03c3) pairs allowed in the window relevant for the operational gravitational wave detectors [$\\mathcal {O}(1\\operatorname{-}100)\\, {\\rm M}_\\odot {}$; Abbott et al. 2019] are mainly constrained by (i) the results from searches for microlensing events on stars in our Galactic neighbourhood and (ii) the effect PBH gas accretion would have on the CMB temperature and ionization history (Carr et al. 2017; Sasaki et al. 2018). Together, the allowed points roughly constitute a sub-plane described as (see fig. 3 in Carr et al. 2017):\n(8)$$\\begin{eqnarray}\r\n(\\mu ,\\sigma)\\in [25\\, {\\rm M}_\\odot ,100\\, {\\rm M}_\\odot ]\\times [0.0,1.0].\r\n\\end{eqnarray}$$It is worth mentioning that these points are ruled out completely if one further takes the following constraints into account (Carr et al. 2017): (iii) the survival of the stellar cluster in Eridanus II and of the entire stellar populations in UFDGs (Brandt 2016; Green 2016; Koushiappas & Loeb 2017; Zhu et al. 2018) and (iv) the survival of wide binaries in the Milky Way (Monroy-Rodr\u00edguez & Allen 2014). These constraints are somewhat less restrictive, as they rely on further astrophysical assumptions (Carr et al. 2017). For instance, the stellar cluster in Eridanus II could have only recently spiralled down into the centre of the galaxy, where dynamical heating by PBHs becomes effective (Brandt 2016). Additionally, wide binaries are in principle hard to detect, yielding some uncertainty in identifying them and consequently in drawing any conclusion on the PBH abundance (Sasaki et al. 2018). Moreover, we emphasize that, most recently, even previous microlensing constraints were called into question as spatial PBH clustering (Garc\u00eda-Bellido & Clesse 2018) and updated galactic rotation curves (Hawkins 2015) tend to relax them.","Citation Text":["Sasaki et al. 2018"],"Functions Text":["The (\u03bc, \u03c3) pairs allowed in the window relevant for the operational gravitational wave detectors","are mainly constrained by (i) the results from searches for microlensing events on stars in our Galactic neighbourhood and (ii) the effect PBH gas accretion would have on the CMB temperature and ionization history"],"Functions Label":["Uses","Uses"],"Citation Start End":[[914,932]],"Functions Start End":[[506,602],[681,894]]}
{"Identifier":"2021MNRAS.507.5882S__Mackereth_et_al._2018_Instance_3","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)"],"Functions Text":["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."],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[2232,2255]],"Functions Start End":[[2256,2406]]}
{"Identifier":"2021AandA...648A.109G__Talon_&_Charbonnel_2005_Instance_1","Paragraph":"The transport of angular momentum (AM) and chemical elements in stars strongly affects their evolution, from pre-main sequence (PMS) to evolved stages. These processes are particularly crucial in the stellar radiative zones, but their modelling remains an open question. In standard evolution models, these stably stratified zones are assumed to be motionless despite early works pointing out the lack of a static solution in uniformly rotating stars (Von Zeipel 1924) and a related large-scale meridional circulation driven by the centrifugal acceleration (Eddington 1925; Sweet 1950). A more complete formalism was then introduced by Zahn (1992) including a self-consistent meridional flow and models of the turbulent transport driven by hydrodynamical instabilities. The importance of additional processes such as internal waves (e.g., Talon & Charbonnel 2005) and magnetic fields (Spruit 2002) has also been investigated. Zahn\u2019s formalism successfully explained a number of observed stellar properties, such as the nitrogen abundances at the surface of red supergiants or the observed blue-to-red supergiant ratio in the Small Magellanic Cloud (Maeder & Meynet 2001). The dynamics of internal gravity waves could possibly explain the flat rotation profile of the solar radiative zone (Talon et al. 2002 and Charbonnel & Talon 2005) inferred through helioseismology (Schou et al. 1998), as well as the lithium dip in solar-like stars (Charbonnel & Talon 2005). Despite these early encouraging results many stellar observations remain unexplained, especially the internal rotation rates revealed by asteroseismology in evolved stars and in intermediate-mass main sequence stars (see Aerts et al. 2019 for a review). The current theoretical understanding of the structure of the differential rotation in stellar radiative zones, crucial for the development of instabilities and thus for the related turbulent transport, is still largely incomplete. In particular, the assumption that the differential rotation is mostly radial, rather than radial and latitudinal, is at the base of Zahn\u2019s formalism, but the validity of this assumption has never been thoroughly tested. In this paper we are particularly interested in the differential rotation and the large-scale meridional flows generated in periods of the stellar life when contraction, expansion, or both processes occur. This is the case for example for PMS stars which are gravitationally contracting before starting their core nuclear reactions, or for subgiant and giant stars which undergo contraction of their core and expansion of their envelope. Within these stars a strong redistribution of AM is expected to happen, producing differential rotation and thus potentially unstable shear flows.","Citation Text":["Talon & Charbonnel 2005"],"Functions Text":["The importance of additional processes such as internal waves (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[839,862]],"Functions Start End":[[770,838]]}
{"Identifier":"2020MNRAS.492.5247S__Sasaki_et_al._2018_Instance_2","Paragraph":"A lognormal mass function:\n(7)$$\\begin{eqnarray}\r\n\\psi (M;\\mu ,\\sigma)=\\frac{1}{M\\sqrt{2\\pi \\sigma ^2}}\\exp {\\left(- \\frac{\\ln ^2(M\/\\mu)}{2\\sigma ^2}\\right)},\r\n\\end{eqnarray}$$where \u03bc > 0 and \u03c3 > 0. The mean PBH mass in the assumption of a lognormal mass function is $\\bar{M}=\\mu \\exp \\left(-\\sigma ^2\/2\\right)$. Such mass function is a good approximation for a large class of PBH formation scenarios, e.g. axion-curvaton, running-mass, and single field double inflation (Green 2016; Kannike et al. 2017). The (\u03bc, \u03c3) pairs allowed in the window relevant for the operational gravitational wave detectors [$\\mathcal {O}(1\\operatorname{-}100)\\, {\\rm M}_\\odot {}$; Abbott et al. 2019] are mainly constrained by (i) the results from searches for microlensing events on stars in our Galactic neighbourhood and (ii) the effect PBH gas accretion would have on the CMB temperature and ionization history (Carr et al. 2017; Sasaki et al. 2018). Together, the allowed points roughly constitute a sub-plane described as (see fig. 3 in Carr et al. 2017):\n(8)$$\\begin{eqnarray}\r\n(\\mu ,\\sigma)\\in [25\\, {\\rm M}_\\odot ,100\\, {\\rm M}_\\odot ]\\times [0.0,1.0].\r\n\\end{eqnarray}$$It is worth mentioning that these points are ruled out completely if one further takes the following constraints into account (Carr et al. 2017): (iii) the survival of the stellar cluster in Eridanus II and of the entire stellar populations in UFDGs (Brandt 2016; Green 2016; Koushiappas & Loeb 2017; Zhu et al. 2018) and (iv) the survival of wide binaries in the Milky Way (Monroy-Rodr\u00edguez & Allen 2014). These constraints are somewhat less restrictive, as they rely on further astrophysical assumptions (Carr et al. 2017). For instance, the stellar cluster in Eridanus II could have only recently spiralled down into the centre of the galaxy, where dynamical heating by PBHs becomes effective (Brandt 2016). Additionally, wide binaries are in principle hard to detect, yielding some uncertainty in identifying them and consequently in drawing any conclusion on the PBH abundance (Sasaki et al. 2018). Moreover, we emphasize that, most recently, even previous microlensing constraints were called into question as spatial PBH clustering (Garc\u00eda-Bellido & Clesse 2018) and updated galactic rotation curves (Hawkins 2015) tend to relax them.","Citation Text":["Sasaki et al. 2018"],"Functions Text":["Additionally, wide binaries are in principle hard to detect, yielding some uncertainty in identifying them and consequently in drawing any conclusion on the PBH abundance"],"Functions Label":["Uses"],"Citation Start End":[[2042,2060]],"Functions Start End":[[1870,2040]]}
{"Identifier":"2019AandA...632A.129W__Feng_&_Wang_2015_Instance_2","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"],"Functions Text":["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","and therefore they are independently reliable ICME indicators"],"Functions Label":["Uses","Uses"],"Citation Start End":[[1189,1205]],"Functions Start End":[[955,1085],[1126,1187]]}
{"Identifier":"2021ApJ...906...57S__Lee_et_al._2016_Instance_1","Paragraph":"Complementary to studies using the integrated emission and angular power spectrum of DM annihilation from a population of Galactic subhalos, in this paper we present a novel strategy using one-point photon statistics to search for the annihilation signature. Our technique takes advantage of the information in the entire population of sources, including both those that are resolved and those that are faint and unresolved. The concept of leveraging the one-point photon-count distribution to search for DM has previously been studied in Dodelson et al. (2009) and Feyereisen et al. (2015) in the context of emission from extragalactic sources and in Lee et al. (2009) and Koushiappas et al. (2010) with application to Galactic subhalos. We introduce a method to search for signatures of DM annihilation from a Galactic subhalo population using the non-Poissonian template fitting (NPTF) framework (Malyshev & Hogg 2011; Lee et al. 2015, 2016; Mishra-Sharma et al. 2017), which has previously been applied to characterize unresolved point sources in the inner Galaxy (Lee et al. 2016; Linden et al. 2016; Leane & Slatyer 2019, 2020a, 2020b; Chang et al. 2020; Buschmann et al. 2020) and at high latitudes (Zechlin et al. 2016; Lisanti et al. 2016; Zechlin et al. 2018). Using simulations, we show that the NPTF can constrain DM annihilation from a population of subhalos in the face of astrophysical background emission. We find that using photon statistics to look for collective emission from a subhalo population can be especially promising when a large number of individual subhalo candidates are identified in point-source catalogs. This establishes a method complementary to the established ones based on characterizing individual resolved point sources as subhalo candidates, as well as those based on using the measured 0-point (overall flux) and two-point (angular power spectrum) statistics to characterize a subhalo population. Moreover, note that our methodology is completely independent of assumptions about, e.g., the location of stellar overdensities, so we are less sensitive to certain uncertainties which can bias dwarf galaxy constraints such as those in modeling tracer populations. Thus, this framework provides an important comparison for the dwarf galaxy analyses as well.","Citation Text":["Lee et al.","2016","Lee et al. 2016"],"Functions Text":["We introduce a method to search for signatures of DM annihilation from a Galactic subhalo population using the non-Poissonian template fitting (NPTF) framework","which has previously been applied to characterize unresolved point sources in the inner Galaxy"],"Functions Label":["Uses","Motivation"],"Citation Start End":[[922,932],[939,943],[1069,1084]],"Functions Start End":[[739,898],[973,1067]]}
{"Identifier":"2018ApJ...868..139W__Schlickeiser_&_Jenko_2010_Instance_2","Paragraph":"By radio continuum surveys of interstellar space and direct in situ measurements in the solar system, it is well established that for many scenarios the background magnetic fields are spatially varying. However, the above research about parallel and perpendicular diffusion only explored the uniform mean magnetic field. One can show that the spatially varying background magnetic fields lead to the adiabatic focusing effect of charged energetic particle transport and introduces correction to the particle diffusion coefficients (see, e.g., Roelof 1969; Earl 1976; Kunstmann 1979; Beeck & Wibberenz 1986; Bieber & Burger 1990; K\u00f3ta 2000; Schlickeiser & Shalchi 2008; Shalchi 2009b, 2011; Litvinenko 2012a, 2012b; Shalchi & Danos 2013; Wang & Qin 2016; Wang et al. 2017b). To explore the influence of adiabatic focusing on particle transport, the perturbation method is frequently used (see, e.g., Beeck & Wibberenz 1986; Bieber & Burger 1990; Schlickeiser & Shalchi 2008; Schlickeiser & Jenko 2010; Litvinenko & Schlickeiser 2013; He & Schlickeiser 2014). To use the perturbation method, since the anisotropic distribution function is an implicit function, by using the iteration method, one can find that the anisotropic distribution function becomes an infinite series of the spatial and temporal derivatives of the isotropic distribution function. Therefore, the governing equation of the isotropic distribution function derived from the Fokker\u2013Planck equation contains infinite terms because of the infinite series of the anisotropic distribution function. By using the truncating method to neglect the higher-order derivative terms, the approximate correction formulas of parallel or perpendicular diffusion coefficients were obtained (see, e.g., Schlickeiser & Shalchi 2008; Schlickeiser & Jenko 2010; He & Schlickeiser 2014). However, the higher-order derivative terms probably also make the correction to the parallel and perpendicular diffusion much like the lower-order derivative ones do. The magnitude of the correction from higher-order derivative term might not necessarily be a higher-order small quantity than the magnitude of the lower-order derivative terms. Therefore, the correction obtained by the previous authors is likely to contain significant errors. In this paper, by considering the higher-order derivative terms, we derive the parallel and perpendicular diffusion coefficients and obtain the correction formulas coming from all order derivative terms by using the improved perturbation method (He & Schlickeiser 2014) and the iteration operation. And for the weak adiabatic focusing limit we evaluate the correction to the parallel diffusion coefficient and compare it with the correction obtained in the previous papers.","Citation Text":["Schlickeiser & Jenko 2010"],"Functions Text":["By using the truncating method to neglect the higher-order derivative terms, the approximate correction formulas of parallel or perpendicular diffusion coefficients were obtained (see, e.g.,"],"Functions Label":["Background"],"Citation Start End":[[1783,1808]],"Functions Start End":[[1563,1753]]}
{"Identifier":"2019AandA...630A..37S__Behar_et_al._2017_Instance_2","Paragraph":"Solar wind velocity distribution moments are described in Behar et al. (2017). The ion density nsw is the moment of order 0, and the ion bulk velocity usw (a vector) appears in the moment of order 1, the flux density \n\n$n_{\\mathrm{sw}} \\ \\underline{\\mathbf{u}_{\\mathrm{sw}}}$\n\n\n\nn\n\nsw\n\n\u2009\n\n\nu\n\nsw\n\n\n_\n\n\n\n. The bulk speed can be defined as the norm of the bulk velocity, that is, \n\n$u_{\\mathrm{sw}} = |\\underline{\\vec{u}_{\\mathrm{sw}}}$\n\n\n\nu\n\nsw\n\n=|\n\n\nu\n\nsw\n\n\n_\n\n\n\n|. However, this bulk speed is representative of single-particle speeds as long as the velocity distribution function is compact (e.g., a Maxwellian distribution). Complex velocity distribution functions were observed by RPC-ICA within the atmosphere of 67P. For instance, partial ring distributions were frequently observed for solar wind protons at intermediate heliocentric distances, when the spacecraft approached the SWIC (Behar et al. 2017). To illustrate the effect of such distorted distributions, a perfect ring (or shell) distribution centered on the origin of the plasma reference frame can be imagined, in which all particles have the same speed of 400 km s\u22121. The norm of the bulk velocity in this case would be 0 km s\u22121, whereas the mean speed of the particles is 400 km s\u22121, which is the relevant speed for SWCX processes. This mean speed, noted Usw, of the particles is calculated by first summing the differential number flux over all angles, and then taking the statistical average (Behar 2018). Over the entire mission, the deceleration of the solar wind using the mean speed of the particles is much more limited than the deceleration shown by the norm of the bulk velocity (Behar et al. 2017): there is more kinetic energy in the solar wind than the bulk velocity vector would let us think. This is the main difference with the paradigm used at previously studied (and more active) comets (Behar et al. 2018b). These complex, nonthermal velocity distribution functions also prevent us from reducing the second-order moment (the stress tensor) to a single scalar value, which, for a Maxwellian distribution, could be identified with a plasma temperature. In the context of 67P and for an important part of the cometary orbit around the Sun, the temperature of the solar wind proton has no formal definition.","Citation Text":["Behar et al. 2017"],"Functions Text":["For instance, partial ring distributions were frequently observed for solar wind protons at intermediate heliocentric distances, when the spacecraft approached the SWIC"],"Functions Label":["Background"],"Citation Start End":[[892,909]],"Functions Start End":[[722,890]]}
{"Identifier":"2020ApJ...895...82V__Fryer_et_al._2018_Instance_2","Paragraph":"The shock is then revived by adding an energy injection following the parameterized method of Fryer et al. (2018). In this model, roughly \n\n\n\n\n\n was deposited into the inner \n\n\n\n\n\n in the first \n\n\n\n\n\n. Some of this energy is lost through neutrino emission and the total explosion energy at late times for this model is \n\n\n\n\n\n. This explosion is then mapped into our three-dimensional calculations, using one million SPH particles. The mapping took place when the supernova shock had moved out of the iron core and propagated into the Si\u2013S rich shell at \n\n\n\n\n\n. We note that our 1D methods employed for modeling the collapse, core bounce, and initial explosion do not capture the full physics of the central engine (for a discussion, see Fryer et al. 2018), and this is a source of uncertainty in our yield calculations. The details of the engine change the shock trajectories, and neutrino chemistry can change Ye values (Saez et al. 2018; Fujimoto & Nagakura 2019). The nature of the shock affects mostly the yields after the shock falls below NSE (before it falls out of NSE, the yields are set by the equilibrium values, not the time-dependent evolution). Our model captures one instance of the range of asymmetric trajectories, and it should be noted that no model at this time is sufficiently accurate to dictate exactly the properties of the asymmetries (Janka et al. 2016). In addition, any model that does not include convection-driven asymmetries from the progenitor star cannot properly capture the asymmetries (Arnett et al. 2015). The 3D explosion model used here also displays stochastic asymmetries, implying that any manner of convective asymmetry could generate similar results. If this behavior is universal, it could have important implications. These points taken together indicate that nucleosynthetic patterns arising from convection-like behavior are robust, regardless of the driver. As discussed below, this increases the utility of NSE nucleosynthesis, particularly of \n\n\n\n\n\n and \n\n\n\n\n\n, as diagnostics of the conditions in the progenitor star.","Citation Text":["Fryer et al. 2018"],"Functions Text":["We note that our 1D methods employed for modeling the collapse, core bounce, and initial explosion do not capture the full physics of the central engine (for a discussion, see","and this is a source of uncertainty in our yield calculations."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[737,754]],"Functions Start End":[[561,736],[757,819]]}
{"Identifier":"2019AandA...627A.130D__Broadhurst_et_al._2019_Instance_1","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"],"Functions Text":["Some of these events may correspond to gravitationally lensed events with magnification factors ranging from a few tens to a few hundreds"],"Functions Label":["Background"],"Citation Start End":[[605,627]],"Functions Start End":[[394,531]]}
{"Identifier":"2019MNRAS.484.1946G__Biesiadzinski_et_al._2012_Instance_1","Paragraph":"Secondly, we focus on galaxies to discuss some remaining solutions. The miscentring (e.g. A1986, A1961) decreases the X-ray luminosity\/SZ signal, as some flux moves outside of the aperture. Sehgal et al. (2013) demonstrated the effect of miscentring on decreasing SZ signal. They also proposed that the miscentring effect causes their lower measured SZ signal compared to Planck due to the finer resolution of ACT. However, as they point out, the miscentring distribution from their sample alone can only explain part of the discrepancy between optical and SZ, unless an unrealistic larger offset exists. Moreover, most miscentred clusters are merging or disturbed clusters with lower X-ray or SZ surface brightness, which makes them more difficult to be detected. The discrepancy of SZ signal is at a level of 10 per\u2009cent in our richness range (Biesiadzinski et al. 2012). The projection (e.g. A750, A1319) will increase the optical richness. As the cluster mass-richness relation is close to a linear relation, projection causes the projected clusters to simply slide up and down the mass-richness relation, without deviating from it (Simet et al. 2017). However, the LX\u2212M or YSZ\u2212M relation is a power-law relation with an index greater than one (\u223c1.6; Rozo et al. 2014a). For example, if we have two clusters with N200 = 80 ($M_{500|N=80}=5\\times 10^{14} \\, \\mathrm{M}_{\\odot}$, bolometric X-ray luminosity LX|N = 80 = 11 \u00d7 1044\u2009erg\u2009s\u22121) projected together, they will be detected as a N = 160 cluster, the corresponding expected mass and X-ray luminosity are $M_{500|N=160}=10\\times 10^{14} \\, \\mathrm{M}_{\\odot}$ and LX|N = 160 = 43 \u00d7 1044\u2009erg\u2009s\u22121, respectively. Though the mass is equal to the sum of two subclusters, the total X-ray luminosity is overestimated by a factor of \u223c2 than the linear combination of 2LX|N = 80 = 22 \u00d7 1044\u2009erg\u2009s\u22121. Thus, projection causes the expected X-ray luminosity or the SZ signal from the summed optical richness to be higher than the actual summed values. The projection fraction of samples extending to much lower richness is around 10 per\u2009cent (Simet et al. 2017), even higher for these richest maxBCG clusters (Fig. 2). The above two effects can act together, especially in super clusters and large-scale filaments (Fig. D1). The contamination of low-mass haloes, whose true halo mass is far below the value suggested by the optical richness, would also dilute and reduce the mean mass of the sample. We note that there is contamination of such low-mass systems based on Fig. A2. These low-mass haloes are mostly blended systems with boosted richness affected by nearby large-scale structure. Thus there is a mixture of halo masses at very high N200, the clean and the blended, and the PDF of M given N200 (or \u03bb) will be asymmetric with a low-mass tail. A skew-normal or Hermite polynomial expansion (Shaw, Holder & Dudley 2010) are good alternatives to mixture modelling. Next, we roughly estimate the contamination fraction. There are four low LX systems beyond the 2\u03c3 line of LX\u2212N relation from Rozo et al. (2014b) in Fig. A2 (and seven beyond the 1.5\u03c3 line towards lower LX versus 1 beyond the 1.5\u03c3 line towards higher LX). Taking these numbers at face value, the contamination is 10\u201315 per\u2009cent.","Citation Text":["Biesiadzinski et al. 2012"],"Functions Text":["Moreover, most miscentred clusters are merging or disturbed clusters with lower X-ray or SZ surface brightness, which makes them more difficult to be detected. The discrepancy of SZ signal is at a level of 10 per\u2009cent in our richness range"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[846,871]],"Functions Start End":[[605,844]]}
{"Identifier":"2021MNRAS.503.5385Z__Barnes_et_al._2001_Instance_1","Paragraph":"\u2018Blind\u2019 searches through H\u2009i surveys performed by large aperture single-dish telescopes provide chances to uncover more H\u2009i absorption systems. Although such surveys usually utilize the non-tracking, drift scan observing strategy, with integration time for each individual source is limited, the collecting areas of the participating telescopes can still achieve considerable sensitivities. Allison, Sadler & Meekin (2014) identified four H\u2009i absorbers, including one previously unknown source, with the archival data of the H\u2009i Parkes All-Sky Survey (H\u2009i PASS; see Barnes et al. 2001), while Darling et al. (2011), Wu et al. (2015), as well as Song et al. (in preparation) have made attempts to perform \u2018blind\u2019 searches for absorption features using part of Arecibo Legacy Fast Arecibo L-Band Feed Array (ALFA) Survey (ALFALFA; see Giovanelli et al. 2005; Haynes et al. 2018) data, respectively, with 10 sources identified in total, including 3 samples remained undetected by other instruments so far, i.e., UGC 00613, CGCG 049-033, and PGC 070403, thus proving the feasibility of searching for new H\u2009i absorbers with massive blind sky surveys. Compared with radio interferometers, large single-dish telescopes such as Arecibo can provide better sensitivities and usually higher spectral resolution, which are all crucial to reveal the characteristics of extragalactic H\u2009i lines. And although the spatial resolution for such instruments are quite limited compared with interferometers, the chance of having two or more H\u2009i absorbing systems lying within the beamwidth with similar redshift is quite low, thus making confusions in source identification unlikely to happen. However, it should be noted that due to temporal variations in spectral baseline commonly seen during drift scans, follow-up observations are needed for reliable characterization of the newly identified absorbers, especially for the weak ones. Also, interferometric mappings are usually required to discern possible fine structures within each absorbing system.","Citation Text":["Barnes et al. 2001"],"Functions Text":["Allison, Sadler & Meekin (2014) identified four H\u2009i absorbers, including one previously unknown source, with the archival data of the H\u2009i Parkes All-Sky Survey (H\u2009i PASS; see"],"Functions Label":["Background"],"Citation Start End":[[566,584]],"Functions Start End":[[391,565]]}
{"Identifier":"2022AandA...666A.134S__Snyder_et_al._2005_Instance_1","Paragraph":"Observation of amino acids and their most essential isomers and potential precursors in the interstellar medium (ISM) should be crucial for revealing the chemistry that may have led to life's origin (Ehrenfreund et al. 2001). In particular, the central question of whether glycine (CH2(NH2)C(O)OH) exists or not in the ISM is one of the most pursued targets in astrochemistry. Although several attempts to observe glycine have been reported (Hollis et al. 2003; Cunningham et al. 2007; Jones et al. 2007; Jim\u00e9nez-Serra et al. 2016, 2020), its detection has never been confirmed (Snyder et al. 2005). Nevertheless, it has been found in the coma of comets 67P\/Churyumov\u2013Gerasimenko through in situ mass spectrometry (Altwegg et al. 2016). The presence of glycine in the volatile cometary material thus strongly suggests the existence of a process capable of generating amino acids in cold environments (Bizzocchi et al. 2020). Hence, as a prerequisite step for its astronomical identification, the rotational spectra of glycinamide (CH2 (NH2)C(O)NH2; Alonso et al. 2018; Kisiel et al. 2022) and aminoacetonitrile (CH2(NH2)CN; Kolesnikov\u00e1 et al. 2017), which are relevant intermediates in the Strecker synthesis of glycine; hydantoin (CH2C(O)NHC(O)NH; Alonso et al. 2017) and hydantoic acid (C(O)OHCH2NHC(O)NH2; Kolesnikov\u00e1 et al. 2019), potential glycine precursors through a different hydrolytic pathway (Ozeki et al. 2017), and the glycine isomer glycolamide (CH2(OH)C(O)NH2); Sanz-Novo et al. 2020) have been recently reported. In addition, Sanz-Novo et al. (2019) very recently carried out a computational study of the potential energy surfaces (PES) corresponding to the formation reactions of several protonated glycine isomers. Surprisingly, the only exothermic process with no net activation barrier led to protonated acetohydroxamic acid [CH3C(O)NH2OH]+. Its formation could therefore be feasible under interstellar conditions. Consequently, the corresponding neutral counterpart, CH3C(O)NHOH, might be a candidate molecule to be searched for in the ISM. Our previous simulations (Sanz-Novo et al. 2019) suggest that this glycine isomer should be searched for in the ISM, eventually, but there is no experimental rotational data available for this molecular system.","Citation Text":["Snyder et al. 2005"],"Functions Text":["In particular, the central question of whether glycine (CH2(NH2)C(O)OH) exists or not in the ISM is one of the most pursued targets in astrochemistry.","its detection has never been confirmed"],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[579,597]],"Functions Start End":[[226,376],[539,577]]}
{"Identifier":"2015AandA...577A.118M__Swift_et_al._2011_Instance_1","Paragraph":"On the other hand, meteoroid rates on the Moon can be related to rates on Earth by taking into account the different gravitational focusing effect between both bodies. The gravitational focusing factor \u03a6 is given by (6)\\begin{equation} \\Phi=1+\\frac{V^{2}_{\\rm esc}}{V^{2}}, \\end{equation}\u03a6=1+Vesc2V2,where Vesc is the escape velocity of the central body and V the meteoroid velocity. The different gravitational focusing effect for the Moon and the Earth is given by the quotient \u03b3 between the gravitational focusing factors for both bodies. For sporadic meteoroids, with an average velocity of 20 km\u2009s-1 (Brown et al. 2002), this velocity-dependent focusing effect is higher for the Earth by a factor of 1.3 (Ortiz et al. 2006), and so \u03b3SPO = 0.77. According to this we have \\begin{eqnarray} &amp;&amp;{\\rm HR}^{\\rm SPO}_{\\rm Moon} =\\gamma^{\\rm SPO}{\\rm HR}^{\\rm SPO}_{\\rm Earth}, \\\\[3mm] &amp;&amp;{\\rm ZHR}^{\\rm ST}_{\\rm Moon} =\\sigma\\gamma^{\\rm ST}{\\rm ZHR}^{\\rm ST}_{\\rm Earth}. \\end{eqnarray}HRMoonSPO=\u03b3SPOHREarthSPO,ZHRMoonST=\u03c3\u03b3STZHREarthST.For the average hourly rate of sporadic events we have HR\\hbox{$^{\\rm SPO} _{\\rm Earth} =10$}SPOEarth=10 meteors\u2009h-1 (Dubietis & Artl 2010). In Eq. (8) an additional factor \u03c3 is included to take into account that the distances from the Earth and the Moon to the meteoric filament will in general be different, and this would give rise to a different density of stream meteoroids for both bodies. If we assume a simple situation where this filament can be approximated as a tube where the meteoroid density decreases linearly from its central axis, the following definition can be adopted for \u03c3, (9)\\begin{equation} \\sigma=\\frac{d_{\\rm Earth}}{d_{\\rm Moon}}, \\end{equation}\u03c3=dEarthdMoon,where dEarth and dMoon are the distance from the center of the meteoric tube to the Earth and the Moon, respectively. On the other hand, the ZHR on Earth at solar longitude \u03bb (which corresponds to the time of detection of the impact flash) can be related to the peak ZHR by means of (Jenniskens 1994) (10)\\begin{equation} {\\rm ZHR}^{\\rm ST}_{\\rm Earth}={\\rm ZHR}^{\\rm ST}_{\\rm Earth}({\\rm max})10^{-b|\\lambda-\\lambda_{\\rm max}|}, \\end{equation}ZHREarthST=ZHREarthST(max)10\u2212b|\u03bb\u2212\u03bbmax|,where ZHR\\hbox{$^{\\rm ST} _{\\rm Earth}$}STEarth(max) is the peak ZHR on Earth (corresponding to the date given by the solar longitude \u03bbmax). The values for the peak ZHR for different meteoroid streams and the corresponding solar longitudes for these maxima can be obtained, for instance, from (Jenniskens 2006). For streams with non-symmetrical ascending and descending activity profiles or with several maxima, Eq. (10) should be modified according to the expressions given in Jenniskens (1994). By putting all these pieces together in Eq. (2), we can write the following expression for the probability parameter: (11)\\begin{equation} p^{\\rm ST}=\\frac{\\gamma^{\\rm ST}\\cos(\\phi)\\sigma {\\rm ZHR}^{\\rm ST}_{\\rm Earth}({\\rm max})10^{-b|\\lambda-\\lambda_{\\rm max}|}}{\\gamma^{\\rm SPO}{\\rm HR}^{\\rm SPO}_{\\rm Earth}+\\gamma^{\\rm ST}\\cos(\\phi)\\sigma {\\rm ZHR}^{\\rm ST}_{\\rm Earth}({\\rm max})10^{-b|\\lambda-\\lambda_{{\\rm max}}|}}\\cdot \\end{equation}pST=\u03b3STcos(\u03c6)\u03c3ZHREarthST(max)10\u2212b|\u03bb\u2212\u03bbmax|\u03b3SPOHREarthSPO+\u03b3STcos(\u03c6)\u03c3ZHREarthST(max)10\u2212b|\u03bb\u2212\u03bbmax|\u00b7However, this formula does not take into account the fundamental fact that only those meteoroids capable of producing impact flashes detectable from Earth should be included in the computations. In fact, by employing only the hourly rates measured on Earth in Eq. (11), which measures the flux of meteor brighter than mag. +6.5, it is implicitly assumed that meteoroids producing meteor events on Earth can also produce detectable impact flashes on the Moon. However, this assumption is incorrect. Thus, for a given meteoroid stream (i.e., for a given meteoroid geocentric velocity), the mass mo of meteoroids giving rise to mag. +6.5 meteors on Earth can be obtained from Eqs. (1) and (2) in Hughes (1987). For instance, this mass yields 5.0 \u00d7 10-8 kg for Perseid meteoroids (Vg = 59 km\u2009s-1), 2.4 \u00d7 10-8 kg for Leonids (Vg = 70 km\u2009s-1), and 5.0 \u00d7 10-6 kg for sporadic meteoroids with an average velocity of 20 km\u2009s-1 (Brown et al. 2002). However, the masses corresponding to impact flashes recorded on the Moon are several orders of magnitudes larger than mo (see, e.g., Ortiz et al. 2006; Yanagisawa et al. 2006; Swift et al. 2011). This means that the minimum kinetic energy Emin to produce a detectable impact flash on the Moon is much higher than the kinetic energy of a portion of the meteoroids included in the computation of hourly rates on Earth. This means that the velocity and mass distribution of meteoroids must be somehow included in Eq. (11) to take into consideration for the computation of the probability parameter p only those meteoroids with a kinetic energy above the threshold kinetic energy given by Emin (the method for the computation of Emin will be explained below). According to Eq. (2) in Bellot Rubio et al. (2000a,b), this can be accomplished by including in Eqs. (2) and (11) the factor (12)\\begin{equation} \\nu=\\left(\\frac{m_{\\rm o}V^{2}}{2}\\right)^{s-1}E^{1-s}_{\\rm min}, \\end{equation}\u03bd=moV22s\u22121Emin1\u2212s,where V is the impact velocity, mo is the mass of a shower meteoroid producing on Earth a meteor of magnitude +6.5, and s is the mass index, which is related to the population index r (the ratio of the number of meteors with magnitude m + 1 or less to the number of meteors with magnitude m or less) by means of the relationship (13)\\begin{equation} s=1+2.5\\log(r). \\end{equation}s=1+2.5log(r).According to the definition of \u03bd, this parameter is different for each meteoroid stream (and for sporadic meteoroids, of course). By taking this into account Eq. (11) should be modified as follows: (14)\\begin{equation} p^{\\rm ST}=\\frac{\\nu^{\\rm ST}\\gamma^{\\rm ST}\\cos(\\phi)\\sigma {\\rm ZHR}^{\\rm ST}_{\\rm Earth}({\\rm max})10^{-b|\\lambda-\\lambda_{\\rm max}|}}{{\\nu^{\\rm SPO}\\gamma^{\\rm SPO}{\\rm HR}^{\\rm SPO}_{\\rm Earth}}{+\\nu^{\\rm ST}\\gamma^{\\rm ST}\\cos(\\phi)\\sigma {\\rm ZHR}^{\\rm ST}_{\\rm Earth}({\\rm max})10^{-b|\\lambda-\\lambda_{\\rm max}|}}}\\cdot \\end{equation}pST=\u03bdST\u03b3STcos(\u03c6)\u03c3ZHREarthST(max)10\u2212b|\u03bb\u2212\u03bbmax|\u03bdSPO\u03b3SPOHREarthSPO+\u03bdST\u03b3STcos(\u03c6)\u03c3ZHREarthST(max)10\u2212b|\u03bb\u2212\u03bbmax|\u00b7If by the time of detection of the impact flash n additional meteoroid streams with significant contributions to the impact rate (and with compatible impact geometry) must be considered, the denominator in Eq. (14) must be modified in the following way, (15)\\begin{equation} p^{\\rm ST}=\\frac{\\nu^{\\rm ST}\\gamma^{\\rm ST}\\cos(\\phi)\\sigma {\\rm ZHR}^{\\rm ST}_{\\rm Earth}({\\rm max})10^{-b|\\lambda-\\lambda_{\\rm max}|}}{{\\nu^{\\rm SPO}\\gamma^{\\rm SPO}{\\rm HR}^{\\rm SPO}_{\\rm Earth}}{+\\nu^{\\rm ST}\\gamma^{\\rm ST}\\cos(\\phi)\\sigma {\\rm ZHR}^{\\rm ST}_{\\rm Earth}({\\rm max})10^{-b|\\lambda-\\lambda_{\\rm max}|}+\\kappa}} , \\end{equation}pST=\u03bdST\u03b3STcos(\u03c6)\u03c3ZHREarthST(max)10\u2212b|\u03bb\u2212\u03bbmax|\u03bdSPO\u03b3SPOHREarthSPO+\u03bdST\u03b3STcos(\u03c6)\u03c3ZHREarthST(max)10\u2212b|\u03bb\u2212\u03bbmax|+\u03ba,where (16)\\begin{equation} \\kappa=\\sum_{i\\,=\\,1}^{n}\\nu^{\\rm ST}_i\\gamma^{\\rm ST}_i \\cos(\\phi_i)\\sigma {\\rm ZHR}^{\\rm ST}_{i,{\\rm Earth}}({\\rm max})10^{-b_i|\\lambda_i-\\lambda_{i,{\\rm max}}|} \\end{equation}\u03ba=\u2211i\u2009=\u20091n\u03bdiST\u03b3iSTcos(\u03c6i)\u03c3ZHRi,EarthST(max)10\u2212bi|\u03bbi\u2212\u03bbi,max|accounts for these n additional streams. The minimum kinetic energy Emin defined above corresponds to the minimum radiated energy Er_min on the Moon detectable from observations on Earth, which in turn is related to the maximum visual magnitude for detectable impacts (mmax). And these values depend on, among other factors, the experimental setup employed. The kinetic energy of the impactor and the radiated energy are linked by the luminous efficiency: (17)\\begin{equation} E_{r\\_{\\rm min}}=\\eta E_{\\rm min}. \\end{equation}Er_min=\u03b7Emin.With our experimental setup, the maximum visual magnitude for detectable impact is mmax ~ 10. The radiated energy can be obtained by integrating the radiated power P defined by the equation (18)\\begin{equation} P=1.36949\\times 10^{-16}10^{(-m+21.1)\/2.5}f \\pi \\Delta \\lambda R^2, \\end{equation}P=1.36949\u00d710-1610(\u2212m+21.1)\/2.5f\u03c0\u0394\u03bbR2,where P is given in Joules, m is the magnitude of the flash, 1.36949 \u00d7 10-16 is the flux density in W\u2009m-2\u2009\u03bcm-1 for a magnitude 21.1 source according to the values given in Bessel (1979), \u0394\u03bb is the width of the filter passband (about 6000 \u00c5 for our devices), and R is the Earth-Moon distance at the instant of the meteoroid impact. The factor f is related to the degree of anisotropy of light emission. Thus, for those impacts where light is isotropically emitted from the surface of the Moon f = 2, while f = 4 if light is emitted from a very high altitude above the lunar surface. As expected, according to Eq. (18) the minimum radiated power for flash detectability (and hence the minimum radiated energy and the minimum meteoroid kinetic energy), which is obtained by using m = mmax, is higher when the distance between Earth and Moon is greater. So, the detectability limit is time-dependent. ","Citation Text":["Swift et al. 2011"],"Functions Text":["However, the masses corresponding to impact flashes recorded on the Moon are several orders of magnitudes larger than mo","This means that the minimum kinetic energy Emin to produce a detectable impact flash on the Moon is much higher than the kinetic energy of a portion of the meteoroids included in the computation of hourly rates on Earth. This means that the velocity and mass distribution of meteoroids must be somehow included in Eq. (11) to take into consideration for the computation of the probability parameter p only those meteoroids with a kinetic energy above the threshold kinetic energy given by Emin"],"Functions Label":["Differences","Uses"],"Citation Start End":[[4365,4382]],"Functions Start End":[[4189,4309],[4385,4878]]}
{"Identifier":"2022ApJ...929....7V__Caprioli_&_Spitkovsky_2014a_Instance_1","Paragraph":"Diffusive shock acceleration (DSA), a process entering the category of first-order Fermi acceleration, is the process by which astrophysical shocks accelerate charged particles to relativistic speeds. DSA requires the magnetic field near the shock front to reflect the particle, leading to repeated shock-crossings with the particle gaining energy at each crossing (e.g., Bell 1978; Blandford & Ostriker 1978; Drury 1983). DSA is a self-sustaining process because the presence of high-energy particles triggers instabilities in the magnetic field, which in turn allow the magnetic field to reflect the particles more efficiently (e.g., Bell 1978, 2004). DSA has been explored numerically using the particle-in-cell (PIC) method, as well as the PIC-hybrid method, whereby the ions are treated as particles and the electrons as a fluid. Simulations of this type seem to indicate that although this process is effective in the case of (quasi-)parallel shocks where the magnetic field is aligned with the direction of motion, it becomes ineffective once the angle between the magnetic field and the shock exceeds approximately 50\u00b0 (Caprioli & Spitkovsky 2014a,2014b, 2014c).\n4\n\n\n4\nNotice that most of the studies on DSA including the present one assume shocks are propagating into a homogeneous medium; effects of a nonmonotonous shock are discussed in Hanusch et al. (2019). Most simulations also assume that the pre-shock medium is fully ionized. Simulations for a partially ionized medium (Ohira 2013, 2016) show an increased injection rate, allowing for DSA even in oblique shocks. Also, studies by Kumar & Reville (2021) demonstrate that DSA can be triggered by the shock emitted waves at highly oblique shocks. However, different results were obtained by van Marle et al. (2018), using a combined PIC-MHD approach. This method, which treats the thermal gas as a fluid but nonthermal ions as individual particles, showed that, given a sufficient injection rate, even for magnetic obliquity up to 70\u00b0, shocks can accelerate particles through the DSA process. Several factors were believed to contribute to this discrepancy between PIC-MHD and PIC or PIC-hybrid methods:1.Time-period covered by the simulation. The DSA occurring in PIC-MHD simulations is only effective after more than 200 ion-gyrotimes (\n\n\n\n\u03c9ci\u22121\n\n) and it takes more than 500 ion-gyrotimes to fully develop whereas published hybrid simulations where limited to scales of the order of 200 ion-gyrotimes. Notice that at the time \u223c200 \n\n\n\n\u03c9ci\u22121\n\n both PIC-hybrid and PIC-MHD simulations find similar results, namely that particles are accelerated by the shock-drift acceleration (SDA).2.Effective number of particles per cell. DSA requires an upstream current of CRs in order to develop. Because only a small fraction of the particles that cross the shock is reflected, a large particle population is required to produce sufficient nonthermal particles in the upstream medium.3.Simulation box size. van Marle et al. (2018) showed that the relevant upstream instabilities operate on long wavelengths, larger than the box size of hybrid simulations.\n","Citation Text":["Caprioli & Spitkovsky 2014a"],"Functions Text":["Simulations of this type seem to indicate that although this process is effective in the case of (quasi-)parallel shocks where the magnetic field is aligned with the direction of motion, it becomes ineffective once the angle between the magnetic field and the shock exceeds approximately 50\u00b0"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1128,1155]],"Functions Start End":[[835,1126]]}
{"Identifier":"2016AandA...593A..95C__Smith_1963_Instance_1","Paragraph":"Galactic fountains or infall are the possible origins for the ECs in this sample. Both scenarios present possible explanations and restrictions. The accretion of low metallicity gas from the intergalactic medium may occur through filamentary structures with the gas cooling into clouds (Fern\u00e1ndez et al. 2012). Subsequently HVCs may be destroyed or fragmented into smaller clouds by drag forces from the halo and phenomena such as Rayleigh-Taylor and Kelvin-Helmholtz instabilities. Star formation can be triggered by these interactions, within clouds that reach sufficient density (Figs. 9 and 10). However, there is evidence that star formation is possible only within dark-matter encapsulated HVCs such as the Smith Cloud (Smith 1963; Heitsch & Putman 2009; Nichols & Bland-Hawthorn 2009; Joung et al. 2012). Christodoulou et al. (1997) argue that without dark matter, HVCs are unable to reach the mass required to collapse. On the other hand, C 932, C 934, and C 939 are located right above the Local spiral arm (Fig. 11), which would be consistent with the chimney scenario. Schlafly et al. (2015) mapped various bubble-like structures vertically along the range 0.3 to 2.8 kpc, which form the Orion superbubble. The expansion of these substructures powered by massive stellar winds and supernovae triggers star formation in various shells and rings, inputting energy to the superbubble (Lee & Chen 2009). The star formation engine in the Galactic fountain may work in a similar way to the infall scenario, through the interaction of a cloud with the surrounding halo environment. However, high-latitude clusters in this study are located at distances from the disc larger than has been expected for a chimney-like event in recent studies (Melioli et al. 2009). Regardless of the scenario, a possible cloud-cloud interaction may be leading the clouds analysed in Fig. 9 to collapse and triggering not only star formation, but also cluster formation. However, the timescale for cloud-cloud collision in a cloud complex appears to be larger than 1 Gyr (Christodoulou et al. 1997). Therefore, additional studies are required to check the presence of dark-matter in HVCs, estimate the cloud-cloud interaction timescales, and provide more insight into the effect of the halo environment on the HVCs and chimney-like events. That is beyond the purpose of this study. ","Citation Text":["Smith 1963"],"Functions Text":["Star formation can be triggered by these interactions, within clouds that reach sufficient density (Figs. 9 and 10). However, there is evidence that star formation is possible only within dark-matter encapsulated HVCs such as the Smith Cloud"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[726,736]],"Functions Start End":[[483,724]]}
{"Identifier":"2021ApJ...917...24Z__Hall_&_Evans_2019_Instance_1","Paragraph":"As Advanced LIGO and Virgo will upgrade while additional interferometers (KAGRA and LIGO-India) come online in the future, the detection horizon would be several hundred megaparsecs for BNS and BH\u2013NS mergers with the localization accuracy reaching \u223c10 deg2 for a network of these 2nd generation GW detectors (Abbott et al. 2018). In contrast, higher localization accuracy would be obtained by networking of 3rd generation GW detectors (Zhao & Wen 2018; Chan et al.2018; Hall & Evans 2019; Maggiore et al. 2020). In these regimes, present and future wide-field-of-view survey projects will be able to cover the sky localization given by GW detections in a few pointings and achieve deep detection depths with relatively short exposure integration times. The detection of GW170817\/AT 2017gfo was made in an i-band filter and subsequently confirmed by larger telescopes (Coulter et al. 2017). Similar to this case, we assume that one can use only one filter to make follow-up searches after GW triggers. With quick response times of follow-up observations after GW events, we also assume that the survey telescopes would not miss the peaks of kilonovae in each band. Therefore, by assuming that all GW detectors in each era can operate normally (i.e., the most optimal situation for detector networks of different generations), the detection rates for target-of-opportunity observations for a given X-band limiting magnitude mX,limit can be expressed as\n15\n\n\n\n\n\n\n\n\nN\n\n\n\n\n\n\n\nX\n\n\n=\n\n\n\u222b\n\n\n0\n\n\n\n\nz\n\n\nmax\n\n\n\n\n\n\n\n\n\n\n\n\u03c1\n\n\n\n\n\n\n\n\n\n\nm\n\n\nX\n,\nlimit\n\n\n\n\n(\nz\n)\n\n\n1\n+\nz\n\n\n\n\n\n\ndV\n(\nz\n)\n\n\ndz\n\n\n\ndz\n,\n\n\nwhere \n\n\n\n\n\n\nz\n\n\nmax\n\n\n\n\n is the maximum detectable distance and \n\n\n\n\n\n\n\n\n\u03c1\n\n\n\n\n\n\n\n\n\n\nm\n\n\nX\n,\nlimit\n\n\n\n\n(\nz\n)\n\n\n represents the cosmological event rate densities of the BH\u2013NS mergers that can be triggered by GW detectors and be detected by telescopes in band X. Here, the astronomical magnitude system we adopted is the AB magnitude system, i.e.,\n\n\n\n\n\n\nm\n\n\n\u03bd\n\n\n=\n\u2212\n2.5\nlog\n(\n\n\nF\n\n\n\u03bd\n\n\n\n\/\n\n3631\n\nJy\n)\n\n\n.","Citation Text":["Hall & Evans 2019"],"Functions Text":["In contrast, higher localization accuracy would be obtained by networking of 3rd generation GW detectors"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[470,487]],"Functions Start End":[[330,434]]}
{"Identifier":"2021AandA...646A..96C__Freundlich_et_al._2019_Instance_1","Paragraph":"Observations of cold molecular gas are more promising since any activity from the AGN (radiation, outflows, or jets) will have an impact on the molecular gas reservoir first, and then on the SFR, which previous studies have focused on. The molecular gas provides an instantaneous measure of the raw fuel from which stars form and can be used as a more direct tracer of potential feedback effects. Over the last decade, large observational efforts have been devoted to map the molecular gas reservoir of galaxies (e.g., Daddi et al. 2010a; Garc\u00eda-Burillo et al. 2012; Bauermeister et al. 2013; Bothwell et al. 2013; Tacconi et al. 2013, 2018; Genzel et al. 2015; Silverman et al. 2015; Decarli et al. 2016; Cicone et al. 2017; Saintonge et al. 2017; Fluetsch et al. 2019; Freundlich et al. 2019). These studies are largely based on observations of carbon monoxide (CO) rotational emission lines, used as a tracer of cold molecular hydrogen H2 (the ground-state rotational transition in particular traces the total molecular gas best), but there are also many studies based on dust emission (e.g., Tacconi et al. 2018 and references therein). The main targets of such CO campaigns have primarily been inactive1 galaxies that mostly lie on the main sequence of star-forming galaxies (e.g., Noeske et al. 2007; Schreiber et al. 2015), where the majority of the cosmic star-formation activity occurs. The fundamental relation between SFR and the molecular gas content of galaxies, the Schmidt-Kennicutt relation, provides precious information about how efficiently galaxies turn their gas into stars (Schmidt 1959; Kennicutt 1989). This star-formation law is usually presented in terms of surface densities, therefore requiring spatially resolved measurements of galaxies. However, in high-redshift studies, an integrated form of this relation with global measurements of the SFR and molecular mass is normally used (e.g., Carilli & Walter 2013; Sargent et al. 2014).","Citation Text":["Freundlich et al. 2019"],"Functions Text":["Over the last decade, large observational efforts have been devoted to map the molecular gas reservoir of galaxies (e.g.,","These studies are largely based on observations of carbon monoxide (CO) rotational emission lines, used as a tracer of cold molecular hydrogen H2 (the ground-state rotational transition in particular traces the total molecular gas best),"],"Functions Label":["Background","Background"],"Citation Start End":[[771,793]],"Functions Start End":[[397,518],[796,1033]]}
{"Identifier":"2020MNRAS.492.5675C__P\u00e9rez-Montero_et_al._2019_Instance_1","Paragraph":"In regarding AGNs, the Te method tends to underestimate the oxygen abundance by an average value of about 0.6 dex in comparison to estimations based on strong-line methods and it produces subsolar O\/H values for most of these objects (Dors et al. 2015; Dors, Freitas-Lemes & \u00c2mores 2020). An alternative method to derive the metallicity or abundances in the nuclear regions of spiral galaxies is the extrapolation of the radial oxygen abundance. Along decades, results based on this indirect method have indicated Z near or slightly above the solar value in nuclear regions (Vila-Costas & Edmunds 1992; Zaritsky, Kennicutt & Huchra 1994; van Zee et al. 1998; Pilyugin, V\u00edlchez & Contini 2004; Gusev et al. 2012; Dors et al. 2015; Zinchenko et al. 2019), in consonance with predictions of chemical evolution models (e.g. M\u00f3lla & D\u00edaz 2005) and with the use of strong-line methods (e.g. Groves, Dopita & Sutherland 2004; Groves, Heckman & Kauffmann 2006; Feltre, Charlot & Gutkin 2016; P\u00e9rez-Montero et al. 2019; Thomas, Kewley & Dopita 2019; Dors et al. 2020). Therefore, Temethod does not seem to work for AGNs. The origin of the discrepancy between Z values calculated via Te method and via strong-line methods, the so-called Teproblem, could be attributed, in part, to the presence of heating\/ionization by gas shock in the narrow-line region (NLR) of AGNs. In fact, Contini (2017) carried out detailed modelling of AGN optical emission lines by using the SUMA code (Contini & Aldrovandi 1983) and suggested the presence of gas shock with low velocity ($v \\: \\lesssim \\: 400 \\: \\rm km \\:s^{-1}$) in a sample of Seyfert 2 nuclei. This result is supported by recent spatially resolved observational studies of Seyfert 2 nuclei, in which the presence of gas outflows with velocity of the order of 100\u2013300 $\\rm km \\: s^{-1}$ have been found (e.g. Riffel, Storchi-Bergmann & Riffel 2017; Riffel, Hekatelyne & Freitas 2018). Moreover, the Te problem can also be originated due to the use of an unappropriate calculation of the ionization correction factor (ICF) for oxygen in AGNs (P\u00e9rez-Montero et al. 2019; Dors et al. 2020).","Citation Text":["P\u00e9rez-Montero et al. 2019"],"Functions Text":["Along decades, results based on this indirect method have indicated Z near or slightly above the solar value in nuclear regions","in consonance with predictions of chemical evolution models","and with the use of strong-line methods"],"Functions Label":["Similarities","Similarities","Similarities"],"Citation Start End":[[984,1009]],"Functions Start End":[[446,573],[754,813],[839,878]]}
{"Identifier":"2017MNRAS.471...80S__Kalugina_et_al._2014_Instance_2","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"],"Functions Text":["Benchmarks by Hochlaf and co-workers","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."],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[1585,1605]],"Functions Start End":[[1466,1502],[1637,1895]]}
{"Identifier":"2021AandA...653A..83V__Kr\u00fchler_et_al._(2015)_Instance_1","Paragraph":"The comparisons show that except for GRBs 060707 and 060605, which have a very high vpeak, the LAE-LGRBs fall in the same parameter space as LyC leakers and follow the correlations between the indirect indicators found by Verhamme et al. (2017). Following their study, we can cut these plots into two regions corresponding to strong LyC leakers (fesc(LyC) > 5%, red rectangle) and weak LyC leakers (fesc(LyC)  5%, blue rectangle). All LAE-LGRBs fall in the category of the weak LyC leakers except for GRB 021004, which appears systematically to agree well with the region of the strong leakers. In panel b of Fig. 12, we also superimpose the distribution of [OIII]\/[OII] ratio for the GRB sample of Kr\u00fchler et al. (2015) and GRB 081121 reported here. For the majority of the GRBs, this ratio is about two, with seven cases at O32 > 4. GRB 021004 has a high value of O32 > 10 and is the strongest LAE of our golden sample (with fesc(Ly\u03b1) of 60%) in agreement with potential high fesc(LyC). Nevertheless, the Ly\u03b1 profile of GRB 021004 is a single peak with no residual flux at the Ly\u03b1 line center. This is not the typical line shape observed for confirmed LyC emitters, which have the tendency to show double- or triple-peak profiles (Verhamme et al. 2017; Rivera-Thorsen et al. 2017; Vanzella et al. 2020). Our shell-model fitting also suggests that the column densities are too high to allow LyC photons to escape. However, this model already predicted a similar HI column density (log(NHI\/cm\u22122)\u2005\u2248\u200519\u221220) for four green pea galaxies out of the five with detected LyC emission reported in Yang et al. (2017). This suggests that LyC emission can escape through holes in the ISM even if the Ly\u03b1 photons probe denser neutral gas. The only LAE-LGRB for which LyC leakage has been detected along the LGRB line of sight (\n\n\n\n\nf\nesc\n\n\n(\nLyC\n)\n\n=\n0\n.\n\n35\n\n\u2212\n0.11\n\n\n+\n0.10\n\n\n\n\n$ f_{\\mathrm{esc}}(\\rm LyC) = 0.35^{+0.10}_{-0.11} $\n\n\n) is GRB 191004B (Vielfaure et al. 2020). However, Fig. 12 (panel c) shows that it does not fall in the high escape fraction region, but in the lower left area. The reasons could be that (i) fesc(LyC) is lower at the scale of the galaxy than along the LGRB line of sight, and (ii) the indicators of strong LyC leakage evolves with redshift. As a comparison, the LyC emitters from Fletcher et al. (2019) are found out of the high escape region (red rectangle). They show lower rest-frame EW(Ly\u03b1) and higher vpeak than the local LyC emitters, whereas their escape fraction of ionizing photons is significantly higher (fesc(LyC) = 15 \u2212 60%). This could also suggest that strong LyC leakers span a wider parameter space than predicted by the study of local LyC emitters. Overall, it is clear that this type of studies is still limited by the poor statistics, and the current results show the difficulty of characterizing LyC leakage based on these properties alone.","Citation Text":["Kr\u00fchler et al. (2015)"],"Functions Text":["In panel b of Fig. 12, we also superimpose the distribution of [OIII]\/[OII] ratio for the GRB sample of","and GRB 081121 reported here.","For the majority of the GRBs, this ratio is about two, with seven cases at O32 > 4. GRB 021004 has a high value of O32 > 10 and is the strongest LAE of our golden sample (with fesc(Ly\u03b1) of 60%) in agreement with potential high fesc(LyC). Nevertheless, the Ly\u03b1 profile of GRB 021004 is a single peak with no residual flux at the Ly\u03b1 line center."],"Functions Label":["Uses","Uses","Compare\/Contrast"],"Citation Start End":[[699,720]],"Functions Start End":[[595,698],[721,750],[751,1095]]}
{"Identifier":"2018MNRAS.475.1160H__Werk_et_al._2013_Instance_2","Paragraph":"Galaxies are surrounded by vast gaseous haloes which extend well beyond the hosts\u2019 stellar components: Early observations of quasar sight lines attributed the presence of absorption at multiple intermittent redshifts to gaseous haloes of intervening galaxies (e.g. Bergeron 1986; Bergeron & Boiss\u00e9 1991; Lanzetta et al. 1995; Tripp, Savage & Jenkins 2000; Chen, Lanzetta & Webb 2001). In the past decade, owing to the rise of large spectroscopic surveys of galaxies with well-determined physical properties (e.g. SDSS), all sky UV surveys (e.g. GALEX), and improved sensitivity of UV spectrographs (e.g. COS), studies of the gaseous haloes of galaxies could systematically connect gas absorption properties to galaxy properties in statistically meaningful samples (e.g. Cooksey et al. 2010; Prochaska et al. 2011; Tumlinson et al. 2013; Liang & Chen 2014; Lehner, Howk & Wakker 2015). The aforementioned gaseous haloes are commonly referred to as the circum-galactic medium (CGM) and are ubiquitous in galaxies regardless of mass or star formation activity: even sub-L* galaxies (Bordoloi et al. 2014), and passive galaxies host a CGM (Thom et al. 2012). The current model of the CGM suggests the presence of a clumpy multiphase medium which extends beyond the virial radius of the host galaxy, with a declining radial density profile, containing a substantial amount of gas and metals (e.g. Werk et al. 2013, 2014; Liang & Chen 2014; Lehner et al. 2014, 2015; Prochaska et al. 2017). Observational studies targeting the CGM of L* galaxies showed that the CGM gas content is comparable to the mass of the interstellar medium (ISM; e.g. Chen et al. 2010; Tumlinson et al. 2011; Werk et al. 2014; Prochaska et al. 2017) and correlates positively with ISM properties (Borthakur et al. 2015). Additionally, CGM observations infer a significant amount of metals (e.g. Werk et al. 2013; Peeples et al. 2014) where CGM metallicities can extend to supersolar metallicities (Prochaska et al. 2017). The clumpy multiphase CGM consists of a warm gas T \u223c 104 \u2212 5\u2009K (clumpy in nature) embedded within a hot diffuse T \u223c 106\u2009K medium (e.g. Heitsch & Putman 2009; Armillotta et al. 2017; Bordoloi et al. 2017). The multiphase structure of the CGM is corroborated by the variety of observed ionic species which survive at a vast range of temperatures: While the warm gas hosts the low ionization species (e.g. H\u2009i, Si\u2009ii, Si\u2009iii, C\u2009ii, C\u2009iv), the hot medium is home for the most highly ionized species (e.g. O\u2009vi, O\u2009vii). Additionally, the spectral line profiles of absorbers in the CGM can be reproduced by invoking a patchy medium (e.g. Stern et al. 2016; Werk et al. 2016), i.e. multiple high density gas clouds contribute to the optical depth along the line of sight thus leaving their kinematic imprint on the absorption line profile. For a review of the CGM, see Putman, Peek & Joung (2012) and Tumlinson, Peeples & Werk (2017).","Citation Text":["Werk et al. 2013"],"Functions Text":["Additionally, CGM observations infer a significant amount of metals (e.g."],"Functions Label":["Background"],"Citation Start End":[[1863,1879]],"Functions Start End":[[1789,1862]]}
{"Identifier":"2019AandA...630A..30L__H\u00e4ssig_et_al._(2015)_Instance_2","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"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[1007,1025]],"Functions Start End":[[835,1005]]}
{"Identifier":"2022MNRAS.513.5245A__Done_&_Jin_2016_Instance_2","Paragraph":"We assume the time-scales we observe here are generated in the corona itself (and note that longer time-scale changes will be driven by the disc outside of the corona) and are made visible by a changing electron temperature and density as a result of local turbulence and coupling to mass accretion rate propagations through the flow from rout to rin. We note that we can discount variations in the seed photon population as the driver for changes in the power spectrum, as the UV emission from the disc is established to be considerably less variable than the corona in NLS1s (Leighly 1999; Smith & Vaughan 2007; Ai et al. 2013; Alston, Vaughan & Uttley 2013; Done & Jin 2016). We assume that the variability generated locally at each radius (r\u03bd) is at the viscous frequency (see Churazov et al. 2001) such that\n(9)$$\\begin{eqnarray}\r\nr_{\\nu } = \\left[ \\frac{2\\pi \\nu }{\\alpha }\\left(\\frac{H}{R}\\right)^{-2}\\right]^{-2\/3} ,\r\n\\end{eqnarray}$$where the frequency is in units of c\/Rg (see e.g. Kato, Fukue & Mineshige 1998; Ar\u00e9valo & Uttley 2006). In the above, \u03b1 and $\\frac{H}{R}$ are the dimensionless viscosity parameter of Shakura & Sunyaev (1973) and scale height of the accretion disc, respectively. We assume that our frequency range of interest, 0.01\u22121 mHz (i.e. the range over which we can practically fit to the data) corresponds to radii between the ISCO (rin) and some radius within the true outer edge of the corona (i.e. rout \u2264 rcorona). The actual frequencies generated in our model therefore depend on the SMBH spin and the combination $\\left(\\frac{H}{R}\\right)^2 \\alpha$ (and somewhat on the SMBH mass \u2013 although here the range is small). Given the reported high spin values for these bright AGNs (Ogle et al. 2004; Fabian et al. 2013; Done & Jin 2016; Kara et al. 2017; Buisson et al. 2018b), we expect the ISCO to sit at \u223c1.25Rg. For the corona at the ISCO to produce variability above our upper frequency limit of 1 mHz requires $\\left(\\frac{H}{R}\\right)^2 \\alpha \\gtrsim 0.01$. We note that for the mass range subtended by our AGN sample (from 106.00 to 106.63\u2009M\u2299, see Table 2), should we instead assume zero spin (rin = 6Rg), the viscous frequency at rin is lower and we have strong curvature in our observed bandpass.","Citation Text":["Done & Jin 2016"],"Functions Text":["Given the reported high spin values for these bright AGNs","we expect the ISCO to sit at \u223c1.25Rg."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1751,1766]],"Functions Start End":[[1654,1711],[1809,1846]]}
{"Identifier":"2015MNRAS.450.3458C__Cichowolski_et_al._2001_Instance_4","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"],"Functions Text":["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"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1012,1035]],"Functions Start End":[[852,1010]]}
{"Identifier":"2020MNRAS.491..560D__Straten_&_Manchester_2004_Instance_1","Paragraph":"As MKT J170456.2\u2013482100\u2009is coincident with a star in a spectroscopic binary, it is possible that the companion star may be a pulsar or pulsating white dwarf, similar to AR Scorpii (Marsh et al. 2016). To investigate this possibility, we performed high-time resolution observations to search for pulsations over a range of periods spanning milliseconds to minutes. We observed the position of MKT J170456.2\u2013482100\u2009on utc 2019 January 30 with the ultra-wide-bandwidth (UWL) receiver deployed on the 64-m Parkes radio telescope in Australia. The output of the receiver is processed by the Medusa Graphics Processing Unit (GPU) cluster which produces 8-bit, full-Stokes filter banks with spectral and temporal resolutions of 1 MHz and 128 \u03bcs respectively over the $3328\\, \\mathrm{MHz}$ band from 0.740 to $4.03\\, \\mathrm{GHz}$. The data were recorded in psrfits format and processed using PSRCHIVE tools (Hotan, van Straten & Manchester 2004). Since our data were strongly affected by radio frequency interference (RFI) we used the clfd8 package described in (Morello et al. 2019) to perform more sophisticated RFI mitigation. Since radio waves are dispersed by free electrons in the interstellar medium along the line of sight, the data need to be corrected for this dispersion delay before any analyses can be made. This dispersion delay can be quantified by the dispersion measure (DM), which is the integrated free electron column density along the line of sight. We dedispersed the data over a range of trial DMs, $0.0 \\le \\mathrm{DM} \\le 30.0$ pc cm$^{-3}$ in steps of 1 pc cm$^{-3}$ to account for the uncertainty in the distance to TYC 8332-2529-1. The DM range was estimated using various distances to the source (Bailer-Jones et al. 2018), and the NE2001 model (Cordes & Lazio 2002). For each trial DM, the resulting dedispersed time series was searched for long and short-period pulsations using both a Fast Folding Algorithm (FFA) and a Fast Fourier Transform (FFT), respectively.","Citation Text":["Hotan, van Straten & Manchester 2004"],"Functions Text":["The data were recorded in psrfits format and processed using PSRCHIVE tools"],"Functions Label":["Uses"],"Citation Start End":[[901,937]],"Functions Start End":[[824,899]]}
{"Identifier":"2016AandA...585A..76W__Neau_et_al._2000_Instance_1","Paragraph":"The median value of the column density of OH in the first quadrant is 3.9 \u00d7 1014 cm-3 and for OH+ it is 0.68 \u00d7 1014 cm-3. In the fourth quadrant, the median column densities amount to 1.7 \u00d7 1014 cm-3 and 0.55 \u00d7 1014 cm-3, respectively. Column densities in excess of 1015 cm-2 (for OH) and above ~1014 cm-2 (for OH+) are rather the exception. The median N(OH) \/N(OH+) ratio over all sightlines and velocity components is 3.3. While the formation of OH+ results from cosmic-ray induced reactions involving atomic and molecular hydrogen and atomic oxygen, the bottleneck expected for the formation of OH via ion-neutral chemistry is the availability of H2. The reaction of OH+ with H2 yields H3O+, via the two hydrogen abstraction reactions OH+(H2,H)H2O+(H2,H)H3O+ (see Appendix C). Then the dissociative recombination of H3O+ yields OH and H2O, with a branching ratio of ~74% to 83% in favor of OH (determined by ion storage ring experiments, Jensen et al. 2000; Neau et al. 2000), while less than ~1% forms O\u2009i\u2009. In the following, we attempt to confirm these predictions, that is, the bottleneck reaction OH+(H2,H)H2O+, and the N(OH) \/N(H2O) ratio. As for the former, a strong anticorrespondence between the column densities of OH and OH+might naively be expected, where the availability of H2 tips the scales in favor of OH, while OH+ traces predominantly atomic gas (Hollenbach et al. 2012, further references therein). But even if a clear anticorrelation between OH and OH+ existed, it would be impossible to observe it. On a given sightline several clouds with high and low molecular hydrogen fractions \\hbox{$f^{\\rm N}_\\HH = N(\\HH)\/(2N(\\HH)+N($}fH2N=N(H2)\/(2N(H2)+N(H\u2009i\u2009)) line up. Even across a single diffuse cloud, the N(OH) \/N(OH+) ratio is expected to vary substantially, depending on the degree of self-shielding of H2 against the interstellar UV radiation field. To quantify the anticorrelation, we normalized the velocity-specific OH column density with the total OH and OH+ reservoir and obtained an abundance ratio r = Nv(OH)\/(Nv(OH) + Nv(OH+) varying from zero (only OH+, no OH) to one, where all the OH+ abundance is exhausted owing to the formation of OH and (see below) water. (The normalization chosen here avoids the divergence of the distribution if OH has no spectral counterpart in OH+.) The resulting distribution (Fig. 12) indeed shows that these extremes are present in the data, although the second case is by an order of magnitude more frequent. We suggest two explanations for this. One reason is that if \\hbox{$f^{\\rm N}_\\HH$}fH2N is too small, OH+ can be efficiently destroyed by the dissociative recombination with free electrons (Appendix C), while the formation of OH+ by the reaction chain H+(O,H)O+(H2,H)OH+ and the secondary, less important path H2+(H2,H)H3+(O,H2)OH+ become less efficient (see Appendix D) because less H2 is available. Another reason is that the fraction \\hbox{$f^{\\rm N}_\\HH$}fH2N is larger in denser gas (cf. Table 3) where column densities are higher and absorption features easier to observe. ","Citation Text":["Neau et al. 2000"],"Functions Text":["determined by ion storage ring experiments,"],"Functions Label":["Background"],"Citation Start End":[[961,977]],"Functions Start End":[[897,940]]}
{"Identifier":"2017AandA...605A..20C__Momjian_&_Sarma_2017_Instance_1","Paragraph":"In principle, when \u0394\u03bdZ>\u03b4\u03bd, with \u03b4\u03bd being the observed full width at half maximun (FWHM) of the line, the complete set of information concerning the magnetic field B can be derived. However, for most of the Zeeman detections, the Zeeman splitting \u0394\u03bdZ turns out to be significantly smaller than \u03b4\u03bd. Indeed, even if we neglect the non-thermal line broadening due to typical processes occurring in star forming regions such as jets (with velocities up to hundreds of km\u2009s-1), outflows, and accreting\/rotating disks (typically \u226410 km\u2009s-1), already the thermal broadening (even at kinetical temperatures lower than 10 K) is definitely larger than the Zeeman splitting (e.g., Frank et al. 2014). For instance, a typical linewidth of the coldest starless cores is ~0.1 km\u2009s-1  (e.g., di Francesco et al. 2007), indeed too broad to allow an observer to unveil the Zeeman effect. The expected B values in star forming regions range from ~100 \u03bcG, for low-mass objects (Li et al. 2014), to ~1 mG, as measured in regions hosting high-mass young stars (e.g., Pillai et al. 2016; Momjian & Sarma 2017). If we consider the N,J = 1, 1 \u2190 0, 1 transition (286.3 GHz), for which the largest Zeeman shift was obtained (Z = 0.87 Hz\/\u03bcG), and we assume that B varies in the 0.1\u20131 mG range, we derive \u0394\u03bdZ\u2243 values in the 87\u2013870 Hz range, that is, 0.3\u20133 m\u2009s-1. Moving to higher frequencies and considering the SO N,J = 2, 2 \u2190 1, 2 transition at 309.5 GHz, we derive \u0394\u03bdZ \u2243 61\u2212610 Hz, that is, 0.06\u22120.59 m\u2009s-1. In conclusion, in such cases, only information about the LoS component of B, B\u2225, is obtained. In fact, the Stokes parameter Q and U spectra are usually too weak to be detected. The analysis is therefore limited to the Stokes V spectrum, which has the shape of the first derivative of the Stokes I spectrum and is proportional to (\u0394\u03bdZ\/\u03b4\u03bd)2 \u00d7 B\u2225. According to Crutcher et al. (1993), Crutcher (2012), and Uchida et al. (2001), the standard procedure is to fit dI\/d\u03bd to the observed V spectrum, with B\u2225 being the free parameter to be determined from the strength of the V spectrum. Furthermore, the direction (toward or away from the observer) of B\u2225 is obtained. Therefore, accurate laboratory (either experimental or theoretical) determination of the g factors might provide the opportunity to complete the missing information. ","Citation Text":["Momjian & Sarma 2017"],"Functions Text":["The expected B values in star forming regions range from ~100 \u03bcG, for low-mass objects",", to ~1 mG, as measured in regions hosting high-mass young stars (e.g.,"],"Functions Label":["Uses","Uses"],"Citation Start End":[[1065,1085]],"Functions Start End":[[870,956],[973,1044]]}
{"Identifier":"2020ApJ...898...92C__Lawson_et_al._2012_Instance_1","Paragraph":"Neutral MF occurs as syn and anti conformers (also denoted Z and E rotamers) with Cs symmetry, with the syn rotamer being more stable (by \u223c20 kJ mol\u22121) and well separated from the anti rotamer by a significant energy barrier of \u223c35 kJ mol\u22121 for anti \n\n\n\n\n\n syn isomerization (Wilmshurst 1957; Curl 1959; Blom & G\u00fcnthard 1981; M\u00fcller et al. 1983; Bagno & Scorrano 1996; Neill et al. 2011, 2012; Zeegers-Huyskens & Kryachko 2011; Ferro-Costas & Mosquera 2013). The zero-kinetic-energy photoelectron spectrum of MF provides an accurate estimate for its adiabatic ionization energy and vibrational frequencies of the radical cation ground state (Waterstradt et al. 1994). The electronic excitation spectrum of MF is very similar to that of acetic acid featuring three valence transitions (Nunes et al. 2010). In contrast to MF and its radical cation (MF+), only very limited information is available for H+MF. Theoretical studies indicate the preference for protonation at the carbonyl oxygen (CO) over the methyl ester (OMe) by around 80 kJ mol\u22121, and the CO-protonated tautomer has four rotamers within 20 kJ mol\u22121 (Zeegers-Huyskens & Kryachko 2011; Ferro-Costas & Mosquera 2013). Numerous mass spectrometric studies have investigated the production of H+MF via various ion\u2013molecule reactions and its fragmentation by metastable decay and collisions (Harrison & Tsang 1976; Benoit & Harrison 1977; Hopkinson et al. 1979; van Baar et al. 1986; Horn et al. 2004; Lawson et al. 2012), and the measured proton affinity (PA) is reported as PA = 782.5 kJ mol\u22121 (Hunter & Lias 1998). However, the site of protonation and the rotational conformation of H+MF could not be established from these mass spectrometric measurements. The dominant fragment channel of H+MF is decarbonylation (loss of CO), whereby the C and O atoms come from the CO group. Significantly, there are no spectroscopic data available for H+MF and its clusters. To this end, our combined spectroscopic and DFT approach provides the first reliable experimental evidence for the preferred protonation site in the isolated molecule and the isomer assignment along with the intermolecular interaction of this prototypical protonated aliphatic ester with nonpolar ligands. There are a few experimental data for H+MF in the condensed phase. However, it is well known that solvation and counter ions can have a drastic effect on both the position and energetics of protonation in the various possible tautomers and rotamers. Previous 1H-NMR spectra of H+MF revealed three OH+ resonances in a super acid solution, demonstrating protonation at the CO group with three rotamers (although their configurations could not be established) (Fraenkel 1961; Olah et al. 1967). IR, Raman, and 1H\/13C-NMR spectra of single crystal salts of H+MF with \n\n\n\n\n\n and \n\n\n\n\n\n show only a single isomer (the most stable (ts) conformer), which forms a strong ionic hydrogen bond (H band) between the OH+ group and the anion (Minkwitz et al. 2000). The characteristic OH stretch frequency is thus largely redshifted and thus not identified in both the IR and Raman spectra of the salts. Also, the IR intensities of the CH stretches are too weak to be observed.","Citation Text":["Lawson et al. 2012"],"Functions Text":["Numerous mass spectrometric studies have investigated the production of H+MF via various ion\u2013molecule reactions and its fragmentation by metastable decay and collisions"],"Functions Label":["Background"],"Citation Start End":[[1459,1477]],"Functions Start End":[[1179,1347]]}
{"Identifier":"2022MNRAS.514.5192L__Borsa_et_al._2021_Instance_1","Paragraph":"We introduce the quantity \u03be to represent a scaled product of equilibrium temperature (Teq) and surface gravity (g):\n(14)$$\\begin{eqnarray}\r\n\\xi = \\left(\\frac{T_{\\text{eq}}}{1000~\\text{K}}\\right)\\left(\\frac{g}{g_{\\rm J}}\\right) ~,\r\n\\end{eqnarray}$$where gJ is the surface gravity of Jupiter. The calculated \u03be values are listed in Table 4, along with the equilibrium temperature, surface gravity, and scale height for each planet. Fig. 6 shows the relative height of sodium for all 10 planets against \u03be. The weight-combined results are shown in dark-blue, together with the unweighted results (light-blue) and literature values using the same HARPS\/HARPS-N data (grey). The salmon-coloured points show recent results from other high-resolution spectrographs: WASP-69b with CARMENES (Khalafinejad et al. 2021), WASP-76b with ESPRESSO (Tabernero et al. 2021) and GRACES (Deibert et al. 2021), and WASP-121b with ESPRESSO (Borsa et al. 2021). We find that hNa is well described by an exponential trend of the form\n(15)$$\\begin{eqnarray}\r\nh_{\\lambda } = a\\text{e}^{-b\\xi } + c ~.\r\n\\end{eqnarray}$$As shown in Fig. 6, we fit this curve to the weight-combined results using the optimize.curve_fit function from scipy. The best-fitting values for the variables are: a = 1.70 \u00b1 1.04, b = 5.04 \u00b1 1.63, and c = 0.113 \u00b1 0.013. The reduced chi-square of the fit across the full sample is $\\chi ^2_\\nu = 3.7$, much of which is skewed by two planets that deviate from the fit by more than 3\u03c3 (WASP-79b and HD 189733 b). The reduced chi-square when excluding these two planets is $\\chi ^2_\\nu = 1.2$. Since the exponential curve is asymptotic to the value of c, our results suggest that planets with \u03be \u2273 1.25 are likely to have an upper limit on hNa of \u223c0.113. The sodium features could potentially be muted due to other atmospheric effects, such as high-altitude clouds and hazes in lower temperature planets, or ionization of most of the sodium in extremely irradiated environments. Therefore, results lower than those given by equation (15) are also possible. The underlying physical processes behind this trend are not discussed in this current work, but these results motivate further observations to confirm the trend and theoretical studies to investigate possible physical mechanisms. Further refinement of this curve will be possible with more observations of low-\u03be planets.","Citation Text":["Borsa et al. 2021)"],"Functions Text":["The salmon-coloured points show recent results from other high-resolution spectrographs:","WASP-121b with ESPRESSO"],"Functions Label":["Uses","Uses"],"Citation Start End":[[918,936]],"Functions Start End":[[668,756],[893,916]]}
{"Identifier":"2017MNRAS.471.3856T__Tempel_et_al._2011_Instance_1","Paragraph":"It is important to note that we have not classified our simulated galaxies by morphology itself in this paper, and more detailed analysis of the morphology, including kinematical decomposition, will be made in future work. In Fig. 4, we show the distributions of our morphology indicators, ETG\u0394s, LTG\u0394s, ETG\u0394t, and LTG\u0394t, with stellar mass. The top panel shows that LTG\u0394t\u2009 dominate at low mass, and ETG\u0394t, which have experienced a major merger within 1 Gyr, are more numerous above M* \u223c 5 \u00d7 1010M\u2299; the ETG\u0394t fraction is \u223c45\u2009per cent and \u223c85\u2009per cent at M* \u223c 5 \u00d7 1010M\u2299 and \u223c1011M\u2299, respectively. Since mergers are strong drivers of morphological change, this implies that our high-mass galaxies are mostly elliptical, which is in good agreement with observational results for ellipticals and S0 galaxies (e.g. Bernardi et al. 2010; Tempel et al. 2011). There is not a perfect correlation between galaxy morphology and tmerge, however, and the details of each individual merger are important. In the lower panel of Fig. 4, galaxies are separated by \u0394SFMS. Galaxies that lie close to the SFMS (LTG\u0394s) are most numerous at lower masses, and constitute a declining fraction of the galaxy population with increasing mass, but the fraction of ETG\u0394s is not perfectly consistent with observations (e.g. Renzini & Peng 2015, \u223c50\u2009per cent and \u223c75\u2009per cent at log\u2009M* = 10.5 and 11). This means that our AGN feedback is probably not strong enough (see TK15a for more discussion). Our predicted distribution of gradients could be affected by these effects. Our galaxies that are classified as both LTG\u0394s and ETG\u0394t should be ellipticals, and their metallicity gradients may be too steep because of the lack of quenching of star formation. Therefore, we think our sample of ETG\u0394s is better for the comparison with observations of SAURON and ATLAS3D. For low-mass galaxies, the lack of resolution could result in underestimating gradients in general. However, our low-mass LTG\u0394s and LTG\u0394t galaxies, it could also overestimate the gradients in the case that the gradients would be measured along small star-forming discs, which are not resolved in our simulations.","Citation Text":["Tempel et al. 2011"],"Functions Text":["Since mergers are strong drivers of morphological change, this implies that our high-mass galaxies are mostly elliptical, which is in good agreement with observational results for ellipticals and S0 galaxies (e.g."],"Functions Label":["Similarities"],"Citation Start End":[[833,851]],"Functions Start End":[[597,810]]}
{"Identifier":"2018MNRAS.479.3254V___2000_Instance_1","Paragraph":"The lifetime of molecular clouds (MCs) remains an active research topic in the study of the interstellar medium and star formation, and most recent studies, both observational and theoretical, place this lifetime at a few times 107 yr for clouds in the 105\u2013106M\u2299 mass range (e.g. Blitz & Shu 1980; Kawamura et al. 2009; Zamora-Avil\u00e9s, V\u00e1zquez-Semadeni & Col\u00edn 2012; Zamora-Avil\u00e9s & V\u00e1zquez-Semadeni 2014; Lee, Miville-Desch\u00eanes & Murray 2016). In addition, several observational studies have suggested that the star formation rate (SFR) of the clouds appears to increase over their lifetimes. For example, studies of young clusters embedded in moderate-mass MCs (\u223c104M\u2299) (e.g. Palla & Stahler 1999, 2000; Da Rio et al. 2010) have shown that their age histograms contain a large majority of young (1\u20132 Myr) objects, but also a tail of older (up to several Myr) ones suggesting an accelerating star formation activity, sometimes followed by a subsequent decline (see also Povich et al. 2016; Schneider et al. 2018). In addition, Kawamura et al. (2009) reported a clear evolutionary process over the lifetime of giant molecular clouds (GMCs; of masses \u223c105\u2013106M\u2299) in the Large Magellanic Cloud, evidenced by the increasing number of massive stars across the sequence of GMC \u2018classes\u2019 proposed by those authors. Finally, on the basis of the large scatter in the observed star formation efficiency (SFE) in Milky Way GMCs, Lee et al. (2016) have concluded that the SFR in those clouds must also be time variable. Numerical simulations of MC formation and evolution also exhibit time varying, increasing SFRs during their early stages (e.g. V\u00e1zquez-Semadeni et al. 2007; Hartmann, Ballesteros-Paredes & Heitsch 2012). Also, in the presence of stellar feedback, at late times the SFRs reach a maximum and begin to decrease again (e.g. V\u00e1zquez-Semadeni et al. 2010; Col\u00edn, V\u00e1zquez-Semadeni & G\u00f3mez 2013). V\u00e1zquez-Semadeni, Gonz\u00e1lez-Samaniego & Col\u00edn (2017) have recently shown that the simulations of Col\u00edn et al. (2013) in fact produce stellar age histograms highly resemblant of the observed ones (Palla & Stahler 1999, 2000; Da Rio et al. 2010), and reproduce observed radial age gradients in clusters (Getman et al. 2014) as well as bottom-heavy stellar initial mass functions (IMFs) in scattered regions of massive star formation (Povich et al. 2016).","Citation Text":["Palla & Stahler","2000"],"Functions Text":["For example, studies of young clusters embedded in moderate-mass MCs (\u223c104M\u2299) (e.g.","have shown that their age histograms contain a large majority of young (1\u20132 Myr) objects, but also a tail of older (up to several Myr) ones suggesting an accelerating star formation activity, sometimes followed by a subsequent decline"],"Functions Label":["Background","Background"],"Citation Start End":[[677,692],[699,703]],"Functions Start End":[[593,676],[725,959]]}
{"Identifier":"2021MNRAS.503.6016K__Martocchia_et_al._2018a_Instance_1","Paragraph":"Observations of splits in the MS and at the eMSTO of young and (at least) moderately massive Magellanic Cloud star clusters have demonstrated the presence of complex stellar populations, suggestive of variations in their stellar rotation properties (Milone et al. 2018a). The effects of stellar rotation on stellar populations have been observed in Galactic open clusters as well (Marino et al. 2018b). Nearly all of these young star clusters are sufficiently massive to host MPs among their MS stars. For example, NGC 1755, a young cluster with a mass of \u223c104\u2009M\u2299 (Milone et al. 2016) or NGC 419, another young massive cluster with an age of 1.72\u2009Gyr and a mass of \u223c104\u2009M\u2299, are comparably massive with respect to other star clusters that show clear evidence of MPs (Martocchia et al. 2018a). Thus, clusters with minimum masses of order 104 M\u2299 tend to show evidence of MPs. However, at present such a mass threshold has yet to be explained theoretically. If we assume that mass is an important driver of MP formation, then NGC 411 should also have shown clear evidence of chemical abundance variations among its RGB stars. However, we have found no such evidence in NGC 411, which challenges the idea that mass may be the only or even the main driver of MP formation. The similarity of the observed RGB width in NGC 411 with that expected for an SSP suggests an absence of significant chemical abundance variations. In fact, for NGC 411, NGC 1718, and NGC 2213 we derive maximum possible helium-abundance variations of \u03b4Y = 0.003 \u00b1 0.001(Y = 0.300), 0.002 \u00b1 0.001(Y = 0.350), and 0.004 \u00b1 0.002(Y = 0.300), respectively. We determined an upper limit to the nitrogen-abundance variation in NGC 411 of \u0394[N\/Fe] = 0.3 dex, although the available data did not allow us to determine useful upper limits for our other sample clusters. Combined with similar results for NGC 419, NGC 1806, and NGC 1846 (Martocchia et al. 2017), our result is indeed inconsistent with mass being the primary driver of MP formation, although it seems likely that a sizeable minimum mass is still required.","Citation Text":["Martocchia et al. 2018a"],"Functions Text":["NGC 419, another young massive cluster with an age of 1.72\u2009Gyr and a mass of \u223c104\u2009M\u2299, are comparably massive with respect to other star clusters that show clear evidence of MPs"],"Functions Label":["Differences"],"Citation Start End":[[766,789]],"Functions Start End":[[588,764]]}
{"Identifier":"2021AandA...656A..44R__Duc_2012_Instance_1","Paragraph":"In this work we aim to explore the presence and properties of LSBGs in the environment of the NGC 1052 group of galaxies. This region is of particular interest in light of the \u201cexotic\u201d properties that recent works claim for some LSBGs found in this region. For instance, van Dokkum et al. (2018a) in [KKS2000] 04 (more commonly known as NGC 1052-DF2) and van Dokkum et al. (2019a) in NGC 1052-DF4 claimed the similarity between the observed baryonic matter and the dynamic matter obtained by the radial velocity dispersion of their GCs and also the stellar component (Emsellem et al. 2019; Danieli et al. 2019). This observational evidence would imply an extreme and unexpected deficit (even lack) of dark matter in these LSBGs. Taking into account that their metallicity content follows the usual stellar mass\u2013metallicity relation for dwarf galaxies (Fensch et al. 2019; Ruiz-Lara et al. 2019), the possibility that they are tidal dwarf galaxies \u2013 galaxies formed by strong interactions with recycled material from massive host galaxies, high in metals, and intrinsically born with a lack of dark matter (see e.g., Duc 2012; Rom\u00e1n et al. 2021) \u2013 is ruled out. The properties of these LSBGs have been the subject of much debate (see e.g., Martin et al. 2018; Ogiya 2018; Famaey et al. 2018; Kroupa et al. 2018; Lewis et al. 2020; Haslbauer et al. 2019; M\u00fcller et al. 2019; Nusser 2019; Montes et al. 2020, 2021). One hypothesis put forward by Trujillo et al. (2019) is that these LSBGs \u201clacking dark matter\u201d could be at a closer distance than initially estimated by van Dokkum et al. (2018a, 2019a). This would not only make the stellar mass smaller (leaving room for a higher M\/L fraction and therefore recovering a certain amount of dark matter, alleviating its strong deficit) but would also resolve the other anomalies that these LSBGs have related to their globular cluster luminosity function (GCLF) (e.g., Shen et al. 2021a), something considered a universal property of any galactic population, including LSBGs (e.g., Jord\u00e1n et al. 2007; Villegas et al. 2010; Rejkuba 2012; Amorisco et al. 2018; Prole et al. 2019b). However, the distance to these galaxies is still the subject of intense debate (van Dokkum et al. 2018b, 2019b; Blakeslee & Cantiello 2018; Monelli & Trujillo 2019; Danieli et al. 2020; Shen et al. 2021b; Zonoozi et al. 2021) and there is currently no broad consensus on these distance values, with the remaining debated range of distances being approximately 12\u201322 Mpc.","Citation Text":["Duc 2012"],"Functions Text":["the possibility that they are tidal dwarf galaxies \u2013 galaxies formed by strong interactions with recycled material from massive host galaxies, high in metals, and intrinsically born with a lack of dark matter (see e.g.,","\u2013 is ruled out."],"Functions Label":["Differences","Differences"],"Citation Start End":[[1116,1124]],"Functions Start End":[[896,1115],[1145,1160]]}
{"Identifier":"2016MNRAS.461.1719C__Fu_et_al._2012_Instance_1","Paragraph":"HATLAS12-00 had already been identified as a candidate gravitationally lensed galaxy as a result of its high submm flux (i.e. F500 > 100 mJy), red Herschel colours and the lack of a bright optical or radio counterpart (see e.g. Negrello et al. 2010 for a discussion of the selection of lens candidates in H-ATLAS and other Herschel surveys). This source was therefore observed spectroscopically in the submm. A CO spectroscopic redshift of 3.26 was first suggested by Z-spec (Bradford et al. 2004) observations, then subsequently confirmed by observations by the CARMA interferometer (Leeuw et al., in preparation) and the Zpectrometer instrument (Harris et al. 2007) on the Greenbank Telescope (Harris et al. 2012; see also Fu et al. 2012). Additional followup observations in the optical, near-IR, submm and other wavelengths were targeted at the lensed z = 3.26 source and the foreground objects responsible for the lensing, resulting in detailed analyses of this lensing system by Fu et al. (2012) and Bussmann et al. (2013). Their conclusions are that the z = 3.26 source HATLAS12-00 is subject to gravitational lensing, with a magnification of 9.6 \u00b1 0.5 in both the submm continuum and CO, and 16.7 \u00b1 0.8 in the K\u2032 band, by two foreground galaxies, one at a spectroscopically determined redshift of 1.22, and another with photometry suggesting that it lies at a similar redshift. The submm photometry of HATLAS12-00 at 890\u2009\u03bcm acquired with the Submillimeter Array (SMA) as part of this programme (Fu et al. 2012; Bussmann et al. 2013) is fully consistent with the 870\u2009\u03bcm and 850\u2009\u03bcm fluxes derived for this source from the LABOCA and SCUBA2 data to be presented here. The spectral energy distribution (SED) of the lensed source, after correcting for the lensing amplification, is well matched by the optically thick SED model for Arp220 from Rangwala et al. (2011), with a lensing-corrected far-IR luminosity of 1.2 \u00b1 0.2 \u00d7 1013\u2009L\u2299, and an implied star formation rate of 1400 \u00b1 300 \u2009M\u2299 yr\u22121. In many ways the unlensed properties of this object match those of the broader population of bright submm selected galaxies first discovered by the SCUBA submm imager (see e.g. Chapman et al. 2005; Clements et al. 2008; Micha\u0142owski, Hjorth & Watson 2010). The unlensed 870\u2009\u03bcm flux of this object would be \u223c7.7 mJy.","Citation Text":["Fu et al. 2012"],"Functions Text":["A CO spectroscopic redshift of 3.26 was first suggested by Z-spec","observations, then subsequently confirmed by observations by","the Zpectrometer instrument","on the Greenbank Telescope"],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[725,739]],"Functions Start End":[[409,474],[498,558],[619,646],[668,694]]}
{"Identifier":"2019MNRAS.490.3875L__Onaka_et_al._1996_Instance_1","Paragraph":"Similar to Chen et al. (2017), we also follow Li & Draine (2001a) to derive upper limits on the abundance of the graphene C24 in the Galactic cirrus and in the diffuse ISM towards l = 44\u00b020\u2032, b = \u22120\u00b020\u2032 based on comparison of the observed IR emission with the calculated emission spectrum of C24. For the Galactic cirrus, the average emission per H has been measured by COBE\/DIRBE (Arendt et al. 1998), COBE\/FIRAS (Finkbeiner, Davis & Schlegel 1999), and Planck (Planck Collaboration XVII 2014). The diffuse ISM towards l = 44\u00b020\u2032, b = \u22120\u00b020\u2032 has been observed by COBE\/DIRBE (Hauser et al. 1998). The Mid-Infrared Spectrograph (MIRS) aboard the Infrared Telescope in Space (IRTS) has also obtained the 4.7\u201311.7$\\, {\\rm \\mu m}$ spectrum for the diffuse ISM towards l = 44\u00b020\u2032, b = \u22120\u00b020\u2032 (Onaka et al. 1996). As shown in Fig. 5(a), even locking up 20\u2009$\\, {\\rm ppm}$ of C\/H \u2013 the upper limit of graphene derived from the interstellar extinction \u2013 all in the specific graphene species C24, the 6.6, 9.8, and 20\u2009$\\, {\\rm \\mu m}$ emission features of C24 would still be hidden by the PAH features at 6.2, 7.7, 8.6, and 11.3\u2009$\\, {\\rm\\, \\mu m}$ and would remain undetected by Spitzer or by the Short Wavelength Spectrometer (SWS) aboard the Infrared Space Observatory (ISO). This is also true for the diffuse ISM towards l = 44\u00b020\u2032, b = \u22120\u00b020\u2032. As illustrated in Fig. 5(b), as much as $\\sim \\, 20\\, {\\rm ppm}$ of C\/H could also be tied up in the C24 graphene while the characteristic 6.6, 9.8, and 20\u2009$\\, {\\rm \\mu m}$ emission features of C24 are still not strong enough to be detected by IRTS. Therefore, for both the Galactic cirrus and the line of sight towards l = 44\u00b020\u2032, b = \u22120\u00b020\u2032, a upper limit of C\/H\u2009\u2272\u200920\u2009$\\, {\\rm ppm}$ is imposed by the COBE\/DIRBE photometric data and the IRTS spectrum. Nevertheless, the actual abundance of C24 could be much lower than 20\u2009$\\, {\\rm ppm}$ since, if graphene is indeed present in the ISM, it could span a wide range of sizes and charging states.","Citation Text":["Onaka et al. 1996"],"Functions Text":["The Mid-Infrared Spectrograph (MIRS) aboard the Infrared Telescope in Space (IRTS) has also obtained the 4.7\u201311.7$\\, {\\rm \\mu m}$ spectrum for the diffuse ISM towards l = 44\u00b020\u2032, b = \u22120\u00b020\u2032"],"Functions Label":["Uses"],"Citation Start End":[[788,805]],"Functions Start End":[[597,786]]}
{"Identifier":"2020ApJ...901...50E__Dame_et_al._2001_Instance_1","Paragraph":"Molecular hydrogen H2 is conventionally believed to make up 17% of the interstellar medium (ISM) by mass in the Galaxy (e.g., Draine 2011, Table 1.2, not including He), but, unlike atomic hydrogen H i (at 60% by mass), H2 is not generally detectable in the cold, diffuse conditions that predominate the volume of the ISM. In order to determine the mass and structure of interstellar H2, emission from other \u201ctrace\u201d molecules must be observed instead. The primary tracer for H2 in current use is the J = (1\u22120) ground-state rotational transition of 12CO at a wavelength of 3 mm. While this tracer has become widely adopted over the years (see, e.g., Heyer et al. 1998; Dame et al. 2001) and has been often defended as a reliable large-scale tracer for molecules in the ISM, it also remains the primary such tracer in common use. There are several reasons to be concerned about this. First, the critical volume density of the 12CO (1\u20130) transition at low optical depths is \n\n\n\n\n\n \n\n\n\n\n\n; this means that gas at lower volume density will be less excited by collisions, thus emitting fewer photons, and may even effectively disappear from sensitivity-limited surveys. Second, in sufficient quantities to avoid the \u201cexcitation\u201d problem, the 12CO signal is optically thick, so that direct calculation of column densities from line profile integrals is not possible. Instead, indirect methods of inferring the quantity of H2 must be used, such as the \u201cX-factor\u201d method (see, e.g., Bolatto et al. 2013, for a review). However, recent evidence from gamma rays and IR surveys has suggested that there are significant regions of the ISM apparently containing molecular gas that is not adequately traced by CO emission (see, e.g., Grenier et al. 2005). The \u201cdark gas\u201d appears not to be hidden optically thick H i (see Murray et al. 2018), and as such is most likely molecular. While several other molecular gas tracers have been studied (see, e.g., Jacob et al. 2019, on CH), these tracers may not be ideal for detecting low-density, diffuse portions of the large-scale ISM in emission owing to their high critical densities. A reliable alternative molecular tracer for diffuse gas is needed to confirm the identity of this \u201cCO-dark\u201d gas, as well as to provide information on its nature, kinematics, column density, and structure.","Citation Text":["Dame et al. 2001"],"Functions Text":["The primary tracer for H2 in current use is the J = (1\u22120) ground-state rotational transition of 12CO at a wavelength of 3 mm. While this tracer has become widely adopted over the years (see, e.g.,","and has been often defended as a reliable large-scale tracer for molecules in the ISM, it also remains the primary such tracer in common use. There are several reasons to be concerned about this."],"Functions Label":["Background","Compare\/Contrast"],"Citation Start End":[[667,683]],"Functions Start End":[[451,647],[685,880]]}
{"Identifier":"2020MNRAS.495.2949F__Stauffer_et_al._2016_Instance_1","Paragraph":"Stellar rotation is known to be related to the activity level; e.g. faster rotators tend to have larger star-spot coverages (e.g. Paper I), and stronger chromospheric emissions (e.g. Stauffer et al. 1997; Douglas et al. 2014; Newton et al. 2017; Paper II) and X-ray emissions (e.g. Pizzolato et al. 2003; Mamajek & Hillenbrand 2008); thus, a correlation between rotation and variation level should also be expected. In fact, Fig. 7 shows that there exists a large scatter in variation amplitude even for stars in the same cluster. It is reasonable to connect the scatter to the stellar rotation rate diversity, considering the fact that there exist noticeable differences in rotation periods between stars of the same temperature in young open clusters (e.g. Stauffer et al. 2016). To understand further on this scenario, we provided variation amplitudes as a function of temperature in Fig. 8, wherein shown are stars with known rotation periods. For comparison purpose, we also displayed the corresponding rotation periods in the upper panels. The periods were collected from the literature (Pleiades: Hartman et al. 2010; Covey et al. 2016; Rebull et al. 2016a; Praesepe: Ag\u00fceros et al. 2011; Delorme et al. 2011; Kov\u00e1cs et al. 2014; Rebull et al. 2017; Hyades: Hartman et al. 2011; Armstrong et al. 2015; Douglas et al. 2016); for a few stars in Praesepe, such as EPIC 211852399 and EPIC 211875458, the periods were obtained by us based on K2 data (see Appendix A). The colour contrast in Fig. 8 denotes the Rossby number, Ro, the normalized rotation period by the convective turnover time (the convective turnover time was estimated by using the correlation reported by Wright et al. 2011; see Paper I\/II for details). It is known that the open cluster members mainly locate in two rotation sequences in the rotation\u2013colour diagram (Barnes 2003), i.e. I (interface) and C (convective) sequence, corresponding to slow and fast rotators, respectively, as shown in the top panels in Fig. 8. The figure shows a tendency that GK-type C sequence stars in Pleiades have larger variation amplitude of EWH\u2009\u03b1 compared to I sequence stars of the same temperature. We found no evident difference between different rotators among M-type Pleiades stars, partially because most of them are very fast rotators that reside in activity saturation regime (when Ro \u2272 0.1; e.g. Paper II). For Praesepe\/Hyades sample, almost all GK-type stars rotate slowly (e.g. Ro \u223c 0.5), located in the I sequence (those stars with blue colours); there is no clear trend among these stars. Among M-type stars, some stars are fast rotators in C sequence (red or green colours), while others are slow rotators (in I sequence), as shown in the upper right panel of Fig. 8. Thus, there is a tendency that faster rotators have larger variation amplitude of EWH\u2009\u03b1. Unlike EWH\u2009\u03b1, TiO2n shows no clear difference among all sample stars with different rotation periods, probably due to its small intrinsic variation amplitude, which could be dominated by observational noise.","Citation Text":["Stauffer et al. 2016"],"Functions Text":["It is reasonable to connect the scatter to the stellar rotation rate diversity, considering the fact that there exist noticeable differences in rotation periods between stars of the same temperature in young open clusters (e.g."],"Functions Label":["Uses"],"Citation Start End":[[759,779]],"Functions Start End":[[531,758]]}
{"Identifier":"2022AandA...666L...5G__Ramos_et_al._2014_Instance_1","Paragraph":"More recently, Garc\u00eda-Bernete et al. (2022) found that the PAH molecules responsible for the 11.3 \u03bcm PAH emission band are more resilient in the hard environments often present in AGN. In particular, the authors found larger 11.3\/7.7 \u03bcm and 11.3\/6.2 \u03bcm PAH ratios in AGN-dominated systems compared to SF galaxies, indicating a larger fraction of neutral PAH molecules (as noted by Smith et al. 2007 using a sample of relatively weak AGN). However, these studies were limited by the spatial resolution (\u223c4\u2033) and the low spectral resolution (R\u2004\u223c\u200460\u2013130) of Spitzer\/InfraRed Spectrograph (IRS). Previous sub-arcsecond angular resolution N-band (\u223c8\u201313 \u03bcm) ground-based spectroscopic studies investigated the 11.3 \u03bcm PAH feature in the nuclear and circumnuclear regions of AGN (e.g., H\u00f6nig et al. 2010; Gonz\u00e1lez-Mart\u00edn et al. 2013; Alonso-Herrero et al. 2014, 2016; Ramos et al. 2014; Esquej et al. 2014; Garc\u00eda-Bernete et al. 2015; Jensen et al. 2017; Esparza-Arredondo et al. 2018). However, these works were unable to provide definitive details regarding the effect of the AGN on the PAH molecules due to limited wavelength coverage and sensitivity. The changes in the PAH properties due to the presence of the AGN might be more prominent in their innermost regions of galaxies. Therefore, the unprecedented combination of high angular and spectral resolution (R\u2004\u223c\u20041500\u2005\u2212\u20053500) in the entire mid-IR range (4.9\u201328.1 \u03bcm) afforded by the James Webb Space Telescope (JWST)\/Mid-Infrared Instrument (MIRI; Rieke et al. 2015; Wells et al. 2015; Wright et al. 2015) is key to investigating PAH properties. In this Letter we report on the first investigation of PAH emission in the nuclear regions of three luminous Seyfert (Sy) galaxies and compare them with emission from SF regions using JWST\/MIRI Medium Resolution Spectrograph (MRS) data. This enables us, for the first time, to characterise the PAH properties of local luminous Sy galaxies (log (Lbol)> 44.46 erg s\u22121)1 at sub-arcsecond scales (\u223c0.45\u2033, \u223c142\u2013245 pc).","Citation Text":["Ramos et al. 2014"],"Functions Text":["Previous sub-arcsecond angular resolution N-band (\u223c8\u201313 \u03bcm) ground-based spectroscopic studies investigated the 11.3 \u03bcm PAH feature in the nuclear and circumnuclear regions of AGN (e.g.,","However, these works were unable to provide definitive details regarding the effect of the AGN on the PAH molecules due to limited wavelength coverage and sensitivity."],"Functions Label":["Background","Motivation"],"Citation Start End":[[861,878]],"Functions Start End":[[592,778],[980,1147]]}
{"Identifier":"2020AandA...633A.147B__Xu_et_al._(2013)_Instance_1","Paragraph":"We used the position and distance of the MCs in the FQS catalogue to describe the structure of the Milky Way in this poorly known portion of the third quadrant. In Fig. 16 we show the distribution of the clouds projected onto the Milky Way plane. At any Galactic longitude the clouds are grouped in three well-separated structures that trace the position of the spiral arms. The dispersion of the positions in each structure is given not only by the width of the arms but also by the uncertainty of the distance determination. Indeed, kinematic distances are affected by significant uncertainties due not only to the assumed Galactic rotation model but also to the unknown possible peculiar motions of the clouds themselves. In Fig. 16 we compare the positions of the clouds derived from FQS data with the expected locations of the spiral arms, adopting a log-periodic spiral form of the Galaxy (Reid et al. 2014). The pitch angle and the distance at a reference position of the three arms are taken from Xu et al. (2013) for the Local arm, Choi et al. (2014) for the Perseus arm, and Hachisuka et al. (2015) for the Outer arm. These authors used the positions of water masers derived from parallax measurements as part of the Bar and Spiral Structure Legacy (BeSSeL) survey to fit the parameters of the spiral arms. We found good agreement between the position of the FQS clouds and the modelled position of the Local and the Perseus arms, while the modelled Outer arm lies further than the observed clouds. We note that while the parameters of the Local and Perseus arms are derived by fitting about 20 water masers located in both the second and third quadrants, covering also the Galactic longitude of the FQS survey, the parameters of the Outer arm are derived from only 5 water masers, 4 of which are located in the second quadrant. Therefore it is not surprising that they might not be fully appropriate to describe the location of the Outer arm in the longitude range of the FQS survey of 220\u00b0\u2005 l\u2004 \u2004240\u00b0. In fact, an alternative model of the Outer arm from Hou & Han (2014), also shown in Fig. 16, seems in better agreement with the FQS results.","Citation Text":["Xu et al. (2013)"],"Functions Text":["The pitch angle and the distance at a reference position of the three arms are taken from","for the Local arm","These authors used the positions of water masers derived from parallax measurements as part of the Bar and Spiral Structure Legacy (BeSSeL) survey to fit the parameters of the spiral arms.","We found good agreement between the position of the FQS clouds and the modelled position of the Local and the Perseus arms,"],"Functions Label":["Uses","Uses","Background","Similarities"],"Citation Start End":[[1005,1021]],"Functions Start End":[[915,1004],[1022,1039],[1128,1316],[1317,1440]]}
{"Identifier":"2015AandA...584A..75V__Essen_et_al._(2014)_Instance_1","Paragraph":"The data presented here comprise quasi-simultaneous observations during secondary eclipse of WASP-33 b around the V and Y bands. The predicted planet-star flux ratio in the V-band is 0.2 ppt, four times lower than the accuracy of our measurements. Therefore, we can neglect the planet imprint and use this band to measure the stellar pulsations, and most specifically to tune their current phases (see phase shifts in von Essen et al. 2014). Particularly, our model for the light contribution of the stellar pulsations consists of eight sinusoidal pulsation frequencies with corresponding amplitudes and phases. Hence, to reduce the number of 24 free parameters and the values they can take, we use prior knowledge about the pulsation spectrum of the star that was acquired during von Essen et al. (2014). As the frequency resolution is 1\/\u0394T (Kurtz 1983), 3.5 h of data are not sufficient to determine the pulsations frequencies. Therefore, during our fitting procedure we use the frequencies determined in von Essen et al. (2014) as starting values plus their derived errors as Gaussian priors. As pointed out in von Essen et al. (2014), we found clear evidences of pulsation phase variability with a maximum change of 2 \u00d7 10-3 c\/d. In other words, as an example after one year time a phase-constant model would appear to have the correct shape with respect to the pulsation pattern of the star, but shifted several minutes in time. To account for this, the eight phases were considered as fitting parameters. The von Essen et al. (2014) photometric follow-up started in August, 2010, and ended in October, 2012, coinciding with these LBT data. We then used the phases determined in von Essen et al. (2014) during our last observing season as starting values, and we restricted them to the limiting cases derived in Sect. 3.5 of von Essen et al. (2014), rather than allowing them to take arbitrary values. The pulsation amplitudes in \u03b4 Scuti stars are expected to be wavelength-dependent (see e.g. Daszy\u0144ska-Daszkiewicz 2008). Our follow-up campaign and these data were acquired in the blue wavelength range. Therefore the amplitudes estimated in von Essen et al. (2014), listed in Table 1, are used in this work as fixed values. This approach would be incorrect if the pulsation amplitudes would be significantly variable (see e.g., Breger et al. 2005). Nonetheless, the short time span of LBT data, and the achieved photometric precision compared to the intrinsically low values of WASP-33\u2019s amplitudes, make the detection of any amplitude variability impossible. ","Citation Text":["von Essen et al. 2014"],"Functions Text":["Therefore, we can neglect the planet imprint and use this band to measure the stellar pulsations, and most specifically to tune their current phases (see phase shifts in"],"Functions Label":["Uses"],"Citation Start End":[[418,439]],"Functions Start End":[[248,417]]}
{"Identifier":"2021AandA...649A.126T__Luck_(2018b)_Instance_2","Paragraph":"Studies of the radial n-capture-to-iron abundance gradients are very scarce so far. We can only search for a broad agreement of our results with several studies of abundance gradients with galactocentric distances (Rgc). da Silva et al. (2016) studied n-capture elements across the Galactic thin disc based on Cepheid variables. Because the Cepheids are young stars, their Rgc may be rather close to their birthplaces and Rmean. da Silva et al. (2016) supplemented their sample of 111 Cepheids with 324 more stars from other studies and found that the [Y\/Fe] distribution is flat throughout the entire disc. In our study, we confirm this finding not only based on the whole thin-disc sample of stars and on a subsample of younger \u22644 Gyr stars, but also add another light s-process dominated element strontium. Like in our study, da Silva et al. (2016) also obtained positive [El\/Fe] radial gradients for La, Ce, Nd, and Eu. The slopes are rather similar. For [Eu\/Fe], they differ just by 0.002 dex kpc\u22121. More recently, Luck (2018b) also investigated the gradients of n-capture element abundance-to-iron ratios with respect to Rgc for a sample of 435 Cepheids. It is interesting to note that the [Ba\/Fe] versus Rgc slope according to this Cepheid sample is also negative, as in our study. [Ba\/Fe] is the only n-capture element-to-iron ratio with a negative radial gradient in our sample of stars and in Luck (2018b). Overbeek et al. (2016) investigated trends of Pr, Nd, and Eu to Fe abundance ratios with respect to Rgc using 23 open clusters. As in our study, they found that these elements have positive linear trends with galactocentric radius (the linear regression slopes are of about +0.04 dex kpc\u22121). They also suggested that the [El\/Fe] relation of Pr and Nd, but not Eu, with the galactocentric radius may not be linear because the [El\/Fe] of these elements appears to be enhanced around 10 kpc and drop around 12 kpc. Because only a small number of stars lie at these large radial distances, we cannot address this question. For the thick-disc stars, the radial abundance-to-iron slopes are negligible, as was found for \u03b1-process elements by Li et al. (2018), even though the production sites of \u03b1-elements and s-processes dominated elements are quite different.","Citation Text":["Luck (2018b)"],"Functions Text":["It is interesting to note that the [Ba\/Fe] versus Rgc slope according to this Cepheid sample is also negative, as in our study. [Ba\/Fe] is the only n-capture element-to-iron ratio with a negative radial gradient in our sample of stars and in"],"Functions Label":["Similarities"],"Citation Start End":[[1403,1415]],"Functions Start End":[[1161,1402]]}
{"Identifier":"2018MNRAS.474.2277D__Oliveira,_Dottori_&_Bica_1998_Instance_1","Paragraph":"There are three possible explanations for the origin of these systems: (1) they formed from the fragmentation of the same molecular cloud (Elmegreen & Elmegreen 1983), (2) they were generated in distinct molecular clouds and then became bound systems after a close encounter leading to a tidal capture (Vallenari, Bettoni & Chiosi 1998; Leon, Bergond & Vallenari 1999), or (3) they are the result of division of a single star-forming region (Goodwin & Whitworth 2004; Arnold et al. 2017). Their subsequent evolution may also have different outcomes. Dynamical models and N-body simulations (see, e.g., Barnes & Hut 1986; de Oliveira, Dottori & Bica 1998, and references therein) have shown that, depending on the initial conditions, a bound pair of clusters may either become unbound, because of significant mass-loss in the early phases of stellar evolution, or merge into a single and more massive cluster on a short time-scale (\u224860\u2009Myr) due to loss of angular momentum from escaping stars (see Portegies Zwart & Rusli 2007). The final product of a merger may be characterized by a variable degree of kinematic and morphologic complexity, mostly depending on the values of the impact parameter of the pre-merger binary system (de Oliveira, Bica & Dottori 2000; Priyatikanto et al. 2016). In some cases, the stellar system resulting from the merger event may show significant internal rotation (in fact, for many years this has been the preferred dynamical route to form rotating star clusters; see Sugimoto & Makino 1989; Makino, Akiyama & Sugimoto 1991; Okumura, Ebisuzaki & Makino 1991; de Oliveira, Dottori & Bica 1998). Merger of cluster pairs has been sometimes invoked to interpret the properties of particularly massive and dynamically complex clusters (e.g. see the study of \u03c9 Centauri by Lee et al. 1999, G1 by Baumgardt et al. 2003 and NGC\u2009 2419 by Br\u00fcns & Kroupa 2011), and, more in general, as an avenue to form clusters with multiple populations with different chemical abundances both in terms of iron and light elements (e.g. van den Bergh 1996; Catelan 1997; Amaro-Seoane et al. 2013; Gavagnin, Mapelli & Lake 2016; Hong et al. 2017).","Citation Text":["de Oliveira, Dottori & Bica 1998"],"Functions Text":["Dynamical models and N-body simulations (see, e.g.,","and references therein) have shown that, depending on the initial conditions, a bound pair of clusters may either become unbound, because of significant mass-loss in the early phases of stellar evolution, or merge into a single and more massive cluster on a short time-scale (\u224860\u2009Myr) due to loss of angular momentum from escaping stars"],"Functions Label":["Background","Background"],"Citation Start End":[[621,653]],"Functions Start End":[[550,601],[655,991]]}
{"Identifier":"2022AandA...667A..82D__Lyne_et_al._(1985)_Instance_1","Paragraph":"Much of the progress in PPS has come from thorough Monte Carlo simulations that generate pulsars and test whether they fulfil the criteria for detection according to geometrical factors and sensitivity issues. It is then possible to develop and optimize a model for the underlying pulsar population, which informs us about the important intrinsic neutron star parameters and distribution, enabling predictions for future surveys. Usually, PPS studies follow two simple approaches. The first one is to take a \u2018snapshot\u2019 of the Galaxy as it appears today, where no assumptions are made regarding the prior evolution of the pulsar population. Instead, this population is generated assuming various distribution functions (typically spatial distribution, spin period P and, \n\n\n\nE\n\u02d9\n\n\n$ \\dot{E} $\n\n\n- or \u1e56), which are optimized to match the observations. Inspired by earlier studies from Taylor & Manchester (1977) and Lyne et al. (1985), Lorimer et al. (2006) applied the snapshot approach to the canonical3 pulsar population to determine best-fitting probability density functions in Galactocentric radius (R), luminosity (L), height with respect to the Galactic plane (z), and the period P for the currently observed population of pulsars. Alternatively, one may consider \u2018evolution\u2019 strategies where the pulsars are evolved from birth up to the present era, starting from an initial spatial distribution, and an initial period and magnetic field distribution. A fine example of the latter genre is the comprehensive study of Faucher-Gigu\u00e8re & Kaspi (2006), which quite successfully reproduced the properties of the main part of the radio pulsar population using a model in which the luminosity has a power-law dependence on P and \u1e56. PPS studies can be also extended to the population of neutron stars observed in other bands, such as X-rays (see for instance Popov et al. 2010) and \u03b3-rays. Indeed, with the broad increase in \u03b3-ray pulsar numbers, a statistical treatment of the \u03b3-ray population in combination with deep radio surveys of the Galactic plane is now feasible. Early works on radio-loud gamma-ray pulsar populations carried out before the Fermi era include Gonthier et al. (2002), (2004), (2007a,b). With the advent of Fermi\/LAT, new studies emerged (see for instance Gonthier et al. 2018; Ravi et al. 2010; Takata et al. 2011; Watters & Romani 2011; Pierbattista et al. 2012), trying to constrain the geometry and the location of the gamma-ray emission sites. Watters & Romani (2011) showed that an initial spin period of P0\u2004=\u200450 ms and a birth rate of one neutron star per 59 yr were required to reproduce the observed \u03b3-ray population. They made the prediction that after ten years of operations, Fermi should detect \u223c120 young \u03b3-ray pulsars, of which about one half would be radio-quiet. Gonthier et al. (2004) included an exponentially decaying magnetic field with a 2.8 Myr timescale and displaying a satisfactory agreement with the P\u2212\u1e56 distribution at that time. Later, more accurate magnetic field decay models were elaborated, especially for magnetars, as presented in Vigan\u00f2 et al. (2013).","Citation Text":["Lyne et al. (1985)"],"Functions Text":["Inspired by earlier studies from Taylor & Manchester (1977) and","Lorimer et al. (2006) applied the snapshot approach to the canonical3 pulsar population to determine best-fitting probability density functions in Galactocentric radius (R), luminosity (L), height with respect to the Galactic plane (z), and the period P for the currently observed population of pulsars."],"Functions Label":["Background","Background"],"Citation Start End":[[914,932]],"Functions Start End":[[850,913],[934,1237]]}
{"Identifier":"2020MNRAS.495L..27H__Cao_et_al._2018_Instance_1","Paragraph":"Independent measurements of distances and redshifts, z, allow astronomers to constrain cosmological models, since they both define the distance\u2013redshift relation. When it was determined that type Ia supernovae (SNIa) could be standardized and therefore used to measure distances, this led to the discovery of the accelerated expansion of our Universe (Riess et al. 1998; Perlmutter et al. 1999). The combination of SNIa (Betoule et al. 2014), baryonic acoustic oscillations (Eisenstein et al. 2005; Alam et al. 2017), and the cosmic microwave background (CMB; Komatsu et al. 2011; Planck Collaboration et al. 2018) led to the emergence of the concordance \u039b cold dark matter (\u039bCDM) model, in which the energy density is dominated by dark energy as a cosmological constant \u039b. In a flat Friedmann\u2013Lema\u00eetre\u2013Robertson\u2013Walker universe, the comoving distance is defined as \n(1)$$\\begin{eqnarray*}\r\nD(z) = \\int _0^z \\frac{c\\mathrm{d}z}{H(z)},\r\n\\end{eqnarray*}$$ where, in the \u039bCDM model, \n(2)$$\\begin{eqnarray*}\r\nH(z) = H_0 \\sqrt{\\Omega _\\mathrm{m} (1+z)^3 + 1-\\Omega _\\mathrm{m}}\r\n\\end{eqnarray*}$$is the Hubble parameter, H0 is the Hubble\u2013Lema\u00eetre constant, and \u03a9m is the matter energy density at the current epoch. The luminosity and angular diameter distances are defined as \n(3)$$\\begin{eqnarray*}\r\nD_\\mathrm{L}(z) = (1+z) D(z)\r\n\\end{eqnarray*}$$(4)$$\\begin{eqnarray*}\r\n{\\mathrm{ and}} \\,D_\\mathrm{A}(z) = \\frac{R}{\\theta }= \\frac{D(z)}{(1+z)}.\r\n\\end{eqnarray*}$$Type Ia supernovae can only be used up to redshift of around 2 (Jones et al. 2013), and there are tensions between direct local measurements of the Hubble\u2013Lema\u00eetre constant and model-dependent estimates using CMB observations (Planck Collaboration et al. 2018). Therefore, independent distance measurements to extragalactic objects are desired. We should emphasize here, that model-independent distance indicators can have various important applications in physical cosmology. In particular, to test different aspects of cosmological models and theories of gravity. For instance, we can use these model-independent measurements to test the FLRW metric (Clarkson, Bassett & Lu 2008; Wiltshire 2009; Shafieloo & Clarkson 2010; L\u2019Huillier & Shafieloo 2017; Shafieloo, L\u2019Huillier & Starobinsky 2018; Cao et al. 2019a; Qi et al. 2019a), to test general relativity and some modified gravity models (Cao et al. 2012; Shafieloo, Kim & Linder 2013a; Cao et al. 2015, 2017b; Qi et al. 2017; L\u2019Huillier, Shafieloo & Kim 2018; Shafieloo et al. 2018; Xu et al. 2018; Chen, Sesana & Conselice 2019; L\u2019Huillier et al. 2020), to test natural constants such as the speed of light (Cao et al. 2018), to test cosmic duality relationships, and to measure cosmic curvature (Shafieloo et al. 2013b; Cao et al. 2019b; Qi et al. 2019a, b; Zheng et al. 2020). Using a combination of such model-independent distance indicators can also be used to measure some key cosmological parameters such as the Hubble Constant (Liao et al. 2015; Suyu et al. 2017; Jee et al. 2019; Liao et al. 2019, 2020; Wong et al. 2019). Amongst the most energetic objects in our Universe are active Galactic nuclei (AGNs). AGNs are the nuclei of massive galaxies that sometimes produce relativistic jets of material launched from near a central supermassive black hole (SMBH). When these jets are not aligned close to our line of sight, AGNs are observed as radio galaxies, whereas if the jet is aligned to within a small angle to our line of sight, they are observed as blazars (Urry & Padovani 1995). Blazars and quasars are amongst the most consistently bright objects in our Universe and can be observed at redshifts as high as \u223c7 (Mortlock et al. 2011). Attempts have been made to measure distances to AGNs in various ways, with some claiming deviations from the expected cosmology at high redshifts (Risaliti & Lusso 2017, 2019; Turner & Shabala 2019). Very long baseline interferometry (VLBI) has also been used to attempt to measure cosmological distances. The approach pioneered by Gurvits, Kellermann & Frey (1999) attempted to measure cosmological parameters by assuming that AGN could be used as a standardizable rod. Vishwakarma (2001) used this data set and compared it with supernovae data, and found it was not possible to differentiate different cosmological models with the VLBI data of Gurvits et al. (1999). Cao et al. (2015) revisited this technique and investigated the evolution of the standard rod by assuming a Planck cosmology. Cao et al. (2017a, c) then introduced a cosmology independent method for calibrating the standard rod and was able to provide reasonable constraints on cosmological parameters. Our approach differs from this by using the speed of light to normalize the rod. This approach was first attempted by Wiik & Valtaoja (2001), which found that that the apparent angular sizes of AGNs maximized at z \u223c 2. In this paper, we demonstrate the method on the famous nearby radio galaxy 3C\u200984 (NGC\u20091275) and discuss possible systematic errors. The source is known to exhibit extremely high energy emission despite not exhibiting strong relativistic effects (Jorstad et al. 2017; Liodakis et al. 2018), and has multiple independent measures of distance (Theureau et al. 2007; Hicken et al. 2009), thus making it an ideal source to test our methodology.","Citation Text":["Cao et al. 2018"],"Functions Text":["For instance, we can use these model-independent measurements","to test natural constants such as the speed of light"],"Functions Label":["Background","Background"],"Citation Start End":[[2624,2639]],"Functions Start End":[[2026,2087],[2570,2622]]}
{"Identifier":"2021ApJ...919...30D__Staguhn_et_al._2014_Instance_2","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"],"Functions Text":["Selecting SMGs from observations at longer wavelengths is thought to favor galaxies at higher redshifts (e.g.,","although it is difficult to compare the redshift distributions in an unbiased way","and account for intrinsic variations of galaxy far-IR spectral energy distributions (SEDs)."],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[802,821]],"Functions Start End":[[650,760],[869,950],[1001,1092]]}
{"Identifier":"2022ApJ...931...70B__Gabrielse_et_al._2012_Instance_2","Paragraph":"RFs can propagate from the magnetotail to Earth over a long distance more than 10 R\nE together with BBFs behind them (Runov et al. 2009; Cao et al. 2010). Studies have suggested that RFs are crucial regions for particle acceleration, pitch-angle evolution, wave\u2013particle interactions, and electromagnetic energy conversion during their Earthward propagation. For instance, rapid increases in energy fluxes of electrons and ions from tens to hundreds of keV are a typical feature of RF events (Khotyaintsev et al. 2011; Liu et al. 2013, 2018c, 2021a, 2022b; Zhou et al. 2018; Liu & Fu 2019; Gabrielse et al. 2021), pitch-angle distribution of suprathermal electrons can evolve dramatically around RFs (Runov et al. 2013; Liu et al. 2020), strong particle and wave activity can occur in the vicinity of RFs (Ono et al. 2009; Zhou et al. 2009, 2014; Fu et al. 2014; Breuillard et al. 2016; Greco et al. 2017; Yang et al. 2017), and RFs are associated with energy conversion from electromagnetic fields to particles (Sitnov et al. 2009; Huang et al. 2015; Khotyaintsev et al. 2017; Liu et al. 2018a, 2022a). The energetic plasma in the vicinity of RFs plays a key role in connecting the magnetotail with the inner magnetosphere because they carry a large amount of energy and can be injected into the inner magnetosphere to affect the ring current and radiation belt (Gabrielse et al. 2012; Duan et al. 2014; Turner et al. 2014). Possible mechanisms responsible for the energization of particles around RFs have been widely investigated based on both spacecraft observations and numerical simulations during the past decade. The strong convection electric field induced by the strong magnetic field gradient of RFs provides significant adiabatic acceleration of the ambient particles (Birn et al. 2004, 2013, 2015; Gabrielse et al. 2012, 2014, 2016; Ganushkina et al. 2013; Liu et al. 2016; Turner et al. 2016). Nonadiabatic effects, caused by particle reflection ahead of the RFs (Zhou et al. 2018), resonance with RFs (Ukhorskiy et al. 2013, 2017), and scattering by wave emissions (Zhou et al. 2009; Greco et al. 2017), are also significant for particle energization. These above studies usually assumed that the RF surface has a planar boundary at a typical thickness comparable to the ion gyroradius and below (Nakamura et al. 2002; Sergeev et al. 2009; Zhou et al. 2009; Schmid et al. 2011; Liu et al. 2013; Vapirev et al. 2013). Divin et al. (2015b) revealed that the RF surface is unstable to instabilities ranging from electron scales to ion scales. Simulation studies found that RFs can be unstable to interchange instability and that finger-like structures on ion\u2013electron hybrid scales can develop at the RF (Vapirev et al. 2013). Such finger-like structures are found to play a role in modulating the electron acceleration process (Wu et al. 2018). Bai et al. (2022) also reported significant ion trapping acceleration at the RF with ion-scale ripples. Unlike these surface structures with ion or ion\u2013electron hybrid scales, Liu et al. (2018b) recently reported that the RF layer has electron-scale density gradients, currents, and electric fields, based on the MMS mission, which consists of four spacecraft separated by 30 km. Such electron-scale ripple structure can be generated by lower hybrid drift instability (Divin et al. 2015b; Pan et al. 2018). Liu et al. (2021c) presented a detailed investigation of energy flux densities at two RFs with\/without the electron-scale surface ripples and indicated that surface ripples may play an important role in the particle dynamics. But how such electron-scale RF structure impacts the electron energization and transport still remains unknown. In this paper, with the aid of observation-based test-particle simulation, we aim to investigate in detail the effect of the front surface ripples on the local electron dynamics.","Citation Text":["Gabrielse et al. 2012"],"Functions Text":["Possible mechanisms responsible for the energization of particles around RFs have been widely investigated based on both spacecraft observations and numerical simulations during the past decade. The strong convection electric field induced by the strong magnetic field gradient of RFs provides significant adiabatic acceleration of the ambient particles"],"Functions Label":["Background"],"Citation Start End":[[1811,1832]],"Functions Start End":[[1426,1779]]}
{"Identifier":"2017ApJ...835...25E__Rutten_1984_Instance_1","Paragraph":"We compare our results with a new reduction of observations from the Lowell Observatory SSS, which is running a long-term stellar activity survey complementary to the MWO HK Project. The SSS observes solar and stellar light with the same spectrograph, with the solar telescope consisting of an exposed optical fiber that observes the Sun as an unresolved source (Hall & Lockwood 1995; Hall et al. 2007). The basic measurement of SSS is the integrated flux in 1 \u212b bandpasses centered on the Ca ii H & K cores from continuum-normalized spectra, \u03d5HK, which can then be transformed to the S-index using a combination of empirical relationships derived from stellar observations:\n7\n\n\n\n\n\nwhere \n\n\n\n\n\n is the continuum flux scale for the Ca ii H & K wavelength region, which converts \u03d5HK to physical flux (erg cm\u22122 s\u22121). \n\n\n\n\n\n is a function of Str\u00f6mgren \n\n\n\n\n\n and is taken from Hall (1996). \n\n\n\n\n\n (simply K in other works) is the conversion factor from the MWO HKP-2 H & K flux (numerator of Equation (1)) to physical flux (Rutten 1984). Ccf is a factor that removes the color term from S and is a function of Johnson \n\n\n\n\n\n (Rutten 1984). Finally, Teff is the effective temperature. See Hall et al. (2007) and Hall & Lockwood (1995) for details on the extensive work leading to this formulation. What is important to realize about this method of obtaining S is that it requires three measurements of solar properties, \n\n\n\n\n\n, \n\n\n\n\n\n, and \n\n\n\n\n\n, along with the determination of one constant, \n\n\n\n\n\n. The solar properties are taken from best estimates in the literature, which vary widely depending on the source used, and can dramatically affect the resulting SSSS for the Sun. Hall et al. (2007) used \n\n\n\n\n\n, \n\n\n\n\n\n, and \n\n\n\n\n\n. The constant \n\n\n\n\n\n was empirically determined to be 0.97 \u00b1 0.11 erg cm\u22122 s\u22121 in Hall et al. (2007) as the value that provides the best agreement between SSSS and \n\n\n\n\n\n from Baliunas et al. (1995) for an ensemble of stars and the Sun. This combination of parameters resulted in a mean SSSS of 0.170 for the Sun using observations covering cycle 23. A slightly different calibration of SSS data in Hall & Lockwood (2004) used a flux scale \n\n\n\n\n\n based on Johnson \n\n\n\n\n\n, set to 0.65 for the Sun, and \n\n\n\n\n\n K. In Table 1 we estimated that this calibration resulted in a mean S = 0.168 for cycle 23. Hall et al. (2009), which included a revised reduction procedure and one year of data with the upgraded camera (see below), found \n\n\n\n\n\n.","Citation Text":["Rutten 1984"],"Functions Text":["(simply K in other works) is the conversion factor from the MWO HKP-2 H & K flux (numerator of Equation (1)) to physical flux"],"Functions Label":["Uses"],"Citation Start End":[[1020,1031]],"Functions Start End":[[893,1018]]}
{"Identifier":"2019MNRAS.485L..78V__Chatterjee_et_al._2017_Instance_2","Paragraph":"The properties of the persistent radio source associated with FRB 121102 may be constrained independently of the Faraday-rotating medium. We assume equipartition between the relativistic gas and magnetic field as is common in synchrotron sources3 (Readhead 1994). The source becomes self-absorbed at $1.5$ GHz for radius $R_{\\rm per} < 0.05$ pc; this is thus the lower bound on the source size. European Very Long Baseline Interferometry (VLBI) Network observations of the source at 5 GHz set an upper bound on the source radius of Rper \u2272 0.35 pc (Marcote et al. 2017). This is consistent with the ${\\approx } 30\\, {{\\rm per\\, cent}}$ amplitude modulations observed in the source at 3 GHz (Chatterjee et al. 2017) being caused by refractive interstellar scintillation in the Milky Way interstellar medium (ISM; Walker 1998). For any radius within the allowed range (0.05  Rper\/pc  0.35), we can determine the equipartition magnetic field, Beq, and the column of relativistic electrons, Nrel, using the standard expressions for synchrotron emissivity and absorption coefficients (Rybicki & Lightman 1979, their equations 6.36 and 6.53). We assume a power-law energy distribution of radiating electrons with somewhat shallow index of b = \u22121.5 that can account for the relatively flat spectrum of the source (Chatterjee et al. 2017). The peak Lorentz factor of the distribution, \u03b3max, is chosen to correspond to the observed spectral break frequency of $\\nu _{\\rm max}=10$ GHz. If the lower Lorentz factor cut-off corresponds to emission at $\\nu _{\\rm min}=1$ GHz,4 then the equipartition magnetic field and electron column thus determined for minimum and maximum source sizes are $B_{\\rm eq}\\approx 140$ mG, $\\gamma _{ \\rm min}\\approx 50$, \u03b3max \u2248 160, $N_{\\rm rel} \\approx 0.95\\, {\\rm pc}\\, {\\rm cm}^{-3}$ for $R_{\\rm per}=0.05$ pc, and $B_{\\rm eq}\\approx 27$ mG, $\\gamma _{ \\rm min}\\approx 120$, \u03b3max \u2248 370, $N_{\\rm rel} \\approx 0.1\\, {\\rm pc}\\, {\\rm cm}^{-3}$ for $R_{\\rm per}=0.35$ pc. The reader can scale the equipartition field to other source sizes using Beq(R) \u221d R\u22126\/7. The total energy contained in the relativistic electrons and the magnetic field (\u2018equipartition energy\u2019) is \u223c1049.1 and \u223c1050.2 erg, respectively. If the relativistic electrons were injected in a one-off event, the synchrotron cooling rates at \u03b3max yield source ages of $14$ yr for R = 0.05 pc and 60 yr for $R=0.35$ pc. The corresponding expansion velocities are $0.011\\, c$ and $0.02\\, c$, respectively.","Citation Text":["Chatterjee et al. 2017"],"Functions Text":["We assume a power-law energy distribution of radiating electrons with somewhat shallow index of b = \u22121.5 that can account for the relatively flat spectrum of the source"],"Functions Label":["Uses"],"Citation Start End":[[1306,1328]],"Functions Start End":[[1136,1304]]}
{"Identifier":"2020MNRAS.497.2941S__Loeb_&_Barkana_2001_Instance_1","Paragraph":"After the cosmological recombination, the universe went into the dark ages during which the density fluctuations in the matter distribution grew, and after reaching a threshold, the matter collapsed to make the first bound objects. The nature of dark matter sets the timeline and characteristics of these first bound objects, which were the hosts for the first sources of light, so it is essential to see the impact of different dark matter models on the observables from the Cosmic Dawn and EoR. Then, the natural question that arises is whether one can use the differences in these observables, estimated for different dark matter models, in order to constrain the nature of dark matter. The present observational probes that allow us to have a peak in this epoch are the absorption spectra of high-redshift quasars (Loeb & Barkana 2001; White et al. 2003; Boera et al. 2019) and the Thomson scattering optical depth of the cosmic microwave background (CMB) radiation photons (Kaplinghat et al. 2003; Komatsu et al. 2011). However, these indirect probes provide very limited and weak constraints on the CD-EoR. The H\u2009i 21-cm line, which arises due to the hyperfine splitting of the ground state of the neutral hydrogen, is a direct and most promising probe to study this period. Motivated by this, a large number of radio interferometers, including the GMRT (Paciga et al. 2013), LOFAR (Ghara et al. 2020; Mertens et al. 2020), MWA (Barry et al. 2019; Li et al. 2019), and PAPER (Kolopanis et al. 2019), are attempting a statistical detection of this signal using the power spectrum statistic. In parallel, there is a complementary approach to detect the sky-averaged global 21-cm signal from the CD-EoR using experiments, e.g. the EDGES (Bowman et al. 2018), DARE (Burns et al. 2017), and SARAS (Singh et al. 2018). The next-generation interferometers like the SKA (Koopmans et al. 2015; Mellema et al. 2015) are expected to see a giant leap in the sensitivity, which will enable them to make tomographic images of the H\u2009i distribution across cosmic time.","Citation Text":["Loeb & Barkana 2001"],"Functions Text":["The present observational probes that allow us to have a peak in this epoch are the absorption spectra of high-redshift quasars"],"Functions Label":["Background"],"Citation Start End":[[819,838]],"Functions Start End":[[690,817]]}
{"Identifier":"2022MNRAS.509.1959S__Ezzeddine_et_al._2019_Instance_1","Paragraph":"However, the transition between the two extremes of modern (metal-rich) and primordial (metal-poor) star formation, and in particular the role of dust coupling and stellar radiation feedback at low metallicity, has thus far received limited exploration. Krumholz (2011) present analytical models for radiation feedback and predict a weak scaling of IMF peak mass with metallicity, while Myers et al. (2011) and Bate (2014, 2019) carry out radiation-hydrodynamic simulations of star formation over a metallicity range from $0.01{-}3\\, Z_{\\rm {\\odot }}$ and find negligible effects on gas fragmentation. However, these studies do not explore lower metallicities, despite available evidence for the existence of a low-metallicity ISM in the past through the discovery of stars with metallicities as low as $10^{-4}\\, \\rm {Z_{\\odot }}$ (Caffau et al. 2011; Starkenburg et al. 2018), as well as several others with $\\rm {[Fe\/H]} \\lt -5$ (Christlieb et al. 2004; Keller et al. 2014; Frebel et al. 2015; Aguado et al. 2017, 2018; Ezzeddine et al. 2019; Nordlander et al. 2019). Coming from the opposite direction, Bromm et al. (2001), Omukai et al. (2005), and Omukai, Hosokawa & Yoshida (2010) consider the thermodynamics of gas of increasing metallicity, and find that dust and metal line cooling permits fragmentation to reach masses \u22721 M\u2299 only once the metallicity exceeds \u223c10\u22123.5 Z\u2299. Dust is a more efficient coolant than metal lines, and permits fragmentation to lower masses at lower metallicity (e.g. Meece, Smith & O\u2019Shea 2014; Chiaki & Yoshida 2020; Shima & Hosokawa 2021), but exactly by how much depends on the poorly known distribution of dust grain sizes in the early Universe (Schneider et al. 2006, 2012; Omukai et al. 2010; Schneider & Omukai 2010; Chiaki et al. 2015). However, the early Universe star formation models are fundamentally misanalogous to the modern ones that consider decreasing metallicity, in that the early Universe models consider dust solely as a coolant that enables fragmentation, whereas the modern ones assign it a more nuanced role, as both a source of cooling and later, once stellar feedback begins, a source of heating \u2013 a changeover that seems crucial to explaining why the IMF in the present-day Universe peaks at ${\\sim}0.2\\, \\rm {M_{\\odot }}$ rather than ${\\sim}10^{-2}\\, \\rm {M_{\\odot }}$ (Kroupa 2001; Chabrier 2003, 2005).","Citation Text":["Ezzeddine et al. 2019"],"Functions Text":["However, these studies do not explore lower metallicities, despite available evidence for the existence of a low-metallicity ISM in the past through the discovery of stars with metallicities as low as $10^{-4}\\, \\rm {Z_{\\odot }}$","as well as several others with $\\rm {[Fe\/H]} \\lt -5$"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1023,1044]],"Functions Start End":[[602,831],[879,931]]}
{"Identifier":"2022MNRAS.509..903N__Murguia-Berthier_et_al._2014_Instance_1","Paragraph":"Binary neutron star (BNS) mergers have long been suspected to produce the central engines of short gamma-ray bursts (sGRBs) (Eichler et al. 1989). The link was firmly established in August 2017, after the combined detection of gravitational waves and an sGRB from the same BNS merger (Abbott et al. 2017a,b,c; Goldstein et al. 2017; Savchenko et al. 2017). The actual origin of the gamma-ray signal is still debated, coming from either an off-axis jet or from the shock breakout of a relativistic cocoon inflated by the jet itself (Kasliwal et al. 2017; Nakar & Piran 2017; Gottlieb, Nakar & Piran 2018a; Gottlieb et al. 2018b; Mooley et al. 2018a; Lundman & Beloborodov 2021). Nevertheless, the multiband observations of a rising afterglow (Alexander et al. 2017; Hallinan et al. 2017; Margutti et al. 2017, 2018; Troja et al. 2017, 2018, 2019; D\u2019Avanzo P. et al. 2018; Lyman et al. 2018; Lamb et al. 2019) together with the detection of superluminal motion (Mooley et al. 2018b; Ghirlanda et al. 2019; Hotokezaka et al. 2019) have settled the presence of a jet that successfully broke out from the surrounding ejecta, observed off-axis with a viewing angle \u03b8obs \u2248 19\u00b0 (Murguia-Berthier et al. 2017; Lamb, Mandel & Resmi 2018; Lazzati et al. 2018; Mooley et al. 2018a; Margutti & Chornock 2020). The information obtainable from afterglow observations is strongly dependent on the angular structure of the emerging jet. This structure might be mainly determined by the launching process (Kathirgamaraju et al. 2019) or it may arise as a consequence of the interaction with the surrounding environment during propagation. Therefore, an understanding of the processes that shape the jet could in principle provide insights into both jet formation, and the post-merger environment. Recent relativistic (magneto-) hydrodynamics simulations (Murguia-Berthier et al. 2014, 2017, 2021b; Geng et al. 2019; Beniamini et al. 2020b; Gottlieb, Levinson & Nakar 2020; Gottlieb & Nakar 2021; Gottlieb et al. 2021a; Gottlieb, Nakar & Bromberg 2021b; Hamidani & Ioka 2021; Lazzati et al. 2021; Nathanail et al. 2021; Pavan et al. 2021; Urrutia et al. 2021) have illustrated the importance of the ambient medium in shaping the jet, thereby highlighting the importance of understanding the remnant structure and ejecta properties. The joint events GW170817 and GRB170817A were followed by an additional electromagnetic transient spanning the spectral bands from UV to optical and IR on time-scales from days to weeks (Abbott et al. 2017b; Arcavi et al. 2017; Cowperthwaite et al. 2017; Drout et al. 2017; Evans et al. 2017; Pian et al. 2017; Smartt et al. 2017; Soares-Santos et al. 2017; Tanvir et al. 2017; Utsumi et al. 2017). The observed properties were consistent with the expectations for a thermal transient powered by the radioactive decay of freshly synthesized r-process elements (a so-called \u2018macronova\u2019 or \u2018kilonova\u2019 e.g. Li & Paczy\u0144ski 1998; Kulkarni 2005; Rosswog 2005; Metzger et al. 2010; Roberts et al. 2011; Kasen, Badnell & Barnes 2013; Yu, Zhang & Gao 2013; Kasen, Fern\u00e1ndez & Metzger 2015; Kasen et al. 2017; Metzger 2017; Perego, Radice & Bernuzzi 2017; Tanaka et al. 2017; Rosswog et al. 2018). Understanding the properties of this signal requires in-depth investigation of all the processes that can unbind material during and after a BNS merger, together with the available amount of free neutrons provided by each ejection channel. By now, several mass-ejection channels have been identified and they differ in terms of launch time, mass, electron fraction Ye, and velocity. During the merger \u223c10\u22123 \u2212 10\u22122 M\u2299 of material are ejected dynamically (Rosswog et al. 1998, 1999; Rosswog & Davies 2002; Oechslin, Janka & Marek 2007; Bauswein, Goriely & Janka 2013; Radice et al. 2018a). On longer time-scales (\u223c1 s) a few 10\u22122 M\u2299 of material can be unbound from the torus surrounding the remnant by the action of nuclear heating (Metzger, Piro & Quataert 2008; Fernandez & Metzger 2013; Fernandez et al. 2015; Just et al. 2015), magnetic (Siegel & Ciolfi 2015; Ciolfi et al. 2017; Siegel & Metzger 2017, 2018; Ciolfi & Kalinani 2020; Murguia-Berthier et al. 2021a), and viscous effects (Shibata, Kiuchi & Sekiguchi 2017; Fujibayashi et al. 2018, 2020; Radice et al. 2018b; Shibata & Hotokezaka 2019). Weak interactions play a key role since they can change the initially extremely low electron fraction and they are therefore of paramount importance for nucleosynthesis and the electromagnetic appearance of a BNS merger (Ruffert et al. 1997; Rosswog & Liebend\u00f6rfer 2003; Dessart et al. 2009; Perego et al. 2014, 2017; Martin et al. 2015, 2018; Miller et al. 2019; Murguia-Berthier et al. 2021a).","Citation Text":["Murguia-Berthier et al. 2014"],"Functions Text":["Recent relativistic (magneto-) hydrodynamics simulations","have illustrated the importance of the ambient medium in shaping the jet, thereby highlighting the importance of understanding the remnant structure and ejecta properties"],"Functions Label":["Motivation","Background"],"Citation Start End":[[1837,1865]],"Functions Start End":[[1779,1835],[2141,2311]]}
{"Identifier":"2016MNRAS.463.2716M__Cho_&_Lazarian_2007_Instance_2","Paragraph":"Returning to the case of HL Tau, where the possible contribution of an infalling envelope is not an issue, how can one reconcile the strong indication of a dominant radial field component in the polarization map with the expectation that the bulk of the mm-wavelength emission originates near the disc mid-plane, where the azimuthal field component dominates? One possibility is that the non-negligible optical depth inferred in the bright emission rings of HL Tau at mm-wavelengths (Jin et al. 2016; Pinte et al. 2016) shifts the emission centroid to finite disc elevations where the magnetic field already has a measurable radial component. However, in view of the very small scale height inferred for the mm-emitting dust in this source, this is unlikely to be the main explanation. Perhaps a more likely possibility is that, even in this comparatively young source, the grains near the mid-plane, which dominate the total intensity, have already grown to sizes that exceed the maximum size $a_\\mathrm{max} = \\lambda \/2\\pi$ for producing polarized emission at wavelength \u03bb (e.g. Cho & Lazarian 2007; for \u03bb = 1.25\u2009mm, amax = 0.2\u2009mm), while the smaller grains (with sizes a amax), which contribute efficiently to the polarized flux, remain suspended at high elevations (where the field is predominantly radial). Another effect that could lower the polarized emission from grains that have settled to the mid-pane is the likelihood that grains become less elongated as they grow (e.g. Hughes et al. 2009), which would tend to reduce the value of the coefficient C in equation (2) (C \u2192 0 as the grain axis ratio \u2192 1).7 This interpretation is supported by the finding in the high-resolution observations of IRAS 4A (Cox et al. 2015) of an average polarization of 15 per cent at 8\u2009mm and 10 per cent at 10\u2009mm, with a peak fractional polarization of \u223c20 per cent. If the intrinsic degree of mm-wavelength polarization in HL Tau is also of the order of 20 per cent, then it may be possible to explain the factor of \u223c10 lower value of P measured in this source at 1.25\u2009mm8 in terms of a dilution of the polarized emission from a \u2272 0.1\u2009mm grains at high disc elevations by weakly polarized emission of larger grains residing near the mid-plane. In this scenario, most of the grains that are responsible for the mm-wavelength flux have settled to the mid-plane and grown to sizes a \u2273 1\u2009mm.9 Although a fraction of these grains may have sizes in excess of 1\u2009mm and would therefore emit less efficiently at that wavelength than a \u2272 1\u2009mm grains (e.g. Miyake & Nakagawa 1993), the mid-plane region should still dominate the total mm-wavelength flux if most of the a \u2273 1\u2009mm grains are concentrated there. Grains of size a \u2272 0.1\u2009mm may be kept at high elevations by turbulent motions that can persist below the wind-driving surface layers (e.g. Simon et al. 2013, 2015; Bai 2015) as well as by the emerging outflows within these layers (see Safier 1993), with porosity effects (e.g. Ormel, Spaans & Tielens 2007) possibly also helping to mitigate gravity's pull towards the mid-plane. This scenario of course needs to be backed by detailed calculations and observational tests. One such test would be to obtain a polarization map of HL Tau at longer ( \u2273 1\u2009cm) wavelengths: if the above picture is correct and the grains in the mid-plane region are aligned, such a map could reveal a stronger (or even dominant) contribution from the azimuthal and (especially if \u039b0 \u226a 1) vertical field components.10 It is, however, conceivable that the large grains in this source are not well aligned because the radiative torque mechanism does not operate efficiently on them: this could happen if the characteristic wavelength of the anisotropic component of the local radiation field were much smaller than the mid-plane grain sizes (e.g. Cho & Lazarian 2007) or if the anisotropic radiation component inside the dust disc were weak due to finite optical depth effects. Note that the possibility of the mid-plane grains not being well aligned provides another reason for why the polarized mm-wavelength emission from this region could be weak.","Citation Text":["Cho & Lazarian 2007"],"Functions Text":["It is, however, conceivable that the large grains in this source are not well aligned because the radiative torque mechanism does not operate efficiently on them: this could happen if the characteristic wavelength of the anisotropic component of the local radiation field were much smaller than the mid-plane grain sizes (e.g."],"Functions Label":["Future Work"],"Citation Start End":[[3812,3831]],"Functions Start End":[[3485,3811]]}
{"Identifier":"2021ApJ...910..124X__Leslie_et_al._2016_Instance_1","Paragraph":"What is the dominant mode of star formation for AGNs in general, and for quasars in particular? The main sequence gives a useful framework for discussing the evolutionary status of AGN host galaxies and their relation to the galaxy population at large. The existing literature in this field, however, is complicated enormously by the diverse strategies of AGN sample selection, the accuracy of the SFR and M* tracers, and the myriad choices of main-sequence prescription. While there is almost unanimous agreement that AGNs of low to moderate luminosity (Lbol \u2272 1045 erg s\u22121) at z \u2248 0\u20133 lie on or below the main sequence (e.g., Shao et al. 2010; Mullaney et al. 2012, 2015; Santini et al. 2012; Rosario et al. 2013a; Shimizu et al. 2015; Ellison et al. 2016; Leslie et al. 2016; Suh et al. 2017; Bernhard et al. 2019; Grimmett et al. 2020; Jackson et al. 2020), no consensus exists for AGNs with Lbol > 1045 erg s\u22121. The situation is particularly contentious at redshifts higher than \u223c0.5, where luminous AGNs have been reported to be above (e.g., Rovilos et al. 2012; Florez et al. 2020; Kirkpatrick et al. 2020), on (e.g., Harrison et al. 2012; Xu et al. 2015; Stanley et al. 2017; Schulze et al. 2019), and below (e.g., Scholtz et al. 2018; Stemo et al. 2020) the main sequence. Fortunately, a better consensus of opinion can be found for luminous (Lbol \u2273 1045 erg s\u22121) AGNs at z \u2272 0.5. Most agree that low-redshift quasars are located largely on and above the main sequence (Husemann et al. 2014; Xu et al. 2015; Zhang et al. 2016; Stanley et al. 2017; Jarvis et al. 2020). Regardless of redshift, it appears that the magnitude of an AGN\u2019s offset from the main sequence correlates positively with its luminosity (Bernhard et al. 2019; Grimmett et al. 2020). For their large sample of z \u2248 0.3 type 1 AGNs with uniform SFRs based on extinction-corrected [O II] \u03bb3727 emission, Zhuang & Ho (2020) demonstrated that SFR systematically rises with increasing Lbol at fixed M*.","Citation Text":["Leslie et al. 2016"],"Functions Text":["While there is almost unanimous agreement that AGNs of low to moderate luminosity (Lbol \u2272 1045 erg s\u22121) at z \u2248 0\u20133 lie on or below the main sequence (e.g.,",", no consensus exists for AGNs with Lbol > 1045 erg s\u22121."],"Functions Label":["Similarities","Differences"],"Citation Start End":[[759,777]],"Functions Start End":[[472,627],[860,916]]}
{"Identifier":"2017MNRAS.464.2545C__L\u00f3pez-Corredoira_&_Molg\u00f3_2014_Instance_1","Paragraph":"One of the fundamental tasks of the Galactic studies is to estimate the structure parameters of the major structure components. Bahcall & Soneira (1980) fit the observations with two structure components, namely a disc and a halo. Gilmore & Reid (1983) introduce a third component, namely a thick disc, confirmed in the earliest Besancon Galaxy Model (Cr\u00e9z\u00e9 & Robin 1983). Since then, various methods and observations have been adopted to estimate parameters of the thin and thick discs and of the halo of our Galaxy. As the quantity and quality of data available continue to improve over the years, the model parameters derived have become more precise, numerically. Ironically, those numerically more precise results do not converge (see table 1 of Chang, Ko & Peng 2011, table 2 of L\u00f3pez-Corredoira & Molg\u00f3 2014 and sections 5 and 6 of Bland-Hawthorn & Gerhard 2016, for a review). The scatters in density law parameters, such as scale lengths, scale heights and local densities of these Galactic components, as reported in the literature, are rather large. At least parts of the discrepancies are caused by degeneracy of model parameters, which in turn can be traced back to the different data sets adopted in the analyses. Those differing data sets either probe different sky areas (Bilir et al. 2006a; Du et al. 2006; Cabrera-Lavers et al. 2007; Ak et al. 2007; Yaz & Karaali 2010; Yaz G\u00f6k\u00e7e et al. 2015), are of different completeness magnitudes and therefore refer to different limiting distances (Karaali et al. 2007), or consist of stars of different populations of different absolute magnitudes (Karaali, Bilir & Hamzaolu 2004; Bilir et al. 2006b; Juri\u0107 et al. 2008; Jia et al. 2014). It should be noted that the analysis of Bovy et al. (2012), using the SEGUE spectroscopic survey, has given a new insight on the thin and thick disc structural parameters. This analysis provides estimate of their scale height and scale height as a function of metallicity and alpha abundance ratio. However, it relies on incomplete data (since it is spectroscopic) with relatively low range of Galactocentric radius as for the thin disc is concerned.","Citation Text":["L\u00f3pez-Corredoira & Molg\u00f3 2014"],"Functions Text":["As the quantity and quality of data available continue to improve over the years, the model parameters derived have become more precise, numerically. Ironically, those numerically more precise results do not converge","table 2 of","The scatters in density law parameters, such as scale lengths, scale heights and local densities of these Galactic components, as reported in the literature, are rather large."],"Functions Label":["Motivation","Motivation","Motivation"],"Citation Start End":[[785,814]],"Functions Start End":[[518,734],[774,784],[885,1060]]}
{"Identifier":"2015AandA...579A.132P__Simha_et_al._(2009)_Instance_1","Paragraph":"A common feature of all previous models is that the relation between the central galaxy stellar mass and the halo mass reaches a maximum at halo masses ~1012 M\u2299. According to Yang et al. (2012), below this threshold the mass accretion of the central galaxy is dominated by star formation. Thus, when the halo mass reaches ~1012 M\u2299 a process takes place that quenches the star formation. Interestingly, this mass scale is very similar to the cold-mode to hot-mode transition scale (Birnboim & Dekel 2003; Kere\u0161 et al. 2005) in the theory of gas accretion, as derived in hydrodynamic simulations, whereas large halos primarily accrete hot gas and low mass halos cold gas. This would suggest that the quenching of central galaxies coincides with the formation of a hot gaseous halo, and thus with a lack of cold gas supply. What would be the fate of satellites? According to Simha et al. (2009), the subhalos also retain their identity for quite some time after accreting a larger halo, so satellites in subhalos less massive than ~1012 M\u2299 do not immediately see the effect of the hot gas in the larger halo and accrete in cold mode. Thus, consistent with the results of Yang et al. (2012) and B\u00e9thermin et al. (2013), satellite galaxies continue to accrete gas and convert it to stars over a rather long period, which according to Simha et al. (2009) is about of 0.5\u22121 Gyr after the merger. The gas accretion declines steadily over this period. Since star formation follows mass accretion with a short delay, satellites should experience quenching in a similar amount of time. This scenario would be consistent with our observations. Indeed, at z ~ 1 when massive halos are just forming via merger, the SF activity in the accreted subhalos is still high. At later epochs, instead, the transition to the hot mode accretion of the satellites and the consequent progressive quenching of their SF activity would lead to the faster decline of their contribution to the CSFH with respect to lower mass halos, which evolve in a cold mode accretion phase maintaining a high SFR. ","Citation Text":["Simha et al. (2009)"],"Functions Text":["According to","the subhalos also retain their identity for quite some time after accreting a larger halo, so satellites in subhalos less massive than ~1012 M\u2299 do not immediately see the effect of the hot gas in the larger halo and accrete in cold mode.","Thus, consistent with the results of Yang et al. (2012) and B\u00e9thermin et al. (2013), satellite galaxies continue to accrete gas and convert it to stars over a rather long period,"],"Functions Label":["Uses","Uses","Similarities"],"Citation Start End":[[872,891]],"Functions Start End":[[859,871],[893,1130],[1131,1309]]}
{"Identifier":"2016MNRAS.457.2480C___2008_Instance_1","Paragraph":"A number of ideas have been put forward to explain the formation and early evolution of the compact Kepler and radial velocity systems, which in cases such as Gliese 581 and HD 69830 appear to contain in excess of \u223c30\u2009M\u2295 of solid material within a few tenths of an au (Lovis et al. 2006; Udry et al. 2007). This concentration of solids close to the star led to classical core accretion models combined with disc-driven migration being developed using population synthesis codes (Alibert et al. 2006). More recent population synthesis calculations that also include prescriptions for planet\u2013planet interactions have also been presented (Ida & Lin 2010). N-body simulations, combined with either hydrodynamic simulations or analytic prescriptions for migration and eccentricity\/inclination damping of planetary growth, have also been used to examine the origins of such systems (Cresswell & Nelson 2006, 2008; Terquem & Papaloizou 2007; McNeil & Nelson 2009, 2010; Hellary & Nelson 2012; Cossou, Raymond & Pierens 2013; Coleman & Nelson 2014; Hands, Alexander & Dehnen 2014). A common outcome of these N-body simulations is the formation of resonant convoys of planets in the presence of convergent migration, an outcome that is not reflected in the Kepler systems. Various ideas have been put forward to explain why the resonances may be unstable, including tidal eccentricity damping followed by separation of the resonance for short-period systems (Terquem & Papaloizou 2007), stochastic migration due to local turbulence (Adams, Laughlin & Bloch 2008; Rein & Papaloizou 2009; Rein 2012) \u2013 a process that is likely to only operate close to the star where the disc can be thermally ionized (Umebayashi & Nakano 1988; Desch & Turner 2015), resonance breaking due to overstable librations (Goldreich & Schlichting 2014), orbital repulsion due to non-linear spiral wave damping in planet coorbital regions (Podlewska-Gaca, Papaloizou & Szuszkiewicz 2012; Baruteau & Papaloizou 2013).","Citation Text":["Cresswell & Nelson","2008"],"Functions Text":["N-body simulations, combined with either hydrodynamic simulations or analytic prescriptions for migration and eccentricity\/inclination damping of planetary growth, have also been used to examine the origins of such systems"],"Functions Label":["Background"],"Citation Start End":[[877,895],[902,906]],"Functions Start End":[[653,875]]}
{"Identifier":"2017MNRAS.470..612F__Feng_etal._2016_Instance_2","Paragraph":"The millimetre bump in M87 as recently observed by the Atacama Large Millimeter\/submillimeter Array can be naturally modelled by the synchrotron emission of the thermal electrons in the ADAF, which is different from the prediction of the jet model. Therefore, it provides an opportunity to explore the accretion process near the BH horizon. In particular, Feng etal. (2016) and Li etal. (2016) both found that the rotation measure predicted from the ADAF is roughly consistent with the observational values. It is still difficult to constrain the BH spin parameter from the modelling of the SED of M87 due to some degeneracy in model parameters (wind parameter, s, magnetic parameter ; Feng etal. 2016). The spin parameter can be better constrained from the jet model if the relativistic jet is indeed powered by the rotating BHs as suggested by MHD simulations and some observations. We find that the dimensionless BH spin parameter should be larger than 0.96 for the lower limit of jet power derived from the X-ray cavities (e.g. Rafferty et al. 2006; Russell etal. 2013b). In this work, we adopt several typical values of parameters (e.g. 0.1, 0.3 and 0.5). The larger value of  will lead to a lower accretion rate near the horizon to explain the observed millimetre bump, and the BHs need to rotate faster to reproduce the observed jet power. The peak of synchrotron emission from the thermal electrons of ADAF will move to the submillimetre waveband if  is too small, which is different from the observed millimetre bump. We adopt the equipartition case of 0.5 in our calculations, where magnetic energy will become dominant if the BH is fast spinning, considering the possible amplification of the magnetic field by the frame dragging effect. For the weaker magnetic case (e.g. 0.5), the BH needs to rotate faster to explain the observed SED and jet power. We find that our results are not sensitive to the viscosity parameter . Therefore, we suggest that the BH should be fast rotating in M87 even after considering the possible uncertainties.","Citation Text":["Feng etal. 2016)"],"Functions Text":["It is still difficult to constrain the BH spin parameter from the modelling of the SED of M87 due to some degeneracy in model parameters (wind parameter, s, magnetic parameter ;"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[686,702]],"Functions Start End":[[508,685]]}
{"Identifier":"2018MNRAS.473.1879R__Freudling_et_al._2011_Instance_1","Paragraph":"Cosmic H\u2009i gas density ($\\Omega _{\\rm {H}\\,\\small {I}}$) as a function of redshift (bottom axis) and lookback time (top axis). All measurements are corrected to the same cosmological parameters. Some DLA measurements adopted a different definition of the cosmic H\u2009i density taking into account neutral gas abundance (\u03a9gas) including helium or contributions from Lyman-\u03b1 absorbers with lower column density ($\\log N({\\rm H\\,{\\small {I}}}) < 20.3$). We have corrected all measurements to a consistent definition of $\\Omega _{\\rm {H}\\,\\small {I}}$. The large red star shows the $\\Omega _{\\rm {H}\\,\\small {I}}$ measurement from this work. The small black square and triangle at z \u223c 0 are the HIPASS and ALFALFA 21-cm emission measurements by Zwaan et al. (2005) and Martin et al. (2010), respectively. The red diamonds are from the Parkes telescope with an H\u2009i stacking technique (Delhaize et al. 2013). The open circle and square are the results from the Arecibo Ultra Deep Survey (AUDS) (Freudling et al. 2011; Hoppmann et al. 2015). Two right-pointing triangles are measured using the WSRT and H\u2009i stacking technique by Rhee et al. (2013). The pink triangle is measured by Lah et al. (2007) applying the GMRT 21-cm emission stacking. The green hexagon denotes the H\u2009i stacking measurement for the COSMOS field using the GMRT (Rhee et al. 2016). The black arrow line is the upper limit constrained using the GMRT H\u2009i stacking (Kanekar et al. 2016). The closed and open diamonds, left-pointing triangle, big circle, open triangles, open diamonds and small circles are damped Lyman-\u03b1 measurements from the HST and the SDSS by Rao et al. (2006, 2017), Neeleman et al. (2016), Prochaska et al. (2005), Noterdaeme et al. (2009), Noterdaeme et al. (2012) and Bird et al. (2017), respectively. The downward triangles and big square at high redshift of z > 2 are ESO UVES and Gemini GMOS measurements of DLAs by Zafar et al. (2013), Crighton et al. (2015), respectively. The black line with a shaded area shows a weighted fit of all $\\Omega _{\\rm {H}\\,\\small {I}}$ measurements and its 95 per cent confidence interval. The solid blue line shows the semi-analytic model prediction of Kim et al. (2015), and the dashed red line shows the mufasa (Dav\u00e9 et al. 2017) simulation.","Citation Text":["Freudling et al. 2011"],"Functions Text":["The open circle and square are the results from the Arecibo Ultra Deep Survey (AUDS)"],"Functions Label":["Uses"],"Citation Start End":[[986,1007]],"Functions Start End":[[900,984]]}
{"Identifier":"2018ApJ...863..194L__Kraft_et_al._1991_Instance_1","Paragraph":"Five XRT observations got only low exposure (i.e., much less than 1 ks) and the X-ray source was therefore undetected in these data sets. Surprisingly, we also found that the source was undetected in a \u201cdeep\u201d observation taken on 2015 August 19 with an exposure time of about 1.6 ks (Table 2). Within a 47\u2033 radius circular region centered at the source position (corresponding to 90% of the encircled energy fraction of XRT at 1.5 keV; Moretti et al. 2005), only one photon (which is located near the edge of the region) was detected in this 1.6 ks observation. Even assuming that only this event is from the source, the inferred count rate is much lower than the measurements in 2010 and 2015\u20132017; for example, seven source counts would have been detected in a 1.6 ks observation with the count rate of 4.4 \u00d7 10\u22123 cts s\u22121 measured five days later (Table 2). Using a Bayesian approach (Kraft et al. 1991), we computed 95% upper limits for all the nondetections. As expected, the upper limits for data with 1 ks are not very much constraining (i.e., a few \u00d710\u22122 cts s\u22121, while the average count rate of the four individual detections is about 4 \u00d7 10\u22123 cts s\u22121). The upper limit for the 1.6 ks data is deeper (i.e., 8.7 \u00d7 10\u22123 cts s\u22121), but still insufficient to clarify whether the low-count-rate measurement is physically or statistically based. For a deeper constraint, we combined all the six XRT observations, and the X-ray source can be marginally detected in the stacked image with \n\n\n\n\n\n cts s\u22121. Although this marginal detection shows a \u223c50% decrease on flux in the period from 2015 February 04 through August 19, the variability is not statistically significant (i.e., less than 2\u03c3). To check whether this variability was seen at other frequencies, we performed a Fermi-LAT analysis with the data collected between 2015 February 04 and August 19, and the \u03b3-ray flux (100 MeV\u2013100 GeV) did not vary significantly. In UV, there are some Ultraviolet\/Optical Telescope (UVOT) images taken simultaneously with the XRT observations. Although the UVOT magnitudes (obtained by aperture photometry using the uvotsource task in HEAsoft v6.22) significantly changed over time (Table 2), this was due to the orbital modulation (Figure 3; will be discussed in the coming sections).","Citation Text":["Kraft et al. 1991"],"Functions Text":["Using a Bayesian approach","we computed 95% upper limits for all the nondetections."],"Functions Label":["Uses","Uses"],"Citation Start End":[[887,904]],"Functions Start End":[[860,885],[907,962]]}
{"Identifier":"2019MNRAS.487.2412G__Weinberg_1993_Instance_1","Paragraph":"There are only three exceptions from the picture discussed above. In Fig. 2 we see that the model with the smallest number of objects, $N=40\\, 000$, does not show abrupt dissolution. This model is characterized with the shortest initial relaxation time (among models considered in this paper), about 1.6 Gyr, and the smallest number of retained BHs, about 30. In such a situation, according to Breen & Heggie (2013), BHs are quickly kicked out from the system and BHS cannot survive for a long time. So, the cluster evolution is not governed by strong energy generation by BHs in BHS. In Fig. 3 we can see that the model with W0 = 3 dissolves extremely fast and the model with W0 = 9 does not show the abrupt dissolution feature. The very fast dissolution of tidally filling models with low King model concentration was already extensively discussed in the literature (e.g. Weinberg 1993; Fukushige & Heggie 1995; Whitehead et al. 2013; Contenta et al. 2015). The dissolution of such clusters is controlled by very strong initial mass-loss powered by stellar\/binary evolution. Relaxation process is not important at all. The situation is much different for the model with W0 = 9. As it can be seen from Fig. 5, the cluster enters the core-collapse phase, so it has to be controlled by the relaxation process. The difference between the model with W0 = 6 and the model with W0 = 9 is connected with the rate of BHS evolution. Model W0 = 9 is initially much denser, so its half-mass radius and half-mass relaxation time are shorter than for the W0 = 6 model. Nevertheless, for both models the BH mass segregation ends nearly at the same time, about 4 Gyr. At that time, the models contain about 50 BHs and 160 BHs for W0 = 9 and W0 = 6, respectively. According to Breen & Heggie (2013) the evolution of BHs is controlled by the energy flow through the cluster half-mass radius, which is proportional to Eb\/Trh \u2248 GM\/Rh\/Trh, where Eb is the cluster binding energy. Models with W0 = 9 have smaller half-mass radius and half-mass relaxation time than models with W0 = 6, so the energy demand is larger for this model and leads to a much faster burning out of BHs \u2013 larger number of dynamical interactions leading to BH removal from the system. After the time of BH mass segregation (about 4 Gyr) the model W0 = 6 contains enough BHs to form a BHS and enters the phase of balanced evolution (Breen & Heggie 2013). Contrarily, the model with W0 = 9 has too small a number of BHs and continues to remove BHs quickly to support the needed energy flow. Finally, it enters the phase when other energy sources connected with ordinary binaries take over and the cluster enters the core-collapse and then core-bounce phases. It is important to note that the tidal field plays an important role in the above picture. The less concentrated model has larger Rh and loses more mass, so it is easier for the BHS to provide the needed energy to support the cluster structure. This phase of evolution ends when the cluster suddenly loses its virial equilibrium.","Citation Text":["Weinberg 1993"],"Functions Text":["The very fast dissolution of tidally filling models with low King model concentration was already extensively discussed in the literature (e.g."],"Functions Label":["Background"],"Citation Start End":[[874,887]],"Functions Start End":[[730,873]]}
{"Identifier":"2015ApJ...815..127W__Borucki_et_al._2010_Instance_1","Paragraph":"Since its launch in March of 2009, the NASA Kepler mission has been monitoring \u223c160,000 stars in order to detect transiting extrasolar planets with high relative photometric precision (\u223c20 ppm in 6.5 hr, Jenkins et al. 2010). In 2013 May, the Kepler main mission ended with the failure of a second reaction wheel; however, the first four years of Kepler data have led to a wealth of planetary discoveries with a total of 4706 announced planet candidates13\n\n13\n\nhttp:\/\/exoplanetarchive.ipac.caltech.edu\/ as of 2015 November 11.\n (Borucki et al. 2010, 2011; Batalha et al. 2013; Burke et al. 2014). The confirmed and candidate exoplanets typically have orbital periods shorter than 1000 days because at least three detected transits are needed for identification by the automated Transit Planet Search algorithm. Therefore, transiting exoplanets with periods longer than \u223c1000 days are easily missed. The detection of short-period planets is further favored because the transit probability decreases linearly with increasing orbital distance. For these reasons, estimates of the statistical occurrence rate of exoplanets tend to focus on orbital periods shorter than a few hundred days (e.g., Dong & Zhu 2013; Fressin et al. 2013; Petigura et al. 2013). Radial velocity (RV) techniques also favor the detection of shorter period orbits. While gas giant planets have been discovered with orbital periods longer than a decade, their smaller reflex velocity restricts detection of sub-Neptune mass planets to orbital radii less than \u223c1 AU (Lovis et al. 2011). In principle, astrometric observations favor longer period orbits; however, high precision needs to be maintained over the correspondingly longer time baselines. For shorter periods, the planets need to be massive enough to introduce a detectable astrometric wobble in the star and Gaia should begin to contribute here (Perryman et al. 2001). Microlensing offers sensitivity to planets in wider orbits and has contributed to our statistical knowledge about occurrence rates of longer period planets (e.g., Gaudi 2010; Cassan et al. 2012) and direct imaging of planets in wide orbits is also beginning to contribute important information (Oppenheimer & Hinkley 2009).","Citation Text":["Borucki et al. 2010"],"Functions Text":["In 2013 May, the Kepler main mission ended with the failure of a second reaction wheel; however, the first four years of Kepler data have led to a wealth of planetary discoveries with a total of 4706 announced planet candidates"],"Functions Label":["Background"],"Citation Start End":[[529,548]],"Functions Start End":[[226,453]]}
{"Identifier":"2019MNRAS.484.1912R__Iyyani_et_al._2016_Instance_1","Paragraph":"Even though most of the energy released by a gamma-ray burst (GRB) is emitted during the prompt emission phase, the emission mechanism is still not understood. In order to determine the emission process typically, the low-energy spectral index, \u03b1, of the GRB spectrum is analysed. This analysis has not given a conclusive answer since the \u03b1-distribution is broad and has not been uniquely explained by a single emission process. For instance, the peak of the distribution has been used as an argument for synchrotron emission since it is close to the expected value. However, a large fraction (${\\sim } 28{{\\ \\rm per\\ cent}}$) was found to be inconsistent with the theoretical limit of \u22122\/3 (\u2018line of death\u2019; Preece et al. 1998; Ghirlanda, Celotti & Ghisellini 2002; Guiriec et al. 2015; Goldstein et al. 2016; Yu et al. 2016); and only specific physical scenarios remain plausible (large emitting radii and Lorentz factors of the flow, Beniamini & Piran 2013; Iyyani et al. 2016; Beniamini, Barniol Duran & Giannios 2018; Burgess et al. 2018). Alternatively, the spectral width or sharpness angle has been used as a tool to characterize the spectrum, which also take into account the high-energy spectral slope (Axelsson & Borgonovo 2015; Yu et al. 2015b). Around 80 per\u2009cent of the bursts were found to have fitted Band spectra that are narrower than what was expected for synchrotron emission. However, direct fitting with a synchrotron model decreases this fraction (Burgess 2017). Therefore, a firm conclusion based on the spectral width or sharpness angle alone cannot either be reached. Yet another approach has been to study the correlation between spectral parameters. A well-studied and strong correlation is the Golenetskii correlation (Golenetskii et al. 1983) which relates the instantaneous flux, F, and peak of the spectrum Epk (Kargatis et al. 1994; Borgonovo & Ryde 2001; Lu et al. 2012). Another prominent correlation is between Epk and \u03b1 which is valid in a fraction of bursts (Crider et al. 1997; Kaneko et al. 2006). Both of these correlations show a variety of behaviours. In some individual pulses, the spectral parameters track each other, while in others the correlation is different during the rise and decay phases of the pulses. The variety of behaviours have complicated the search for physical explanations.","Citation Text":["Iyyani et al. 2016"],"Functions Text":["However, a large fraction (${\\sim } 28{{\\ \\rm per\\ cent}}$) was found to be inconsistent with the theoretical limit of \u22122\/3 (\u2018line of death\u2019;","and only specific physical scenarios remain plausible (large emitting radii and Lorentz factors of the flow,"],"Functions Label":["Differences","Differences"],"Citation Start End":[[961,979]],"Functions Start End":[[567,708],[828,936]]}
{"Identifier":"2018MNRAS.480.1639D__cent,_Ricci_et_al._2015_Instance_1","Paragraph":"Six out of the eight sources (75 per cent of the sample) show evidence for high obscuration by cold gas (column density NH\u2273 1023 cm\u22122). The number of AGN with NH\u2273 1024 cm\u22122 is in the range 2\u20134 when 90 per cent confidence errors on the column density are considered (see Table 4). Although our sample is definitely too small to make any strong conclusion in terms of statistical incidence of obscured AGN in dual systems, we note that the fraction of CT AGN in our sample of dual AGN is 25\u201350 per cent, i.e. higher than the fraction of hard X-ray selected CT AGN in isolated systems in a similar luminosity interval (BAT: 27 \u00b1 4 per cent, Ricci et al. 2015, NuSTAR: \u223c4 per cent, Marchesi et al. 2018). Recently, it has been suggested (Ricci et al. 2017) that AGN in multiple systems in a late stage of merging (with a projected distance d \u2264 10 kpc) are often type 2 AGN due to the huge amount of gas and dust obscuration channelled towards the nuclear region by the close galaxy encounter. Furthermore, in a sample of 44 UltraLuminous IR Galaxies (ULIRG) from the GOALS survey (Sanders et al. 2003) in the local universe, it has been demonstrated that the fraction of CT AGN in late mergers is higher (65$^{+12}_{-13}$ per cent) than in local isolated hard X-ray selected AGN (27$^{+4}_{-4}$ per cent, Ricci et al. 2017), while it is marginally higher than isolated AGN in early mergers (35$^{+13}_{-12}$ per cent). The increasing of the obscuration with the disturbed or interacting morphology of the galaxy host has been also observed up to higher redshift (up to z = 1.5) using X-ray Chandra data from deep surveys (Kocevski et al. 2015). These authors have analysed 154 heavily obscured AGN and compared their morphologies with control samples composed of moderately obscured and unobscured AGN. They found that the fraction of galaxies undergoing mergers is higher in AGN with NH above 3 \u00d7 1023 cm\u22122 with respect to unobscured AGN with NH  1022 cm\u22122 (7.8$^{+1.9}_{-1.3}$ versus 21.5$^{+4.2}_{-3.3}$ per cent). X-rays (Kocevski et al. 2015; Ricci et al. 2017) as well as more recent mid-IR colour selection (Satyapal et al. 2017) strongly suggest that AGN in late state of merging (10 kpc) are highly absorbed (but see also Ellison et al. 2017), with NH above 1024 cm\u22122. All these observational results suggest that CT AGN in dual systems correspond to a phase where the supermassive black hole is still accreting gas during an early stage of an interaction or merging.","Citation Text":["Ricci et al. 2015"],"Functions Text":["Although our sample is definitely too small to make any strong conclusion in terms of statistical incidence of obscured AGN in dual systems, we note that the fraction of CT AGN in our sample of dual AGN is 25\u201350 per cent, i.e. higher than the fraction of hard X-ray selected CT AGN in isolated systems in a similar luminosity interval (BAT: 27 \u00b1 4 per cent,"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[638,655]],"Functions Start End":[[280,637]]}
{"Identifier":"2016ApJ...827...58S__Stern_et_al._2014_Instance_1","Paragraph":"The multiwavength UV to mid-IR SED of J1224+5555 supports the scenario of a highly obscured AGN, based on the empirically derived templates of Assef et al. (2010). These templates consist of a set of three galaxy templates and one AGN template. The galaxy templates, E, Sbc, and Im, are based on templates from Coleman et al. (1980) with wavelength ranges extended using stellar models from Bruzual & Charlot (2003) and added dust and polycyclic aromatic hydrocarbon components from Devriendt et al. (1999). The AGN template is based on the average type 1 AGN template from Richards et al. (2006) with efforts to remove host galaxy contamination and create a template for an unobscured AGN. This method has since been used to successfully uncover highly obscured and even Compton-thick AGNs in several other works thus far (Chung et al. 2014; Hainline et al. 2014; Lansbury et al. 2014; Stern et al. 2014; Assef et al. 2015; Chen et al. 2015; Lansbury et al. 2015). Figure 7 shows the best-fit linear combination of the three galaxy templates to the observed SDSS, Two Micron All Sky Survey (2MASS), and WISE photometry of the target galaxy, where the best fit is obtained through a \n\n\n\n\n\n minimization. The multiwavelength fluxes were obtained from the NASA\/IPAC Extragalactic Database (NED).17\n\n17\n\nhttps:\/\/ned.ipac.caltech.edu\/\n\n The maximum wavelength of the Assef et al. (2010) templates is 30 \u03bcm; therefore, we cannot use the IRAS  60 \u03bcm flux in our fits. As discussed in Section 5, the 60 \u03bcm flux for J1224+5555 is likely overestimated, and its use is in any case questionable. While the SDSS \n\n\n\n\n\n, 2MASS \n\n\n\n\n\n, and WISE W1 bands are well fit by a combination of the Sbc and Im templates, the combined model flux is much lower than the observed fluxes in the W2\u2013W4 bands. In fact, Table 2 of Assef et al. (2010) shows that at z = 0 and z = 0.1 no galaxy template has a W1\u2013W2 color greater than 0.2 (Vega). However, when we add in the AGN template with a variable level of obscuration (red line) in Figure 7, we see that the observed multiwavelength photometry of the galaxy is better fit with a combination of a young and old stellar population and highly obscured AGNs (red line) dominating the WISE bands, with a \n\n\n\n\n\n value almost a factor of 7 times lower than the best fit obtained using the galaxy templates alone. As an additional test, we investigated the effect of adding variable additional extinction to the galaxy templates alone to determine if a comparable fit can be obtained without invoking the presence of an obscured AGN contribution. However, as can be seen from Figure 7, the best-fit model cannot adequately fit the SED in the WISE bands and results in a \n\n\n\n\n\n value that is significantly worse than the model that includes an AGN (Figure 7). We note that the improvement in fit when the AGN template is added cannot be attributed simply to having additional model parameters in the fit. In the first model without an AGN, we achieve a \n\n\n\n\n\n of 1323.29, with three model parameters. While the addition of an AGN template more than doubles the number of model parameters to 8, \n\n\n\n\n\n is reduced by a factor of \n\n\n\n\n\n, indicating that a model that includes an AGN template produces a significant improvement to the fit of the data.","Citation Text":["Stern et al. 2014"],"Functions Text":["This method has since been used to successfully uncover highly obscured and even Compton-thick AGNs in several other works thus far"],"Functions Label":["Background"],"Citation Start End":[[887,904]],"Functions Start End":[[691,822]]}
{"Identifier":"2017ApJ...845..160P__Camenzind_1986a_Instance_1","Paragraph":"It is a pressing question as to how the radiation that is observed in relativistic jets in active galactic nuclei (AGNs) is generated (e.g., Blandford & K\u00f6nigl 1979; Marscher 1980; Zensus 1997; Laing & Bridle 2002; Honda 2010; Levinson & Rieger 2011; Mo\u015bcibrodzka et al. 2011; Ito et al. 2013; Mason et al. 2013; Potter & Cotter 2013; Hovatta et al. 2014; Scott & Stewart 2014; Shih & Stockton 2014; Wang et al. 2014; Turner & Shabala 2015; Asada et al. 2016; Hirotani et al. 2016; Koay et al. 2016; Khabibullin et al. 2016; Prieto et al. 2016). Although there is a common consensus that the emitters are energetic particles, how these particles are accelerated to such high energies, how they dissipate their energy, and how they are transported with the jets themselves are still the subject of feverish investigation (e.g., Blandford & Eichler 1987). It is argued that relativistic outflows from black holes are associated with accretion flows (Blandford 1976; Fender et al. 2004; Meier 2005; Ferreira et al. 2006; Trump et al. 2011; Pu et al. 2012; Wu et al. 2013; Ishibashi et al. 2014; Sbarrato et al. 2014). In the case of collimated relativistic jets, magnetic fields must play an important role (Camenzind 1986a, 1986b, 1987; Fendt & Greiner 2001; Vlahakis & K\u00f6nigl 2004; Komissarov et al. 2007; Lyubarsky 2009; Nakamura & Asada 2013; Homan et al. 2015), and it is argued that jets are powered at the expense of the black hole, wherein energy is extracted from a reservoir of rotational energy from the black hole itself, either by electromagnetic means (Blandford & Znajek 1977; Komissarov 2004, 2005; Toma & Takahara 2014) or through magnetohydrodynamical processes (Phinney 1983; Takahashi et al. 1990; Koide et al. 2002; McKinney & Gammie 2004; Hawley & Krolik 2006). Models have been proposed for both of these cases, and they both in principle possess certain testable predictions. In particular, for the latter, numerical GRMHD simulations (e.g., McKinney 2006) and analytical GRMHD studies (e.g., Takahashi et al. 1990; Pu et al. 2016) consistently show the presence of a stagnation or separation surface (a separatrix). This surface separates the (inner) inflow region from the (outer) outflow region, both of which follow the same global, black-hole-threading magnetic field lines. The relatively slow radial velocities near the stagnation surface imply a high concentration of fluid particles. If energetic particles are injected in the vicinity of the stagnation surface or near the black hole event horizon, they must accumulate in high concentrations near the stagnation surface, provided that the cooling timescale is not significantly shorter than the dynamical timescale of the jet fluid flow. This surface is a unique feature of relativistic GRMHD jets and in contrast to an ideal force-free magnetic jet (e.g., McKinney & Narayan 2007; Tchekhovskoy et al. 2008; Broderick & Loeb 2009).","Citation Text":["Camenzind 1986a"],"Functions Text":["In the case of collimated relativistic jets, magnetic fields must play an important role"],"Functions Label":["Background"],"Citation Start End":[[1205,1220]],"Functions Start End":[[1115,1203]]}
{"Identifier":"2016AandA...586A..81H__Massey_(2002)_Instance_1","Paragraph":"\n\nTable 2\n\nCatalogue description.\n\n\n\nColumn\n Description\n\n\n\n\n\n\n\n1\nSource number\n\n\n2\u22123\nX-ray coordinates, right ascension and declination (epoch 2000.0)\n\n\n4\nUncertainty of X-ray position [\u2032\u2032]. For XMM-Newton positions taken from Sturm et al. (2013c) the 1\u03c3  error includes a systematic uncertainty of 0.5\u2032\u2032.\n\n\n5\nOrigin of the X-ray coordinate (A: ASCA, C: Chandra, E: Einstein, I: Integral, N: XMM-Newton, R: ROSAT, S: Swift, X: RXTE). When no reliable position could be determined from the non-imaging RXTE collimator-instruments, a radius of 30\u2032 for the position error indicates the size of the field of view.\n\n\n6\nReference for source discovery.\n\n\n7\nIdentification of optical counterpart with emission-line star from Meyssonnier & Azzopardi (1993). The negative number indicates a star found in the catalogue of Murphy & Bessell (2000).\n\n\n8\u221215\nFlags indicating different source properties. For their description see Table 3.\n\n\n16\nConfidence class (values 1\u22126, see Table 4).\n\n\n17\u221218\nOptical coordinates, right ascension and declination (epoch 2000.0) for the identified counterpart from Zaritsky et al. (2002), or \u2013 when not available there \u2013 from Massey (2002).\n\n\n19\u221226\nThe Magellanic Clouds Photometric Survey (MCPS): U, error(U), B, error(B), V, error(V), I, error(I) [mag] from Zaritsky et al. (2002).\n\n\n27\u221232\nColour indices U\u2212B, error(U\u2212B), B\u2212V, error(B\u2212V), V\u2212I, error(V\u2212I) [mag] derived from MCPS photometry.\n\n\n33\u221234\nReddening-free Q-value (Q = U\u2212B\u22120.72 \u00d7 (B\u2212V)) and error(Q) [mag].\n\n\n35\nNear-IR counterpart to the optical star from the Two Micron All Sky Survey (2MASS, Skrutskie et al. 2006).\n\n\n36\u221241\nNear-IR magnitudes with corresponding errors: J, error(J), H, error(H), K, error(K) [mag].\n\n\n42\u221245\nNear-IR colour indices, J\u2212H, error(J\u2212H), H\u2212K, error(H\u2212K) [mag].\n\n\n46\u221253\nSpitzer IRAC fluxes at 3.6, 4.5, 5.8 and 8.0 \u03bcm [mag] with respective errors (from the SAGE project, for a description see Meixner et al. 2006).\n\n\n54\nAngular distance between X-ray and optical position [\u2032\u2032].\n\n\n55\nAngular distance between optical and near-IR position [\u2032\u2032].\n\n\n56\nNeutron star spin period [s] inferred from X-rays.\n\n\n57\nOrbital period [days] (see flags for origin).\n\n\n58\u221259\nMaximum and minimum X-ray flux [erg cm-2 s-1] when available in the 0.2\u221210 keV band. Fluxes in the SMC XMM-Newton catalogue of Sturm et al. (2013c) are given for the 0.2 to 4.5 keV band. To convert them into the 0.2\u221210 keV band, we multiplied them by a factor of 2.6 assuming a standard power law with photon index 0.9 (Haberl et al. 2008) and a column density of 1021 cm-2 (solar abundance). For Swift XRT count rates we used a flux conversion factor of 1.1 \u00d7  10-10 erg cm-2 cts-1\n\n\n60\nFlag for minimum flux: 1 for a non-detection with an upper limit; \u22121 when unknown; 0 for detection.\n\n\n61\u221262\nReferences for maximum and minimum X-ray flux.\n\n\n63\nX-ray variability factor (ratio of maximum to minimum flux).\n\n\n64\nEquivalent width of the H\u03b1 line [\u00c5] (minimum value if more than one measurement is available).\n\n\n65\nMaximum equivalent width of the H\u03b1 line [\u00c5].\n\n\n66\nReferences for the H\u03b1 measurements.\n\n\n67\nComments with key references.\n\n\n\n","Citation Text":["Massey (2002)"],"Functions Text":["Optical coordinates, right ascension and declination (epoch 2000.0) for the identified counterpart from Zaritsky et al. (2002), or \u2013 when not available there \u2013 from"],"Functions Label":["Uses"],"Citation Start End":[[1148,1161]],"Functions Start End":[[983,1147]]}
{"Identifier":"2019ApJ...875...90L__Velli_et_al._2015_Instance_1","Paragraph":"When energy flows from the interior of the Sun outward into the solar atmosphere, why is the Sun\u2019s outer atmosphere, the corona, much hotter than the inner atmosphere, the underlying chromosphere and photosphere? This is the long-standing problem of the coronal heating, which is one of the eight key mysteries in modern astronomy (Kerr 2012). For about 80 yr since the discovery of the extremely hot corona around the late 1930s (Grotian 1939; Edlen 1945), people have worked hard on addressing this issue, and great advances have been made in observation and theoretical studies (Parnell & De Moortel 2012; Amari et al. 2015; Arregui 2015; Cargill et al. 2015; De Moortel & Browning 2015; Jess et al. 2015; Klimchuk 2015; Longcope & Tarr 2015; Peter 2015; Schmelz & Winebarger 2015; Velli et al. 2015; Wilmot-Smith 2015). Especially during recent decades, high-resolution observations of solar super-fine structures indicate that small spicules, minor hot jets along small-scale magnetic channels from the low atmosphere upwards to the corona, petty tornados and cyclones, and small explosive phenomena such as mini-filament eruptions and micro- and nano-flares\u2014all of these small-scale magnetic activities contribute greatly to coronal heating (De Pontieu et al. 2011; 2018; Zhang & Liu 2011; Parnell & De Moortel 2012; Klimchuk 2015; Peter 2015; Schmelz & Winebarger 2015; Henriques et al. 2016; Li et al. 2018a). Additionally, contributions of MHD waves to heating the corona have been observationally illustrated (van Ballegooijen et al. 2011; Jess et al. 2015; Kubo et al. 2016; Morton et al. 2016; Soler et al. 2017; Morgan & Hutton 2018). Meanwhile, with the progress of observational studies, two groups of theoretical models, magnetic reconnection models and magnetohydrodynamic wave models, have traditionally attempted to explain coronal heating, but so far no models can address the key mystery perfectly (van Ballegooijen et al. 2011; Arregui 2015; Cargill et al. 2015; Peter 2015; Velli et al. 2015; Wilmot-Smith 2015). Maybe we do not need to intentionally take to heart such the classical dichotomy, because waves and reconnections may interact with each other (De Moortel & Browning 2015; Velli et al. 2015). Additionally, statistical studies may look at coronal heating from a comprehensive perspective. Li et al. (2018b) found that the long-term variation of the heated corona, which is represented by coronal spectral irradiances, and that of small-scale magnetic activity are in lockstep, indicating that the corona should statistically be effectively heated by small-scale magnetic activity. Observational and theoretical model studies through heating channels and modes, and statistical studies by means of heating effect (performance of the heated corona), both suggest that coronal heating originates from small-scale magnetic activity.","Citation Text":["Velli et al. 2015"],"Functions Text":["For about 80 yr since the discovery of the extremely hot corona around the late 1930s (Grotian 1939; Edlen 1945), people have worked hard on addressing this issue, and great advances have been made in observation and theoretical studies"],"Functions Label":["Background"],"Citation Start End":[[785,802]],"Functions Start End":[[344,580]]}
{"Identifier":"2021AandA...653A..85S__However,_Genovali_et_al._(2014)_Instance_1","Paragraph":"Figure 6 shows the orbital eccentricities as a function of [M\/H] for the metal-rich disc sample. The solid lines correspond to the required eccentricity (see Eq. (2)) for different values of ISM radial metallicity gradients: \u22120.10 dex kpc\u22121 (black), \u22120.07 dex kpc\u22121 (our measured gradient for young stars in Table 1; see also Minchev et al. 2018, red), \u22120.04 dex kpc\u22121 (orange), and \u22120.06 dex kpc\u22121 (Cepheids analysis from Genovali et al. 2014, green). For the three first cases, we assumed ISM[M\/H](R\u2299) = 0.0 to estimate Rbirth from the stellar metallicity. However, Genovali et al. (2014) have their own zero point, defined as: [Fe\/H]\u2004=\u2004\u22120.06\u2005*\u2005Rg\u2005+\u20050.57, with a clear shift in the relation compared to the other ones assumed in this work. The impact of the ISM gradient value and the zero-point assumption on the derived Rbirth, and therefore on the required eccentricity to reach the solar vicinity without the need for churning, is clearly observed. As described in Hayden et al. (2020), given the measured [M\/H] and eccentricity, stars lying to the left are able to reach the solar neighbourhood through blurring, while the stars to the right of the line are possible candidates to have migrated through churning. This is the case for most of the SMR stars (70% of the SMR stars lie below the line that corresponds to the Cepheids analysis); they are therefore likely to have been brought to the solar neighbourhood by churning, which is in close agreement with previous studies (e.g., Kordopatis et al. 2015a; Wojno et al. 2016). However, it is worth noting that the observed metallicity distribution function in Fig. 2 peaks around 0.2 dex, which is higher than previous reported solar vicinity MDFs (see e.g., Fuhrmann et al. 2017). A possible ignored bias towards more metal-rich objects in the sample selection could be pulling the percentage of possible migrators to higher values. Among the entire distribution, our churned candidates with [M\/H] > \u2005+\u20050.1 comprise around 17% of the sample. If we constrain the number of migrators to only stars with [M\/H] > \u2005+\u20050.25, the global percentage decreases to 8% of the sample.","Citation Text":["Genovali et al. 2014"],"Functions Text":["Cepheids analysis from","green"],"Functions Label":["Uses","Uses"],"Citation Start End":[[423,443]],"Functions Start End":[[400,422],[445,450]]}
{"Identifier":"2016ApJ...830...28P__Zank_et_al._2015_Instance_1","Paragraph":"Our basic model for acceleration is shown in Figure 1. The left panel shows a cartoon of the model for acceleration, transport, and radiation at the X-reconnection site in the corona. Acceleration can be by the second order Fermi (or stochastic) acceleration by turbulence (see, e.g., Petrosian & Liu 2004), by a standing shock produced by the downflow from the X-reconnection site (see, e.g., Guo & Giacalone 2012), or in the merging of islands shown to arise in PIC simulations during reconnection (Drake et al. 2006, 2013; Le et al. 2012; Oka et al. 2010). The latter model has been invoked as a possible mechanism in the downstream of the CME shock (Le Roux et al. 2015; Zank et al. 2015). The right panel shows a similar cartoon for acceleration in the environment of the CME. The (red) rectangles show the cross section of the box within which particles are accelerated. For our purposes here the relevant parameters of the acceleration and transport of particles are the momentum (p) and pitch angle cosine (\u03bc) diffusion rates \n\n\n\n\n\n and \n\n\n\n\n\n, and direct energy gain \n\n\n\n\n\n and loss \n\n\n\n\n\n rates, where E is the particle energy.1\n\n1\nHere and in what follows we neglect the effects of the third diffusion coefficient \n\n\n\n\n\n, which are generally small (see Schlickeiser 1989; Petrosian & Liu 2004).\n For an isotropic pitch angle distribution the evolution of the volume-integrated energy spectrum, \n\n\n\n\n\n, is governed by the so-called leaky-box model kinetic equation (see Petrosian 2012)\n1\n\n\n\n\n\nwhere \n\n\n\n\n\n is a source term, \n\n\n\n\n\n, and the direct acceleration coefficients for stochastic and shock accelerations are given as\n2\n\n\n\n\n\nHere \n\n\n\n\n\n is the shock velocity, \n\n\n\n\n\n is the coefficient of spatial diffusion (along the field lines), \n\n\n\n\n\n, \n\n\n\n\n\n is the Lorentz factor, and \u03b6 depends on the shock compression ratio and other factors (see, e.g., Steinacker et al. 1988). The energy-loss rate at low energies (for both electrons and ions) is dominated by Coulomb collisions with background particles (mainly electrons).2\n\n2\nCoulomb collisions also cause pitch angle scattering, and therefore spatial diffusion along the field lines, with a rate that is comparable to the energy-loss rate at nonrelativistic energies but decreases as \n\n\n\n\n\n. They also cause energy diffusion, which is negligible in the cold target case (\n\n\n\n\n\n), but can be comparable to the energy-loss rate for a hot target as E approaches kT (see Petrosian & Kang 2015), where T is the background plasma temperature.\n At higher energies (not relevant for the discussion here) inelastic interactions (synchrotron, inverse Compton, and bremsstrahlung) for electrons and (nuclear line excitation, neutron and pion productions) for ions become important.","Citation Text":["Zank et al. 2015"],"Functions Text":["The latter model has been invoked as a possible mechanism in the downstream of the CME shock"],"Functions Label":["Uses"],"Citation Start End":[[675,691]],"Functions Start End":[[560,652]]}
{"Identifier":"2017ApJ...850..197P__Langer_2012_Instance_1","Paragraph":"In order to explode as an ECSN, several ingredients need to be in place. Nomoto (1984) argued that stars with helium cores between 2.0 and \n\n\n\n\n\n (which corresponds roughly to initial masses between 8 and \n\n\n\n\n\n) would explode as an ECSN. His models, however, did not develop a second dredge-up, which can significantly reduce the mass of the helium core and diminishes the predictive power of this criterion (see Podsiadlowski et al. 2004; Poelarends et al. 2008). Since then, several authors (Siess 2007; Poelarends et al. 2008; Doherty et al. 2010, 2015; Jones et al. 2013) have established precise initial mass ranges for ECSNe to occur, although these mass ranges are highly sensitive to the adopted convection criteria, overshooting, and mass-loss prescriptions (Poelarends et al. 2008; Doherty et al. 2010; Langer 2012). Jones et al. (2013) produced several detailed models, computed all the way to electron captures on 24Mg and 20Ne and found that CO cores with masses over \n\n\n\n\n\n are able to reach densities high enough for this to occur (\n\n\n\n\n\n g cm\u22123). If the CO core is massive enough, neon will ignite off-center (Jones et al. 2013; Schwab et al. 2016), but Jones et al. (2014) also found that the upper boundary for ECSNe is affected by uncertainties regarding the progression or stalling of the neon flame. There seems to be consensus, however, that the mass of the CO core is a reliable indicator for the final fate of stars in this mass range. How this translates into the initial mass of the star depends on the adopted convection criterion, with the Schwarzschild criterion producing more massive cores than the Ledoux criterion for the same initial mass. Inclusion of overshooting will also lead to larger cores (see Siess 2007). However, although the initial mass range for ECSNe is therefore quite sensitive to the adopted convection criteria, this does not seem to be the case for the final MCO, as most authors find similar values for MCO at which neon ignites, even though they treat convection differently. In the context of binary evolution, this provides an additional reason to adopt MCO as an indicator for whether the star explodes as an ECSN or not, as the CO core is generally not eroded by mass transfer, in contrast to the He core.","Citation Text":["Langer 2012"],"Functions Text":["Since then, several authors","have established precise initial mass ranges for ECSNe to occur, although these mass ranges are highly sensitive to the adopted convection criteria, overshooting, and mass-loss prescriptions"],"Functions Label":["Background","Background"],"Citation Start End":[[814,825]],"Functions Start End":[[466,493],[577,767]]}
{"Identifier":"2022AandA...666A.112L__Cormier_et_al._2015_Instance_2","Paragraph":"Local dwarf galaxies were the focus of large Herschel and Spitzer surveys (e.g., The Dwarf Galaxy Survey, DGS; Madden et al. 2006). Studies on both resolved and integrated-galaxy scales have highlighted some distinctively unique observational signatures of star-forming low-metallicity dwarf galaxies. A non-linear relation of the dust-to-gas mass (D\/G) with metallicity is observed, with extremely low dust masses observed for the lowest metallicity galaxies (Z \u2264 0.1 Z\u2299; Herrera-Camus et al. 2012; Fisher et al. 2014; R\u00e9my-Ruyer et al. 2015; Galliano et al. 2018, 2021; Cigan et al. 2021). Furthermore, the hard radiation fields in star-forming dwarf galaxies, along with their lower dust abundance, result in extended ionized gas regions prominent on global galaxy scales (Hunter et al. 2011; Cormier et al. 2015, 2019). The consequence is often a largely photodissociated molecular phase, existing in clumps which are difficult to observe with the usual molecular gas tracer, CO (1-0) (Cormier et al. 2014; Hunt et al. 2015; Accurso et al. 2017b), beckoning the presence of a CO-dark molecular gas phase (Grenier et al. 2005; R\u00f6llig et al. 2006; Wolfire et al. 2010; Glover & Clark 2012; Bolatto et al. 2013; Accurso et al. 2017a; Madden et al. 2020). Other emission lines, however, such as the far-infrared [C ii]\u03bbl58 \u00b5m line, are strikingly enhanced on global scales in dwarf galaxies (e.g., Cormier et al. 2015, 2019; Cigan et al. 2016; Lebouteiller et al. 2017; Jameson et al. 2018), making the [C ii]\u03bbl58 \u00b5m line a potential tool for tracing star formation activity (Malhotra et al. 2001; Papadopoulos et al. 2007; Pineda et al. 2014; De Looze et al. 2014; Olsen et al. 2015; Herrera-Camus et al. 2015, Herrera-Camus et al. 2018; Carniani et al. 2018; Matthee et al. 2019; Izumi et al. 2021; Fujimoto et al. 2021) and a tracer of the total H2 in galaxies, near and far (Poglitsch et al. 1995; Wolfire et al. 2010; Pineda et al. 2013; Nordon & Sternberg 2016; Fahrion et al. 2017; \nAccurso et al. 2017b; Zanella et al. 2018; Madden et al. 2020; Schaerer et al. 2020; Tacconi et al. 2020).","Citation Text":["Cormier et al. 2015"],"Functions Text":["Other emission lines, however, such as the far-infrared [C ii]\u03bbl58 \u00b5m line, are strikingly enhanced on global scales in dwarf galaxies (e.g.,"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1398,1417]],"Functions Start End":[[1256,1397]]}
{"Identifier":"2022AandA...660A.135V__Nissen_2016_Instance_1","Paragraph":"The first study, to our knowledge, to notice the net increase in the abundance of slow (s) neutron capture elements in young stellar populations is D\u2019Orazi et al. (2009), in which the abundance of barium in young star clusters was seen to be higher than in the older ones. Maiorca et al. (2011, 2012) added a few more elements with important s-process contributions (yttrium, zirconium, lanthanum, and cerium), confirming the increasing trend towards younger ages. Subsequently, a number of works have attempted to both clarify the origin of this increase (see, e.g., Bisterzo et al. 2014; Mishenina et al. 2015; Trippella et al. 2016; Magrini et al. 2018; Spina et al. 2018; Busso et al. 2021) and to use their abundances to estimate the ages of stars, often using neutron capture s-process elements in combination with other elements with opposite behaviours, such as \u03b1 elements \u2013 that we indicate as chemical clocks \u2013 and thus maximising the dependence of the relationship with age (see, e.g., Tucci Maia et al. 2016; Nissen 2016; Feltzing et al. 2017; Fuhrmann et al. 2017; Slumstrup et al. 2017; Titarenko et al. 2019). Once the existence of a relationship between age and chemical clocks was established (see, e.g., Spina et al. 2016; Delgado Mena et al. 2019; Jofr\u00e9 et al. 2020), the next steps were the following: (i) to clarify the applicability of these relationships with luminosity class (dwarf or giant) (see, e.g., Tucci Maia et al. 2016; Slumstrup et al. 2017; Casamiquela et al. 2021), metallicity (see, e.g., Feltzing et al. 2017; Delgado Mena et al. 2019; Casali et al. 2020), and population type (thin disc, thick disc, halo) (see, e.g., Titarenko et al. 2019; Nissen et al. 2020; Tautvai\u0161ien\u0117 et al. 2021), or even in dwarf galaxies (Sk\u00falad\u00f3ttir et al. 2019; ii) to calibrate them with a sample of stars with reliable age determination, which are usually open star clusters (OCs), solar twins, or targets with asteroseismic observations. Finally, it is essential to understand whether these relationships are valid throughout the Galactic disc, or whether they are necessarily position-dependent. For the first time, Casali et al. (2020) applied the relations derived from a large sample of solar-like stars located in the solar neighbourhood and noted that they fail to reproduce the ages of star clusters in the inner disc. They concluded that the relationship between age and chemical clocks is not universal and that it varies with galactocentric position. Later, Magrini et al. (2021b) suggested that the differences in the relationships between age and chemical clocks in different parts of the Galactic disc are due to the strong dependence on the metallicity of the yields of low-mass stars, which produce s-process elements during the final stages of their evolution. Casamiquela et al. (2021) used red clump stars in open clusters to investigate the age dependence of several abundance ratios, including those that contain s-process and \u03b1 elements. They found that the relationship between [Y\/Mg] and ages outlined by open clusters is similar to the one found using solar twins in the solar neighbourhood. They also found that the abundance ratios involving Ba are those with the highest correlation with age. However, they also note that as one moves away from the solar neighbourhood, the dispersion increases and is in agreement with the findings of Casali et al. (2020), which attributed this to the spatial variation of the star formation history along the galactocentric radius.","Citation Text":["Nissen 2016"],"Functions Text":["Subsequently, a number of works have","and to use their abundances to estimate the ages of stars, often using neutron capture s-process elements in combination with other elements with opposite behaviours, such as \u03b1 elements \u2013 that we indicate as chemical clocks \u2013 and thus maximising the dependence of the relationship with age (see, e.g.,"],"Functions Label":["Background","Background"],"Citation Start End":[[1021,1032]],"Functions Start End":[[465,501],[695,996]]}
{"Identifier":"2018AandA...616A..43S__Breeveld_&_Puchnarewicz_1998_Instance_1","Paragraph":"Within the same scenario, Smith et al. (2004, 2005) have theoretically predicted that, when observed in polarized light, NLS1s should have a low polarization fraction, the polarization angle should change monotonically from one wing of the BELs to the other, and the BELs should be significantly broader than those observed in direct light (Smith et al. 2004, 2005). We assume that polarized BELs are scattered into the line of sight by material that is close to coplanar with the BLR (i.e. the torus). This prediction has been searched for in RQ-NLS1s, but many objects have not been detected in polarized light, and none of the detected show any significant broadening of the BELs (e.g. Goodrich 1989; Breeveld & Puchnarewicz 1998; Kay et al. 1999). On the other hand, only one RL-NLS1 (PKS 2004-447) has been the object of a spectropolarimetry study (Baldi et al. 2016), and the three predicted features have been observed that have a FWHM in the polarized H\u03b1 of \u2248 9000 km s\u22121. We propose for the first time a simple scenario that can resolve such apparent tension between the two classes of NLS1s. We speculate that in those NLS1s that are intrinsically without a jet (a fraction of the RQ-NLS1) a significant amount of polarized light could be due to polar scattering of the BELs, by material above the BLR, as observed for type II AGN (Antonucci 1983; Miller & Antonucci 1983; Antonucci & Miller 1985). A polar-scattered component could easily overwhelm the planar-scattering contribution, predicted to result in low polarization fractions for face-on objects because of the axi-symmetry of the scattering material. In RL-NLS1s, on the contrary, the jet could evacuate the intervening material, resulting in a lower degree of polarization in which only the broadened component of the BEL is left. The evacuation of the polar material by a jet closely aligned withthe line of sight has been already proposed by Ghisellini & Sbarrato (2016) to account for the lack of high-z blazar parent population. This simple speculative scenario can be tested through spectropolarimetry of a larger sample of RL-NLS1s.","Citation Text":["Breeveld & Puchnarewicz 1998"],"Functions Text":["We assume that polarized BELs are scattered into the line of sight by material that is close to coplanar with the BLR (i.e. the torus). This prediction has been searched for in RQ-NLS1s, but many objects have not been detected in polarized light, and none of the detected show any significant broadening of the BELs (e.g."],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[704,732]],"Functions Start End":[[367,688]]}
{"Identifier":"2019MNRAS.490.4975R__Neilson_et_al._2012_Instance_1","Paragraph":"Fig. 6 displays the MLR for the investigated CCs whose intrinsic stellar parameters were derived from the best-fitting models listed in Table 1. These data are compared with the predicted canonical (no overshooting, no mass-loss) MLR (Bono et al. 2000; the solid lines) and with the relations obtained by increasing the zero-point of the canonical MLR by 0.25\u2009dex (the dashed lines) and 0.5\u2009dex (the dotted lines) to reproduce the effect of mild and full overshooting,5 respectively (see Chiosi et al. 1993; Bono et al. 1999, for details). Inclusion of mass-loss and\/or rotation would produce a similar increase in the Cepheid luminosity level at fixed mass (see Neilson et al. 2012, for details). As the light curves of $OGLE\\_CEP\\_LMC\\_2019$ are best reproduced adopting a different value of the helium content (see above), in Fig. 6 we also show the MLR for Y = 0.30, Z = 0.008 (the green lines). Note that this relation is slightly more luminous than those calculated for the standard LMC elemental composition (Y = 0.25, Z = 0.008). According to the location of the variables in the ML plane, the canonical MLR is not strictly satisfied, as the points are spread between the canonical and full overshooting predictions. Even if at this stage we cannot disentangle the role of overshooting, mass-loss, and rotation in producing the quoted excess luminosity, at fixed mass, the detected dispersion might indicate a combination of these different non-canonical phenomena. Indeed, if only overshooting were efficient, one would in principle expect the same amount of excess luminosity for all stellar masses (within small uncertainties). Rotation produces similar effects as overshooting because it implies a larger He burning core and a brighter luminosity at fixed mass (see e.g. Anderson et al. 2016) On the other hand, if the mass-loss process were efficient, this could be inferred from the predicted deviation of the best-fitting stellar mass from the value corresponding to the canonical MLR. Such a deviation is represented in Fig. 7 as a function of the canonical mass (top) and of the pulsation period (bottom) for the CCs in our sample. We note that the expected mass differences range from $0{{\\ \\rm per\\ cent}}$ to almost $\\sim 50{{\\ \\rm per\\ cent}}$ and are not clearly correlated with the pulsation period or the stellar mass.","Citation Text":["Neilson et al. 2012"],"Functions Text":["Inclusion of mass-loss and\/or rotation would produce a similar increase in the Cepheid luminosity level at fixed mass (see","for details)."],"Functions Label":["Uses","Uses"],"Citation Start End":[[663,682]],"Functions Start End":[[540,662],[684,697]]}
{"Identifier":"2017MNRAS.472..876O__Rodono_et_al._1985_Instance_1","Paragraph":"Leitzinger et al. (2014) briefly compared the mass estimates of stellar mass ejection events and the estimated X-ray energies of their associated flares with the solar scaling. We extend their comparison here. Specifically, we consider the events observed as blueshifted extra emissions in Balmer lines on the active M dwarfs AD\u2009Leo (Houdebine et al. 1990), AT\u2009Mic (Gunn et al. 1994), V374\u2009Peg (Vida et al. 2016) and the pre-main-sequence star DZ\u2009Cha (Guenther & Emerson 1997). Published estimates of their minimum masses are 7.7 \u00d7 1017, \u223c1015, \u223c1016 and 1.4 \u00d7 1018\u20267.8 \u00d7 1019\u2009g, respectively. Since no simultaneous X-ray observations of the associated flares are available, we estimate the corresponding X-ray flare energies based on the assumption that the published U-band energies are of comparable magnitude (cf. Hawley & Pettersen 1991). Doing so, we find EX \u2248 EU \u223c 2 \u00d7 1032\u2009erg for AD\u2009Leo (Rodono et al. 1985; Hawley & Pettersen 1991), 3 \u00d7 1031\u2009erg for AT\u2009Mic (Gunn et al. 1994) and 1.2\u20262.5 \u00d7 1035\u2009erg for DZ\u2009Cha (Guenther & Emerson 1997), respectively. The recently published event at V374\u2009Peg has an H\u03b3 flare energy of \u223c4 \u00d7 1030\u2009erg, which corresponds to \u223c1032\u2009erg in X-rays (Butler, Rodono & Foing 1988). We add two other events, which are, however, more uncertain in their interpretation as real CME events. Doyle et al. (1988) observed a strong increase in neutral hydrogen column density during a flare on the active M dwarf YZ\u2009CMi, which may be interpreted as a rising filament obscuring parts of the flare region. The estimated mass of this event is 3 \u00d7 1017\u2009g and the estimated X-ray flare energy EX \u2248 EU \u223c 8 \u00d7 1030\u2009erg, since the observed X-ray flare emission was likely partly absorbed by the neutral material (Doyle et al. 1988). The second one is a long-decay flare on the young, active M dwarf AU\u2009Mic, which was interpreted as an eruptive event with a CME mass in the order of 1020\u2009g by Cully et al. (1994). However, there is a different model of this strong flare (\u223c3 \u00d7 1035\u2009erg) including post-eruptive energy release, which does not involve a CME (Katsova, Drake & Livshits 1999). Other studies that observed stellar CME events (cf. Section 1) did not provide mass estimates and had to be excluded.","Citation Text":["Rodono et al. 1985"],"Functions Text":["Doing so, we find EX \u2248 EU \u223c 2 \u00d7 1032\u2009erg for AD\u2009Leo"],"Functions Label":["Uses"],"Citation Start End":[[897,915]],"Functions Start End":[[844,895]]}
{"Identifier":"2022MNRAS.510.5088B__Faisst_et_al._2017_Instance_2","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"],"Functions Text":["Evidence for offset dust continuum emission relative to the rest-UV has been identified in several high-redshift Lyman-break galaxies (LBGs;"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1368,1386]],"Functions Start End":[[1204,1344]]}
{"Identifier":"2017ApJ...844...17P__Xie_et_al._2004_Instance_1","Paragraph":"Observations by a single coronagraph allow us to obtain CME 2D images projected on the plane of the sky. Using the image sequences of CME expansions, we can determine CME 2D parameters such as speed and angular width. Such observations do have some limitations, mostly the projection effect. Kwon et al. (2015) showed that the apparent width of CMEs does not represent an accurate measure of the CME ejecta size. In previous studies, geometrical models based on the assumption of different cone shapes were developed to obtain CME 3D parameters from single spacecraft observations (Xie et al. 2004; Xue et al. 2005; Michalek 2006). Na et al. (2013) compared the 3D parameters of halo CMEs obtained from the cone shapes (elliptical-cone, ice-cream-cone, and asymmetric-cone models) and found that their 3D speeds are comparable to one another but their angular widths and inclination angles are quite different. We note that stereoscopic observations by twin spacecraft STEREO make it possible to obtain the CME 3D parameters using geometrical triangulation methods (Mierla et al. 2008, 2009; Liewer et al. 2009; Temmer et al. 2009; Liu et al. 2010a, 2010b) and a flux rope method (Chen et al. 2000; Thernisien et al. 2009). Recently, the STEREO CME Analysis Tool (StereoCAT), based on a triangulation method from multiple spacecraft observations, was developed by the Community Coordinated Modeling Center (CCMC) for real-time CME analysis carried out by the CCMC\/SWRC forecasting team. CME 3D parameters can be determined by selecting two coronagraphs from SOHO\/LASCO C3 and STEREO\/SECCHI COR2 images using the model. The triangulation algorithm and the performance of the CME measurements from StereoCAT are described in Sections 3.1 and 3.2 of Mays et al. (2015). Lee et al. (2015) showed, using 44 halo CMEs from 2010 to 2011, that the CME 3D parameters (speed, angular width, and source location) obtained from StereoCAT are similar to those obtained from the graduated cylindrical shape flux rope model (Thernisien et al. 2009; Thernisien 2011). Jang et al. (2016) showed, using 306 LASCO front-side halo CMEs from 2009 to 2013, that around 20% of the CME 2D speeds in the LASCO CME catalog are underestimated compared with the 3D ones obtained from StereoCAT.","Citation Text":["Xie et al. 2004"],"Functions Text":["In previous studies, geometrical models based on the assumption of different cone shapes were developed to obtain CME 3D parameters from single spacecraft observations"],"Functions Label":["Background"],"Citation Start End":[[582,597]],"Functions Start End":[[413,580]]}
{"Identifier":"2017ApJ...850...75S__Bugaev_et_al._2016_Instance_2","Paragraph":"A realistic EoS that is able to reproduce the properties of compact astrophysical objects has to fulfill several requirements. The possibility of including many particle species, which is known as multicomponent character, is of crucial importance for modeling the NS interiors, which in even the simplest treatment include neutrons, protons, and electrons, while more advanced descriptions have to account for the presence of hyperons (Schaffner-Bielich et al. 2002). Therefore, the grand canonical ensemble is the natural choice for the formulation of such an EoS. Another element of the realistic phenomenological hadronic EoS corresponds to the short-range repulsive interaction of the hard core nature between particles  (Andronic et al. 2006; Bugaev et al. 2016). Analysis of the particle yields produced in relativistic A+A collisions within statistical (thermal) models, i.e., the Hadron Resonance Gas (HRG) model (Bugaev et al. 2016), shows the importance of the particle hard core repulsion. In this approach every particle species is defined as a rigid sphere with a fixed radius estimated from experimental data analysis. These radii do not exceed 0.5 fm  (Andronic et al. 2017; Sagun et al. 2017a). Note that the hard core of hadrons in phenomenological EoSs is important in order to suppress thermal excitations of the hadronic spectrum and provide deconfinement of the color degrees of freedom expected at high temperatures\/densities  (Satz 2012). Another requirement to the phenomenological EoS is related to its causal behavior when the speed of sound cannot exceed the speed of light. At sufficiently high densities this condition is violated by the hard core repulsion. As was shown by Sagun et al. (2017a), introducing the induced surface tension (IST) of particles to the model with the hard core repulsion between an arbitrary number of hadron species makes the EoS significantly softer and extends its causality range up to 7.5 normal nuclear densities, where formation of the quark-gluon plasma is expected. The IST is the key element of this approach (Sagun et al. 2014), as it allows us to account for the hard core repulsion between constituents in the most accurate way, and to properly reproduce the virial expansion of the multicomponent EoS. Recently, the IST EoS was used to describe the experimental data of hadron multiplicities measured at AGS, SPS, RHIC, and LHC energies of nuclear collisions (Sagun et al. 2017b), as well as the nuclear matter properties (Sagun et al. 2014). In this work the focus is on the application of IST EoS to the study of NS properties.","Citation Text":["Bugaev et al. 2016"],"Functions Text":["Analysis of the particle yields produced in relativistic A+A collisions within statistical (thermal) models, i.e., the Hadron Resonance Gas (HRG) model","shows the importance of the particle hard core repulsion."],"Functions Label":["Background","Background"],"Citation Start End":[[923,941]],"Functions Start End":[[770,921],[944,1001]]}
{"Identifier":"2022MNRAS.509.3599T__Du_et_al._2015_Instance_3","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"],"Functions Text":["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","where \u211344 is the 5100 \u00c5 luminosity in units of $10^{44}\\, {\\rm erg\\, s^{-1}}$ and $M_7=M_{\\bullet }\/10^7\\, \\mathrm{M}_{\\odot }$."],"Functions Label":["Uses","Uses"],"Citation Start End":[[2039,2053]],"Functions Start End":[[1900,2037],[2056,2184]]}
{"Identifier":"2016AandA...592A..74S__Sobolewska_&_Papadakis_(2009)_Instance_2","Paragraph":"In Fig. B.1 we plot the soft X-ray light curves for our candidate highly variable AGN using available X-ray data taken by the satellite missions Einstein, ROSAT, XMM, Suzaku and Swift. The count rates were obtained from different archives including HEASARC, the XMM Science Archive, the Swift UKSSDC and from our own Swift XRT data analysis, and for upper limits the 1SXPS catalogue (Evans et al. 2014) and the XMM upper limit server2 were queried. The count rates of the different satellites were converted to fluxes between 0.2\u20132.0\u2009keV using PIMMS3 assuming a power law with a photon index of 1.7 as a spectral shape taking into account Galactic extinction as given by Willingale et al. (2013). Sobolewska & Papadakis (2009) found a positive correlation between flux and spectral slope for a sample of bright RXTE AGN in the 2\u201310\u2009keV band. This could affect the relative fluxes seen in our sample which are plotted in Fig. B.1. We have attempted to quantify this for the different detectors used in the creation of our light curves. The sample of Sobolewska & Papadakis (2009) showed spectral changes with observed power-law slope varying between 1.0 and 2.0 (see their Fig. 7). For a typical Galactic absorption of 3 \u00d7 1020 cm-2 the change from slope of 1.0 to 2.0 would alter our estimated fluxes by \u221214% (ROSAT), \u221213% (XMM-Newton), +7% (Swift-XRT), +76% (Suzaku), +25% (Einstein-IPC). The change is large for Suzaku observations since in this case we use the count rate between 2.0\u201310.0\u2009keV and extrapolate it to the soft band. All other satellites are sensitive in the soft band and hence the fluxes are less dependent upon the assumed spectral index. Six sources within our sample (XMMSL1\u2009J024916.6-041244, J034555.1-355959, J045740.0-503053, J051935.5-323928, J070841.3-493305, and J193439.3+490922) display variation in flux of a factor of ten or greater between at least one pair of XMM and Swift observations, on timescales of months to years. The ratio between the soft X-ray flux observed with Swift and that observed with XMM for the remaining sources is typically a factor of a few. We observed the two TDE candidates with XRT, and found that both had faded significantly, following expectations from previous and later fluxes and upper limits (Figs. 1p and h). ","Citation Text":["Sobolewska & Papadakis (2009)"],"Functions Text":["The sample of","showed spectral changes with observed power-law slope varying between 1.0 and 2.0 (see their Fig. 7). For a typical Galactic absorption of 3 \u00d7 1020 cm-2 the change from slope of 1.0 to 2.0 would alter our estimated fluxes by \u221214% (ROSAT), \u221213% (XMM-Newton), +7% (Swift-XRT), +76% (Suzaku), +25% (Einstein-IPC). The change is large for Suzaku observations since in this case we use the count rate between 2.0\u201310.0\u2009keV and extrapolate it to the soft band. All other satellites are sensitive in the soft band and hence the fluxes are less dependent upon the assumed spectral index."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1049,1078]],"Functions Start End":[[1035,1048],[1079,1657]]}
{"Identifier":"2015ApJ...805...44S__Tanaka_&_Haiman_2009_Instance_1","Paragraph":"Pair-instability supernovae (PI SNe) are the most energetic thermonuclear explosions known and can be detected near the edge of the observable universe. They have now been studied by several groups for their potential to probe the properties of the first stars and galaxies (Greif et al. 2008, 2010, 2011, 2012; Johnson et al. 2009, 2014; Turk et al. 2009; Stacy et al. 2010, 2012; Clark et al. 2011; Hosokawa et al. 2011; Smith et al. 2011; Glover 2013; Jeon et al. 2012; Susa 2013; Pawlik et al. 2011, 2013; Wise et al. 2012; Whalen 2013; Hirano et al. 2014). They can also shed light on the origins of supermassive black holes and early cosmological reionization and chemical enrichment (Mackey et al. 2003; Whalen et al. 2004, 2008a, 2010; Abel et al. 2007; Smith & Sigurdsson 2007; Wise & Abel 2008; Alvarez et al. 2009; Tanaka & Haiman 2009; Smith et al. 2009; Park & Ricotti 2011, 2012, 2013; Volonteri 2012; Agarwal et al. 2012; Ritter et al. 2012; Whalen & Fryer 2012; Choi et al. 2013; Latif et al. 2013a, 2013b; Reisswig et al. 2013; Schleicher et al. 2013; Johnson et al. 2014; Chiaki et al. 2013; Safranek-Shrader et al. 2014). For example, detections of both PI and core-collapse (CC) SNe at high redshift could be roughly binned by mass, thereby building up a simple Pop III IMF over time as enough events are discovered (de Souza et al. 2013, 2014). If the Pop III initial mass function (IMF) proves to be top heavy, this could account for the origins of SMBHs because enough 200\u2013300 \n\n\n\n\n\n  seed black holes might be formed at \n\n\n\n\n\n 20 for a few to reach 109\n\n\n\n\n\n  by \n\n\n\n\n\n 7. If not, alternatives for SMBH seeds must be found, such as BHs forming by direct collapse in atomically cooled halos at slightly later epochs, z \u223c 10\u201315 (Johnson et al. 2012, 2013b). PI SN candidates such as SN 2007bi (Gal-Yam et al. 2009; Kozyreva et al. 2014) and SN 2213\u20131745 (Cooke et al. 2012) have now been discovered at z = 0.126 and 2.05, respectively.","Citation Text":["Tanaka & Haiman 2009"],"Functions Text":["They can also shed light on the origins of supermassive black holes and early cosmological reionization and chemical enrichment"],"Functions Label":["Motivation"],"Citation Start End":[[826,846]],"Functions Start End":[[562,689]]}
{"Identifier":"2022MNRAS.511.4946N__Orr_&_Browne_1982_Instance_1","Paragraph":"In the standard AGN unification theory, the anisotropic radio emission produced by relativistic jets points to a description where properties like radio morphology and radio spectral index depend on orientation. Two widely used orientation indicators in the radio regime are the (a) core-to-lobe flux density ratio (R), which is the ratio of the flux densities of the core and lobe in radio wavelengths, and (b) radio-to-optical ratio of the quasar core (RI), which is the ratio of flux densities of the core in radio and optical wavelengths (See equation 5 and 6 for the exact definition of R and RI parameters). A high R-value suggests a small viewing angle to the jet axis, while a low R-value indicates more of an edge-on view to the quasar (Kapahi & Saikia 1982; Orr & Browne 1982; Morisawa & Takahara 1987; Morganti et al. 1997). Studies have shown that R correlates with core radio luminosity (Hardcastle & Worrall 2000) as well with the core optical luminosity (Kharb & Shastri 2004). Factors like age and environment are believed to introduce the scatter in the correlation involving the R parameter suggesting that R might not be the best measure for orientation. The R parameter depends on both the core\u2019s radio flux as well as the radio flux of the lobes arising from the quasar jets. These jets are suggested to be affected by environmental factors. Wills & Brotherton (1995) and Kharb, Lister & Cooper (2010) showed that the radio-to-optical ratio (RI) of the quasar core might be a better measure of studying the orientation. The RI parameter measures the core-boosting factor by normalizing the core\u2019s radio flux by the core\u2019s optical flux. Although with some scatter in the correlation, Kimball et al. (2011) demonstrated a strong correlation between R and RI parameter suggesting that the two parameters are good indicators of quasar orientation. The scatter in the correlation supports the idea that other factors like age, size, environment, and luminosity also influence these two measurements.","Citation Text":["Orr & Browne 1982"],"Functions Text":["A high R-value suggests a small viewing angle to the jet axis, while a low R-value indicates more of an edge-on view to the quasar"],"Functions Label":["Background"],"Citation Start End":[[768,785]],"Functions Start End":[[614,744]]}
{"Identifier":"2020AandA...639A..46B__\u0160tver\u00e1k_et_al._(2009)_Instance_5","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)"],"Functions Text":["These findings by",", 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."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1488,1509]],"Functions Start End":[[1470,1487],[1509,1749]]}
{"Identifier":"2021AandA...654L...5P__Beaug\u00e9_&_Nesvorn\u00fd_2013_Instance_1","Paragraph":"Armstrong et al. (2020) recently announced the discovery of TOI-849b, a planet having a size comparable to the one of Neptune (\n\n\n\n\nR\npl\n\n=\n3\n.\n\n45\n\n\u2212\n0.12\n\n\n+\n0.16\n\n\n\n\nR\n\u2295\n\n\n\n$ R_{\\mathrm{pl}} = 3.45^{+0.16}_{-0.12}\\,R_{\\oplus} $\n\n\n), but an anomalously larger mass (\n\n\n\n\nM\npl\n\n=\n40\n.\n\n8\n\n\u2212\n2.5\n\n\n+\n2.4\n\n\n\n\nM\n\u2295\n\n\n\n$ M_{\\mathrm{pl}} = 40.8^{+2.4}_{-2.5}\\,M_{\\oplus} $\n\n\n) and a density similar to the one of the Earth (\u03c1\u2004=\u20045.5\u2005\u00b1\u20050.8 g cm\u22123). This is the densest Neptune-sized planet discovered to date. Interior structure models suggest that for this kind of planet, any H\/He envelope would consist of no more than 3.9% of the total mass (Armstrong et al. 2020). TOI-849b orbits around a late G-type star, with an orbital period of P\u2004=\u200418.4 h. Its equilibrium temperature is Teq\u2004=\u20041800 K (for a Bond Albedo 0.3). With such properties, TOI-849b represents one of the few planets populating the hot Neptune desert, a region on the radius-orbital distance plane characterised by a surprising deficit of Neptune-sized planets on very short orbits (e.g. Lecavelier Des Etangs 2007; Beaug\u00e9 & Nesvorn\u00fd 2013; Mazeh et al. 2016). Growing evidence suggests that the evaporation of hot Neptunes due to stellar irradiation represents a major process in shaping the desert (Owen 2019). Some of the planets within the desert or at its lower-radius border could thus be the remnant cores of larger gas-rich progenitors that lost most of their atmosphere (e.g. Lecavelier des Etangs et al. 2004). Alternatively, orbital migration may have also played a part in sculpting the desert, with different classes of planets forming differently (e.g. Batygin et al. 2016), or following different dynamical tracks (e.g. Matsakos & K\u00f6nigl 2016). Even so, the origin of this key feature in the demographics of close-in planets remains debated (Zahnle & Catling 2017; Owen & Lai 2018), and investigating the past history of planets such as TOI-849b contributes to the global effort towards disentangling the important mechanisms at the root of the desert. Since the mass of the planet is larger than the threshold value for runaway gas accretion (roughly 10\u2005\u2212\u200520\u2006M\u2295, Mizuno et al. 1978; Rafikov 2006; Piso et al. 2015), TOI-849b might have been a gas giant before undergoing extreme mass loss via thermal self-disruption, collisions with other giant planets, or it was prevented from accreting gas because of gap openings in the protoplanetary disc, or because of late formation (Armstrong et al. 2020). In their work, Armstrong et al. (2020) found that their estimated photoevaporation rates cannot provide the mass-loss rates needed to remove a roughly 280\u2006M\u2295 envelope from a Jupiter-like gas giant. Nevertheless, in their estimations the planetary atmosphere is not self-consistently monitored, and neither is the luminosity emitted by the host star relative to its rotational history. This is critical, considering how recent works emphasise the interplay between all these elements in driving the evolution of close-in worlds (e.g. Owen & Alvarez 2016; Kubyshkina et al. 2018; King & Wheatley 2021; Pezzotti et al. 2021).","Citation Text":["Beaug\u00e9 & Nesvorn\u00fd 2013"],"Functions Text":["With such properties, TOI-849b represents one of the few planets populating the hot Neptune desert, a region on the radius-orbital distance plane characterised by a surprising deficit of Neptune-sized planets on very short orbits (e.g."],"Functions Label":["Background"],"Citation Start End":[[1077,1099]],"Functions Start End":[[813,1048]]}
{"Identifier":"2022AandA...657A..50G__Qian_&_Wasserburg_2003_Instance_1","Paragraph":"Carbon-enhanced metal-poor (CEMP) stars form an important class of metal-poor giants, sub-giants, and dwarfs, with a large fraction of them showing enhanced abundances of heavy elements (see Beers & Christlieb 2005; Frebel 2018 for a general review). Among the different types of CEMP stars, the CEMP-s stars are enriched with products of s-process nucleosynthesis, the CEMP-r stars are enriched with the products of r-process nucleosynthesis, and CEMP-r\/s stars are enriched with products of i-process nucleosynthesis. Understanding the diverse abundance patterns exhibited by different groups of CEMP stars that are believed to be associated with different formation mechanisms has been a challenge. In Goswami et al. (2021), we present a detailed analysis and discussion on the classification criteria of CEMP stars, as well as the formation scenarios of CEMP stars put forward by different authors (Cowan & Rose 1977; Hill et al. 2000; Qian & Wasserburg 2003; Cohen et al. 2003; Jonsell et al. 2006; Campbell & Lattanzio 2008; Campbell et al. 2010; Stancliffe et al. 2011; Herwig et al. 2011; Doherty et al. 2015; Abate et al. 2016; Jones et al. 2016; Banerjee et al. 2018; Clarkson et al. 2018; Denissenkov et al. 2017; C\u00f4t\u00e9 et al. 2018). In this paper, we report an extremely metal-poor carbon-enhanced star, HE 1005\u20131439, whose surface chemical composition is found to be enriched with both s-process and i-process nucleosynthesis that forms a new class of object with a distinct abundance pattern. The peculiar abundance pattern, observed for the first time in a CEMP star, was investigated based on a parametric-model-based analysis that revealed almost equal contributions from both the s-process and the i-process to its surface chemical composition. We examined various production mechanisms and formation scenarios for this object. A formation scenario involving effective proton ingestion episodes (PIEs) triggering i-process nucleosynthesis followed by s-process asymptotic giant branch (AGB) nucleosynthesis with limited third-dredge-up (TDU) episodes seems to be most promising for this type of object.","Citation Text":["Qian & Wasserburg 2003"],"Functions Text":["In Goswami et al. (2021), we present a detailed analysis and discussion on the classification criteria of CEMP stars, as well as the formation scenarios of CEMP stars put forward by different authors"],"Functions Label":["Background"],"Citation Start End":[[940,962]],"Functions Start End":[[702,901]]}
{"Identifier":"2022ApJ...935....7L__Takeuchi_&_Kono_2020_Instance_1","Paragraph":"Among the GRB empirical correlations, the Amati correlation is a very popular one, which connects the spectral peak energy in the GRB cosmological rest frame and the isotropic equivalent radiated energy (E\n\np\n \u2212 E\niso; Amati et al. 2002; Amati 2006a, 2006b; Amati et al. 2008, 2009; Amati & Della Valle 2013). Recently, we proposed an improved Amati correlation (Liu et al. 2022) by using the Gaussian copula, which is a powerful statistical tool capable of describing the dependence structures between multivariate random variables, and has been applied to various fields by the astronomical community (Benabed et al. 2009; Jiang et al. 2009; Koen 2009; Scherrer et al. 2010; Takeuchi 2010; Yuan et al. 2018; Qin et al. 2020; Takeuchi & Kono 2020). In Liu et al. (2022), by choosing the spatially flat \u039bCDM model with \u03a9m0 = 0.30 and H\n0 = 70 km s\u22121 Mpc\u22121 as the fiducial model, we utilize the low-redshift (z  1.4) GRB data to calibrate the standard and improved Amati correlations, and then extrapolate the results to the high-redshift GRB data to achieve the GRB Hubble diagram, where \u03a9m0 is the present dimensionless matter density parameter. Using these calibrated GRBs to constrain the flat \u039bCDM model, we found that the improved Amati correlation can give results well consistent with the fiducial model, while the standard one cannot. Thus, in Liu et al. (2022), the reliability of the improved Amati correlation was ascertained with a fiducial model, but its cosmological application was not carried out. In this work, we will fill this gap. In order to obtain the Hubble diagram of the latest A220 GRB samples (Khadka et al. 2021) model-independently, we use the Pantheon SN Ia data (Scolnic et al. 2018) to calibrate the standard and improved Amati correlations, and then use these calibrated GRB data to constrain the \u039bCDM and wCDM models. Besides the GRB data, the H(z) data set is also added in our analysis to obtain a tight constraint on model parameters. The rest of the paper is organized as follows: Section 2 introduces the improved Amati correlation briefly, and standardizes the GRB samples by using the method of the low-redshift calibration. Section 3 studies the constraints on the \u039bCDM and wCDM models from the GRB data and the GRB + H(z) data. Our conclusions are summarized in Section 4.","Citation Text":["Takeuchi & Kono 2020"],"Functions Text":["Recently, we proposed an improved Amati correlation","by using the Gaussian copula, which is a powerful statistical tool capable of describing the dependence structures between multivariate random variables, and has been applied to various fields by the astronomical community"],"Functions Label":["Background","Background"],"Citation Start End":[[727,747]],"Functions Start End":[[310,361],[380,602]]}
{"Identifier":"2021ApJ...919..140S__Bartos_et_al._2017_Instance_2","Paragraph":"Resonant dynamical friction may have applications beyond the relaxation of IMBHs examined in this paper. It may affect all objects in stellar clusters much more massive than the individual constituents of the disk, if present, including massive stars, stellar mass black holes (BHs), or the center of mass of massive binaries. Furthermore, it is also expected to operate in any type of disk with a high number of particles, including active galactic nucleus (AGN) accretion disks. Previously, it has been argued that stars and BHs crossing the disk on low-inclination orbits get captured by Chandrasekhar dynamical friction into the disk (Bartos et al. 2017; Panamarev et al. 2018; Tagawa et al. 2020). An interesting implication is that, if BHs settle into the disk, they interact dynamically and form BH\u2013BH binaries efficiently, and frequent dynamical interactions and gas effects drive the BHs to merger, producing gravitational waves (GWs) detectable by LIGO, VIRGO, and KAGRA (McKernan et al. 2014, 2018; Bartos et al. 2017; Leigh et al. 2018; Yang et al. 2019; Tagawa et al. 2020, 2021; Samsing et al. 2020). Mergers are also facilitated by Lidov\u2013Kozai oscillations in anisotropic systems (Heisler & Tremaine 1986; Petrovich & Antonini 2017; Hamilton & Rafikov 2019). The results in this paper show that resonant dynamical friction may accelerate the capture of objects in the accretion disks by a factor proportional to the SMBH mass over the local disk mass for large orbital inclinations. Pressure and viscosity in a gaseous disk do not inhibit the orbit-averaged torque from the IMBH, which leads to realignment and the warping of the disk (Bregman & Alexander 2012). Thus, RDF may efficiently catalyze the alignment of the orbital planes of BHs even in low-luminosity AGN or Seyfert galaxies with relatively small disk masses, which may not be possible for Chandrasekhar dynamical friction. In fact, this mechanism extends the scope of the \u201cAGN merger channel\u201d for GW source populations even beyond low-luminosity AGN and Seyfert galaxies, as it may organize BHs into disks also in nonactive galaxies with nuclear stellar disks.","Citation Text":["Bartos et al. 2017"],"Functions Text":["An interesting implication is that, if BHs settle into the disk, they interact dynamically and form BH\u2013BH binaries efficiently, and frequent dynamical interactions and gas effects drive the BHs to merger, producing gravitational waves (GWs) detectable by LIGO, VIRGO, and KAGRA"],"Functions Label":["Background"],"Citation Start End":[[1010,1028]],"Functions Start End":[[703,980]]}
{"Identifier":"2019AandA...631A.167D__George_et_al._2013_Instance_1","Paragraph":"The applicability of mid-IR and far-IR FSL as diagnostic tools of the ISM has received a boost thanks to the publication of samples of nearby galaxies observed with the infrared space observatory (ISO) and Herschel (e.g. Brauher et al. 2008; Farrah et al. 2013; Sargsyan et al. 2014; Kamenetzky et al. 2014; Cormier et al. 2015; Cigan et al. 2016; Herrera-Camus et al. 2016; Fern\u00e1ndez-Ontiveros et al. 2016; Zhao et al. 2016; D\u00edaz-Santos et al. 2017; Zhang et al. 2018). At high redshift (z\u2004\u2273\u20041), the far-IR FSL conveniently shift into the (sub)millimetre atmospheric windows. The most popular line is clearly [CII] 158 \u03bcm, followed by the [CI] 370, 609 \u03bcm lines (e.g. Walter et al. 2011; Bothwell et al. 2017). Both single dish submillimetre telescopes and interferometers have also detected the [NII] 122 \u03bcm and 205 \u03bcm lines at high redshift (Ferkinhoff et al. 2011, 2015; Combes et al. 2012; Nagao et al. 2012; Decarli et al. 2012, 2014; B\u00e9thermin et al. 2016; Pavesi et al. 2016; Lu et al. 2017; Tadaki et al. 2019; Novak et al. 2019). After a slow start, the [OIII] 88 \u03bcm line is quickly becoming a popular line to confirm redshifts of galaxies in the epoch of reionization (z\u2004\u2273\u20046), where it shifts into the submillimetre atmospheric windows below 500 GHz (Ferkinhoff et al. 2010; Inoue et al. 2016; Carniani et al. 2017; Marrone et al. 2018; Vishwas et al. 2018; Hashimoto et al. 2018, 2019; Walter et al. 2018; Tamura et al. 2019; Tadaki et al. 2019; Novak et al. 2019). Deep Herschel\/SPIRE spectroscopy has also revealed a number of FSL detections, either in individual objects (Valtchanov et al. 2011; Coppin et al. 2012; George et al. 2013; Uzgil et al. 2016; Rigopoulou et al. 2018; Zhang et al. 2018), or in stacked spectra (Wardlow et al. 2017; Wilson et al. 2017; Zhang et al. 2018). These include the only detections of the [OI] 63 \u03bcm line at high redshift reported thus far. This [OI]3P2\u22123P1 line is arguably the best tracer for the star-forming gas, as it traces the very dense (\n\n\n\nn\ncrit\nH\n\n\n$ n^{\\mathrm{H}}_{\\mathrm{crit}} $\n\n\n = 5\u00d7105 cm\u22123) neutral gas (one caveat being the frequeny presence of self-absorption observed in local ultra-luminous IR galaxies, Rosenberg et al. 2015). Like the [CII] line, the [OI] 63 \u03bcm line also shows a \u201cdeficit\u201d in the most luminous far-IR sources, though with a higher scatter (Graci\u00e1-Carpio et al. 2011; Cormier et al. 2015; D\u00edaz-Santos et al. 2017). Surprisingly, this bright FSL has not been frequently observed with ALMA, probably because it is only observable in the highest frequency bands. At least as surprising is that the fainter, but more accessible [OI]3P1\u22123P0 transition at \u03bbrest = 145 \u03bcm (\n\n\n\nn\ncrit\nH\n\n\n$ n^{H}_{\\mathrm{crit}} $\n\n\n = 9.5\u00d7104 cm\u22123) has only recently been detected at high redshifts (Novak et al. 2019 report a tentative detection in a z = 7.5 quasar). Also in nearby galaxies, this [OI] 145 \u03bcm line has not been observed very frequently as in most cases, it is fainter than the nearby [CII] 158 \u03bcm line. After initial detections with ISO (Malhotra et al. 2001; Brauher et al. 2008), Herschel has now detected [OI] 145 \u03bcm in significant samples of nearby galaxies (Spinoglio et al. 2015; Cormier et al. 2015; Fern\u00e1ndez-Ontiveros et al. 2016; Herrera-Camus et al. 2018), and recently in a z = 6.5 lensed quasar (Yang et al. 2019).","Citation Text":["George et al. 2013"],"Functions Text":["Deep Herschel\/SPIRE spectroscopy has also revealed a number of FSL detections, either in individual objects"],"Functions Label":["Background"],"Citation Start End":[[1631,1649]],"Functions Start End":[[1478,1585]]}
{"Identifier":"2019ApJ...886...34F__Sahijpal_&_Goswami_1998_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":["Sahijpal & Goswami 1998"],"Functions Text":["If 26Al was introduced into the solar system at the earliest stage of the disk evolution (e.g.,","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."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1148,1171]],"Functions Start End":[[1052,1147],[1174,1375]]}
{"Identifier":"2021MNRAS.504.5575K__Cohen_et_al._2018_Instance_1","Paragraph":"First detected as a radio source during the Vermilion River Observatory Sky Survey (Dickel et al. 1967) and the Ohio Sky Survey (Kraus 1977), OJ 287 has been studied extensively in the radio regime. Its relativistic jet is pointing at us with an average viewing angle of \u223c2\u00b0 (Jorstad et al. 2005; Agudo et al. 2012) and shows remarkable short-time-scale variability interpreted as a turbulent injection process and\/or a clumpy accretion disc structure (Agudo et al. 2012) or as a binary-induced wobble (Valtonen & Pihajoki 2013; Dey et al. 2021). The inner jet is the source of highly variable \u03b3-rays (e.g. Agudo et al. 2011; Hodgson et al. 2017) and displays strong and variable radio polarization (e.g. Aller et al. 2014; Cohen et al. 2018; Myserlis et al. 2018). Though only faintly detected in the very-high-energy (VHE) band (>100 GeV; Mukherjee et al. 2017; O\u2019Brien 2017), OJ 287 is a well-known X-ray emitter and has been detected with most major X-ray observatories, including Einstein (Madejski & Schwartz 1988), EXOSAT (Sambruna et al. 1994), ROSAT (Comastri, Molendi & Ghisellini 1995), BeppoSAX (Massaro et al. 2003), ASCA (Idesawa et al. 1997), the Neil Gehrels Swift observatory (Swift hereafter; Massaro et al. 2008), Suzaku (Seta et al. 2009), XMM\u2013Newton (Ciprini et al. 2007), and most recently with NuSTAR (Komossa et al. 2020a). These observations established OJ 287 as a bright and variable X-ray source, and allowed single-component X-ray spectral fits. Its (0.5\u201310) keV X-ray spectrum was interpreted as a mix of synchrotron and inverse-Compton (IC) emission, with the former more variable (Urry et al. 1996). OJ 287 is found most of the time in a rather flat spectral state with a photon index \u0393x \u2248 1.5\u20131.9, though \u0393x is as steep as 2.6 during one ROSAT observation (Urry et al. 1996) in the band (0.1\u20132.4) keV. Its steepest state, with an equivalent \u0393x = 2.8, (better fit with a logarithmic parabolic model) was detected with XMM\u2013Newton in the band (0.3\u201310) keV, which caught the source right at the peak of one of the brightest X-ray outbursts measured so far (Komossa et al. 2020a). Imaging with the Chandra X-ray observatory has revealed a long, curved X-ray jet consisting of multiple knots and extending out to 20\u2009arcsec or a de-projected scale of >1 Mpc, and bright central emission (Marscher & Jorstad 2011).","Citation Text":["Cohen et al. 2018"],"Functions Text":["The inner jet","and displays strong and variable radio polarization (e.g."],"Functions Label":["Background","Background"],"Citation Start End":[[724,741]],"Functions Start End":[[547,560],[647,704]]}
{"Identifier":"2017ApJ...849...52N__Merloni_&_Fabian_2002_Instance_1","Paragraph":"The bolometric luminosity of 3C 84 is about 0.4% of the Eddington luminosity. Thus, the accretion flow of 3C 84 is likely to be a radiatively inefficient accretion flow (RIAF: Narayan & Yi 1995) rather than a standard disk (Shakura & Sunyaev 1973). However, we note that 3C 84 has a cold (\n\n\n\n\n\n\nT\n\n\ne\n\n\n\u223c\n\n\n10\n\n\n4\n\n\n\n\n K) disk-like accretion flow, as identified by FFA of the emission from the counter jet in the parsec scale (Walker et al. 2000) and inhomogeneous gas distribution around the black hole (Fujita et al. 2016). A number of theoretical studies predicted that the accretion flow components of hot geometrically thick (RIAF-like) and cold geometrically thin can coexist in either horizontal or vertical stratification (e.g., Miller & Stone 2000; Merloni & Fabian 2002; Liu et al. 2007; Ho 2008; Liu & Taam 2013). The measured Faraday rotation can be caused by such an RIAF-like component. We thus estimate the accretion rate of the RIAF-like component using the measured RM. For a simplicity, we assume that the RIAF-like component is a quasi-spherical Bondi accretion flow with a power-law density profile. We can calculate the accretion rate, following the formulation as follows (Quataert & Gruzinov 2000; Marrone et al. 2006; Kuo et al. 2014).\n\n\n\n\n\n\n\n\n\nM\n\n\n\u02d9\n\n\n\n\n=\n\n\n1.3\n\u00d7\n\n\n10\n\n\n\u2212\n10\n\n\n\n\n\n[\n\n1\n\u2212\n\n\n\n(\n\n\n\nr\n\n\nout\n\n\n\n\/\n\n\n\nr\n\n\nin\n\n\n\n)\n\n\n\n\u2212\n(\n3\n\u03b2\n\u2212\n1\n)\n\n\/\n\n2\n\n\n\n]\n\n\n\n\u2212\n2\n\n\/\n\n3\n\n\n\n\n\n\n\n\n\u00d7\n\n\n\n\n\n\n\n\n\nM\n\n\nBH\n\n\n\n\n8.0\n\u00d7\n\n\n10\n\n\n8\n\n\n\n\nM\n\n\n\u2299\n\n\n\n\n\n\n\n\n\n4\n\n\/\n\n3\n\n\n\n\n\n\n\n\n\n2\n\n\n3\n\u03b2\n\u2212\n1\n\n\n\n\n\n\n\n\u2212\n2\n\n\/\n\n3\n\n\n\n\nr\n\n\nin\n\n\n7\n\n\/\n\n6\n\n\n\n\n\n\n\n\n\nRM\n\n\nrad\n\n\n\nm\n\n\n\u2212\n2\n\n\n\n\n\n\n\n\n\n2\n\n\/\n\n3\n\n\n.\n\n\n\n\n\nFor an inner effective radius \n\n\n\n\n\n\nr\n\n\nin\n\n\n\n\n of 1 pc (\n\n\n\n\n1.3\n\u00d7\n\n\n10\n\n\n4\n\n\n\n\nR\n\n\ns\n\n\n\n\n), where the hotspot is located, the observed RM implies an accretion rate of \n\n\n\n\n\u223c\n4.3\n\u00d7\n\n\n10\n\n\n\u2212\n2\n\n\n\n\nM\n\n\n\u2299\n\n\n\n\n yr\u22121 and \n\n\n\n\n\u223c\n8.6\n\u00d7\n\n\n10\n\n\n\u2212\n2\n\n\n\n\nM\n\n\n\u2299\n\n\n\n\n yr\u22121 for \n\n\n\n\n\u03b2\n=\n0.5\n\n\n and \n\n\n\n\n\u03b2\n=\n1.5\n\n\n, which are corresponding to convection-dominated accretion flow (CDAF: Narayan et al. 2000; Quataert & Gruzinov 2000) and advection-dominated accretion flow (ADAF: Ichimaru 1977; Narayan & Yi 1995), respectively. Here we assumed the outer effective radius \n\n\n\n\n\n\nr\n\n\nout\n\n\n\n\n of \n\n\n\n\n\n\n10\n\n\n5\n\n\n\n\nR\n\n\ns\n\n\n\n\n (\u223c8 pc), which is approximately the same with the Bondi radius of 8.6 pc (Fujita et al. 2016). The derived accretion rate is roughly consistent with that estimated from the bolometric luminosity with a black hole mass of \n\n\n\n\n8\n\u00d7\n\n\n10\n\n\n8\n\n\n\n\nM\n\n\n\u2299\n\n\n\n\n and a radiative efficiency of 10% (\n\n\n\n\n\n\nM\n\n\n\u02d9\n\n\n\u223c\n\n\nL\n\n\nbol\n\n\n\n\/\n\n(\n0.1\n\n\nc\n\n\n2\n\n\n)\n\u2243\n7.1\n\n\u00d7\n\n\n10\n\n\n\u2212\n2\n\n\n\n\nM\n\n\n\u2299\n\n\n\n\n yr\u22121).","Citation Text":["Merloni & Fabian 2002"],"Functions Text":["A number of theoretical studies predicted that the accretion flow components of hot geometrically thick (RIAF-like) and cold geometrically thin can coexist in either horizontal or vertical stratification (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[759,780]],"Functions Start End":[[527,737]]}
{"Identifier":"2022MNRAS.513.1623C__Koyama,_Taruya_&_Hiramatsu_2009_Instance_1","Paragraph":"Two-point statistics are central to many of the leading cosmological analyses of the large-scale structure searching for deviations from \u039bCDM (Simpson et al. 2013; Song et al. 2015; Amon et al. 2018; Abbott et al. 2019a; Chudaykin, Dolgikh & Ivanov 2021; Lee et al. 2022; Muir et al. 2021; Tr\u00f6ster et al. 2021; Vazsonyi et al. 2021), and a great deal of effort has gone into accurately modelling the non-linear matter power spectrum in modified gravity and dark energy cosmologies \u2013 a theoretical ingredient essential to extract the cosmological information locked in small scales (e.g. Koyama, Taruya & Hiramatsu 2009; Brax & Valageas 2012; Takahashi et al. 2012; Heitmann et al. 2014; Zhao 2014; Casarini et al. 2016; Mead et al. 2016; Cusin, Lewandowski & Vernizzi 2018; Cataneo et al. 2019; Winther et al. 2019; Euclid Collaboration 2021; Ramachandra et al. 2021). However, non-linear gravitational clustering converts the nearly Gaussian initial density field (Planck Collaboration VI 2020) to a late-time density field with significant non-Gaussian features that these standard analyses are unable to access (Bernardeau et al. 2002). Non-Gaussian statistics, such as the bispectrum (Brax & Valageas 2012; Munshi 2017; Yamauchi, Yokoyama & Tashiro 2017; Crisostomi, Lewandowski & Vernizzi 2020; Bose et al. 2020b), higher order weak lensing spectra (Munshi & McEwen 2020), the halo mass function (Lam & Li 2012; Cataneo et al. 2016; Hagstotz et al. 2019; McClintock et al. 2019; Bocquet et al. 2020), the void size function (Perico et al. 2019; Verza et al. 2019; Contarini et al. 2021) and Minkowski functionals (Kratochvil et al. 2012; Fang, Li & Zhao 2017), respond strongly to modified gravity and dark energy through the induced changes in the higher moments of the cosmic density field, and their remarkable complementarity to traditional two-point functions leads to tighter joint constraints on the extra non-standard parameters (Shirasaki et al. 2017; Peel et al. 2018; Sahl\u00e9n 2019; Liu et al. 2021).","Citation Text":["Koyama, Taruya & Hiramatsu 2009"],"Functions Text":["Two-point statistics are central to many of the leading cosmological analyses of the large-scale structure searching for deviations from \u039bCDM","and a great deal of effort has gone into accurately modelling the non-linear matter power spectrum in modified gravity and dark energy cosmologies \u2013 a theoretical ingredient essential to extract the cosmological information locked in small scales (e.g."],"Functions Label":["Background","Background"],"Citation Start End":[[587,618]],"Functions Start End":[[0,141],[334,586]]}
{"Identifier":"2019MNRAS.482.5597A__Bergeron_&_Stasi\u0144ska_1986_Instance_1","Paragraph":"There have been many efforts in the past that have used the optical emission lines to understand the AGN and the host galaxy properties (e.g. Baldwin, Wampler & Burbidge 1981; Tadhunter et al. 1998). While broad and high ionization lines tend to arise from the immediate vicinity of the central massive black hole, low ionization lines arise further out (e.g. Richardson et al. 2014), tracing the host-galaxy dynamics, chemical composition and gas distribution. It is known that low ionization gases trace the H\u2009i\u2009distribution in the galactic discs (e.g. Lin & Zou 2001). Indeed, Mg ii is known as a good tracer of high column density H\u2009i\u2009gas in damped Lyman \u03b1 absorbers (e.g. Bergeron & Stasi\u0144ska 1986; P\u00e9roux et al. 2004). [O ii] \u03bb3727 is widely used to trace star formation in galaxy surveys, particularly for redshifts $z$ \u2273 0.4 (e.g. Lilly et al. 1996; Hippelein et al. 2003). In optically selected quasars, [O ii] emission in excess of the base level that is expected from non-stellar photoionization indicates star formation accompanying quasar activity (e.g. Ho 2005). These lines will hence allow a better understanding of the distribution and kinematical properties of the atomic gas. Recently, Shen & M\u00e9nard (2012) and Khare et al. (2014) studied the properties of [O ii] \u03bb3727 emission from a large sample of associated Mg ii absorbers, in order to understand the origin of associated Mg ii absorption systems relative to the AGN. Khare et al. (2014) find a high excess [O ii] emission in the composite quasar spectrum (constructed in the quasar rest frame) for the sample of quasars that show outflows, i.e. the systems for which the Mg ii absorption is blueshifted relative to the AGN. Further, the authors find that the presence of an associated Mg ii absorption enhances [O ii] emission in the composite quasar spectrum. However the excess emission flux in the [O ii] line does not depend on the strength of the Mg ii absorption line (Khare et al. 2014). The excess [O ii] line flux could be originating either from the host galaxy, or the parent AGN; in the former case, the excess [O ii] emission could indicate higher star formation in the host galaxy. The association of high [O ii] line emission with the occurrence of outflows suggests either strong stellar outbursts in the host galaxy, or strong AGN jet\u2013gas interactions.","Citation Text":["Bergeron & Stasi\u0144ska 1986"],"Functions Text":["Indeed, Mg ii is known as a good tracer of high column density H\u2009i\u2009gas in damped Lyman \u03b1 absorbers (e.g."],"Functions Label":["Background"],"Citation Start End":[[677,702]],"Functions Start End":[[572,676]]}
{"Identifier":"2020AandA...638A..34J__Hardcastle_(2018)_Instance_1","Paragraph":"We stress that the age values discussed here should be taken with care as they are based on simulations which rely on a number of assumptions. One of these is the magnetic field which is typically computed using the often unrealistic equipartition conditions. For example, a factor two difference in the magnetic field value translates directly into about a factor two age difference. Another important parameter is the assumed total source age distribution which is assumed to be different in the various simulations. For example, in the case of the modelling presented in Brienza et al. (2017) and Godfrey et al. (2017), the distribution of ton was taken to be a truncated log-normal distribution between 20 and 200 Myr, with a median of 30 Myr. Hardcastle (2018) on the other hand explored two types of models, with ages being either uniformly distributed between 0 and 1000 Myr, or linearly distributed in log space between 1 and 1000 Myr. These latter authors found that the uniform-age models efficiently explain the size and luminosity statistics of bright sources observed by LOFAR, while the log-uniform models were a better representation of the faint population. Similarly, Shabala et al. (2020) showed that power-law age-distribution models (i.e. models where the radio-source population is dominated by short-lived sources) can explain the observed properties of both active and remnant and\/or restarted LOFAR sources. These findings are consistent with the modelling assumptions of Brienza et al. (2017) and Godfrey et al. (2017), as the majority of short-lived sources will simply never grow large or luminous enough to make it into the observable LOFAR sample1. Modelling by both Shabala et al. (2020) and Hardcastle (2018) confirms the picture in which the remnant lobes fade quickly below the LOFAR detection limit. This fading is more rapid for older sources. Future studies, including modelling of the spectral index and follow-up observations, will help us to put tighter constraints on the age of the candidate restarted radio galaxies and the time that passed between the two bursts of activity.","Citation Text":["Hardcastle (2018)","Hardcastle (2018)"],"Functions Text":["on the other hand explored two types of models, with ages being either uniformly distributed between 0 and 1000 Myr, or linearly distributed in log space between 1 and 1000 Myr.","These latter authors found that the uniform-age models efficiently explain the size and luminosity statistics of bright sources observed by LOFAR, while the log-uniform models were a better representation of the faint population.","Modelling by both Shabala et al. (2020) and","confirms the picture in which the remnant lobes fade quickly below the LOFAR detection limit."],"Functions Label":["Differences","Compare\/Contrast","Similarities","Similarities"],"Citation Start End":[[748,765],[1722,1739]],"Functions Start End":[[766,943],[944,1173],[1678,1721],[1740,1833]]}
{"Identifier":"2022MNRAS.511.2105K__Husemann_et_al._2019_Instance_2","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"],"Functions Text":["This can be mitigated by using integral field spectroscopy (IFS) which is being increasingly used for extragalactic studies (e.g.","although there could still be projection effects with the IFS data."],"Functions Label":["Background","Background"],"Citation Start End":[[1561,1581]],"Functions Start End":[[1347,1476],[1635,1702]]}
{"Identifier":"2017AandA...600A..47D__Becker_2015_Instance_1","Paragraph":"To check for the validity of this idea, the case of some PACWBs already investigated in X-rays deserves to be examined. Some systems, such as WR\u2009140 (Williams et al. 1990; Pollock et al. 2005) or Cyg\u2009OB2\u2009#8A (Harnden et al. 1979; De Becker et al. 2006), are indeed known to be both well-studied PACWBs and bright thermal X-ray emitters, due to their colliding winds. However, other examples should be commented. The first is HD\u2009167971. It consists of a hierarchical triple system with a 3.3-day O-type binary and a third later-type companion evolving on a 21-yr orbit (De Becker et al. 2012; Ibanoglu et al. 2013). Recently, it has been demonstrated that its thermal X-ray spectrum is dominated by the colliding-wind region in the close binary, with only a weak\/moderate contribution coming from the wind-wind interaction in the wide orbit (De Becker 2015). However, it has also been demonstrated that HD\u2009167971 presents a non-thermal radio emission modulated with a period of 21 yr, thus providing compelling evidence for a colliding-wind region origin (Blomme et al. 2007). In addition, HD\u2009167971 is the brightest O-type synchrotron radio emitter included to date in the catalogue of PACWB. Another relevant example is HD\u200993129A. It is known as an early O-type wide binary with undetermined period (probably more than a century), with non-thermal radio emission recently imaged in Long Baseline Array observations (Benaglia et al. 2015). The direct imaging of this emission region points to a clear and significant synchrotron radio emission coincident with the colliding-wind region. However, the investigation in X-rays by Gagn\u00e9 et al. (2011) did not point to any strong over-luminosity attributable to an X-ray bright wind-wind interaction region. These examples demonstrate that the association between bright thermal X-ray emission and bright non-thermal radio emission is not necessarily clear. As a result, the quest for new members of the catalogue should by no means be restricted to objects known to display a bright thermal soft X-ray spectrum. ","Citation Text":["De Becker 2015"],"Functions Text":["Recently, it has been demonstrated that its thermal X-ray spectrum is dominated by the colliding-wind region in the close binary, with only a weak\/moderate contribution coming from the wind-wind interaction in the wide orbit"],"Functions Label":["Background"],"Citation Start End":[[841,855]],"Functions Start End":[[615,839]]}
{"Identifier":"2019ApJ...883..174X__Lattimer_&_Prakash_2000_Instance_1","Paragraph":"Both the magnitude and slope of nuclear symmetry energy contribute to the pressure of NS matter. For example, the pressure of npe matter in NSs at \u03b2 equilibrium at density \u03c1 and isospin asymmetry \u03b4 is explicitly\n2\n\n\n\n\n\nThe first term is the SNM pressure \n\n\n\n\n\n, while the last two terms are the isospin-asymmetric pressure \n\n\n\n\n\n from nucleons and electrons, separately. At the saturation density \u03c10, P0 vanishes, and the electron contribution is also negligible, leaving the total pressure completely determined by the slope of the symmetry energy. Both P0 and Pasy increase with density with rates determined separately by the respective density dependences of the SNM EOS and the symmetry energy. In the region around \u03c10 \u223c 2.5\u03c10, Pasy dominates over P0 using most EOSs available. At higher densities, the SNM pressure P0 dominates, while Pasy also plays an important role, depending on the high-density behaviors of nuclear symmetry energy (Li & Steiner 2006). The exact transition of dominance from Pasy to P0 depends on the stiffnesses of both the SNM EOS and the symmetry energy. It is also well known that the radius R1.4 of canonical NSs is essentially determined by the pressure at densities around \u03c10 \u223c 2.5\u03c10 (Lattimer & Prakash 2000), while the maximum mass of NSs is determined by the pressure at higher densities reached in the core. Thus, knowledge of the density dependence of nuclear symmetry energy is important for understanding measurements of the masses and especially the radii of NSs. Moreover, the critical densities for forming hyperons (Sumiyoshi & Toki 1994; Lee 1996; Kubis & Kutschera 2003; Provid\u00eancia et al. 2019), \u0394(1232) resonances (Drago et al. 2014; Cai et al. 2015; Zhu et al. 2016; Sahoo et al.2018; Ribes et al. 2019), kaon condensation (Odrzywolek & Kutschera 2009), and the quark phase (Di Toro et al. 2010; Wu & Shen 2019) are also known to depend sensitively on the high-density nuclear symmetry energy. Information about the latter is thus a prerequisite for exploring the evolution of the NS matter phase diagram in the isospin dimension. Once Esym(\u03c1) is better determined and, hopefully, with more astrophysical data, it will be interesting to introduce extra model parameters characterizing the physics associated with the exotic particles and\/or new phases predicted to appear in superdense neutron-rich matter. With the very limited data available and the expensive computational costs of simultaneously inferring a lot more than the six EOS parameters we already have in the minimum NS model consisting of only nucleons and two leptons, our goals in this work are conservative and practical. However, inferring new physics parameters associated with the exotic degrees of freedom and new phases in superdense neutron-rich matter from astrophysical data by extending the model used in the present work are high on our working agenda.","Citation Text":["Lattimer & Prakash 2000"],"Functions Text":["It is also well known that the radius R1.4 of canonical NSs is essentially determined by the pressure at densities around \u03c10 \u223c 2.5\u03c10","while the maximum mass of NSs is determined by the pressure at higher densities reached in the core."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1220,1243]],"Functions Start End":[[1086,1218],[1246,1346]]}
{"Identifier":"2018MNRAS.478.4657P__Gadotti_2011_Instance_1","Paragraph":"We have studied a sample of 263 LSB galaxies observed by Green Bank Telescope (Schneider et al. 1992) which are in overlap with the SDSS footprint. We have performed two-component bulge-disc decomposition of 263 galaxies in the SDSS g, r and, i bands and investigated their structural properties in detail. We have found that $60\\hbox{ per cent}$ LSBs in our specific sample are bulgeless, while $40\\hbox{ per cent}$ are with bulges. Some of the LSBs are associated with significant bulge component with B\/T> 0.1. Since LSBs are known to be dwelling in less-dense environment (Rosenbaum & Bomans 2004), mergers and interactions are unlikely to have led the bulge formation. We also have $15 \\hbox{ per cent}$ barred galaxies in our sample. Our findings of bulges and bars suggest a considerable on-going evolution in the local LSB galaxies and the bars might as well be playing a role in the bulge growth (Laurikainen et al. 2007; Gadotti 2011; Cheung et al. 2013). The interesting fact about our sample is that they are not the class of giant LSB galaxies, in fact, most of our LSBs are faint blue and gas-rich and roughly half of them are hosting bars and bulges. Since LSB galaxies are dark matter-dominated, discs are known to be stable against bar formation (Ostriker & Peebles 1973; Efstathiou, Lake & Negroponte 1982; Christodoulou, Shlosman & Tohline 1995; Cervantes Sodi, Li & Park 2015; Algorry et al. 2017), as shown in numerical simulations of stellar discs with dark matter dominance at all radii (Saha 2014). Question arises how these faint blue LSBs are making bars and bulges. One possibility is that LSB discs are embedded in dark matter haloes that are spinning (Jimenez et al. 1998; Vitvitska et al. 2002; Kim & Lee 2013) which might be promoting bar formation provided spin parameter is not too high (Saha & Naab 2013; Cervantes-Sodi et al. 2013; Long, Shlosman & Heller 2014; Collier, Shlosman & Heller 2018). Whether these bars lead to the formation of bulges or other processes such as minor mergers being involved, needs further and detailed investigation.","Citation Text":["Gadotti 2011"],"Functions Text":["Our findings of bulges and bars suggest a considerable on-going evolution in the local LSB galaxies and the bars might as well be playing a role in the bulge growth"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[931,943]],"Functions Start End":[[740,904]]}
{"Identifier":"2021MNRAS.508.4429C__McClintock_et_al._1976_Instance_1","Paragraph":"Vela X-1 (4U 0900\u221240) is an eclipsing high-mass X-ray binary (HMXB) discovered during rocket borne X-ray observations in 1967 (Chodil et al. 1967). It is located at a distance of \u223c2.0 kpc (Sadakane et al. 1985; Nagase 1989) in the Vela constellation. Recent estimates using Gaia data infer distance of $2.42^{+0.19}_{-0.17}$ kpc (Bailer-Jones et al. 2018). The system consists of a massive B0.5Ib supergiant HD 77581 (Brucato & Kristian 1972; Hiltner, Werner & Osmer 1972; Jones & Liller 1973; Vidal, Wickramasinghe & Peterson 1973) having mass of about \u223c23 $\\rm {M_\\odot }$ and radius of \u223c34 $\\rm {R_\\odot }$ (Van Paradijs et al. 1976; Joss & Rappaport 1984; Nagase 1989; van Kerkwijk et al. 1995) and a neutron star with mass \u223c1.8 $\\rm {M_\\odot }$ (Van Paradijs et al. 1976; Nagase 1989; Barziv et al. 2001; Rawls et al. 2011). The orbital period of the binary system is about 9 d (Hiltner et al. 1972; Forman et al. 1973; Vidal et al. 1973; Watson & Griffiths 1977; van Kerkwijk et al. 1995). Due to the close proximity of about 1.7 $\\rm {R_\\star }$ (Conti 1978; Quaintrell et al. 2003) between the neutron star and its companion, the neutron star is immersed in the dense stellar wind of the donor star having typical mass-loss rate of about $\\dot{M} \\mathrm{\\sim 10^{-6} ~{M_\\odot } \\, yr^{-1}}$ (Hutchings 1974; Dupree et al. 1980; Nagase et al. 1986; Sako et al. 1999). A fraction of the stellar wind is captured and channelled along the strong magnetic field (\u223c2.7 \u00d7 1012\u2009G; Kretschmar et al. 1996; Coburn et al. 2002; Kreykenbohm et al. 2002) of the neutron star on to the magnetic poles, producing regular X-ray pulsations caused by the spin period \u223c283\u2009s (Rappaport & McClintock 1975; McClintock et al. 1976) of the neutron star. Although Vela X-1 is known to be a persistent source having luminosity of about 4 \u00d7 1036\u2009erg\u2009s\u22121 (McCray et al. 1984; Sadakane et al. 1985; Nagase et al. 1986; Kreykenbohm et al. 2002), it shows a plethora of X-ray variabilities such as sudden flares lasting a few minutes to several hours wherein the luminosity increases by several folds within very short time-scales of a few tens of seconds (Lapshov et al. 1992; Staubert et al. 2004; Kreykenbohm et al. 2008). Occurrence of sudden flares in this system are not so well understood and is believed to be due to enhanced accretion rate due to variabilities in the stellar wind from the companion star (Nagase et al. 1983; Haberl & White 1990) or accretion of clumpy stellar wind (Staubert et al. 2004; Ducci et al. 2009; F\u00fcrst et al. 2010; Odaka et al. 2013). Some studies suggest that sudden flares might be related to formation of transient accretion disc (Inoue et al. 1984; Taam & Fryxell 1989; Haberl & White 1990; Kreykenbohm et al. 2008). Another bizarre manifestation seen in Vela X-1 is occurrence of abrupt \u2018off-states\u2019 wherein X-ray pulsations cessation (within less than the pulse period) is observed for several tens of minutes at a time (Inoue et al. 1984; Lapshov et al. 1992; Kreykenbohm et al. 1999, 2008; Doroshenko, Santangelo & Suleimanov 2011; Sidoli et al. 2015). These states are poorly understood and might be caused by changes in the accretion rate due to variabilities in the stellar wind (Lapshov et al. 1992; Coburn et al. 2002). Some earlier studies also suggest that \u2018off-states\u2019 might be associated with formation of transient accretion discs (Inoue et al. 1984) or the accretion is choked due to the sudden onset of propeller effect (Kreykenbohm et al. 2008). It has also been suggested that the onset of these \u2018off-states\u2019 can be caused due to transition from the higher luminosity Compton cooling regime to the lower luminosity radiative cooling regime (Shakura, Postnov & Hjalmarsdotter 2013) or due to unstable hydrodynamic flows in the vicinity of the neutron star (Manousakis & Walter 2015a). Recent numerical studies suggest formation of temporary accretion discs in wind-fed X-ray pulsars (Karino, Nakamura & Taani 2019; El Mellah, Sundqvist & Keppens 2019a; El Mellah et al. 2019b) but conclusive evidence of their existence has been elusive. Interestingly, Liao et al. (2020) infer presence of temporary accretion disc in Vela X-1 during an extended low state lasting at least 30 ks that was accompanied by unusual spin-up event and similar Fe K\u03b1 fluxes compared to the preceding flaring period.","Citation Text":["McClintock et al. 1976"],"Functions Text":["A fraction of the stellar wind is captured and channelled along the strong magnetic field","of the neutron star on to the magnetic poles, producing regular X-ray pulsations caused by the spin period \u223c283\u2009s","of the neutron star."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[1696,1718]],"Functions Start End":[[1377,1466],[1552,1665],[1720,1740]]}
{"Identifier":"2021ApJ...920..145H__Damone_et_al._2018_Instance_3","Paragraph":"Over the past decade, many attempts to address this issue have been carried out, such as from the perspective of conventional nuclear physics and even exotic physics beyond the standard BBN framework (Angulo et al. 2005; Cyburt et al. 2008, 2016; Boyd et al. 2010; Pospelov & Pradler 2010; Fields 2011; Kirsebom & Davids 2011; Wang et al. 2011; Broggini et al. 2012; Coc et al. 2012, 2013, 2014; Cyburt & Pospelov 2012; Kang et al. 2012; Voronchev et al. 2012; Bertulani et al. 2013; Hammache et al. 2013; He et al. 2013; Kusakabe et al. 2014; Pizzone et al. 2014; Yamazaki et al. 2014; Hou et al. 2015, 2017; Famiano et al. 2016; Damone et al. 2018; Hartos et al. 2018; Luo et al. 2019; Rijal et al. 2019; Clara & Martins 2020). However, despite the fact some solutions using exotic physics have succeeded in resolving this issue, it appears there is still no universally accepted solution in the academic community since validations of these mysterious exotic physics are beyond the capabilities of current science. Conversely, it seems more worthwhile to exclude any potential possibility of resolving the 7Li discrepancy from the perspective of nuclear physics. It is known that the majority of the primordial 7Li production arises from the decay of 7Be by electron capture during the 2 months after BBN stops. Thus, for the solution of the Li problem, reactions involving 7Be could be more significant than those involving 7Li. Therefore, many reactions that potentially destroy 7Be were investigated to solve this discrepancy over past 10 yr (Kirsebom & Davids 2011; Broggini et al. 2012; Hammache et al. 2013; Hou et al. 2015; Hartos et al. 2018). Meanwhile, enormous efforts have been made to refine the reaction rates of key BBN reactions in the past 20 yr (Smith et al. 1993; Descouvemont et al. 2004; Serpico et al. 2004; Cyburt & Davids 2008; Neff 2011; Pizzone et al. 2014; Tumino et al. 2014; Hou et al. 2015; Barbagallo et al. 2016; Iliadis et al. 2016; Kawabata et al. 2017; Lamia et al. 2017, 2019; Damone et al. 2018; Rijal et al. 2019; Mossa et al. 2020), but the probability of solving or alleviating the 7Li problem by improving our knowledge of relevant nuclear reaction rates still cannot be eliminated. Recent experiments for key nuclear reactions like 7Be(n,p)7Li and 7Be(d,p)24He allow for a reduction of the 7Li production by about 12% (Damone et al. 2018; Rijal et al. 2019) compared to previous calculations. At present, nuclear uncertainties cannot rule out that some of the reactions destroying 7Li are indeed more efficient than those currently used (Boyd et al. 2010; Chakraborty et al. 2011).","Citation Text":["Damone et al. 2018"],"Functions Text":["Recent experiments for key nuclear reactions like 7Be(n,p)7Li and 7Be(d,p)24He allow for a reduction of the 7Li production by about 12%","compared to previous calculations."],"Functions Label":["Differences","Differences"],"Citation Start End":[[2364,2382]],"Functions Start End":[[2227,2362],[2403,2437]]}
{"Identifier":"2015AandA...584A.103S__Potekhin_et_al._2013_Instance_3","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":["Potekhin et al. 2013"],"Functions Text":["We shall adopt here the EoS of the BSk21 model","as a representative example of contemporary EoS for the complete NS structure,"],"Functions Label":["Uses","Uses"],"Citation Start End":[[1785,1805]],"Functions Start End":[[1675,1721],[1828,1906]]}
{"Identifier":"2022ApJ...929..186L__Lanzoni_et_al._2013_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":["Lanzoni et al. 2013"],"Functions Text":["Our group is addressing this problem by combining a variety of complementary perspectives:","(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"],"Functions Label":["Uses","Uses"],"Citation Start End":[[793,812]],"Functions Start End":[[0,90],[612,791]]}
{"Identifier":"2016MNRAS.455.4426V__Smith_&_Dwek_1998_Instance_1","Paragraph":"It is interesting that the tracers of atomic and molecular hydrogen become maximum at distances \u223c0.85 kpc, where no dust cloud is inferred from the Swift\/XRT data (see Fig. 5). A fiducial dust cloud at the distance of \u223c0.85 kpc would produce an X-ray ring with angular size \u223c9 arcmin on MJD 57205.5. As this would fall well within the Swift\/XRT FOV (see Fig. 2), its absence should be related to the dust properties of the respective cloud. We estimated therefore that the maximum grain size should be \u226b 0.2\u2009\u03bcm. A dust cloud composed by large grains ( \u223c 1\u2009\u03bcm) would suppress the scattered intensity even at 1 keV, while it would significantly attenuate X-rays at \u22640.5 keV (e.g. Smith & Dwek 1998; Corrales & Paerels 2015). Alternatively, the absence of the X-ray ring might indicate a much smaller number of grains per hydrogen atom at \u223c0.85 kpc compared to the other clouds. The analysis presented so far, as in most studies, does not take into account possible variations of the radial profiles in the azimuthal direction. We remind that the profiles shown in Figs 3 and B1 were created by summing the counts of each annulus, i.e. integrating over the azimuthal angle. Although this method is sufficient for probing the average properties of the dust clouds, such as position and average column density, it cannot provide information about the spatial inhomogeneity of the dust clouds and\/or their inclination with respect to the LOS. The photon statistics of the Swift\/XRT observations are sufficient for quantifying the azimuthal variations of the ring intensity (see also Fig. 2). Fig. 12 shows the azimuthal variation of the scattered photons for Swift\/XRT observation 2. There are three important features that need to be mentioned: (i) the radial profile of ring 5 is highly variable along the azimuthal direction, with the peak number of counts changing by a factor of \u223c5 above the background level (see sub-panels for 180\u00b0\u2013216\u00b0 and 252\u00b0\u2013288\u00b0); (ii) azimuthal variations are also present in the rings 3 and 4, yet the variability amplitude is lower when compared to the ring 5. The peak number of counts of the fourth ring is for all azimuth angles larger compared to that of the third ring; (iii) the azimuthal behaviour of the rings 1 and 2 is the most intriguing, as the second peak (corresponding to the ring 2) becomes prominent for azimuth angles 0\u00b0\u201336\u00b0 and 324\u00b0\u2013360\u00b0 while the first peak dominates for azimuth angles of 216\u00b0\u2013324\u00b0. Obviously, this information is lost from the azimuthal integrated radial profiles shown in Fig. 3. Similar features appear in all Swift\/XRT observations with their significance varying according to the photon statistics of the observation.","Citation Text":["Smith & Dwek 1998"],"Functions Text":["A dust cloud composed by large grains ( \u223c 1\u2009\u03bcm) would suppress the scattered intensity even at 1 keV, while it would significantly attenuate X-rays at \u22640.5 keV (e.g."],"Functions Label":["Uses"],"Citation Start End":[[678,695]],"Functions Start End":[[512,677]]}
{"Identifier":"2019AandA...624A..60L__Dubernet_et_al._2016_Instance_1","Paragraph":"One of the main uncertainties in elemental abundances is the quality of the adopted atomic data used in spectral synthesis calculations (Bigot & Thevenin 2008). The Belgian repository of fundamental atomic data and stellar spectra (BRASS) aims to provide astronomers with quality information for the large amount of atomic data available for high-resolution optical spectroscopy, in an attempt to help reduce systematic input errors in quantitative spectroscopy from atomic data and line selection (Lobel et al. 2017). Previously, we retrieved and cross-matched a large quantity of atomic data from several major atomic databases such as the Vienna Atomic Line Database (VALD3; Ryabchikova et al. 2015), the National Institute of Standards and Technology Atomic Spectra Database (NIST ASD; Kramida et al. 2018), and providers within the Virtual Atomic and Molecular Data Centre (VAMDC; Dubernet et al. 2016), in preparation for quality assessment work (Laverick et al. 2018; hereafter Paper 1). In this work the atomic data of seemingly \u201cunblended\u201d spectral lines are quality assessed against several benchmark dwarf stars, including the Sun, spanning late F-type to early K-type stars, for the spectral range 4200\u20136800 \u00c51. Unblended spectral lines for stars of \u223cG2V spectral type are identified in a homogeneous manner using both the observed benchmark spectra, and the theoretical input line list of Paper I. Astrophysical oscillator strengths2 are derived for these unblended lines, using two commonly utilised methods, to gauge the reliability of the spectral line for quantitative spectroscopy and to produce \u201cbenchmark\u201d log(g\u2006f) values for quality assessment of atomic data. The literature log(g\u2006f) values are then compared against these benchmark log(g\u2006f) values to determine which literature values reliably reproduce the stellar spectra of cool dwarf stars, and thus whether the values can be recommended for spectroscopic modelling. This paper presents three main sets of results:","Citation Text":["Dubernet et al. 2016"],"Functions Text":["Previously, we retrieved and cross-matched a large quantity of atomic data from several major atomic databases such as the","and providers within the Virtual Atomic and Molecular Data Centre (VAMDC;","in preparation for quality assessment work"],"Functions Label":["Uses","Uses","Uses"],"Citation Start End":[[886,906]],"Functions Start End":[[519,641],[812,885],[909,951]]}
{"Identifier":"2020ApJ...897..177P__Lister_et_al._2019_Instance_1","Paragraph":"Relativistic jets are the manifestation of the extreme processes that occur within the central regions of galaxies (see Blandford et al. 2019 for a review). Active galactic nuclei (AGNs) hosting relativistic jets closely aligned to the line of sight are called blazars. Due to their peculiar orientation, the relativistic amplification of the nonthermal jetted radiation (Doppler boosting; see, e.g., Rybicki & Lightman 1979) leads to the observation of a number of interesting phenomena. A few examples are detection at all accessible frequencies (e.g., Abdo et al. 2011), observation of temporal and spectral variability (Gaidos et al. 1996; Acciari et al. 2011; Fuhrmann et al. 2014; Paliya et al. 2017b), and superluminal motion and high brightness temperature (Scheuer & Readhead 1979; Lister et al. 2019). The optical and radio emissions detected from blazars are found to be significantly polarized (e.g., Fan et al. 2008; Itoh et al. 2016). The flux enhancement also makes blazars a dominating class of \u03b3-ray emitters in the extragalactic high-energy sky (Ajello et al. 2020) and one of the very few astrophysical source classes detected at cosmic distances (e.g., Romani et al. 2004; Sbarrato et al. 2013). Blazars are classified as flat-spectrum radio quasars (FSRQs) and BL Lac objects based on their optical spectroscopic properties. FSRQs are characterized by broad emission lines (rest-frame equivalent width >5 \u212b), whereas BL Lac sources exhibit weak or no emission lines in their optical spectra, thereby making it challenging to detect their redshift (Stickel et al. 1991). BL Lac objects are known to exhibit a negative or mildly positive evolution compared to the strong positive evolution noticed in FSRQs (Ajello et al. 2012, 2014). Altogether, FSRQs dominate the known population of high-redshift (z \u2273 3) blazars and are found to be much more luminous than the BL Lac population (e.g., Ajello et al. 2009; Ackermann et al. 2017; Paliya et al. 2019d).","Citation Text":["Lister et al. 2019"],"Functions Text":["A few examples are","and superluminal motion and high brightness temperature"],"Functions Label":["Background","Background"],"Citation Start End":[[791,809]],"Functions Start End":[[489,507],[709,764]]}
{"Identifier":"2020MNRAS.493.4868L__Lazarian_&_Hoang_2007b_Instance_1","Paragraph":"Recently, polarized (sub)millimetre emission has been detected in an increasing number of discs by Atacama Large Millimeter\/submillimeter Array (ALMA) with its high sensitivity and angular resolution. However, the origin of disc polarization remains uncertain, since grains do not have to be aligned with just the magnetic field (Kataoka et al. 2017; Yang et al. 2019). They may also be aligned in the direction of the radiative anisotropy (Lazarian & Hoang 2007a; Tazaki, Lazarian & Nomura 2017) or the drift velocity of the grains relative to the ambient gas (Gold 1952; Lazarian 1995; Lazarian & Hoang 2007b). Furthermore, even spherical grains can produce polarized emission by self-scattering of large grains in an anisotropic radiation field (Kataoka et al. 2015; Yang et al. 2016, 2017; Stephens et al. 2019). The scattering interpretation of the disc polarization is favoured in several targets (e.g. Stephens et al. 2014, 2017; Kataoka et al. 2016; Bacciotti et al. 2018; Dent et al. 2019; Girart et al. 2018; Harris et al. 2018; Hull et al. 2018; Lee et al. 2018). One way to gauge the effects of scattering and identify polarization from aligned grains would be to observe at multiple wavelengths since the efficiency for scattering for grains of given sizes decreases rapidly with the wavelength in the optically thin and small-particle (or Rayleigh scattering) limit. Indeed, in the disc of Class I protostar BHB 07-11, Alves et al. (2018) detected polarization with ALMA at three wavebands (Bands 3, 6, and 7 or \u223c 3, 1.3, and 0.87\u2009mm, respectively) with consistent polarization orientations across three bands and increasing polarization fraction with wavelength, which is generally not expected for scattering-induced polarization. The rather high mean polarization fractions (\u223c7.9, 5.3, and 3.5\u2009${{\\ \\rm per\\ cent}}$ for Bands 3, 6, and 7 respectively) are also higher than those typically produced in models of scattering-induced disc polarization ($\\sim \\!1 {{\\ \\rm per\\ cent}}$). At least for this well-studied source, scattering is unlikely the main mechanism for producing the observed multiwavelength disc polarization and aligned grains are favoured.","Citation Text":["Lazarian & Hoang 2007b"],"Functions Text":["They may also be aligned in the direction of","or the drift velocity of the grains relative to the ambient gas"],"Functions Label":["Background","Background"],"Citation Start End":[[588,610]],"Functions Start End":[[370,414],[497,560]]}
{"Identifier":"2021MNRAS.504.4626K__Kraljic_et_al._2020b_Instance_4","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"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[1950,1970]],"Functions Start End":[[1593,1911]]}
{"Identifier":"2017ApJ...839...26D__Conroy_et_al._2006_Instance_1","Paragraph":"One might hope that these global observations would strongly constrain the SFHs of individual galaxies, but this connection is not easily established. For example, average SFHs can be inferred by integrating the main sequence SFR over time (e.g., Leitner 2012), but this approach leads to inconsistencies (Leja et al. 2015). Instead, the most successful theoretical models link the growth of stellar mass to the growth of the dark matter halos that galaxies inhabit, for example, via subhalo abundance matching (e.g., Kravtsov et al. 2004; Conroy et al. 2006; Behroozi et al. 2013b; Moster et al. 2013), halo occupation distributions (e.g., Peacock & Smith 2000; Seljak 2000; Hearin et al. 2016), semi-analytic models (Kauffmann et al. 1993; Somerville et al. 2001; Guo et al. 2011), or other assumptions (Bouch\u00e9 et al. 2010; Dav\u00e9 et al. 2012; Lilly et al. 2013; Tacchella et al. 2013; Mitra et al. 2017). One important conclusion from these models is that there has to be significant scatter between halo and galaxy masses (and thus growth histories) in order to explain observations (More et al. 2009; Behroozi et al. 2013b; Reddick et al. 2013; Gu et al. 2016). As a result, even models that agree on global constraints can lead to orthogonal interpretations of the evolution of individual galaxies. A good example for such disagreement is the \u201crapid quenching\u201d7\n\n7\nThe term \u201cquenching\u201d is somewhat ambiguous. In this paper, we use it to mean the cessation of star formation, without any presumption as to whether the decrease happens quickly or slowly, and whether it happens due to a diminishing gas supply or other physical processes.\n framework where galaxies follow the main sequence until they sharply fall below the main sequence (Peng et al. 2012; Wetzel et al. 2013; Tacchella et al. 2016b; Tinker et al. 2016) and the \u201cstochastic\u201d framework where correlated scatter and the central limit theorem lead to the main sequence (Kelson 2014; Kelson et al. 2016).","Citation Text":["Conroy et al. 2006"],"Functions Text":["Instead, the most successful theoretical models link the growth of stellar mass to the growth of the dark matter halos that galaxies inhabit, for example, via subhalo abundance matching (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[540,558]],"Functions Start End":[[325,517]]}
{"Identifier":"2020ApJ...891..111K__Howe_&_Burrows_2015_Instance_1","Paragraph":"To explore magma effects on sub-Neptune atmospheres, we assume the atmosphere equilibrates with a well-stirred magma ocean. Our model makes the following simplifications: (a) we consider only the elements Fe, Mg, Si, O, and H. Chemically reduced carbon compounds may also contain H; for simplicity, we omit consideration of them here. We also restrict ourselves to the range of magma elemental compositions for which SiO2 is a major constituent. (b) We set 2000 K \u2264 Tmai \u2264 3000 K, so that the magma\u2013atmosphere interface is molten but the vapor pressure of the magma is small relative to the total atmospheric pressure (Fegley et al. 2016; Sossi & Fegley 2018). This Tmai is at the low end of the Tmai output by thermal evolution models for multi-Gyr-old sub-Neptunes (e.g., Bodenheimer & Lissauer 2014; Howe & Burrows 2015; Vazan et al. 2018a). We consider lower Tmai in Section 4.4. (c) We ignore Fe3+ (i.e., Fe2O3). We expect Fe3+ will be a minor constituent of a magma ocean equilibrated with a H2-dominated atmosphere. Equal thermodynamic activities of FeO and Fe2O3 in magma at 3000 K require fO2 = 28 bar, much higher than expected from thermal dissociation of steam at any P and T below 1 kilobar and 3500 K. (d) We assume that metal (if present) is pure Fe for these calculations; in reality, the metal will be an Fe-dominated alloy. We may be (slightly) underestimating fO2 by doing this. (e) We track nonideal behavior of both H2 and H2O (Appendix C), but we assume ideal mixing of H2 and H2O. This is a valid assumption both under mineral-free conditions in the atmosphere for T > 650 K, and also at the magma\u2013atmosphere interface given our model assumptions (Seward & Franck 2019; Bali et al. 2013; Soubiran & Militzer 2015). We ignore joint-solubility effects. (f) We use a single value of gravitational acceleration g, corresponding to 1.2\u00d7 the bare-rock radius, to convert from bottom-of-atmosphere pressure to atmosphere column mass. (g) We ignore the effect of dissolved volatiles on core mass.","Citation Text":["Howe & Burrows 2015"],"Functions Text":["This Tmai is at the low end of the Tmai output by thermal evolution models for multi-Gyr-old sub-Neptunes (e.g.,"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[803,822]],"Functions Start End":[[661,773]]}
{"Identifier":"2017ApJ...848..126S__Luo_et_al._2017_Instance_1","Paragraph":"In Figure 2, we show the Chandra 0.3\u20138 keV images with SDSS contours overlaid for the six mergers in our sample. In Table 4 we list the number of counts detected in the 0.3\u20138 keV and 2\u20138 keV bands, and list the detection threshold, PB, associated with each position. There are a number of PB thresholds employed in the literature to define source detection significance. Based on weak sources in Chandra deep field images, the threshold for PB adopted based on a balance between reliability and completeness varies from 0.002 to 0.007 (Xue et al. 2011, 2016; Luo et al. 2017). Based on even the lowest of all these thresholds, all sources are detected. Note that the detection thresholds adopted in the literature for sources at known locations are often PB  0.01, which corresponds to \n\n\n\n\n\n (e.g., Lansbury et al. 2014). We note that due to insufficient counts, coupled with the fact the exact location of the galaxy nuclei is uncertain in the advanced mergers in our sample (see SDSS images in Figure 1), it is impossible to quantify any possible spatial offsets between the positions of the detected Chandra sources and the actual galactic nuclei. It is therefore not possible to determine using these observations alone if some of the mergers host offset AGNs, another unambiguous signature of a galaxy merger that can probe AGN triggering through galaxy interactions (Barrows et al. 2016, 2017). We note that for sources within the 2\u2032 of the boresight, the absolute accuracy of source locations on the ACIS-S chip has a 90% uncertainty radius of \n\n\n\n\n\n,18\n\n18\n\nhttp:\/\/cxc.cfa.harvard.edu\/proposer\/POG\/pdf\/MPOG.pdf\n\n indicating that the positions of all detected targets are consistent with the locations of the SDSS knots, suggesting that the detected sources likely correspond to the nuclei of the mergers. J1036+0221, J1045+3519 Gal 1, J1221+1137 Gal 1, and J1306+0735 Gal 2 were detected in the hard band using the same detection threshold PB  0.002 using the 2\u20138 keV source counts. Note that there are insufficient counts to perform a spectral analysis or to obtain reliable HRs for our targets. We therefore list in Table 6 the observed X-ray luminosity, assuming a simple power-law model with \n\n\n\n\n\n, corrected for Galactic absorption.","Citation Text":["Luo et al. 2017"],"Functions Text":["Based on weak sources in Chandra deep field images, the threshold for PB adopted based on a balance between reliability and completeness varies from 0.002 to 0.007"],"Functions Label":["Uses"],"Citation Start End":[[559,574]],"Functions Start End":[[371,534]]}
{"Identifier":"2021MNRAS.507.5053E__Johnston_et_al._2006_Instance_2","Paragraph":"Multiwavelength observations of the GC indicate that the number of pulsars in the central few parsecs should be high (Wharton et al. 2012) and conditions are highly favourable for relativistic binaries (Faucher-Gigu\u00e8re & Loeb 2011). The dense nuclear star cluster surrounding Sgr A* (see e.g. Genzel, Eisenhauer & Gillessen 2010, for a review) contains a majority of older late-type stars, but contrary to expectations, massive young main-sequence stars (Ghez et al. 2003) and possible neutron star progenitors such as Wolf\u2013Rayet stars (Paumard et al. 2001). The presence of neutron stars is further indicated by large numbers of X-ray binaries, possible pulsar wind nebulae, X-ray features such as the \u2018cannonball\u2019 and compact radio variables (Muno et al. 2005; Wang, Lu & Gotthelf 2006; Zhao, Morris & Goss 2013, 2020). Despite this only six radio pulsars have been discovered within half a degree of Sgr A* (Johnston et al. 2006; Deneva, Cordes & Lazio 2009; Eatough et al. 2013c; Shannon & Johnston 2013) even after many dedicated searches at multiple wavelengths (Kramer et al. 1996a, 2000; Klein et al. 2004; Klein 2005; Deneva 2010; Macquart et al. 2010; Eatough et al. 2013a; Siemion et al. 2013). Hyperstrong scattering of radio waves in the GC has been the principal explanation for the scarcity of detected pulsars (Cordes & Lazio 1997, 2002; Lazio & Cordes 1998a,b), however, scatter broadening measurements of PSR J1745\u22122900 in Spitler et al. (2014) and Bower et al. (2014) appear to contest this.1 Other authors have noted that the lack of GC pulsars is expected under a certain set of conditions and considering the sensitivity limits of existing pulsar surveys (Chennamangalam & Lorimer 2014; Liu & Eatough 2017; Rajwade, Lorimer & Anderson 2017). Alternatively, the scarcity of detected pulsars might be caused by a more complex scattering structure towards the GC (Cordes & Lazio 1997; Lazio & Cordes 1998a, b; Johnston et al. 2006; Schnitzeler et al. 2016; Dexter et al. 2017).","Citation Text":["Johnston et al. 2006"],"Functions Text":["Alternatively, the scarcity of detected pulsars might be caused by a more complex scattering structure towards the GC"],"Functions Label":["Background"],"Citation Start End":[[1929,1949]],"Functions Start End":[[1764,1881]]}
{"Identifier":"2017ApJ...844...14L__Lim_et_al._2016_Instance_1","Paragraph":"In order to see whether the CN- and HK\u2032-strong stars in our study are also enhanced in Fe and s-process elements, we have compared our results with high-resolution spectroscopy by M15. In Figure 4, our \u03b4CN and \u03b4HK\u2032 indices are plotted with [Na\/Fe] and [Ca\/Fe] abundances, respectively, for 33 common stars. In general, the strength of the CN band is correlated with the N and Na abundances, while the CH band is affected by C abundance (Sneden et al. 1992; Smith et al. 1996; Marino et al. 2008). The upper panel of Figure 4 also shows a strong correlation between [Na\/Fe] and the \u03b4CN index, which is in good agreement with previous studies (Sneden et al. 1992; Lim et al. 2016).4\n\n4\nCareful inspection of the upper panel of Figure 4 also shows the possibility that the s-poor and s-rich groups are probably separated on this diagram, with the s-rich stars more enhanced in both [Na\/Fe] and \u03b4CN. This would imply that the variations in light elements would be present in each group with different s-process elements abundances, which has already been found in other GCs with s-process element and Fe variations, such as M2 and M22 (Marino et al. 2011; Yong et al. 2014). More samples of stars, however, are needed to confirm this trend in NGC 5286.\n The CN-weak (blue) and CN-strong (red) subpopulations are almost identical to the s-poor (triangles) and s-rich (squares) groups, respectively. In addition, the \u03b4HK\u2032 index is understandably correlated with the [Ca\/Fe] abundance with a few exceptions (see the lower panel of Figure 4). According to this comparison, the difference in \u03b4HK\u2032 index between CN-weak and CN-strong stars (\u223c0.094) is equivalent to 0.09 dex in \u0394[Ca\/Fe] and 0.15 dex in \u0394[Fe\/H]. These comparisons confirm that our results from low-resolution spectroscopy are consistent with those from high-resolution spectroscopy by M15. Consequently, the later-generation stars in NGC 5286 show the enhancements not only in light elements (CN) but also in heavy elements (Fe and Ca) and s-process elements, although the presence of Fe spread requires further investigations (see Mucciarelli et al. 2015; Lee 2016).","Citation Text":["Lim et al. 2016"],"Functions Text":["The upper panel of Figure 4 also shows a strong correlation between [Na\/Fe] and the \u03b4CN index, which is in good agreement with previous studies"],"Functions Label":["Similarities"],"Citation Start End":[[662,677]],"Functions Start End":[[497,640]]}
{"Identifier":"2019MNRAS.485.4841R__Creminelli_et_al._2010_Instance_2","Paragraph":"Although, the standard form of Press\u2013Schechter mass function with $f(\\nu)=\\sqrt{{2}\/{\\pi }} \\nu \\mathrm{ e}^{-\\frac{\\nu }{2}}$ which discussed in Press & Schechter (1974) and Bond et al. (1991) can provide a good approximation of the predicted number density of haloes, it fails by predicting approximation too many low-mass haloes and too few high-mass ones (Sheth & Tormen 1999, 2002; Lima & Marassi 2004). Thus, in this study we apply another well-known fitting formula which first proposed in Sheth & Tormen (1999): \n(24)\r\n\\begin{eqnarray*}\r\nf(\\nu)=0.2709\\sqrt{\\dfrac{2}{\\pi }}(1+1.1096\\nu ^{0.6})\\mathrm{ exp}(-\\dfrac{0.707 \\nu ^2}{2})\\,\\, .\r\n\\end{eqnarray*}\r\nIn a Gaussian density field, \u03c3 is given by \n(25)\r\n\\begin{eqnarray*}\r\n\\sigma ^2(R)=\\dfrac{1}{2 \\pi ^2}{\\int _0}^\\infty k^2 P(k) W^2(kR) \\, {\\rm d}k\\,\\, ,\r\n\\end{eqnarray*}\r\nwhere R = (3M\/4\u03c0\u03c1m0)1\/3 is the radius of the spherical overdense region, W(kR) is the Fourier transform of a spherical top-hat profile with radius R and P(k) is the linear power spectrum of density fluctuations (Peebles 1993). To obtain the value of \u03c3, we follow the procedure presented in Abramo et al. (2007a). Following on Ade et al. (2016), we use the normalization of matter power spectrum \u03c38 = 0.815 for \u039bCDM cosmology. The number density of virialized haloes above a certain value of mass M at zc, the collapse redshift obtained by \n(26)\r\n\\begin{eqnarray*}\r\nN(\\: M,z)={\\int _0}^\\infty \\dfrac{{\\rm d}n(z)}{{\\rm d}M^{\\prime }}\\, {\\rm d}M^{\\prime }\\,\\,. \r\n\\end{eqnarray*}\r\nThe above limit of integration in equation (26) is $M=10^{18}\\, \\mathrm{ M}_{\\rm \\odot}\\, \\mathrm{ h}^{-1}$ which such gigantic structures could not in practice be observed. Now we can calculate the number density of virialized haloes in both homogeneous and clustered DE scenarios using equations (23) and (26). In this way the total mass of a halo is equal to the mass of pressureless matter perturbations. However, the virialization of dark matter perturbations in the non-linear regime cannot be independent from the properties of DE (Lahav et al. 1991; Maor & Lahav 2005; Creminelli et al. 2010; Basse, Bjlde & Wong 2011). Thus, in clustered DE scenarios, we should consider the contribution of perturbated DE components to the total mass of the haloes (Creminelli et al. 2010; Basse et al. 2011; Batista & Pace 2013). Based on the behaviour of wde(z), DE can reduce or enhance the total mass of the virialized halo. One can obtain \u03f5(z), the ratio of DE mass to be taken into account with respect to the mass of dark matter, from: \n(27)\r\n\\begin{eqnarray*}\r\n\\epsilon (z)=\\dfrac{m_{\\rm DE}}{m_{\\rm DM}}\\,\\, ,\r\n\\end{eqnarray*}\r\nwhere the value of mDE depends on what we consider as the mass of DE component. When one only considers the contribution of the perturbations of DE, the mDE takes the form \n(28)\r\n\\begin{eqnarray*}\r\n{m_{\\rm DE}}^{\\mathrm{ Perturbed}}=4 \\pi \\bar{\\rho }_{\\rm DE}{\\int _0}^{R_{\\rm vir}} \\, {\\rm d}R R^2 \\delta _{\\rm DE}(1+3{c_{\\rm eff}}^2)\\,\\,. \r\n\\end{eqnarray*}\r\nIn the other hand, if we assume both DE contributions of perturbation and background level, the total mass of DE in virialized haloes takes this new form\n(29)\r\n\\begin{eqnarray*}\r\n{m_{\\rm DE}}^{\\mathrm{ Total}}=4 \\pi \\bar{\\rho }_{\\rm DE}{\\int _0}^{R_{\\rm vir}} {\\rm d}R R^2 [(1+3 w_{\\rm DE})+ \\delta _{\\rm DE}(1+3{c_{\\rm eff}}^2)]. \r\n\\end{eqnarray*}\r\nThe quantities inside a spherical collapsing region in the framework of the top-hat profile, evolve only with cosmic time. Thus from equation (28) one can find \n(30)\r\n\\begin{eqnarray*}\r\n\\epsilon (z)=\\dfrac{\\Omega _{\\rm DE}}{\\Omega _{\\rm DM}}\\dfrac{\\delta _{\\rm DE}}{1+\\delta _{\\rm DM}}\\,\\, \r\n\\end{eqnarray*}\r\nand from equation (29) we can obtain \n(31)\r\n\\begin{eqnarray*}\r\n\\epsilon (z)=\\dfrac{\\Omega _{\\rm DE}}{\\Omega _{\\rm DM}}\\dfrac{1+3 w_{\\rm DE}+\\delta _{\\rm DE}}{1+\\delta _{\\rm DM}}\\,\\, .\r\n\\end{eqnarray*}\r\nThe mass of dark matter also is obtained from (see also Batista & Pace 2013): \n(32)\r\n\\begin{eqnarray*}\r\n{m_{\\rm DM}}=4 \\pi \\bar{\\rho }_{\\rm DM}{\\int _0}^{R_{\\rm vir}} \\, {\\rm d}R R^2 (1+ \\delta _{\\rm DM})\\,\\,. \r\n\\end{eqnarray*}\r\nIn Fig. 5 we plot the evolution of \u03f5(z) using equation (30) as the definition of DE mass. We observe that, at high redshift, where the role of DE is less important, \u03f5 for both of parametrizations becomes negligible. This parameter has a greater value in the case of parametrization (2).","Citation Text":["Creminelli et al. 2010"],"Functions Text":["Thus, in clustered DE scenarios, we should consider the contribution of perturbated DE components to the total mass of the haloes"],"Functions Label":["Uses"],"Citation Start End":[[2272,2294]],"Functions Start End":[[2141,2270]]}
{"Identifier":"2015AandA...576L..16P__Chandler_&_Sjouwerman_2014_Instance_1","Paragraph":"To constrain the size of outflows that we could have missed, we performed simple simulations. We placed artificial, unipolar secondary sources next to a primary point source model representing Sgr A* and compared the closure phases obtained from the resulting artificial visibility data with the observations. We considered two geometries: a single-point source and a jet composed of ten equally spaced point sources (knots) with equal fluxes. We probed four orientations for the simulated outflows (see Fig. 2): along the major axis, along the minor axis of the beam, the jet direction claimed by Li et al. (2013), and the jet direction claimed by Yusef-Zadeh et al. (2012). We used total fluxes of 0.2 Jy and 0.55 Jy for the artificial sources; these values ensure that our simulated outflows are sufficiently faint to not violate the constraints given by the known recent brightness evolution of Sgr A* (0.2 Jy from the mean variability of \u224815% from June 2013 to February 2014 at 41 GHz, with 0.55 Jy corresponding to the strongest variation in the same period, Chandler & Sjouwerman 2014). For each simulation setup, we measured the average of absolute values of the closure phases for each triangle. We varied the distances of the model sources (for the jet model: the largest distance) from Sgr A* until we found a critical distance at which the absolute values of the simulated closure phases exceeded those of the observations by more than the 1\u03c3 error at all triangles. We summarize our results in Table 1. As expected, the critical distances are smaller for brighter outflows. Jet-like structures lead to larger critical distances than equally luminous single, compact sources. As a consequence of the very elongated beam, the critical distances for sources located along the major axis of the beam are larger by a factor of \u22487 than for those located along the minor axis. In a few cases (denoted \u201cN\/A\u201d in Table 1), the absolute values of the simulated closure phases were similar to those of the observations for all distances of the model sources, meaning that we were unable to identify a critical distance. Overall, our observations limit the extension of asymmetric (in the observer frame) jet-like outflows from Sgr A* to projected distances of \u22482.5 mas along the major axis and \u22480.4 mas along the minor axis. ","Citation Text":["Chandler & Sjouwerman 2014"],"Functions Text":["We used total fluxes of 0.2 Jy and 0.55 Jy for the artificial sources; these values ensure that our simulated outflows are sufficiently faint to not violate the constraints given by the known recent brightness evolution of Sgr A* (0.2 Jy from the mean variability of \u224815% from June 2013 to February 2014 at 41 GHz, with 0.55 Jy corresponding to the strongest variation in the same period,"],"Functions Label":["Uses"],"Citation Start End":[[1065,1091]],"Functions Start End":[[676,1064]]}
{"Identifier":"2015ApJ...798..104S___2000_Instance_1","Paragraph":"The Vela Molecular Ridge Cloud-D (hereafter VMR-D) (260\u00c2\u00b0 \u00e2\u0089\u00b2 \u00e2\u0084\u0093 \u00e2\u0089\u00b2 264\u00c2\u00b0; |b| \u00e2\u0089\u00b2 1\u00c2\u00b0) is part of a giant molecular complex located along the Galactic plane (260\u00c2\u00b0 \u00e2\u0089\u00b2 \u00e2\u0084\u0093 \u00e2\u0089\u00b2 272\u00c2\u00b0; |b| \u00e2\u0089\u00b2 3\u00c2\u00b0; Murphy & May 1991) and is then well suited to represent a typical star-forming region (SFR) of our Galaxy. For this reason a subregion of this cloud has been the subject of many previous papers, dealing with different observational aspects of the star formation (SF), such as the presence of outflows (Wouterloot & Brand 1999; Elia et\u00c2 al. 2007), jets (Lorenzetti et\u00c2 al. 2002; Giannini et\u00c2 al. 2005, 2013), and clustering (Massi et\u00c2 al. 2000). The continuum submillimeter emission in the VMR-D cloud was surveyed by Massi et\u00c2 al. (2007), who catalogued 29 resolved dust cores and also obtained a further list of 26 unresolved candidate cores. More recently, thanks to the opportunity offered by the Spitzer Space Telescope, the VMR-D region was observed with the IRAC (\u00ce\u00bb = 3.6, 4.5, 5.8, 8.0\u00e2\u0080\u0089\u00ce\u00bcm) and MIPS (\u00ce\u00bb = 24, 70\u00e2\u0080\u0089\u00ce\u00bcm) focal-plane instruments, obtaining in this way six mosaics, covering about 1.2\u00c2 deg2, which have been analyzed to produce a merged photometric Spitzer-IRAC point-source catalog (hereinafter Spitzer-PSC) complemented with MIPS photometry (Strafella et\u00c2 al. 2010, hereinafter Paper\u00c2 I). Further observational progress was made when the BLAST experiment (Pascale et\u00c2 al. 2008) mapped the whole Vela Molecular Ridge in the far-IR (FIR) spectral region (\u00ce\u00bb = 250, 350, 500\u00e2\u0080\u0089\u00ce\u00bcm), complementing in this way the Spitzer spectral coverage toward long wavelengths. These observations were discussed by Olmi et\u00c2 al. (2009), who obtained a catalog of dense cores\/clumps in the VMR-D cloud. The last important observational progress was made with the Herschel Space Observatory, which, surveying the Galactic plane in the framework of the Hi-GAL key project, partially mapped the VMR-D region in the FIR spectral range (\u00ce\u00bb = 70, 160, 250, 350, 500\u00e2\u0080\u0089\u00ce\u00bcm) with a spatial resolution and sensitivity almost twice that of BLAST. Here we also analyze these observations for the first time, thanks to the support of the Hi-GAL collaboration that provided us with the corresponding calibrated maps. These have been used to extract five single-band photometries that constitute another important spectral extension of our information about this region.","Citation Text":["Massi et\u00c2 al. 2000"],"Functions Text":["For this reason a subregion of this cloud has been the subject of many previous papers, dealing with different observational aspects of the star formation (SF), such as","and clustering"],"Functions Label":["Background","Background"],"Citation Start End":[[623,641]],"Functions Start End":[[306,474],[607,621]]}
{"Identifier":"2018AandA...615L..16F__Gal_et_al._(2014)_Instance_1","Paragraph":"Nitrogen chemistry mainly consists of three competing processes: (i) the conversion of atomic N to N2 in the gas phase, (ii) destruction of N2, for instance, via photodissociation and reaction with He+, and (iii) freeze-out of atomic N and N2 onto dust grains followed by surface reactions (e.g., Daranlot et al. 2012; Li et al. 2013). The conversion of atomic N into N2 has been proposed to occur by slow neutral-neutral reactions, such as NO + N and CN + N (Herbst & Klemperer 1973; Daranlot et al. 2012). According to the pseudo-time-dependent gas-phase astrochemical model under dense cloud conditions (104 cm\u22123, 10 K, 10 mag) by Le Gal et al. (2014), the conversion of atomic N into N2 takes an order of Myr, depending on assumed elemental abundances. In the gas-ice model of Daranlot et al. (2012), under the similar physical conditions, the conversion of atomic N to N2 takes ~5 \u00d7 105 yr, and it occurs after the significant fraction of nitrogen is frozen out. On the other hand, N2 mainly forms via the reactions NH2 + N and NH + N around the transition from atomic to molecular nitrogen in the models of Furuya & Aikawa (2018) and Furuya & Persson (2018), in which the dynamical evolution of molecular clouds is considered. NH2 and NH are mainly formed via photodesorption of NH3 ice, followed by photodissociation in the gas phase. In this case, the formation rate of N2 from atomic N is, roughly speaking, similar to the freeze-out rate of atomic N. Considering that interstellar ices, at least water ice, are already abundant in molecular clouds with relatively low line-of-sight visual extinction (e.g., ~3 mag for the Taurus dark clouds, Whittet 1993), it may not be surprising that the transition from atomic to molecular nitrogen occurs in the parent cloud of L1544 or in the outer regions of L1544. It should be noted that the N2-dominant region could be larger than the regions traced by N2H+ and NH3 emission, since their abundances are controlled not only by N2, but also by CO; the catastrophic CO freeze-out, which occurs in the late stage of the interstellar ice formation at high densities (\u2273105 cm\u22123; e.g., Pontoppidan 2006), causes their abundances to be enhanced (e.g., Aikawa et al. 2005).","Citation Text":["Le Gal et al. (2014)"],"Functions Text":["According to the pseudo-time-dependent gas-phase astrochemical model under dense cloud conditions (104 cm\u22123, 10 K, 10 mag) by","the conversion of atomic N into N2 takes an order of Myr, depending on assumed elemental abundances. In the gas-ice model of Daranlot et al. (2012), under the similar physical conditions, the conversion of atomic N to N2 takes ~5 \u00d7 105 yr, and it occurs after the significant fraction of nitrogen is frozen out."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[634,654]],"Functions Start End":[[508,633],[656,967]]}
{"Identifier":"2018MNRAS.481..138B__Mullaney_et_al._2013_Instance_1","Paragraph":"To mitigate this problem, several studies have suggested to use the width of the [O\u2009iii]\u03bb5007\u2009\u00c5 emission line (hereafter [O\u2009iii]) originating in the narrow-line region (NLR) as a surrogate for \u03c3\u22c6, assuming that the NLR is gravitationally bound to the bulge and thus, that the gas kinematics follows the bulge potential (e.g. Terlevich, Diaz & Terlevich 1990; Whittle 1992; Nelson & Whittle 1996; Nelson 2000; Boroson 2003; Shields et al. 2003; Greene & Ho 2005; Netzer & Trakhtenbrot 2007; Salviander et al. 2007; Salviander & Shields 2013). However, while the [O\u2009iii] emission line is a prominent line that can be easily measured in AGNs out to large distances, it is also known to often have asymmetric line profiles due to non-gravitational gas kinematics such as outflows, infalls, or interaction with radio jets. In particular, it is known to often display a blue wing (e.g. Heckman et al. 1981; De Robertis & Osterbrock 1984; Whittle 1985; Wilson & Heckman 1985; Mullaney et al. 2013; Woo et al. 2016), generally interpreted as a signature of outflows with dust preferentially hiding one cone behind the stellar disc. For that reason, some studies have excluded the [O\u2009iii] blue wing, as well as any radio sources and galaxies undergoing tidal interactions. The MBH was found to scale with the width of the [O\u2009iii] line (\u03c3[O\u2009iii]), albeit with a large scatter (e.g. Nelson & Whittle 1996; Greene & Ho 2005). Other studies have suggested the use of different emission lines, such as [S\u2009ii]\u03bb\u03bb6716, 6731 (e.g. Komossa & Xu 2007; Ho 2009) that have a lower ionization potential and do not suffer from substantial asymmetries, or mid-infrared lines (e.g. Dasyra et al. 2008, 2011), but the scatter is comparable to that of the core of the [O\u2009iii] line. While all studies confirm the original findings by Nelson & Whittle (1996), i.e. a moderately strong correlation between \u03c3\u22c6 and \u03c3[O\u2009iii] but with real scatter, the origin of the scatter remains unclear. No dependencies have been found with AGN luminosity, host-galaxy morphology, star formation rate, or local environment (Greene & Ho 2005; Rice et al. 2006).","Citation Text":["Mullaney et al. 2013"],"Functions Text":["In particular, it is known to often display a blue wing (e.g.","generally interpreted as a signature of outflows with dust preferentially hiding one cone behind the stellar disc. For that reason, some studies have excluded the [O\u2009iii] blue wing"],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[969,989]],"Functions Start End":[[818,879],[1009,1189]]}
{"Identifier":"2021AandA...653A..36M__Goulding_&_Alexander_(2009)_Instance_6","Paragraph":"The SFG sample was constructed using the Great Observatories All-Sky LIRG Survey (GOALS sample, Armus et al. 2009), from which we extracted 158 galaxies, with data from Inami et al. (2013), who report the fine-structure lines at high resolution in the 10\u2005\u2212\u200536\u2006\u03bcm interval, and Stierwalt et al. (2014), who include the detections of the H2 molecular lines and the PAH features at low spectral resolution. For those galaxies in the GOALS sample that have a single IRAS counterpart, but more than one source detected in the emission lines, we have added together the line or feature fluxes of all components, to consistently associate the correct line or feature emission to the total IR luminosity computed from the IRAS fluxes. To also cover lower luminosity galaxies, as the GOALS sample only includes luminous IR galaxies (LIRGs) and ultra-luminous IR galaxies (ULIRGs), we included 38 galaxies from Bernard-Salas et al. (2009) and Goulding & Alexander (2009), to reach the total sample of 196 galaxies with IR line fluxes in the 5.5\u2005\u2212\u200535\u2006\u03bcm interval in which an AGN component is not detected. For the Bernard-Salas et al. (2009), Goulding & Alexander (2009), and the GOALS samples, we excluded all the composite starburst-AGN objects identified as those with a detection of [NeV] either at 14.3 or 24.3 \u03bcm. It is worth noting that the original samples from Goulding & Alexander (2009) and Bernard-Salas et al. (2009) have spectra solely covering the central region of the galaxies. To estimate the global SFR, we corrected the published line fluxes of the Spitzer spectra by multiplying them by the ratio of the continuum reported in the IRAS point source catalogue to the continuum measured on the Spitzer spectra extracted from the CASSIS database (Lebouteiller et al. 2015). We assumed here that the line emission scales (at first order) with the IR brightness distribution. In particular, we considered the continuum at 12\u2006\u03bcm for the [NeII]12.8\u2006\u03bcm and [NeIII]15.6\u2006\u03bcm lines, and the continuum at 25 \u03bcm for the [OIV]25.9\u2006\u03bcm, [FeII]26\u2006\u03bcm, [SIII]33.5\u2006\u03bcm, and [SiII]34.8\u2006\u03bcm lines. This correction was not needed for the AGN sample and the GOALS sample because of the greater average redshift of the galaxies in the 12MGS and GOALS samples. In particular, the 12MGS active galaxy sample has a mean redshift of 0.028 (Rush et al. 1993), while the GOALS sample has a mean redshift of 0.026. The galaxies presented by Bernard-Salas et al. (2009) have instead an average redshift of 0.0074, while the sample by Goulding & Alexander (2009) has an average redshift of 0.0044. For the other lines in the 10\u2005\u2212\u200536\u2006\u03bcm interval, Goulding & Alexander (2009) did not report a detection, and we used the data presented in Bernard-Salas et al. (2009) for a total of 15 objects. Both Bernard-Salas et al. (2009) and Goulding & Alexander (2009) reported data from the high-resolution Spitzer-IRS spectra. Data in the 50\u2005\u2212\u2005205\u2006\u03bcm interval were taken from D\u00edaz-Santos et al. (2017). For the GOALS sample, 20 starburst galaxies were taken from Fern\u00e1ndez-Ontiveros et al. (2016), and 23 objects were taken from the ISO-LWS observations of Negishi et al. (2001). As a result, we obtained a total sample of 193 objects. Lastly, the PAH features\u2019 fluxes were measured from the low-resolution Spitzer-IRS spectra by Brandl et al. (2006), including 12 objects from the sample of Bernard-Salas et al. (2009) and 179 objects from Stierwalt et al. (2014).","Citation Text":["Goulding & Alexander (2009)"],"Functions Text":["Both Bernard-Salas et al. (2009) and","reported data from the high-resolution Spitzer-IRS spectra."],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[2800,2827]],"Functions Start End":[[2763,2799],[2828,2887]]}
{"Identifier":"2015ApJ...803...79L__G\u00fcsten_et_al._2006_Instance_1","Paragraph":"The number counts of the most massive clusters have the potential to constrain cosmological parameters. Unfortunately, these systems also tend to be most affected by the effects described above. Massive systems will produce the greatest gravitational lensing shear, and in our hierarchical universe, they are also commonly disrupted by recent merging activity. In this work, we aim to better understand how these considerations affect the observed SZE signals using high-resolution submillimeter and radio imaging of a representative sample of SZE-selected clusters.We present new observations at 345 GHz (19 2 resolution) with the Large APEX Bolometer Camera (LABOCA; Siringo et al. 2009) on the Atacama Pathfinder EXperiment (APEX; G\u00fcsten et al. 2006) telescope25\n\n25\nThis publication is based on data acquired with the Atacama Pathfinder Experiment (APEX). APEX is a collaboration between the Max-Planck-Institut f\u00fcr Radioastronomie, the European Southern Observatory, and the Onsala Space Observatory.\n and at 2.1 GHz (5\u2033 resolution) with the Australia Telescope Compact Array (ATCA) of a sample of massive SZE-selected galaxy clusters. We call the project \u201cLASCAR,\u201d the LABOCA\/ACT Survey of Clusters at All Redshifts, in honor of the Lascar volcano near the ACT site in northern Chile. We use these data to measure the properties of the clusters\u2019 spatially resolved SZE increment signals, and quantify the degree of background and foreground radio and infrared galaxy contamination. Section 2 describes our cluster sample. Section 3 presents observations and data reduction techniques. Section 4 assesses the SZE contamination by point sources. Section 5 uses the point-source subtracted multi-wavelength SZE maps to place constraints on cluster peculiar velocities. In Section 6 we discuss our results in the context of previous work, and in Section 7, we conclude. In our calculations, we assume a flat \u039bCDM cosmology with \n\n\n\n\n\n, \n\n\n\n\n\n, and \n\n\n\n\n\n (Komatsu et al. 2011).","Citation Text":["G\u00fcsten et al. 2006"],"Functions Text":["In this work, we aim to better understand how these considerations affect the observed SZE signals using high-resolution submillimeter and radio imaging of a representative sample of SZE-selected clusters.We present new observations at 345 GHz (19 2 resolution) with the Large APEX Bolometer Camera","on the Atacama Pathfinder EXperiment (APEX","telescope"],"Functions Label":["Uses","Uses","Uses"],"Citation Start End":[[734,752]],"Functions Start End":[[361,659],[690,732],[754,763]]}
{"Identifier":"2019ApJ...875...90L__Velli_et_al._2015_Instance_2","Paragraph":"When energy flows from the interior of the Sun outward into the solar atmosphere, why is the Sun\u2019s outer atmosphere, the corona, much hotter than the inner atmosphere, the underlying chromosphere and photosphere? This is the long-standing problem of the coronal heating, which is one of the eight key mysteries in modern astronomy (Kerr 2012). For about 80 yr since the discovery of the extremely hot corona around the late 1930s (Grotian 1939; Edlen 1945), people have worked hard on addressing this issue, and great advances have been made in observation and theoretical studies (Parnell & De Moortel 2012; Amari et al. 2015; Arregui 2015; Cargill et al. 2015; De Moortel & Browning 2015; Jess et al. 2015; Klimchuk 2015; Longcope & Tarr 2015; Peter 2015; Schmelz & Winebarger 2015; Velli et al. 2015; Wilmot-Smith 2015). Especially during recent decades, high-resolution observations of solar super-fine structures indicate that small spicules, minor hot jets along small-scale magnetic channels from the low atmosphere upwards to the corona, petty tornados and cyclones, and small explosive phenomena such as mini-filament eruptions and micro- and nano-flares\u2014all of these small-scale magnetic activities contribute greatly to coronal heating (De Pontieu et al. 2011; 2018; Zhang & Liu 2011; Parnell & De Moortel 2012; Klimchuk 2015; Peter 2015; Schmelz & Winebarger 2015; Henriques et al. 2016; Li et al. 2018a). Additionally, contributions of MHD waves to heating the corona have been observationally illustrated (van Ballegooijen et al. 2011; Jess et al. 2015; Kubo et al. 2016; Morton et al. 2016; Soler et al. 2017; Morgan & Hutton 2018). Meanwhile, with the progress of observational studies, two groups of theoretical models, magnetic reconnection models and magnetohydrodynamic wave models, have traditionally attempted to explain coronal heating, but so far no models can address the key mystery perfectly (van Ballegooijen et al. 2011; Arregui 2015; Cargill et al. 2015; Peter 2015; Velli et al. 2015; Wilmot-Smith 2015). Maybe we do not need to intentionally take to heart such the classical dichotomy, because waves and reconnections may interact with each other (De Moortel & Browning 2015; Velli et al. 2015). Additionally, statistical studies may look at coronal heating from a comprehensive perspective. Li et al. (2018b) found that the long-term variation of the heated corona, which is represented by coronal spectral irradiances, and that of small-scale magnetic activity are in lockstep, indicating that the corona should statistically be effectively heated by small-scale magnetic activity. Observational and theoretical model studies through heating channels and modes, and statistical studies by means of heating effect (performance of the heated corona), both suggest that coronal heating originates from small-scale magnetic activity.","Citation Text":["Velli et al. 2015"],"Functions Text":["Meanwhile, with the progress of observational studies, two groups of theoretical models, magnetic reconnection models and magnetohydrodynamic wave models, have traditionally attempted to explain coronal heating, but so far no models can address the key mystery perfectly"],"Functions Label":["Motivation"],"Citation Start End":[[1997,2014]],"Functions Start End":[[1648,1918]]}
{"Identifier":"2019MNRAS.488.4638L__Drabek-Maunder_et_al._2016_Instance_1","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"],"Functions Text":["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"],"Functions Label":["Differences"],"Citation Start End":[[753,779]],"Functions Start End":[[579,734]]}
{"Identifier":"2018AandA...619A.177B__Mayne_(2010)_Instance_1","Paragraph":"The wealth of precise all-sky data from the Gaia Data Release 2 (DR2) revealed a new feature in the Herzsprung\u2013Russell diagram (HRD), namely a gap in the mid-M dwarf main sequence (Jao et al. 2018). The gap appears at a magnitude MG \u223c 10 and colour GBP \u2013 GRP \u223c 2.3\u22122.5 in the Gaia filter system. It is observed in optical and near-infrared colour-magnitude diagrams (CMDs), indicating that it is not specific to the Gaia photometry and not due to an atmospheric feature that would depend on the wavelength. Jao et al. (2018) suggest the feature is linked to the onset of full convection in M dwarfs. Interestingly, Mayne (2010) was the first to suggest the existence of an observational signature for the transition between fully and partly convective structures for pre-main sequence stars and predicted that it would result in a HRD gap. This author explored signatures of this transition in young clusters and linked the growth of a radiative core to rapid change in effective temperature caused by changes in the dominant energy transport mechanism and ignition of hydrogen burning. Following the announcement of the HRD gap discovery by Jao et al. (2018), MacDonald & Gizis (2018) proposed an explanation based on standard stellar evolution models. They suggest that the observed feature is due to the complex interplay between production of 3He and its transport by convection. More specifically, they predict a fast change in the luminosity over a narrow mass range of \u223c0.31 M\u2299\u22120.34 M\u2299 characterised by the presence of a convective core and a convective envelope that ultimately merge. During this merging process, the central 3He abundance increases, causing an increase in luminosity and thus an observable feature in the luminosity function. In this paper, we re-examine the explanation suggested by MacDonald & Gizis (2018), since at first sight it is not clear why a sudden increase of the luminosity due to the merging of the convective zones of the core and the envelope would create a dip in the luminosity function, and thus a gap in the HRD. In this analysis, we confirm that the best explanation for the observed feature is linked to the property of 3He nuclear production and destruction, and to its mixing. We also find that a change in the energy transport from convection to radiation does not induce structural changes that could be visible. Regarding the very details of the process, however, we disagree with MacDonald & Gizis (2018) and propose an alternative explanation.","Citation Text":["Mayne (2010)"],"Functions Text":["Interestingly,","was the first to suggest the existence of an observational signature for the transition between fully and partly convective structures for pre-main sequence stars and predicted that it would result in a HRD gap. This author explored signatures of this transition in young clusters and linked the growth of a radiative core to rapid change in effective temperature caused by changes in the dominant energy transport mechanism and ignition of hydrogen burning."],"Functions Label":["Background","Background"],"Citation Start End":[[615,627]],"Functions Start End":[[600,614],[628,1086]]}
{"Identifier":"2015AandA...584A..76S__Bernstein_et_al._(1995)_Instance_1","Paragraph":"The only gas-phase process, which is predicted to efficiently lead to products, is the process involving ionized methanimine. According to the model by Vuitton et al. (2007), the amount in the upper atmosphere of Titan of ionized methanimine is small, but not negligible. The products of the reaction CH2NH + CH2NH+ all have a mass-to-charge ratio of 57, where an important contribution is given by the abundant carbocation C4H\\hbox{$_{9}^{+}$}+9. In this condition, it is difficult to see if any of these species is present in small amounts in the ionosphere of Titan. Interestingly, we have also investigated the further reaction of 1+ with a third molecule of methanimine. The formation of the species (CH2NH)\\hbox{$_{3}^{+}$}+3, starting from two molecules of methanimine and the ionic species CH2NH+, is a strongly exothermic reaction, being \\hbox{$\\Delta H_{0}^{\\circ} = -305$}\u0394H0\u25e6=\u2212305 kJ\/mol at CCSD(T) level. The optimized structure of (CH2NH)\\hbox{$_{3}^{+}$}+3 is shown in Fig. 10. We can conclude, therefore, that polymerization of methanimine in the gas-phase at low temperatures may well be initiated by the presence of an ionized molecule. As for the experiment on ice by Bernstein et al. (1995), it is well known that reaction barriers possibly present in gas-phase reactions are not significantly reduced when moving to ice-mediated reactions (see, for instance, Rimola et al. 2014). Normally, the tunneling effect is invoked to explain the observed reactivity, but in this case the reaction barriers are so high that it is difficult to think that a reaction sequence starting with dimerization of neutral methanimine molecules can account for the observed formation of hexamethylenetetramine or polymethylenimine. The reaction must start by involving a radical or an ionized species and not two neutral closed shell molecules. This was already noted by Vinogradoff et al. (2012), who suggested that the reaction between two neutral methanimine molecules is mediated by formic acid, which acts as a proton donor, and by Cottin et al. (2001), who irradiated mixed ice with protons. Notably, Vinogradoff et al. (2013) failed to see hexamethylenetetramine and polymethylenimine formation starting from pure CH2=NH ice. Since in the experiment by Bernstein et al. (1995) the ice was irradiated by VUV photons at the Lyman alpha wavelength, and we now know that methanimine can be efficiently ionized by those photons, we can also argue that ionization of several methanimine molecules can instead trigger the process in cold interstellar ices. This statement is in line with what is known for gas-phase polymerization of olefin (El-Shall 2008; Cottin et al. 2001). Alternatively, external strong energy sources that can induce local nonequilibrium conditions, could promote neutral-neutral dimerization. For instance, in a study by Zhou et al. (2010), in which acetylene ices were irradiated with energetic electrons, electronically excited acetylene molecules were invoked to account for the experimental observation of benzene formation. ","Citation Text":["Bernstein et al. (1995)"],"Functions Text":["As for the experiment on ice by","it is well known that reaction barriers possibly present in gas-phase reactions are not significantly reduced when moving to ice-mediated reactions"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1187,1210]],"Functions Start End":[[1155,1186],[1212,1359]]}
{"Identifier":"2019AandA...629A.134G__Leitherer_et_al._2010_Instance_1","Paragraph":"We used the combination of models from STARBURST99 and our model for the contribution from stripped stars to represent the radiation of a full stellar population in which stripped stars are formed. We made our models publicly available on the STARBURST99 online interface2, providing the addition from stripped stars to the spectral energy distribution, the emission rates of H\u202fI-, He\u202fI-, and He\u202fII-ionizing photons, and the high-resolution UV and optical spectra. In this study, we compare the contribution from stripped stars to the version of STARBURST99 that uses the initial mass function from Kroupa (2001) with mass limits of 0.1\u2006M\u2299 and 100\u2006M\u2299, that is, the same as what we employed for the stripped stars in the population. We chose to compare to the STARBURST99 models that use non-rotating stellar evolutionary models from the Geneva grids (Levesque et al. 2012) and spectral models from WM-BASIC for OB-stars (Pauldrach et al. 2001; Leitherer et al. 2010), CMFGEN for WR stars (Hillier & Miller 1998; Smith et al. 2002), and BASEL v3.1 for later type main sequence stars, cooler stars, and red supergiants (Lejeune et al. 1997). When comparing our models of various metallicity with STARBURST99, we used the STARBURST99 models of Z\u2004=\u20040.014, 0.008, 0.002, and 0.001 together with our models with metallicity of Z\u2004=\u20040.014, 0.006, 0.002, and 0.0002, respectively (see also Table A.1). We consider that the difference in metallicity between the two model sets is small and expect that the difference when using models with exactly the same metallicity is also small. We consider that the combination of the stripped star models and STARBURST99 is a good assumption for radiation with wavelengths shorter than \u223c5000 \u00c5. For the model to be accurate at longer wavelengths, we would need to decrease the radiation from giant stars to compensate for stars that we assume have become stripped. Giant stars emit their radiation primarily at wavelengths longer than \u223c5000 \u00c5. We expect the decrement of radiation at these long wavelengths to be at maximum about 30% (as also suggested by Eldridge et al. 2017, see e.g., their Figs. 5 and 15). This is approximately the fraction of massive stars that get stripped (Sana et al. 2012) and thus the fraction of giant stars that should be missing. This topic is beyond the scope of this paper, but we hope to address it more in detail at a later stage.","Citation Text":["Leitherer et al. 2010"],"Functions Text":["We chose to compare to the","spectral models from WM-BASIC for OB-stars"],"Functions Label":["Uses","Uses"],"Citation Start End":[[944,965]],"Functions Start End":[[732,758],[877,919]]}
{"Identifier":"2019MNRAS.485.3288I__Tsao,_Silberberg_&_Barghouty_1998_Instance_1","Paragraph":"Coming back to the flux values given in Table 1, it is tempting to explain the deficiency of the measured fluxes at energies \u223c68 and \u223c78 keV compared to the fluxes measured at 1.157 MeV as a consequence of an additional component of the flux at 1.157 MeV, which is generated by the interaction of low-energy cosmic rays (LECR) with abundant circumstellar species like Fe, Mn, and Cr in the environment of Cas A. Siegert et al. (2015) attempted to interpret these discrepant line fluxes of ${}^{44}_{}\\mathrm{Ti}$\u2009 at different energies exactly in this way. However, considering that little or nothing is known about the LECR fluxes in the Cas A SNR and uncertainties in the generally small cross-sections for the relevant reactions to produce either ${}^{44}_{}\\mathrm{Ti}$\u2009 or excited ${}^{44}_{}\\mathrm{Ca}^{*}$ (Silberberg, Tsao & Barghouty 1998; Tsao, Silberberg & Barghouty 1998), the validity of such an interpretation of the discrepant measurements of ${}^{44}_{}\\mathrm{Ti}$\u2009 fluxes at 68 and 78 keV and at 1.157 MeV by different instruments and by the same instrument (SPI) may be doubted. Additionally, if one assumes that the production of the 1.157 MeV line emission is supported in Cas A by the enhanced flux of LECRs, then immediately a problem arises related to the absence of other lines emission from much more abundant elements like oxygen, carbon, and nitrogen with very large, well-measured cross-sections for excitation. Indeed, simulations of the excitation emission lines produced by LECRs in Cas A have shown that lines at 4.43, 6.13, 6.9, and 7.1 MeV are expected to be strongest (Summa et al. 2011) for the expected spectrum of LECR (Berezhko et al. 2003) that are interacting with the CSM near the Cas A SNR. However, because no gamma-ray line emission at 4.43, 6.13, 6.9, and 7.1 MeV is observed from the Cas A SNR, we consider this explanation of the excitation origin of the high 1.157 MeV ${}^{44}_{}\\mathrm{Ti}$\u2009 line flux compared to that at 68 and 78 keV as very problematic.","Citation Text":["Tsao, Silberberg & Barghouty 1998"],"Functions Text":["However, considering that little or nothing is known about the LECR fluxes in the Cas A SNR and uncertainties in the generally small cross-sections for the relevant reactions to produce either ${}^{44}_{}\\mathrm{Ti}$\u2009 or excited ${}^{44}_{}\\mathrm{Ca}^{*}$","the validity of such an interpretation of the discrepant measurements of ${}^{44}_{}\\mathrm{Ti}$\u2009 fluxes at 68 and 78 keV and at 1.157 MeV by different instruments and by the same instrument (SPI) may be doubted."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[850,883]],"Functions Start End":[[557,813],[886,1098]]}
{"Identifier":"2015AandA...577A..43S__Odstr\u010dil_&_Karlick\u00fd_(1997)_Instance_1","Paragraph":"The initialization of solar flares remains an unsolved problem. Early ideas on how the initialization might occur were described by Norman & Smith (1978). They argued that flare process cannot start in the entire flare volume at one instant, and proposed that the flare onset was localized in a small part of an active region, from which the energy release extends as dissipation spreading process throughout the flare volume. Two types of agents that may lead to this kind of a dissipation process were addressed: electron beams and shock waves. These agents can trigger flares at large distances from their initial locations, causing sympathetic (simultaneous) flares or leading to a sequential flare energy release in one active region (Liu et al. 2009; Zuccarello et al. 2009). These triggering processes were numerically studied by Karlick\u00fd & Jungwirth (1989) and Odstr\u010dil & Karlick\u00fd (1997). Karlick\u00fd & Jungwirth (1989) assumed that electron beams, penetrating into the current sheet in the magnetic reconnection region, generate Langmuir waves. Then, using the particle-in-cell model, the authors studied the effects of these electrostatic waves on the plasma system. Sufficiently strong Langmuir waves were found to be able to generate ion-sound waves through the three-wave decay process (B\u00e1rta & Karlick\u00fd 2000). These ion-sound waves increase electrical resistivity in the current sheet system, which results in the onset of the energy dissipation. Thus, the electron beams are able to cause magnetic reconnection. Odstr\u010dil & Karlick\u00fd (1997) studied the mechanism for the flare trigger by shock waves. They used a 2D magnetohydrodynamic model with the MHD shock wave propagating towards the current sheet. A portion of the shock wave passed through the sheet, and the rest was reflected. Nothing occurred at the very beginning of the wave-current sheet interaction. However, after some time, specific plasma flows around the current sheet were formed, which led to the start of magnetic reconnection. This shows that for reconnection to be triggered, the enhanced electrical resistivity as well as the plasma flows are important. ","Citation Text":["Odstr\u010dil & Karlick\u00fd (1997)"],"Functions Text":["These triggering processes were numerically studied by Karlick\u00fd & Jungwirth (1989) and"],"Functions Label":["Background"],"Citation Start End":[[869,895]],"Functions Start End":[[782,868]]}
{"Identifier":"2019MNRAS.482.3656R__Kashi_&_Soker_2011_Instance_1","Paragraph":"Several evolutionary channels have been proposed that lead to a SNIa explosion. For a comprehensive review, see Livio & Mazzali (2018) and Wang (2018). Among these, the two classical scenarios are the single- and the double-degenerate channels. In the single-degenerate channel a white dwarf (WD) in a binary system accretes mass from a non-degenerate donor until it grows near the Chandrasekhar limit (Whelan & Iben 1973; Han & Podsiadlowski 2004; Nomoto & Leung 2018). In the double-degenerate channel two WDs in a close binary system merge due to angular momentum loss caused by the emission of gravitational waves (GWs) and the resulting merger has a mass near the Chandrasekhar limit (Whelan & Iben 1973; Iben & Tutukov 1984; Liu, Wang & Han 2018). Additional evolutionary channels for SNIa include the double-detonation mechanism (Woosley & Weaver 1986; Livne & Arnett 1995; Shen et al. 2012), the violent merger model (Pakmor et al. 2010; Sato et al. 2016), the core-degenerate channel (Sparks & Stecher 1974; Livio & Riess 2003; Kashi & Soker 2011; Wang et al. 2017) and a mechanism which involves the collision of two WDs (Benz, Thielemann & Hills 1989; Aznar-Sigu\u00e1n et al. 2013; Kushnir et al. 2013). In the double-detonation scenario a WD accumulates helium-rich material on its surface, which is compressed and ultimately detonates. The compression wave propagates towards the centre of the WD and a second detonation occurs near the centre of its carbon--oxygen core. In the violent merger model, the detonation of the WD core is initiated during the early stages of the merger. This can happen, for example, due to compressional heating by accretion from the disrupted secondary or due to a preceeding detonation of accreted helium (alike the double-detonation scenario) that is ignited dynamically (Guillochon et al. 2010; Pakmor et al. 2010, 2011, 2012, 2013; Kashyap et al. 2015; Sato et al. 2015, 2016). In the core-degenerate scenario a WD merges with the hot core of an asymptotic giant branch star during (or after) a common envelope (CE) phase. Finally, the evolutionary phase involving the collision of two WDs requires a tertiary star which brings the two WDs to collide due to the Kozai--Lidov mechanism, or dynamical interactions in a dense stellar system, where this kind of interaction is more likely to happen.","Citation Text":["Kashi & Soker 2011"],"Functions Text":["Additional evolutionary channels for SNIa include","the core-degenerate channel"],"Functions Label":["Background","Background"],"Citation Start End":[[1037,1055]],"Functions Start End":[[754,803],[965,992]]}
{"Identifier":"2020MNRAS.497.4231R__Brown,_Vasil_&_Zweibel_2012_Instance_1","Paragraph":"Application of linear theory to investigate IGWs has been done extensively in the field of oceanography (Sutherland 2010). The general idea is to linearize hydrodynamical equations and introduce a wave-like ansatz to solve the linearized equations. In the context of stellar parameters, one of the earliest works was done by Press (1981), which focused on solar-like stars. In our work, we follow similar steps, starting with the linearized hydrodynamical equations in the anelastic approximation:\n(8)$$\\begin{eqnarray*}\r\n\\nabla \\cdot \\overline{\\rho } \\boldsymbol {v} = 0,\r\n\\end{eqnarray*}$$(9)$$\\begin{eqnarray*}\r\n\\frac{\\partial \\boldsymbol {v}}{\\partial t} = - \\nabla \\left(\\frac{P}{\\overline{\\rho }}\\right) - C \\overline{g} \\boldsymbol {\\hat{r}}\r\n\\end{eqnarray*}$$(10)$$\\begin{eqnarray*}\r\n\\frac{\\partial T}{\\partial t} = - v_\\mathrm{ r} \\left(\\frac{\\partial \\overline{T}}{\\partial r} - (\\gamma - 1) \\overline{T} h_\\rho \\right)\r\n\\end{eqnarray*}$$We ignore the pressure term in equation (4), which leads to a set of equations that conserve energy (Brown, Vasil & Zweibel 2012). We also do not consider rotational, thermal diffusion, or viscous effects. In cylindrical coordinates, this 2D analysis allows us to set the z-derivatives to zero. The three equations shown above can then be reduced to one second-order differential equation with vr(r) as the evolving term:\n(11)$$\\begin{eqnarray*}\r\n0 &amp;=&amp; \\frac{\\partial ^2 \\alpha }{\\partial r^2} + \\left(\\frac{N^2}{\\omega ^2} - 1 \\right) \\frac{m^2}{r^2} \\alpha + \\left[ -\\overline{\\rho }^{-1\/2}\\frac{\\partial ^2 \\left(\\overline{\\rho }^{1\/2}\\right)}{\\partial r^2} + \\frac{\\partial h_{\\rho }}{\\partial r} \\right]\\alpha \\nonumber \\\\\r\n&amp;&amp;+\\, \\frac{1}{4r^2} \\alpha,\r\n\\end{eqnarray*}$$where $\\alpha = v_\\mathrm{ r} \\overline{\\rho }^{1\/2} r^{3\/2}$. We have used a wave ansatz of the form vr(r, \u03b8, z) \u221d vr(r)eim\u03b8e\u2212i\u03c9t, where m is the horizontal wavenumber,1 \u03c9 is the angular frequency, and \u03b8 is the angular coordinate in the equatorial plane. Generally, in the radiation zone, there will be regions where the oscillatory term (OT), (N2\/\u03c92 \u2212 1)m2\/r2, dominates and regions where the density term (DT), $[ -\\overline{\\rho }^{-1\/2}\\partial ^2 (\\overline{\\rho }^{1\/2})\/\\partial r^2 + \\partial h_{\\rho }\/\\partial r]$, dominates. When the ratio of OT to DT is less than 1, an IGW loses its wave-like behaviour and the approximate radius where this ratio is exactly equal to 1 is called the turning point. The importance of this will be discussed in Section 4.","Citation Text":["Brown, Vasil & Zweibel 2012"],"Functions Text":["We ignore the pressure term in equation (4), which leads to a set of equations that conserve energy"],"Functions Label":["Uses"],"Citation Start End":[[1049,1076]],"Functions Start End":[[948,1047]]}
{"Identifier":"2019ApJ...886...14F__Urquhart_et_al._2014_Instance_1","Paragraph":"Figure 6(a) shows the 1.3 mm continuum and C18O(J = 2\u20131) distribution around the Papillon Nebula. The extended 1.3 mm continuum emission around the Papillon Nebula YSO roughly agrees with the distribution of H\u03b1 emission and has no counterparts of molecular gas, indicating that the continuum emission seems to be dominated by free\u2013free emission from the ionized gas (see also Paper I). On the other hand, the filamentary dust clumps along the C18O emission, a tracer of cold\/dense molecular gas, are considered to be dominated by thermal dust emission. Two major local maxima, hereafter MMS-1 and MMS-2, are found in the 1.3 mm continuum image in Figures 6(a)\u2013(c). Because we detected extended emission in 1.3 mm and C18O along the north\u2013south direction, the two sources are considered to be dense cores embedded in the filamentary cloud. The peak column densities and total masses of both sources deduced from the 1.3 mm continuum emission are \u223c1 \u00d7 1024 cm\u22122, and \u223c2 \u00d7 102 M\u2299, respectively, with an assumption of the dust emissivity, \n\n\n\n\n\n of 1 cm2 g\u22121 for protostellar envelopes (e.g., Ossenkopf & Henning 1994), a dust-to-gas ratio of 3.0 \u00d7 10\u22123 (Herrera et al. 2013; Gordon et al. 2014), and a uniform dust temperature of 20 K, which is typically adapted\/estimated for galactic high-mass star-forming dust clumps (e.g., Urquhart et al. 2014; Yuan et al. 2017). Note that the core boundaries deriving their total masses were set as the 30% intensity level of each continuum peak to avoid contaminations from the parental filamentary clouds. Although these sources were not cataloged as point sources in the infrared observations (e.g., Chen et al. 2010; Carlson et al. 2012), their detection in the 1.3 mm is strongly suggestive of their protostellar nature. We have identified outflow wings that have a velocity span of \u223c30 km s\u22121 in 12CO(J = 2\u20131) toward MMS-1 and MMS-2. Figures 6(b)\u2013(e) show the distributions and the line profiles of the outflow wings. We extracted the spectra and chose the velocity ranges using the following procedure. We defined the outer boundaries of the outflow spectra (i.e., maximum relative velocity) above the 2\u03c3 level of the velocity-smoothed spectra with a smoothing kernel of 9 ch (=1.8 km s\u22121). We extracted the spectra at positions \u223c1\u2033 away from the millimeter sources in the east direction as the references without the outflow contaminations and defined the velocities below 2\u03c3 intensity levels as the inner edge of the outflow spectra (see Figures 6(d) and (e)). The outflow and YSO properties are listed in Table 2. The size of the outflow is as small as \u223c0.1 pc, and escaped detection with our lower-resolution study (Paper I). It is likely that the launching points of the outflows coincide with the 1.3 mm continuum peaks. Because the positional offsets between the blueshifted and redshifted lobes are quite small toward MMS-1 and MMS-2, the outflow orientations may be close to pole-on or very small. The dynamical time (td) of the outflows is roughly estimated to be 104 yr from a ratio of 0.1 pc\/20 km s\u22121. We roughly estimated the mechanical forces of the outflow lobes (Fout) using the following equation: Fout = MoutVout\/td (e.g., Beuther et al. 2002), where Vout and Mout are the outflow velocity and mass, respectively. The estimated Fout is \u223c10\u22123\u201310\u22122 M\u2299 km s\u22121 yr\u22121, depending on the assumption of the outflow inclination angle (i = 30\u00b0\u201370\u00b0). The Fout and the envelope mass derived from millimeter dust continuum observations are consistent with those of Galactic high-mass protostars (e.g., Beuther et al. 2002) and the N159W-South region (TFH19). A large amount of the surrounding gas with a mass of \u223c102 M\u2299 and the nondetection of infrared emission suggest that MMS-1 and MMS-2 are in an extremely young phase of high-mass star formation, with an age of \u2272104 yr. We note that the Papillon Nebula YSO is more evolved than the two sources; however, the age is as young as \u223c0.1 Myr (Paper I) judging from the combination of the spectral energy distribution fitting using the Spitzer\/Herschel data and the compactness of the H i region (see also Chen et al. 2010). Although certain time differences are found, the three high-mass protostellar systems with separations of \u223c1\u20132 pc are formed within the order of \u223c0.1 Myr. We discuss the possible processes of high-mass star formation taking place in the N159E-Papillon region in Section 4.1.","Citation Text":["Urquhart et al. 2014"],"Functions Text":["The peak column densities and total masses of both sources deduced from the 1.3 mm continuum emission are \u223c1 \u00d7 1024 cm\u22122, and \u223c2 \u00d7 102 M\u2299, respectively, with an assumption","and a uniform dust temperature of 20 K, which is typically adapted\/estimated for galactic high-mass star-forming dust clumps (e.g.,"],"Functions Label":["Uses","Uses"],"Citation Start End":[[1325,1345]],"Functions Start End":[[839,1010],[1193,1324]]}
{"Identifier":"2017ApJ...850...97B__Tamburro_et_al._2009_Instance_2","Paragraph":"The H i mass fraction of every gas particle in the baryonic runs is calculated based on the particle\u2019s temperature and density and the cosmic UV background radiation flux while including a prescription for self-shielding of H2 and dust shielding in both H i and H2 (Christensen et al. 2012). This allows for the straightforward calculation of the total H i mass of each simulated galaxy. We create mock H i data cubes only for the 42 halos that contain \n\n\n\n\n\n. Specifically, we create mock data cubes that mimic ALFALFA observations (Haynes et al. 2011). After specifying a viewing angle (see below), our code considers the line-of-sight velocity of each gas particle. The velocity of each particle is tracked in the simulation by solving Newton\u2019s equations of motion, but any turbulent velocity of the gas is not taken into account. Velocity dispersions in dwarf galaxies can be on the order of the rotational velocity, \u223c10\u201315 km s\u22121 (e.g., Stanimirovi\u0107 et al. 2004; Tamburro et al. 2009; Oh et al. 2015). Dispersions are thought to be driven at least partially by thermal velocities or supernovae (Tamburro et al. 2009; Stilp et al. 2013a, 2013b). In our simulations, supernovae inject thermal energy, and the thermal state of the H i gas needs to be considered in the mock H i linewidth for a realistic comparison to observations. To account for the thermal velocity, the H i mass of each gas particle is assumed to be distributed along the line-of-sight in a Gaussian distribution with a standard deviation given by the thermal velocity dispersion, \n\n\n\n\n\n, where T is the temperature of the gas particle. After this thermal broadening is calculated, a mock H i data cube can be generated by specifying the spatial and velocity resolution. For all of our galaxies, we adopt a spatial resolution of 54 pixels across 2Rvir. In practice, this corresponds to a range of \u223c1 kpc resolution in our lowest-mass galaxies up to \u223c9 kpc resolution in our most massive galaxies. However, the spatial resolution plays no role in our study, since measurements of the VF are based on spatially unresolved H i data. For the velocity resolution, we match the ALFALFA specification of 11.2 km s\u22121 (two-channel boxcar-smoothed).","Citation Text":["Tamburro et al. 2009"],"Functions Text":["Dispersions are thought to be driven at least partially by thermal velocities or supernovae"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1100,1120]],"Functions Start End":[[1007,1098]]}
{"Identifier":"2015ApJ...801...88S__Liu_et_al._2008_Instance_1","Paragraph":"There are several different effects that have been considered to explain the observed offset in the [O iii]\u03bb5007\/H\u03b2 vs. [N ii]\u03bb6584\/H\u03b1 BPT diagram among high-redshift galaxies. First, there are the physical parameters describing the H ii regions contributing to the integrated line ratios from galaxies. These include the ionization parameter, ionizing spectrum of the stars illuminating the H ii-region gas, and the electron density. Early work highlighting the issue of the high-redshift offset in the [O iii]\u03bb5007\/H\u03b2 vs. [N ii]\u03bb6584\/H\u03b1 BPT diagram focused on these parameters (e.g., Shapley et al. 2005; Brinchmann et al. 2008; Liu et al. 2008), and they have been revisited more recently by Kewley et al. (2013), Steidel et al. (2014), and Masters et al. (2014). Systematically higher ionization parameters (which appear to apply in high-redshift galaxies; Nakajima et al. 2013), harder ionizing spectra, and higher electron densities (Shirazi et al. 2014), all tend to shift the locus of galaxies in the [O iii]\u03bb5007\/H\u03b2 vs. [N ii]\u03bb6584\/H\u03b1 diagram toward higher [O iii]\/H\u03b2 and [N ii]\/H\u03b1 values. Next, there is the role of different types of pressure in determining the internal structure and dynamics of H ii regions. Yeh et al. (2013) and Verdolini et al. (2013) suggest that radiation pressure is significant in high-redshift H ii regions, as compared with gas pressure associated with stellar winds, and that the effects of radiation pressure can lead to [O iii]\/H\u03b2 line ratios in excess of the \u201cmaximum starburst\u201d limit of Kewley et al. (2001). Possible contamination by weak AGNs has also been suggested as a way to shift galaxy emission-line ratios into the \u201ccomposite\u201d region of the [O iii]\u03bb5007\/H\u03b2 vs. [N ii]\u03bb6584\/H\u03b1 BPT diagram (e.g., Wright et al. 2010), in between the curves of Kauffmann et al. (2003) and Kewley et al. (2001). Finally, both Masters et al. (2014) and Steidel et al. (2014) consider gas-phase abundance ratios\u2014specifically, the N\/O ratio\u2014which can affect where galaxies fall in the [O iii]\u03bb5007\/H\u03b2 vs. [N ii]\u03bb6584\/H\u03b1 BPT diagram. If the relationship between N\/O and \n\n\n\n\n\n evolves out to high redshift, then distant galaxies will shift relative to local ones in the [O iii]\u03bb5007\/H\u03b2 vs. [N ii]\u03bb6584\/H\u03b1 BPT diagram.","Citation Text":["Liu et al. 2008"],"Functions Text":["Early work highlighting the issue of the high-redshift offset in the [O iii]\u03bb5007\/H\u03b2 vs. [N ii]\u03bb6584\/H\u03b1 BPT diagram focused on these parameters (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[631,646]],"Functions Start End":[[435,585]]}
{"Identifier":"2022MNRAS.510.3479B__Perucho_et_al._2010_Instance_1","Paragraph":"Our simulations, despite probing quite different jet velocities and resolution levels, point to a very unstable nature of the structure resulting from the jet\u2013wind\u2013orbit interaction in HMMQ. Although this was already suggested, our numerical calculations strongly endorse this possibility. Our results, not being an extensive exploration but just two characteristic examples, do not settle the issue for other parameter choices. Nevertheless, the simulations probe fiducial cases and their outcomes suggest that the jet is unlikely to leave unscathed the region close to the binary, say $z\\lesssim 10\\, a$. Even for smaller \u03c7j values and longer Porb values than those explored here, we propose that jet\u2013wind\u2013orbit interaction could induce instability growth, which could produce precession\u2013like patterns in the jets at larger scale. Even if the jets are so powerful that significant mixing and disruption does not occur on scales $\\sim 1-10\\, a$, and \u03a6 is initially rather small, the growth of Kelvin\u2013Helmholtz (KH) instabilities is expected (see also, e.g. Perucho 2019, in the context of extragalactic jets). In fact, the jet\u2013wind interaction alone can already perturb the jet (Perucho & Bosch-Ramon 2008; Perucho et al. 2010), and stellar wind clumping, not considered in the present simulations, should also render the jet more prone to disruption (Perucho & Bosch-Ramon 2012). In addition to KH instabilities, Rayleigh\u2013Taylor and Richtmyer\u2013Meshkov instabilities could also develop at the dynamical jet\u2013wind contact discontinuity, particularly due to the Coriolis force, as in the similar case studied by Bosch-Ramon et al. (2015) with a pulsar wind instead of a jet. These additional instabilities would produce non-linear perturbations that would couple to the KH instability, enhancing the disruptive effects of all these processes.5 We note that jet precession-like features caused by KH instabilities may or may not follow the orbital period, as these instabilities could span a broad range of wavelengths. Therefore, in addition to the cases explored here, \u03c7j values even smaller than \u03b8j (see Bosch-Ramon & Barkov 2016, for a discussion on the \u03c7j \u2013 \u03b8j relation) could also lead to distorted jets on scales \u226ba. Interestingly, this has been observed in radio in Cyg X-3, but not yet in Cyg X-1 (see,e.g. Mioduszewski et al. 2001; Miller-Jones et al. 2004; Tudose et al. 2007, for related radio observations of Cyg X-3). Interestingly, regardless of the \u03c7j value, if the jet gets disrupted on scales $\\lesssim 10\\, a$, the strong pressure drop outwards of the embedding medium can make the jet\u2013wind mixed flow turn into a collimated structure (see Bosch-Ramon & Barkov 2016; see also Millas et al. 2019 for the case of the HMMQ SS 433, a different but related situation -see Section 1-). In the case of a weaker wind-orbit effect on the jet, the jet inertia, and the density drop of the medium, may allow the jet to keep some coherence despite KH instability growth, at least until interacting with the surrounding medium on much larger scales (see, e.g. Marti et al. 1996; Gallo et al. 2005; Russell et al. 2007; Sell et al. 2015 for Cyg X-1 radio observations, and Bordas et al. 2009; Bosch-Ramon et al. 2011; Yoon et al. 2011 for simulations, of the jet-medium interaction). To properly study the evolution of orbit-affected jets at z \u226b a in HMMQ, devoted simulations are planned.","Citation Text":["Perucho et al. 2010"],"Functions Text":["In fact, the jet\u2013wind interaction alone can already perturb the jet"],"Functions Label":["Background"],"Citation Start End":[[1208,1227]],"Functions Start End":[[1111,1178]]}
{"Identifier":"2016AandA...589A.132B__Sarangi_&_Cherchneff_2013_Instance_1","Paragraph":"Extensive modelling efforts have been undertaken to explain observations and the formation of dust in local and high redshift SNe. While some studies use the classical nucleation theory (CNT) to describe dust formation under equilibrium conditions in shocked environments (Kozasa et al. 1989; Todini & Ferrara 2001; Nozawa et al. 2003; Schneider et al. 2004), other studies describe the formation of both molecules and dust clusters from the shocked gas by assuming a chemical kinetic approach under non-equilibrium conditions and the subsequent coalescence and coagulation of these clusters to form dust grains (Cherchneff & Lilly 2008; Cherchneff & Dwek 2009; 2010; Sarangi & Cherchneff 2013; 2015, hereafter SC15). On the other hand, the processing of dust in SNRs has received less attention. The thermal sputtering of dust by the forward shock has been studied by Nozawa et al. (2006), and sputtering by the RS in local and primitive SNe was modelled by Nozawa et al. (2007), Bianchi & Schneider (2007), and Nath et al. (2008). Sputtering by the RS in Type II-b SNe with application to the Cas A SNR was studied by Nozawa et al. (2010). Finally, Silvia et al. (2010; 2012) proposed a hydrodynamic study of an ejecta clump, which was processed by the RS, and followed thermal sputtering as the clump was crossed and gradually disrupted by the shock. All studies show that dust is strongly reprocessed in SNRs. However, the final mass of surviving dust is not well constrained because of several factors. Firstly, all studies use as initial conditions for sputtering in the SNR the dust masses and size distributions that result from applying CNT. The chemical types and masses of the dust that are used as initial conditions may consequently be in error using this formalism, as discussed by Cherchneff (2014). Furthermore, some of these studies assume a homogeneous ejecta in the SNR phase (Nozawa et al. 2007; 2010; Bianchi & Schneider 2007; Nath 2008), which implies that very harsh conditions hold in the post-shock gas for dust survival because the RS crosses and reprocesses the ejecta at very high velocities. By considering clumpy ejecta, the RS velocity will be greatly reduced in the dense clumps (Silvia et al. 2010; Biscaro & Cherchneff 2014). ","Citation Text":["Sarangi & Cherchneff 2013"],"Functions Text":["other studies describe the formation of both molecules and dust clusters from the shocked gas by assuming a chemical kinetic approach under non-equilibrium conditions and the subsequent coalescence and coagulation of these clusters to form dust grains"],"Functions Label":["Background"],"Citation Start End":[[668,693]],"Functions Start End":[[360,611]]}
{"Identifier":"2018ApJ...866...48U__Plagge_et_al._2013_Instance_1","Paragraph":"RX J1347.5\u20131145 is one of the most luminous X-ray galaxy clusters and is located at a redshift of z = 0.451. It was thought to be a relaxed cluster when it was discovered in the ROSAT all sky survey (Schindler et al. 1997). Komatsu et al. (1999) made the first measurements of the Sunyaev\u2013Zel\u2019dovich effect (SZE: Sunyaev & Zeldovich 1972) toward this cluster with the James Clerk Maxwell Telescope at 350 GHz as well as with the 45 m Nobeyama Radio Telescope at 21 and 43 GHz. A higher angular resolution observation of the SZE was performed by Komatsu et al. (2001) using the Nobeyama Bolometer Array and they found a prominent substructure which has no counterpart in the soft X-ray image from ROSAT. The presence of the substructure has been confirmed by Chandra and XMM-Newton (e.g., Allen et al. 2002; Gitti & Schindler 2004) as well as by more recent SZE measurements (Mason et al. 2010; Korngut et al. 2011; Plagge et al. 2013; Adam et al. 2014; Kitayama et al. 2016). Allen et al. (2002) measured the mean temperature of the ICM to be over 10 keV, which is relatively high compared to other typical clusters. Kitayama et al. (2004) and Ota et al. (2008) found a very hot (>20 keV) component of the ICM in this cluster. In addition, the radial profile and spatial distribution of the ICM temperature indicate that the temperature drops to \u223c6 keV toward the cluster center so that the cool core is formed (e.g., Allen et al. 2002; Ota et al. 2008; Kreisch et al. 2016). A disturbed morphology is further supported by radio synchrotron observations (e.g., Ferrari et al. 2011) and gravitational lensing maps (e.g., K\u00f6hlinger & Schmidt 2014). The total mass of RX J1347.5\u20131145 within r200 is estimated to be \u223c1.5 \u00d7 1015 h\u22121 \n\n\n\n\n\n  using weak-lensing analysis, where r200, the radius within which the mean mass density is 200 times the critical density of the universe, is 1.85 h\u22121 Mpc (Lu et al. 2010) for this galaxy cluster.18\n\n18\nThey adopted the Hubble constant of 70 km s\u22121 Mpc\u22121.\n\n","Citation Text":["Plagge et al. 2013"],"Functions Text":["The presence of the substructure has been confirmed by Chandra and XMM-Newton","as well as by more recent SZE measurements"],"Functions Label":["Background","Background"],"Citation Start End":[[915,933]],"Functions Start End":[[703,780],[831,873]]}
{"Identifier":"2021ApJ...912..106Y__Minchev_et_al._2013_Instance_1","Paragraph":"Our analysis on the LAMOST-RC stars by dissecting the MAPs shows that the chemical bimodality is observed throughout the Galactic disk, and the high- and low-[\u03b1\/Fe] sequences are corresponding to the thick and thin disks of the Milky Way, respectively. How to explain the formation mechanism of the stellar thin and thick disks is beyond the scope of this paper, but our results provide some observational constraints to the model of the chemodynamical evolution of the Milky Way disk. Our flared vertical profiles for the thin and thick disks are in good agreement with the prediction of the thin+thick flaring disk model (L\u00f3pez-Corredoira & Molg\u00f3 2014), and are consistent with the number simulations of the chemodynamical evolution in Galactic disks formed in the cosmological context (Minchev et al. 2013, 2014, 2015, 2017), as well as the cosmological zoom simulation of VINTERGATAN (Agertz et al. 2021). These model simulations suggest that the vertical flaring trends are a natural consequence of inside-out, upside down growth coupled with disk flaring (see also Bird et al. 2013; Garc\u00eda de la Cruz et al. 2021), which allows for the low-[\u03b1\/Fe] stars to exist several kpc above the disk\u2019s midplane. As analyzed by B16, the exponential flaring profiles for the low-[\u03b1\/Fe] MAPs suggests that radial migration played an important role in the formation and evolution of the thin disk. Radial migration of stars via cold torquing, also known as \u201cchurning,\u201d by a bar and spiral waves (Minchev et al. 2013) then allows for the populations to spatially overlap in the solar neighborhood. Similar to the flared thin disk, the flaring profile for the high-[\u03b1\/Fe] MAP indicates the radial migration has occurred in the formation of the thick disk as suggested by model simulations (e.g., Sch\u00f6nrich & Binney 2009; Minchev et al. 2015; Li et al. 2018). Of course, we cannot rule out the other formation scenarios of the thick disk, such as the accreted gas from satellites (Brook et al. 2004), accreted stars from galaxy mergers (Abadi et al. 2003), or from disk-crossing satellites heating up the thin disk (Read et al. 2008). On the other hand, the broken exponential radial profiles for the thin and thick disks cannot be explained by any model of the galactic disks. In fact, nearly all the models we mentioned above present a single-exponential profile decreasing with the increasing of R (e.g., Minchev et al. 2015; Li et al. 2018; Agertz et al. 2021). And the smooth downtrend of radial profile in the outer disk R > Rpeak, as shown in Figure 10, means that there is no cut-off of the stellar component at R = 14\u201315 kpc as stated by Ruphy et al. (1996), which is also discovered by B16.","Citation Text":["Minchev et al. 2013"],"Functions Text":["Our flared vertical profiles for the thin and thick disks","and are consistent with the number simulations of the chemodynamical evolution in Galactic disks formed in the cosmological context"],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[789,808]],"Functions Start End":[[486,543],[656,787]]}
{"Identifier":"2020AandA...640A.121G__Nakajima_1983_Instance_1","Paragraph":"The planetary physical properties that can be derived from measurements with the two most successful exoplanet detection techniques, that is, with the radial velocity method and the transit method, are the planet radius (if the planet transits its star), mass (a lower limit if the planet does not transit its star), and the orbital period and distance, eccentricity, and inclination (if the planet transits its star) (see Perryman 2018, and references therein). The upper atmospheres (above optically thick clouds) of giant exoplanets and those of smaller planets around M stars can be probed using spectroscopy during primary and\/or secondary transits (Bean et al. 2010; Swain et al. 2009, 2008; Tinetti et al. 2007; Nakajima 1983). It appears to be virtually impossible to characterize the (lower) atmospheresand surfaces of small, Earth-like planets in the habitable zones of solar-type stars (B\u00e9tr\u00e9mieux & Kaltenegger 2014; Misra et al. 2014; Kaltenegger & Traub 2009), in particular because the light of the parent star is refracted while traveling through the lower atmosphere of its planet and emerges forever out of reach of terrestrial telescopes (Garc\u00eda Mu\u00f1oz et al. 2012). The (lower) atmosphere and surface of a planet are crucial for determining the habitability of a planet, as they hold information about cloud composition, trace gases in disequilibrium and probably most importantly, liquid surface water (see, e.g., Schwieterman et al. 2018; Kiang et al. 2007a,b, and references therein). For such a characterization of terrestrial-type planets, direct observations of the thermal radiation that they emit or of the light of their parent star that they reflect are required. The numerical results that we present in this paper concern the reflected starlight. Because of the huge distances involved, any measured reflected starlight pertains to the (illuminated and visible part of the) planetary disk. It therefore is a disk-integrated signal.","Citation Text":["Nakajima 1983"],"Functions Text":["The upper atmospheres (above optically thick clouds) of giant exoplanets and those of smaller planets around M stars can be probed using spectroscopy during primary and\/or secondary transits"],"Functions Label":["Background"],"Citation Start End":[[719,732]],"Functions Start End":[[463,653]]}
{"Identifier":"2018AandA...617A.116L__Hendecourt_&_Allamandola_(1986)_Instance_1","Paragraph":"Information about the composition and structure of astrophysical ices, and in particular the presence of solid methanol, is obtained by comparing astrophysical observations in the infrared to laboratory data (Dartois et al. 1999; Pontoppidan et al. 2003; Boogert et al. 2008; Bottinelli et al. 2010). The main infrared absorption features used for solid CH3OH identification in interstellar ices are the 3.53 \u03bcm (2833 cm\u22121) and 9.74 \u03bcm (1027 cm\u22121) bands. They correspond to the CH3 symmetric stretching mode and the CO\nstretching mode, respectively. The laboratory data most widely used for comparison to observations are those provided by d\u2019Hendecourt & Allamandola (1986), Hudgins et al. (1993), and Kerkhof et al. (1999). The first reference gives IR spectra and band strengths of methanol ice at 10 K in the 2.5\u201320 \u03bcm range. The second reference presents a thorough infrared study providing a compendium of band strengths and optical constants in the 2.5\u2013200 \u03bcm interval for ices of astrophysical relevance. For pure methanol, optical constants are given for an ice grown at 10 K and subsequently warmed to 50 K, 75 K, 100 K, and 120 K. Band strengths at 10 K are also given. As far as we know, these are the only optical constants available in the literature for solid methanol. The third reference reports band strength variations of methanol absorptions in the mid-IR region with dilution in H2O ice, or in H2O and CO2 ice mixtures. Additional laboratory works provide infrared band strengths of this species at 10 K. Specifically, the data in Palumbo et al. (1999) correspond to the mid-IR (MIR), and the data in Sandford & Allamandola (1993) and Gerakines et al. (2005) to the near-IR (NIR). Finally, a recent paper by Bouilloud et al. (2015) presents a compilation and new measurements of MIR band strengths of methanol ice at low temperatures. Since methanol ice densities were unknown, the uncertainty on band-strength determination is very large. In particular, a discrepancy among the literature results up to 40% for the strength of the 9.74 \u03bcm band has been pointed out.","Citation Text":["d\u2019Hendecourt & Allamandola (1986)"],"Functions Text":["The laboratory data most widely used for comparison to observations are those provided by","The first reference gives IR spectra and band strengths of methanol ice at 10 K in the 2.5\u201320 \u03bcm range."],"Functions Label":["Uses","Uses"],"Citation Start End":[[640,673]],"Functions Start End":[[550,639],[725,828]]}
{"Identifier":"2019AandA...622A.106M__Lanz_et_al._(2010)_Instance_1","Paragraph":"The standard single-frequency detection methods for point sources in the CMB and far IR are based on wavelet techniques (Vielva et al. 2003; Barnard et al. 2004; Gonz\u00e1lez-Nuevo et al. 2006) or on the matched filter (or MF hereafter, Tegmark & de Oliveira-Costa 1998; Herranz et al. 2002; Barreiro et al. 2003; L\u00f3pez-Caniego et al. 2006, see also Herranz & Vielva 2010 for a review.). Wavelets are well suited for the detection of compact sources due to their good position-scale determination properties, whereas the MF is the optimal linear detector-estimator because it provides the maximum signal-to-noise (S\/N) amplification for a source with a known shape (usually the point-spread function, or PSF hereafter, of the telescope) embedded in statistically homogeneous and spatially correlated noise. By default, these techniques are applicable only to single-frequency sky images: even for multiwavelength observatories such as the Herschel Space Observatory (Pilbratt et al. 2010) or Planck (Tauber et al. 2010), the standard detection pipelines have produced individual source catalogs for each frequency band (see e.g., Planck Collaboration VII 2011; Planck Collaboration XXVIII 2014; Planck Collaboration XXVI 2016; Maddox et al. 2018). The next logical step is to boost the signal of faint sources by combining the different bands into a single detection, that is, \u201cmultifrequency detection\u201d. Most of the blind component separation algorithms that are used for diffuse components in microwave and far IR astronomy can not deal with the high diversity of spectral behaviors associated to the different populations of extragalactic compact sources (see for example Leach et al. 2008). However, over the last few years a number of multifrequency compact source detection techniques have been proposed in the literature (Herranz & Sanz 2008; Herranz et al. 2009; Lanz et al. 2010, 2013; Planck Collaboration Int. LIV 2018). A review on the topic can be found in Herranz et al. (2012). In particular, if the spatial profile and the spectral energy distribution (SED) of the sources are known, and if the cross-power spectrum is known, or can be estimated from the data, the optimal linear detection method is the matched multifilter (or MMF hereafter, Herranz et al. 2002). Lanz et al. (2010) also showed that the MMF can be generalized for the case where the SED of the sources is not known. This generalization outperforms the single-frequency MF in terms of S\/N and can be used to infer the spectral index of synchrotron-dominated radio sources, as shown in Lanz et al. (2013). However, in this paper we will incorporate a specific SED to the MMF in order to derive a photometric redshift estimation of dusty galaxies and high-redshift star forming galaxies detected in the IR part of the spectrum1. We will do so by applying the multifrequency MMF filter to the first and second data releases of the Herschel Astrophysical Terahertz Large Area Survey (the Herschel-ATLAS or H-ATLAS, Eales et al. 2010), the largest single key project carried out in open time with the Herschel Space Observatory. We restrict our multifrequency analysis to the three wavelength bands covered by the SPIRE instrument aboard Herschel (Griffin et al. 2010), centered around 250, 350 and 500 \u03bcm. As discussed in Hopwood et al. (2010), Lapi et al. (2011), Gonz\u00e1lez-Nuevo et al. (2012), Pearson et al. (2013) and Donevski et al. (2018), the SPIRE bands are ideal for capturing the peak in the SED corresponding to dust emission of star-forming galaxies at z\u2004\u223c\u20042, that is redshifted from its rest-frame wavelength around 70\u2013100 \u03bcm to the SPIRE wavelengths: This is the redshift range where galaxies have formed most of their stars. At higher redshifts, dusty star-forming galaxies (DSFGs) occupy the most massive halos and are among the most luminous objects found at z\u2004\u2273\u20044 (Micha\u0142owski et al. 2014; Oteo et al. 2016; Ikarashi et al. 2017). These high-redshift DSFGs have markedly red colors as seen by SPIRE, with rising flux densities from 250 to 500 \u03bcm (the so-called \u201c500 \u03bcm-risers\u201d), and have received a great deal of attention in the recent years (see for example Ivison et al. 2016; Negrello et al. 2017; Strandet et al. 2017). The DSFGs, and particularly the 500 \u03bcm risers uncovered by Herschel, are providing much insight into the early star forming history of the universe. However, sensitivity and limited angular resolution severely constrain the power of this type of objects as astrophysical probes. The sensitivity of SPIRE allows for the direct detection of only the brightest, and thus rarest objects, at the bright end of the luminosity function. By means of our multifrequency MMF technique, we intend to enhance the detectability and statistical significance of very faint red objects in the H-ATLAS source catalog and so expand the list of reliable 500 \u03bcm-riser candidates.","Citation Text":["Lanz et al. 2010"],"Functions Text":["However, over the last few years a number of multifrequency compact source detection techniques have been proposed in the literature"],"Functions Label":["Background"],"Citation Start End":[[1867,1883]],"Functions Start End":[[1691,1823]]}
{"Identifier":"2020ApJ...901...21B__Cane_2000_Instance_1","Paragraph":"In general, the most important effect of the ICME passage in generating major FDs is ascribable to the energetic interplanetary shock\/sheath region, although it has been proposed that the MC effect could even be dominant (e.g., Sanderson et al. 1990). Intense shocks are associated with the fast ICME propagation and may overcome the MC effect in modulating the GCR intensity, as turbulent magnetic fluctuations within the sheath can influence the GCR propagation significantly (e.g., Wibberenz et al. 1998). However, the minimum intensity occurs after the arrival of the MC (Badruddin 1986; Zhang & Burlaga 1988). For slower ICMEs, the effect of shock and turbulent sheath on the GCR intensity could be similar to that associated with the MC and sometimes it could be negligible. Therefore, if no interplanetary shock is found leading an ICME and a weak magnetic field turbulence is observed in the sheath region, the FD evolution mostly depends on the magnetic configuration of the ICME. A statistical study was performed by Richardson & Cane (2011) on more than 300 ICMEs, showing that 80% of them is associated with an FD, detected by the anticoincidence guard data on the International Monitoring Platform (IMP-8; Cane 2000), and that the minimum GCR intensities occur within the MC. They also concluded that ICMEs containing MCs cause deeper FDs on average with respect to ICMEs that do not have any MC structure. Despite numerous attempts to relate the properties of FDs with those of ICMEs at 1 au (Richardson et al. 1996; Belov et al. 2001, 2014; Belov 2008; Dumbovi\u0107 et al. 2012) there are significant gaps in our understanding of their underlying physical mechanisms. A renewed interest in studying GCR FDs has been fostered by the most recent S\/C observations in the near-Earth space at the lower rigidities with respect to NMs, for which FDs appear larger allowing for the study of the fine structure of the decrease formation. For instance, different data sets have been used to study GCRs and the associated FDs recorded by a radiation monitor on board LISA Pathfinder (LPF; Armano et al. 2018), the Electron Proton Helium INstrument detector on board the Solar and Heliospheric Observatory (SoHO) and the Chandra X-ray observatory (Heber et al. 2015; Dumbovi\u0107 et al. 2018), and the anticoincidence shield of the International Gamma-Ray Astrophysics Laboratory\u2019s spectrometer (Jordan et al. 2011). FDs have also been observed at several solar distances by the Radiation Assessment Dosimetry dose rates on the Mars Science Laboratory (Guo et al. 2018; von Forstner et al. 2018, 2020; Papaioannou et al. 2019), the Cassini\u2019s Magnetosphere Imaging Instrument and Low Energy Magnetospheric Measurement System measurements (Roussos et al. 2018) and beyond (Witasse et al. 2017; Winslow et al. 2018).","Citation Text":["Cane 2000"],"Functions Text":["A statistical study was performed by Richardson & Cane (2011) on more than 300 ICMEs, showing that 80% of them is associated with an FD, detected by the anticoincidence guard data on the International Monitoring Platform (IMP-8;","and that the minimum GCR intensities occur within the MC."],"Functions Label":["Background","Background"],"Citation Start End":[[1219,1228]],"Functions Start End":[[990,1218],[1231,1288]]}
{"Identifier":"2021ApJ...915...86A__Troja_et_al._2017_Instance_1","Paragraph":"In addition to confirming the origin of some short GRBs, combining data from observations of GW170817 and GRB 170817A allowed for the inference of basic properties of short GRB jets. These include the isotropic equivalent luminosity of the jet, determined through a redshift measurement made possible by the optical follow-up of the GW localization (Abbott et al. 2017a; Goldstein et al. 2017), and the geometry of the GRB jets (Williams et al. 2018; Farah et al. 2020; Mogushi et al. 2019). The precise mechanism by which the jet is launched is still unknown, although it is typically believed to be either neutrino-driven or magnetically driven (Nakar 2007, but see also Liu et al. 2015 and references therein). Indeed, the scientific debate about the emission profile of the jet and the subsequent gamma-ray production mechanism of GRB 170817A is still ongoing (Hallinan et al. 2017; Kasliwal et al. 2017; Lamb & Kobayashi 2017; Troja et al. 2017; Gottlieb et al. 2018b; Lazzati et al. 2018; Gill & Granot 2018; Mooley et al. 2018; Zhang et al. 2018; Ghirlanda et al. 2019). It is generally believed that there are symmetric polar outflows of highly relativistic material that travel parallel to the total angular momentum of the binary system (Aloy et al. 2005; Kumar & Zhang 2014; Murguia-Berthier et al. 2017). These jets are thought to be collimated and roughly axisymmetric, emitting preferentially in a narrow opening angle due to a combination of outflow geometry and relativistic beaming. The data from extensive multi-wavelength observation campaigns that ran for nearly 20 months following the merger (Fong et al. 2019; Makhathini et al. 2020; Troja et al. 2020) are in agreement with a structured jet model, in which the energy and bulk Lorentz factor gradually decrease with angular distance from the jet symmetry axis (e.g., Dai & Gou 2001; Lipunov et al. 2001; Rossi et al. 2002; Zhang & M\u00e9sz\u00e1ros 2002; Ghirlanda et al. 2019; Beniamini et al. 2020). Further, according to one of the models proposed, as the jet drills through the surrounding merger ejecta it inflates a mildly relativistic cocoon due to interactions between the material at the edge of the jet and the ejecta (Lazzati et al. 2017; Gottlieb et al. 2018a). In this case, it is possible that the cocoon alone could produce the gamma-rays observed from GRB 170817A (Gottlieb et al. 2018b). Additional joint detections of GRBs and GWs will significantly aid our understanding of the underlying energetics (Lamb & Kobayashi 2017; Wu & MacFadyen 2018; Burns et al. 2019), jet geometry (Farah et al. 2020; Mogushi et al. 2019; Biscoveanu et al. 2020; Hayes et al. 2020), and jet ignition mechanisms (Veres et al. 2018; Ciolfi et al. 2019; Zhang 2019) of binary neutron star (BNS) coalescences.","Citation Text":["Troja et al. 2017"],"Functions Text":["Indeed, the scientific debate about the emission profile of the jet and the subsequent gamma-ray production mechanism of GRB 170817A is still ongoing"],"Functions Label":["Motivation"],"Citation Start End":[[932,949]],"Functions Start End":[[714,863]]}
{"Identifier":"2021ApJ...912..163B__Davis_et_al._2018_Instance_1","Paragraph":"Braukmuller et al. (2018) proposed that all elements fall into one of four categories based on their condensation temperature: refractory elements (50% condensation temperature, Tc,50 > 1400 K), which exhibit approximately uniform enrichments in their Si-normalized concentrations in CC chondrites compared to CI chondrites by a factor of \u223c1\u20131.4; main component elements (1300 K  Tc,50  1400 K), which have approximately the same Si-normalized elemental abundances in CC chondrites as CI chondrites (differ by a factor of \u223c0.8\u20131.1); slope-volatile elements (800 K  Tc,50  1300 K), which exhibit monotonically decreasing Si-normalized concentrations with decreasing Tc,50 compared to CI chondrites; and plateau volatile elements (Tc,50  800 K), which display uniform depletions in Si-normalized concentrations compared to CI chondrites by a factor of \u223c0.1\u20130.7 that are characteristic of each CC chondrite group. Given their uniform nature with Tc,50 and comparatively well-constrained isotopic and elemental compositions, we chose to focus on the concentrations of refractory, main component, and plateau volatile elements in this study. For the refractory and main component elements in CC chondrites, we examine the elemental and isotopic compositions of Ti and Cr, respectively, because these are lithophile elements whose isotopic compositions have been measured precisely for a large number of chondrites and their components (Trinquier et al. 2007, 2009; Qin et al. 2010; Olsen et al. 2016; Van Kooten et al. 2016; Gerber et al. 2017; Davis et al. 2018; Zhu et al. 2019; Schneider et al. 2020; Williams et al. 2020). For CC iron meteorites, we examine the isotopic compositions of Mo and Ni, respectively, because these are siderophile elements (so are therefore present in appreciable concentrations in iron meteorites, unlike Ti and Cr) whose compositions have also been relatively well studied in a number of iron meteorites as well as chondrites and their components (Burkhardt et al. 2011; Budde et al. 2016; Kruijer et al. 2017; Bermingham et al. 2018; Nanne et al. 2019; Budde et al. 2019; Worsham et al. 2019; Brennecka et al. 2020; Spitzer et al. 2020). For the plateau volatile elements, we examine the elemental compositions of six elements (Bi, Ag, Pb, Zn, Te, and Sn) that exhibit a number of desirable properties: their concentrations have been relatively well constrained in CC chondrites; they show a range of lithophile, siderophile, and chalcophile behaviors; their concentrations do not appear to be strongly dependent on redox state; they show minimal variability among NC chondrite groups. Our reasoning for not considering the isotopic compositions of these elements is discussed in Section 2.3. The adopted isotopic and chemical composition of each element used in this study in CC chondrites, CC iron meteorites, CAIs, CI chondrites, and NC chondrites are included in Table 1. Uncertainties on elemental concentrations have not been routinely reported throughout the literature, although these values are typically \u00b15 wt% (e.g., Lodders 2003; Palme et al. 2014). CAIs can be categorized into six groups based on their compositions (Stracke et al. 2012). For the purposes of this study, we adopt the composition of type I CAIs as the representative value of refractory objects because they are seemingly the most abundant type and lack the characteristic elemental depletions of other CAI groups (e.g., Stracke et al. 2012; Brennecka et al. 2020). We also focus largely on ordinary chondrites (OC) as representative NC meteorites rather than enstatite chondrites (EC) or Rumuruti chondrites (RC). This is because EC chondrites formed under more reducing conditions than OC and RC chondrites, which introduced a compositional signature for some elements to EC chondrites that is not present in OC, RC, or CC chondrites (presumably due to their formation in more oxidizing environments) so is not representative of large-scale mixing in the disk (Alexander 2019b). Additionally, the isotopic compositions of RC chondrites are sparsely measured compared to OC and EC chondrites. NC meteorites could have experienced a number of processes (e.g., mixing, chondrule formation, volatile loss, the addition of refractory materials, etc.) that gave these meteorites their specific chemical and isotopic signatures (Alexander 2019b). We do not explore these processes in this study and simply adopt the measured elemental and isotopic compositions of NC chondrites as potential end-members for the compositions of CC meteorites.","Citation Text":["Davis et al. 2018"],"Functions Text":["For the refractory and main component elements in CC chondrites, we examine the elemental and isotopic compositions of Ti and Cr, respectively, because these are lithophile elements whose isotopic compositions have been measured precisely for a large number of chondrites and their components"],"Functions Label":["Motivation"],"Citation Start End":[[1540,1557]],"Functions Start End":[[1137,1429]]}
{"Identifier":"2022MNRAS.509.3427D__D\u00edaz_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":["D\u00edaz et al. 2017"],"Functions Text":["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","and later confirmed by others teams"],"Functions Label":["Background","Background"],"Citation Start End":[[895,911]],"Functions Start End":[[622,774],[797,832]]}
{"Identifier":"2016AandA...595A..16T__Eggleton_et_al._(2006)_Instance_1","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)"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[886,908]],"Functions Start End":[[731,885]]}
{"Identifier":"2018AandA...614A..72V__Bocquet_et_al._2015_Instance_1","Paragraph":"Cluster cosmology requires jointly modelling the physical parameters describing the evolution of the intra-cluster medium (ICM) alongwith the impact of selection procedure. While the first self-consistent methods have followed a backward modelling of the recovered cosmology-dependent, mass function (e.g. Vikhlinin et al. 2009), more recent studies moved to a forward approach whose likelihood includes physical quantities such as luminosity, temperature, or gas fraction (e.g. Mantz et al. 2014, 2015). Cluster number counts from Sunyaev-Z\u2019eldovich surveys are routinely modelled in terms of the signal-to-noise ratio (S\/N) or the Compton parameter of the detections, which can be related to the cluster mass via scaling relations from X-ray, lensing, or velocities (e.g. Vanderlinde et al. 2010; Hasselfield et al. 2013; Benson et al. 2013; Bocquet et al. 2015; Planck Collaboration XXIV 2016). In this context, we are developing a cosmological analysis method (ASpiX) based on X-ray cluster number counts that does not explicitly rely either on cluster mass determinations or physical quantities. This method consists in the modelling of the multidimensional distribution of a set of directly measurable X-ray clusters quantities, namely: count rates (CRs), hardness ratios (HRs), and apparent size (rc), which are all cosmology independent. This method is particularly suited to rather shallow survey-type data, when the number of collected X-ray photons is too low to enable detailed spectral and morphological analyses. Thanks to its modularity, the ASpiX method considerably eases the process by simultaneously fitting in the observed parameter space, the effect of cosmology, selection, and cluster physics. Depending on the volume surveyed, that is, the number of clusters involved in the analysis, the number of parameters that may be fitted can increase from a few to 15 or more, including in particular scatter and evolution in the scaling relations. This method cannot rival approaches including deep pointed X-ray observations along with ancillary data from other wavebands and, fundamentally, faces the same uncertainties as to the observable-mass transformation. However, the method allows the inclusion of the vast majority of the detected clusters even when only a few tens of photons are available. Furthermore, when cosmological simulations are produced at a significantly high rate, the method will allow us to totally bypass any mass estimate or scaling-relation related formalism; instead, it will solely rely on the simulations by comparing the observed and simulated parameter distributions (Pierre et al. 2017). In the end, neitherassumptions based on the hydrostatic equilibrium nor any modelling of the mass function will be necessary.","Citation Text":["Bocquet et al. 2015"],"Functions Text":["Cluster number counts from Sunyaev-Z\u2019eldovich surveys are routinely modelled in terms of the signal-to-noise ratio (S\/N) or the Compton parameter of the detections, which can be related to the cluster mass via scaling relations from X-ray, lensing, or velocities (e.g."],"Functions Label":["Background"],"Citation Start End":[[844,863]],"Functions Start End":[[505,773]]}
{"Identifier":"2022AandA...661A..71H__Liu_et_al._2020_Instance_1","Paragraph":"While using the Pad\u00e9(2,1) approximation, we checked the influence of the quasar relation on the cosmological constraints. We first provide the constraints considering a free quasar relation. Then combining that with the following two special cases, the effect of the slope parameter \u03b3 is studied. One case is the fitting parameters q0 and j0 in the case of fixed parameters \u03b4, \u03b3, and \u03b2. The other is to refit q0 and j0 by changing the setting of the fixed parameter \u03b3 within a 1\u03c3 error (0.005). The fixed parameters \u03b4, \u03b3, and \u03b2 are given in terms of the best fitting results from 1598 quasars in a flat \u039bCDM model. By substituting Eqs. (25) into (26), and replacing dL with Eq. (6), we obtained the best fits, that is \n\n\n\n\n\u03a9\nm\n\n=\n0\n.\n\n65\n\n\u2212\n0.19\n\n\n+\n0.16\n\n\n\n\n$ \\Omega_{m} = 0.65^{+0.16}_{-0.19} $\n\n\n, \u03b4\u2004=\u20040.23\u2006\u2005\u00b1\u2005\u20060.01, \u03b3\u2004=\u20040.62\u2006\u2005\u00b1\u2005\u20060.01, and \u03b2\u2004=\u20047.60\u2006\u2005\u00b1\u2005\u20060.28. This result is consistent with previous research within a 1\u03c3 error (Lusso & Risaliti 2016; Melia 2019; Salvestrini et al. 2019; Khadka & Ratra 2020a, 2022; Liu et al. 2020; Wei & Melia 2020). Here, it is worth noting that the 1\u03c3 errors of the \u03b3 best fit from previous research are both larger than 0.01. Figure 4 shows the fitting results with a free quasar relation. The best fits are \u03b4\u2004=\u20040.23\u2006\u2005\u00b1\u2005\u20060.01, \u03b3\u2004=\u20040.64\u2006\u2005\u00b1\u2005\u20060.01, \n\n\n\n\u03b2\n=\n7\n.\n\n17\n\n\u2212\n0.27\n\n\n+\n0.28\n\n\n\n\n$ \\beta = 7.17^{+0.28}_{-0.27} $\n\n\n, q0\u2004=\u2004\u22120.67\u2006\u2005\u00b1\u2005\u20060.04, and \n\n\n\n\nj\n0\n\n=\n2\n.\n\n43\n\n\u2212\n0.46\n\n\n+\n0.51\n\n\n\n\n$ j_{0} = 2.43^{+0.51}_{-0.46} $\n\n\n. This result is in line with the previous results obtained by the logarithmic polynomials. The corresponding q0\u2005\u2212\u2005j0 projection is also plotted in Fig. 5 as a gray contour. In addition, we also show the confidence contours of the other two special cases. (1) When we chose \u03b4\u2004=\u20040.23, \u03b3\u2004=\u20040.62, and \u03b2\u2004=\u20047.60 (best fits from 1598 quasars in a flat \u039bCDM model), the results are q0\u2004=\u2004\u22120.82\u2006\u2005\u00b1\u2005\u20060.05 and \n\n\n\n\nj\n0\n\n=\n4\n.\n\n99\n\n\u2212\n0.63\n\n\n+\n0.69\n\n\n\n\n$ j_{0} = 4.99^{+0.69}_{-0.63} $\n\n\n, which is represented by the red contours in Fig. 5. There exists more than a 4\u03c3 tension with the flat \u039bCDM. (2) In only changing \u03b3\u2004=\u20040.625, the new result is q0\u2004=\u2004\u22120.58\u2006\u2005\u00b1\u2005\u20060.03 and \n\n\n\n\nj\n0\n\n=\n1\n.\n\n19\n\n\u2212\n0.31\n\n\n+\n0.34\n\n\n\n\n$ j_{0} = 1.19^{+0.34}_{-0.31} $\n\n\n which is described by the blue contours of Fig. 5. We find that the \u03b3 value changes from 0.62 to 0.625 (within a 1\u03c3 error of 0.01), and the 4\u03c3 tension disappears. This result is consistent with the \u039bCDM model within a 1\u03c3 level. We find that the quasar relation can obviously affect the cosmographic constraint, especially the slope parameter \u03b3.","Citation Text":["Liu et al. 2020"],"Functions Text":["This result is consistent with previous research within a 1\u03c3 error","Here, it is worth noting that the 1\u03c3 errors of the \u03b3 best fit from previous research are both larger than 0.01."],"Functions Label":["Similarities","Differences"],"Citation Start End":[[1018,1033]],"Functions Start End":[[862,928],[1054,1165]]}
{"Identifier":"2021MNRAS.507.2766S__Sumiyoshi_et_al._2005_Instance_2","Paragraph":"In order to make a linear analysis, first we have to prepare the PNS models as a background. The PNS properties depend on not only the density and pressure profiles but also the distributions of temperature (or entropy per baryon) and electron fraction inside the PNS, while such profiles can be determined only via the numerical simulation of the core-collapse supernova explosion. In this study, as in Sotani & Sumiyoshi (2019), we particularly adopt the profiles obtained via the numerical simulations performed by solving the general relativistic neutrino-radiation hydrodynamics under the spherical symmetry. In the simulations, hydrodynamics and neutrino transfer in general relativity are solved simultaneously (Yamada 1997; Yamada, Janka & Suzuki 1999; Sumiyoshi et al. 2005). To describe the neutrino transfer, the Boltzmann equation is directly solved with the multi-angle and multi-energy neutrino distributions for four species, \u03bde, $\\bar{\\nu }_\\mathrm{ e}$, \u03bd\u03bc\/\u03c4, and $\\bar{\\nu }_{\\mu \/\\tau }$, i.e. we implement six species of neutrinos by assuming \u03bc-type and \u03c4-type (anti-)neutrinos have identical distributions. For the collision term associated with neutrino emission, absorption, and scattering with leptons, nucleons, and nuclei handles, the basic neutrino reactions are adopted (Bruenn 1985; Sumiyoshi et al. 2005). The metric adopted in the numerical code is given by\n(1)$$\\begin{eqnarray*}\r\n\\mathrm{ d}s^2 = -\\mathrm{ e}^{2\\Phi (t,m_\\mathrm{ b})}\\mathrm{ d}t^2 + \\mathrm{ e}^{2\\Lambda (t,m_\\mathrm{ b})}\\mathrm{ d}m_\\mathrm{ b}^2 + r^2(t,m_\\mathrm{ b})(\\mathrm{ d}\\theta ^2 + \\sin ^2\\theta \\mathrm{ d}\\mathrm{ }\\phi ^2), \\nonumber\\\\\r\n\\end{eqnarray*}$$where t and mb denote the coordinate time and the baryon mass coordinate, respectively (Misner & Sharp 1964). In addition, mb is related to the circumference radius (r) via the baryon mass conservation, while the metric functions, \u03a6(t, mb) and \u039b(t, mb), are evolved together with hydrodynamical variables in the numerical simulations (Yamada 1997). The numerical simulations for core-collapse supernovae have been done with 255 grid points in the radial mass coordinate, 6 grid points in the neutrino angle, and 14 grid points in the neutrino energy. The rezoning of radial mesh is made during the simulations to resolve the accreting matter. We remark that the radial grids of mass coordinate are non-uniformly arranged to cover not only the dense region inside the central object but also the region for accreting matter.","Citation Text":["Sumiyoshi et al. 2005"],"Functions Text":["For the collision term associated with neutrino emission, absorption, and scattering with leptons, nucleons, and nuclei handles, the basic neutrino reactions are adopted"],"Functions Label":["Uses"],"Citation Start End":[[1312,1333]],"Functions Start End":[[1128,1297]]}
{"Identifier":"2020AandA...641A.126B__Jones_&_Hardee_1979_Instance_1","Paragraph":"Many low-luminosity active galactic nuclei (LLAGN) display prominent jets and compact cores that are sources of highly nonthermal continuum radio emission (see, e.g., Heeschen 1970; Wrobel & Heeschen 1991). The observational signatures of the compact cores have been reproduced using models that produce self-absorbed synchrotron emission in the jet (Falcke & Biermann 1995; Falcke et al. 2004) or in a magnetized accretion flow (Narayan et al. 1998; Yuan et al. 2003; Broderick & Loeb 2006; Moscibrodzka et al. 2009; Dexter et al. 2009; see also Falcke et al. 2001). This radiation is emitted by relativistic electrons gyrating around magnetic field lines. In the optically thin limit, the emission is significantly polarized (Jones & Hardee 1979), an effect that has been observed in higher-luminosity AGN sources (Gabuzda et al. 1996; Gabuzda & Cawthorne 2000; Lyutikov et al. 2005). The polarized emission from an accreting AGN can therefore yield information about the magnetic-field morphology of the source, which may be crucial to the evolution of the accretion flow of the AGN. The Event Horizon Telescope (EHT) is a worldwide millimeter-wavelength array capable of resolving the black-hole shadow (Goddi et al. 2017; Event Horizon Telescope Collaboration 2019); this is a characteristic feature of the radio-frequency emission from optically thin AGN at the scale of the event horizon (Falcke et al. 2000; Broderick & Narayan 2006), although the black-hole shadow may be obscured or exaggerated in certain accretion scenarios (see Gralla et al. 2019 and Narayan et al. 2019). The EHT can also determine the polarization state of such emission: Johnson et al. (2015) report 1.3 mm observations (230 GHz) that indicate partially ordered magnetic fields within a region of about six Schwarzschild radii around the event horizon of Sagittarius A* (Sgr A*), the supermassive black hole in the center of the Milky Way. Bower et al. (2003) reported stable long-term behavior and short-term variability in Sgr A* rotation measure, implying a complex inner region (within 10 Schwarzschild radii) in which both emission and propagation effects are important to the observed polarization. Hada et al. (2016) studied the central black hole in the galaxy M 87, and observed a bright feature with (linear) polarization degree of 0.2 at 86 GHz at the jet base. Observations in infrared by Gravity Collaboration (2018) were consistent with a model in which a relativistic \u201chot spot\u201d of material, orbiting near the innermost stable circular orbit (ISCO) of Sgr A* in a poloidal magnetic field, emits polarized synchrotron radiation.","Citation Text":["Jones & Hardee 1979"],"Functions Text":["In the optically thin limit, the emission is significantly polarized","an effect that has been observed in higher-luminosity AGN sources"],"Functions Label":["Background","Similarities"],"Citation Start End":[[728,747]],"Functions Start End":[[658,726],[750,815]]}
{"Identifier":"2022MNRAS.513.4464T__Hopkins_et_al._2020_Instance_1","Paragraph":"\nGalactic winds: Galactic winds driven by CRs have often been simulated in two limits: a diffusion-dominated regime, due possibly to \u2018extrinsic confinement\u2019, where CRs are scattered by extrinsic turbulence, and\/or due to various wave damping mechanisms (e.g. ion neutral damping) and streaming-dominated \u2018self confinement\u2019, where CRs are confined by Alfven waves they produce via the gyroresonant streaming instability. In the diffusive \u2018extrinsic confinement\u2019 case, CRs do not heat the gas.19 In the streaming dominated \u2018self confinement\u2019 case, CR transport heats gas at a rate vA \u00b7 \u2207Pc. The diffusive case fits \u03b3 ray observations better, because CRs can propagate out of the galaxy faster (Chan et al. 2019). It is also much better at driving winds, because the CRs do not suffer strong energy losses via Alfven wave heating (Wiener et al. 2017b; Hopkins et al. 2020). However, we expect self-confinement to be very strong at the \u223cGeV energies where CR energy peaks (Kulsrud & Pearce 1969; Farmer & Goldreich 2004; Wiener et al. 2013), while extrinsic compressible turbulence is strongly damped at small scales, and unlikely to efficiently scatter \u223cGeV CRs (Yan & Lazarian 2002). Thus, CR winds should be streaming dominated and relatively inefficient. The CR staircase changes these dichotomies by changing the structure of the wind. We have seen how CR pressure can build up in streaming dominated simulations, due to trapping at bottlenecks. This increases mass outflow rates, similar to the effect of increased opacity in radiative outflows. In CR streaming simulations of isothermal winds where the CR acoustic instability arose, Quataert et al. (2022a) found an increase in wind mass loss rates by an order of magnitude, compared to analytic models without a CR staircase, illustrating the potential impact of CR staircases. High-resolution cosmological zoom simulations of CR staircases are actually well within reach. As seen in Appendix Section B, all that is required is that the diffusion length $l_{\\rm diff} \\sim \\kappa \/c_{\\rm s} \\sim 2 \\, {\\rm kpc} \\, \\left(\\frac{\\kappa }{10^{29} {\\rm cm^2 s^{-1}}} \\right)\\left(\\frac{c_{\\rm s}}{150 \\, {\\rm km \\, s^{-1}}} \\right)^{-1}$ is resolved. However, to date only the FIRE collaboration has implemented the two moment method (capable of dealing with CR streaming) in such simulations, and \u2013 in contrast to, for instance, van de Voort et al. (2021) \u2013 the plasma \u03b2 in their winds is too high for the acoustic instability to develop (Hopkins et al. 2020). But alternate setups where CR staircases appear are certainly numerically feasible.","Citation Text":["Hopkins et al. 2020"],"Functions Text":["It is also much better at driving winds, because the CRs do not suffer strong energy losses via Alfven wave heating"],"Functions Label":["Uses"],"Citation Start End":[[849,868]],"Functions Start End":[[711,826]]}
{"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)"],"Functions Text":["We repeated the plasma pressure calculation presented by","and Fuselier et al. (2012) for the new ENA energy spectrum."],"Functions Label":["Uses","Uses"],"Citation Start End":[[57,80]],"Functions Start End":[[0,56],[81,140]]}
{"Identifier":"2018ApJ...859..115W__Luo_et_al._2016_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":["Luo et al. 2016"],"Functions Text":["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.,"],"Functions Label":["Similarities"],"Citation Start End":[[685,700]],"Functions Start End":[[481,643]]}
{"Identifier":"2019MNRAS.489.4429R__Lada_et_al._2008_Instance_1","Paragraph":"In the currently accepted model of star formation, most stars form in families of star clusters within giant molecular clouds (GMCs). However, clusters and their families can be quite diverse, and there is no consensus yet on a general model for cluster formation. An additional complication is that once stars are born, they destroy most of the evidence relating to the initial stages of their formation. In this sense, determining the early stages of cluster formation in GMCs may be key to understanding many distinct and complex processes involved in such a phenomenon. The Pipe Nebula, a nearby (d = 130 pc) GMC with a total mass of 7.9 \u00d7 103\u2009M\u2299 (Lada, Lombardi & Alves 2010) and very little star formation activity (Brooke et al. 2007; Forbrich et al. 2009, 2010) has played a workhorse role in a number of previous studies (e.g. Alves, Lombardi & Lada 2007; Muench et al. 2007; Alves, Franco & Girart 2008; Lada et al. 2008; Rathborne et al. 2009; Rom\u00e1n-Z\u00fa\u00f1iga, Lada & Alves 2009; Rom\u00e1n-Z\u00fa\u00f1iga et al. 2010; Peretto et al. 2012; Frau et al. 2015; Hasenberger et al. 2018). Many of these works have the common goal of determining how GMCs are organized in the stages that precede their collapse and the consequent conversion of gas into stars. Interestingly, despite the fact that it is almost starless, the column density peaks1 in the Pipe appear to show a clear imprint of clustering, as well as segregation by density and mass (e.g. Rom\u00e1n-Z\u00fa\u00f1iga et al. 2010; Alfaro & Rom\u00e1n-Z\u00fa\u00f1iga 2018, hereafter RAL10 and AR18, respectively). The Pipe was also the first cloud where a \u2018core mass function\u2019 (see Alves et al. 2007; Rathborne et al. 2009) was constructed, leading to many subsequent studies that have discussed what appears to be a clear homology to its stellar counterpart (e.g. Alves et al. 2007; Goodwin et al. 2008; Swift & Williams 2008; Kainulainen et al. 2009). These findings are suggestive of a primordial organization of clouds towards the formation of groups and clusters, evident in the properties of density peaks.","Citation Text":["Lada et al. 2008"],"Functions Text":["The Pipe Nebula,","has played a workhorse role in a number of previous studies (e.g.","Many of these works have the common goal of determining how GMCs are organized in the stages that precede their collapse and the consequent conversion of gas into stars."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[914,930]],"Functions Start End":[[574,590],[770,835],[1079,1248]]}
{"Identifier":"2021ApJ...916...70Z__Millon_et_al._2020_Instance_1","Paragraph":"For the uncertainty from the Fermat potential, which is determined by lens modeling, source host galaxies without dazzling AGNs are conducive to the reconstruction of the mass distribution of the lens galaxies in lensed FRB systems. Simulations show that the Fermat potential component contributes approximately \u223c0.8% uncertainty on D\u0394t (Li et al. 2018b). Actually, in lens modeling, except for the contamination from the light of the dazzling AGN in the source, mass-sheet degeneracy (MSD; a family of mass density profiles that could reproduce the same lensing observables, e.g., image positions and relative fluxes, but yields different measured values of H0) also plays an important role in leading to the loss of precision and accuracy. In recent years, this issue has been intensively investigated. Generally, there are two ways, which include applying theoretical priors (make radial-mass density profile assumptions) and appealing nonlensing data (e.g., stellar kinematics) to break the MSD. For instance, the TDCOSMO collaboration (Millon et al. 2020) has achieved \u223c2% in the inference of H0 from time-delay cosmography under the assumption that the radial-mass distribution of the lens can be described by a power-law mass profile or a composite of a dark-matter halo (Navarro et al. 1997) and baryon matter. Meanwhile, Gomer & Williams (2020) tested the effects of the power-law assumption and found that the power-law assumption would introduce significant bias in the recovery of H0. In practice, as suggested in Li et al. (2018b), a high-quality optical\/IR image of the source\u2013lens system could be very helpful to avoid choosing the wrong models. That is, the model we choose to characterize the mass distribution of the lens according to the high-resolution image might be more complex than the simple power-law one, and thus could significantly reduce the bias. However, it is difficult to choose an exactly right model for the lens, thus the uncertainty of the Fermat potential might be larger than 0.8% in real observations. For example, using eight time-delay galaxy lenses and more flexible modeling methods, a precision of 4.97% on H0 was achieved (Denzel et al. 2021). With spatially resolved kinematics and external observations to break MSD, a 5% precision of H0 was inferred by assuming that the deflector of TDCOSMO and the Sloan Lens ACS (SLACS) lenses are drawn from the same population (Birrer et al. 2020). It is promising that by increasing the size of samples and applying the hierarchical framework introduced by Birrer et al. (2020), a precision of 1.5% or 1.2% on H0 will be achieved without assumptions on the radial-mass profile of lens galaxies (Birrer & Treu 2021). More recently, Ding et al. (2021) implemented realistic simulations in lens modeling based on HST WFC3 observations from transient sources (e.g., supernovae, gamma-ray bursts, FRBs, and GWs) to compare the precision of H0 inferred from the transient case and the lensed AGN case. They found that, compared with traditional lensed quasars, the precision for inferring the Shapiro delay and the geometric delay could be improved by a factor of 3.8 and 4.7 for lensed transient systems, respectively. It means that the precision of the Fermat potential reconstruction could be improved by a factor of about 6. Furthermore, the lensed transient system also facilitates the determination of higher signal-to-noise stellar kinematics of the main deflector, and thus its mass density profile, which in turn plays a key role in breaking the mass-sheet degeneracy. With the abovementioned investigations taken into consideration, we first take 0.8% relative uncertainty on the Fermat potential on the basis of simulations presented in Ding et al. (2021) and Li et al. (2018b). In addition, we also consider a more conservative 1.5% relative uncertainty on Fermat potential for comparison.","Citation Text":["Millon et al. 2020"],"Functions Text":["For instance, the TDCOSMO collaboration","has achieved \u223c2% in the inference of H0 from time-delay cosmography under the assumption that the radial-mass distribution of the lens can be described by a power-law mass profile or a composite of a dark-matter halo","and baryon matter."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[1041,1059]],"Functions Start End":[[1000,1039],[1061,1277],[1300,1318]]}
{"Identifier":"2017AandA...602A..26L__Robrade_&_Schmitt_2005_Instance_1","Paragraph":"Along with the flare related changes, we also examine possible orbital variations by dividing the overall quiescent spectra in periods dubbed as \u201cQuiet 1\u201d and \u201cQuiet 2\u201d (the time bins for the spectra are as shown in Fig. 1) and model the pn, MOS and RGS X-ray spectra. We specifically determine the temperatures, emission measures and abundances relative to solar values (Grevesse & Sauval 1998) with simultaneous iterative global XSPEC fits to the combination of EPIC and RGS (RGS+PN or RGS+MOS) spectra with variable-APEC (VAPEC; Smith et al. 2001) plasma models. As is often observed (G\u00fcdel et al. 2001; Robrade & Schmitt 2005; Lalitha et al. 2013), we require multi-temperature components to achieve an adequate description of the observed coronal spectra. We use combinations of two, three and four temperature components and find that a three temperature component model leads to an adequate description of the data. We fit each of these spectra in the full 0.2\u201310 keV energy range. For fitting the RGS spectra, the temperature and the abundances of elements like carbon, nitrogen, oxygen, neon and iron are allowed to vary freely and independently, however, the abundances are fixed among the different APEC temperature components. For fitting EPIC-MOS or pn spectra we allow the magnesium, sulphur, and silicon abundances to vary along with the oxygen, neon and iron abundance. However, the carbon and nitrogen abundances are fixed to values obtained from the RGS, which is more sensitive to strong individual lines of these elements. In Table 1, we summarise the results of this fitting procedure along with the 90% confidence range errors. Table 1 shows that the quiescent state is characterised by dominant plasma components at ~3, ~7.5 and ~20 MK, while during the flare the coronal temperature bins increase to ~3.5, ~12 and ~32 MK. During the flare, a pronounced enhancement of the emission measure at 2.8 keV is present, indicating the rise of emission measure at a higher temperature. The quiescent bins (quiet 1 and quiet 2) does not show any significant difference in the coronal properties when compared to the overall quiescent time-bin; suggesting no changes in coronal properties with orbital variation. ","Citation Text":["Robrade & Schmitt 2005"],"Functions Text":["As is often observed",", we require multi-temperature components to achieve an adequate description of the observed coronal spectra."],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[607,629]],"Functions Start End":[[566,586],[651,760]]}
{"Identifier":"2021MNRAS.504.3316B__than_2000_Instance_5","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)"],"Functions Text":["May & Stevenson (2020) claimed that the differences between the retrieved phase offsets is the result of temporal binning which was not used by","and May & Stevenson (2020) but was used by Mendon\u00e7a et al. (2018), Morello et al. (2019), and this work."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[963,986]],"Functions Start End":[[819,962],[987,1091]]}
{"Identifier":"2018MNRAS.478.4986K__Greig_&_Mesinger_2017_Instance_1","Paragraph":"Various observational techniques, such as measuring the Gunn\u2013Peterson optical depth from QSO spectra or the prevalence of Ly \u03b1 emission in high-redshift galaxies, have placed very tight constraints on the volume filling fraction of neutral hydrogen in the intergalactic medium (IGM) towards the end of reionization at z\u223c 6 (Becker et al. 2001; Fan et al. 2006; Totani et al. 2006; McQuinn et al. 2007; McQuinn et al. 2008; Ota et al. 2008; Ouchi et al. 2010; Bolton et al. 2011; Mortlock et al. 2011; Ono et al. 2012; Becker & Bolton 2013; Chornock et al. 2013; Robertson et al. 2013; Schroeder, Mesinger & Haiman 2013; Caruana et al. 2014; Pentericci et al. 2014; Schenker et al. 2014; Tilvi et al. 2014; McGreer, Mesinger & D\u2019Odorico 2015; Mesinger et al. 2015; Mitra, Choudhury & Ferrara 2015, 2018; Sobacchi & Mesinger 2015; Greig & Mesinger 2017). While the timing of the end of reionization is rather well constrained by observations and modelling (e.g. Fan et al. 2006; Choudhury et al. 2015) as is the photoionization rate of neutral hydrogen in the post-reionization Universe (Bolton & Haehnelt 2007; Calverley et al. 2011; Wyithe & Bolton 2011; Becker & Bolton 2013), much remains uncertain about the onset and extent of the process. This uncertainty is primarily driven by the current lack of understanding of which sources reionized the Universe and the difficulty of observing these systems deep into the epoch of reionization. Various classes of objects have been proposed as the sources of reionization including dwarf galaxies, mini-haloes, massive galaxies, active galactic nuclei, accretion shocks, globular clusters, stellar mass black holes, and dark matter annihilation and decay (Couchman & Rees 1986; Shapiro & Giroux 1987; Haiman & Loeb 1998; Madau, Haardt & Rees 1999; Ricotti 2002; Madau et al. 2004; Ricotti & Ostriker 2004; Mapelli, Ferrara & Pierpaoli 2006; Dopita et al. 2011; Mirabel et al. 2011; Katz & Ricotti 2013, 2014; Madau & Haardt 2015).","Citation Text":["Greig & Mesinger 2017"],"Functions Text":["Various observational techniques, such as measuring the Gunn\u2013Peterson optical depth from QSO spectra or the prevalence of Ly \u03b1 emission in high-redshift galaxies, have placed very tight constraints on the volume filling fraction of neutral hydrogen in the intergalactic medium (IGM) towards the end of reionization at z\u223c 6"],"Functions Label":["Background"],"Citation Start End":[[829,850]],"Functions Start End":[[0,322]]}
{"Identifier":"2021AandA...649A.126T__Luck_(2018b)_Instance_1","Paragraph":"Studies of the radial n-capture-to-iron abundance gradients are very scarce so far. We can only search for a broad agreement of our results with several studies of abundance gradients with galactocentric distances (Rgc). da Silva et al. (2016) studied n-capture elements across the Galactic thin disc based on Cepheid variables. Because the Cepheids are young stars, their Rgc may be rather close to their birthplaces and Rmean. da Silva et al. (2016) supplemented their sample of 111 Cepheids with 324 more stars from other studies and found that the [Y\/Fe] distribution is flat throughout the entire disc. In our study, we confirm this finding not only based on the whole thin-disc sample of stars and on a subsample of younger \u22644 Gyr stars, but also add another light s-process dominated element strontium. Like in our study, da Silva et al. (2016) also obtained positive [El\/Fe] radial gradients for La, Ce, Nd, and Eu. The slopes are rather similar. For [Eu\/Fe], they differ just by 0.002 dex kpc\u22121. More recently, Luck (2018b) also investigated the gradients of n-capture element abundance-to-iron ratios with respect to Rgc for a sample of 435 Cepheids. It is interesting to note that the [Ba\/Fe] versus Rgc slope according to this Cepheid sample is also negative, as in our study. [Ba\/Fe] is the only n-capture element-to-iron ratio with a negative radial gradient in our sample of stars and in Luck (2018b). Overbeek et al. (2016) investigated trends of Pr, Nd, and Eu to Fe abundance ratios with respect to Rgc using 23 open clusters. As in our study, they found that these elements have positive linear trends with galactocentric radius (the linear regression slopes are of about +0.04 dex kpc\u22121). They also suggested that the [El\/Fe] relation of Pr and Nd, but not Eu, with the galactocentric radius may not be linear because the [El\/Fe] of these elements appears to be enhanced around 10 kpc and drop around 12 kpc. Because only a small number of stars lie at these large radial distances, we cannot address this question. For the thick-disc stars, the radial abundance-to-iron slopes are negligible, as was found for \u03b1-process elements by Li et al. (2018), even though the production sites of \u03b1-elements and s-processes dominated elements are quite different.","Citation Text":["Luck (2018b)"],"Functions Text":["More recently,","also investigated the gradients of n-capture element abundance-to-iron ratios with respect to Rgc for a sample of 435 Cepheids."],"Functions Label":["Background","Background"],"Citation Start End":[[1020,1032]],"Functions Start End":[[1005,1019],[1033,1160]]}
{"Identifier":"2015MNRAS.451.2544P__Leitet_et_al._2011_Instance_1","Paragraph":"Observations are now probing galaxies in the middle of the reionization epoch, when the gas in the intergalactic medium was transformed from its initially neutral state into a hot, ionized plasma (e.g. McLure et al. 2011; Finkelstein et al. 2012; Bouwens et al. 2014). Most likely stars in galaxies are responsible for this transformation, although this heavily depends on the fraction of ionizing photons produced by the stars that make it into the intergalactic medium, the so-called escape fraction fesc. The escape fraction is a key parameter in studies of the contribution of the observed galaxy population to reionization (e.g. Bouwens et al. 2012; Robertson et al. 2013), semi-analytic modelling of reionization (e.g. Choudhury, Haehnelt & Regan 2009; Pritchard, Loeb & Wyithe 2010; Santos et al. 2010; Mesinger, Furlanetto & Cen 2011; Raskutti et al. 2012; Mitra, Ferrara & Choudhury 2013; Shull et al. 2012) and numerical simulations of reionization (e.g. Iliev et al. 2006; Trac & Cen 2007; Ciardi et al. 2012). A large effort is going into determining the escape fraction observationally. Except for two objects (Leitet et al. 2011, 2013), in the local Universe no ionizing radiation has been detected directly (Leitherer et al. 1995; Deharveng et al. 2001), although some objects show indirect evidence of photon leakage (Heckman et al. 2011; Zastrow et al. 2011). The lack of detections may be partly due to selection bias (Bergvall et al. 2013), but the objects from which radiation is detected have very low escape fractions, fesc  4\u2009per\u2009cent. At z \u223c 1, no objects with leaking ionizing photons have been detected (Bridge et al. 2010; Siana et al. 2010), but at z \u223c 3, the highest redshift at which the opacity of the intergalactic medium for ionizing photons is approximately less than unity, ionizing photons have been detected in \u223c10\u2009per\u2009cent of the observed objects (Nestor et al. 2013). Attempts to constrain the escape fraction with numerical simulations find ranges between fesc  10\u2009per\u2009cent (Razoumov & Sommer-Larsen 2006, 2007; Gnedin, Kravtsov & Chen 2008; Paardekooper et al. 2011; Kim et al. 2013) and fesc > 80\u2009per\u2009cent (Wise & Cen 2009; Razoumov & Sommer-Larsen 2010), with likely a strong mass and redshift dependence (Yajima, Choi & Nagamine 2010; Wise et al. 2014). Due to the opacity of the intergalactic medium, we need to mostly rely on numerical simulations to learn about the escape fraction during the epoch of reionization.","Citation Text":["Leitet et al. 2011"],"Functions Text":["Except for two objects","in the local Universe no ionizing radiation has been detected directly","although some objects show indirect evidence of photon leakage"],"Functions Label":["Background","Background","Background"],"Citation Start End":[[1124,1142]],"Functions Start End":[[1100,1122],[1151,1221],[1270,1332]]}
{"Identifier":"2019MNRAS.482.5651M__Schweizer_&_Middleditch_1980_Instance_1","Paragraph":"Therefore, the kinetics characteristics of the star could be the only piece to judge whether or not the SM star is the surviving companion of SN 1006. If its space velocity is significantly different from the other stars in the remnant of SN 1006, the probability to be the surviving companion would become high. Otherwise, the probability becomes low. We check the proper motion of the stars within 5\u2009arcmin of the remnant centre from Gaia DR2, as shown in Fig. 19. From the figure, it seems that there is not difference between the SM star and other stars in the remnant of SN 1006 in the aspect of proper motion, i.e. the proper motion of the SM star only slightly deviates from the median value of the proper motions of the stars at the direction of the SNR centre of SN 1006, and such a proper motion disfavours the SM star as the surviving companion of SN 1006 (Schweizer & Middleditch 1980; Burleigh et al. 2000). So, a 3D space velocity is helpful to judge the nature of the SM star. However, unfortunately, some data of the SM star in Gaia DR2 are so uncertain that we cannot use them to constrain its 3D space velocity, otherwise we could obtain a complete wrong conclusion.5 For example, the parallax of the SM star is \u03d6 = 0.0736 \u00b1 0.1244, and then \u03c3\u03d6\/\u03d6 = 1.69 which is much larger than the threshold value of 0.2 for distance estimation from GAIA DR2 data (Astraatmadja & Bailer-Jones 2016; Katz et al. 2018). The distance of the SM star from this parallax is much larger than all the previous measurements from spectrum by at least a factor of 2 (see summary in Burleigh et al. 2000). Considering that some other astrometric measurements of the SM star are also very uncertain, we applied the measurements in the previous literatures as the distance of the SM star. Based on a radial velocity of $-13\\pm 17\\, {\\rm km^{\\rm -1}}$ and a distance of 2.07 \u00b1 0.18 kpc (Schweizer & Middleditch 1980; Winkler et al. 2003; Kerzendorf et al. 2018), we can calculate the UVW velocities of the SM star, i. e. $U=-5.2\\pm 14\\, {\\rm km^{\\rm -1}}$, $V=197\\pm 10\\, {\\rm km^{\\rm -1}}$, and $W=3.1\\pm 5\\, {\\rm km^{\\rm -1}}$. The V value of the SM star is smaller than that of a normal disc star. We then transform these velocities into the Galactic rotational velocity at a Galactocentric distance of \u223c6.67 kpc, i.e. $V_{\\rm c}=196\\pm 12\\, {\\rm km^{\\rm -1}}$, which is smaller than the Galactic rotational velocity of the disc stars at the Galactocentric distance by $50\\pm 19\\, {\\rm km^{\\rm -1}}$ (Huang et al. 2016). This velocity difference is marginally consistent with the predicted orbital velocity here (see Fig. 7). In addition, the smaller rotational velocity of the SM star may explain its small proper motion shown in Fig. 19. So, the SM star is still possible to be the surviving companion of SN 1006.","Citation Text":["Schweizer & Middleditch 1980"],"Functions Text":["From the figure, it seems that there is not difference between the SM star and other stars in the remnant of SN 1006 in the aspect of proper motion, i.e. the proper motion of the SM star only slightly deviates from the median value of the proper motions of the stars at the direction of the SNR centre of SN 1006, and such a proper motion disfavours the SM star as the surviving companion of SN 1006"],"Functions Label":["Similarities"],"Citation Start End":[[868,896]],"Functions Start End":[[467,866]]}
{"Identifier":"2022AandA...666A.134S__Rodr\u00edguez-Almeida_et_al._2021b_Instance_1","Paragraph":"A more puzzling task was the resolution of the fine and hyperfine structure. First, the A\u2013E split structure attributed to the methyl internal rotation motion can be generally used as a molecular \u201cfingerprint\u201d to search for a molecule in an astronomical line survey (Cernicharo et al. 2016; Belloche et al. 2017). Moreover, it is known that for interstellar searches using surveys that done at centimeter wavelengths (i.e., those conducted with telescopes such as GBT and Yebes 40m), the recognition of hyperfine patterns in the observed spectra significantly help in the proper identification of interstellar molecules (McCarthy & McGuire 2021). As an example, the analysis of the hyperfine structure of several N-bearing species, such as N-protonated isocyanic acid (H2NCO+, Rodr\u00edguez-Almeida et al. 2021b) which has been detected toward G+0693-0.027 along with ethyl isocyante (C2H5NCO, Kolesnikov\u00e1 et al. 2018), was required for its conclusive line-by-line identification. In this context, sensitive data at centimeter wavelengths are crucial because in TMC-1 and in G+0.693-0.027 the excitation temperatures of the molecules are very low, and hence the brightest lines of this relatively complex species fall at those frequencies. These interstellar sources have recently been suggested to be among the best targets to search for large molecular systems (McGuire et al. 2021; Rivilla et al. 2021b), further supporting centimeter-wave studies of as-yet-undiscovered interstellar candidates. Hence, we employed our narrowband LA-MB-FTMW spectrometer (Le\u00f3n et al. 2021) to completely resolve both fine and hyperfine structures. We initially used the SPFIT\/SPCAT program package (Pickett 1991) to analyze hyperfine patterns for the previously measured A and E lines of Z-AHA (see Fig. 3). This scrutiny was followed by measuring additional hyperfine components of several a- and b-type R-branch lines. Once we completed the analysis, new R-branch transitions belonging to the E-conformer were identified.","Citation Text":["Rodr\u00edguez-Almeida et al. 2021b"],"Functions Text":["As an example, the analysis of the hyperfine structure of several N-bearing species, such as N-protonated isocyanic acid (H2NCO+,",") which has been detected toward G+0693-0.027","was required for its conclusive line-by-line identification."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[776,806]],"Functions Start End":[[646,775],[806,851],[915,975]]}
{"Identifier":"2021ApJ...912...92F__Feng_et_al._2021_Instance_1","Paragraph":"The sudden brightening of NGC 2617 provides a good opportunity to investigate the properties of CL-AGNs. In general, the CL processes can be interpreted as changes in obscuration or accretion rate. The obscuration model can be well applied to some X-ray selected CL-AGNs (e.g., Bianchi et al. 2005). The crossing time for an intervening object orbiting outside a BLR can be estimated as (LaMassa et al. 2015)\n2\n\n\n\n\n\nwhere rorb is the orbital radius of the foreground object on a circular, Keplerian orbit around the central black hole, M8 is the black hole mass in units of 108M\u2299, and rsrc is the true size of the BLR. The above equation is derived by evaluating the time needed for this object to travel the length of an arc that corresponds to the projected size of the BLR. The observational constraints of rdt and rdt\/rBLR are discussed in estimating tcross for the double-peaked BEL CL-AGN NGC 3516, where rdt is a dust torus radius (Feng et al. 2021, and references therein). rdt\/rBLR \u223c 4\u20135 is obtained for most reverberation-mapped AGNs with the H\u03b2 BLR and dust torus lag measurements. NGC 3516 has rdt\/rBLR = 7. Thus, it seems appropriate to use rdt\/rBLR = 5 in the CL-AGN NGC 2617. NGC 2617 has averages of rBLR = 6.15 lt-day and M8 = 0.22 for H\u03b1 and H\u03b2, and we obtain tcross = 7 yr. NGC 2617 is Seyfert 2 in 1994 and on 2003 December 30, and Seyfert 1 on 2013 April 25 (see Figure 2 in Shappee et al. 2014). Our observations show a Seyfert 1 type of optical spectra from 2019 October to 2020 May. The poor sampling of long-term spectroscopic observations makes it impossible to give the reliable durations of the appearance and \u201cdisappearance\u201d of BELs. A long-term monitoring with a good sampling, e.g., twice a year, may give a reliable duration of \u201cdisappearance\u201d of BELs in the CL processes, especially, the continuous and complete process of changing of Seyfert 1 \u2192 Seyfert 2 \u2192 Seyfert 1. This will be important to test the obscuration model of CL-AGNs.","Citation Text":["Feng et al. 2021"],"Functions Text":["The observational constraints of rdt and rdt\/rBLR are discussed in estimating tcross for the double-peaked BEL CL-AGN NGC 3516, where rdt is a dust torus radius"],"Functions Label":["Uses"],"Citation Start End":[[939,955]],"Functions Start End":[[777,937]]}
{"Identifier":"2021MNRAS.506.3330W__Gronow_et_al._2020_Instance_1","Paragraph":"The majority of observational evidence based on studies of SN Ia populations thus points towards a DD scenario for most, if not all, SNe Ia. However, finding self-consistent progenitor and explosion models that recreate the observed luminosity function as well as correlations between luminosity, light-curve parameters, and host galaxy properties has proven difficult. In particular, simulations based around explosions of MCh WDs (linked strongly with the SD scenario but also with many DD scenarios) find difficulty in reproducing the light curves of \u2018normal\u2019 SNe Ia as well as \u2018peculiar\u2019 objects (Ropke et al. 2007; Sim et al. 2013; see Maoz et al. 2014; Blondin et al. 2017; Jha, Maguire & Sullivan 2019 for overviews). In recent years, attention directed towards explosions of sub-MCh WDs triggered by double detonations (primarily related to a DD scenario) has led to promising results (e.g. Shen, Toonen & Graur 2017; Shen et al. 2018; Townsley et al. 2019; Gronow et al. 2020; Shen et al. 2021) although they still struggle to match observations at late times in the light-curve evolution (Gronow et al. 2021; Shen et al. 2021). An additional factor in support of the sub-MCh model is that the SN luminosity is related to the mass of the primary WD, which itself is likely to be related to its age (although this relation is probably complicated by other factors such as accretion rate, metallicity, and the composition of the companion), thereby providing an explanation for observed relation between light-curve stretch and stellar age (Rigault et al. 2013, 2020; Rose, Garnavich & Berg 2019; Nicolas et al. 2021). Other proposed scenarios include hybrid models in which standard CO WDs merge with hybrid helium-CO WDs (Zenati, Toonen & Perets 2019). With many models showing promising similarities to observations but each subject to its own drawbacks, it is becoming accepted that more than one progenitor scenario may contribute significantly to the overall population of \u2018normal\u2019 SNe Ia; detailed observations are thus required in order to place constraints on the relative fractions of each possible progenitor channel.","Citation Text":["Gronow et al. 2020"],"Functions Text":["In recent years, attention directed towards explosions of sub-MCh WDs triggered by double detonations (primarily related to a DD scenario) has led to promising results (e.g."],"Functions Label":["Background"],"Citation Start End":[[966,984]],"Functions Start End":[[725,898]]}
{"Identifier":"2021AandA...653A.129C__Coutens_et_al._(2016)_Instance_1","Paragraph":"Figure 15 shows the comparison between pairs of molecular abundances, HC(O)NH2 and HNCO, CH3NCO and HNCO, and CH3NCO and HC(O)NH2. It is already known from previous observations that there is a correlation between HNCO and HC(O)NH2 (L\u00f3pez-Sepulcre et al. 2015, 2019; Allen et al. 2020). In the top-left panel of Fig. 15 the best power-law fits derived by L\u00f3pez-Sepulcre et al. (2015), X[HC(O)NH2] = 0.04 \u00d7 X[HNCO]0.93, and by Qu\u00e9nard et al. (2018), X[HC(O)NH2] = 32.14 \u00d7 X[HNCO]1.29, are compared with the one derived here, X[HC(O)NH2] = 0.006 \u00d7 X[HNCO]0.73 (with a Pearson coefficient of 0.99 and a P-value 0.05, indicating a strong positive correlation). Thus, the sample of sources discussed in this work, which also includes HMCs and a shock-dominated molecular cloud, is in agreement with the correlation found previously for low- and intermediate-mass pre-stellar and protostellar objects, which holds across several orders of magnitude in abundance. Based on this tight correlation, it has been proposed that the two species are chemically related and that the formation of HC(O)NH2 might occur through H-addition to solid-phase HNCO (e.g. Tielens & Hagen 1982; Charnley et al. 2004). Experimental works first suggested that this process is not efficient (Noble et al. 2015; Fedoseev et al. 2015), while recent works revised this possibility and found that a correlation between these two molecular species can be understood by H-abstraction and addition reactions (e.g. Nguyen et al. 2011; Haupa et al. 2019; Suhasaria & Mennella 2020). Moreover, hydrogenation of NO combined with UV-photon exposure and radical-radical reactions on grains has been suggested as the main formation pathways for both HNCO and HC(O)NH2 (e.g. Jones et al. 2011; Fedoseev et al. 2016; Ligterink et al. 2018; Dulieu et al. 2019). Coutens et al. (2016) found that the deuteration (D\/H ratio) of HC(O)NH2 in IRAS 16293 B is similar to that of HNCO, in agreement with the hypothesis that both species are chemically related via grain-surface reactions. Gas-phase formation routes have also been proposed (see e.g. NH2 + H2CO, Barone et al. 2015; Skouteris et al. 2017). Laboratory experiments by Mart\u00edn-Dom\u00e9nech et al. (2020) show that both HNCO and HC(O)NH2 could form upon UV photoprocessing or electron irradiation of ice samples, indicating that energetic processing (like UV photons and cosmic rays) of ISM CO-rich ices could form both species, without the need of a chemical link and\/or a similar precursor between the two. This was predicted by the chemical modelling of Qu\u00e9nard et al. (2018), who showed that the formation of HC(O)NH2 at different temperature regimes is governed by different chemical processes. While at low temperatures the formation of HC(O)NH2 is driven by gas-phase formation via the reaction NH2 + H2CO \u2192 HC(O)NH2 + H, at high temperature its formation occurs on the surface of dust grains via radical-radical addition reactions. Moreover, they showed that for HNCO grain-surface and gas-phase reactions are equally efficient at low temperature, while at high temperatures the gas-phase formation predominates and the small fraction formed on grains is released into the gas phase via thermal desorption. Rimola et al. (2018) also showed via theoretical quantum chemical computations that HC(O)NH2 can form on grain surfaces starting from CN, which can quickly react with water-rich amorphous ices. Thus, the correlation between HNCO and HC(O)NH2 is mainly due to a similar response to the temperature of the two molecules, and not to a direct chemical link. In fact, the increase of the temperaturetriggers processes on the ice-mantle of grains, such as thermal evaporation. Moreover, as discussed above, other processes, like UV photons, cosmic rays, and shocks, could help both on the formation of these molecules on grain surfaces and on their desorption in the gas.","Citation Text":["Coutens et al. (2016)"],"Functions Text":["found that the deuteration (D\/H ratio) of HC(O)NH2 in IRAS 16293 B is similar to that of HNCO, in agreement with the hypothesis that both species are chemically related via grain-surface reactions."],"Functions Label":["Background"],"Citation Start End":[[1816,1837]],"Functions Start End":[[1838,2035]]}
{"Identifier":"2019ApJ...875...68A__Clemens_&_Alexander_2002_Instance_1","Paragraph":"Associations between ULXs and star clusters have also been studied in interacting galaxies. The latter are known to host a higher average number (>5) of ULXs. Therefore, they are good candidates in which to examine the properties of the population of ULXs. Poutanen et al. (2013) extensively examined the significant associations between ULXs and stellar clusters in the Antennae galaxies. Using data from HST and the Very Large Telescope supplemented with theoretical stellar isochrones, they estimated the ages of these clusters as 6 Myr. It was discussed that these ULXs were probably ejected from the cluster in the evolutionary process; thus these sources might be high-mass X-ray binaries instead of intermediate-mass black holes. Another well-known interacting galaxy is NGC 4490\/NGC 4485 at a distance of 7.8 Mpc (Tully 1988). NGC 4490 is a late-type spiral galaxy and NGC 4485 is an irregular galaxy. Their linear sizes are 15 kpc for NGC 4490 and 5.6 kpc for NGC 4485. Radio observations show that star formation in NGC 4490 has been ongoing at a constant rate of \u223c4.7 M yr\u22121 (Clemens et al. 1999). Also, NGC 4490 has a giant H i envelope that probably originated from the star formation (Clemens & Alexander 2002). Our aim in this study is to identify the possible optical counterparts of ULXs in this galaxy pair and to investigate their associations with star groups or clusters. Previously, three ULXs in this pair were detected by ROSAT HRI observations (Roberts & Warwick 2000). Later, using a Chandra ACIS-S observation, three more ULXs were identified by Roberts et al. (2002). The calculated unabsorbed luminosities of six ULXs (in the 0.5\u20138 keV band) fall into the range (2.6\u20134.9) \u00d7 1039 erg s\u22121. In addition, Roberts et al. noted that these ULXs appear to be spatially coincident with the star formation regions in the pair. Further, Fridriksson et al. (2008) searched for long-term variability of 38 X-ray sources in this galaxy pair using three Chandra observations. Eight of these sources were classified as ULXs in the luminosity range (0.6\u20133) \u00d7 1039 erg s\u22121. One of them is a transient ULX detected in a single observation (ID 4726). Gladstone & Roberts (2009) investigated spectral and temporal features of seven ULXs (except for ULX X-5\u2014this source was ignored because of its low luminosity of LX \u223c 6 \u00d7 1038 erg s\u22121, which was given in Table 5 of Fridriksson et al. 2008) using the same Chandra and XMM-Newton data sets. The LX values of these sources are given in the range (0.9\u20134) \u00d7 1039 erg s\u22121 within the 0.5\u22128 keV energy band. Six of these seven sources (except the transient one, which is X-7 in this paper) were classified as ULXs by Swartz et al. (2011).","Citation Text":["Clemens & Alexander 2002"],"Functions Text":["Also, NGC 4490 has a giant H i envelope that probably originated from the star formation"],"Functions Label":["Background"],"Citation Start End":[[1199,1223]],"Functions Start End":[[1109,1197]]}
{"Identifier":"2019AandA...628A.110M__Kryukova_et_al._(2012)_Instance_1","Paragraph":"Deriving the completeness limits of the WISE photometry is mandatory to assess the reliability of our catalogue of starless cores. We examined the histograms of the number of mid-infrared (MIR) sources versus magnitude; taking into account the effects of the cuts required to fulfil the criteria of Koenig et al. (2012), rough completeness limits are [3.6] ~ 14, [4.6] ~ 12, [12] ~ 9 and [22] ~ 7. These values are 1\u20133 mag brighter than the sensitivity limits quoted in the WISE Explanatory Supplement3 for the relevant sky region. Once converted into flux units and, for example, compared with the models of Class I and Class 0 sources of 0.5 M\u2299 by Whitney et al. (2004), it can be seen that the completeness limits at 3.6 and 4.6\u03bcm are faint enough to detect such objects taking into account a distance of 700 pc and a further foreground reddening up to AV = 20. Even in the worst case of edge-on discs, these objects would be detectable at 3.6 and 22\u03bcm. Furthermore, the completeness limit at 22 \u03bcm is faint enough to allow detection of Class I and Class 0 sources of even-lower-mass central objects. Alternatively, one can compute the bolometric luminosity following Kryukova et al. (2012). Starting from our completeness limit at 22 \u03bcm, after conservatively dereddening it by AV = 20, we assumed a spectral index \u03b3 = \u22122 (see Table 3 for definition) to compute the MIR luminosity from Eq. (6) of Kryukova et al. (2012). Equation (7) of Kryukova et al. (2012) then yields Lbol ~ 1.7\u20132.8 L\u2299, depending on whether the NIR flux is neglected (which may be the case) or extrapolated from \u03b3 = \u22122. A comparison with the birthline of Palla & Stahler (1993) indicates a mass of ~ 0.4\u20130.5 M\u2299 for the central protostar. For the sake of comparison, we can roughly estimate the completeness limit in central masses of the Herschel protostellar cores in Giannini et al. (2012) using their quoted completeness limit at 70 \u03bcm of 0.21 Jy and following Dunham et al. (2008). By using Eq. (2) of Dunham et al. (2008), scaled to a distance of 700 pc, we found that the flux density at 70 \u03bcm translates into a bolometric luminosity of the central (proto)star Lbol ~ 0.28 L\u2299 (we note that Dunham et al. 2008 indicate this luminosity as Lint). We highlight the fact that the 70 \u03bcm emission is in principle a more sensitive protostellar tracer than WISE. However, this contrasts with the much lower number of protostellar cores found by Giannini et al. (2012), which may be due to a poorer effective sensitivity because of their selection criteria.","Citation Text":["Kryukova et al. (2012)"],"Functions Text":["Alternatively, one can compute the bolometric luminosity following"],"Functions Label":["Uses"],"Citation Start End":[[1171,1193]],"Functions Start End":[[1104,1170]]}
{"Identifier":"2021AandA...647A.186S__Shebalin_et_al._1983_Instance_1","Paragraph":"Chandrasekhar & Fermi (1953) assumed that \u201cturbulent motions are isotropic\u201d, and they adopted \n\n$f\\,{=}\\,1\/\\sqrt{3}$f\u2009=\u20091\/3\n. If the field strength is weak, turbulent motions will drag the field lines in random directions and turbulence will be isotropic (super-Alfv\u00e9nic turbulence). However, there is overwhelming observational evidence that magnetic fields in the ISM have well-defined directions, indicating that turbulence is sub, trans-Alfv\u00e9nic and hence turbulent properties are highly anisotropic (see, for example, Montgomery & Turner 1981; Shebalin et al. 1983; Higdon 1984; Sridhar & Goldreich 1994; Goldreich & Sridhar 1995, 1997). Heyer et al. (2008), using CO data, found that velocity structures in Taurus are highly anisotropic. In the same region, Goldsmith et al. (2008) reported the existence of highly anisotropic density structures, which are aligned parallel to the mean field, known as striations. Striations have also been observed in the Polaris Flare (Panopoulou et al. 2015) and Musca (Cox et al. 2016; Tritsis & Tassis 2018), and they are formed due to magnetosonic waves (Tritsis & Tassis 2016) in sub-Alfv\u00e9nic turbulence (Beattie & Federrath 2020). More evidence for ordered magnetic fields in molecular clouds can be found in Franco et al. (2010), Franco & Alves (2015), Pillai et al. (2015), Hoq et al. (2017), and Tang et al. (2019). Stephens et al. (2011) explored the magnetic field properties of 52 star forming regions in our Galaxy and concluded that more than 80% of their targets exhibit ordered magnetic fields. The diffuse atomic clouds in our Galaxy are preferentially aligned with the magnetic field (Clark et al. 2014), implying the importance of the magnetic field in their formation. Planck Collaboration Int. XXXV (2016) studied a larger sample of molecular clouds in the Goult Belt and concluded that density structures align parallel or perpendicular to the local mean field direction. This is also consistent with sub, trans-Alfv\u00e9nic turbulence (e.g., Soler et al. 2013). In addition, Mouschovias et al. (2006), using Zeeman data, concluded that turbulence in molecular clouds is slightly sub-Alfv\u00e9nic as well. All these lines of evidence indicate that ISM turbulence is sub, trans-Alfv\u00e9nic, and hence anisotropic.","Citation Text":["Shebalin et al. 1983"],"Functions Text":["However, there is overwhelming observational evidence that magnetic fields in the ISM have well-defined directions, indicating that turbulence is sub, trans-Alfv\u00e9nic and hence turbulent properties are highly anisotropic"],"Functions Label":["Background"],"Citation Start End":[[549,569]],"Functions Start End":[[284,503]]}
{"Identifier":"2021AandA...650A.155Z__Oh_et_al._2012_Instance_3","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"],"Functions Text":["However, the question of whether galactic bars can significantly affect AGN activity is still under debate"],"Functions Label":["Background"],"Citation Start End":[[903,917]],"Functions Start End":[[756,862]]}
{"Identifier":"2015ApJ...811..129B__Pasachoff_et_al._2009_Instance_1","Paragraph":"An alternative possibility to coronal heating is that the release and dissipation of magnetic energy actually takes place in the chromosphere. In this scenario, the chromospheric plasma is directly heated to coronal temperatures, rather than by thermal conduction fronts that are driven into the lower atmosphere as a consequence of heating localized in the corona (Hansteen et al. 2010). One manifestation of this process might be the type II spicules (De Pontieu et al. 2007) that have been suggested to play a role in supplying mass and energy to the corona (De Pontieu et al. 2009, 2011; Moore et al. 2011). N. E. Raouafi et al. (2015, in preparation) have argued that type II spicules have much more in common with so-called \u201cclassical\u201d spicules (Beckers 1968, 1972; Pasachoff et al. 2009) than do type I spicules. In particular, type II and classical spicules are both commonly observed in quiet Sun and coronal hole regions, whereas type I spicules appear to be exclusively confined to active regions. Furthermore, Pereira et al. (2013) artificially coarsened Hinode\/SOT observations of type II spicules to demonstrate that their lifetimes and ejection speeds are consistent with the earlier, ground-based observations of classical spicules (\u223c5 minutes and 25 km s\u22121). (N. E. Raouafi et al. 2015, in preparation) suggest that the term \u201cclassical spicules\u201d be used for the earlier objects (see Sterling 2000, for a review) and the terms type I and II spicules be reserved for spicular phenomena observed during the era of Hinode observations. The key properties of type II spicules are their faster velocities (30\u2013110 km s\u22121) and shorter lifetimes (50\u2013150 s) compared with type I spicules. Type II spicules are typically observed in Ca ii emission by SOT before fading out of that passband and appearing in the SDO\/AIA 304 \u212b channel, indicating that some degree of heating takes place as they rise. There is some observational evidence for a transition region or coronal component of their emission, visible as a bright, moving front in the AIA 171 \u212b channel (e.g., De Pontieu et al. 2011). The two possibilities for producing this warmer emission are: (1) pre-existing coronal material is shock-heated as the upflowing spicular material rams into it (Klimchuk 2012; Petralia et al. 2014a); or (2) the tip of the spicule is heated by some in situ process, which must be impulsive because the cooler emission quickly disappears from view (De Pontieu et al. 2007).","Citation Text":["Pasachoff et al. 2009"],"Functions Text":["N. E. Raouafi et al. (2015, in preparation) have argued that type II spicules have much more in common with so-called \u201cclassical\u201d spicules","than do type I spicules. In particular, type II and classical spicules are both commonly observed in quiet Sun and coronal hole regions, whereas type I spicules appear to be exclusively confined to active regions."],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[772,793]],"Functions Start End":[[612,750],[795,1008]]}
{"Identifier":"2019MNRAS.488.1728N__Kaiser_et_al._1995_Instance_1","Paragraph":"The mass estimate for each cluster was obtained by fitting a reduced tangential shear profile predicted by a projected Navarro\u2013Frenk\u2013White (NFW) profile (e.g. Bartelmann 1996) to the observed ellipticities. We derive the best-fitting profile parameters R200 and c200 by minimizing the merit function \n(4)\r\n\\begin{eqnarray*}\r\n\\chi ^{2}\\!=\\!\\sum _{i=1}^{N}{\\frac{\\left|g_{i}(\\theta _{i},\\beta _i;R_{200},c_{\\mathrm{NFW}})\\!-\\! \\tilde{\\varepsilon }_{\\mathrm{t},i}(\\theta _{i})\\right|^{2}}{\\tilde{\\sigma }_{\\!i}^{2}\\left(1\\!-\\!\\left| g_{i}(\\theta _{i},\\beta _i;R_{200},c_{\\mathrm{NFW}})\\right|^{2}\\right)^{2}}}. \r\n\\end{eqnarray*}\r\nHere gi(\u03b8i, \u03a3crit, i;; R200, cNFW) is the model prediction for galaxy i and $\\tilde{\\varepsilon }_{\\mathrm{t},i}$ the observed ellipticity times 1.08 for the same galaxy. The factor 1.08 is the multiplicative shear calibration bias of the used KSB+ pipeline (Kaiser et al. 1995; Erben et al. 2001) to convert from measured to true ellipticity. This calibration bias has an uncertainty of \u223c5 per cent. This uncertainty is a dominant source of systematic uncertainty in the mass measurements. Each shear profile was centred on the BCG, using distances in the range of 0.2\u20134.2 Mpc for the fitting procedure. We minimized the \u03c72 on a grid of R200 and c200. Finally, we used the mass\u2013concentration relation described by Bhattacharya et al. (2013) to put priors on the concentration parameter to break the degeneracies in the profile models. The initial mass estimates from equation (4) are biased. In evaluating the NFW shear profile we make use of the ratio in equation (2) when averaging the value of \u03b2 over the reference catalogue sources. However, $\\frac{\\gamma (\\langle \\beta \\rangle)}{1-\\kappa (\\langle \\beta \\rangle)}\\ne \\left\\langle \\frac{\\gamma (\\beta)}{1-\\kappa (\\beta)}\\right\\rangle$. Given the finite width of the \u03b2 distribution that are averaged over when calculating \u03b2i from a reference catalogue (equation 3), we find a biased point estimate for \u03b2i. Especially in the inner regions of the cluster, this would model the shear profile incorrectly. We estimate the final masses by correcting for the averaging over \u03b2 in two subsequent iterations. We utilize the best-fitting mass estimate from the zeroth iteration to predict the reduced shear, g, at the projected distance \u03b8 from the cluster centre, and \u03b2k. We then introduce $\\beta _i^{\\prime }$, which satisfies the equation: \n(5)\r\n\\begin{eqnarray*}\r\ng(\\beta _{i}^{\\prime })=\\!\\frac{\\sum _{k=1}^{N}w_kg(\\theta _{i},\\beta _k)}{\\sum _{k=1}^{N}w_k}\\frac{1}{v_{b}(c_{1},c_{2})}. \r\n\\end{eqnarray*}\r\n","Citation Text":["Kaiser et al. 1995"],"Functions Text":["The factor 1.08 is the multiplicative shear calibration bias of the used KSB+ pipeline","to convert from measured to true ellipticity."],"Functions Label":["Uses","Uses"],"Citation Start End":[[886,904]],"Functions Start End":[[798,884],[925,970]]}
{"Identifier":"2018AandA...618L...3M__Greaves_&_Rice_(2010)_Instance_1","Paragraph":"One way to explain the observations is to postulate that the cores of planets are formed in the very first Myr of the protoplanetary disk evolution, or even in the embedded phase while the disk is still forming. The disks for which masses have been measured have ages >1 Myr, and a general trend of declining disk mass with ages older than 1 Myr has been observed (e.g., Barenfeld et al. 2016). Thus, it is possible that disks were massive enough to form the cores of planets at younger ages. This idea has been suggested by Greaves & Rice (2010), Williams (2012), and Najita & Kenyon (2014), among others. In this scenario, the vast majority of the material composing planets must already be in the form of planetesimals, for rocky planet formation, and of already-formed planetary cores. The latter condition is necessary as gas giant planets need to accrete gas from the gas-rich disk. Assuming a gas-to-dust ratio of 100, disks at \u223c1\u20133 Myr have just the right amount of gas mass to explain the population of gas giants (see Fig. 1). Thus, cores must already be in place at this age. Scenarios have been proposed to explain that pebble accretion can form planetesimals very early (0.1 Myr) in disks, when these are massive and possibly gravitationally unstable (Booth & Clarke 2016). However, it is expected that the formation of planetary cores is highly inefficient, with \u223c350 M\u2295 of pebbles needed to grow the core of Jupiter from half a lunar mass to 20 M\u2295 (e.g., Morbidelli et al. 2016). Thus, such an inefficiency would imply that disks were initially \u223c10\u2013100 times more massive than is observed at ages >1 Myr. This would also imply that the vast majority of disks were initially gravitationally unstable. One issue of this scenario is that, if early disks are \u223c10\u2013100 times more massive than observed for disks older than 1 Myr, an extremely efficient gas-removal mechanism has to be found, consistent with the observations of Class 0 outflows (Frank et al. 2014).","Citation Text":["Greaves & Rice (2010)"],"Functions Text":["Thus, it is possible that disks were massive enough to form the cores of planets at younger ages. This idea has been suggested by","Williams (2012), and Najita & Kenyon (2014), among others. In this scenario, the vast majority of the material composing planets must already be in the form of planetesimals, for rocky planet formation, and of already-formed planetary cores. The latter condition is necessary as gas giant planets need to accrete gas from the gas-rich disk."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[525,546]],"Functions Start End":[[395,524],[548,888]]}
{"Identifier":"2021MNRAS.507.5882S__Mackereth_et_al._2018_Instance_1","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"],"Functions Text":["The bimodality has now been observed in some simulations"],"Functions Label":["Similarities"],"Citation Start End":[[983,1004]],"Functions Start End":[[906,962]]}
{"Identifier":"2020MNRAS.498.5116I__Andrews_et_al._2012_Instance_1","Paragraph":"It is now largely understood that the dust-to-gas mass ratio \u03f5 in protoplanetary discs is unlikely to match the canonical value in the interstellar medium of 0.01 (Mathis, Rumpl & Nordsieck 1977). For example, recent observations of CO isotopologues and dust continuum emission have found disc-averaged \u03f5 values to be much higher (\u223c0.2; see Ansdell et al. 2016), and simulations have demonstrated that discs can be formed with significant dust enrichment (Lebreuilly, Commer\u00e7on & Laibe 2020). There is also an observed discrepancy between the radial extents of the gas and dust components of protoplanetary discs, with dust discs being significantly more compact (Andrews et al. 2012; P\u00e9rez et al. 2012, 2015b). In some cases, such as IM Lup, the gas disc may be 10 times larger in radial extent than the mm dust disc (e.g. Cleeves et al. 2016). This discrepancy can explained by a combination of grain growth and radial drift. Essentially, the radial velocity, vr, of dust particles has two components, namely drag, vdrag, and drift, vdrift, and is given by\n(1)$$\\begin{eqnarray*}\r\nv_{\\mathrm{r}}=v_{\\mathrm{drag}}+v_{\\mathrm{drift}}=\\frac{v_{\\mathrm{g}, \\mathrm{r}}}{1+{St}^{2}}+\\frac{1}{{St}+{St}^{-1}} \\frac{1}{\\rho _{\\mathrm{g}} \\Omega } \\frac{\\partial P}{\\partial R},\r\n\\end{eqnarray*}$$where vg, r is the radial velocity of the gas, \u03c1g is the gas mass surface density, \u03a9 is the Keplerian angular velocity, and P is the pressure (Weidenschilling 1977). The Stokes number, St, determines the degree of coupling to the gas. For a vertically isothermal disc (as used here), this is expressed simply as\n(2)$$\\begin{eqnarray*}\r\nSt = \\frac{\\pi a\\rho _\\mathrm{s}}{2\\Sigma _\\mathrm{g}},\r\n\\end{eqnarray*}$$where a is grain size, \u03c1s is intrinsic bulk grain density, and \u03a3g is the gas surface density. Examining equations (1) and (2 ), we can see that the drag term decreases as grain size increases. However, the drift term peaks at St = 1, which means that the grain size for which radial drift is fastest depends on the surface density, and therefore location within the disc.","Citation Text":["Andrews et al. 2012"],"Functions Text":["There is also an observed discrepancy between the radial extents of the gas and dust components of protoplanetary discs, with dust discs being significantly more compact"],"Functions Label":["Background"],"Citation Start End":[[664,683]],"Functions Start End":[[493,662]]}
{"Identifier":"2018ApJ...858...13H___2009b_Instance_1","Paragraph":"It is widely accepted that an SN Ia is a thermonuclear explosion resulting from a binary system, of which one star is necessarily a degenerate carbon-oxygen (\n\n\n\n\n\n) white dwarf (WD) (Hoyle & Fowler 1960) near the Chandrasekhar mass (\n\n\n\n\n\n). Current research considers various progenitor configurations and final outcomes. Depending on whether or not both stars are WDs, the progenitor system is called double-degenerate (DD) or single-degenerate (SD). In addition to progenitor configuration, proposed scenarios are distinguishable by the ignition mechanism and other characteristics. In the case of a DD progenitor system, a dynamical merger or a violent collision between the WDs is capable of releasing enough heat to trigger an ignition. This process can end up as a SN Ia, a highly magnetized WD (MWD), or an accretion-induced collapse (AIC; Iben & Tutukov 1984; Webbink 1984; Benz et al. 1990; Rasio & Shapiro 1994; Segretain et al. 1997; Yoon et al. 2007; Lor\u00e9n-Aguilar et al. 2009; Wang et al. 2009a, 2009b; Isern et al. 2011; Pakmor et al. 2011). Another class of SNe Ia scenarios is the double-detonation of a sub-\n\n\n\n\n\n WD with accretion from a helium (\n\n\n\n\n\n) companion. A detonation in the surface helium layer causes a secondary detonation in the core (Weaver & Woosley 1980; Nomoto 1982a; Livne 1990; Woosley & Weaver 1994; Hoeflich & Khokhlov 1996; Kromer et al. 2010; Woosley & Kasen 2011). Finally, there is the \n\n\n\n\n\n explosion scenario, where the WD progenitor accretes material from a companion star and nuclear surface burning to C\/O leads to an increase of the WD mass. With increasing WD mass the electron gas in the central region becomes increasingly relativistic, which leads to faster compressional heat release, the rising of the central temperatures, and the triggering of a central C\/O deflagration front when the mass of the progenitor approaches \n\n\n\n\n\n. (Hoyle & Fowler 1960; Sugimoto & Nomoto 1980; Nomoto 1982b; Hoeflich & Stein 2002; Piersanti et al. 2003). It is likely that the dynamical merger, \n\n\n\n\n\n explosion, and double-detonation channels all contribute to the SN Ia population because of the \u201cstellar amnesia\u201d effect (Hoeflich (2006), and references therein). This can happen in either a SD system, where the donor star is a main-sequence star, a red giant, etc., or in a DD system with another WD being the donor (Whelan & Iben 1973; Piersanti et al. 2003).","Citation Text":["Wang et al.","2009b"],"Functions Text":["This process can end up as a SN Ia, a highly magnetized WD (MWD), or an accretion-induced collapse"],"Functions Label":["Background"],"Citation Start End":[[992,1003],[1011,1016]],"Functions Start End":[[744,842]]}
{"Identifier":"2018ApJ...853...34Z__Giebels_et_al._2007_Instance_2","Paragraph":"Several well-studied TeV blazars show rich spectral behavior in X-rays, which may represent the general behavior of the synchrotron peak of all AGN jets. The X-ray spectra are usually curved (Massaro et al. 2004) and can only locally be fitted by a power law. The spectral variation with flux can be complex (Zhang et al. 2002; Cui 2004). Generally, the spectrum hardens when the flux increases (e.g., Gliozzi et al. 2006; Xue et al. 2006; Tramacere et al. 2009), but photon indices can saturate at higher fluxes (Xue & Cui 2005; Giebels et al. 2007). The synchrotron peak usually moves to higher frequencies with increasing flux during outbursts (e.g., Pian et al. 1998), but no correlation between the break energy and the flux exists when a broken power law is adopted to fit the X-ray spectra (Xue & Cui 2005; Giebels et al. 2007; Garson et al. 2010). A cooling break in the spectrum of emitting particles cannot explain these features (Wierzcholska & Wagner 2016), and some special particle acceleration processes may be involved (Madejski & Sikora 2016). There are also energy-dependent lags between the variations of different energy bands. In some flares, soft bands lag behind hard bands (e.g., Zhang et al. 2002), while lags in the opposite direction can also happen (e.g., Ravasio et al. 2004; Sato et al. 2008). Hysteresis in the HR (hardness ratio)\u2013flux diagram is often used as a diagnostic of lags. Clockwise loops (e.g., Acciari et al. 2009; Kapanadze et al. 2016) in the HR\u2013flux plane are a sign of soft lags while counterclockwise loops (e.g., Tramacere et al. 2009) are a sign of hard lags. The same source can exhibit both clockwise and counterclockwise loops; the observed patterns are further complicated by the superposition of flares at different timescales (Cui 2004). The above knowledge of TeV blazars in the X-ray regime comes from studies focusing on timescales of hours to weeks. We will extend this kind of analysis to much smaller timescales in this paper.","Citation Text":["Giebels et al. 2007"],"Functions Text":["The synchrotron peak usually moves to higher frequencies with increasing flux during outbursts","but no correlation between the break energy and the flux exists when a broken power law is adopted to fit the X-ray spectra"],"Functions Label":["Background","Background"],"Citation Start End":[[814,833]],"Functions Start End":[[552,646],[673,796]]}
{"Identifier":"2022AandARv..30....6M__Magliocchetti_et_al._2018b_Instance_2","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"],"Functions Text":["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.,","but see further in this section for different points of view)"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1773,1799]],"Functions Start End":[[1451,1750],[1838,1899]]}
{"Identifier":"2016ApJ...817...12P__Kleeorin_et_al._2000_Instance_1","Paragraph":"Large-scale magnetic fields with strength of the order of 1\u201310 \u03bcG have been observed in disk galaxies (e.g., Beck et al. 1996; Fletcher 2010; Beck 2012; Beck & Wielebinski 2013; Van Eck et al. 2015). The origin of these fields can be explained through mean-field dynamo theory (Ruzmaikin et al. 1988; Beck et al. 1996; Brandenburg & Subramanian 2005a; Kulsrud & Zweibel 2008). The conservation of magnetic helicity is one of the key constraints in these models, and also leads to the suppression of the \u03b1-effect. The operation of the mean-field dynamo automatically leads to the growth of magnetic helicity of opposite signs between the large-scale and small-scale magnetic fields (Pouquet et al. 1976; Gruzinov & Diamond 1994; Blackman & Field 2002). To avoid catastrophic suppression of the dynamo action (\u03b1-quenching), the magnetic helicity due to the small-scale magnetic field should be removed from the system (Blackman & Field 2000, 2001; Kleeorin et al. 2000). Mechanisms suggested to produce these small-scale magnetic helicity fluxes are: advection of magnetic fields by an outflow from the disk through the galactic fountain or wind (Shukurov et al. 2006; Sur et al. 2007; Chamandy et al. 2014), magnetic helicity flux from anisotropy of the turbulence produced by differential rotation (Vishniac & Cho 2001; Subramanian & Brandenburg 2004, 2006; Sur et al. 2007; Vishniac & Shapovalov 2014), and through diffusive flux (Kleeorin et al. 2000, 2002; Brandenburg et al. 2009; Mitra et al. 2010; Chamandy et al. 2014). The outflow of magnetic helicity from the disk through dynamo operation leads to the formation of a corona (Blackman & Field 2000). According to Taylor's hypothesis, an infinitely conducting corona would resistively relax to force-free field configurations under the constraint of global magnetic helicity conservation (Woltjer 1960; Taylor 1974; Finn & Antonsen 1983; Berger & Field 1984; Mangalam & Krishan 2000). In this paper, we include advective and diffusive fluxes in a simple semi-analytic model of a galactic dynamo that transfers magnetic helicity outside the disk and consequently builds up a force-free corona in course of time. We first solve the time-dependent dynamo equations by expressing them as separable in variables r and z. The radial part of the dynamo equation is solved using an eigenvector expansion constructed using the steady-state solutions of the dynamo equation. The eigenvalues of the z part of the solution are obtained by solving a fourth-order algebraic equation, which primarily depends upon the turbulence parameters and the magnetic helicity fluxes. Once the dynamo solutions are written out as parametric functions of these parameters, the evolution of the mean magnetic field is computed numerically by simultaneously solving the dynamical equations for \u03b1-quenching and the growth of large-scale coronal magnetic helicity. Since the large-scale magnetic field lines cross the boundary between the galactic disk and the corona, the magnetic helicity of the large-scale magnetic field in the disk volume is not well defined. Hence we use the concept of gauge-invariant relative helicity (Finn & Antonsen 1983; Berger & Field 1984; Berger 1985) to estimate the large-scale magnetic helicity in the disk and the corona. Here the gauge-invariant relative helicity for the cylindrical geometry is calculated using the prescription given in Low (2006, 2011). We then investigate the dependence of the saturated mean magnetic field strength and its geometry on the magnetic helicity fluxes within the disk and the corresponding evolution of the force-free field in the corona.","Citation Text":["Kleeorin et al. 2000","Kleeorin et al. 2000"],"Functions Text":["To avoid catastrophic suppression of the dynamo action (\u03b1-quenching), the magnetic helicity due to the small-scale magnetic field should be removed from the system","Mechanisms suggested to produce these small-scale magnetic helicity fluxes are:","and through diffusive flux"],"Functions Label":["Background","Background","Background"],"Citation Start End":[[946,966],[1432,1452]],"Functions Start End":[[752,915],[969,1048],[1404,1430]]}
{"Identifier":"2020ApJ...896L..21R___2015_Instance_1","Paragraph":"Top left: the revised, CO(J = 1 \u2192 0)-based Mgas from VLASPECS confirm that z = 2\u20133 galaxies detected in the ASPECS survey (green circles; tentative detections are marked with a plus sign) closely follow the \u201cstar formation law\u201d (i.e., Mgas\u2013SFR relation) at high redshift. CO-detected main-sequence galaxies at similar redshifts from the PHIBBS1\/2 surveys (typically based on CO J = 3 \u2192 2, but using a metallicity-dependent conversion factor; Tacconi et al. 2018) and local galaxies from the xCOLD GASS CO(J = 1 \u2192 0) survey (Saintonge et al. 2017) are shown for comparison. Bottom left: same as the top left panel, but plotting the depletion time tdep against Mgas. All samples cover a similar range in tdep, but the average tdep for the (higher Mgas) high-z samples appear lower. Top right: the r31 brightness temperature ratio of VLASPECS galaxies (green circles) is similar to that of strongly lensed z \u223c 3 Lyman-break galaxies (red triangles; Riechers et al. 2010), z > 2 main-sequence galaxies from the PHIBSS survey (gray crosses; Bolatto et al. 2015), and z > 2 dusty star-forming galaxies (DSFGs; blue squares; compilation from Sharon et al. 2016; including data from Danielson et al. 2011; Ivison et al. 2011; Riechers et al. 2011a, 2011b, 2013; Thomson et al. 2012; Fu et al. 2013; Sharon et al. 2013, 2015; other DSFGs shown as light gray squares are from Nayyeri et al. 2017; Dannerbauer et al. 2019; Harrington et al. 2019; Leung et al. 2019; Sharon et al. 2019) and clustered DSFGs (dark gray squares; Bussmann et al. 2015; G\u00f3mez-Guijarro et al. 2019), but \u223c2 times higher on average than BzK-selected main-sequence galaxies at z \u223c 1.5 (magenta crosses; Daddi et al. 2015). Nearby galaxy samples from the xCOLD GASS survey (Lamperti et al. 2020) and two studies of infrared-luminous galaxies (Yao et al. 2003; Papadopoulos et al. 2012) are shown for comparison. Dashed lines and shaded regions indicate mean\/median values and spread for high-z samples with >2 galaxies or clusters, with the same color coding as the symbols. Dashed\u2013dotted lines indicate mean values for the low-z samples. Bottom right: same as the top right panel, but shown as binned histograms in r31 (excluding upper limits) and across the full redshift range, and only including samples for which mean\/median values are indicated in the top right panels.","Citation Text":["Sharon et al.","2015"],"Functions Text":["Top right: the r31 brightness temperature ratio of VLASPECS galaxies (green circles) is similar to that of","and z > 2 dusty star-forming galaxies (DSFGs; blue squares; compilation from Sharon et al. 2016; including data from"],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[1291,1304],[1311,1315]],"Functions Start End":[[780,886],[1058,1174]]}
{"Identifier":"2021MNRAS.502.5038N__Gosset_et_al._1994_Instance_1","Paragraph":"Several searches for high-frequency signals were performed in the past for WR stars, notably in the context of searches for compact companions (e.g. Marchenko et al. 1994). However, there were few cases of reported and confirmed periodicities. Focusing on our sample, the following detections were published. Blecha, Schaller & Maeder (1992) claimed a detection of pulsations with a 627-s period and a 5-mmag peak-to-peak amplitude in WR\u200940. However, at the corresponding frequency of 138\u2009d\u22121, the TESS high-cadence data (sector 10) show no signs of such a signal, confirming previous negative reports (Gosset et al. 1994; Marchenko et al. 1994; Martinez et al. 1994; Schneider et al. 1994), including that by the discovery team itself (Bratschi & Blecha 1996). Other low-frequency detections for that star (e.g. Antokhin et al. 1995) certainly correspond to red-noise stochastic variability. Another photometric campaign detected a 6.828\u2009d\u22121 signal in WR\u200966 (Antokhin et al. 1995), but with strong daily aliasing. It was subsequently confirmed in an independent, less aliased dataset (largest peak at 5.815\u2009d\u22121: Rauw et al. 1996). These values are close (but not identical) to the frequency of the largest TESS peak, or its daily alias. To assess the compatibility between datasets, we have simulated a signal composed of the three main frequencies detected by TESS but sampled as in Rauw et al. (1996). The resulting periodogram appears similar to the one reported in the literature, even if the passbands are different6: the old and new datasets therefore appear fully compatible. Note also that the presence of several frequencies casts further doubt on the hypothesis of the literature signal being the orbital period of a compact companion (Antokhin et al. 1995). Finally, a 25-min signal with 2.6-mmag amplitude was reported by Bratschi & Blecha (1996) for WR\u200978 but considered as a transient event, as it was observed only during one night. Because WR\u200978 was not observed with a 2-min cadence, we cannot check for the presence of such a frequency in the TESS data.","Citation Text":["Gosset et al. 1994"],"Functions Text":["However, at the corresponding frequency of 138\u2009d\u22121, the TESS high-cadence data (sector 10) show no signs of such a signal, confirming previous negative reports"],"Functions Label":["Similarities"],"Citation Start End":[[603,621]],"Functions Start End":[[442,601]]}
{"Identifier":"2018MNRAS.475.1160H__Werk_et_al._2013_Instance_1","Paragraph":"Galaxies are surrounded by vast gaseous haloes which extend well beyond the hosts\u2019 stellar components: Early observations of quasar sight lines attributed the presence of absorption at multiple intermittent redshifts to gaseous haloes of intervening galaxies (e.g. Bergeron 1986; Bergeron & Boiss\u00e9 1991; Lanzetta et al. 1995; Tripp, Savage & Jenkins 2000; Chen, Lanzetta & Webb 2001). In the past decade, owing to the rise of large spectroscopic surveys of galaxies with well-determined physical properties (e.g. SDSS), all sky UV surveys (e.g. GALEX), and improved sensitivity of UV spectrographs (e.g. COS), studies of the gaseous haloes of galaxies could systematically connect gas absorption properties to galaxy properties in statistically meaningful samples (e.g. Cooksey et al. 2010; Prochaska et al. 2011; Tumlinson et al. 2013; Liang & Chen 2014; Lehner, Howk & Wakker 2015). The aforementioned gaseous haloes are commonly referred to as the circum-galactic medium (CGM) and are ubiquitous in galaxies regardless of mass or star formation activity: even sub-L* galaxies (Bordoloi et al. 2014), and passive galaxies host a CGM (Thom et al. 2012). The current model of the CGM suggests the presence of a clumpy multiphase medium which extends beyond the virial radius of the host galaxy, with a declining radial density profile, containing a substantial amount of gas and metals (e.g. Werk et al. 2013, 2014; Liang & Chen 2014; Lehner et al. 2014, 2015; Prochaska et al. 2017). Observational studies targeting the CGM of L* galaxies showed that the CGM gas content is comparable to the mass of the interstellar medium (ISM; e.g. Chen et al. 2010; Tumlinson et al. 2011; Werk et al. 2014; Prochaska et al. 2017) and correlates positively with ISM properties (Borthakur et al. 2015). Additionally, CGM observations infer a significant amount of metals (e.g. Werk et al. 2013; Peeples et al. 2014) where CGM metallicities can extend to supersolar metallicities (Prochaska et al. 2017). The clumpy multiphase CGM consists of a warm gas T \u223c 104 \u2212 5\u2009K (clumpy in nature) embedded within a hot diffuse T \u223c 106\u2009K medium (e.g. Heitsch & Putman 2009; Armillotta et al. 2017; Bordoloi et al. 2017). The multiphase structure of the CGM is corroborated by the variety of observed ionic species which survive at a vast range of temperatures: While the warm gas hosts the low ionization species (e.g. H\u2009i, Si\u2009ii, Si\u2009iii, C\u2009ii, C\u2009iv), the hot medium is home for the most highly ionized species (e.g. O\u2009vi, O\u2009vii). Additionally, the spectral line profiles of absorbers in the CGM can be reproduced by invoking a patchy medium (e.g. Stern et al. 2016; Werk et al. 2016), i.e. multiple high density gas clouds contribute to the optical depth along the line of sight thus leaving their kinematic imprint on the absorption line profile. For a review of the CGM, see Putman, Peek & Joung (2012) and Tumlinson, Peeples & Werk (2017).","Citation Text":["Werk et al. 2013"],"Functions Text":["The current model of the CGM suggests the presence of a clumpy multiphase medium which extends beyond the virial radius of the host galaxy, with a declining radial density profile, containing a substantial amount of gas and metals (e.g."],"Functions Label":["Background"],"Citation Start End":[[1392,1408]],"Functions Start End":[[1155,1391]]}
{"Identifier":"2015MNRAS.454..193B__Langhoff_et_al._1998_Instance_1","Paragraph":"Theoretical quantum chemical calculations help in narrowing down candidates for much more expensive laboratory experiments. Considering the high cost, time consumption and other constraints faced in laboratory, theoretical computational study can propose selected PAHs for which laboratory spectroscopy can most usefully be performed. Density Functional Theory (DFT) has been used rigorously to calculate the harmonic frequencies and intensities of vibrational modes of PAHs in various forms including size, composition and charge states (Langhoff 1996; Bauschlicher & Langhoff 1997; Bauschlicher et al. 1997; Langhoff et al. 1998; Hudgins, Bauschlicher & Allamandola 2001; Hudgins, Bauschlicher & Sandford 2004; Pathak & Rastogi 2005, 2006, 2007; Pathak & Sarre 2008; Candian, Sarre & Tielens 2014). In this work, DFT in combination with a B3LYP functional and a 6\u2013311G** basis set has been used to optimize the molecular structures of PAHs. The optimized geometry is used to obtain the vibrational frequencies of various modes at the same level of theory. Theoretical calculations tend to overestimate the frequency compared to experiments (Langhoff 1996). The use of a larger basis set, e.g. 6\u2013311G**, generally reduces the overestimation compared to smaller basis sets (Langhoff 1996). Use of a larger basis set compared to a smaller one also shows good agreement with experiment. However, the use of a large basis set does not support use of a single scaling factor for all of the vibrational modes (Langhoff 1996). In order to evaluate the mode-dependent scaling factors, calibration calculations were made for selected PAHs, both neutral and ionized. On comparing the theoretical frequencies with matrix isolated spectroscopic experimental data (Hudgins & Allamandola 1995a,b; Hudgins & Sandford 1998a,b), three different scaling factors have been determined. The scaling factors obtained are 0.974 for the C\u2013H out-of-plane (oop) mode, 0.972 for the C\u2013H in-plane and C-C stretching modes and 0.965 for the C\u2013H stretching mode. Gaussian line shapes of 30\u2009cm\u2212 1 FWHM are used to plot the computationally obtained spectra. Our sample includes deuteronated pyrene (DC16H$_{10}^+$), deuteronated perylene (DC20H$_{12}^+$) and deuteronated coronene (DC24H$_{12}^+$). Isomers of DC16H$_{10}^+$ and DC20H$_{12}^+$ have also been included. The data presented here were produced using gamess quantum chemistry suite of programs (Schmidt et al. 1993).","Citation Text":["Langhoff et al. 1998"],"Functions Text":["Density Functional Theory (DFT) has been used rigorously to calculate the harmonic frequencies and intensities of vibrational modes of PAHs in various forms including size, composition and charge states"],"Functions Label":["Background"],"Citation Start End":[[610,630]],"Functions Start End":[[335,537]]}
{"Identifier":"2015MNRAS.450.2404G__Fornasa_&_S\u00e1nchez-Conde_2015_Instance_1","Paragraph":"The extragalactic sky from any direction and at all frequencies is filled with radiation from discrete sources and from a diffuse (or unresolved) component known as the cosmic background. This pervasive radiation, discovered only in relatively recent times, is one of the most fundamental observables from the Universe, as it carries crucial information on the integrated radiation emitted over the entire cosmic history. The first extragalactic background to be detected was the cosmic X-ray background (CXB; Giacconi et al. 1962), a few years before the discovery of the cosmic microwave background (CMB; Penzias & Wilson 1965), which is much brighter in terms of energy density, and is truly diffuse. The extragalactic \u03b3-ray background (EGB) was first detected by the SAS2 satellite in the 1970's (Fichtel et al. 1977), and for a long time its nature was poorly understood. The latest Fermi Large Area Telescope (Fermi-LAT) survey data are now providing strong evidence towards an origin mostly due to integrated radiation from blazars and radio galaxies (e.g. Inoue 2014; Ajello et al. 2015; Fornasa & S\u00e1nchez-Conde 2015). Blazars are a special type of extragalactic sources showing extreme observational properties, such as rapid and large amplitude variability, superluminal motion, and strong non-thermal emission across the entire electromagnetic spectrum. These sources are thought to be active galactic nuclei (AGN) that host a jet pointing almost directly to the observer, within which relativistic particles move and radiate by losing their energy in a magnetic field (Blandford & Rees 1978; Urry & Padovani 1995). Among extragalactic objects blazars are known to accelerate particles to the highest observed energies and therefore are considered as prime candidates for multimessenger astronomy. Padovani & Resconi (2014), on the basis of a joint positional and energetic diagnostic, have suggested a possible association between blazars (BL Lacs) and some neutrino events reported by the IceCube collaboration (IceCube Collaboration 2014). It has also been suggested that blazars could be sites where ultra-high energy cosmic rays are generated (e.g. Zhang, Zhao & Cao 2014), although a firm association has so far been elusive despite the continuous improvement of the available data (P. Auger Collaboration 2014).","Citation Text":["Fornasa & S\u00e1nchez-Conde 2015"],"Functions Text":["The latest Fermi Large Area Telescope (Fermi-LAT) survey data are now providing strong evidence towards an origin mostly due to integrated radiation from blazars and radio galaxies (e.g."],"Functions Label":["Background"],"Citation Start End":[[1096,1124]],"Functions Start End":[[877,1063]]}
{"Identifier":"2017MNRAS.468.2590S__Fassnacht_et_al._2002_Instance_2","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"],"Functions Text":["The unprecedented quality of the light curves combined with new curve-shifting algorithms","lead to time delays with typically \u223c3 per cent accuracy"],"Functions Label":["Uses","Uses"],"Citation Start End":[[996,1017]],"Functions Start End":[[817,906],[939,994]]}
{"Identifier":"2017AandA...599A..13Y__Wimmer-Schweingruber_et_al._1997_Instance_1","Paragraph":"Figure 2 shows solar wind plasma and magnetic field measurements for a CIR that occurred between July 26 and 27, 2003 (days of year 207\u2013208). Following Chotoo et al. (2000), Richardson et al. (1993), we marked four regions in the plot: the slow wind region (S), the compressed slow wind region (S\u2032), the compressed fast wind region (F\u2032), and the fast wind itself (F). Throughout four regions, the mean charge states of iron measured by ACE\/SWICS lies around 11+, consistent with typical values in the solar wind (Lepri et al. 2001). The stream interface (S\u2032-F\u2032) is indicated by the vertical line in Fig. 2 and is characterized by a drop of the O7+\/O6+ abundance ratio measured with SWICS in the bulk solar wind (Wimmer-Schweingruber et al. 1997, 1999). The leading (S-S\u2032) and trailing edge (F\u2032-F) of the CIR were determined by the total pressure (Jian et al. 2006). Bu\u010d\u00edk et al. (2009) found that CIR boundaries can be well defined when the total pressure exceeds 50 pPa (indicated by the horizontal dashed line in Fig. 2), which is slightly higher than that in the background solar wind, which typically is 20\u221230 pPa, according to Jian et al. (2006). The total pressure P was obtained from the sum of plasma and magnetic field pressure, that is, \\hbox{$P=n_{\\rm p}v^{2}_{\\rm th}m+B^2\/2\\mu_0$}P=npvth2m+B2\/2\u03bc0, where np and vth are the proton density and thermal speed, respectively, and B is the magnitude of the magnetic field. Because SOHO has no magnetometer, we used magnetic field data from ACE\/MAG (which is also around L1). Comparing plasma parameters (bulk speed, thermal speed, and proton density) measured by PM with those of SWEPAM, we see that the physical conditions at SOHO and ACE were almost the same, and that the time difference between passages of the CIR boundaries is less than ten minutes. The CIR shown in Fig. 2 was bounded by a reverse shock (vertical line separating F\u2032 from F). We clearly see that the suprathermal He++ intensity peaks inside the decelerated and compressed fast-wind region (F\u2032), close to the reverse shock. In contrast, suprathermal particles are very rare in the S and S\u2032 regions. After passage of the reverse shock, suprathermal particles continue to be observable for more than one day. They are commonly believed to be the sunward particles accelerated by the reverse shock far beyond the Earth orbit. In other words, the observer saw the duration of the CIR particle event, which was longer than that of the CIR compression region itself. The background level shown in green was estimated using the method described above. The signal-to-noise ratio (S\/N) in the F and F\u2032 regions is higher than 100, confirming that our observations are due to real He++ particles. ","Citation Text":["Wimmer-Schweingruber et al. 1997"],"Functions Text":["The stream interface (S\u2032-F\u2032) is indicated by the vertical line in Fig. 2 and is characterized by a drop of the O7+\/O6+ abundance ratio measured with SWICS in the bulk solar wind"],"Functions Label":["Uses"],"Citation Start End":[[712,744]],"Functions Start End":[[533,710]]}
{"Identifier":"2019AandA...631A..88Y__Bohren_&_Huffman_(1998)_Instance_3","Paragraph":"Starting from the four aforementioned materials, we consider several composition mixtures and grain structures. For the sake of comparison, we first consider compact grains of purely a-Sil, a-C, or a-C:H. Subsequently, according to K\u00f6hler et al. (2015), we consider compact grains made of two thirds a-Sil and one third a-C (Mix 1) or one third a-C:H (Mix 2), in terms of volume fractions. These allow reproduction of the mass fractions derived by Jones et al. (2013) for the diffuse ISM. The effect of porosity is tested for the Mix 1 mixture, with a porosity degree of 50% (Mix 1:50). We also evaluate theimpact of the presence of a water ice mantle on compact Mix 1 grains (Mix 1:ice). We further consider two material compositions defined in Pollack et al. (1994) based on depletion measurements: (i) 21% a-Sil and 79% a-C (Mix 3); and (ii) 8% a-Sil, 30% a-C, and 62% water ice (Mix 3:ice). The various grain compositions are summarised in Table 1. For each grain composition, we derive the absorption and scattering efficiencies Qabs and Qsca, respectively, and the asymmetry factor of the phase function g = \u27e8cos\u03b8\u27e9. To allow fast calculations, we make the major assumption that the grains are spherical and compute their optical properties using the Mie theory (Mie 1908; Bohren & Huffman 1983) with the Fortran 90 version of the BHMIE routine given in Bohren & Huffman (1998). For grains consisting of two or three materials, we first derive effective optical constants following the Maxwell Garnett mixing rule (Maxwell Garnett 1904; Bohren & Huffman 1998). Indeed, we assume that in Mix 1 grains, for example, carbon appears as proper inclusions in the silicate matrix rather than assuming a completely random inhomogeneous medium. Mishchenko et al. (2016a,b) performed exhaustive studies of the applicability of the Maxwell Garnett mixing rule to heterogeneous particles. These latter authors showed that this rule can provide accurate estimates of the scattering matrix and absorption cross-section of heterogeneous grains at short wavelengths (typically up to the visible for a 0.1 \u03bcm grain and to the mid-infrared(MIR) for a 10 \u03bcm grain) if twocriteria are met: both the size parameter of the inclusions and the refractive index contrast between the host material and the inclusions have to be small. Moreover, Mishchenko et al. (2016a) demonstrated that the extinction and asymmetry-parameter errors of the Maxwell Garnett mixing rule are significantly smaller than the scattering-matrix errors, remaining small enough for most typical applications and in particular the kind of applications we perform here. It is however well known that this kind of mixing rule systematically underestimates the absorption efficiency in the FIR to millimetre wavelength range, the implications of which are discussed in Sect. 3.2. We perform our computations with the emc routine of V. Ossenkopf3. For Mix 1 and Mix 2, we assume a matrix of a-Sil with inclusions of a-C or a-C:H, and for Mix 3 a matrix of a-C with inclusions of a-Sil. For grains surrounded by an ice mantle, the optical properties are derived with the core-mantle Mie theory using the BHCOAT routine given in Bohren & Huffman (1998).","Citation Text":["Bohren & Huffman (1998)"],"Functions Text":["For grains surrounded by an ice mantle, the optical properties are derived with the core-mantle Mie theory using the BHCOAT routine given in"],"Functions Label":["Uses"],"Citation Start End":[[3177,3200]],"Functions Start End":[[3036,3176]]}
{"Identifier":"2021MNRAS.504.3316B__than_2000_Instance_2","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)"],"Functions Text":["The more contentious issue is that of the phase curve\u2019s phase offset and nightside temperature.","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."],"Functions Label":["Differences","Differences"],"Citation Start End":[[392,415]],"Functions Start End":[[296,391],[416,611]]}
{"Identifier":"2022MNRAS.515.3299G__Tagawa_et_al._2020a_Instance_1","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"],"Functions Text":["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","then get captured in the disc by hydrodynamic drag as they cross the disc (e.g."],"Functions Label":["Background","Background"],"Citation Start End":[[759,778]],"Functions Start End":[[281,462],[607,686]]}
{"Identifier":"2019ApJ...873...89M__Bourrier_et_al._2016_Instance_1","Paragraph":"Ultraviolet observations of the hot Jupiter HD 209458 b have found up to a 15% occultation in the wings of hydrogen Ly\u03b1, with effective Doppler shifts of up to \u00b1150 km s\u22121, significantly larger than the 5% occultation observed at optical (Vidal-Madjar et al. 2003, 2008; Ben-Jaffel 2007; Ehrenreich et al. 2008). The high occultation and large velocity are indicative of a fast and extended component of the atmosphere, interpreted as an escaping planetary wind. Similar outflows have been reported for another hot and one warm Jupiter, HD 189733 b (Lecavelier Des Etangs et al. 2010; Bourrier et al. 2013) and 55 Cnc b (Ehrenreich et al. 2012\n)5\n\n5\nAlong with a nondetection for a super Earth, 55 Cnc e, placing an upper limit on its mass loss.\n, and for one hot Neptunian planet, GJ 435 b (Kulow et al. 2014; Ehrenreich et al. 2015; Bourrier et al. 2016; Lavie et al. 2017). Interestingly, the outflow from GJ 435 b is asymmetric both temporally and spectrally,6\n\n6\nRedshifted occultation of (0.7 \u00b1 3.6)% pretransit and (8.0 \u00b1 3.1)% posttransit. Blueshifted occultation of (17.6 \u00b1 5.2)% pretransit and (47.2 \u00b1 4.1)% posttransit (Ehrenreich et al. 2015).\n suggesting a cometary tail-like outflow moving rapidly away from the star. Tentative detections indicate that metals may be present in these escaping winds, including oxygen (Vidal-Madjar et al. 2004; Ben-Jaffel & Sona Hosseini 2010), magnesium (Vidal-Madjar et al. 2013), and carbon and silicon (Linsky et al. 2010; Loyd et al. 2017). Additionally, hydrogen H\u03b1 absorption has been seen in HD 189733 b\u2019s transmission spectra (Jensen et al. 2012), but its relation to hydrodynamic escape is still uncertain (Barnes et al. 2016). Recently, the outflow from Wasp-107 b was detected in the 1083 nm line of excited neutral helium (Spake et al. 2018). This line, predicted for exoplanet atmospheres by Seager & Sasselov (2000), and in their outflows by Oklop\u010di\u0107 & Hirata (2018), provides an opportunity for ground-based observations.","Citation Text":["Bourrier et al. 2016"],"Functions Text":["Similar outflows have been reported for","and for one hot Neptunian planet, GJ 435 b"],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[835,855]],"Functions Start End":[[463,502],[748,790]]}
{"Identifier":"2016MNRAS.463.3637R__Mirabel,_Dottori_&_Lutz_1992_Instance_1","Paragraph":"In MS, the tidal tails host the formation of overdensities similar to TDGs (about 3 in each tail are visible in Fig. 3). Their associated phantom DM further favours their growth, while no such structure is visible in the stellar nor gaseous component of the NA model. The formation of TDGs in simulations is sensitive to the resolution (Wetzstein, Naab & Burkert 2007) and the truncation of the DM halo (Bournaud et al. 2003). By conducting our comparisons at the same resolution in Newton and MOND, and by testing the formation of substructures with much more extended haloes (Appendix C), we ensure that the differences we detect have a physical origin. Observations of the Antennae galaxies report only one TDG candidate, at the tip of the southern tail (Mirabel, Dottori & Lutz 1992), but the exact nature of this structure is still questioned (Hibbard et al. 2001, see also Bournaud et al. 2004). It could be either an unbound object or a forming TDG still out of equilibrium. For the models we consider, the Newtonian framework does not allow for the formation of TDGs, while the MOND does. However we note that, in the absence of efficient shielding from the rest of the galaxy, star-forming regions in the tidal tails are more sensitive to ultraviolet radiation of extragalactic origin. A stronger radiation, e.g. at higher redshift or in a denser galactic environment, could potentially prevent the formation of TDGs and thus reduce the differences (in the young stellar component) between Newtonian and Milgromian cases in this context. Our simulations show however that the old stellar component is likely to remain more clumpy in MOND, as long as the gaseous contribution to the local gravitational potential is negligible over the stellar one. Leaving this issue aside, with the specific models considered here, the Milgromian runs tend to slightly overproduce TDGs given the absence of unambiguously defined ones in the Antennae, whilst the Newtonian runs might potentially slightly underproduce them if the observed TDG candidate turns out to be real. However, the uniqueness of our initial conditions has not been established and it is possible that other sets of parameters could reproduce the same morphology with a different number of TDGs. A much larger simulation sample including more interacting systems would be necessary to reach a clear conclusion on this particular topic.","Citation Text":["Mirabel, Dottori & Lutz 1992"],"Functions Text":["Observations of the Antennae galaxies report only one TDG candidate, at the tip of the southern tail","but the exact nature of this structure is still questioned"],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[758,786]],"Functions Start End":[[656,756],[789,847]]}
{"Identifier":"2015ApJ...805..115V__Robitaille_et_al._2006_Instance_1","Paragraph":"To classify the bursts in our models, we use the remaining mass in the envelope to define the boundary between the embedded and optically visible phases. Namely, we assume that the optically visible Class II begins when less than 10% of the initial core mass is left in the envelope. The boundary between the deeply embedded Class 0 phase and the partly embedded Class I phase is defined as the time when 50% of the initial core mass is left in the envelope. Our adopted classification scheme is based on physical properties of a young stellar object, such as envelope and disk masses (e.g., Robitaille et al. 2006; Dunham et al. 2010, 2014), rather than on observational signatures, such as submillimeter luminosities or effective temperatures (e.g., Andr\u00e9 et al. 1993; Chen et al. 1995). Classifications relying upon physical properties are usually referred in the literature as \u201cstages\u201d, whereas those using observational signatures are called \u201cclasses\u201d. For simplicity here we use the term \u201cclass\u201d to refer to both the physical stages and observational classes. Our adopted definition of physical stages was extensively investigated in Dunham et al. (2010). They found that there is not always a one-to-one correspondence between physical stage defined by the envelope mass and observational class defined by the submillimeter luminosity or effective temperature due to the effects of geometry and extinction. In reality the exact point at which to set the class boundaries is somewhat uncertain, which could shift the duration of the embedded phase in our models by a factor of order unity in either direction. We disentangle the disk and infalling envelope on our numerical grid using the algorithm described in Vorobyov (2011), which is based on the disk-to-envelope transition density of \n\n\n\n\n\n g cm\u22122 and the velocity field in the infalling envelope. Varying the value of \n\n\n\n\n\n by a factor of 5 results in changes of the estimated onset time of different phases by only a few per cent.","Citation Text":["Robitaille et al. 2006"],"Functions Text":["Our adopted classification scheme is based on physical properties of a young stellar object, such as envelope and disk masses (e.g.","rather than on observational signatures, such as submillimeter luminosities or effective temperatures"],"Functions Label":["Uses","Differences"],"Citation Start End":[[592,614]],"Functions Start End":[[459,590],[643,744]]}
{"Identifier":"2017MNRAS.464..183N__Iannuzzi_&_Dolag_2012_Instance_1","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":["Iannuzzi & Dolag 2012"],"Functions Text":["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."],"Functions Label":["Uses"],"Citation Start End":[[854,875]],"Functions Start End":[[609,829]]}
{"Identifier":"2015ApJ...807...91C__Moffat_1969_Instance_1","Paragraph":"The standard photometric analysis (see Dalessandro et al. 2008a, 2008b) has been performed on the \u201cflc\u201d images, which are corrected for flat field, bias, dark counts, and charge transfer efficiency. These images have been further corrected for \u201cPixel-Area-Map\u201d10\n\n10\nFor more details see the ACS Data Handbook.\n with standard IRAF procedures. By using the DAOPHOT II ALLSTAR and ALLFRAME packages (Stetson 1987), we performed an accurate photometric analysis of each image. First of all, we modeled the point-spread function (PSF) by using a sample of \u223c200 bright but not saturated stars. The model has been chosen on the basis of a \u03c72 test and, in every image, the best fit is provided by a Moffat function (Moffat 1969). Then we performed a source detection analysis, setting a 3\u03c3 detection limit, where \u03c3 is the standard deviation of the measured background. Once a list of stars was obtained, we performed a PSF-fitting in each image. In the resulting catalog we included only stars present at least in half the images for each filter. For each star, we homogenized the magnitudes estimated in different images, and their weighted mean and standard deviation have been finally adopted as the star mean magnitude and its related photometric error (see Ferraro et al. 1991, 1992). However, in order to perform variability studies, for each source we also kept the homogenized magnitude measured in each frame in both filters. Then, instrumental magnitudes have been calibrated to the VEGAMAG system, cross-correlating11\n\n11\nWe used CataXcorr, a code that is specifically developed to perform accurate astrometric solutions. It has been developed by P. Montegriffo at INAF\u2013Osservatorio Astronomico di Bologna. This package is available at http:\/\/davide2.bo.astro.it\/paolo\/Main\/CataPack.html, and has been successfully used in a large number of papers by our group in the past ten years.\n our catalog with that by Anderson et al. (2008), using the \u223c7600 stars in common.","Citation Text":["Moffat 1969"],"Functions Text":["The model has been chosen on the basis of a \u03c72 test and, in every image, the best fit is provided by a Moffat function"],"Functions Label":["Uses"],"Citation Start End":[[709,720]],"Functions Start End":[[589,707]]}
{"Identifier":"2018ApJ...864...51W__Lee_et_al._2014_Instance_1","Paragraph":"Understanding how, when, and where galaxies quench their star formation is one of the most muddled, outstanding problems in galaxy formation. Different models for galaxy formation and evolution often make very different assumptions and\/or predictions. Among them, semi-analytic galaxy formation models (SAMs; e.g., White & Frenk 1991; Kang et al. 2005; Bower et al. 2006; Croton et al. 2006; Bower et al. 2008; Somerville et al. 2008; Parry et al. 2009; Guo et al. 2011; Lu et al. 2011, 2014; Gonzalez-Perez et al. 2014; Lee et al. 2014; Henriques et al. 2015; Ruiz et al. 2015; Somerville et al. 2015) and hydrodynamical simulations (e.g., Katz et al. 1992; Springel et al. 2001, 2005; Springel & Hernquist 2003; Kere\u0161 et al. 2009; Angulo et al. 2012; Vogelsberger et al. 2014a, 2014b; Schaye et al. 2015) are two powerful tools to trace galaxy formation and evolution in cosmological volumes. SAMs are phenomenological models that use an approximate\/empirical formula to describe all baryonic processes relevant to galaxy formation, such as gas accretion, heating and cooling, star formation, feedback from stars and AGN, mergers, and stripping due to tides and ram pressure. Because of their different locations within host halos, centrals and satellites in SAMs are assumed to undergo different quenching processes (see, e.g., Henriques et al. 2017). For example, the stripping of hot and cold gas associated with galaxies is assumed to only act on satellites. In addition, the efficiency of radio AGN feedback is assumed to depend on the associated hot gas mass, which may be very different between centrals and satellites. Hydrodynamical simulations, on the other hand, evolve the dark matter and baryonic components in a self-consistent way, though some assumptions have also to be adopted to model subgrid physics (see, e.g., Springel & Hernquist 2003). More importantly, centrals and satellites are not treated differently a priori, and their differences, if any, result directly from a complicated interaction between the baryonic content of a galaxy and its environment, while the (subgrid) modeling of star formation and feedback processes carry no knowledge of this environment.","Citation Text":["Lee et al. 2014"],"Functions Text":["Different models for galaxy formation and evolution often make very different assumptions and\/or predictions. Among them, semi-analytic galaxy formation models (SAMs; e.g.,","are two powerful tools to trace galaxy formation and evolution in cosmological volumes. SAMs are phenomenological models that use an approximate\/empirical formula to describe all baryonic processes relevant to galaxy formation, such as gas accretion, heating and cooling, star formation, feedback from stars and AGN, mergers, and stripping due to tides and ram pressure."],"Functions Label":["Background","Background"],"Citation Start End":[[521,536]],"Functions Start End":[[142,314],[807,1177]]}
{"Identifier":"2020MNRAS.493.2373D__Haghi_&_Amiri_2016_Instance_1","Paragraph":"The lack of detection of supersymmetric particles at Large Hadron Collider led alternative candidates for dark matter to rapidly gaining attention. Apart from other particles, modifications of the gravitational action also offer a new avenue to account for the dark matter. Among many theories, MOG has been successfully tested with galaxies and cluster of galaxies showing the capability to explain the rotation curves of spiral galaxies, the gravitational lensing, the Sunyaev\u2013Zeldovich effect, and the X-ray emission in galaxy clusters without resorting to any dark matter component (Brownstein & Moffat 2006; Moffat & Toth 2013; Moffat & Rahvar 2013, 2014; Moffat 2016; Banerjee et al. 2017; De Martino & De Laurentis 2017; Lopez Armengol & Romero 2017). Nevertheless, a recent study of the dynamics of dwarf spheroidals orbiting around the Milky Way has shown some inconsistencies with previous results (Haghi & Amiri 2016). Since dwarfs are supposed to be dominated by dark matter, any theory that aims to replace it with a modification of the gravitational potential must also be able to account for the dynamics of these galaxies. Here, we have investigated the dynamics of stars in Antlia II in the framework of MOG theory. Antlia II is a recently discovered low-surface-brightness dwarf galaxy being \u223c100 more diffuse than ultra diffuse galaxies. It has a very wide core \u223c3 kpc, and it is supposed to be strongly dominated by dark matter (Broadhurst et al. 2019; Torrealba et al. 2019). Due to its own nature, it represents an ideal candidate to test alternative theories of gravity such as MOG. Therefore, we have predicted the dispersion velocity profile by solving the spherically symmetric Jeans equation, and assuming that the dynamics of stars is determined by the modified potential well in equation (2) arising in MOG weak field limit, and the stellar density profile is well described by the Plummer model in equation (5). Finally, we have projected the solution of Jeans equation along the line of sight to compare it with the data.","Citation Text":["Haghi & Amiri 2016"],"Functions Text":["Nevertheless, a recent study of the dynamics of dwarf spheroidals orbiting around the Milky Way has shown some inconsistencies with previous results"],"Functions Label":["Differences"],"Citation Start End":[[909,927]],"Functions Start End":[[759,907]]}
{"Identifier":"2015MNRAS.451.2544P__Mesinger,_Furlanetto_&_Cen_2011_Instance_1","Paragraph":"Observations are now probing galaxies in the middle of the reionization epoch, when the gas in the intergalactic medium was transformed from its initially neutral state into a hot, ionized plasma (e.g. McLure et al. 2011; Finkelstein et al. 2012; Bouwens et al. 2014). Most likely stars in galaxies are responsible for this transformation, although this heavily depends on the fraction of ionizing photons produced by the stars that make it into the intergalactic medium, the so-called escape fraction fesc. The escape fraction is a key parameter in studies of the contribution of the observed galaxy population to reionization (e.g. Bouwens et al. 2012; Robertson et al. 2013), semi-analytic modelling of reionization (e.g. Choudhury, Haehnelt & Regan 2009; Pritchard, Loeb & Wyithe 2010; Santos et al. 2010; Mesinger, Furlanetto & Cen 2011; Raskutti et al. 2012; Mitra, Ferrara & Choudhury 2013; Shull et al. 2012) and numerical simulations of reionization (e.g. Iliev et al. 2006; Trac & Cen 2007; Ciardi et al. 2012). A large effort is going into determining the escape fraction observationally. Except for two objects (Leitet et al. 2011, 2013), in the local Universe no ionizing radiation has been detected directly (Leitherer et al. 1995; Deharveng et al. 2001), although some objects show indirect evidence of photon leakage (Heckman et al. 2011; Zastrow et al. 2011). The lack of detections may be partly due to selection bias (Bergvall et al. 2013), but the objects from which radiation is detected have very low escape fractions, fesc  4\u2009per\u2009cent. At z \u223c 1, no objects with leaking ionizing photons have been detected (Bridge et al. 2010; Siana et al. 2010), but at z \u223c 3, the highest redshift at which the opacity of the intergalactic medium for ionizing photons is approximately less than unity, ionizing photons have been detected in \u223c10\u2009per\u2009cent of the observed objects (Nestor et al. 2013). Attempts to constrain the escape fraction with numerical simulations find ranges between fesc  10\u2009per\u2009cent (Razoumov & Sommer-Larsen 2006, 2007; Gnedin, Kravtsov & Chen 2008; Paardekooper et al. 2011; Kim et al. 2013) and fesc > 80\u2009per\u2009cent (Wise & Cen 2009; Razoumov & Sommer-Larsen 2010), with likely a strong mass and redshift dependence (Yajima, Choi & Nagamine 2010; Wise et al. 2014). Due to the opacity of the intergalactic medium, we need to mostly rely on numerical simulations to learn about the escape fraction during the epoch of reionization.","Citation Text":["Mesinger, Furlanetto & Cen 2011"],"Functions Text":["The escape fraction is a key parameter in studies of the contribution of the observed galaxy population to reionization","semi-analytic modelling of reionization (e.g."],"Functions Label":["Background","Background"],"Citation Start End":[[810,841]],"Functions Start End":[[508,627],[679,724]]}
{"Identifier":"2021MNRAS.500.1817L__Abbott_et_al._2020b_Instance_1","Paragraph":"Since the errors of the LIGO-estimated rates are dominated by Poisson statistics (Abbott et al. 2020a,b), we approximate the PDF for the expected number of detections $\\mathcal {N}=\\mathcal {R}VT$ (from the surveyed space\u2013time volume VT) by $\\mathrm{d}P\/\\mathrm{d}\\mathcal {N}\\propto \\mathcal {N}^{k-1\/2}\\mathrm{e}^{-\\mathcal {N}}\/k!$, where k = 1 for each of the relevant cases ($\\mathcal {R}_{190814}$, $\\mathcal {R}_{170817}$, and $\\mathcal {R}_{190425}$), and the factor of $\\mathcal {N}^{-1\/2}$ is from Jeffrey\u2019s prior (Abbott et al. 2020a). From the median values of $\\bar{\\mathcal {R}}_{190814}=7\\rm \\, Gpc^{-3}\\, yr^{-1}$ (Abbott et al. 2020a), $\\bar{\\mathcal {R}}_{\\rm 170817}=760\\rm \\, Gpc^{-3}\\, yr^{-1}$, and $\\bar{\\mathcal {R}}_{\\rm 190425}=460\\rm \\, Gpc^{-3}\\, yr^{-1}$ (Abbott et al. 2020b), we obtain the effective surveyed space\u2013time volumes $VT=1.2\/\\bar{\\mathcal {R}}$ for each of these three events (\u20181.2\u2019 is the median of $\\mathrm{d}P\/\\mathrm{d}\\mathcal {N}$). We consider both GW170817 and GW190425 as bNS mergers, because the component masses of GW190425 are not far from those of GW170817 and the nature of the merging objects makes little practical difference in our model. Thus, the PDF of the total bNS merger rate from the sum of the two is given by a convolution of the two individual PDFs\n(1)$$\\begin{eqnarray*}\r\n{\\mathrm{d}P\\over \\mathrm{d}\\mathcal {R}_{\\rm bns}} = \\int _0^{\\mathcal {R}_{\\rm bns}} \\mathrm{d}\\mathcal {R}_1 {\\mathrm{d}P\\over \\mathrm{d}\\mathcal {R}_1} \\left.{\\mathrm{d}P\\over \\mathrm{d}\\mathcal {R}_2}\\right|_{\\mathcal {R}_{\\rm bns}-\\mathcal {R}_1},\r\n\\end{eqnarray*}$$where we have written $\\mathcal {R}_{1} = \\mathcal {R}_{170817}$, $\\mathcal {R}_{2} = \\mathcal {R}_{190425}$ for brevity. We then calculate the PDF for the inverse of the total bNS merger rate $\\mathrm{d}P\/\\mathrm{d}\\mathcal {R}_{\\rm bns}^{-1}=\\mathcal {R}_{\\rm bns}^2\\mathrm{d}P\/\\mathrm{d}\\mathcal {R}_{\\rm bns}$. Finally, the PDF of the rate ratio $\\beta =\\mathcal {R}_{190814}\/\\mathcal {R}_{\\rm bns}$ is given by\n(2)$$\\begin{eqnarray*}\r\n{\\mathrm{d}P\\over \\mathrm{d}\\beta } = \\int _0^\\infty {\\mathrm{d}\\mathcal {R}_{3}\\over \\mathcal {R}_3} {\\mathrm{d}P\\over \\mathrm{d}\\mathcal {R}_{3}} \\left.{\\mathrm{d}P\\over \\mathrm{d}\\mathcal {R}_{\\rm bns}^{-1}}\\right|_{\\beta \/\\mathcal {R}_3},\r\n\\end{eqnarray*}$$where we have written $\\mathcal {R}_{3} = \\mathcal {R}_{190814}$ for brevity. We find the 90 per cent confidence interval for the rate ratio to be in the range $0.064\\, \\rm {per\\, cent}\\lt \\beta \\lt 2.8\\, \\rm {per\\, cent}$.","Citation Text":["Abbott et al. 2020","b"],"Functions Text":["Since the errors of the LIGO-estimated rates are dominated by Poisson statistics","we approximate the PDF for the expected number of detections $\\mathcal {N}=\\mathcal {R}VT$ (from the surveyed space\u2013time volume VT) by $\\mathrm{d}P\/\\mathrm{d}\\mathcal {N}\\propto \\mathcal {N}^{k-1\/2}\\mathrm{e}^{-\\mathcal {N}}\/k!$, where k = 1 for each of the relevant cases ($\\mathcal {R}_{190814}$, $\\mathcal {R}_{170817}$, and $\\mathcal {R}_{190425}$), and the factor of $\\mathcal {N}^{-1\/2}$ is from Jeffrey\u2019s prior"],"Functions Label":["Uses","Uses"],"Citation Start End":[[82,100],[102,103]],"Functions Start End":[[0,80],[106,523]]}
{"Identifier":"2016AandA...591A..91P__Halpern_et_al._2014_Instance_2","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"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[991,1010]],"Functions Start End":[[877,989]]}
{"Identifier":"2015ApJ...804..130C__Bertschinger_1985_Instance_3","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"],"Functions Text":["The initial perturbation can be compensated or uncompensated, which leads to different void growth scenarios"],"Functions Label":["Background"],"Citation Start End":[[2188,2205]],"Functions Start End":[[2078,2186]]}
{"Identifier":"2015ApJ...807..101K__Sui_et_al._2005_Instance_1","Paragraph":"In a collisional thick-target model, the HXRs are produced by collisional bremsstrahlung during the passage of non-thermal electrons through denser plasma regions, in which the electrons are stopped completely by Coulomb collisions (Brown 1971; Brown et al. 2009; Kontar et al. 2011). Using the electron distribution parameters derived from RHESSI spectroscopy, the power delivered by non-thermal electrons above low-energy cutoff (ELC) can be calculated by the expression\n2\n\n\n\n\n\nwhere ELC is the low-energy cutoff, Fe is the total number of electrons per second above ELC in units of 1035 electrons s\u22121, and \u03b4 is the electron spectral index (Fletcher et al. 2013). An accurate determination of the low-energy cutoff to non-thermal electron distributions is crucial for the calculation of power and consequently non-thermal energy in solar flares (Sui et al. 2005; Veronig et al. 2005). In general flares are thought to have low-energy cutoffs close to or in the region where the emission is dominated by thermal bremsstrahlung (Ireland et al. 2013). We further emphasize that in flares with multiple HXR sub-peaks during the impulsive phase (like the present one), the determination of ELC is rather illusive during the peak emission as it is difficult to distinguish the signals of flare accelerated electrons against dominant thermal bremsstrahlung. This often results in the overestimation of ELC during HXR peak phases. For this event, we find that ELC varies in the range of 20\u201332 keV with relatively higher values (\u223c32 keV) during the second peak (see Figure 13(d)). Therefore, in order to get an estimation of power (and consequently energy) of the non-thermal electrons, we have taken ELC as 25 keV, which is the average of ELC over the flare time interval. The estimated non-thermal electron flux (Fe) above 25 keV is plotted in Figure 14(b). In Figure 14(c), we have plotted the power (Pnth) and energy (Enth) contained in non-thermal electron beams. The plots of Fe and Pnth clearly indicate noticeable enhancement in the particle rate and consequently the power of flare-accelerated electrons during the two HXR peaks.","Citation Text":["Sui et al. 2005"],"Functions Text":["An accurate determination of the low-energy cutoff to non-thermal electron distributions is crucial for the calculation of power and consequently non-thermal energy in solar flares"],"Functions Label":["Motivation"],"Citation Start End":[[848,863]],"Functions Start End":[[666,846]]}
{"Identifier":"2015AandA...583A..77L__Busso_et_al._2001_Instance_1","Paragraph":"During the past decade, significant information has been gathered on the chemical compositions of post-asymptotic giant branch (AGB) stars in the Milky Way. This has led to the discovery of a class of post-AGB stars that have C\/O > 1 and display extreme enrichment in the abundances of the elements heavier than Fe produced by slow neutron captures (the s-process, Van Winckel & Reyniers 2000; Reyniers & Van Winckel 2003; Reyniers et al. 2004). Since AGB stars can become C rich and have been confirmed both theoretically and observationally as the main stellar site for the s-process (see, e.g., Busso et al. 2001), it is natural to interpret these post-AGB observations as the signature of the nucleosynthesis and mixing events that occurred during the preceeding AGB phase. These events are currently identified as (i) the mixing of protons into the radiative He-rich intershell leading to the formation of a thin region rich in the main neutron source 13C (the 13C \u201cpocket\u201d); (ii) proton-ingestion episodes (PIEs) inside the convective thermal pulses (TPs); and (iii) the third dredge-up (TDU), which carries C and s-process elements from the He-rich intershell to the convective envelope and to the stellar surface. Since the details of all these processes are very uncertain (see discussion in, e.g., Busso et al. 1999; Herwig 2005; Campbell & Lattanzio 2008), observations of post-AGB stars provide strong constraints. Recent observations of the chemical composition of four low-metallicity ([Fe\/H] from \u22121.15 to \u22121.34), s-process-rich, C-rich post-AGB stars in the Large (J050632, J052043, and J053250) and Small (J004441) Magellanic Clouds (LMC and SMC, respectively) have provided a challenge to AGB s-process models (De Smedt et al. 2012, 2014; van Aarle et al. 2013)1. Since we know the distance of these stars, it is possible to determine from the observed luminosity that their initial stellar mass was in the range 1\u20131.5 M\u2299. Stellar AGB models in this range of mass and metallicity can produce the high observed abundances of the s-process elements, such as Zr and La (1  [Zr\/Fe]  2 and 1  [La\/Fe]  3), together with [Pb\/La] \u2243 1, if a deep TDU is assumed after a last TP. Instead, negative [Pb\/La] values are observed as upper limits (De Smedt et al. 2014). Here we test different possible modifications of the current AGB s-process scenario to explain the neutron-capture abundance pattern observed in the MC post-AGB stars. ","Citation Text":["Busso et al. 2001"],"Functions Text":["Since AGB stars can become C rich and have been confirmed both theoretically and observationally as the main stellar site for the s-process (see, e.g.,","it is natural to interpret these post-AGB observations as the signature of the nucleosynthesis and mixing events that occurred during the preceeding AGB phase."],"Functions Label":["Background","Background"],"Citation Start End":[[598,615]],"Functions Start End":[[446,597],[618,777]]}
{"Identifier":"2020AandA...643A.148G__Martin_et_al._1998_Instance_1","Paragraph":"Figure 11 shows that the distribution of the spectral index \u03b1 (see Table A.1) is bimodal suggesting the existence of two populations (disc-bearing and discless stars) with approximately the same number of sources. We verified that the bimodality of the distribution of spectral indices is not an artefact caused by the different number of points and wavelength range of the photometric data available for each star to compute the spectral indices. A similar result was also observed for the Chamaeleon I star-forming region (see Fig. 11 of Luhman et al. 2008). Previous studies suggested that the early disappearence of circumstellar discs could be related to environmental effects imposed by the presence of massive stars that produce strong UV radiation, stellar winds, and supernova explosions (see e.g. Walter et al. 1994; Martin et al. 1998). However, we see no dependency of the spectral index on the position of the stars in our sample and any nearby OB star surrounding the Lupus clouds to support this scenario. As shown in Fig. 10 we note the existence of only a few stars older than 10 Myr in our sample which could be potential contaminants (as expected based on the performance of our classifier, see Table 1). Thus, the hypothesis of contamination by older field stars does not explain the bimodality of spectral indices in our sample given that the two populations have the same number of stars, similar ages and are younger than the potential contaminants from the Sco-Cen association. Alternatively, we investigated the dependency of the spectral indices on the age and colour of the stars. Figure 12 shows that the age distribution of the Class II and Class III stars in Lupus overlap. The median age of Class III stars inferred from the Baraffe et al. (2015) and Siess et al. (2000) stellar models is about 3 Myr which yields a rough estimate of the typical disc lifetime in the Lupus association. However, it should also be noted from Fig. 12 that we observe an excess of Class III stars at cooler temperatures (red colours) suggesting that the survival time of circumstellar discs may also depend on other stellar parameters. For example, Galli et al. (2015) used an empirical disc evolution model to determine the lifetime of circumstellar discs in Lupus in terms of the mass of the star. According to their model the average lifetime of a circumstellar disc around a star with 0.1 M\u2299 is of the order of 1 Myr which could explain the early disapperance of circumstellar discs for some stars in our sample.","Citation Text":["Martin et al. 1998"],"Functions Text":["Previous studies suggested that the early disappearence of circumstellar discs could be related to environmental effects imposed by the presence of massive stars that produce strong UV radiation, stellar winds, and supernova explosions (see e.g.","However, we see no dependency of the spectral index on the position of the stars in our sample and any nearby OB star surrounding the Lupus clouds to support this scenario."],"Functions Label":["Background","Differences"],"Citation Start End":[[827,845]],"Functions Start End":[[561,806],[848,1020]]}
{"Identifier":"2016MNRAS.457.2569M__Binney_&_Piffl_2015_Instance_1","Paragraph":"The top-down dynamical approach consists in producing ab initio simulations of Milky Way-like galaxies in a cosmological context. This approach can be useful to understand some general features of galaxy formation (e.g. Minchev et al. 2014). However, it is not flexible enough to produce an acceptable model for the wide range of extremely detailed data soon to be available for our own Galaxy. On the other side, the bottom-up approach for dynamical modelling consists in starting from the actual Galactic data, rather than from simulations, in order to construct a model of the Galaxy. To avoid the redundancy and computational waste of representing the orbits of every single particle in the model, one can use a phase-space distribution function (DF) to represent each population of constituent particles (typically, various stellar populations and dark matter; see e.g. Binney & Piffl 2015; Piffl, Penoyre & Binney 2015). The model-building generally starts from the assumptions of dynamical equilibrium and axisymmetry. These assumptions allow us to make use of Jeans\u2019 theorem constraining the DF to depend only on three integrals of motion, which can typically be chosen to be the radial, azimuthal, and vertical action variables. However, one should remember, especially when modelling the stellar populations of the Galactic disc, that the Galaxy is obviously not axisymmetric, as it harbours a central bar as well as spiral arms. Such perturbations can obviously be treated through perturbation theory, whose foundations in the case of flat 2D discs have been laid down by Kalnajs (1971). For instance, following up on the work of Binney & Lacey (1988) who derived the orbit-averaged Fokker\u2013Planck equation for a 2D stellar disc, recent investigations (e.g. Fouvry, Binney & Pichon 2015) have focused on the long-term secular evolution of such a flat disc by means of diffusion through action space at resonances, producing ridges in action space. Here, we are rather interested in the present-day perturbed DF in the action-angle space of the unperturbed Hamiltonian, in the presence of a 3D spiral arm perturber, which could be fitted to a snapshot of the Galaxy taken by current and upcoming large surveys. Our philosophy is thus closer to that of McMillan (2013), except that the shape of the perturbed DF will be computed directly from the linearized Boltzmann equation. Moreover, in this paper, we will first concentrate only on the response away from the main resonances, the extremely interesting effects expected at resonances, as well as the effect of resonance overlaps of multiple perturbers (e.g. Quillen 2003; Minchev & Famaey 2010), being the subject of further analytical work.","Citation Text":["Binney & Piffl 2015"],"Functions Text":["To avoid the redundancy and computational waste of representing the orbits of every single particle in the model, one can use a phase-space distribution function (DF) to represent each population of constituent particles (typically, various stellar populations and dark matter; see e.g."],"Functions Label":["Uses"],"Citation Start End":[[875,894]],"Functions Start End":[[588,874]]}
{"Identifier":"2019AandA...631A..35B__Bridges_et_al._(1996)_Instance_4","Paragraph":"The collision velocity dependence of the coefficient of restitution between particles was observed in experiments (Bridges et al. 1996; Higa et al. 1996) and is discussed in the literature (e.g., Ram\u00edrez et al. 1999; Zhang & Vu-Quoc 2002). However, the experiments by Hei\u00dfelmann et al. (2010), used in the present paper to support our assumption of a constant coefficient of restitution, do not see a variation of the coefficient of restitution between particles at low collision velocities (\u2264 1 cm s\u22121). This discrepancy in results might originate in the nature of the collisions studied in these different experiments: Bridges et al. (1996) and Higa et al. (1996) performed collisions of a particle with a flat surface, while Hei\u00dfelmann et al. (2010) observed particle-particle collisions in a free-floating environment. The latter is an experimental environment very similar to NanoRocks. In such inter-particle collisions in free-floating environments, other physical effects lead to a different behavior of the energy dissipation during collisions. In particular, the damping behavior of a large plate or surface is expected to differ from that of a same-sized particle, so that the velocity dependence of the coefficient of restitution might be an effect of the experimental setup in Bridges et al. (1996) and Higa et al. (1996). Colwell et al. (2016) and Brisset et al. (2018) studied collisions between a round cm-sized particle and a flat surface of fine grains. They also observed an increase of the coefficient of restitution with decreasing collision velocity. While the composition of the target surface was different than in Bridges et al. (1996) and Higa et al. (1996) (fine granular material vs. solid ice), the similar behavior of the coefficient of restitution supports the fact that particle-surface collisions are very different from particle-particle collisions, and coefficients of restitution are only velocity dependent for collisions with particles with very different sizes (a much larger particle can be approximated as a target surface).","Citation Text":["Bridges et al. (1996)"],"Functions Text":["While the composition of the target surface was different than in","and Higa et al. (1996) (fine granular material vs. solid ice), the similar behavior of the coefficient of restitution supports the fact that particle-surface collisions are very different from particle-particle collisions, and coefficients of restitution are only velocity dependent for collisions with particles with very different sizes (a much larger particle can be approximated as a target surface)."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1639,1660]],"Functions Start End":[[1573,1638],[1661,2065]]}
{"Identifier":"2022AandA...667A.131B__Izumi_et_al._(2016)_Instance_1","Paragraph":"Molecular line ratio diagnostics are often used to investigate the physics and chemistry of the ISM in all of these environments. For example, as the gas chemistry located in the central, nuclear regions of galaxies is believed to be dominated by X-rays produced by the AGN, in so-called X-ray dominated regions (XDRs), the molecular content of the ISM surrounding such nuclei will greatly differ from that in starburst regions (Usero et al. 2004; Garc\u00eda-Burillo et al. 2010). Hence, line ratios of specific molecules have been proposed as indicators of certain energetic or physical processes, for example HCN\/HCO+ as a tracer of AGNs (Loenen et al. 2007), HCN\/HNC as a mechanical heating tracer (Hacar et al. 2020), and HCN\/CO as a density tracer (Leroy et al. 2017). In particular, the \u201csubmillimeter-HCN diagram\u201d, first proposed in Izumi et al. (2013) and later expanded upon in Izumi et al. (2016), is a very notable example of the use of molecular line ratios as a probe of AGN-galaxies; this diagram makes use of two line ratios, HCN(4\u22123)\/HCO+(4\u22123) and HCN(4\u22123)\/CS(7\u22126), where all of the molecules involved are considered tracers of dense gas. Izumi et al. (2016) observed a clear trend that AGNs, including the Seyfert composite galaxy NGC 1068, tend to show higher HCN\/HCO+ and\/or HCN\/CS than in SB galaxies as long as the observations were at high enough spatial resolutions to separate energetically discrete regions. Izumi et al. (2016) propose a scenario where it is the high temperature that is responsible for the HCN enhancement, whereby neutral-neutral reactions with high reaction barriers are enhanced (Harada et al. 2010), thus leading to the possible enhancement of HCN and the depletion of HCO+ via newly available formation and destruction paths, respectively. However, while of course higher gas temperatures are expected in AGN-dominated regions, these are not unique to these environments, as starburst regions and\/or regions where outflows dominate can also harbour high enough temperatures for such enhancement to occur. Additionally, the higher temperatures could increase HCN excitation, relative to HCO+ and CS, without necessarily changing their relative abundances (Imanishi et al. 2018a). Finally, infrared radiative pumping is also a possible explanation of the HCN intensity enhancement relative to HCO+ and CS. Infrared pumping is a result of the emission of 14 \u03bcm infrared photons due to the presence of hot dust around AGN. These photons vibrationally excite HCN to the \u03bd2\u2004=\u20041 state. Upon de-exciting from this state back to the vibrational ground state, \u03bd\u2004=\u20040, the HCN line intensities are thus pumped to higher fluxes (Imanishi et al. 2018a). However, we note that it is also not unlikely that the 12 \u03bcm infrared photons can similarly vibrationally excite HCO+, thus nullifying the extent of this effect (Imanishi et al. 2016).","Citation Text":["Izumi et al. (2016)"],"Functions Text":["In particular, the \u201csubmillimeter-HCN diagram\u201d, first proposed in Izumi et al. (2013) and later expanded upon in","is a very notable example of the use of molecular line ratios as a probe of AGN-galaxies; this diagram makes use of two line ratios, HCN(4\u22123)\/HCO+(4\u22123) and HCN(4\u22123)\/CS(7\u22126), where all of the molecules involved are considered tracers of dense gas."],"Functions Label":["Background","Background"],"Citation Start End":[[883,902]],"Functions Start End":[[770,882],[904,1150]]}
{"Identifier":"2018MNRAS.475.3419A__Davis_et_al._1999_Instance_1","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"],"Functions Text":["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","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."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1531,1548]],"Functions Start End":[[1365,1529],[1551,1815]]}
{"Identifier":"2019ApJ...886...34F__Bonal_et_al._2010_Instance_1","Paragraph":"Here, we report Li\u2013Be\u2013B and Al\u2013Mg isotopic compositions of seven CAIs from Sayh al Uhaymir (SaU) 290 (CH) and one CAI from Isheyevo (CH\/CB) chondrites, which are one of the most pristine (unmetamorphosed) meteorites in our collections (e.g., Bischoff et al. 1993; Weisberg et al. 2001; Krot et al. 2002). The CH and CH\/CB (hereafter: CH\u2013CB) chondrites and their components are characterized by enrichments in 15N (up to 1,100\u2030; Murty et al. 2007; Ivanova et al. 2008; Briani et al. 2009; Bonal et al. 2010), which is similar to the characteristics of comets (F\u00fcri & Marty 2015). In addition, Van Kooten et al. (2016) and Olsen et al. (2016) proposed that Mg and Cr isotopic compositions of bulk CH\u2013CB chondrites require significant amounts (20%\u201350%) of primordial molecular cloud matter in their precursor material. Hence, CH\u2013CB chondrites may have accreted a significant amount of outer solar system materials. Most CAIs in CH\u2013CB chondrites also have a unique absence of excess 26Mg derived from the decay of the short-lived radionuclide 26Al (Kimura et al. 1993; Weber et al. 1995; Krot et al. 2008a), which is unlikely attributed to thermal metamorphism and\/or aqueous alteration in the parent bodies (Kimura et al. 1993; Krot et al. 2002; Krot et al. 2008b; Zhang & Hsu 2009). Thus, CH\u2013CB CAIs may also have distinctive information about the origin of 10Be in the solar protoplanetary disk. However, because of their small sizes, it is difficult to measure B isotopic compositions of CAIs in CM, CO, and CH\u2013CB chondrites using a conventional mass spectrometer. To overcome this problem, we have developed a protocol for high accuracy and high spatial resolution measurement techniques using a NanoSIMS 50 (Fukuda et al. 2018). Based on a newly obtained data set in this study, together with the data of previous studies, we will discuss the origin of 10Be in the early solar system and its implications for the astronomical setting of CAI formation in the solar protoplanetary disk.","Citation Text":["Bonal et al. 2010"],"Functions Text":["The CH and CH\/CB (hereafter: CH\u2013CB) chondrites and their components are characterized by enrichments in 15N (up to 1,100\u2030;"],"Functions Label":["Background"],"Citation Start End":[[488,505]],"Functions Start End":[[305,427]]}
{"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"],"Functions Text":["Once metallicity starts to increase, low- and intermediate-mass stars release carbon and massive stars start to eject more carbon than oxygen"],"Functions Label":["Background"],"Citation Start End":[[1301,1319]],"Functions Start End":[[1158,1299]]}
{"Identifier":"2019ApJ...875L..31H__Leary_et_al._2006_Instance_2","Paragraph":"The recent detection of gravitational-wave (GW) emission from a merging neutron star binary (Abbott et al. 2017d) and merging black hole binaries (BHBs; Abbott et al. 2016a, 2016b, 2017a, 2017b, 2017c; The LIGO Scientific Collaboration & The Virgo Collaboration 2018) by the Laser Interferometer Gravitational-Wave Observatory (LIGO)\/Virgo have ushered in an exciting new era of GW astrophysics. The astrophysical origin of the detected mergers is currently under debate, with numerous explanations proposed. These explanations can be very roughly divided into two main categories: mergers due to isolated binary evolution (e.g., Belczynski et al. 2016; de Mink & Mandel 2016; Mandel & de Mink 2016; Marchant et al. 2016), and mergers due to dynamical interactions (e.g., Portegies Zwart & McMillan 2000; Wen 2003; O\u2019Leary et al. 2006, 2009, 2016; Antonini & Perets 2012; Kocsis & Levin 2012; Antonini et al. 2014; Antonini & Rasio 2016; Rodriguez et al. 2016; VanLandingham et al. 2016; Askar et al. 2017; Arca-Sedda & Gualandris 2018; Fragione & Kocsis 2018; Hoang et al. 2018; Randall & Xianyu 2018; Arca-Sedda & Capuzzo-Dolcetta 2019). Orbital eccentricity has been explored as a way to distinguish between these merger channels in both the LIGO\/Virgo and Laser Interferometer Space Antenna (LISA) frequency bands. In contrast to mergers from isolated binary evolution, merging binaries from dynamical channels have been shown to have measurable eccentricities when they enter the LISA and\/or LIGO\/Virgo band, and can potentially be used as a way to distinguish between channels (e.g., O\u2019Leary et al. 2009; Cholis et al. 2016; Gond\u00e1n et al. 2018; Lower et al. 2018; Randall & Xianyu 2018; Rodriguez et al. 2018; Samsing 2018; Zevin et al. 2019). Unlike LIGO\/Virgo, which can only detect merging BHBs in the final inspiral phase before merger, LISA will be able to detect eccentric stellar-mass BHBs for long timescales before they merge in the LIGO\/Virgo band (e.g., O\u2019Leary et al. 2006; Breivik et al. 2016; Nishizawa et al. 2016; Chen & Amaro-Seoane 2017; Nishizawa et al. 2017; D\u2019Orazio & Samsing 2018; Kremer et al. 2019; Samsing & D\u2019Orazio 2018). This provides us with invaluable insight into the dynamical evolution of eccentric binaries leading up to the merger, which has important implications about the astrophysical context in which merging binaries evolve.","Citation Text":["O\u2019Leary et al. 2006"],"Functions Text":["Unlike LIGO\/Virgo, which can only detect merging BHBs in the final inspiral phase before merger, LISA will be able to detect eccentric stellar-mass BHBs for long timescales before they merge in the LIGO\/Virgo band (e.g.,"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1971,1990]],"Functions Start End":[[1750,1970]]}
{"Identifier":"2022AandA...659A...5Y__Mills_et_al._2018_Instance_2","Paragraph":"Since its discovery more than five decades ago (Cheung et al. 1968), ammonia (NH3) has been a most valuable molecule for investigating the physical properties of molecular clouds (e.g., Ho & Townes 1983). While thermally excited transitions in the centimeter-wavelength inversion transitions of ammonia are regarded as a reliable thermometer of molecular clouds (e.g., Walmsley & Ungerechts 1983; Danby et al. 1988), ammonia masers have attracted attention since the first detection of maser action in the (J, K) = (3,3) metastable (J = K) line toward the massive star-forming region W33 (Wilson et al. 1982). Subsequent observations have led to the detection of new metastable ammonia masers, including 15NH3 (3,3) (Mauersberger et al. 1986), NH3 (1,1) (Gaume et al. 1996), NH3 (2,2) (Mills et al. 2018), NH3 (5,5) (Cesaroni et al. 1992), NH3 (6,6) (Beuther et al. 2007), NH3 (7,7), NH3 (9,9), and NH3 (12,12) (Henkel et al. 2013). These have led to the discovery of metastable maser lines in 22 different regions (Mauersberger et al. 1986, 1987; Wilson & Henkel 1988; Wilson et al. 1990; Pratap et al. 1991; Cesaroni et al. 1992; Wilson & Schilke 1993; Mangum & Wootten 1994; Kraemer & Jackson 1995; Zhang & Ho 1995; Zhang et al. 1999; Walsh et al. 2007; Hunter et al. 2008; Galv\u00e1n-Madrid et al. 2009; Brogan et al. 2011; Urquhart et al. 2011; Walsh et al. 2011; Wang et al. 2012; Henkel et al. 2013; Hoffman & Joyce 2014; McEwen et al. 2016; Mills et al. 2018; Hogge et al. 2019; Mei et al. 2020; Towner et al. 2021). Compared with the metastable ammonia masers, detected non-metastable (J > K) ammonia maser transitions are more numerous. The first highly excited non-metastable ammonia maser was detected by Madden et al. (1986) in the (J, K) = (9,6) and (6,3) lines. Thereafter, many other NH3 non-metastable inversion transition lines have been identified as masers, including the (5,3), (5,4), (6,1), (6,2), (6,4), (6,5), (7,3), (7,4), (7,5) (7,6), (8,3), (8,4), (8,5), (8,6), (9,3), (9,4), (9,5), (9,7), (9,8), (10,7), (10,8), (10,9), and (11,9) transitions (e.g., Mauersberger et al. 1987, 1988; Walsh et al. 2007; Henkel et al. 2013; Mei et al. 2020). Except for the NH3 (3,3) masersproposed to be associated with four supernova remnants (McEwen et al. 2016), almost all the other ammonia masers are detected in high-mass star-forming regions (HMSFRs). However, while many HMSFRs host water (H2O), hydroxyl (OH), or methanol (CH3OH) masers, ammonia masers are quite rare in these sources, and the role that the environment of a young high-mass star plays in their excitation remains unclear. Therefore, dedicated searches for ammonia masers in HMSFRs are indispensable in regard to their overall incidence and association with different environments, which can provide additional constraints on the pumping mechanism of ammonia masers.","Citation Text":["Mills et al. 2018"],"Functions Text":["These have led to the discovery of metastable maser lines in 22 different regions"],"Functions Label":["Background"],"Citation Start End":[[1445,1462]],"Functions Start End":[[933,1014]]}
{"Identifier":"2016ApJ...832...52F__Harris_et_al._2012_Instance_1","Paragraph":"However, most Herschel sources are not SMGs; they are, instead, less luminous dusty star-forming galaxies at lower redshifts (z  2; Casey et al. 2012, 2014). To select Herschel sources that are likely to be SMGs, we chose only the subsample that satisfies the following criteria: (1) flux density peak at 350 \u03bcm (\n\n\n\n\n\n and S500 S350; i.e., \u201c350 \u03bcm peakers\u201d), (2) S500 > 20 mJy, and (3) >3\u03c3 detections in all three SPIRE bands. Criterion 1 is essentially a photometric redshift selection because emission from dusts at T = 35 K would peak at 350 \u03bcm if redshifted to z \u223c 2.5. This is confirmed by the blind carbon monoxide (CO J = 1\u20130) survey of a subsample of the brightest 350 \u03bcm peakers (S350 \u2265 115 mJy), which has shown a strikingly similar redshift distribution as 850 \u03bcm selected SMGs (zCO = 2.5 \u00b1 0.8; Harris et al. 2012). But note that most of these bright sources are strongly lensed and they do not overlap with our sample. Criterion 2 is introduced to ensure that the Rayleigh\u2013Jeans extrapolation would give S850 > 3 mJy, the classic definition of an SMG, given a typical power-law slope of 3.5 for a modified blackbody with a frequency-dependent absorption cross section (\n\n\n\n\n\n). Criterion 3 ensures that all of the sources we considered are statistically significant. This is necessary because the image depth varies substantially from field to field, ranging from \u03c3500 = 15 mJy beam\u22121 for the large HeLMS and HerS fields (confusion noise included; Oliver et al. 2012; Viero et al. 2014) to confusion limited with \u03c3500 = 6.8 mJy beam\u22121 for the deeper HerMES fields (Nguyen et al. 2010). Given the range of observed far-IR SEDs at z = 2 from Casey et al. (2012), our color selection and the high threshold on the 500 \u03bcm flux density ensure that \u223c95% of our sample would be classified as SMGs if they were observed at 870 \u03bcm. Nevertheless, we should keep in mind that the Herschel-selected SMGs are a subsample of SMGs and they likely cover a smaller range of dust temperatures than 870 \u03bcm selected SMGs (e.g., Hwang et al. 2010; Magnelli et al. 2012). Only 70,823 Herschel sources remained after this selection. The average surface density of 92 deg\u22122 is five times lower than the observed 870 \u03bcm source count above S870 \u2273 3 mJy (\u223c500 deg\u22122; Coppin et al. 2006; Wei\u00df et al. 2009). This is not surprising given that almost half of the total Herschel area is only 10%\u201320% complete at S500 = 20 mJy. Note that this incompleteness in the Herschel catalogs is not a concern for compiling a sample of SMG\u2212QSO pairs.","Citation Text":["Harris et al. 2012"],"Functions Text":["This is confirmed by the blind carbon monoxide (CO J = 1\u20130) survey of a subsample of the brightest 350 \u03bcm peakers (S350 \u2265 115 mJy), which has shown a strikingly similar redshift distribution as 850 \u03bcm selected SMGs (zCO = 2.5 \u00b1 0.8;"],"Functions Label":["Similarities"],"Citation Start End":[[808,826]],"Functions Start End":[[575,807]]}
{"Identifier":"2019ApJ...885...81S__Grogin_et_al._2011_Instance_1","Paragraph":"We measured the apparent axial ratio on the IF814W-band data, which correspond to the rest-frame V band for galaxies at z \u223c 0.4 and the rest-frame B band for those at z \u223c 0.8. Such differences in the rest-frame wavelength could cause some biases in the morphological analysis due to the color differences between bulge and disk, the blue star-forming regions\/clumps, the dust extinction effect, and so on (e.g., Windhorst et al. 2002; Huertas-Company et al. 2009; Wuyts et al. 2012; Vika et al. 2013; Murata et al. 2014; Mager et al. 2018). In order to check the effects of the morphological K-correction, we used publicly available HST\/ACS VF606W-band data over a 0.05 deg2 region in the COSMOS field from the CANDELS survey (Grogin et al. 2011; Koekemoer et al. 2011). With the VF606W-band data, we can measure the apparent axial ratio of galaxies at 0.2  z  0.6 in the rest-frame B band and investigate to what extent the difference in the rest-frame wavelength affects the measurements. There are 92 main-sequence and 51 passively evolving galaxies with VF606W  25 at 0.2  z  0.6 in the region, and we measured the apparent axial ratio of these galaxies on the VF606W-band data in the same way. In Figure 17, we compare the apparent axial ratios b\/a measured on the VF606W-band data with those measured on the IF814W-band data. The differences between the VF606W and IF814W bands are also summarized in Table 4. The apparent axial ratios measured on the VF606W- and IF814W-band data agree well with each other for both main-sequence and passively evolving populations. The average values of (b\/a)F606W \u2212 (b\/a)F814W are \u22120.006 and \u22120.007 for main-sequence and passively evolving galaxies, respectively. These systematic offsets from zero are slightly larger than the averages of the measurement errors but much smaller than the dispersion around the mean value. When we use only bright subsamples with VF606W  22, the results do not significantly change, although the average offsets and measurement errors become slightly smaller. Since these systematic offsets are much smaller than the bin width of 0.1 in the distribution of b\/a we used, the morphological K-correction does not significantly affect the distribution of the apparent axial ratio.","Citation Text":["Grogin et al. 2011"],"Functions Text":["In order to check the effects of the morphological K-correction, we used publicly available HST\/ACS VF606W-band data over a 0.05 deg2 region in the COSMOS field from the CANDELS survey"],"Functions Label":["Uses"],"Citation Start End":[[727,745]],"Functions Start End":[[541,725]]}
{"Identifier":"2021ApJ...912..163B__Kleine_et_al._2020_Instance_1","Paragraph":"Shortly after the ignition of the Sun, the solar system contained a colossal cloud of dust and gas known as the protoplanetary disk. Within just a few million years, trillions of submillimeter-scale solids formed from this disk and coalesced to create a spectrum of planetary bodies that included asteroids, comets, and ultimately the planets that we recognize today. The mechanisms by which these early solids formed and evolved in the disk as well as the processes by which they interacted with each other and the remaining gas and dust governed the chemical and physical properties adopted by these bodies. As such, these early processes played pivotal roles in establishing the long-term thermochemical evolutions of the different planetary bodies found throughout our solar system. Aggregates of these early solids, as well as fragments of the first planetary bodies that formed in the solar system, exist on Earth today in the form of meteorites. Through novel high-resolution and high-sensitivity techniques, these samples have recently provided several new insights into the evolution of the protoplanetary disk and the formation of the first solids. For instance, a number of measurements have revealed that the stable isotopic compositions of all meteorites fall into two distinct families (Warren 2011; Scott et al. 2018; Kruijer et al. 2019; Kleine et al. 2020; Bermingham et al. 2020). This dichotomy has been used to argue that these extraterrestrial rocks each originate from one of two reservoirs of material that existed in the protoplanetary disk that were spatially separated by a large feature (possibly Jupiter; Kruijer et al. 2017; Brasser & Mojzsis 2020; Lichtenberg et al. 2021) that hindered the exchange of material between these regions. This restricted motion prevented disk-wide mixing and compositional homogenization, allowing the distinct composition of each reservoir to form and be retained over much of the disk\u2019s lifetime. One key observation supporting this dichotomy is the \u03b594Mo and \u03b595Mo values (where \u03b5 denotes normalized isotopic concentration in parts per 10,000) of iron meteorites (fragments of the metallic cores of melted asteroids), rocky achondrites (fragments of the mantles and crusts of melted asteroids), and chondrites (fragments of unmelted asteroids), which together form two parallel lines when plotted against each other (Budde et al. 2016; Kruijer et al. 2017; Budde et al. 2019). Additionally, the isotopic compositions of a suite of other elements (including \u03b548Ca, \u03b550Ti, \u03b554Cr, \u03b562Ni, \u03b5100Ru, and \u25b317O values, where \u25b3 denotes the normalized non-mass-dependent isotopic concentration in parts per 1000) all form pairs of clusters when plotted against each other, further supporting the existence of separated reservoirs composed of material with isotopically distinct compositions in the protoplanetary disk (Warren 2011; Fischer-Godde & Kleine 2017; Bermingham et al. 2018; Schiller et al. 2018; Nanne et al. 2019; Worsham et al. 2019). In all of these cases, the same meteorite groups systematically share similar isotopic signatures and fall into the same families. As such, all meteorites can be primarily categorized as either noncarbonaceous (NC) or carbonaceous (CC) based on their isotopic compositions (named after the type of chondrite found within each family; Budde et al. 2016).","Citation Text":["Kleine et al. 2020"],"Functions Text":["For instance, a number of measurements have revealed that the stable isotopic compositions of all meteorites fall into two distinct families"],"Functions Label":["Background"],"Citation Start End":[[1354,1372]],"Functions Start End":[[1159,1299]]}
{"Identifier":"2016MNRAS.457.2433P__Nolan_et_al._2012_Instance_2","Paragraph":"From the result of the \u03c72 minimization, we found that the minimized \u03c72 values agree with the expected values, i.e. the computed \u03c72 are typically in the range of ($\\mathrm{d.o.f.}-\\sqrt{2 \\mathrm{d.o.f.}}$, $\\mathrm{d.o.f.}+\\sqrt{2 \\mathrm{d.o.f.}}$), where d.o.f. is the number of degrees of freedom. This means that the fits describe the observed data rather well. The only exception is with \u03c72 \u2248 20, which occurs for nearby AGN, z  0.2, and for the highest energy band, E > 10 GeV. Note that there is a strong contribution of the source, Mrk 421, in the first redshift interval at high energies, E > 10 GeV for quiescent states. Mrk 421 is a very hard spectrum \u03b3-ray source with a photon index of \u22481.77 and its semiminor and semimajor axes at 68 per cent confidence are of 0$_{.}^{\\circ}$0067 as derived in Nolan et al. (2012). Semiminor and semimajor axes of many 2FGL sources are derived with one order of magnitude higher uncertainties than those for Mrk 421 in the 2FGL catalogue (Nolan et al. 2012). We noted that the discrepancy between the observation and model is particularly strong in the annular bin, r  0$_{.}^{\\circ}$05, for the redshift interval z  0.2 and for the highest energy band, E > 10 GeV. If we exclude photons from Mrk 421, then the minimized \u03c72 value is 7.5 and is consistent with the expected one. In the limit of a large number of counts in each bin, the likelihood is given by $\\mathcal {L}=\\text{exp}(-\\chi ^{2}\/2)$, so that minimizing \u03c72 is equivalent to maximizing the likelihood, $\\mathcal {L}$. We found that the inclusion of a pair halo component in the model does not improve the likelihood value sufficiently to establish the presence of this pair halo component in the data. Therefore, we derived the one-sided 95 per cent upper limit on the fraction of photons attributable to a pair halo component by fitting the normalization of this component, for which we increase its normalization until the maximum likelihood decreases by 2.71\/2 in logarithm. The computed upper limits are between 2 and 6 per cent depending on energy band and redshift interval. These upper limits are stronger than those obtained before. Note that the model for a point-like source used in the likelihood analysis is considered to be precisely established, however, the number of photons recorded during flaring states is close to those numbers of photons recorded during quiescent states for each of these redshift intervals. The expression, such as equation (1), leads to more conservative upper limits on the fraction of photons attributable to a pair halo component, since it takes the error bars assigned to the model into account. If the point-like source model is considered well established, then the error bars shown in Table 3 would decrease by a factor of \u22481.5.","Citation Text":["Nolan et al. 2012"],"Functions Text":["Semiminor and semimajor axes of many 2FGL sources are derived with one order of magnitude higher uncertainties than those for Mrk 421 in the 2FGL catalogue"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[987,1004]],"Functions Start End":[[830,985]]}
{"Identifier":"2016MNRAS.463.2716M__Cho_&_Lazarian_2007_Instance_1","Paragraph":"Returning to the case of HL Tau, where the possible contribution of an infalling envelope is not an issue, how can one reconcile the strong indication of a dominant radial field component in the polarization map with the expectation that the bulk of the mm-wavelength emission originates near the disc mid-plane, where the azimuthal field component dominates? One possibility is that the non-negligible optical depth inferred in the bright emission rings of HL Tau at mm-wavelengths (Jin et al. 2016; Pinte et al. 2016) shifts the emission centroid to finite disc elevations where the magnetic field already has a measurable radial component. However, in view of the very small scale height inferred for the mm-emitting dust in this source, this is unlikely to be the main explanation. Perhaps a more likely possibility is that, even in this comparatively young source, the grains near the mid-plane, which dominate the total intensity, have already grown to sizes that exceed the maximum size $a_\\mathrm{max} = \\lambda \/2\\pi$ for producing polarized emission at wavelength \u03bb (e.g. Cho & Lazarian 2007; for \u03bb = 1.25\u2009mm, amax = 0.2\u2009mm), while the smaller grains (with sizes a amax), which contribute efficiently to the polarized flux, remain suspended at high elevations (where the field is predominantly radial). Another effect that could lower the polarized emission from grains that have settled to the mid-pane is the likelihood that grains become less elongated as they grow (e.g. Hughes et al. 2009), which would tend to reduce the value of the coefficient C in equation (2) (C \u2192 0 as the grain axis ratio \u2192 1).7 This interpretation is supported by the finding in the high-resolution observations of IRAS 4A (Cox et al. 2015) of an average polarization of 15 per cent at 8\u2009mm and 10 per cent at 10\u2009mm, with a peak fractional polarization of \u223c20 per cent. If the intrinsic degree of mm-wavelength polarization in HL Tau is also of the order of 20 per cent, then it may be possible to explain the factor of \u223c10 lower value of P measured in this source at 1.25\u2009mm8 in terms of a dilution of the polarized emission from a \u2272 0.1\u2009mm grains at high disc elevations by weakly polarized emission of larger grains residing near the mid-plane. In this scenario, most of the grains that are responsible for the mm-wavelength flux have settled to the mid-plane and grown to sizes a \u2273 1\u2009mm.9 Although a fraction of these grains may have sizes in excess of 1\u2009mm and would therefore emit less efficiently at that wavelength than a \u2272 1\u2009mm grains (e.g. Miyake & Nakagawa 1993), the mid-plane region should still dominate the total mm-wavelength flux if most of the a \u2273 1\u2009mm grains are concentrated there. Grains of size a \u2272 0.1\u2009mm may be kept at high elevations by turbulent motions that can persist below the wind-driving surface layers (e.g. Simon et al. 2013, 2015; Bai 2015) as well as by the emerging outflows within these layers (see Safier 1993), with porosity effects (e.g. Ormel, Spaans & Tielens 2007) possibly also helping to mitigate gravity's pull towards the mid-plane. This scenario of course needs to be backed by detailed calculations and observational tests. One such test would be to obtain a polarization map of HL Tau at longer ( \u2273 1\u2009cm) wavelengths: if the above picture is correct and the grains in the mid-plane region are aligned, such a map could reveal a stronger (or even dominant) contribution from the azimuthal and (especially if \u039b0 \u226a 1) vertical field components.10 It is, however, conceivable that the large grains in this source are not well aligned because the radiative torque mechanism does not operate efficiently on them: this could happen if the characteristic wavelength of the anisotropic component of the local radiation field were much smaller than the mid-plane grain sizes (e.g. Cho & Lazarian 2007) or if the anisotropic radiation component inside the dust disc were weak due to finite optical depth effects. Note that the possibility of the mid-plane grains not being well aligned provides another reason for why the polarized mm-wavelength emission from this region could be weak.","Citation Text":["Cho & Lazarian 2007"],"Functions Text":["Perhaps a more likely possibility is that, even in this comparatively young source, the grains near the mid-plane, which dominate the total intensity, have already grown to sizes that exceed the maximum size $a_\\mathrm{max} = \\lambda \/2\\pi$ for producing polarized emission at wavelength \u03bb (e.g.","; for \u03bb = 1.25\u2009mm, amax = 0.2\u2009mm)"],"Functions Label":["Uses","Uses"],"Citation Start End":[[1082,1101]],"Functions Start End":[[786,1081],[1101,1134]]}
{"Identifier":"2015MNRAS.454.2003P__Paardekooper_&_Papaloizou_2009_Instance_1","Paragraph":"In both of these approaches, disc\u2013planet interactions causing planet migration are modelled using analytical formulae for the disc torque experienced by the planet. This gravitational torque consists of two components. The differential Lindblad torque results from the angular momentum exchange between the planet and the spiral density waves it generates inside the disc. For sufficiently low-mass planets, it scales linearly with disc mass and planet mass and as the inverse square of the disc aspect ratio (Tanaka, Takeuchi & Ward 2002). Although its sign depends on the density and temperature gradients inside the disc, the differential Lindblad torque is generally negative for typical disc models and is therefore responsible for inward migration. The corotation torque is due to the torque exerted by the material located in the coorbital region of the planet. It is composed of a barotropic part which scales with the vortensity (i.e. the ratio between the vertical component of the vorticity and the disc surface density) gradient (Goldreich & Tremaine 1979) plus an entropy-related part which scales with the entropy gradient (Baruteau & Masset 2008; Paardekooper & Papaloizou 2008). A negative vortensity (resp. entropy) gradient gives rise to a positive vortensity (resp. entropy) related corotation torque. It has been shown that for mildly positive surface density gradients or negative entropy gradients, a positive corotation torque can eventually counteract the effect of a negative differential Lindblad torque, which may stall or even reverse migration (Masset, D'Angelo & Kley 2006; Paardekooper & Papaloizou 2009). In isothermal discs, the corotation torque is a non-linear process generally referred as the horseshoe drag and whose amplitude is controlled by advection and diffusion of vortensity inside the horseshoe region. In non-isothermal discs, the corotation torque is also powered by singular production of vortensity due to an entropy discontinuity on downstream separatrices (Masset & Casoli 2009; Paardekooper et al. 2010). In the absence of any diffusion processes inside the disc, vortensity and entropy gradients across the horseshoe region tend to flatten through phase mixing, which causes the two components of the horseshoe drag to saturate. Consequently, desaturating the horseshoe drag requires that some amount of viscous and thermal diffusions are operating inside the horseshoe region. In that case, the amplitude of the horseshoe drag depends on the ratio between the diffusion time-scales and the horseshoe libration time-scale and its optimal value, also referred as the fully unsaturated horseshoe drag, is obtained when the diffusion time-scales are approximately equal to half the horseshoe libration time (e.g. Baruteau & Masset 2013). In the limit where the diffusion time-scales become shorter than the U-turn time-scale, the corotation torque decreases and approaches the value predicted by linear theory. Therefore, the corotation torque can be considered as a linear combination of the fully unsaturated horseshoe drag and the linear corotation torque with coefficients depending on the ratio between the diffusion time-scales and the horseshoe libration time-scale. Corotation torque formulae as a function of viscosity and thermal diffusivity were recently proposed by Paardekooper, Baruteau & Kley (2011) and Masset & Casoli (2010).","Citation Text":["Paardekooper & Papaloizou 2009"],"Functions Text":["It has been shown that for mildly positive surface density gradients or negative entropy gradients, a positive corotation torque can eventually counteract the effect of a negative differential Lindblad torque, which may stall or even reverse migration"],"Functions Label":["Background"],"Citation Start End":[[1604,1634]],"Functions Start End":[[1321,1572]]}
{"Identifier":"2021AandA...651A..87O__Brunthaler_et_al._2021_Instance_2","Paragraph":"To complement our study, we also analyzed GLOSTAR continuum images toward sites with maser emission. A full description of the GLOSTAR continuum data calibration and imaging is given in Brunthaler et al. (2021), while the full analysis of continuum images of Cygnus X will be presented in a forthcoming paper. Here, we briefly discuss the imaging strategy. The calibration and imaging of the continuum data was performed with the Obit package (Cotton 2008). The 2 GHz bandwidth was first rearranged into nine frequency subbands, which were used to image each pointing individually. Then, for each frequency subband the pointings were combined into large individual mosaics to cover the entire observed area. Finally, we combined the different frequencies to obtain the image at the reference frequency, which has circular beams of 19\u2033 and 1.5\u2033 in the D and Bconfiguration, respectively. Continuum and methanol line maps from Effelsberg observations have also been obtained as part of the GLOSTAR survey (Brunthaler et al. 2021, Rugel et al., in prep.) We note that continuum images were constructed for Effelsberg data, the VLA D configuration, the VLA B configuration, a combination of the VLA D and B (D+B) configurations, and a combination of the VLA D configuration and Effelsberg observations. The central frequency of these images is 5.8 GHz. Here, we only use B-configuration continuum maps to study the region of the investigated methanol maser positions and D+B maps of the region around DR21 (see Sect. 4.5). Methanol line data from Effelsberg were also inspected to look for flux variations in the VLA-detected masers (Sect. 4.4). The noise in the continuum images is not uniform, but rather varies across the mapped region, and can be high around strong sources with complex or extended emission. We locally measured the noise in regions close to the maser locations, resulting in 1\u03c3 values in the range from 0.056 to 0.43 mJy beam\u22121 for B-configuration images. For the D configuration, the 1\u03c3 rms noise ranges from 0.10 to 2.6 mJy beam\u22121. The higher values measured in D-configuration data are due to bright extended emission, which is present across the Cygnus X region, and are resolved out by the array in the B configuration. The highest local rms noise occurs around the strong radio source, DR21, a compact HII region.","Citation Text":["Brunthaler et al. 2021"],"Functions Text":["Continuum and methanol line maps from Effelsberg observations have also been obtained as part of the GLOSTAR survey"],"Functions Label":["Uses"],"Citation Start End":[[1004,1026]],"Functions Start End":[[887,1002]]}
{"Identifier":"2021AandA...652A..98G__Buta_et_al._(2009)_Instance_1","Paragraph":"On the CND scales displayed in Fig. 12, the CO(3\u20132) emission in this highly inclined (i\u2004=\u200459\u00b0; Appendix D) barred Seyfert 2 galaxy is detected at every single position inside r\u2004\u2004200 pc. However, a sizeable fraction of the molecular gas appears concentrated in a nuclear ring of \u223c200 pc (deprojected) radius. Molecular gas in this nuclear ring is feeding an active star formation episode detected at optical as well as near and mid IR wavelengths. The molecular gas ring is the likely signature of the gas response to the \u223c140\u2033 (16 kpc) long stellar bar at its ILR region. The bar, detected in the NIR by Quillen et al. (1997) and Buta et al. (2009), shows a prominent boxy-shape morphology. At the larger radii imaged inside the ALMA FOV (Fig. 10), molecular gas shows a two-arm spiral structure that is connected to the nuclear ring. Closer to the Seyfert 2 nucleus (r\u2004\u200450 pc), molecular gas probed by CO is concentrated in an asymmetric ringed disk of \u223c30\u2005\u2212\u200540 pc radius located around the AGN and oriented along \u223c160\u00b0. The molecular disk shows a similar orientation to the extended component of the 351 GHz continuum emission detected by ALMA (PAGauss\u2004\u223c\u2004160\u00b0\u2005\u2212162\u00b0; Table 4). The CO disk appears to be partly incomplete: its southwest hemisphere is on average a factor of 3 weaker than its northeast counterpart. The V\u2005\u2212\u2005H map shows a dusty ring feature in excellent correspondence with the morphology of the CO ringed disk. There is nevertheless significant molecular gas inside the ring: CO emission is detected toward the position of the AGN defined by the position of the ALMA 351 GHz continuum point source. This implies values for the molecular gas mass and H2 column densities of Mgas[r\u2004\u2264\u20047\u2005\u2212\u20059 pc] \u2004\u223c\u2004(1.5\u2005\u2212\u20052.9)\u00d7105\u2006M\u2299 and N(H2)\u223c(3.5\u2005\u2212\u20053.6)\u00d71022 cm\u22122, respectively. The galaxy shows a bright ionization cone southwest of the nucleus oriented along PAout\u2004\u223c\u2004235\u00b0\u2005\u2212245\u00b0 and characterized by blueshifted velocities, which betray an outflow (Morris et al. 1985; Storchi-Bergmann & Bonatto 1991; Davies et al. 2016; Ricci et al. 2018a). The northeast ionization cone appears to be obscured by the host galaxy. The orientation of the 70\u2005\u2212\u200590 pc-size disk detected by ALMA in continuum and CO is therefore equatorial (torus-like).","Citation Text":["Buta et al. (2009)"],"Functions Text":["The bar, detected in the NIR by Quillen et al. (1997) and","shows a prominent boxy-shape morphology."],"Functions Label":["Background","Background"],"Citation Start End":[[630,648]],"Functions Start End":[[572,629],[650,690]]}
{"Identifier":"2022AandA...662A..34D__Hyde_&_Bernardi_2009_Instance_1","Paragraph":"The marginalized probability for Mgal and Mhalo is shown in Fig. 6. The contours correspond to the 68% and 95% intervals and the black dot marks the position of the best model. We note that the best model is outside the 68% region of the marginalized probability. The N-dimensional likelihood near the best model forms a shallow valley that extends toward the 68% confidence region. The best models clearly prefer a relatively narrow region in the Mgal\u2005\u2212\u2005MDM space. An interesting result from this figure, is that a baryon-only model (i.e., MDM\u2004=\u20040) does not perform much worse than a model with a DM halo, although the data prefers a model with a DM halo. The relatively weaker dependence with the DM component may be a consequence of the small size of the Einstein ring, since at this radius the baryon component dominates the projected mass (see Table 1 below). The inferred mass of the galaxy (baryons) in the best model (5.4\u2005\u00d7\u20051010 M\u2299) is marginally consistent with the estimated stellar mass derived from the velocity dispersion \u2013 stellar mass correlation (Hyde & Bernardi 2009; Zahid et al. 2016; Cannarozzo et al. 2020). Assuming the velocity dispersion in M\u00f6rtsell et al. (2020), and the models in Zahid et al. (2016), Cannarozzo et al. (2020), the predicted stellar mass is \u22482.5\u2005\u00d7\u20051010 M\u2299. This is about a factor two less than the stellar mass inferred from the lens model. However, a factor two uncertainty is typical in mass estimations from the velocity dispersion (see the references above). Since the galaxy and halo masses are partially degenerate (see Fig. 6), it is possible that the galaxy mass is lower than the one at the maximum of the likelihood. Another possibility is that the gas mass is similar to that of the stars, increasing the baryonic mass by a factor two compared with the stellar mass. This is in principle feasible based on the gas to stellar mass ratios, which is close to 1 for early type galaxies (similar to the lens considered in this work) and stellar masses \u223c1010 M\u2299 (Calette et al. 2018). However, in the central part of the galaxy the stellar component is still expected to dominate, especially if a bulge is present, as suggested by the data. An alternative way of estimating the contribution from the gas is by comparing with galaxies of similar morphological type. According to Casasola et al. (2020), and considering a morphological type for the lens galaxy between Sa and Sb, the gas fraction for this type of galaxy is \u224820%. In Schruba et al. (2011), gas and molecular surface densities as high as O(100) M\u2299 pc\u22122 can be found in areas with large star formation rates. In the most extreme cases of star formation rates, gas surface densities of \u22481000 M\u2299 pc\u22122 can be found. However, even at these extreme star formation rates, the contribution from the gas to the baryonic mass is still subdominant with respect to the stellar contribution from our fiducial model.","Citation Text":["Hyde & Bernardi 2009"],"Functions Text":["The inferred mass of the galaxy (baryons) in the best model (5.4\u2005\u00d7\u20051010 M\u2299) is marginally consistent with the estimated stellar mass derived from the velocity dispersion \u2013 stellar mass correlation"],"Functions Label":["Similarities"],"Citation Start End":[[1063,1083]],"Functions Start End":[[865,1061]]}
{"Identifier":"2020ApJ...892...53A__Connolly_et_al._2018_Instance_1","Paragraph":"These new limits, in conjunction with the inconsistency of isotropic flux interpretations, leave no room for an astrophysical interpretation of AAEs in the context of the standard model for time windows as short as 103 s. However, it has been shown that these events can be explained using physics beyond the standard model, as many models suggest that AAEs lend support for axionic dark matter, sterile neutrinos, supersymmetry, or heavy dark matter (Anchordoqui et al. 2018; Connolly et al. 2018; Dudas et al. 2018; Fox et al. 2018; Huang 2018; Abdullah et al. 2019; Anchordoqui & Antoniadis 2019; Borah et al. 2019; Chauhan & Mohanty 2019; Cherry & Shoemaker 2019; Chipman et al. 2019; Cline et al. 2019; Collins et al. 2019; Esmaili & Farzan 2019; Esteban et al. 2019; Heurtier et al. 2019a, 2019b; Hooper et al. 2019). Many of these models, excluding the axionic dark matter explanation (Esteban et al. 2019) or those heavy dark matter scenarios that are tuned to prevent signatures in IceCube (Hooper et al. 2019), can be constrained by this nonobservation at IceCube. Dedicated tests to quantify these constraints are beyond the scope of this work and may be the focus of a future study. In addition to explanations that incite new physics, it has recently been suggested that AAEs could be explained by downward-going CR-induced EASs that reflected off of subsurface features in the Antarctic ice (Shoemaker et al. 2019). Another possible explanation could be coherent transition radiation from the geomagnetically induced air shower current, which could mimic an upgoing air shower (Motloch et al. 2017; de Vries & Prohira 2019). Explaining these anomalous events with systematic effects or confirming the need for new physics requires a deeper understanding of ANITA\u2019s detection volume. Efforts such as the HiCal radio frequency pulser, which has flown alongside ANITA on the last two flights (Prohira et al. 2018), are already underway to try to characterize the various properties of the Antarctic ice surface.","Citation Text":["Connolly et al. 2018"],"Functions Text":["These new limits, in conjunction with the inconsistency of isotropic flux interpretations, leave no room for an astrophysical interpretation of AAEs in the context of the standard model for time windows as short as 103 s. However, it has been shown that these events can be explained using physics beyond the standard model, as many models suggest that AAEs lend support for axionic dark matter, sterile neutrinos, supersymmetry, or heavy dark matter","Many of these models,","can be constrained by this nonobservation at IceCube."],"Functions Label":["Motivation","Motivation","Motivation"],"Citation Start End":[[477,497]],"Functions Start End":[[0,450],[824,845],[1021,1074]]}
{"Identifier":"2020MNRAS.494.3413T__Shidatsu_&_Done_2019_Instance_2","Paragraph":"The existence of winds is shown by blueshifted absorption lines from highly ionized ions. These are only seen in soft state but not in hard state (Ponti et al. 2012), anticorrelated with the radio jet which is seen in the hard state but not in the soft. This was thought to be evidence that the wind was magnetically driven by the same field as was responsible for the jet, but in a different geometric configuration (Miller et al. 2012). However, in Tomaru et al. (2019, hereafter Paper I) we show instead that thermally driven winds can explain this switch (see also Done, Tomaru & Takahashi 2018; Shidatsu & Done 2019). Thermal driving produces a wind by irradiation from the central source heating the surface of accretion disc up to the Compton temperature ($T_\\text{IC} \\sim 10^7 \\!-\\!10^8\\, \\text{K}$), which is hot enough for its thermal energy to overcome the gravity at large radii. The characteristic radius at which the wind can be launched is called the Compton radius, defined by RIC = \u03bcmpGM\/kTIC \u223c 105 \u2212 106Rg (Begelman, McKee & Shields 1983). Paper I show the first modern radiation hydrodynamic simulations of thermal (and thermal-radiative) winds designed to investigate the switch in wind properties between the hard and soft states changing illumination spectra. These simulations were tailored to the BHB system H1743\u2212332, where there is Chandra high-resolution data in both states giving detailed spectral information on the wind or its absence (Miller et al. 2012; Shidatsu & Done 2019). They incorporate radiation force on the electrons, both bound and free, as they show that this is important factor driving the escape of the thermal wind in the fairly high Eddington fraction (L\/LEdd \u223c 0.2\u20130.3), fairly low Compton temperature (TIC \u223c 0.1 \u00d7 108 K) characteristics of the soft state. The only other modern hydrodynamic simulation of thermal winds (e.g. Luketic et al. 2010; Higginbottom & Proga 2015; Higginbot et al. 2016) has not included radiation pressure, which is important in setting the velocity structure for L \u2265 0.3LEdd as required here (Paper I).","Citation Text":["Shidatsu & Done 2019"],"Functions Text":["These simulations were tailored to the BHB system H1743\u2212332, where there is Chandra high-resolution data in both states giving detailed spectral information on the wind or its absence"],"Functions Label":["Uses"],"Citation Start End":[[1488,1508]],"Functions Start End":[[1283,1466]]}
{"Identifier":"2022AandA...659A.124H__Liu_et_al._(2013b)_Instance_2","Paragraph":"Combining different samples from various instruments at different redshifts therefore inevitably introduces ENLR size\u2013luminosity relations with different slopes \u03b1 depending on the details of target selection and analysis approaches. Slopes ranging from \u03b1\u2004=\u20040.22\u2005\u00b1\u20050.04 (Greene et al. 2012), \u03b1\u2004=\u20040.25\u2005\u00b1\u20050.02 (Liu et al. 2013b), \u03b1\u2004\u223c\u20040.3\u20130.4 (Hainline et al. 2013; Chen et al. 2019a), to \u03b1\u2004\u223c\u20040.5 (Bennert et al. 2002; Husemann et al. 2014) are reported in the literature. The slopes solely inferred from the CARS data are consistent with those reported by Greene et al. (2012) and Liu et al. (2013b) and are therefore on the shallower side of previous estimates. Nevertheless, the scatter in the observed relation is significant and measured slope variations might be entirely attributed to the observationally induced biases as discussed above. A slope of \u03b1\u2004=\u20040.5 is reminiscent of the BLR size-luminosity relation, but would require a constant ionization parameter U that demands a constant density with radius. This is not observed for the ENLR on kiloparsec scales (e.g., Bennert et al. 2006; Kakkad et al. 2018) and more detailed photoionization calculations are required to predict the shallower slopes inferred for most studies (Dempsey & Zakamska 2018). We cannot study the radial variations of U as our snapshot MUSE observations are not deep enough to map the electron density given the too low S\/N of the [S\u202fII] doublet on kpc scales. However, the photoionization calculations do not take into account variable ionizing flux from AGN on 105 yr time scales (Schawinski et al. 2015) and the various geometrical intersections of the ionizing radiation field with the gas distribution of the galaxies. The CARS survey is least biased with regard to RENLR given the narrow redshift range and large dynamic range offered by MUSE (see Fig. 13). Therefore, the CARS survey is one of the best data set to explore the origin of the significant scatter in ENLR size\u2013luminosity relation and search for additional factors or more fundamental parameters controlling the ENLR size.","Citation Text":["Liu et al. (2013b)"],"Functions Text":["The slopes solely inferred from the CARS data are consistent with those reported by Greene et al. (2012) and","and are therefore on the shallower side of previous estimates."],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[578,596]],"Functions Start End":[[469,577],[597,659]]}
{"Identifier":"2019AandA...630A.151P__Gallerani_et_al._(2011)_Instance_1","Paragraph":"Capturing the full information content of the cosmic large-scale structure requires a field-based approach to infer the entire three-dimensional cosmic large-scale structure from observations. This poses a particular challenge for the analyses of Ly-\u03b1 forest observations, which provide sparse inherently one-dimensional information along the lines of sight. Various approaches to perform three-dimensional density reconstructions from one-dimensional Ly-\u03b1 forests have been proposed in the literature (e.g. Kitaura et al. 2012; Cisewski et al. 2014; Stark et al. 2015a; Ozbek et al. 2016; Horowitz et al. 2019). Gallerani et al. (2011) and Kitaura et al. (2012) proposed a Gibbs sampling scheme to jointly infer density and velocity fields and corresponding power-spectra. However, these approaches assume matter density amplitudes to be log-normally distributed. The log-normal distribution reproduces one- and two-point statistics but fails to reproduce higher-order statistics associated with the filamentary dark matter distribution. In an attempt to extrapolate information from one-dimensional quasar spectra into the three-dimensional volume, Cisewski et al. (2014) applied a local polynomial smoothing method. Ozbek et al. (2016) and Stark et al. (2015a) employed a Wiener filtering approach to reconstruct the three-dimensional density field between lines of sight of Ly-\u03b1 forest data. In order to reproduce higher-order statistics, Horowitz et al. (2019) recently used a large-scale optimization approach to fit a gravitational structure growth model to Ly-\u03b1 data, showing that this approach allows recovering the more filamentary structure of the cosmic web. Although the approach improves over linear and isotropic Wiener filtering approaches, it shows systematic deviations of reconstructed matter power-spectra and underestimates density amplitudes at scales corresponding to the mean separation between lines of sight (Horowitz et al. 2019).","Citation Text":["Gallerani et al. (2011)"],"Functions Text":["and Kitaura et al. (2012) proposed a Gibbs sampling scheme to jointly infer density and velocity fields and corresponding power-spectra.","However, these approaches assume matter density amplitudes to be log-normally distributed. The log-normal distribution reproduces one- and two-point statistics but fails to reproduce higher-order statistics associated with the filamentary dark matter distribution."],"Functions Label":["Background","Compare\/Contrast"],"Citation Start End":[[613,636]],"Functions Start End":[[637,773],[774,1038]]}
{"Identifier":"2016AandA...586A..16B__Mol_&_Defrise_(2004)_Instance_1","Paragraph":"The problem we are facing consists in the reconstruction of two maps of a given region of an astrophysical object: the former consists in point sources, such as stars, of very high intensity, the latter of smooth structures surrounding these sources. The standard approach, i.e. simply solving Eq. (2), fails since the presence of the point sources destroys all the information about the smooth structures. Hence, the main idea is to consider the object x as the sum of two components, namely x = xP + xE, where xP represents the point sources and xE the extended source. This approach was proposed for the first time in De Mol & Defrise (2004), assuming that the positions of the point sources are known. Denoting by pj the position of the jth source, we can write (4)\\begin{equation} \\label{2:point} x_{\\rm P}=\\sum_{j=1}^qc_j\\delta(p_j) , \\end{equation}xP=\u2211j=1qcj\u03b4(pj),where q is the total number of sources, cj is the jth unknown intensity, and \u03b4(pj) is the delta function centred in pj. Thus, xP is a vector with zero entries except in the q positions corresponding to the known locations of the sources. Then, instead of computing f1 on the whole object x in Eq. (2), we only regularize the extended source, since the structure induced on xP already works as a regularization. This requires a slightly modification in the computation of f0. We introduce the vector c = (c1,c2,...,cq)t containing the intensities of the sources, and we define \\hbox{$\\overline{x} = (c^t,x_{\\rm E}^t )^t$}x=(ct,xEt)t and the matrix \\hbox{$\\mathcal{H}=[\\overline{H}, H]$}\u210b=[H,H], where \\hbox{$\\overline{H}=[h_{p_1},h_{p_2},\\dots,h_{p_q}]$}H=[hp1,hp2,...,hpq], with hj denoting the jth column of H. We are hence led to solve(5)\\begin{equation} \\label{2:solve} \\tilde{x} =\\arg\\min_{x\\in\\mathcal{C}} KL(\\mathcal{H} \\overline x+b;g)+\\beta f_1\\left(x_{\\rm E}\\right) . \\end{equation}\u02dcx=argminx\u2208\ud835\udc9eKL(\u210bx+b;g)+\u03b2f1(xE).It may happen that \\hbox{$\\tilde x$}\u02dcx has a loss in contrast, even using the optimal value for \u03b2. To overcome this difficulty, we then propose the use of the inexact Bregman procedure, which permits the use of an overestimation of the regularization parameter, and at the same time allows us to obtain a contrast enhancement. ","Citation Text":["De Mol & Defrise (2004)"],"Functions Text":["Hence, the main idea is to consider the object x as the sum of two components, namely x = xP + xE, where xP represents the point sources and xE the extended source. This approach was proposed for the first time in","assuming that the positions of the point sources are known."],"Functions Label":["Uses","Uses"],"Citation Start End":[[621,644]],"Functions Start End":[[407,620],[646,705]]}
{"Identifier":"2022ApJ...926..155S__Umehata_et_al._2017_Instance_1","Paragraph":"However, the sensitivity limits of most current submillimeter surveys only allow for the study of the most extremely star-bursting systems (e.g., Asboth et al. 2016; Geach et al. 2017; Simpson et al. 2019), which may represent merely a tip of an iceberg of dust-obscured star formation in the early Universe. A potential population of more typical dusty star-forming galaxies at high-z is still to be found (Wang et al. 2017). The Atacama Large Millimetre\/submillimetre Array (ALMA) has now opened a new avenue to refine our understanding of dusty galaxies at high redshifts, enabling to uncover faint SMGs down to a flux level of 0.1\u20131 mJy. Several ALMA blind surveys have been performed and allowed to detect and characterize the faint SMGs across cosmic times (e.g., Aravena et al. 2016; Wang et al. 2016b; Dunlop et al. 2017; Umehata et al. 2017; Franco et al. 2018; Williams et al. 2019; Yamaguchi et al. 2019; Gonz\u00e1lez-L\u00f3pez et al. 2020). Based on the ALMA survey of GOODS-South field over an area of 69 arcmin2, Franco et al. (2018) found that 20% of the 1.1 mm sources are not detected with HST down to a depth of H \u223c 28 mag, and suggested that they are massive main-sequence star-forming galaxies at z > 4 (see also, Yamaguchi et al. 2019; Umehata et al. 2020). A similar fraction of HST-dark galaxies has also been found in the ALMA [C ii] survey of main-sequence galaxies at 4.4  z  5.9 (\u223c14%, Gruppioni et al. 2020). Conversely, the existence of such HST-dark galaxies can also be uncovered by focusing the reddest galaxies in the IRAC and H bands (H-[4.5 \u03bcm] > 4.0), namely H-dropouts (Huang et al. 2011; Caputi et al. 2012; Wang et al. 2016a). Follow-up continuum observations with ALMA of a sample of 63 H-dropouts have yielded detections of 39 sources down to an 870 \u03bcm flux density of 0.6 mJy (Wang et al. 2019). They further suggested that the ALMA-detected H-dropouts are the bulk populations of massive (M\n\u22c6 \u2273 1010.3\nM\n\u2299) star-forming galaxies at z > 3 with the contribution to star formation rate density an order of magnitude higher than that of equivalently massive LBGs. To uncover the nature of H-dropouts and the critical role they play in the cosmic evolution of massive star-forming galaxies, we need to explore the fainter population that might have even fainter (sub)millimeter fluxes.","Citation Text":["Umehata et al. 2017"],"Functions Text":["Several ALMA blind surveys have been performed and allowed to detect and characterize the faint SMGs across cosmic times (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[830,849]],"Functions Start End":[[642,769]]}
{"Identifier":"2018ApJ...860...24P__Warmuth_2015_Instance_2","Paragraph":"Figure 13 shows the temporal evolution of the density, \u03c1, plasma flow velocity, vx, position of the wave crest, PosA, phase speed, vw, and magnetic field component in the z-direction, Bz, for the primary waves in every different case of initial amplitude, \u03c1IA. In Figure 13(a), we observe that the amplitude of the density remains approximately constant at their initial values until the time when the shock is formed and the density amplitude of the primary wave starts decreasing (see Vr\u0161nak & Luli\u0107 2000), i.e., at t \u2248 0.03 (blue), t \u2248 0.04 (red), and t \u2248 0.055 (green). For the case of \u03c1IA = 1.3 (magenta), a decrease of the amplitude of the primary can hardly be observed, as expected for low-amplitude wave (Warmuth 2015). One can see that the larger the initial amplitude, \u03c1IA, the stronger the decrease of the primary wave\u2019s amplitude, which is consistent with observations (Warmuth & Mann 2011; Muhr et al. 2014; Warmuth 2015). The amplitudes decrease to values of \u03c1 \u2248 1.6 (blue), \u03c1 \u2248 1.5 (red), and \u03c1 \u2248 1.4 (green) until the primary wave starts entering the CH. Due to the fact that the waves with larger initial amplitude enter the CH earlier than those with small initial amplitude, we can see in Figure 13(a) that the tracking of the parameters of the faster waves stops at an earlier time than the one for the slower waves. A similar behavior to the one of the density, \u03c1, can be observed for the plasma flow velocity, vx, in Figure 13(b) and the magnetic field component, Bz, in Figure 13(e). Here, the amplitudes decrease from vx = 0.75, Bz = 1.9 (for \u03c1IA = 1.9, blue), vx = 0.6, Bz = 1.7 (for \u03c1IA = 1.7, red), vx = 0.45, Bz = 1.5 (for \u03c1IA = 1.5, green), and vx = 0.27, Bz = 1.3 (for \u03c1IA = 1.3, magenta) to vx = 0.55, Bz = 1.6 (for \u03c1IA = 1.9, blue), vx = 0.46, Bz = 1.5 (for \u03c1IA = 1.7, red), vx = 0.36, Bz = 1.4 (for \u03c1IA = 1.5, green), and vx = 0.25, Bz = 1.25 (for \u03c1IA = 1.3, magenta). Figure 13(c) shows how the primary waves propagate in the positive x-direction. In all four cases of different initial amplitude, \u03c1IA, the phase speed decreases slighty (consistent with observations; see Warmuth et al. 2004 and Warmuth 2015) until the waves enter the CH at different times, i.e., the values for the phase speed start at vw \u2248 2.2 (for \u03c1IA = 1.9, blue), vw \u2248 1.9 (for \u03c1IA = 1.7, red), vw \u2248 1.7 (for \u03c1IA = 1.5, green), and vw \u2248 1.4 (for \u03c1IA = 1.3, magenta) and decrease to vw \u2248 1.5 (for \u03c1IA = 1.9, blue), vw \u2248 1.39 (for \u03c1IA = 1.7, red), vw \u2248 1.2 (for \u03c1IA = 1.5, green), and vw \u2248 1.13 (for \u03c1IA = 1.3, magenta).","Citation Text":["Warmuth 2015"],"Functions Text":["One can see that the larger the initial amplitude, \u03c1IA, the stronger the decrease of the primary wave\u2019s amplitude, which is consistent with observations"],"Functions Label":["Similarities"],"Citation Start End":[[922,934]],"Functions Start End":[[729,881]]}
{"Identifier":"2018ApJ...866...93L___2012c_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.","2012c"],"Functions Text":["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","but never clearly observed. These observations, for the first time, show a direct link between the cigar distribution and betatron cooling."],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[776,785],[793,798]],"Functions Start End":[[565,774],[825,964]]}
{"Identifier":"2018AandA...616L...2K__Frew_et_al._(2016)_Instance_1","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)"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[731,749]],"Functions Start End":[[576,730]]}
{"Identifier":"2019MNRAS.489.4669S__Bigiel_et_al._2010_Instance_2","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"],"Functions Text":["Similar deviations are observed in","outer parts of HSB spiral galaxies"],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[1862,1880]],"Functions Start End":[[1719,1753],[1826,1860]]}
{"Identifier":"2022ApJ...928...51B__Yadav_et_al._2016_Instance_1","Paragraph":"Relatively few such global convective dynamo studies have been conducted in the domain of M-dwarf stars. Early work by Browning (2008) considering lower mass FC M dwarfs found that the deep CZ could support very strong nonaxisymmetric fields, which strongly quenched the star\u2019s differential rotation. Later, more turbulent simulations of FC M-dwarf stars led to a number of interesting results. Yadav et al. (2015a) found strong, axisymmetric fields which were statistically steady in time and recovered many of the observed characteristics of M-dwarf surface fields. A somewhat slower rotating model (Yadav et al. 2015b) revealed that flux concentration by merging downflow lanes could lead to the formation of large, persistent high-latitude starspots in these stars. A still slower rotating model (Yadav et al. 2016) built large-scale, axisymmetric, cycling magnetic fields of somewhat lower amplitude which did not eliminate the star\u2019s differential rotation, reminiscent of the distributed \u03b1\u03a9-type dynamos prevalent in solar-like contexts. Brown et al. (2020) presented the first simulation of a stratified, rotating FC star whose computational domain extended to r = 0, finding a preference for hemispheric dynamo states. Similar models were recently explored by K\u00e4pyl\u00e4 (2021) for a broad range of rotation rates, which yielded a variety of results. At slow rotation, differential rotation was antisolar, and dynamo action dipole dominant. Moderate rotation rates yielded solar-like profiles, which were then magnetically quenched in the fastest rotators, yielding nonaxisymmetric dynamos. In Bice & Toomre (2020; hereafter BT20), we presented an exploration of the influence exerted by a tachocline in more massive, shell-convecting M-dwarf stars as a contributing factor to the break in observed magnetic activity across the tachocline divide. Our models produced a wide variety of field configurations, nearly all of which led to quenching of the differential rotation to a significant degree. We found that including a tachocline in models of early M-dwarf stars led to their surface fields being more favorable for rapid stellar spin-down, which may contribute to the formation of the tachocline divide.","Citation Text":["Yadav et al. 2016"],"Functions Text":["A still slower rotating model","built large-scale, axisymmetric, cycling magnetic fields of somewhat lower amplitude which did not eliminate the star\u2019s differential rotation, reminiscent of the distributed \u03b1\u03a9-type dynamos prevalent in solar-like contexts."],"Functions Label":["Background","Background"],"Citation Start End":[[801,818]],"Functions Start End":[[770,799],[820,1043]]}
{"Identifier":"2017ApJ...836L...4S__Scholer_&_Burgess_2007_Instance_2","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"],"Functions Text":["kinetic instabilities such as the Buneman and modified two-stream instability (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[1111,1133]],"Functions Start End":[[943,1027]]}
{"Identifier":"2016AandA...588A..25M__Weiss_&_Ferguson_2009_Instance_1","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"],"Functions Text":["Lastly, post-AGB stellar evolution models, computed with updated physics in a reduced mass range","show a strong disagreement with the previous grids."],"Functions Label":["Differences","Differences"],"Citation Start End":[[737,758]],"Functions Start End":[[624,720],[761,812]]}
{"Identifier":"2019AandA...628A.110M__Kryukova_et_al._(2012)_Instance_2","Paragraph":"Deriving the completeness limits of the WISE photometry is mandatory to assess the reliability of our catalogue of starless cores. We examined the histograms of the number of mid-infrared (MIR) sources versus magnitude; taking into account the effects of the cuts required to fulfil the criteria of Koenig et al. (2012), rough completeness limits are [3.6] ~ 14, [4.6] ~ 12, [12] ~ 9 and [22] ~ 7. These values are 1\u20133 mag brighter than the sensitivity limits quoted in the WISE Explanatory Supplement3 for the relevant sky region. Once converted into flux units and, for example, compared with the models of Class I and Class 0 sources of 0.5 M\u2299 by Whitney et al. (2004), it can be seen that the completeness limits at 3.6 and 4.6\u03bcm are faint enough to detect such objects taking into account a distance of 700 pc and a further foreground reddening up to AV = 20. Even in the worst case of edge-on discs, these objects would be detectable at 3.6 and 22\u03bcm. Furthermore, the completeness limit at 22 \u03bcm is faint enough to allow detection of Class I and Class 0 sources of even-lower-mass central objects. Alternatively, one can compute the bolometric luminosity following Kryukova et al. (2012). Starting from our completeness limit at 22 \u03bcm, after conservatively dereddening it by AV = 20, we assumed a spectral index \u03b3 = \u22122 (see Table 3 for definition) to compute the MIR luminosity from Eq. (6) of Kryukova et al. (2012). Equation (7) of Kryukova et al. (2012) then yields Lbol ~ 1.7\u20132.8 L\u2299, depending on whether the NIR flux is neglected (which may be the case) or extrapolated from \u03b3 = \u22122. A comparison with the birthline of Palla & Stahler (1993) indicates a mass of ~ 0.4\u20130.5 M\u2299 for the central protostar. For the sake of comparison, we can roughly estimate the completeness limit in central masses of the Herschel protostellar cores in Giannini et al. (2012) using their quoted completeness limit at 70 \u03bcm of 0.21 Jy and following Dunham et al. (2008). By using Eq. (2) of Dunham et al. (2008), scaled to a distance of 700 pc, we found that the flux density at 70 \u03bcm translates into a bolometric luminosity of the central (proto)star Lbol ~ 0.28 L\u2299 (we note that Dunham et al. 2008 indicate this luminosity as Lint). We highlight the fact that the 70 \u03bcm emission is in principle a more sensitive protostellar tracer than WISE. However, this contrasts with the much lower number of protostellar cores found by Giannini et al. (2012), which may be due to a poorer effective sensitivity because of their selection criteria.","Citation Text":["Kryukova et al. (2012)"],"Functions Text":["Starting from our completeness limit at 22 \u03bcm, after conservatively dereddening it by AV = 20, we assumed a spectral index \u03b3 = \u22122 (see Table 3 for definition) to compute the MIR luminosity from Eq. (6) of"],"Functions Label":["Uses"],"Citation Start End":[[1400,1422]],"Functions Start End":[[1195,1399]]}
{"Identifier":"2020ApJ...903....6M__Bueno_2010_Instance_1","Paragraph":"In this paper, we consider resonance scattering on a two-level atom with an infinitely sharp and unpolarized lower level. As for the frequency redistribution, we consider both CRD and angle-averaged PRD. In an unmagnetized one-dimensional spherically symmetric atmosphere, the polarized radiation field is axially symmetric and is described by \n\n\n\n\n\n. We also take into account the effects of radial velocity fields. We solve the concerned transfer equation in the CMF and in the nonrelativistic regime of velocity fields. To efficiently handle the sphericity effects, we solve the spherical transfer equation in the (p, z) coordinate system (Hummer & Rybicki 1971), which is also called the tangent-ray method. Here z is the distance along the tangent rays and p is impact parameter (see Figure 1 in Megha et al. 2019). Following Frisch (2007), we express the Stokes vector components in terms of their irreducible components. From here on we present all the basic equations in the irreducible basis (see Sampoorna & Trujillo Bueno 2010 for details). The CMF polarized PRD transfer equation for a spherically symmetric medium in (p, z) representation under the nonrelativistic limit is given by (see also Equation (17) of Megha et al. 2019)\n1\n\n\n\n\n\nwhere x denotes the nondimensional frequency and \n\n\n\n\n\n denotes the irreducible Stokes vector with \u201c+\u201d and \u201c\u2212\u201d referring to the outgoing and incoming rays, respectively. The radial optical depth is defined as d\u03c4r = \u2212\u03c7l(r)dr, where r is the radial distance and \u03c7l(r) is the line averaged absorption coefficient. In the CMF the monochromatic optical depth along the tangent ray is given by d\u03c4 = [\u03c6(x) + \u03b2c]d\u03c4r\/\u03bc, where \u03c6(x) is the line absorption profile function and \u03b2c = \u03c7c(r)\/\u03c7l(r) with \u03c7c(r) denoting the continuum absorption coefficient. The direction cosine of the tangent ray about the radius vector of the intercepting radial shell is given by \n\n\n\n\n\n. In Equation (1), \n\n\n\n\n\n denotes the CMF term, which has the form\n2\n\n\n\n\n\nwhere \u03c7(r, x) = \u03c7l(r)\u03c6(x) + \u03c7c(r), and\n3\n\n\n\n\n\nThe symbol V denotes the ratio of radial (vr) to the thermal (\n\n\n\n\n\n) velocities. In Equation (1), \n\n\n\n\n\n represents the irreducible CMF total source vector and is given by\n4\n\n\n\n\n\nwhere the line source vector is of the form\n5\n\n\n\n\n\nHere \u03f5 gives the probability of destruction of photons by inelastic collisions, B\u03bd0 is the Planck function at the line center frequency \u03bd0, and \n\n\n\n\n\n. The continuum is assumed to be unpolarized. Therefore, the continuum source vector is given by \n\n\n\n\n\n. The frequency-averaged PRD mean intensity vector is given by\n6\n\n\n\n\n\nwhere \n\n\n\n\n\n is the 2 \u00d7 2 nonmagnetic angle-averaged PRD matrix (Domke & Hubeny 1988; Bommier 1997), and\n7\n\n\n\n\n\nThe Rayleigh phase matrix \n\n\n\n\n\n in the irreducible basis is given in Appendix A of Frisch (2007). It is useful to rewrite Equation (7) as\n8\n\n\n\n\n\nwhere\n9\n\n\n\n\n\n\n","Citation Text":["Sampoorna & Trujillo Bueno 2010"],"Functions Text":["From here on we present all the basic equations in the irreducible basis (see","for details)."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1006,1037]],"Functions Start End":[[928,1005],[1038,1051]]}
{"Identifier":"2015MNRAS.450..763M__Rieke_et_al._2004_Instance_1","Paragraph":"In this work, we focus on the 56 most massive (M* \u2265 1011\u2009M\u2299) galaxies at 1.4 \u2264 z \u2264 2, 31 of which have spectroscopic redshifts. Two more objects entered the original sample, but they have been excluded from this study, because their WFC3\/HST images are not available, or too noisy to perform SB fitting, due to the proximity of saturated stars. The sample was culled from the K-selected (K(Vega)  22) multiband catalogue of Daddi et al. (2007a, hereafter D07), including data from all the available filters in GOODS-S, i.e. HST\/ACS optical, F435W (B), F606W (V), F775W (I), and F850LP(z), VLT\/ISAAC near-IR, J, H, K, and Spitzer\/Infrared Array Camera (IRAC), 3.6, 4.5, 5.8, and 8.0\u2009\u03bcm (for more details on the data sets see Giavalisco et al. 2004). The optical\/near-IR photometry was then complemented with the 24\u2009\u03bcm catalog (Daddi et al., in preparation), built as summarized in Section 2.1 from the Multi-Band Imaging Photometer for Spitzer (MIPS) images (Rieke et al. 2004). Far-IR fluxes were extracted (Daddi et al., in preparation) from the publicly released PACS 70\u2013160\u2009\u03bcm data from Herschel GOODS (Elbaz et al. 2011), and SPIRE 250\u2009\u03bcm data from Herschel Multi-tiered Extragalactic Survey (HerMES; Oliver et al. 2010). For objects without spectroscopic information, we used photometric redshifts from the public GOODS-MUSIC catalogue, which agree well with the spectroscopic ones (\u0394z\/(1 + z) \u2243 0.03) for galaxies at z  2 (cf. Grazian et al. 2006, 2007). Although the uncertainties on photometric redshifts could result in the inclusion in the sample of a few lower redshift contaminants, we avoided applying colour criteria to pre-select high-z objects, in favour of completeness. The HST\/WFC3\/F160W H-band image mosaic (>5\u03c3 point source sensitivity for H160  27.7, AB system) drizzled to a pixel scale of 0.06 arcsec, was exploited to study galaxy morphology in the optical rest frame with a very high resolution (FWHM \u223c 0.18\u2009arcsec \u2243 1.5 kpc). For more details on the observations and data reduction see Grogin et al. (2011) and Koekemoer et al. (2011).","Citation Text":["Rieke et al. 2004"],"Functions Text":["The optical\/near-IR photometry was then complemented with the 24\u2009\u03bcm catalog",", built as summarized in Section 2.1 from the Multi-Band Imaging Photometer for Spitzer (MIPS) images"],"Functions Label":["Uses","Uses"],"Citation Start End":[[958,975]],"Functions Start End":[[749,824],[855,956]]}
{"Identifier":"2018MNRAS.473.2000T__Noutsios_et_al._2011_Instance_2","Paragraph":"The launch of the Fermi Gamma-ray Space Telescope has spurred on the search for pulsars in \u03b3-rays (Grenier & Harding 2015), yielding over 2001 detections and triggering multiwavelength observations. While pulsars are common targets in the X-rays, they are very challenging targets in the optical and very few of them have been identified (see Mignani et al. 2016, and references therein). Here, we report on Large Binocular Telescope (LBT) observations of an isolated pulsar, PSR\u2009J2043+2740 (Taylor, Manchester & Lyne 1993), detected by both AGILE (Pellizzoni et al. 2009) and Fermi (Abdo et al. 2010; Noutsios et al. 2011). It was discovered as a radio pulsar (Ray et al. 1996) and later on as an X-ray source by XMM\u2013Newton (Becker et al. 2004), although X-ray pulsations have not yet been found. PSR\u2009J2043+2740 is one of the very few non-recycled pulsars older than 1\u2009Myr detected in \u03b3-rays, with a characteristic age \u03c4c = 1.2\u2009Myr, inferred from the values of its spin period Ps = 0.096\u2009s and its derivative $\\dot{P}_{\\rm s} = 1.27 \\times 10^{-15}$\u2009s\u2009s\u22121 (Ray et al. 1996). This also yields a rotational energy loss rate $\\dot{E}_{\\rm rot} = 5.6 \\times 10^{34}$\u2009erg\u2009s\u22121 and a surface dipolar magnetic field Bs = 3.54 \u00d7 1011 G.2 Although PSR\u2009J2043+2740 does not have a very large spin-down power compared to young (\u223c1\u201310 kyr) pulsars (\u223c1036\u20131038\u2009erg\u2009s\u22121), it is still a factor of 2 larger than that of middle aged \u03b3-ray pulsars (\u223c0.1\u20130.5\u2009Myr), such as Geminga, PSR\u2009B0656+14 and PSR\u2009B1055\u221252, all detected in the optical, thanks to their distances \u2272 500\u2009pc (Abdo et al. 2013). The distance to PSR\u2009J2043+2740 is uncertain owing to the lack of a radio parallax measurement. The radio dispersion measure (DM = 21.0\u2009\u00b1\u20090.1\u2009pc cm\u22123; Ray et al. 1996) gives a distance of 1.8\u2009\u00b1\u20090.3\u2009kpc from the NE2001 model of the Galactic-free electron density (Cordes & Lazio 2002). A slightly smaller distance (1.48\u2009kpc) is inferred from the model of Yao, Manchester & Wang (2017). The hydrogen column density towards the pulsar obtained from the X-ray spectral fits (NH \u2272 3.6 \u00d7 1020\u2009cm\u22122; Abdo et al. 2013) suggests a distance of a few hundred pc (He, Ng & Kaspi 2013), although these estimates depend on the model X-ray spectrum. Such a distance would make PSR\u2009J2043+2740 a viable target for deep optical observations, never carried out until now, and might be compatible with the debated association (Noutsios et al. 2011) with the Cygnus Loop supernova remnant (SNR) at $540^{+100}_{-80}$\u2009pc (Blair, Sankrit & Raymond 2005).","Citation Text":["Noutsios et al. 2011"],"Functions Text":["Such a distance would make PSR\u2009J2043+2740 a viable target for deep optical observations, never carried out until now, and might be compatible with the debated association"],"Functions Label":["Motivation"],"Citation Start End":[[2383,2403]],"Functions Start End":[[2211,2381]]}
{"Identifier":"2017AandA...607A...9M__Weideman_1994_Instance_1","Paragraph":"Having a line list that is as complete as possible is crucial for doing proper radiative transfer computations in high temperature atmospheric environments. However, computing the line opacities for a large number of lines can be computationally challenging. For example, the ExoMol line list of CH4 contains on the order of 1010 lines (Yurchenko & Tennyson 2014). Computing the exact pressure and temperature broadened Voigt profile for each of these lines is computationally extremely demanding. There are several approximate methods for computing Voigt profiles available in the literature (e.g. Huml\u00edcek 1982; Weideman 1994; Zaghloul & Ali 2012). Also, much effort is spend succesfully on making these approximate methods faster and more accurate (see e.g. Poppe & Wijers 1990; Letchworth & Benner 2007). These methods all focus on obtaining a given accuracy of the exact shape of the Voigt profile by applying mathematical approximations to decrease the computation time. While these methods still require significant computation time, they are now routinely applied to compute Voigt profiles in many applications. The fastest code able to compute Voigt profiles of large numbers of lines at this moment is the HELIOS-k code (Grimm & Heng 2015). This code is able to compute ~ 105 lines per second on a dedicated NVIDIA K20 GPU based machine. This implies that the computation of 1010 lines still requires on the order of one day for a single point in pressure temperature space. Usually a grid of pressure and temperature points is required. Thus, there is the need for an even faster method. In addition, we have to make sure that there are no systematic errors in the computations because a small systematic error in a single line can become large when computed for 1010 lines. Thus, we seek a method that is statistically exact, preserves the integrated opacities and the average shape of the lines, and computes the line profile accurately for the stronger lines. ","Citation Text":["Weideman 1994"],"Functions Text":["Computing the exact pressure and temperature broadened Voigt profile for each of these lines is computationally extremely demanding. There are several approximate methods for computing Voigt profiles available in the literature (e.g."],"Functions Label":["Background"],"Citation Start End":[[614,627]],"Functions Start End":[[365,598]]}
{"Identifier":"2019MNRAS.490.2155S__Blake_et_al._2016_Instance_1","Paragraph":"Modern optical imaging surveys measure the positions and ellipticities of millions of galaxies; from them, the galaxy overdensity field as well as the gravitational lensing shear field can be derived. The two-point auto and cross-correlations of these two fields are the two-point correlation functions of cosmic shear, galaxy\u2013galaxy lensing and galaxy clustering. A joint analysis of these correlation functions can break degeneracies between cosmological and nuisance parameters, leading to tighter cosmological constraints (Joachimi & Bridle 2010). Several earlier studies have indeed considered such joint analyses (Cacciato et al. 2013; Mandelbaum et al. 2013; More et al. 2015; Kwan et al. 2017; Nicola, Refregier & Amara 2017), albeit very few in a modified gravity context. Among the latter, Joudaki et al. (2018) recently performed a combined analysis of cosmic shear tomography, galaxy\u2013galaxy lensing tomography, and redshift space multipole power spectra using data from KiDS-450 (\u223c450\u2009deg2 of cosmic shear data from the KiDS survey) and two overlapping spectroscopic surveys, the 2-degree Field Lensing Survey4 (2dFLenS; Blake et al. 2016) and the Baryon Oscillation Spectroscopic Survey5 (BOSS; Dawson et al. 2013). They found that none of the extended cosmologies considered were simultaneously favoured in a model selection sense and able to resolve the discordance with Planck, except for an evolving dark energy component with a time-dependent w0 \u2212 wa equation of state. Amon et al. (2018) presented a measurement of EG, a statistic combining measurements of weak gravitational lensing, galaxy clustering, and redshift space distortions, proposed as a consistency test of General Relativity (Zhang et al. 2007). They determined the value of EG using data from the KiDS, 2dFLenS, BOSS, and Galaxy And Mass Assembly6 (GAMA; Driver et al. 2009, 2011; Liske et al. 2015) surveys; their results show that measurements of the EG statistic cannot be conducted as consistency checks of General Relativity until the aforementioned tension in cosmological parameters is resolved, and their EG measurements favour a lower matter density cosmology than the CMB. Recently, DES Collaboration (2019) presented a combined analysis of galaxy clustering and weak gravitational lensing from the first-year data of the Dark Energy Survey, targeting modifications of the metric potentials that would be a signal of modified gravity. They found that their constraints are compatible with a cosmological constant scenario.","Citation Text":["Blake et al. 2016"],"Functions Text":["Among the latter, Joudaki et al. (2018) recently performed a combined analysis of cosmic shear tomography, galaxy\u2013galaxy lensing tomography, and redshift space multipole power spectra using data from KiDS-450 (\u223c450\u2009deg2 of cosmic shear data from the KiDS survey) and two overlapping spectroscopic surveys, the 2-degree Field Lensing Survey4 (2dFLenS;"],"Functions Label":["Background"],"Citation Start End":[[1133,1150]],"Functions Start End":[[782,1132]]}
{"Identifier":"2016AandA...596A..59F__Sobotka_et_al._1993_Instance_1","Paragraph":"Sunspot light bridges can be formed by the fragmentation of the umbra in the decay phase or during the merging of different magnetized areas during the formation of the sunspot in complex active regions (Bray & Loughhead 1964; Garcia de La Rosa 1987). At the last stages of a sunspot, the photospheric-like conditions are recovered as a consequence of the splitting of the umbra, and a granulation pattern similar to that of the quiet Sun is found in light bridges (Vazquez 1973), although the light-bridge convection cells differ significantly from normal granules (Lagg et al. 2014). Understanding the decay of sunspots requires the study of the magnetic and dynamical structure of light bridges. Previous studies have classified light bridges based on their morphological properties. Light bridges separating two umbral regions are called strong light bridges (e.g., Sobotka et al. 1993, 1994; Jur\u010d\u00e1k et al. 2006; Rimmele 2008). Their brightness is similar to that of the penumbra, and they typically appear between two regions with the same polarity. Faint light bridges (e.g., Lites et al. 1991; Sobotka et al. 1993) are elongated bright structures that penetrate the umbra. They are composed of rows of bright grains with sizes comparable to umbral dots. Light bridges can exhibit different internal structures. Most of them show several segments along their length that resemble granules (Berger & Berdyugina 2003), while some light bridges are similar to the bright filaments seen in the penumbra (Lites et al. 2004). All types show a weakened magnetic field strength relative to the surrounding umbra and more strongly inclined field lines (e.g. Beckers & Schr\u00f6ter 1969; Lites et al. 1991; Leka 1997), which in some case even exhibit a polarity reversal (Lagg et al. 2014). Observations are consistent with a reduced field strength in the photosphere, with a magnetic canopy extending from either side of the light bridge and merging at the top (Jur\u010d\u00e1k et al. 2006; Lagg et al. 2014). The velocity field measured in light bridges shows evidence of their convective origin (Rimmele 1997; Hirzberger et al. 2002; Lagg et al. 2014). Recently, Schlichenmaier et al. (2016) have also found evidence that the umbral magnetic field is wrapped around light bridges. The authors defined a new type, the plateau light bridge, with Y-shaped dark canals that resemble penumbral grains and suggest the presence of inclined magnetic fields. Several dynamic phenomena have been detected at the chromosphere above light bridges, including recurrent plasma ejections visible in H\u03b1 (Roy 1973; Asai et al. 2001), Ca\u2009ii\u2009\u2009H (Shimizu et al. 2009), and other bands (AIA 1600 and 1700 \u00c5, IRIS 1330 and 1400 \u00c5, Toriumi et al. 2015), brightness enhancements in the TRACE 1600 \u00c5 band (Berger & Berdyugina 2003), or fan-shaped jets observed in H\u03b1 (Robustini et al. 2016). All these processes are considered to be caused by the interaction of the light bridge magnetic field with the surrounding atmosphere through magnetic reconnection events. ","Citation Text":["Sobotka et al. 1993","Sobotka et al. 1993"],"Functions Text":["Previous studies have classified light bridges based on their morphological properties. Light bridges separating two umbral regions are called strong light bridges (e.g.,","Their brightness is similar to that of the penumbra, and they typically appear between two regions with the same polarity. Faint light bridges (e.g.,","are elongated bright structures that penetrate the umbra. They are composed of rows of bright grains with sizes comparable to umbral dots."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[870,889],[1101,1120]],"Functions Start End":[[699,869],[932,1081],[1122,1260]]}
{"Identifier":"2021AandA...655A..12T__Tang_et_al._2017b_Instance_5","Paragraph":"Using the RADEX3 non local thermodynamic equilibrium (LTE) modeling program (van der Tak et al. 2007) with collisional rate coefficients from Wiesenfeld & Faure (2013), we modeled the relation between the gas kinetic temperature and the measured average of para-H2CO 0.5 \u00d7 [(322\u2013221 + 321\u2013220)\/303\u2013202] ratios, adopting a 2.73 K background temperature, an average observational linewidth of 4.0 km s\u22121, and column densities N(para-H2CO) = 2.7 \u00d7 1012 and 3.7 \u00d7 1012 cm\u22122 for N113 and N159W, respectively. The results are shown in Fig. 5. The values of the para-H2CO column density were obtained with APEX data (beam size ~30\u2033; Tang et al. 2017b), which cover similar regions. Different column densities of para-H2CO only weakly affect derived kinetic temperatures (see Fig. 3 in Tang et al. 2017b or Fig. 4 in Tang et al. 2018a; this was also shown in Fig. 13 and discussed in Sect. 4.3.1 of Mangum & Wootten 1993) as long as all lines are optically thin. Considering that the relation between the gas temperature and the para-H2CO line ratio may vary at different spatial densities (see Fig. 2 in Tang et al. 2017b), we modeled it at spatial densities 104, 105, and 106 cm\u22123 in Fig. 5. It appears that Tkin at n(H2) = 105 cm\u22123 is consistently lower than values at 104 and 106 cm\u22123 by \u227223% and \u227234%, respectively, for Tkin \u2272 100 K. Local thermodynamic equilibrium (LTE) is a good approximation for the H2CO level populations under optically thin and high-density conditions (Mangum & Wootten 1993; Tang et al. 2017a,b, 2018b). Following the method applied by Tang et al. (2017b) in their Eq. (2), we plot the relation between the LTE kinetic temperature, TLTE, and the para-H2CO (3\u20132) line ratio in Fig. 5. Apparently, TLTE agrees well with Tnon-LTE at volume densities n(H2) ~ 105 cm\u22123 as long as Tkin \u2272 100 K. Previous observations show that para-H2CO (3\u20132) is sensitive to gas temperature at density 105 cm\u22123 (Ginsburg et al. 2016; Immer et al. 2016; Tang et al. 2017b). The spatial density measured with para-H2CO (303\u2013202) and C18O (2\u20131) in N113 and N159W is n(H2) ~ 105 cm\u22123 on a size of ~30\u2033 (Tang et al. 2017b). Therefore, here we adopt 105 cm\u22123 as an averaged spatial gas density in the N113 and N159W regions.","Citation Text":["Tang et al. (2017b)"],"Functions Text":["Following the method applied by","in their Eq. (2), we plot the relation between the LTE kinetic temperature, TLTE, and the para-H2CO (3\u20132) line ratio in Fig. 5."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1558,1577]],"Functions Start End":[[1526,1557],[1578,1705]]}
{"Identifier":"2021ApJ...907...15H__Geiss_1982_Instance_1","Paragraph":"Figure 2 shows the ratios of the relative fluences (element\/Mg) measured in Genesis bulk SW (Figure 2(a)) and regime targets (Figure 2(b)) after normalization to the respective solar abundance ratios (Table 7 in Appendix A.5). The bulk SW data for K and Fe in silicon (marked by asterisks in Figure 2(a)) are the preliminary fluences from Rieck (2015), Rieck et al. (2016), and Burnett et al. (2017). If SW had the composition of the solar photosphere, the ordinate value in Figure 2 would be 1 for all elements. The bulk SW shows a fractionation of elemental abundances correlating with FIP, as previously observed by spacecraft (e.g., Geiss 1982; Bochsler 2009). However, the better accuracy and element coverage of Genesis data compared to in situ measurements allows further insights into this FIP-related fractionation. The low-FIP elements (Figure 2(a)) are consistent with a flat pattern, at least below 7 eV, as many pre-Genesis studies suggested. This consistency would make these SW data easier to apply to cosmochemistry, as they would suggest that low-FIP elements are unfractionated relative to each other and hence their SW abundance ratios equal photospheric ratios. However, Figure 2(a) is also consistent with a monotonic increase of the low-FIP elements with decreasing FIP, which could be viewed as a trend continuing that of the high-FIP elements between Ar and C discussed below. The Na abundance given here is about 20% higher than that of Burnett et al. (2017). Na fluences derived from backside depth profiling of diamond-like-C collectors are roughly a factor of 2 lower than ours, which are based on Si collectors (Rieck 2015; Rieck et al. 2016). Jurewicz et al. (2019) propose that diffusion of surface contamination Na might have enhanced fluences from Si collectors. Because the amounts of surface contamination is highly variable, the Na fluences from Si should show significant scatter; however, replicate analyses of bulk and regime samples agree to within 2%\u201311% (Table 4, Appendix A.1). More analyses are required to resolve this important discrepancy. A \u223c20% lower fractionation of K compared to that of Na, if real, would be surprising although within 2\u03c3, both values still overlap.","Citation Text":["Geiss 1982"],"Functions Text":["The bulk SW shows a fractionation of elemental abundances correlating with FIP, as previously observed by spacecraft (e.g.,","However, the better accuracy and element coverage of Genesis data compared to in situ measurements allows further insights into this FIP-related fractionation."],"Functions Label":["Similarities","Compare\/Contrast"],"Citation Start End":[[637,647]],"Functions Start End":[[513,636],[665,824]]}
{"Identifier":"2022MNRAS.517.4986G__Chambers_1999_Instance_1","Paragraph":"Wang et al. (2017), motivated by the fact that previous studies on high-eccentricity migration focused only on the total efficiency of HJs formed, analyzed the efficiency of each high-eccentricity mechanism, trying to understand which mechanism is the dominant one in HJ formation. They considered multiplanetary systems containing from two to five planets with equal masses (1MJ), initially in circular and near-coplanar orbits with a host star of $1\\, {\\rm M}_\\odot$ and $1\\, {\\rm R}_\\odot$. Moreover, various initial semimajor axis conditions were considered, depending on the initial mutual separation between the planets. Their numerical simulations were performed using the classical version of the Mercury code (Chambers 1999), which does not include general relativity (GR) effects and tidal interaction with the central star. They studied how the initial number of planets, the spatial separation between them, and the location of the inner planet influence the efficiencies of high-eccentricity mechanisms. As a result, they found that the Kozai\u2013Lidov mechanism plays the most important role in HJ production. The restriction made in the previous study, regarding not including the contribution of GR and realistic initial mass configurations (unequal masses), makes us wonder about the implications of these contributions on the efficiency of each high-eccentricity mechanism in the activation of the HJ production process. The main effect of GR is to cause the apsidal precession of planetary orbit, where the rate of precession is faster for planets close to the star and with eccentric orbits due to the term a5\/2(1 \u2212 e2) in the denominator, being a and e the planet\u2019s semimajor axis and eccentricity, respectively (Einstein 1916; Misner, Thorne & Wheeler 1973). Studies related to the influence of GR on the dynamic evolution of systems with more than one massive planet have been carried out, but only considering some particular systems (e.g. Adams & Laughlin 2006; Migaszewski & Go\u017adziewski 2009; Veras & Ford 2010; Zhang, Hamilton & Matsumura 2013; Marzari & Nagasawa 2019, 2020). Therefore, the aim of our study is to obtain a broader and more detailed description of the efficiency of each high-eccentricity mechanism in the activation of the HJ production process, including the contribution of GR and different initial planetary mass configurations. The analysis is made by solving the numerical simulation of the exact equations of motion, in the context of general N-body problem. Several initial conditions are considered, changing the initial mass and number of planets, the semimajor axis of the inner planet and the location of the other planets in the system.","Citation Text":["Chambers 1999"],"Functions Text":["Their numerical simulations were performed using the classical version of the Mercury code","which does not include general relativity (GR) effects and tidal interaction with the central star. They studied how the initial number of planets, the spatial separation between them, and the location of the inner planet influence the efficiencies of high-eccentricity mechanisms. As a result, they found that the Kozai\u2013Lidov mechanism plays the most important role in HJ production. The restriction made in the previous study, regarding not including the contribution of GR and realistic initial mass configurations (unequal masses), makes us wonder about the implications of these contributions on the efficiency of each high-eccentricity mechanism in the activation of the HJ production process."],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[719,732]],"Functions Start End":[[627,717],[735,1434]]}
{"Identifier":"2018AandA...620A..46M__Bergin_et_al._2000_Instance_1","Paragraph":"Astrochemical models have always dedicated special attention to molecular oxygen. With a cosmic abundance twice that of C, atomic O is the third most abundant element in space. In dense clouds, standard gas phase chemical models therefore suggest a comparable ratio of CO and O2 after times \u2265 3 \u00d7 105 yr (e.g. Woodall et al. 2007), where O2 is supposed to be formed especially via OH + O \u2192 O2 + H. The OH here can be formed by the dissociative recombination of H3O+, H3O+ +e\u2212 \u2192 OH + 2H. However, observations with the Submillimeter Wave Astronomy Satellite (SWAS) by Goldsmith et al. (2000) towards Orion and with Odin by Larsson et al. (2007) towards \u03c1 Oph A showed a significant difference between model predictions and measurements. The O2 abundances found were more than 100 times smaller than those predicted by models (Goldsmith et al. 2000). Better agreement with observations can be obtained if freeze-out of O atoms onto dust grains is taken into account in gas-grain chemical models (Bergin et al. 2000; Viti et al. 2001), with consequence surface production of H2O and O2, which may trap a significant fraction of oxygen, leaving only some atomic O and CO in the gas phase. Observations conducted by Liseau et al. (2012) led to a O2 column density of N(O2) = 5.5 \u00d7 1015 cm\u22122 with an upper limit of abundance of N(O2)\/N(H2) \u223c 5 \u00d7 10\u22128 in warm gas (T > 50 K) and to N(O2) = 6 \u00d7 1015 cm\u22122 with a little higher abundance in cold gas (T  30 K). Liseau et al. (2012) stated that detecting gas phase O2 might be so difficult because the O2 abundance is transient in \u03c1 Oph A and O2 is no longer detectable after \u223c2 \u00d7 105 yr. A relatively large amount of O2 has only been found with Herschel in Orion as reported by Goldsmith et al. (2011). This source is quite warm (\u2265180 K), leading to a grain temperature of \u2265100 K. At this temperature the grains are warm enough to desorb H2O ice and keep a large amount of oxygen with a big fraction in the form of O2 in the gas phase. Another explanation for the high O2 abundance found by Goldsmith et al. (2000) is that low-velocity C-shocks might be responsible for the increase of molecular oxygen in the gas phase.","Citation Text":["Bergin et al. 2000"],"Functions Text":["Better agreement with observations can be obtained if freeze-out of O atoms onto dust grains is taken into account in gas-grain chemical models",", with consequence surface production of H2O and O2, which may trap a significant fraction of oxygen, leaving only some atomic O and CO in the gas phase."],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[994,1012]],"Functions Start End":[[849,992],[1031,1184]]}
{"Identifier":"2015MNRAS.451.4328K__Coddington_1994_Instance_1","Paragraph":"We estimate the DPL parameters in equation (9) via Monte Carlo simulations. We compare the real structure function of an AGN to simulations computed from mock light curves generated using the DPL model of equation (9). We generate \u2018mock\u2019 light curves using the Timmer & K\u00f6nig (1995) method. To create a single mock light curve, pseudo-random numbers are generated using the Fast Mersenne Twister SFMT19937 generator seeded with hardware-generated random numbers (generated using Intel RDRAND instruction) to ensure that the random number sequences are free of artificial correlations (Coddington 1994) induced by poor random seed choices. At this intermediate stage, the mock light curve is oversampled by a factor of 10 i.e. we generate points at 10 \u00d7 the required cadence in order to avoid sampling issues. To include low-frequency modes that are not adequately characterized by the length of the observed light curve, the intermediate mock light curve is much longer than is ultimately required to make the final mock; mock light curves generated in this manner are capable of exhibiting low-frequency modes longer than the length of the observed data. Fast Fourier Transforms (FFTs) are most efficient for data sequences that are a power of 2; for this reason, we pick the intermediate (including the oversampling) to be of length 223. This results in the intermediate mock light curve being between 15 \u00d7 to 45 \u00d7 the length required for the final mock light curve depending on the actual length of the observed light curve. We pick a uniformly distributed random segment of the intermediate overly long light curve that has the same length as the observed light curve and generate another stream of uncorrelated Gaussian random deviates to simulate the white-noise properties of Kepler instrumentation noise. After adding this \u2018measurement noise\u2019, we set data points corresponding to the unobserved cadences in the real light curve to 0. This procedure creates a final mock light curve with identical sampling and noise properties to the real light curve. Fig. 8 shows the true light curve (orange) along with an example mock light curve (light green) for the Sy 1 AGN kplr006932990 illustrating what the mock light curves look like for the best-fitting DPL parameters for this object.","Citation Text":["Coddington 1994"],"Functions Text":["To create a single mock light curve, pseudo-random numbers are generated using the Fast Mersenne Twister SFMT19937 generator seeded with hardware-generated random numbers (generated using Intel RDRAND instruction) to ensure that the random number sequences are free of artificial correlations","induced by poor random seed choices."],"Functions Label":["Uses","Uses"],"Citation Start End":[[585,600]],"Functions Start End":[[291,583],[602,638]]}
{"Identifier":"2016MNRAS.461.1719C__Fu_et_al._2012_Instance_3","Paragraph":"HATLAS12-00 had already been identified as a candidate gravitationally lensed galaxy as a result of its high submm flux (i.e. F500 > 100 mJy), red Herschel colours and the lack of a bright optical or radio counterpart (see e.g. Negrello et al. 2010 for a discussion of the selection of lens candidates in H-ATLAS and other Herschel surveys). This source was therefore observed spectroscopically in the submm. A CO spectroscopic redshift of 3.26 was first suggested by Z-spec (Bradford et al. 2004) observations, then subsequently confirmed by observations by the CARMA interferometer (Leeuw et al., in preparation) and the Zpectrometer instrument (Harris et al. 2007) on the Greenbank Telescope (Harris et al. 2012; see also Fu et al. 2012). Additional followup observations in the optical, near-IR, submm and other wavelengths were targeted at the lensed z = 3.26 source and the foreground objects responsible for the lensing, resulting in detailed analyses of this lensing system by Fu et al. (2012) and Bussmann et al. (2013). Their conclusions are that the z = 3.26 source HATLAS12-00 is subject to gravitational lensing, with a magnification of 9.6 \u00b1 0.5 in both the submm continuum and CO, and 16.7 \u00b1 0.8 in the K\u2032 band, by two foreground galaxies, one at a spectroscopically determined redshift of 1.22, and another with photometry suggesting that it lies at a similar redshift. The submm photometry of HATLAS12-00 at 890\u2009\u03bcm acquired with the Submillimeter Array (SMA) as part of this programme (Fu et al. 2012; Bussmann et al. 2013) is fully consistent with the 870\u2009\u03bcm and 850\u2009\u03bcm fluxes derived for this source from the LABOCA and SCUBA2 data to be presented here. The spectral energy distribution (SED) of the lensed source, after correcting for the lensing amplification, is well matched by the optically thick SED model for Arp220 from Rangwala et al. (2011), with a lensing-corrected far-IR luminosity of 1.2 \u00b1 0.2 \u00d7 1013\u2009L\u2299, and an implied star formation rate of 1400 \u00b1 300 \u2009M\u2299 yr\u22121. In many ways the unlensed properties of this object match those of the broader population of bright submm selected galaxies first discovered by the SCUBA submm imager (see e.g. Chapman et al. 2005; Clements et al. 2008; Micha\u0142owski, Hjorth & Watson 2010). The unlensed 870\u2009\u03bcm flux of this object would be \u223c7.7 mJy.","Citation Text":["Fu et al. 2012"],"Functions Text":["The submm photometry of HATLAS12-00 at 890\u2009\u03bcm acquired with the Submillimeter Array (SMA) as part of this programme","is fully consistent with the 870\u2009\u03bcm and 850\u2009\u03bcm fluxes derived for this source from the LABOCA and SCUBA2 data to be presented here."],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[1503,1517]],"Functions Start End":[[1386,1501],[1541,1672]]}
{"Identifier":"2020AandA...635A..47H___2005_Instance_1","Paragraph":"Finally, the multiphase nature of galactic outflows implies that measurements of the outflow properties based on a single gas phase can lead to misleading conclusions (for a discussion, see e.g., Cicone et al. 2018b). Historically, systematic studies of galactic outflows in nearby and high-z galaxies have focused on the ionized gas \u2013 for example, as observed as broad wing emission in the spectra of the H\u03b1, [O\u202fIII] or Pa\u03b1 lines \u2013 (e.g., Heckman et al. 1990; Rupke & Veilleux 2013a; Woo et al. 2016; Harrison et al. 2016; F\u00f6rster Schreiber et al. 2019; Ramos Almeida et al. 2019) and the atomic phase \u2013 based on the Na D or Mg II lines in absorption (e.g., Heckman et al. 2000; Rupke et al. 2002, 2005; Weiner et al. 2009; Roberts-Borsani & Saintonge 2019). The molecular component of outflows, on the other hand, has been much more difficult to study. Great progress was made with the Herschel Space Observatory using the OH 119 \u03bcm line in absorption to study molecular outflows in Seyfert and luminous infrared galaxies (Fischer et al. 2010; Sturm et al. 2011; Veilleux et al. 2013; Bolatto et al. 2013; Spoon et al. 2013; George et al. 2014; Stone et al. 2016; Gonz\u00e1lez-Alfonso et al. 2017; Zhang et al. 2018). More recently, the advent of powerful millimeter-wave interferometers such as the Atacama Large Millimeter\/submillimeter Array (ALMA) and the NOrthern Extended Millimeter Array (NOEMA) are rapidly increasing the number of molecular outflows detected based on observations of the CO line (e.g., Combes et al. 2013; Sakamoto et al. 2014; Garc\u00eda-Burillo et al. 2014; Leroy et al. 2015; Feruglio et al. 2015; Morganti et al. 2015; Dasyra et al. 2016; Pereira-Santaella et al. 2016, 2018; Veilleux et al. 2017; Fluetsch et al. 2019; Lutz et al. 2020). At high-z, so far only a handful of large-scale, molecular outflows have been studied in QSOs (e.g., Cicone et al. 2015; Vayner et al. 2017; Feruglio et al. 2017; Carniani et al. 2017; Fan et al. 2018; Brusa et al. 2018), sub-millimeter galaxies (e.g., Spilker et al. 2018), and main-sequence, star-forming galaxies (e.g., Herrera-Camus et al. 2019).","Citation Text":["Rupke et al.","2005"],"Functions Text":["Historically, systematic studies of galactic outflows in nearby and high-z galaxies have focused on the ionized gas","and the atomic phase \u2013 based on the Na D or Mg II lines in absorption (e.g."],"Functions Label":["Background","Background"],"Citation Start End":[[680,692],[699,703]],"Functions Start End":[[218,333],[582,657]]}
{"Identifier":"2022AandA...663A.105P__Finoguenov_et_al._2010_Instance_1","Paragraph":"Cluster radio relics are usually found in the outskirts of merging galaxy clusters. They exhibit elongated morphologies and high degrees of polarisation above 1 GHz (up to 70%, Ensslin et al. 1998; Bonafede et al. 2014; Loi et al. 2019; de Gasperin et al. 2022). The resolved spectral index in radio relics shows a gradient: it steepens towards the cluster centre and flattens towards the outskirts. Their size can reach up to \u223c2 Mpc, and high-resolution observations have revealed filamentary structures within relics themselves (Di Gennaro et al. 2018; Rajpurohit et al. 2020, 2022a,b; de Gasperin et al. 2022). The Largest Linear Sizes (LLS) and radio powers of relics are correlated, as well as the integrated spectral index and the radio power (van Weeren et al. 2009b; Bonafede et al. 2012; de Gasperin et al. 2014). Relics trace ICM shock waves with relatively low (M\u2004 \u20043) Mach numbers (Finoguenov et al. 2010; Akamatsu et al. 2013; Shimwell et al. 2015; Botteon et al. 2016). The acceleration of electrons is believed to proceed via diffusive shock acceleration (DSA) in the ICM (Ensslin et al. 1998; Roettiger et al. 1999), in which particles scatter back and forth across the shock front gaining energy at every crossing. Nevertheless, this mechanism has been shown to be rather inefficient in accelerating electrons from the thermal pool (Vazza & Br\u00fcggen 2014; Vazza et al. 2016; Botteon et al. 2020a; Br\u00fcggen & Vazza 2020; see Brunetti & Jones 2014 for a review). Recently, it has been suggested that seed electrons could originate from the tails and lobes (driven by AGN outflows) of cluster radio galaxies (Bonafede et al. 2014; van Weeren et al. 2017; Stuardi et al. 2019), which alleviates the requirements of high acceleration efficiencies at cluster shocks (e.g., Markevitch et al. 2005; Kang et al. 2012, 2017; Botteon et al. 2016; Eckert et al. 2016). In some cases, double relics have been detected on opposite sides of the cluster centre (e.g., Rottgering et al. 1997; van Weeren et al. 2010, 2012b; Bonafede et al. 2012; de Gasperin et al. 2015a). In these clusters it is possible to constrain the merger history, providing important information about the formation processes of relics.","Citation Text":["Finoguenov et al. 2010"],"Functions Text":["Relics trace ICM shock waves with relatively low (M\u2004 \u20043) Mach numbers"],"Functions Label":["Background"],"Citation Start End":[[894,916]],"Functions Start End":[[823,892]]}
{"Identifier":"2019MNRAS.490.3061V__Myers_2009_Instance_1","Paragraph":"On the other hand, the line-shift-absence conundrum of ZP74 is easily explained through geometrical considerations. Essentially, the arguments leading to this conundrum assume that the collapse is roughly spherically symmetric and monolithic, so that the infall motions are coherent, and directed towards a single collapse centre, at the geometrical centre of the cloud. This assumption is inconsistent with our current understanding of the structure of MCs, which are known to be far from spherically symmetric, and instead consist of an intricate and inhomogeneous network of filaments and clumps within them (e.g. Bally et al. 1987; Feitzinger et al. 1987; Gutermuth et al. 2008; Juvela, Pelkonen & Porceddu 2009; Myers 2009; Andr\u00e9 et al. 2010; Henning et al. 2010; Men\u2019shchikov et al. 2010; Molinari et al. 2010; Arzoumanian et al. 2011; Busquet et al. 2013). The central clumps (\u2018hubs\u2019) appear to accrete from the filaments, while in turn the filaments seem to accrete radially from their surroundings (Schneider et al. 2010; Kirk et al. 2013; Peretto et al. 2014; Gong et al. 2018; Lu et al. 2018; Williams et al. 2018; Shimajiri et al. 2019). Thus, the geometry is far from being spherically symmetric, and therefore the accreting gas is not isotropically distributed around the collapse centres (the hubs). In addition, the velocity field is highly complex and chaotic (e.g. G\u00f3mez & V\u00e1zquez-Semadeni 2014; Zamora-Avil\u00e9s, Ballesteros-Paredes & Hartmann 2017; G\u00f3mez, V\u00e1zquez-Semadeni & Zamora-Avil\u00e9s 2018), so there is no reason to expect a systematic redshift of the absorption lines produced in the gas surrounding the hubs. Instead, the accretion flow is most directly observed as velocity-centroid gradients along the filaments, directed towards the hubs. Indeed, synthetic CO observations of simulations of the regime often show only marginal or no evidence for infall profiles, due to the chaotic motions and perhaps velocity crowding effects, although the line profiles do look similar to observed ones (e.g. Heitsch et al. 2009; Heiner, V\u00e1zquez-Semadeni & Ballesteros-Paredes 2015; Clarke et al. 2018). Nevertheless, recent dedicated searches for evidence of infall signatures in CO lines from GMCs have met with success. For example, Schneider et al. (2015) have found the classical combination of self-absorbed and blue-skewed optically thick lines (12CO (3\u20132)) together with centrally peaked optically thin (13CO (1\u20130)) lines, indicating collapse in the molecular gas surrounding IRDC G28.37+0.07, while Barnes et al. (2018) have measured shifts between the lines of 12CO (tracing gas in the outer parts of the cloud) and 13CO (tracing gas deeper into the cloud) in the pc-scale, massive clumps of the CHaMP survey, finding systematic velocity differentials between the two lines that imply an average mass accretion time-scale of \u223c16 Myr, consistent with the time-scales we discuss in this paper (cf. Section 7.1 and Fig. 13).","Citation Text":["Myers 2009"],"Functions Text":["This assumption is inconsistent with our current understanding of the structure of MCs, which are known to be far from spherically symmetric, and instead consist of an intricate and inhomogeneous network of filaments and clumps within them (e.g."],"Functions Label":["Differences"],"Citation Start End":[[717,727]],"Functions Start End":[[371,616]]}
{"Identifier":"2019AandA...630A..98S__Cenko_et_al._2012a_Instance_1","Paragraph":"The close approach of a star to a supermassive black hole (SMBH) can lead to the destruction of the stellar body in a process known as a tidal disruption event (TDE; Hills 1975). Gravitationally bound material returns to the black hole and is accreted, giving rise to a flare whose electromagnetic signature peaks in the extreme ultraviolet (EUV) band (Rees 1988; Ulmer 1999). These flares were first detected in the soft X-ray band by ROSAT (Komossa et al. 2004; Bade et al. 1996; Komossa & Greiner 1999; Komossa & Bade 1999), later by XMM-Newton and Chandra (Esquej et al. 2007; Saxton et al. 2012a, 2017; Maksym et al. 2010; Lin et al. 2015) (see review by Komossa 2017), and also in the UV band by GALEX (Gezari et al. 2006, 2008, 2009). In recent years, large-area optical surveys have detected candidate TDEs emitting at temperatures of a few \u00d7104 K (van Velzen et al. 2011; Cenko et al. 2012a; Gezari et al. 2012; Arcavi et al. 2014; Holoien et al. 2016), ostensibly too cool to be coming from an accretion disc (e.g. Bonning et al. 2007). This optical radiation has been interpreted as being due to reprocessing of the accretion radiation by an optically thick screen (Metzger & Stone 2016; Roth & Kasen 2018; Dai et al. 2018) or to emission from shocks (Piran et al. 2015). Super-Eddington accretion in the initial phase of the disruption causes a large-scale, radiation-driven outflow of material from the central engine (Strubbe & Quataert 2009), which Metzger & Stone (2016) showed would initially completely absorb the radiation from the central engine and convert it into optical\/UV photons with an effective temperature similar to that observed. In this model, the screen density is expected to drop after a few months to the point where the inner thermal radiation would become visible, with the delay time and ratio of X-ray to optical\/UV flux depending on the line of sight (Metzger & Stone 2016; Dai et al. 2018). Observationally, the evidence for differences in the X-ray and UV\/optical timescales is mixed. The X-rays may have lagged the UV by \u223c32 days in ASASSN-14li (Pasham et al. 2017), but broadly fell on the same timescale (Brown et al. 2017), as they did in 2MASX 0740-85 (Saxton et al. 2017). In SDSS J1201+30 the UV flux did not change, while the X-rays dropped by a factor of 100 (Saxton et al. 2012a), whereas in ASASSN-15oi the X-ray luminosity was quite low (LX\u2004\u223c\u20041041 ergs s\u22121) at the peak of the optical flare, but 200\u2013300 days later had increased to LX\u2004\u223c\u20041042ergs s\u22121 (Holoien et al. 2018; Gezari et al. 2017).","Citation Text":["Cenko et al. 2012a"],"Functions Text":["In recent years, large-area optical surveys have detected candidate TDEs emitting at temperatures of a few \u00d7104 K"],"Functions Label":["Background"],"Citation Start End":[[881,899]],"Functions Start End":[[742,855]]}
{"Identifier":"2017ApJ...835..154H__Lee_et_al._2014_Instance_1","Paragraph":"The simplest explanation for these metal abundances, i.e., that they reflect the yields of normal core-collapse SNe (averaged over the stellar initial mass function (IMF)), fails to predict anything like the observed stellar abundances of the extremely metal-poor stars ([Fe\/H]  \u22123). Of course, at these low metallicities, the number of progenitor SNe enriching the ISM may be small, so models typically allow for metal-poor or metal-free progenitor stars, and an arbitrary mix of progenitor stellar masses (i.e., assuming the abundances might come from just one or at most a few SNe with individual explosion and progenitor parameters fitted to the observations). However, even with these degrees of freedom, the models still often fail to explain the abundances of certain individual species at the order-of-magnitude level (see, e.g., Nomoto et al. 2006; Heger & Woosley 2010; Lee et al. 2014; Placco et al. 2015), although they undoubtedly explain many of the observed abundance ratios. For the lowest metallicity stars observed ([Fe\/H]  \u22124), and in particular for the CEMP stars, these remaining discrepancies have led to more \u201cexotic\u201d models with a number of free parameters, invoking a mix of normal\/faint SNe and hypernovae (with variable explosion energies of \u223c1051\u20131054 erg); jets, prior \u201cfailed explosions,\u201d and fallback episodes; rotation and adjustable mixing layers allowing for a tunable stellar abundance profile in the progenitor stars; and pollution of the stars via companions (e.g., Tominaga et al. 2007; Ishigaki et al. 2014; Takahashi et al. 2014; Abate et al. 2015). These additions can improve the agreement with observations; however, there is still no consistent theoretical scenario that simultaneously explains most of the observed stars, and even the best-fit models for many individual stars can still have order-of-magnitude discrepancies with certain outlier elements (see Tominaga et al. 2014; Frebel et al. 2015; Placco et al. 2015, and references therein).","Citation Text":["Lee et al. 2014"],"Functions Text":["However, even with these degrees of freedom, the models still often fail to explain the abundances of certain individual species at the order-of-magnitude level (see, e.g.,",", although they undoubtedly explain many of the observed abundance ratios."],"Functions Label":["Differences","Similarities"],"Citation Start End":[[880,895]],"Functions Start End":[[665,837],[916,990]]}
{"Identifier":"2015MNRAS.446.1140T__Murray_et_al._2010_Instance_2","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"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[843,861]],"Functions Start End":[[554,792]]}
{"Identifier":"2015ApJ...799..149J___2014_Instance_3","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)"],"Functions Text":["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"],"Functions Label":["Uses"],"Citation Start End":[[1193,1213]],"Functions Start End":[[1039,1192]]}
{"Identifier":"2021MNRAS.502.1312C__Recchia,_Blasi_&_Morlino_2016_Instance_1","Paragraph":"The dynamical importance of CRs is even more uncertain. This is in part because most early work on this question focused only on galactic conditions similar to those found locally (Jokipii 1976; Badhwar & Stephens 1977; Ghosh & Ptuskin 1983; Chevalier & Fransson 1984; Boulares & Cox 1990; Ko, Dougherty & McKenzie 1991; Ptuskin 2001), and\/or focused largely on the question of how and whether CRs can drive galactic winds originating in the ionized, low-density medium found several scale heights above galactic planes (Ipavich 1975; Breitschwerdt, McKenzie & Voelk 1991; Zirakashvili et al. 1996; Ptuskin et al. 1997; Zirakashvili & V\u00f6lk 2006; however, for an exception see Breitschwerdt, McKenzie & Voelk 1993). More recent numerical and analytic models have continued in this vein (e.g. Everett et al. 2008; Jubelgas et al. 2008; Samui, Subramanian & Srianand 2010; Wadepuhl & Springel 2011; Uhlig et al. 2012; Booth et al. 2013; Pakmor et al. 2016; Simpson et al. 2016; Recchia, Blasi & Morlino 2016, 2017; Ruszkowski, Yang & Zweibel 2017; Pfrommer et al. 2017; Buck et al. 2019), rather than address the question of whether CRs represent a significant contribution to the support of the neutral material that dominates the total mass budget and occupies at least $\\sim 50{{\\ \\rm per\\ cent}}$ of the volume (e.g. Dekel et al. 2019) near the mid-plane. Indeed, the vast majority of published simulations that include CR transport do not resolve the neutral phase or galactic scale heights (\u223c100 pc), and those that do (e.g. Hanasz et al. 2013; Salem & Bryan 2014; Salem, Bryan & Corlies 2016; Chan et al. 2019) generally assume that CR transport in the neutral ISM is identical to that in the ionized ISM (though see Farber et al. 2018), an assumption that is almost certainly incorrect (e.g. Zweibel 2017; Xu & Lazarian 2017; Krumholz et al. 2020). Only a few published models attempt to address the question of CR pressure support in the neutral ISM for non-Solar neighbourhood (mostly starburst or Galactic Centre) conditions (e.g. Thompson et al. 2006; Socrates, Davis & Ramirez-Ruiz 2008; Lacki, Thompson & Quataert 2010; Lacki et al. 2011; Crocker et al. 2011; Crocker 2012; Lacki 2013; Yoast-Hull, Gallagher & Zweibel 2016; Yoast-Hull & Murray 2019; Krumholz et al. 2020).","Citation Text":["Recchia, Blasi & Morlino 2016"],"Functions Text":["More recent numerical and analytic models have continued in this vein (e.g.","rather than address the question of whether CRs represent a significant contribution to the support of the neutral material that dominates the total mass budget and occupies at least $\\sim 50{{\\ \\rm per\\ cent}}$ of the volume","near the mid-plane."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[975,1004]],"Functions Start End":[[715,790],[1086,1311],[1337,1356]]}
{"Identifier":"2021ApJ...919..140S__Bartos_et_al._2017_Instance_1","Paragraph":"Resonant dynamical friction may have applications beyond the relaxation of IMBHs examined in this paper. It may affect all objects in stellar clusters much more massive than the individual constituents of the disk, if present, including massive stars, stellar mass black holes (BHs), or the center of mass of massive binaries. Furthermore, it is also expected to operate in any type of disk with a high number of particles, including active galactic nucleus (AGN) accretion disks. Previously, it has been argued that stars and BHs crossing the disk on low-inclination orbits get captured by Chandrasekhar dynamical friction into the disk (Bartos et al. 2017; Panamarev et al. 2018; Tagawa et al. 2020). An interesting implication is that, if BHs settle into the disk, they interact dynamically and form BH\u2013BH binaries efficiently, and frequent dynamical interactions and gas effects drive the BHs to merger, producing gravitational waves (GWs) detectable by LIGO, VIRGO, and KAGRA (McKernan et al. 2014, 2018; Bartos et al. 2017; Leigh et al. 2018; Yang et al. 2019; Tagawa et al. 2020, 2021; Samsing et al. 2020). Mergers are also facilitated by Lidov\u2013Kozai oscillations in anisotropic systems (Heisler & Tremaine 1986; Petrovich & Antonini 2017; Hamilton & Rafikov 2019). The results in this paper show that resonant dynamical friction may accelerate the capture of objects in the accretion disks by a factor proportional to the SMBH mass over the local disk mass for large orbital inclinations. Pressure and viscosity in a gaseous disk do not inhibit the orbit-averaged torque from the IMBH, which leads to realignment and the warping of the disk (Bregman & Alexander 2012). Thus, RDF may efficiently catalyze the alignment of the orbital planes of BHs even in low-luminosity AGN or Seyfert galaxies with relatively small disk masses, which may not be possible for Chandrasekhar dynamical friction. In fact, this mechanism extends the scope of the \u201cAGN merger channel\u201d for GW source populations even beyond low-luminosity AGN and Seyfert galaxies, as it may organize BHs into disks also in nonactive galaxies with nuclear stellar disks.","Citation Text":["Bartos et al. 2017"],"Functions Text":["Previously, it has been argued that stars and BHs crossing the disk on low-inclination orbits get captured by Chandrasekhar dynamical friction into the disk"],"Functions Label":["Background"],"Citation Start End":[[639,657]],"Functions Start End":[[481,637]]}
{"Identifier":"2015AandA...573A.138B__Deubner_(1975)_Instance_1","Paragraph":"Solar-like oscillations are excited stochastically by motions in the convective envelope of stars. For low-mass stars, they are found in all evolutionary states, between the main sequence and horizontal branch of helium-core burning stars (e.g. Leighton et al. 1962; Frandsen et al. 2002; Carrier et al. 2003; Hekker et al. 2009; Chaplin et al. 2011; Huber et al. 2011; Kallinger et al. 2012; Mosser et al. 2013) and were even detected in the M5 super giant \u03b1\u2009Her (Moravveji et al. 2013). These very characteristic oscillations lead to a nearly regular spaced comb-like pattern in the power spectrum. It was shown by Deubner (1975) that this ridge structure is governed by the degree of oscillation modes and resembles the predictions made by Ando & Osaki (1975). Empirically, the frequency patterns of solar-like oscillations were described through scaling relations by Kjeldsen & Bedding (1995). Since then, these relations have been tested and revised from large sample studies based on high-precision space photometry of red giants in clusters and in eclipsing binaries besides single stars (e.g. Corsaro et al. 2012b; Frandsen et al. 2013; Kallinger et al. 2010). Indications of non-radial oscillation modes were found in the variations of the absorption lines of bright red giants (Hekker et al. 2006; Hekker & Aerts 2010), but were firmly established in a large set of red giants observed with the CoRoT satellite (De Ridder et al. 2009). The identification of dipole mixed modes extended the sensitivity of the seismic analyses towards the core of evolved stars also (Beck et al. 2011; Bedding et al. 2011; Mosser et al. 2011). The analysis of solar-like oscillations enabled us to unravel many open questions on stellar structure and evolution, such as constraining the internal rotational gradient (Elsworth et al. 1995; Beck et al. 2012; Deheuvels et al. 2012) or determining the evolutionary status in terms of nuclear burning of a given red giant star (Bedding et al. 2011; Mosser et al. 2011). ","Citation Text":["Deubner (1975)"],"Functions Text":["It was shown by","that this ridge structure is governed by the degree of oscillation modes and resembles the predictions made by Ando & Osaki (1975)."],"Functions Label":["Background","Background"],"Citation Start End":[[617,631]],"Functions Start End":[[601,616],[632,763]]}
{"Identifier":"2017AandA...606A.113G__Carollo_et_al._(2013)_Instance_2","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)"],"Functions Text":["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."],"Functions Label":["Background"],"Citation Start End":[[571,592]],"Functions Start End":[[593,948]]}
{"Identifier":"2022MNRAS.515.2914E__Lada_&_Lada_2003_Instance_1","Paragraph":"Also addressing the evolution of a stellar cluster is work by Parker & Goodwin (2009) and Parker et al. (2009) in which N-body simulations are used to analyse the prevalence of planets susceptible to the Kozai effect from a binary companion, and the stability of binaries in a dense cluster environment respectively. Parker & Goodwin (2009) found that around $20{{\\ \\rm per\\ cent}}$ of all exoplanets should at one point in their lives be in the presence of a binary companion that has been sufficiently inclined by its cluster environment such that Kozai cycles can occur. This is an intriguing finding that bolsters the idea that a binary companion perturbed by the cluster can have an appreciable effect on the planets, with the caveat that the authors only examined the evolution of binary orbits, as their simulations lacked planetary bodies. Additionally, they particularly focused on very dense clusters similar to Orion, with a half-mass radius of only 0.1 pc. This is not typical for embedded clusters, which have typical half mass radii of \u223c0.8 pc (Lada & Lada 2003), and therefore would have less frequent interactions between cluster stars and a particular binary. Nevertheless, they showed that a dense cluster environment can significantly alter the architecture of a stellar binary which in turn can affect a protoplanetary disc or mature planet system. Parker et al. (2009) focused on similarly dense cluster environments, but instead explored the longevity of moderately wide (${\\sim }10^3 \\, \\text{au}$) to ultrawide (${\\gt}10^4 \\, \\text{au}$) binaries. They found that cluster environments strip away all ultra-wide binary companions, and that the denser clusters, with half mass radii of 0.1\u20130.2 pc do not retain any binaries with separations ${\\gt}10^3 \\, \\text{au}$. The less dense clusters, with half mass radii 0.4\u20130.8 pc, do retain some of these moderately wide binaries. The authors noted that as ultrawide binaries are often stripped in only a few cluster crossing times, these very separated binaries may form in isolation An alternate channel is that ultrawide binaries are formed during cluster dissolution (Kouwenhoven et al. 2010).","Citation Text":["Lada & Lada 2003"],"Functions Text":["This is not typical for embedded clusters, which have typical half mass radii of \u223c0.8 pc",", and therefore would have less frequent interactions between cluster stars and a particular binary."],"Functions Label":["Differences","Differences"],"Citation Start End":[[1059,1075]],"Functions Start End":[[969,1057],[1076,1176]]}
{"Identifier":"2020MNRAS.495.4508E__Heinke_et_al._2014_Instance_1","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"],"Functions Text":["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."],"Functions Label":["Background"],"Citation Start End":[[764,782]],"Functions Start End":[[505,690]]}
{"Identifier":"2022MNRAS.511.5797M__Liu_&_Lai_2019_Instance_1","Paragraph":"A variety of formation channels have been proposed for binary black holes (BBHs; see e.g. Mapelli 2021 for a recent review): BBH mergers can be the outcome of isolated binary evolution via common envelope (Tutukov & Yungelson 1973; Bethe & Brown 1998; Portegies Zwart & Yungelson 1998; Belczynski, Kalogera & Bulik 2002; Belczynski et al. 2008, 2016a; Dvorkin et al. 2016, 2018; Eldridge & Stanway 2016; Mapelli et al. 2017, 2019; Stevenson, Berry & Mandel 2017; Kruckow et al. 2018; Spera et al. 2019; Belczynski et al. 2020; Klencki et al. 2021; Olejak, Belczynski & Ivanova 2021; Tanikawa et al. 2021a), stable mass transfer (Giacobbo, Mapelli & Spera 2018; Neijssel et al. 2019; Bavera et al. 2021; Gallegos-Garcia et al. 2021; Shao & Li 2021), or chemically homogeneous evolution (de Mink & Mandel 2016; Mandel & de Mink 2016; Marchant et al. 2016; du Buisson et al. 2020; Riley et al. 2021). Alternatively, BBHs can form dynamically in triples (e.g. Antonini, Toonen & Hamers 2017; Silsbee & Tremaine 2017; Fragione & Silk 2020; Arca Sedda, Li & Kocsis 2021a; Vigna-G\u00f3mez et al. 2021), multiples (e.g. Fragione & Kocsis 2019; Liu & Lai 2019, 2021; Hamers & Safarzadeh 2020), young star clusters (YSCs; Banerjee, Baumgardt & Kroupa 2010; Mapelli 2016; Banerjee 2017, 2021; Di Carlo et al. 2019, 2020a; Kumamoto, Fujii & Tanikawa 2019, 2020), globular clusters (GCs; Portegies Zwart & McMillan 2000; Tanikawa 2013; Samsing, MacLeod & Ramirez-Ruiz 2014; Rodriguez, Chatterjee & Rasio 2016; Askar et al. 2017; Fragione & Kocsis 2018; Hong et al. 2018; Choksi et al. 2019; Kamlah et al. 2022), and nuclear star clusters (NSCs; Antonini & Rasio 2016; Petrovich & Antonini 2017; Antonini, Gieles & Gualandris 2019; Arca Sedda 2020; Arca Sedda et al. 2020; Fragione, Loeb & Rasio 2020). Furthermore, gas torques in active galactic nucleus (AGN) discs trigger the formation of BBHs and speed up their mergers (e.g. Bartos et al. 2017; Stone, Metzger & Haiman 2017; McKernan et al. 2018; Yang et al. 2019; Ishibashi & Gr\u00f6bner 2020; Tagawa, Haiman & Kocsis 2020). Finally, primordial black holes (BHs), born from gravitational collapses in the early Universe, might also pair up and merge via gravitational wave (GW) emission (e.g. Carr & Hawking 1974; Carr, K\u00fchnel & Sandstad 2016; Sasaki et al. 2016; Ali-Ha\u00efmoud, Kovetz & Kamionkowski 2017; Clesse & Garc\u00eda-Bellido 2017; De Luca et al. 2021).","Citation Text":["Liu & Lai 2019"],"Functions Text":["Alternatively, BBHs can form dynamically in","multiples (e.g."],"Functions Label":["Background","Background"],"Citation Start End":[[1132,1146]],"Functions Start End":[[898,941],[1092,1107]]}
{"Identifier":"2015AandA...584A..76S__Bernstein_et_al._(1995)_Instance_2","Paragraph":"The only gas-phase process, which is predicted to efficiently lead to products, is the process involving ionized methanimine. According to the model by Vuitton et al. (2007), the amount in the upper atmosphere of Titan of ionized methanimine is small, but not negligible. The products of the reaction CH2NH + CH2NH+ all have a mass-to-charge ratio of 57, where an important contribution is given by the abundant carbocation C4H\\hbox{$_{9}^{+}$}+9. In this condition, it is difficult to see if any of these species is present in small amounts in the ionosphere of Titan. Interestingly, we have also investigated the further reaction of 1+ with a third molecule of methanimine. The formation of the species (CH2NH)\\hbox{$_{3}^{+}$}+3, starting from two molecules of methanimine and the ionic species CH2NH+, is a strongly exothermic reaction, being \\hbox{$\\Delta H_{0}^{\\circ} = -305$}\u0394H0\u25e6=\u2212305 kJ\/mol at CCSD(T) level. The optimized structure of (CH2NH)\\hbox{$_{3}^{+}$}+3 is shown in Fig. 10. We can conclude, therefore, that polymerization of methanimine in the gas-phase at low temperatures may well be initiated by the presence of an ionized molecule. As for the experiment on ice by Bernstein et al. (1995), it is well known that reaction barriers possibly present in gas-phase reactions are not significantly reduced when moving to ice-mediated reactions (see, for instance, Rimola et al. 2014). Normally, the tunneling effect is invoked to explain the observed reactivity, but in this case the reaction barriers are so high that it is difficult to think that a reaction sequence starting with dimerization of neutral methanimine molecules can account for the observed formation of hexamethylenetetramine or polymethylenimine. The reaction must start by involving a radical or an ionized species and not two neutral closed shell molecules. This was already noted by Vinogradoff et al. (2012), who suggested that the reaction between two neutral methanimine molecules is mediated by formic acid, which acts as a proton donor, and by Cottin et al. (2001), who irradiated mixed ice with protons. Notably, Vinogradoff et al. (2013) failed to see hexamethylenetetramine and polymethylenimine formation starting from pure CH2=NH ice. Since in the experiment by Bernstein et al. (1995) the ice was irradiated by VUV photons at the Lyman alpha wavelength, and we now know that methanimine can be efficiently ionized by those photons, we can also argue that ionization of several methanimine molecules can instead trigger the process in cold interstellar ices. This statement is in line with what is known for gas-phase polymerization of olefin (El-Shall 2008; Cottin et al. 2001). Alternatively, external strong energy sources that can induce local nonequilibrium conditions, could promote neutral-neutral dimerization. For instance, in a study by Zhou et al. (2010), in which acetylene ices were irradiated with energetic electrons, electronically excited acetylene molecules were invoked to account for the experimental observation of benzene formation. ","Citation Text":["Bernstein et al. (1995)"],"Functions Text":["Since in the experiment by","the ice was irradiated by VUV photons at the Lyman alpha wavelength, and we now know that methanimine can be efficiently ionized by those photons, we can also argue that ionization of several methanimine molecules can instead trigger the process in cold interstellar ices."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[2260,2283]],"Functions Start End":[[2233,2259],[2284,2556]]}
{"Identifier":"2020MNRAS.498.1319M__MacGregor_et_al._2017_Instance_1","Paragraph":"While most systems with exo-Kuiper belts do not have known planetary mass companions, in a few of these it has been possible to directly image one, thus enabling the study of planet\u2013disc interactions in more detail. There are well-known examples such as \u03b2 Pic with a massive planet possibly warping the disc (Mouillet et al. 1997; Lagrange et al. 2012, 2019; Matr\u00e0 et al. 2019); HR 8799 with four giant planets creating a scattered disc and possibly replenishing its warm dust closer in (e.g. Marois et al. 2010; Booth et al. 2016; Zurlo et al. 2016; Read et al. 2018; Wilner et al. 2018; Geiler et al. 2019; Faramaz et al., in preparation); HD 95086\u2019s axisymmetric disc implying a low eccentricity of its 4MJup planet (Rameau et al. 2016; Su et al. 2017); and Fomalhaut having a narrow and eccentric planetesimal belt (Kalas, Graham & Clampin 2005; Acke et al. 2012; Boley et al. 2012; MacGregor et al. 2017), implying that its candidate companion on an eccentric orbit has a low mass (\u223cEarth or super-Earth) and is not sculpting the belt (Quillen 2006; Kalas et al. 2008; Chiang et al. 2009; Beust et al. 2014; Faramaz et al. 2015), or is not a compact object but rather the dusty aftermath of a recent planetesimal collision (Gaspar & Rieke 2020). Some exo-Kuiper belt host systems even have companions in the brown dwarf (BD) or low stellar mass regime, suggesting that their likely formation through gravitational instability (Boss 1997, 2003, 2011; Vorobyov 2013) is compatible with the formation of massive Kuiper belt analogues, e.g. HR 2562 (Konopacky et al. 2016), HD 193571 (Musso Barcucci et al. 2019), HD 92536 (Launhardt et al. 2020), and HD 206893 (Milli et al. 2017). The last one is the subject of this paper. For even more massive companions, Yelverton et al. (2019) found a significant lower detection rate of debris discs around binaries, with no discs detected in binaries with separations between 25 and 135 au (comparable to typical debris disc radii; Matr\u00e0 et al. 2018b). This is likely due to dynamical perturbation inhibiting planetesimal formation or clearing any debris disc formed near those separations.","Citation Text":["MacGregor et al. 2017"],"Functions Text":["While most systems with exo-Kuiper belts do not have known planetary mass companions, in a few of these it has been possible to directly image one, thus enabling the study of planet\u2013disc interactions in more detail. There are well-known examples such as","and Fomalhaut having a narrow and eccentric planetesimal belt"],"Functions Label":["Background","Background"],"Citation Start End":[[887,908]],"Functions Start End":[[0,253],[757,818]]}
{"Identifier":"2017AandA...599A..97H__Gratton_et_al._2012_Instance_1","Paragraph":"Amongst the oldest stellar systems known to exist in the Milky Way (MW) are metal-poor globular clusters (GCs). These accumulations of stars do not seem to have undergone substantial star formation for extended periods. Given the limited quality of the available data, for a long time color-magnitude diagrams (CMDs) of GCs appeared to be narrow and could be readily described by a single isochrone. These observations have justified the establishment of the long-lasting paradigm that considers CGs as prime examples of simple stellar populations (SSPs), that is, the results of very short bursts of star formation in their natal clouds. However, improved photometric precision indicates the presence of sub-populations in the cluster CMDs that are inconsistent with the SSP assumption, for a number of luminous GCs in a variety of bandpasses. Thus, early detections of chemical abundance variations (e.g., Cohen 1978) could be more easily explained in a scenario involving several populations. Moreover, in recent years evidence has grown supporting the statement that GCs are generally composed of two or three chemically distinct populations. These subpopulations are separated by a few tens to hundreds of Myr in age and show vastly varying abundances of light elements such as C, N, O, Na, Mg, and Al (see, e.g., Carretta et al. 2009b; Gratton et al. 2012, and references therein). Theoretical considerations (see, e.g., D\u2019Ercole et al. 2008, 2011) imply that GCs could have lost the majority of the initial stellar content of the first population, which consequently should have ended up in the Galactic halo. In fact, numerous studies found metal-poor GCs to be consistent with the abundance trends of the MW halo at equally low metal content (e.g., Pritzl et al. 2005; Koch et al. 2009; Koch & McWilliam 2014; Villanova et al. 2016). We address this scenario by adding NGC 6426 to the short list of metal-poor clusters with available information on detailed chemical abundances. There are only two GCs in the Harris catalog (Harris 1996, 2010 edition) more metal poor than NGC 6426. At 12.9 \u00b1 1.0 Gyr, the cluster is the oldest in the age compilation by Salaris & Weiss (2002). At a galactocentric distance of Rgc = 14.4 kpc and a galactic latitude of 16.23\u00b0 it is located in the transition region between inner and outer halo. Previous studies found consistent [Fe\/H]1 values: \u22122.20 \u00b1 0.17 dex (Zinn & West 1984), \u22122.33 \u00b1 0.15 (Hatzidimitriou et al. 1999), and \u22122.39 \u00b1 0.04 dex (Dias et al. 2015). The latter value originates from the very first spectroscopic analysis of NGC 6426 at low resolution, which also stated [Mg\/Fe] = 0.38 \u00b1 0.06 dex. To date, there has been no study further addressing the detailed metal content of this cluster. ","Citation Text":["Gratton et al. 2012"],"Functions Text":["Moreover, in recent years evidence has grown supporting the statement that GCs are generally composed of two or three chemically distinct populations. These subpopulations are separated by a few tens to hundreds of Myr in age and show vastly varying abundances of light elements such as C, N, O, Na, Mg, and Al (see, e.g.,"],"Functions Label":["Background"],"Citation Start End":[[1342,1361]],"Functions Start End":[[996,1318]]}
{"Identifier":"2022MNRAS.516.6194C__Tran_et_al._2001_Instance_1","Paragraph":"In this paper, we focus on the bright elliptical galaxy NGC 5813. This is the central dominant member of a subgroup, hereafter referred to as the NGC 5813 group with an extensive diffuse X-ray emission. While part of the well-isolated NGC 5846 group ($z$ = 0.006578; e.g. Mahdavi, Trentham & Tully 2005a; Machacek et al. 2011) with a projected separation of \u223c740\u2009kpc, the two show no signs of an interaction between them. NGC 5813 itself has been assumed to be dynamically old (Emsellem et al. 2007) with no evidence of a recent major merger in its history, as indicated by the lack of any significant disturbances to its dusty circumnuclear disc (Tran et al. 2001). Regarding its atmosphere, the NGC 5813 group has a rather well defined regular morphology, consisting of three collinear pairs of cavities and associated shock fronts that are products of three distinct outburst events in the central AGN\u2019s history (Randall et al. 2011). Such cavities often appear as a result of AGN-ICM or IGrM interaction with the radio lobes produced by the AGN, displacing the surrounding X-ray gas (e.g. McNamara & Nulsen 2012; Barai et al. 2016; Yang, Gaspari & Marlow 2019; Gastaldello et al. 2021). It has been estimated that an energy of 1.5 \u00d7 1056 and $4\\times 10^{57}\\,{\\rm erg}$ has been released by the most and second most recent outbursts, respectively, suggesting that the most recent outburst might still be ongoing. These features are correlated with the group\u2019s observed radio emission (Giacintucci et al. 2011), and can serve as clear indicators of AGN feedback being present. Previous studies of the group have also found moderate turbulent velocities of \u223c175\u2009km\u2009s\u22121 and a 3D Mach number of the order of 0.4, using resonant scattering (e.g. de Plaa et al. 2012; Ogorzalek et al. 2017), further indicating the presence of such a mechanism. These characteristics make NGC 5813, a very promising candidate for the study of the effects AGN feedback can have on the distribution of elements in a low-mass system. Additionally, this target has some of the deepest Chandra data of any galaxy group and, similar to M49 has been found to have an explicit anticorrelation between its metal abundance and the location of its radio lobes (Randall et al. 2015). That work already suggests that this result is dependent on the Fe-bias; employing a two-temperature fit brings the abundance in the region of the extended radio lobes in better agreement to their surrounding medium, although the central region appears to remain under-enriched.","Citation Text":["Tran et al. 2001"],"Functions Text":["NGC 5813 itself has been assumed to be dynamically old","with no evidence of a recent major merger in its history, as indicated by the lack of any significant disturbances to its dusty circumnuclear disc"],"Functions Label":["Background","Background"],"Citation Start End":[[648,664]],"Functions Start End":[[422,476],[500,646]]}
{"Identifier":"2018AandA...620A..80M__Coutens_et_al._2018_Instance_2","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"],"Functions Text":["These values are in the range of those observed towards the molecular clouds in the Galactic centre but lower than in Orion KL"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1066,1085]],"Functions Start End":[[938,1064]]}
{"Identifier":"2020MNRAS.497L..56Y__Verde,_Treu_&_Riess_2019_Instance_1","Paragraph":"Over the past decade, direct measurements of the Hubble constant have achieved few percent in precision (Freedman 2017). Among the conducted measurements, the Supernova H0 for the Equation of State (SH0ES) team challenged the well-believed Hubble constant (H0) value inferred from the Planck cosmic microwave background (CMB) measurements assuming a flat Lambda cold dark matter (\u039bCDM) model. In detail, the latest result from SH0ES is H0 = 74.03 \u00b1 1.42 km s\u22121 Mpc\u22121 (Riess et al. 2019), which differs from thePlanck result H0 = 67.4 \u00b1 0.5 km s\u22121 Mpc\u22121 by 4.4\u03c3 (Planck Collaboration VI 2018). Recently, the Carnegie-Chicago Hubble Program (CCHP) also presented a new and independent determination of H0 parameter based on a calibration of the Tip of the Red Giant Branch (TRGB) applied to Type Ia supernovae (SNIa) (Freedman et al. 2019, 2020). They find a value of H0 = 69.6 \u00b1 2.5 km s\u22121Mpc\u22121, which is in the middle and consistent with the SH0ES and Planck values. A different analysis and calibration of the TRGB method was performed by the SH0ES team (Reid, Pesce & Riess 2019; Yuan et al. 2019) resulting in a slightly higher H0 value. To come to a robust conclusion, independent H0 probes with accuracy better than $2{{\\ \\rm per\\ cent}}$ are crucial (Verde, Treu & Riess 2019). Among the possible independent probes, the time-delay strong lensing (TDSL) measurements, such as from the H0 Lenses in COSMOGRAIL\u2019s Wellspring (H0LiCOW) collaboration (Bonvin et al. 2017; Suyu et al. 2017; Birrer et al. 2019; Wong et al. 2019), are the most precise to date. The latest constraint from a joint analysis of six gravitationally lensed quasars with measured time delays (Wong et al. 2019) indicates for a flat \u039bCDM, $H_0=73.3^{+1.7}_{-1.8}$ km s\u22121 Mpc\u22121, a $2.4{{\\ \\rm per\\ cent}}$ precision measurement, which are in agreement with local measurements from SNIa, but in 3.1\u03c3 tension with CMB. The forecasts of future 40 TDSL measurements suggest that the H0 would be constrained at $\\mathcal {O}(1){{\\ \\rm per\\ cent}}$ level (Shajib, Treu & Agnello 2018; Yildirim, Suyu & Halkola 2019). A more optimistic forecast for dark energy studies can be found in Shiralilou et al. (2019). On the other hand, one of the main obstacles for the lensing mass modelling, or to determine precise H0 value, is the mass-sheet degeneracy (Schneider & Sluse 2014; Xu et al. 2016). These issues in the H0LiCOW analysis have been frequently discussed in the literature (Sonnenfeld 2018; Kochanek 2019; Pandey, Raveri & Jain 2019). No direct evidence of bias or errors is found from a comparison of self-consistency among the individual lenses  (Millon et al. 2019; Liao et al. 2020). Considering the fact that both the Planck and H0LiCOW\u2019s H0 values are based on general relativity (GR) plus \u039bCDM model, it inspires us to question the concordance cosmology model and investigate the modified gravity (MG).","Citation Text":["Verde, Treu & Riess 2019"],"Functions Text":["To come to a robust conclusion, independent H0 probes with accuracy better than $2{{\\ \\rm per\\ cent}}$ are crucial"],"Functions Label":["Motivation"],"Citation Start End":[[1257,1281]],"Functions Start End":[[1141,1255]]}
{"Identifier":"2022MNRAS.517.2801W__Gallo_et_al._2014_Instance_2","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"],"Functions Text":["The underlying cause of the split tracks (see","for a clustering analysis into the statistical robustness of this split) for BHXBs is not known"],"Functions Label":["Background","Background"],"Citation Start End":[[1341,1358]],"Functions Start End":[[1266,1311],[1400,1495]]}
{"Identifier":"2016ApJ...818..141V__Sellwood_2014b_Instance_1","Paragraph":"Historically, the study of orbits in potentials has focused on periodic orbits. In systems like disk galaxies small perturbations to closed periodic orbits (e.g., the epicyclic and vertical perturbations of circular orbits) provided a good analytic description of most orbits. Self-consistent distribution functions are thought to be \u201cparented\u201d by stable periodic orbits (Arnold 1978). Early works (e.g., Contopoulos & Papayannopoulos 1980) identified and characterized the stability properties of the periodic orbit families in rapidly rotating bars. The most important periodic families in two-dimensional bars were identified as the prograde x1 family, which is elongated along the major axis of the bar, and the prograde stable x2 and unstable x3 families, which are elongated perpendicular to the bar (primarily found at small radii). The retrograde x4 (stable) orbit family is also elongated perpendicular to the bar at small radii, but becomes rounder as it extends to larger radii (for detailed description of orbit families and how they are identified see Contopoulos & Grosbol 1989; Sellwood & Wilkinson 1993; Binney & Tremaine 2008; Sellwood 2014b). In the frame of reference rotating with the bar, all of these families are characterized by a 1:2 resonance between the tangential oscillation frequency (\u03a9\u03d5) and the radial or epicyclic frequency (\u03a9R). Indeed, studies of orbits in 2D N-body bars largely confirmed the picture arising from the study of periodic orbits and showed that many regular orbits elongated along the bar were parented by x1 orbits, a small fraction were parented by retrograde x4 orbits (Sparke & Sellwood 1987), and none were parented by prograde x2 orbits. The realization that bars can also undergo buckling instabilities (Combes & Sanders 1981; Raha et al. 1991), which makes them develop substantial vertical thickness and peanut-shaped morphologies, led to the study of periodic orbits in three-dimensional bars (Pfenniger 1984; Martinet & de Zeeuw 1988; Pfenniger & Friedli 1991; Skokos et al. 2002a, 2002b). It was shown that the appearance of specific morphological features in images of bars, such as the X-shape and peanut features seen in edge-on bars and the boxy\/rectangular isophotes and \u201cansae\u201d of face-on bars, could be explained by orbits trapped around specific periodic orbit families (Patsis et al. 2002, 2003, 2010). The introduction of a third dimension did not drastically change the picture of the nature of periodic orbits, and it was found that 3D bars are composed primarily of vertical bifurcations (resonances) of the x1 family and a few additional families (e.g., Pfenniger & Friedli 1991; Skokos et al. 2002a, 2002b). Most studies of periodic orbits in analytic potentials consider prograde x2 orbits (but not retrograde x4 orbits) to be another fundamental building block of bars (Skokos et al. 2002a; Binney & Tremaine 2008). In our study we do not find any orbits parented by the periodic prograde x2 orbit in our initial bar model, but we do find orbits that are parented by the periodic retrograde x4 orbit\u2014a result that is consistent with previous studies (Sparke & Sellwood 1987; Pfenniger & Friedli 1991; Voglis et al. 2007).","Citation Text":["Sellwood 2014b"],"Functions Text":["The retrograde x4 (stable) orbit family is also elongated perpendicular to the bar at small radii, but becomes rounder as it extends to larger radii (for detailed description of orbit families and how they are identified see","In the frame of reference rotating with the bar, all of these families are characterized by a 1:2 resonance between the tangential oscillation frequency (\u03a9\u03d5) and the radial or epicyclic frequency (\u03a9R)."],"Functions Label":["Background","Background"],"Citation Start End":[[1144,1158]],"Functions Start End":[[840,1064],[1161,1362]]}
{"Identifier":"2020MNRAS.498.2575W__Tampo_et_al._2020_Instance_1","Paragraph":"The sample of DES RETs shows a preference for low-metallicity, strongly star-forming host environments. The PDF of their metallicities displays a strong similarity to the hosts of SESNe, as well as LGRBs. There is a clear difference to the PDF of SNe II, which follow SDSS field galaxies. The preference for low-metallicity systems is not as strong as for LGRBs or SLSNe, but the highest metallicities found in all three samples are very similar at around solar metallicity. This result is suggestive of a stripped-envelope, massive-star origin for RETs. The population of RET hosts lies, on average, between CCSNe and LGRBs\/SLSNe in terms of both star formation and metallicity. A loose correlation exists between the luminosity and rarity of events, and the host galaxy conditions required for their formation \u2212 on average, rarer events occur in more extreme environments. The approximate rate of RETs (\u226510\u22126 Mpc\u22123\u2009yr\u22121; Drout et al. 2014; P18; Coppejans et al. 2020; Ho et al. 2020; Tampo et al. 2020, although the definition of RET varies in the above calculations) is ${\\sim}1{{\\ \\rm per\\ cent}}$ of the CCSN rate (Horiuchi et al. 2011; Li et al. 2011; Strolger et al. 2015), which itself is divided into the more common SNe II and sub-dominant SESNe (Kelly & Kirshner 2012; Frohmaier et al., submitted). At ${\\sim}1{{\\ \\rm per\\ cent}}$ of the CCSN rate, RETs are more common than SLSNe (${\\sim}0.01{-}0.05{{\\ \\rm per\\ cent}}$ of CCSNe; McCrum et al. 2015; Prajs et al. 2017; Frohmaier et al., in preparation) and LGRBs (intrinsically ${\\sim}0.08{{\\ \\rm per\\ cent}}$ when accounting for beaming; Graham & Schady 2016). These figures place the rate of DES RETs between extreme objects (SLSNe, LGRBs) and more common SNe (SNe II, SESNe) in terms of rate, matching the location of RET hosts in the various host galaxy parameter spaces presented in Section 5. While stressing rates are uncertain and host galaxy parameters span wide ranges for all transients, they are both linked to the respective transients\u2019 progenitor channels. While it is likely that RETs are a heterogeneous population comprising several progenitor scenarios, it is reasonable to infer from the rates and the host properties that RETs are linked to very massive stars, potentially stripped of their envelopes, and possibly sharing some of the extreme properties of SLSN or LGRB progenitors such as rapid rotation and low metallicity. This hypothesis can be extended to posit that some RETs represent an intermediate and\/or precursory step in the late stages of evolution of a massive star that is close to forming a SLSN or LGRB, whereby the initial collapse of the star occurs leading to shock breakout and subsequent cooling driving the RET light curve (P18), but the progenitor is sufficiently different to the progenitors of LGRBs and SLSNe such that the central engine either does not form or has properties that differ from the central engines of LGRBs or SLSNe, hence the lack of longer-term light curves.","Citation Text":["Tampo et al. 2020"],"Functions Text":["The approximate rate of RETs (\u226510\u22126 Mpc\u22123\u2009yr\u22121"],"Functions Label":["Uses"],"Citation Start End":[[986,1003]],"Functions Start End":[[875,921]]}
{"Identifier":"2019MNRAS.485..189O__Ogiya_&_Mori_2014_Instance_1","Paragraph":"Finally, we emphasize that although the parameter space covered by the DASH library is vast, it is by no means exhaustive. One obvious shortcoming, as discussed above, is that the DASH simulations are inadequate to describe major mergers with ${\\cal M}\\lesssim 100$. In those cases, dynamical friction due to the host, and self-friction due to tidally stripped material, cause the orbit of the subhalo to decay, exposing it to stronger tides. Another degree of freedom not covered here is the inner density slope of dark matter haloes. It is well known that observations of dwarf galaxies often suggest that their haloes have constant density cores, rather than the steep r\u22121-cusps predicted by dark matter-only simulations (e.g. Burkert 1995; Gentile et al. 2004; Oh et al. 2011; Hayashi & Chiba 2015). Such cores can be created within the CDM paradigm by a variety of baryonic processes (e.g. El-Zant, Shlosman & Hoffman 2001; Inoue & Saitoh 2011; Pontzen & Governato 2012; Ogiya & Mori 2014), and have a dramatic impact on the tidal evolution of subhaloes (Pe\u00f1arrubia et al. 2010; Errani, Pe\u00f1arrubia & Tormen 2015; Ogiya 2018). In addition, baryons modify the potentials of host- and subhaloes through the bulges and discs that they form at the halo centres, and these also strongly impact the tidal fields (Errani et al. 2017; Garrison-Kimmel et al. 2017). Finally, in the DASH simulations presented here, the host halo is assumed to be spherically symmetric, which allows us to completely specify each orbit with only two parameters (energy and angular momentum). Cosmological simulations, though, indicate that dark matter haloes are expected to be triaxial systems (e.g. Jing & Suto 2002; Allgood et al. 2006; Hayashi, Navarro & Springel 2007), consistent with the shapes of the gravitational potentials of galaxies and clusters as inferred from a variety of observations (e.g. Oguri et al. 2005; Corless & King 2007; Law & Majewski 2010). Triaxial systems have a much richer variety of orbits, which is likely to impact the tidal evolution of subhaloes.","Citation Text":["Ogiya & Mori 2014"],"Functions Text":["Such cores can be created within the CDM paradigm by a variety of baryonic processes (e.g."],"Functions Label":["Background"],"Citation Start End":[[976,993]],"Functions Start End":[[804,894]]}
{"Identifier":"2021MNRAS.507.4316B__Tan_et_al._2014_Instance_1","Paragraph":"One of the main debates in the star formation community is whether the massive young stellar objects are a scaled-up version of low-mass young stellar objects (YSOs) where disc-accretion plays the dominating role for gaining the stellar mass. For the formation of low-mass stars, bipolar outflows driven by the accretion discs are proposed to be the basic formation mechanism theoretically (Shu, Adams & Lizano 1987), and are also verified observationally (e.g. Bontemps et al. 1996; Richer et al. 2000; Arce et al. 2007, and references therein). However, on the other hand, understanding of the formation mechanism of massive stars is still elusive (Tan et al. 2014). Two major competing models for massive star formation are (i) core accretion via disc (McKee & Tan 2003) and (ii) competitive accretion (Bonnell et al. 2001). The most obvious way to distinguish between these two models might be the detection of the accretion disc around massive protostars. But a direct detection of accreting disc is difficult because the accretion disc is small and short-lived, and also because of complicated gas dynamics at that scale (Kim & Kurtz 2006). Here, the study of the properties of molecular outflows that are the manifestation of disc-accretion in young sources, could help us to improve our understanding of the underlying formation process (Shepherd & Churchwell 1996; Beuther et al. 2002; Molinari et al. 2002; Arce et al. 2007). If it is assumed that massive stars do form via an accretion disc similar to the low-mass stars, they should generate massive and powerful outflows (see de Villiers et al. 2014, and references therein). A few recent studies (e.g. de Villiers et al. 2014; Li et al. 2018) indeed found the applicability of the same scaling between outflow activity and clump masses for both low-mass and massive objects, suggesting a similar formation mechanism. However, an extensive study of molecular outflows toward massive star-forming regions is still lacking owing to their large distances and high level of clustering. Recent interferometric observations with the ALMA enable us to target such regions thanks to its high spatial resolution and sensitivity. Studies of outflows associated with Galactic massive star-forming regions may provide us clues to better understand the launching mechanism of molecular outflows, and hence, the underlying star formation mechanism.","Citation Text":["Tan et al. 2014"],"Functions Text":["However, on the other hand, understanding of the formation mechanism of massive stars is still elusive"],"Functions Label":["Background"],"Citation Start End":[[651,666]],"Functions Start End":[[547,649]]}
{"Identifier":"2019AandA...630A..26M__Wozniakiewicz_et_al._(2012)_Instance_1","Paragraph":"The majority of the particles collected by Stardust are olivine and pyroxene silicates with solar isotopic compositions, which suggests an origin in our solar system rather than an interstellar provenance. These polymineralic particles dominate those made of a single mineral even down to sizes smaller than 100 nm, indicating that the dust composition is surprisingly consistent at different scales and that the smallest subunits of the dust may be as small as tens of nanometers (H\u00f6rz et al. 2006; Zolensky et al. 2006). The sizes of these smallest single mineral impactors are similar to those of the nanocrystals determined by Rietmeijer (1993). As discussed above, they might also be existing in MIDAS dust particles and might be fused into the 100 nm features. Price et al. (2010) and Wozniakiewicz et al. (2012) investigated the sizes of particles smaller than 10 \u03bcm that impacted the aluminum foils of the Stardust probe. The distribution peaks at about 175 nm, but if we assume that the particles areagglomerates of smaller subunits, as indicated by their common polymineralic nature, then the subunit size distribution would peak at sizes below 100 nm (Price et al. 2010). A study of over 450 particles that do not seem to be agglomerates, that is, those that show single mineral impactors of silicate or sulfide, found geometric mean sizes of \n\n$532^{741}_{-310}$\n\n\n\n\n532\n\n\u2212310\n\n+741\n\n\n\n nm for the silicate particles and \n\n$406^{491}_{-222}$\n\n\n\n\n406\n\n\u2212222\n\n+491\n\n\n\n nm for the sulfides (Wozniakiewicz et al. 2013). These sizes are notably larger than the 175 nm (or less) found for the whole dataset. This large spread of subunit sizes could indicate a size distribution with a large width. No fits of these sizedistributions are available, but the figures in Wozniakiewicz et al. (2012) and Price et al. (2010) indicate that the differential sizes may follow a log-normal distribution. When we assume that the smallest subunit sizes are possibly between tens and hundreds of nanometers, the subunit size range found for MIDAS smallest features would be encompassed. The determination of the size distributions for the small Stardust particles and a detailed comparison to the distributions obtained for comet 67P could be the work of an interesting future project.","Citation Text":["Wozniakiewicz et al. (2012)"],"Functions Text":["Price et al. (2010) and","investigated the sizes of particles smaller than 10 \u03bcm that impacted the aluminum foils of the Stardust probe. The distribution peaks at about 175 nm, but if we assume that the particles areagglomerates of smaller subunits, as indicated by their common polymineralic nature, then the subunit size distribution would peak at sizes below 100 nm"],"Functions Label":["Uses","Uses"],"Citation Start End":[[791,818]],"Functions Start End":[[767,790],[819,1161]]}
{"Identifier":"2019MNRAS.484.3307M__Hobbs_et_al._2005_Instance_1","Paragraph":"For most of our models, we can already determine, or at least extrapolate, the final neutron star properties quite well, barring the possibility of late-time fallback. Except for the $12.5 \\, \\mathrm{M}_\\odot$ model, mass outflow already dominates over mass accretion on to the PNS, and the PNS mass has practically stabilized at its final value. Correcting for the binding energy of the neutron stars, we obtain gravitational masses between $1.22 $ and $1.44 \\, \\mathrm{M}_\\odot$, which is compatible with the distribution of observed neutron star masses (\u00d6zel & Freire 2016; Antoniadis et al. 2016; Tauris et al. 2017). While the neutron star kicks are still growing at the end of the simulations due to the long-range gravitational tug by the asymmetric ejecta, the subsequent acceleration of the neutron star can be smoothly extrapolated to obtain tentative final values in all but one case. The extrapolated kicks range from $11$ to $695 \\, \\mathrm{km}\\, \\mathrm{s}^{-1}$. Thus, the most extreme, ECSN-like models with the smallest helium cores can reproduce the very low kicks required to explain some double neutron star systems and pulsars in globular clusters (Tauris et al. 2017), while the models with higher He core masses are compatible with the typical kicks of young pulsars (Arzoumanian, Chernoff & Cordes 2002; Hobbs et al. 2005; Ng & Romani 2007). If the extrapolated kick of $1236 \\, \\mathrm{km}\\, \\mathrm{s}^{-1}$ for the $18 \\, \\mathrm{M}_\\odot$ model of M\u00fcller et al. (2017a) is included, the 3D coconut-fmt models roughly span the full range of observed kick velocities. We see tentative evidence for a correlation of the kick velocity with the explosion energy as proposed by Janka (2017) and Vigna-G\u00f3mez et al. (2018, as a refinement of earlier ideas for progenitor-dependent kicks by Bray & Eldridge 2016). Our models suggest that this correlation may not be a tight one, however, and that the kicks may scatter between zero and an upper limit that scales with the explosion energy. Low kicks can be achieved in more energetic explosions if the explosion geometry is bipolar rather than unipolar, as has already been noted in 2D by Scheck et al. (2006). Such a bipolar explosion occurs in one of our seven simulations (the $12 \\, \\mathrm{M}_\\odot$ model). Although there is some concern that the bipolarity may be connected to the grid geometry, we find unipolar models even in cases where we do not include strong aspherical seed perturbations in the convective O shell that break grid alignment; this suggests that the possibility of bipolar neutrino-driven explosions with low kicks is real in 3D. We also find a loose correlation between the neutron star mass and the kick velocity, which is in line with current observations, and partly theoretical expectations, of double neutron stars (Tauris et al. 2017), but cannot make as strong a case for this correlation based on our simulations. An investigation of a larger suite of supernova simulations of ultra-stripped stars is needed to confirm this hypothesis.","Citation Text":["Hobbs et al. 2005"],"Functions Text":["while the models with higher He core masses are compatible with the typical kicks of young pulsars"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1328,1345]],"Functions Start End":[[1191,1289]]}
{"Identifier":"2016ApJ...827...31F__Lu_&_Yuan_1998_Instance_1","Paragraph":"While treated as fully relativistic under strong gravity, we note that the current model is time-independent based on axisymmetric plasma. This assumption makes it impossible for us to predict any temporal nature of the soft excess considered in this work, e.g., spectral time variabilities associated with shock compression and cooling effects. The downstream plasma properties are numerically solved by considering adiabatic (nonradiative) Rankine\u2013Hugoniot jump conditions as a pure mathematical discontinuity with no energy\/mass loss. Hence, most of the heat generated at the shock front is advected with the downstream plasma. A more realistic shock process, on the other hand, is most likely accompanied by radiative cooling to some degree in which the post-shock plasma temperature may stay comparatively as cool as that of the upstream one, as in the isothermal shocks (e.g., Lu & Yuan 1998; Das et al. 2003; Fukumura & Tsuruta 2004; Fukumura & Kazanas 2007b). Radiative dissipation at the shock front could therefore drastically change the subsequent downstream plasma condition, which in turn alters the Comptonization process. In reality, furthermore, accreting plasma may be characterized by a two-temperature gas between electrons and ions (e.g., Shapiro et al. 1976; Mahadevan 1998; Manmoto 2000) unless the Coulomb coupling between the two is very efficient, whereas in this work we prescribed a single-fluid approximation for simplicity. Becker et al. (2011) have considered a particle transport process (e.g., bulk advection, spatial diffusion, and particle escape) via the effects of the first-order Fermi acceleration across a standing shock. In a more self-consistent scenario such a calculation of diffusive shock acceleration should be incorporated to reflect the energetic outflows\/jets from the shock front. Although all these micro-physics should be addressed and incorporated into more sophisticated calculations of GRMHD simulations for completeness, this is beyond the scope of this paper.","Citation Text":["Lu & Yuan 1998"],"Functions Text":["A more realistic shock process, on the other hand, is most likely accompanied by radiative cooling to some degree in which the post-shock plasma temperature may stay comparatively as cool as that of the upstream one, as in the isothermal shocks (e.g.,"],"Functions Label":["Differences"],"Citation Start End":[[883,897]],"Functions Start End":[[631,882]]}
{"Identifier":"2019ApJ...885...93F__McClintock_et_al._2006_Instance_1","Paragraph":"The vertical structure of an accretion disk has been investigated in detail by Abramowicz et al. (1997). A general accurate expression for the vertical hydrodynamical equilibrium that is valid both for a thin and a slim disk has been derived in their work. The hydrostatic balance in the vertical direction of the accretion disk is described by\n19\n\n\n\n\n\n\n\n1\n\n\n\u03c1\n\n\n\n\n\n\n\u2202\np\n\n\n\u2202\nz\n\n\n\n+\n\n\n\n\u2202\n\u03c8\n\n\n\u2202\nz\n\n\n\n+\n\n\nv\n\n\nr\n\n\n\n\n\n\u2202\n\n\nv\n\n\nz\n\n\n\n\n\u2202\nr\n\n\n\n+\n\n\nv\n\n\nz\n\n\n\n\n\n\u2202\n\n\nv\n\n\nz\n\n\n\n\n\u2202\nz\n\n\n\n=\n0\n,\n\n\nwhere p is the pressure, \u03c1 is the density of the gas, and \u03c8 is the gravitational potential (Abramowicz et al. 1997). As we focus on how the general properties of the disk structure is affected by the radiation-driven outflows, we simplify the expression of the hydrostatic balance in the vertical direction as\n20\n\n\n\n\n\n\n\n1\n\n\n\u03c1\n\n\n\n\n\n\n\u2202\np\n\n\n\u2202\nz\n\n\n\n+\n\n\n\n\u2202\n\u03c8\n\n\n\u2202\nz\n\n\n\n=\n0\n,\n\n\nwhere we assume the terms of \u2202vz\/\u2202r and \u2202vz\/\u2202z to be negligible. It has been pointed out that the widely used approximation for the vertical component of gravity, GMz\/R3, is only valid for the thin disk with H\/R \u226a 1, and a more accurate expression of vertical gravity should be adopted for a slim disk (McClintock et al. 2006; Gu & Lu 2007; Jiao et al. 2009; Gu 2012; Cao & Gu 2015). As discussed in Cao & Gu (2015), there are upper limits on the radiation flux, frad, and half-thickness, H, for a slim disk, above which the radiation force will overwhelm the vertical gravity. The maximal flux can be calculated with\n21\n\n\n\n\n\n\nf\n\n\nrad\n\n\n=\nq\n(\nH\n)\n=\n\n\n\nGM\n\u03c1\nH\n\n\n\n\n\n\nr\n\n\n2\n\n\n+\n\n\nH\n\n\n2\n\n\n\n\n\n\n\n\n\n\n\n\nr\n\n\n2\n\n\n+\n\n\nH\n\n\n2\n\n\n\n\n\u2212\n\n\nr\n\n\ns\n\n\n\n\n\n\n2\n\n\n\n\n\n.\n\n\nHow the gas is blown away from the disk surface by the radiation force is still quite unclear, and, therefore, it is assumed that the gas at the disk surface will be accelerated into outflows by the radiation force when the radiation force overwhelms the vertical gravity. Then, the mass accretion rate in the disk decreases due to outflows. This self-adjustment mechanism leads to a maximal radiation flux\/thickness of the disk (see the detailed discussion in Cao & Gu 2015). Thus, we assume that the outflows are triggered in the disk where the condition\n22\n\n\n\n\n\n\nQ\n\n\nrad\n\n\n\u2212\n\n\n\u2265\n2\n\n\nf\n\n\nrad\n\n\nmax\n\n\n,\n\n\nis satisfied, where \n\n\n\n\n\n\nf\n\n\nrad\n\n\nmax\n\n\n\n\n is the maximal radiation flux and then the mass accretion rate is self-adjusted by the outflows to maintain\n23\n\n\n\n\n\n\nQ\n\n\nrad\n\n\n\u2212\n\n\n\u2261\n2\n\n\nf\n\n\nrad\n\n\nmax\n\n\n.\n\n\nThis relation is adopted in the energy of Equation (15) for the region in the disk where outflows are driven. Otherwise, the outflows are suppressed, and the disk is the same as the conventional slim disk without outflows, which is described by a set of equations for a slim disk (see Chapter 7 in Kato et al. 2008), i.e., the continuity equation:\n24\n\n\n\n\n\n\nM\n\n\n\u02d9\n\n\n=\n\u2212\n2\n\u03c0\nr\n\u03a3\n\n\nv\n\n\nr\n\n\n,\n\n\nwhere \n\n\n\n\n\n\nM\n\n\n\u02d9\n\n\n\n\n remains constant radially, and the angular momentum equation is\n25\n\n\n\n\n\n\nv\n\n\nr\n\n\n(\n\u03a9\n\n\nr\n\n\n2\n\n\n\u2212\n\n\nj\n\n\nin\n\n\n)\n=\n\u2212\n\u03b1\n\n\nrc\n\n\ns\n\n\n2\n\n\n.\n\n\nExcept for Equations (23)\u2013(25), the radial-component of the momentum equation, the equation of state, and energy equation are the same as the disk with outflows (see Equations (4) and (14)\u2013(18)).","Citation Text":["McClintock et al. 2006"],"Functions Text":["It has been pointed out that the widely used approximation for the vertical component of gravity, GMz\/R3, is only valid for the thin disk with H\/R \u226a 1, and a more accurate expression of vertical gravity should be adopted for a slim disk"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1152,1174]],"Functions Start End":[[914,1150]]}
{"Identifier":"2017AandA...605A..88L__Ceccarelli_et_al._2000_Instance_1","Paragraph":"Altogether, the approximately thirty molecules recently detected have confirmed the chemical complexity in the nebula, and generated our interest for the present study. Of these species, we will focus our attention on the seventeen species listed by molecular families in Table 1. As can be seen in this table, the WHISPER survey allowed the detection of some organic molecules in the Horsehead nebula, such as formaldehyde (H2CO) and methanol (CH3OH), which constitute key species in the likely synthesis of more complex organic molecules such as some prebiotic molecules (Bernstein et al. 2002; Mu\u00f1oz Caro et al. 2002; Garrod et al. 2008). Because they are detected in a wide variety of interstellar sources \u2013 in hot cores (Sutton et al. 1995; Ceccarelli et al. 2000), dark clouds (Bergman et al. 2011), shocked regions (e.g. Sakai et al. 2012; Codella et al. 2012; Tafalla et al. 2010) and even in comets (Mumma & Charnley 2011; Cordiner et al. 2015) \u2013 it is of prime importance to understand well how these precursor molecules form. H2CO is commonly thought to form both in the gas-phase and on grain surfaces, while CH3OH is believed to be only formed on grain surfaces (Garrod et al. 2006; Geppert et al. 2006). Guzman et al. (2013) reported the observations of these two molecules toward the Horsehead nebula in both the PDR and Core positions. Unable to reproduce the observed abundances of either H2CO or CH3OH at the PDR position with only pure gas-phase models, they concluded that, for this region, both species are formed on grain surfaces and then photodesorbed into the gas phase. On the other hand, at the Core position, a pure gas-phase model can reproduce the observed H2CO abundance, while photodesorption of ices is still needed to explain the observed abundance of CH3OH. Other organic molecules were reported in the Horsehead nebula as first detections in a PDR environment, including HCOOH (formic acid), CH2CO (ketene), CH3CHO (acetaldehyde), and CH3CCH (propyne) (Guzman et al. 2014). Their abundances were found to be higher at the PDR position than at the Core, revealing that complex organic chemistry is also occurring in UV-illuminated neutral gas (Guzman et al. 2014). Of these molecules, some \u2013 HCOOH, CH2CO, and CH3CHO \u2013 have now also been detected in the Orion bar PDR (Cuadrado et al. 2016, 2017). ","Citation Text":["Ceccarelli et al. 2000"],"Functions Text":["Because they are detected in a wide variety of interstellar sources \u2013 in hot cores","it is of prime importance to understand well how these precursor molecules form."],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[746,768]],"Functions Start End":[[642,724],[956,1036]]}
{"Identifier":"2016MNRAS.458.3760S__Pettini_et_al._1994_Instance_1","Paragraph":"Plots of various characteristics of sub-DLAs and DLAs in the literature. Blue star points mark the detections presented in this work. The Xs mark detections from P\u00e9roux et al. (2011a, 2012). The points mark detections from previous studies in the literature. References: Pettini et al. 2000, Lacey et al. 2003, Junkkarinen et al. 2004, Chen et al. 2005, Burbidge et al. 1996, Chun et al. 2006, Rao et al. 2005, Boisse et al. 1998, Le Brun et al. 1997, Gharanfoli et al. 2007, Rao et al. 2006, Steidel et al. 2002, Lanzetta et al. 1997, Lanzetta et al. 1995, Bergeron & Boisse 1991, Deharveng et al. 1995, Meiring et al. 2007, Cristiani et al. 1987, Bowen et al. 2005, Schulte-Ladbeck et al. 2005, Bowen et al. 2001, Rosenberg et al. 2006, Steidel et al. 1997, Bergeron 1986, Prochaska & Wolfe 1997, Weatherly et al. 2005, Lu et al. 1993, Lu et al. 1997, Djorgovski et al. 1996, Christensen et al. 2004, Rao & Turnshek 2000, Ellison et al. 2005, Pettini et al. 1994, Francis et al. 1996, D'Odorico et al. 2002, Francis & Williger 2004, Francis et al. 2001, Kulkarni et al. 2005, Cherinka et al. 2009, Noterdaeme et al. 2009 and Fynbo et al. 2010. See table 4 of P\u00e9roux et al. (2011a) for the full list of measurements. For the SFR plots, colour indicates the emission from which the value was determined: blue for [O\u2009ii], red for H\u03b1, green for H\u03b2, purple for [O\u2009iii], and black for Ly\u03b1. SFRs have not been dust corrected. The solid red line in the SFR versus zabs plot is an arbitrarily normalized Madau relationship from Madau & Dickinson (2014). Qualitatively, the trend in host galaxy SFR with absorber redshift is consistent with this trend. The blue outliers are the interacting pair in the field of Q1436-0051. The dashed vertical lines indicate the column density boundary between sub-DLAs and DLAs. The outlier in b at \u223c180 kpc is a host galaxy that is part of a cluster environment reported in Francis et al. (2004) and references therein. Grey arrows indicate limits for non-detections. The error bars on Q1436-0051 zabs = 0.9281 correspond to the range of possible H\u2009i column densities and Zn metallicities. See the text for details.","Citation Text":["Pettini et al. 1994"],"Functions Text":["The points mark detections from previous studies in the literature. References:"],"Functions Label":["Uses"],"Citation Start End":[[945,964]],"Functions Start End":[[191,270]]}
{"Identifier":"2020AandA...639A..48S__Shulyak_et_al._2004_Instance_1","Paragraph":"In order to predict emission and transmission spectra of HJs we utilized the \u03c4-REx (Tau Retrieval for Exoplanets) software package (Waldmann et al. 2015b,a). This package uses up-to-date molecular cross sections based on line lists provided by EXOMOL6 project (Tennyson & Yurchenko 2012) and HITEMP (Rothman et al. 2010). We additionally used the HCN line list after Harris et al. (2006) provided by the EXOCLIME project7 and generated with the HELIOS-K code (Grimm & Heng 2015). These cross sections are precomputed on a grid of temperatures and pressures and are stored in binary opacity tables that are available for a number of spectral resolutions. The continuum opacity includes Rayleigh scattering on molecules and collisionally-induced absorption due to H2 \u2013H2 and H2\u2013He either after Abel et al. (2011, 2012) or Borysow et al. (2001); Borysow (2002); Borysow & Frommhold (1989), respectively. We extended the public version of \u03c4-REx8 by incorporating additional opacity sources essential for the atmospheres of UHJs. In particular, bound-free and free\u2013free transitions of H\u2212 become one of the major continuum opacity contributors for the temperatures hotter than about 2000 K. The cross sections of H\u2212 are from John (1988). We also included opacity due to free\u2013free transitions of He\u2212 as well as Rayleigh scattering on H\u202fI atoms and Thomson scattering on free electrons. The He\u2212 cross sections are those originally from John (1968) using the polynomial fit by Carbon et al. (1969). Rayleigh scattering on H\u202fI is calculated after Dalgarno (1962). All relevant numerical routines were extracted from the\u202fLLMODELS stellar model atmosphere code (Shulyak et al. 2004). At the temperatures of HJs the H\u2212 is the dominant continuum opacity source that impacts the observed spectra of these planets (Arcangeli et al. 2018). However, at millibar pressures the He\u2212 opacity can become comparable or even stronger than that of H\u2212 for wavelengths longer than 1.6 \u03bcm, as shown in Fig. 3 (third panel from the bottom), where we display examples of continuum opacity coefficient at different altitudes in the atmosphere of a HJ with Teq = 3000 K. At even smaller pressures, electron scattering and Rayleigh scattering on H\u202fI atoms also become important contributors to the continuum opacity at particular wavelengths (top panel in Fig. 3). However, the contribution of H\u202fI and e\u2212 on the transmission spectra is marginal because their opacity is only strong at low pressures that are hardly probed by transmission spectroscopy. Thus, among all continuum opacity sources, only H\u2212 and He\u2212 significantly contribute to the predicted amplitude of the transmission spectra, as shown in the bottom plot of Fig. 3. When both continuum and line opacity are included, the effect of He\u2212 on predicted spectra is diluted by a much stronger opacity in molecular lines while H\u2212 contribution is still significant. Nevertheless, as can be seen from the second plot (from bottom) in Fig. 3, in optically thick layers the He\u2212 opacity could still be stronger than, for example, collision-induced absorption due to H2 -H2 and H2-He. We thus conclude that accurate calculation of atmospheric opacity requires He\u2212, H\u202fI, and e\u2212 opacity included especially at low pressures, following modern stellar atmosphere codes. However, observed transmission and emission spectra of HJs are hardly affected by these three opacity sources. Finally, we updated HELIOS with He\u2212 and e\u2212 opacity and found out that this has little impact on the atmospheric temperature structure (with a modification in local temperature of at most \u0394T \u2248 10 K) and thus can be ignored; the original version of HELIOS already includes H\u202fI Rayleigh scattering.","Citation Text":["Shulyak et al. 2004"],"Functions Text":["All relevant numerical routines were extracted from the\u202fLLMODELS stellar model atmosphere code"],"Functions Label":["Uses"],"Citation Start End":[[1650,1669]],"Functions Start End":[[1554,1648]]}
{"Identifier":"2021MNRAS.507.4389G__Masters_et_al._2011_Instance_1","Paragraph":"Erwin (2018) showed that, in a sample drawn from the Spitzer Survey of Stellar Structure in Galaxies (S4G), the bar fraction is constant over a range of (g \u2212r) colours and gas fractions. Their bar fraction does not increase, but rather decreases for stellar masses higher than \u223c 109.7M\u2299. These results are in contrast to many SDSS-based studies cited above. Erwin (2018) argues that this apparent contradiction can be explained if SDSS-based studies miss bars in low-mass blue galaxies. In Figs 5 and 6, we showed that the newly detected bars in GZD (compared to GZ2) are weak bars in low-mass blue galaxies. Nevertheless, the \u2018combined\u2019 bar fraction in Fig. 6 is not constant over (g \u2212r) colour and agrees well with Masters et al. (2011) for redder colours [(g \u2212r) colour > 0.5]. Additionally, our \u2018combined\u2019 bar fraction remains roughly constant over stellar mass. As mentioned before, we conclude that strong bars drive the trends of bar fraction with (g \u2212r) colour, stellar mass, and SFR observed in other studies (Nair & Abraham 2010b; Masters et al. 2011, 2012; Vera et al. 2016; Cervantes Sodi 2017). However, the addition of weak bars in low-mass blue galaxies is insufficient to resolve the apparent disagreement between Erwin (2018) and many SDSS-based studies (Masters et al. 2011, 2012; Vera et al. 2016; Cervantes Sodi 2017; Kruk et al. 2018), which instead seems likely to be due to the very different sample selection of the S4G and SDSS galaxy samples. For example, the median stellar mass of the sample used in Erwin (2018) is \u223c109.6M\u2299 (based on their Fig. 4 and the bins in the top left-hand panel of their Fig. 5). However, the median stellar mass of our sample is 1010.6M\u2299. As stellar mass correlates with many parameters (including bar length), this can have major consequences. Additionally, as Erwin (2018) notes, there is also the issue of resolution to consider. With an r-band FWHM of 1.18 arcsec from DECaLS (Dey et al. 2019) and a mean redshift of 0.036, the mean linear resolution of our sample is approximately 834 pc, which is higher than the 165 pc of Erwin (2018). This explains why they observe many sub-kpc bars, while we do not. These differences in stellar mass and resolution will manifest themselves in the conclusions, so a more detailed analysis is needed for a proper comparison with Erwin (2018).","Citation Text":["Masters et al. (2011)"],"Functions Text":["Nevertheless, the \u2018combined\u2019 bar fraction in Fig. 6 is not constant over (g \u2212r) colour and agrees well with","for redder colours [(g \u2212r) colour > 0.5]"],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[717,738]],"Functions Start End":[[609,716],[739,779]]}
{"Identifier":"2015ApJ...814...73S__Soria_et_al._2004_Instance_1","Paragraph":"At a distance of 4.8 Mpc (Karachentsev et al. 2002), NGC 5408 X-1 is one of the best studied ULXs. It has been observed by several of the current generation of X-ray satellites, on a multitude of occasions. Observations include: an XMM-Newton large programme (e.g., Pasham & Strohmayer 2012); Swift XRT monitoring (Kaaret & Feng 2009; Gris\u00e9 et al. 2013); and eight Chandra exposures, which we re-analyze here. The flux variability of the source rules out an X-ray supernova remnant and confirms that it is powered by accretion onto a compact object (Kaaret et al. 2003; Soria et al. 2004). It persistently displays a distinct soft ultraluminous two component X-ray spectrum in XMM-Newton data (Sutton et al. 2013) at an average 0.3\u201310 keV unabsorbed luminosity of \n\n\n\n\n\n (Strohmayer 2009; although we note that they fit the high energy spectrum with a soft power-law, so may over-estimate the intrinsic luminosity, cf. Middleton et al. 2014). Additional soft residuals have been detected in the XMM-Newton spectra which can be well modeled as thermal plasma emission (Strohmayer & Mushotzky 2009; Miller et al. 2013; Middleton et al. 2014). It has previously been assumed that these were the result of diffuse star formation related emission in the host galaxy (Strohmayer et al. 2007; Strohmayer & Mushotzky 2009; Miller et al. 2013). However, we know from observational studies of galaxies that the X-ray luminosity of such emission is correlated with star formation rate (\n\n\n\n\n\n\/\n\n\n\n\n\n; Mineo et al. 2012), and Middleton et al. (2014) contend that the luminosity of the putative thermal plasma emission (\n\n\n\n\n\n calculated from Miller et al. 2013) greatly exceeds that inferred from star formation, even over the entirety of NGC 5408 (\n\n\n\n\n\n calculated from a 24 \u03bcm flux density of \n\n\n\n\n\n, Dale et al. 2005). Instead, Middleton et al. (2014) show that the putative plasma emission features could actually be commensurate with broadened, blueshifted absorption in a partially ionized, optically thin medium. Such a medium would be expected to occur in a super-Eddington wind, as it becomes optically thin at large distances (\n\n\n\n\n\n) from the central black hole.","Citation Text":["Soria et al. 2004"],"Functions Text":["The flux variability of the source rules out an X-ray supernova remnant and confirms that it is powered by accretion onto a compact object"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[570,587]],"Functions Start End":[[410,548]]}
{"Identifier":"2015AandA...579A..46M__Ilyushin_et_al._2010_Instance_1","Paragraph":"The Hamiltonian used in the present work is the so-called RAM (rho axis method) internal-rotation Hamiltonian based on the work of Kirtman (Kirtman 1962), Lees and Baker (Lees & Baker 1968), and Herbst et al. (Herbst et al. 1984). Since rather complete descriptions of this method, which takes its name from the choice of axis system, have been presented several times (Hougen et al. 1994; Kleiner 2010) we do not repeat this general description here. The main advantage of the RAM Hamiltonian is its general approach that simultaneously takes into account the A- and E-symmetry species and all the torsional levels, intrinsically taking the intertorsional interactions into account within the rotation-torsion manifold of energy levels. This method was successfully applied to a number of molecules containing a C3v rotor and Cs frame, including the main isotopolog of acetaldehyde (Smirnov et al. 2014). As for the main isotopolog (Smirnov et al. 2014) we employed the RAM36 (rho-axis-method for 3- and 6-fold barriers) code that uses the RAM approach for the molecules with the C3v top attached to a molecular frame of Cs or C2v symmetry and having 3- or 6-fold barriers to internal rotation, respectively (Ilyushin et al. 2010, 2013). The Hamiltonian in the RAM36 program is presented by the following expression: (1)\\begin{eqnarray} \\nonumber  H &amp;=&amp; (1\/2) \\sum_{knpqrs} B_{knpqrs0} \\left[P^{2k}P_z^nP_x^pP_y^qp_{\\alpha}^r \\cos(3s{\\alpha}) \\right.\\\\[2mm] \\nonumber  &amp;&amp;\\left. + \\cos(3s\\alpha) p_{\\alpha}^rP_y^qP_x^pP_z^nP^{2k}\\right] \\\\[2mm] \\nonumber &amp;&amp; + (1\/2) \\sum_{knpqrt} B_{knpqr0t} \\left[P^{2k}P_z^nP_x^pP_y^qp_{\\alpha}^r \\sin(3t\\alpha) \\right.\\\\ [2mm]  &amp;&amp;\\left. + \\sin(3t\\alpha) p_{\\alpha}^rP_y^qP_x^pP_z^nP^{2k}\\right] \\end{eqnarray}H=(1\/2)\u2211knpqrsBknpqrs0[P2kPznPxpPyqp\u03b1rcos(3s\u03b1)+cos(3s\u03b1)p\u03b1rPyqPxpPznP2k]+(1\/2)\u2211knpqrtBknpqr0t[P2kPznPxpPyqp\u03b1rsin(3t\u03b1)+sin(3t\u03b1)p\u03b1rPyqPxpPznP2k]where the Bknpqrst are fitting parameters; p\u03b1 is the angular momentum conjugate to the internal rotation angle \u03b1; and Px,Py,Pz are projections on the x,y,z axes of the total angular momentum P. In the case of a C3v top and Cs frame (as is appropriate for acetaldehyde), the allowed terms in the torsion-rotation Hamiltonian must be totally symmetric in the group G6 (and also must be Hermitian and invariant to the time reversal operation). Since all individual operators p\u03b1,Px,Py,Pz,P2,cos(3s\u03b1) and sin(3t\u03b1) used in Eq. (1) are Hermitian, all possible terms provided by Eq. (1) will automatically be Hermitian. The particular term to be fit is represented in the input file with a set of k,n,p,q,r,s,t integer indices that are checked by the program for conformity with time reversal and symmetry requirements, to prevent accidental introduction of symmetry-forbidden terms into the Hamiltonian. ","Citation Text":["Ilyushin et al. 2010"],"Functions Text":["As for the main isotopolog","we employed the RAM36 (rho-axis-method for 3- and 6-fold barriers) code that uses the RAM approach for the molecules with the C3v top attached to a molecular frame of Cs or C2v symmetry and having 3- or 6-fold barriers to internal rotation, respectively"],"Functions Label":["Uses","Uses"],"Citation Start End":[[1210,1230]],"Functions Start End":[[906,932],[955,1208]]}
{"Identifier":"2021MNRAS.500.2336Y__Matonick_&_Fesen_1997_Instance_2","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"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[1264,1285]],"Functions Start End":[[1041,1171]]}
{"Identifier":"2022AandA...658A.194P__Khata_et_al._2020_Instance_3","Paragraph":"The stellar photospheric parameters we collected from literature for the benchmark stars are summarized in Table A.1. Although most benchmark stars have v sini  2 km s\u22121 (Reiners et al. 2018), there are two stars with larger values: J07558+833 (12.1 km s\u22121) and J13005+056 (16.4 km s\u22121). These stars are useful to investigate the performance of the algorithms when dealing with higher rotational velocities. The literature values were derived with different methods. These methods include: interferometry to estimate the stellar radius and Teff (Boyajian et al. 2012; S\u00e9gransan et al. 2003; von Braun et al. 2014; Berger et al. 2006; Newton et al. 2015), synthetic model fitting using BT-Settl models to determine Teff (Gaidos et al. 2014; L\u00e9pine et al. 2013; Gaidos & Mann 2014; Mann et al. 2015) and log g (L\u00e9pine et al. 2013), empirical relations to derive stellar mass in the form of mass-luminosity relations (Mann et al. 2015; Khata et al. 2020; Boyajian et al. 2012; Berger et al. 2006; S\u00e9gransan et al. 2003), along with the mass-magnitude relations (Maldonado et al. 2015), mass-radius relations (von Braun et al. 2014), mass\u2013Teff relations (Gaidos & Mann 2014; Gaidos et al. 2014), empirical relations to derive the stellar radius in the form of mass-radius relations (Maldonado et al. 2015) and Teff\u2013radius relations (Gaidos & Mann 2014; Gaidos et al. 2014; Houdebine et al. 2019), pEW measurements to determine Teff (Maldonado et al. 2015; Neves et al. 2014; Newton et al. 2015) and [Fe\/H] (Maldonado et al. 2015; Neves et al. 2014; Gaidos et al. 2014; Mann et al. 2015), the definition of spectral indices such as the H2O-K2 index to estimate Teff (Rojas-Ayala et al. 2012), as well as the combination of the H2O-K2 index with pEWs to derive [Fe\/H] (Rojas-Ayala et al. 2012; Khata et al. 2020), the stellar radius and Teff (Khata et al. 2020), and spectral curvature indices for the determination of Teff (Gaidos & Mann 2014). Additionally, [Fe\/H] was derived by using color-magnitude metallicity relations (Dittmann et al. 2016), atomic line strength relations (Gaidos & Mann 2014), and spectral feature relations (Terrien et al. 2015). Terrien et al. (2015) used K-band magnitudes and the Dartmouth Stellar Evolution Program (Dotter et al. 2008) to derive the stellar radius, whereas Mann et al. (2015) employed the Boltzmann equation with Teff determined from synthetic model fits. Last, but not least, Houdebine et al. (2019) derived Teff from photometric colors. For more details on the individual methods, we refer to the descriptions in the corresponding works.","Citation Text":["Khata et al. 2020"],"Functions Text":["the stellar radius and Teff"],"Functions Label":["Background"],"Citation Start End":[[1837,1854]],"Functions Start End":[[1808,1835]]}
{"Identifier":"2020AandA...644A..97C__Leroy_et_al._2013_Instance_3","Paragraph":"Major nearby galaxy cold gas mapping surveys (Regan et al. 2001; Wilson et al. 2009; Rahman et al. 2011; Leroy et al. 2009; Donovan Meyer et al. 2013; Bolatto et al. 2017; Sorai et al. 2019; Sun et al. 2018) have focused on observations of the molecular gas (through CO lines). Despite a few notable exceptions (e.g. Alatalo et al. 2013; Saintonge et al. 2017), these surveys observed mainly spiral or infrared-bright galaxies (i.e. galaxies with significant star formation) and have furthered our understanding of how star formation happens, rather than how it stops. This boils down to quantifying the relation between molecular gas and star formation rate (SFR), which appears nearly linear in nearby discs (Kennicutt 1998; Bigiel et al. 2008; Leroy et al. 2013; Lin et al. 2019). This relationship is often parametrised via the ratio between the SFR and the molecular gas mass (Mmol), which is called the molecular star formation efficiency (SFE = SFR\/Mmol = 1\u2215\u03c4dep), where the inverse of the SFE is the depletion time, \u03c4dep. The depletion time indicates how much time is necessary to convert all the available molecular gas into stars at the current star formation rate. On kpc scales and in the discs of nearby star-forming galaxies, \u03c4dep is approximately constant around 1\u20132 Gyr (Bigiel et al. 2011; Rahman et al. 2012; Leroy et al. 2013; Utomo et al. 2017), and it appears to weakly correlate with many galactic properties such as stellar mass surface density or environmental hydrostatic pressure (Leroy et al. 2008; Rahman et al. 2012). Nevertheless, small but important deviations for a constant SFE have been noticed, which can be the first hints of star formation quenching. In some galaxies, the depletion time in the centres appear shorter (Leroy et al. 2013; Utomo et al. 2017) or longer (Utomo et al. 2017) with respectto their discs. These differences may correlate with the presence of a bar or with galaxy mergers (Utomo et al. 2017; see also Muraoka et al. 2019) and do not seem to be related to unaccounted variation in the CO-to-H2 conversion factor (Leroy et al. 2013; Utomo et al. 2017). Spiral arm streaming motions have also been observed to lengthen depletion times (Meidt et al. 2013; Leroy et al. 2015).","Citation Text":["Leroy et al. 2013"],"Functions Text":["In some galaxies, the depletion time in the centres appear shorter"],"Functions Label":["Background"],"Citation Start End":[[1756,1773]],"Functions Start End":[[1688,1754]]}
{"Identifier":"2022MNRAS.515.2698A__In_2018_Instance_1","Paragraph":"In the last few years, similar (though not identical) experimental developments to our method have been reported. In 2017, Wehres et al. (2017) described the design, construction, and operation of two laboratory broad-band emission spectrometers for gas-phase characterization of large molecules. The first is based on a Schottky-barrier diode heterodyne receiver spectrometer operating between 80 and 110\u2009GHz, whilst the second uses cryo-cooled SIS technology and operates at higher frequency between 270 and 290\u2009GHz (Wehres et al. 2018). Gas phase pyridine and methyl cyanide features were successfully detected and matching analytic simulations were presented for the observed spectral transitions. In 2018, a new technique combining a Terahertz radiometer (41\u201349) GHz and a vacuum chamber was developed to observe the generation of cold plasma and UV photochemistry in the gas phase under low pressure as described by Tanarro et al. (2018). However, in the context of the experiment described here, it is important to note that the spectrometer developments were not coupled to a system where the gas-phase molecules originated from the solid state. More recently, Yocum et al. (2019) reported on an experiment using a THz source and hot-electron bolometer to detect the gas-phase absorption spectra of simple molecular species such as H2O, D2O, and CH3OH desorbing into the gas-phase from ices grown in an ultra-high vacuum (UHV) chamber. This experimental set-up (SubLIME), combined with a Fourier-transform Infrared spectrometer, gave the spectroscopic insight on the ultraviolet photolysis and warm up of methanol ice sample, for the first time observed at those frequencies in a laboratory environment (Yocum et al. 2021). This showed the potential of the use of submillimetre\/far-infrared technique to identify molecules in complex gas mixtures. One challenge of this technique was the sensitivity of absorption spectroscopy to the desorbing molecules. The combination of temperature programmed desorption and microwave spectroscopy techniques was discussed by Theul\u00e9 et al. (2020). In these experiments, desorption of water, deuterium, methanol, and ammonia was measured using a hot electron bolometer and Si-diode technology combined with Fourier transform spectrometers. While Theule et al. observed the desorption through a waveguide cavity, Yocum et al. (2019) detected the desorption above a metal substrate as seen in standard TPD studies.","Citation Text":["Tanarro et al. (2018)"],"Functions Text":["In 2018, a new technique combining a Terahertz radiometer (41\u201349) GHz and a vacuum chamber was developed to observe the generation of cold plasma and UV photochemistry in the gas phase under low pressure as described by","However, in the context of the experiment described here, it is important to note that the spectrometer developments were not coupled to a system where the gas-phase molecules originated from the solid state."],"Functions Label":["Background","Differences"],"Citation Start End":[[922,943]],"Functions Start End":[[702,921],[945,1153]]}
{"Identifier":"2022AandA...665A..25C__Lambrechts_&_Johansen_2012_Instance_1","Paragraph":"Various theoretical studies employing numerical simulations have been performed to investigate planet formation through core accretion or gravitational instability in the low-stellar-mass regime. Payne & Lodato (2007) assessed planet formation around brown dwarfs adapting models for higher stellar masses based on core accretion. Through Monte Carlo simulations, they found that Earth-like planets can form in this condition and the planet mass depends strongly on the disk mass. However, none of their simulations showed a planetary rocky core accreting a gaseous envelope to form a giant planet. A way to overcome the radial drift barrier is a rapid rocky core growth. Pebble accretion is a mechanism able to speed up significantly the giant planet formation process (i.e., Lambrechts & Johansen 2012; Bitsch et al. 2015). Liu et al. (2020) carried out a theoretical study on planet formation driven by pebble accretion in the (sub)stellar mass range between 0.01 and 0.1 M\u2299. First, they calculated the initial masses of protoplanets by extrapolating previous numerical simulations conducted in previous literature. Next, they performed a population synthesis study to track the growth and migration of a large sample of protoplanets under the influence of pebble accretion. Their results show that, around a 0.01 M\u2299 brown dwarf, planets can grow up to 0.1\u22120.2 M\u2295, while, around 0.1 M\u2299 stars, planets can reach a maximum mass of 2\u22123 M\u2295. Findings from this study show that even pebble accretion does not seem to be sufficient to form gas giants around VLM stars and brown dwarfs. Miguel et al. (2020) used a population synthesis approach based on planetesimal accretion to explore planet formation in the stellar mass range between 0.05 and 0.25 M\u2299. They let the synthetic population of planetary systems evolve for 108 yr. The authors find that to form planets with masses higher than 0.1 M\u2295 they need stars of at least 0.07 M\u2299, implying that planet formation around brown dwarf may not be a usual outcome. Then, stars with masses higher than 0.15 M\u2299 are necessary to form planets more massive than the Earth. Therefore, from all of these studies, we conclude that core accretion model currently cannot explain the presence of gas giants around VLM stars or brown dwarfs. Either the core accretion theory is incomplete, or another mechanism for planet formation is needed. This is confirmed by Lodato et al. (2005), who discussed the origin of the 5 MJup planet detected around the 25 MJup brown dwarf 2MASSW J1207334-393254 (Chauvin et al. 2005). They found that the core accretion mechanism is far too slow to generate such a planet in less than 107 yr, the estimated age of the system. Therefore, the authors proposed gravitational instabilities arising during the early phases of the disk lifetime as a viable possibility for the formation of the planet.","Citation Text":["Lambrechts & Johansen 2012"],"Functions Text":["Pebble accretion is a mechanism able to speed up significantly the giant planet formation process (i.e.,"],"Functions Label":["Background"],"Citation Start End":[[777,803]],"Functions Start End":[[672,776]]}
{"Identifier":"2021MNRAS.505..435S__Deming_et_al._2013_Instance_1","Paragraph":"Detections of ionic, atomic, and molecular species in exoplanetary atmospheres serve as a unique and strong diagnostic of those chemical and dynamical processes driving their formation and evolution. Their detection and abundance measurements could act as indicators of planetary formation scenarios and reveal connections to the primordial protoplanetary disc and the host star (Williams & Cieza 2011; Mordasini et al. 2016; Madhusudhan et al. 2017). Furthermore, discoveries of atmospheric chemical species allow us to better understand various thermodynamical processes and chemistry, winds in the upper atmosphere (Goodman 2009; Snellen et al. 2010; Brogi et al. 2016; Madhusudhan et al. 2016; Wyttenbach et al. 2020), and to probe planetary interiors and various bulk properties through their abundances (Kite et al. 2016; Thorngren & Fortney 2019; Madhusudhan et al. 2020). A whole host of ions, atoms, and molecules have been detected through a variety of, often complementary, techniques, such as differential spectrophotometry using low-to-mid resolution spectroscopy (e.g. Gibson et al. 2012, 2017; Deming et al. 2013; Kreidberg et al. 2014; Kirk et al. 2016; Nortmann et al. 2016), and high resolution spectroscopic techniques (e.g. Redfield et al. 2008; Snellen et al. 2008; Rodler, Lopez-Morales & Ribas 2012; Birkby et al. 2013; Hoeijmakers et al. 2015, 2018, 2020; Brogi et al. 2016; Birkby et al. 2017; \u017d\u00e1k et al. 2019; Ehrenreich et al. 2020). To date, ionic species such as Fe\u2009ii and Ti\u2009ii (Hoeijmakers et al. 2019), atomic absorption from Na, K, H \u03b1, and He (e.g. Redfield et al. 2008; Sedaghati et al. 2016; Casasayas-Barris et al. 2017; Spake et al. 2018; Chen et al. 2020; Seidel et al. 2020), and molecules such as H2O, CH4, and CO (e.g. Konopacky et al. 2013; Brogi et al. 2014; Fraine et al. 2014; Barman et al. 2015; Sing et al. 2016) have been detected through the aforementioned techniques. Needless to say that this list of detected constituents is by no means exhaustive, nor that of methods employed to detect exoplanetary atmospheres. For instance, high-resolution imaging instruments such as SPHERE (Beuzit et al. 2019) and GRAVITY (Gravity Collaboration et al. 2017), both at the VLT (ESO\u2019s Very Large Telescope), through combination with low-dispersion spectroscopy, have facilitated direct measurements of exoplanetary atmospheres (Samland et al. 2017; Gravity Collaboration et al. 2020).","Citation Text":["Deming et al. 2013"],"Functions Text":["A whole host of ions, atoms, and molecules have been detected through a variety of, often complementary, techniques, such as differential spectrophotometry using low-to-mid resolution spectroscopy (e.g."],"Functions Label":["Background"],"Citation Start End":[[1109,1127]],"Functions Start End":[[880,1082]]}
{"Identifier":"2021AandA...655A.104P__Draine_&_Li_2007_Instance_1","Paragraph":"This last, more technical aspect gave birth to a branch of research devoted to the problem of fitting a galaxy SED (Arnouts et al. 1999; Bolzonella et al. 2000; Cid Fernandes et al. 2005; Ilbert et al. 2006; Ocvirk et al. 2006; Tojeiro et al. 2007; Fritz et al. 2007, 2017; Franzetti et al. 2008; Pappalardo et al. 2010; Han & Han 2014; Leja et al. 2017; Weaver et al. 2021). Historically, models reproducing the evolution of stellar populations and dust emission have been developed separately, with reliable prescriptions to build synthetic spectra of different stellar populations on one side (Bruzual & Charlot 1993; Bressan et al. 1994; Worthey 1994; Fioc & Rocca-Volmerange 1997; Leitherer et al. 1999; Vazdekis 1999; Charlot & Fall 2000; Maraston 2005; Fioc & Rocca-Volmerange 2019), and a thorough comprehension of dust grains physics on the other (Draine & Lee 1984; Dale 2001; Takeuchi et al. 2003, 2005; Zubko et al. 2004; Draine & Li 2007; da Cunha et al. 2010; Silva et al. 2011; Asano et al. 2013; Calura et al. 2014; Zhukovska 2014; Mancini et al. 2015; Schneider et al. 2016; Popping et al. 2017; Aoyama et al. 2017; De Vis et al. 2017, 2019; Ginolfi et al. 2018; Graziani et al. 2019; Burgarella et al. 2019; Nanni et al. 2019; De Looze et al. 2020; Galliano et al. 2021). The two branches finally converged, performing simultaneous fitting of both components (Devriendt et al. 1999; Groves et al. 2008; da Cunha et al. 2008; Noll et al. 2009; Silva 2009; Graziani et al. 2019). Two of these methods, CIGALE (Noll et al. 2009; Boquien et al. 2019) and MAGPHYS (da Cunha et al. 2008), are based on the assumption that the radiation produced during the star formation process by the stellar and nebular components is partially absorbed by the dust and then re-emitted in the infrared part of the spectrum, satisfying the energy conservation. New families of codes have also been produced in an attempt to combine spectral and photometric analyses, recovering, through the implementation of photoionisation codes, emission lines feature; see for example Prospect (Robotham et al. 2020) and Prospector (Leja et al. 2017). Recently, even machine learning methods have been involved in this challenge, with neural network algorithms showing promising results (Simet et al. 2021).","Citation Text":["Draine & Li 2007"],"Functions Text":["Historically, models reproducing the evolution of stellar populations and dust emission have been developed separately,","and a thorough comprehension of dust grains physics on the other"],"Functions Label":["Background","Background"],"Citation Start End":[[934,950]],"Functions Start End":[[376,495],[791,855]]}
{"Identifier":"2017ApJ...850...97B__Tamburro_et_al._2009_Instance_1","Paragraph":"The H i mass fraction of every gas particle in the baryonic runs is calculated based on the particle\u2019s temperature and density and the cosmic UV background radiation flux while including a prescription for self-shielding of H2 and dust shielding in both H i and H2 (Christensen et al. 2012). This allows for the straightforward calculation of the total H i mass of each simulated galaxy. We create mock H i data cubes only for the 42 halos that contain \n\n\n\n\n\n. Specifically, we create mock data cubes that mimic ALFALFA observations (Haynes et al. 2011). After specifying a viewing angle (see below), our code considers the line-of-sight velocity of each gas particle. The velocity of each particle is tracked in the simulation by solving Newton\u2019s equations of motion, but any turbulent velocity of the gas is not taken into account. Velocity dispersions in dwarf galaxies can be on the order of the rotational velocity, \u223c10\u201315 km s\u22121 (e.g., Stanimirovi\u0107 et al. 2004; Tamburro et al. 2009; Oh et al. 2015). Dispersions are thought to be driven at least partially by thermal velocities or supernovae (Tamburro et al. 2009; Stilp et al. 2013a, 2013b). In our simulations, supernovae inject thermal energy, and the thermal state of the H i gas needs to be considered in the mock H i linewidth for a realistic comparison to observations. To account for the thermal velocity, the H i mass of each gas particle is assumed to be distributed along the line-of-sight in a Gaussian distribution with a standard deviation given by the thermal velocity dispersion, \n\n\n\n\n\n, where T is the temperature of the gas particle. After this thermal broadening is calculated, a mock H i data cube can be generated by specifying the spatial and velocity resolution. For all of our galaxies, we adopt a spatial resolution of 54 pixels across 2Rvir. In practice, this corresponds to a range of \u223c1 kpc resolution in our lowest-mass galaxies up to \u223c9 kpc resolution in our most massive galaxies. However, the spatial resolution plays no role in our study, since measurements of the VF are based on spatially unresolved H i data. For the velocity resolution, we match the ALFALFA specification of 11.2 km s\u22121 (two-channel boxcar-smoothed).","Citation Text":["Tamburro et al. 2009"],"Functions Text":["Velocity dispersions in dwarf galaxies can be on the order of the rotational velocity, \u223c10\u201315 km s\u22121 (e.g.,"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[968,988]],"Functions Start End":[[834,941]]}
{"Identifier":"2021ApJ...923L..22A__Dvorkin_&_Barausse_2017_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":["Dvorkin & Barausse 2017"],"Functions Text":["Orbiting SMBHBs produce a stochastic GW background (GWB"],"Functions Label":["Background"],"Citation Start End":[[781,804]],"Functions Start End":[[484,539]]}
{"Identifier":"2021ApJ...921..107B__Bellovary_et_al._2021_Instance_1","Paragraph":"A black hole mass this large in Leo I is not expected from extrapolation of any of the standard black hole\u2212host galaxy correlations. Of course, these small systems do not necessarily need to follow the trends seen in normal galaxies, but the black hole mass reported here does stand out. L\u00fctzgendorf et al. (2015) explore extrapolations of black hole correlations down to globular cluster scales, and using a velocity dispersion of 12 km s\u22121, Leo I has a black hole mass a factor of 100 more than the extrapolated trends. On the numerical side, van Wassenhove et al. (2010) consider different scenarios for formation of a black hole in Milky Way satellites and place the likelihood of one of them having a black hole around the size found here to be below 1%, but this result also depends on the initial seed mass (see also Bellovary et al. 2021). Runaway mergers of stellar mass black holes are unlikely to produce such a black hole in such a small galaxy, since the required initial mass function to reach the ratios seen in the models might be more top-heavy than what chemical abundances and star formation history studies suggest. An alternative explanation for the abnormally large central black hole may come from the recent study on Leo I\u2019s star formation history from Ruiz-Lara et al. (2020). The authors identify a period of quenching from z = 1\u20132 followed by reignition until almost the present day, when ram pressure stripping may have shut it down as it fell into the Milky Way. While the authors speculate that this reignition at intermediate redshifts could be due to a past merger with a smaller dwarf, this could also be consistent with gas accretion and potential active galactic nuclei feedback, lending support to the high MBH values presented here. Amaro-Seoane et al. (2014) also suggest that dwarf systems may in fact have significantly larger black holes compared to the host galaxy\u2013black hole relationships. Having a larger sample of black hole limits measured in dwarf galaxies will be important to explore.","Citation Text":["Bellovary et al. 2021"],"Functions Text":["On the numerical side, van Wassenhove et al. (2010) consider different scenarios for formation of a black hole in Milky Way satellites and place the likelihood of one of them having a black hole around the size found here to be below 1%, but this result also depends on the initial seed mass (see also"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[824,845]],"Functions Start End":[[522,823]]}
{"Identifier":"2016AandA...592A.157M__Mernier_et_al._2015_Instance_1","Paragraph":"As seen in Fig. 6 (right), the large MOS-pn discrepancy in the Ni\/Fe abundance ratio prevents us from deriving a precise measurement. This discrepancy is worrying, but can be explained by imperfections in the cross-calibration of the two instruments. Alternatively, and perhaps more likely, the high energy band around the Ni-K transitions is significanly affected by the instrumental background (as the flux of the cluster emission sharply decreases at high energies). This hard particle background (already mentioned in Sect. 3.1) has a different spectral shape in MOS and pn, which might even vary with time, thus between observations. In particular, an instrumental line (Cu K\u03b1) is known to affect pn at a rest-frame energy of ~8 keV (Mernier et al. 2015). Despite our efforts to carefully estimate the background, that line might interfere with the Ni-K line in several observations, making a proper modelling of the Ni-K line impossible, and hence, boosting the Ni absolute abundance in pn. In this context, it can be instructive to compare our Ni\/Fe measurements with those of Suzaku, which has a lower relative hard particle background. Sato et al. (2007b) (A\u20091060) and Tamura et al. (2009) (Perseus) reported ratios of ~1.3 \u00b1 0.4 and ~1.11 \u00b1 0.19, respectively (after rescaling to the proto-solar values). Although these measurements might be also be affected by further uncertainties (e.g. the choice of the spectral modelling, Sect. 4.3), they appear to be consistent with the Ni\/Fe average ratio measured with MOS is this work, favouring our above supposition that MOS is more trustworthy than pn for measuring Ni\/Fe. However, in order to be conservative, we prefer to retain the pn value as a possible result and, therefore, we keep large systematic uncertainties for Ni\/Fe. We finally note that, unsurprisingly, Ni\/Fe cannot be constrained in the cool objects (Fig. 6, left) because the gas temperature is too low to excite Ni-K transitions. ","Citation Text":["Mernier et al. 2015"],"Functions Text":["In particular, an instrumental line (Cu K\u03b1) is known to affect pn at a rest-frame energy of ~8 keV","Despite our efforts to carefully estimate the background, that line might interfere with the Ni-K line in several observations, making a proper modelling of the Ni-K line impossible, and hence, boosting the Ni absolute abundance in pn."],"Functions Label":["Uses","Uses"],"Citation Start End":[[739,758]],"Functions Start End":[[639,737],[761,996]]}
{"Identifier":"2020AandA...643A..93H__Lin_et_al._2016a_Instance_1","Paragraph":"In recent years, type Ia supernovae (SNe Ia; Amanullah et al. 2010; Suzuki et al. 2012; Betoule et al. 2014; Scolnic et al. 2018) have been widely employed to test cosmic isotropy. Antoniou & Perivolaropoulos (2010) searched for the preferred direction of anisotropy for the Union2 sample by adopting the hemisphere comparison (HC) method (Schwarz & Weinhorst 2007). They found a maximum accelerating expansion rate, which corresponds to a preferred direction of anisotropy. After that, Mariano and Perivolaropoulos (Mariano & Perivolaropoulos 2012) found a possible preferred anisotropic direction at the 2\u03c3 level using the Union2 sample, but by employing the dipole fitting (DF) method. Since then, these two methods have been widely used to explore the cosmic anisotropy (Cai & Tuo 2012; Cai et al. 2013; Zhao et al. 2013; Li et al. 2013; Heneka et al. 2014; Bengaly et al. 2015; Andrade et al. 2018; Sun & Wang 2019) by investigating observational data of, for instance, the Union2.1 sample (Yang et al. 2014; Javanmardi et al. 2015; Lin et al. 2016a), the Joint Light-Curve Analysis (JLA) sample (Lin et al. 2016b; Chang et al. 2018a; Wang & Wang 2018), the Pantheon sample (Sun & Wang 2018), gamma-ray bursts (GRBs; Wang & Wang 2014), galaxies (Zhou et al. 2017), as well as gravitational wave and fast radio bursts (Qiang et al. 2019; Cai et al. 2019). Using HC and DF methods, Zhao et al. (2019) studied the cosmic anisotropy via the Pantheon sample. They found that the SDSS sample plays a decisive role in the Pantheon sample. It may imply that the inhomogeneous distribution has a significant effect on the cosmic anisotropy (Chang et al. 2018b). This opinion was also presented by Sun & Wang (2019). Their conclusions show that the effect of redshift on the result is weak and there is a negligible anisotropy when making a redshift tomography. Deng & Wei (2018a) tested the cosmic anisotropy with the Pantheon sample, but by using the following three methods: the HC method, the DF method, and Healpix1 (G\u00f3rski et al. 2005). They also performed a cross check. There are two preferred directions from the HC method. In adopting the DF method and Healpix, they found no noticeable anisotropy. They also compared the HC method with the DF method by using the JLA sample (Deng & Wei 2018b) and found that the results of these two methods have not always been approximately coincident with each other. In order to better test the cosmic isotropy, the best way would be to add new samples with a relatively homogeneous distribution.","Citation Text":["Lin et al. 2016a"],"Functions Text":["Since then, these two methods have been widely used to explore the cosmic anisotropy","by investigating observational data of, for instance, the Union2.1 sample"],"Functions Label":["Background","Background"],"Citation Start End":[[1038,1054]],"Functions Start End":[[689,773],[921,994]]}
{"Identifier":"2018MNRAS.473.3810Y__Mitrushchenkov_et_al._2017_Instance_1","Paragraph":"The lack of data on inelastic processes due to collisions with neutral hydrogen atoms has been a major limitation on modelling of F-, G- and K-star spectra in statistical equilibrium, and thus to reliably proceeding beyond the assumption of local thermodynamic equilibrium (LTE) in analysis of stellar spectra and the determination of elemental abundances. This problem has been well documented, e.g. see Lambert (1993); Barklem (2016a) and references therein. Significant progress has been made in recent times through detailed full-quantum scattering calculations, based on quantum chemical data, for the cases of simple atoms such as Li, Na, Mg and Ca (Belyaev & Barklem 2003; Barklem, Belyaev & Asplund 2003; Belyaev et al. 2010; Barklem et al. 2010; Belyaev et al. 2012; Barklem et al. 2012; Mitrushchenkov et al. 2017). These calculations have demonstrated the importance of the ionic-covalent curve crossing mechanism leading naturally to charge transfer processes (mutual neutralization and ion-pair production), in addition to excitation and de-excitation processes. The importance of this mechanism has allowed various simplified model approaches to be developed, which may be used in cases where suitable quantum chemistry data are not been available. In particular a semi-empirical model has been employed for Al, Si, Be and Ca (Belyaev 2013a,b; Belyaev, Yakovleva & Barklem 2014b; Yakovleva, Voronov & Belyaev 2016; Belyaev et al. 2016), and a theoretical model based on a two-electron asymptotic linear combinations of atomic orbitals (LCAO) approach, has also been employed for Ca (Barklem 2016b, 2017). Comparisons of the two methods show quite good agreement and reasonable agreement with the full quantum calculations is found, particularly for the most important processes with the largest rates (Barklem 2016b, 2017; Mashonkina, Sitnova & Belyaev 2017; Mitrushchenkov et al. 2017). Thus, the model approaches provide a useful route for obtaining estimates of the rates for these processes for many elements of astrophysical interest.","Citation Text":["Mitrushchenkov et al. 2017"],"Functions Text":["Significant progress has been made in recent times through detailed full-quantum scattering calculations, based on quantum chemical data, for the cases of simple atoms such as Li, Na, Mg and Ca"],"Functions Label":["Background"],"Citation Start End":[[797,823]],"Functions Start End":[[461,654]]}
{"Identifier":"2019MNRAS.482.3803P__Schaerer_&_Vacca_1998_Instance_1","Paragraph":"The broad blue WR bump feature around 4686 \u00c5 was searched in all the star-forming regions in the galaxies. A clear detection of the broad blue WR bump was made only in two star-forming regions hosted by SBS 1222+614 as shown in Fig. 8. This detection is in good agreement with Shirazi & Brinchmann (2012), who have also reported the WR features in SBS 1222+614. The broad blue bump consists of a blend of C\u2009iii\/C\u2009iv \u03bb4650, 4658, N\u2009iii \u03bb4634, 4640, [Ar\u2009iv] \u03bb4711, 4740, and He\u2009ii \u03bb4686 emission lines. This detection generally indicates a good number (102\u2013105) of young WR stars in the galaxy (e.g. Kunth & Sargent 1981; Kunth & Schild 1986). The blue bump appears mainly due to the presence of late-type WN (WNL) and early-type WC (WCE) stars (Schaerer & Vacca 1998). The red WR bump around 5808 \u00c5 is also expected in the WR galaxies. We also possibly identified the red bump feature in the SBS 1222+614 (# a + b) region as shown in Fig. 9. The red WR bump appears mainly due to the presence of emission lines [N\u2009ii] \u03bb5795 and He\u2009i \u03bb5875 from the WCE stars. The red bump is rarely detected in WR galaxies as it contains very weak emission line and is expected in high-mettalicity region (Guseva, Izotov & Thuan 2000). Although we could not detect the WR features in the star-forming regions of other galaxies in our sample, the optical SDSS spectrum for the brightest regions of knot #b and #a in IC 3521 and CGCG 0410-023, respectively, shows a detection of WR features (Brinchmann et al. 2008). The SDSS spectrum also shows the detection of the broad WR blue bump having a relatively lower strength in the knot #b of CGCG 038-051. These detections are missed out in this study most likely due to the low SNR in the blue part of the spectra. The starburst ages estimated to be very young (\u22646 Myr; see Table 3) for these star-forming regions including those present in SBS 1222+614 are consistent with the detection of the WR features, which appears only during the very early periods of star formation. Although a few other star-forming regions such as the knot #a in IC 3521 and CGCG 038-051 and knot #c in CGCG 041-023 showing a very young starburst of \u22646 Myr are expected to have the WR features, it is not clear from the present data if these regions also host WR stars or not. Some star-forming regions such as knot #c and b in CGCG 038-051 and CGCG 041-023, respectively, show the starburst of ages of \u223c10 Myr, indicating that they have probably completed their WR phases. Overall, the detections of the WR features from our own observations and those from the SDSS data from at least one star-forming region in each galaxy in the sample suggest that the sample galaxies are undergoing young massive star formation phase having a significant population of WR stars.","Citation Text":["Schaerer & Vacca 1998"],"Functions Text":["The blue bump appears mainly due to the presence of late-type WN (WNL) and early-type WC (WCE) stars"],"Functions Label":["Background"],"Citation Start End":[[744,765]],"Functions Start End":[[642,742]]}
{"Identifier":"2020AandA...640A.133R___et_al._2017_Instance_1","Paragraph":"The full sample for which the circumstellar CO line emission will be modeled consists of the \u223c180 C-, M-, and S-type AGB stars analyzed in Sch\u00f6ier & Olofsson (2001), Gonz\u00e1lez Delgado et al. (2003), and Ramstedt et al. (2006) together with additional sources presented in Danilovich et al. (2015). In this initial paper, the new interferometric data for the southern M- and C-type stars are presented. Some of the available sample statistics for the full \u223c180 star DEATHSTAR sample are shown in Fig. 1. The distance distribution (Fig. 1, middle) is compared with the estimated distribution in the solar neighborhood (Jura 1990; Jura & Kleinmann 1992; Jura et al. 1993). The estimated distribution is derived from 2MASS and ground-based observations (Jura & Kleinmann 1990) and assumes a smooth distribution of 40 C-type stars kpc\u22122, a scale height of 200 pc, and that there are a third as many S-type as C-type stars. For the full \u223c180 star DEATHSTAR sample, the C-type stars from Sch\u00f6ier & Olofsson (2001) are all brighter than K = 2 mag. The M-type sample consists of the non-Mira stars from the General Catalog of Variable Stars (GCVS; Samus\u2019 et al. 2017) with quality flag 3 in the IRAS 12, 25, and 60 \u03bcm bands and 60 \u03bcm flux \u22733 Jy with the addition of the Mira stars in Gonz\u00e1lez Delgado et al. (2003). The S-type sample also consists of stars that have good quality flux measurements in the IRAS 12, 25, and 60 \u03bcm bands, that are found in the General Catalog of Galactic S stars, and that are detected in Tc and are hence intrinsic. The completeness of the S-type sample is discussed in Ramstedt et al. (2009) and it is thought to be complete out to 600 pc. Furthermore, stars of all three spectral types are only included in the sample if they are detected in CO radio line emission, which could be reproduced under the assumption of spherical symmetry. Sources that show strongly asymmetric line profiles when observed with single-dish telescopes, or with known CO-detached shells, are hence not included (e.g., R Scl, U Ant, EP Aqr, and \u03c01 Gru). In this paper, we also exclude stars that have previously been observed with ALMA; however, they will be included in the future analysis. The lower panel of Fig. 1 shows that the stellar and wind parameters of the full \u223c180 star sample cover the ranges expected for AGB stars. As expected, fewer stars are found at the high end. The sample is biased to mass-losing stars since only stars that are previously detected in CO radio line emission are included. It is also likely that the full range of AGB masses is not covered simply because higher mass stars are rare. It is our assessment that the full \u223c180 star DEATHSTAR sample is representative of Galactic mass-losing AGB stars and covers the relevant ranges of wind and stellar parameters to provide the necessary constraints for theoretical models.","Citation Text":["Samus\u2019 et al. 2017"],"Functions Text":["The M-type sample consists of the non-Mira stars from the General Catalog of Variable Stars (GCVS;","with quality flag 3 in the IRAS 12, 25, and 60 \u03bcm bands and 60 \u03bcm flux \u22733 Jy"],"Functions Label":["Uses","Uses"],"Citation Start End":[[1138,1156]],"Functions Start End":[[1039,1137],[1158,1234]]}
{"Identifier":"2021MNRAS.504.2168G__Steiner_et_al._2011_Instance_1","Paragraph":"Finally, we attempt to characterize the reflection component using the full 2\u201335\u2009keV spectra with a sophisticated model [M4: ${\\tt{\\rm constant}}$*${\\tt{\\rm TBabs}}$*(${\\tt{\\rm simplr}}$*${\\tt {\\rm kerrbb2}}$+${\\tt{\\rm kerrconv}}$*(${\\tt{\\rm ireflect}}$*${\\tt{\\rm simplc}}$)), to evaluate the impact on the spin measurement and understand the changes of the accretion flow and the interaction between the disc and the corona. This model features a self-consistent treatment of the thermal, Compton scattering and the reflection component: ${\\tt {\\rm kerrbb2}}$ describes the thermal component and supplies the seed photons for ${\\tt{\\rm simplr}}$ (a modified version of ${\\tt{\\rm simpl}}$, Steiner et al. 2011) to generate the Compton component; while a portion of the Compton component will escape to reach an observer, the remains (refer as ${\\tt{\\rm simplc}}$, Steiner et al. 2011) will strike back to the disc to generate the reflected component. The reflection fraction Rref in ${\\tt{\\rm ireflect}}$ (Magdziarz & Zdziarski 1995), defined as the ratio of the Compton photons striking back to the disc to that escaping to infinity, is restricted to negative value thereby only the reflected component is returned by ${\\tt{\\rm ireflect}}$. It is linked to the reflection constant parameter x in ${\\tt{\\rm simplr}}$ via the relation x = 1 + |Rref| (Gou et al. 2011). We set the elemental abundance to unity and the iron abundance AFe to five times the solar abundance (Bharali et al. 2019; Buisson et al. 2019; Xu et al. 2020). The disc temperature Tin is fixed at the value returned by ${\\tt{\\rm diskbb}}$ (M1, refer to Gou et al. 2011). The ionization parameter \u03be is fixed at 1000 (i.e. log(\u03be) = 3, Xu et al. 2020; Buisson et al. 2019), as it is difficult to be constrained. Finally we use ${\\tt{\\rm kerrconv}}$ (Brenneman & Reynolds 2006) to apply relativistic effects assuming an unbroken emissivity profile with index q = 3. The key parameters in ${\\tt {\\rm kerrbb2}}$ and ${\\tt{\\rm kerrconv}}$ are linked together.","Citation Text":["Steiner et al. 2011"],"Functions Text":["This model features a self-consistent treatment of the thermal, Compton scattering and the reflection component: ${\\tt {\\rm kerrbb2}}$ describes the thermal component and supplies the seed photons for ${\\tt{\\rm simplr}}$ (a modified version of ${\\tt{\\rm simpl}}$,","to generate the Compton component;"],"Functions Label":["Uses","Uses"],"Citation Start End":[[690,709]],"Functions Start End":[[426,689],[711,745]]}
{"Identifier":"2019MNRAS.484.1100M__Dekel_&_Birnboim_2006_Instance_1","Paragraph":"The stream velocity is proportional to the halo virial velocity, which can be related to the sound speed of gas at the virial temperature, yielding Mb \u223c 1. The density contrast is obtained by assuming pressure equilibrium between hot gas at Tb \u223c Tv, the virial temperature of an NFW halo, and cold gas at Ts \u223c 104K, set by the steep drop in the cooling rate below that temperature (Sutherland & Dopita 1993). If both the halo and the stream are roughly isothermal, then this ratio is constant throughout the halo. In practice, the stream temperature can be as high as \u223c3 \u00d7 104\u2009K due to photoheating from the UV background (e.g. Goerdt et al. 2010), while the post-shock temperature in the hot halo near Rv may only be \u223c0.5Tv (Dekel & Birnboim 2006, P18). Finally, the stream radius is related to its velocity and density through the mass accretion rate along the stream, ${\\dot{M}}_{\\rm s}\\simeq \\pi R_{\\rm s}^2 \\rho _{\\rm s}V_{\\rm s}$. Cosmological simulations suggest that this is typically a fixed fraction of the total accretion rate on to the halo virial radius (Danovich et al. 2012), which is constrained by cosmology (e.g. Dekel et al. 2013). The final expressions, including uncertainties in model parameters, are\n(42)\r\n\\begin{eqnarray*}\r\nM_{\\rm b}\\simeq 0.75-2.25,\r\n\\end{eqnarray*}\r\n(43)\r\n\\begin{eqnarray*}\r\n\\delta \\simeq (10-100)\\times M_{12}^{2\/3} (1+z)_3^{-1},\r\n\\end{eqnarray*}\r\n(44)\r\n\\begin{eqnarray*}\r\n\\frac{R_{\\rm s}}{R_{\\rm v}} \\simeq (0.01-0.1)\\times \\left(\\delta _{75} M_{\\rm b,1.5}\\right)^{-1\/2},\r\n\\end{eqnarray*}\r\nwhere $M_{12}=M_{\\rm v}\/10^{12}\\, \\mathrm{M}_\\odot$, (1 + $z$)3 = (1 + $z$)\/3, \u03b475 = \u03b4\/75, and Mb,1.5 = Mb\/1.5. We note that \u03b4 \u223c 10 implies that the background gas has a temperature of ${\\sim } 3\\times 10^5\\, {\\rm K}$, which would make it thermally unstable unless it is illuminated by a very hard ionizing background (e.g. Efstathiou 1992, though see also Binney, Nipoti & Fraternali 2009, who argue that buoyancy effects may stabilize the halo against thermal instability even without a photoionizing source). This implies that \u03b4 \u2273 30 is likely the most physically plausible regime.","Citation Text":["Dekel & Birnboim 2006"],"Functions Text":["while the post-shock temperature in the hot halo near Rv may only be \u223c0.5Tv",", P18)."],"Functions Label":["Uses","Uses"],"Citation Start End":[[726,747]],"Functions Start End":[[649,724],[747,754]]}
{"Identifier":"2021AandA...653A..36M__Goulding_&_Alexander_(2009)_Instance_3","Paragraph":"The SFG sample was constructed using the Great Observatories All-Sky LIRG Survey (GOALS sample, Armus et al. 2009), from which we extracted 158 galaxies, with data from Inami et al. (2013), who report the fine-structure lines at high resolution in the 10\u2005\u2212\u200536\u2006\u03bcm interval, and Stierwalt et al. (2014), who include the detections of the H2 molecular lines and the PAH features at low spectral resolution. For those galaxies in the GOALS sample that have a single IRAS counterpart, but more than one source detected in the emission lines, we have added together the line or feature fluxes of all components, to consistently associate the correct line or feature emission to the total IR luminosity computed from the IRAS fluxes. To also cover lower luminosity galaxies, as the GOALS sample only includes luminous IR galaxies (LIRGs) and ultra-luminous IR galaxies (ULIRGs), we included 38 galaxies from Bernard-Salas et al. (2009) and Goulding & Alexander (2009), to reach the total sample of 196 galaxies with IR line fluxes in the 5.5\u2005\u2212\u200535\u2006\u03bcm interval in which an AGN component is not detected. For the Bernard-Salas et al. (2009), Goulding & Alexander (2009), and the GOALS samples, we excluded all the composite starburst-AGN objects identified as those with a detection of [NeV] either at 14.3 or 24.3 \u03bcm. It is worth noting that the original samples from Goulding & Alexander (2009) and Bernard-Salas et al. (2009) have spectra solely covering the central region of the galaxies. To estimate the global SFR, we corrected the published line fluxes of the Spitzer spectra by multiplying them by the ratio of the continuum reported in the IRAS point source catalogue to the continuum measured on the Spitzer spectra extracted from the CASSIS database (Lebouteiller et al. 2015). We assumed here that the line emission scales (at first order) with the IR brightness distribution. In particular, we considered the continuum at 12\u2006\u03bcm for the [NeII]12.8\u2006\u03bcm and [NeIII]15.6\u2006\u03bcm lines, and the continuum at 25 \u03bcm for the [OIV]25.9\u2006\u03bcm, [FeII]26\u2006\u03bcm, [SIII]33.5\u2006\u03bcm, and [SiII]34.8\u2006\u03bcm lines. This correction was not needed for the AGN sample and the GOALS sample because of the greater average redshift of the galaxies in the 12MGS and GOALS samples. In particular, the 12MGS active galaxy sample has a mean redshift of 0.028 (Rush et al. 1993), while the GOALS sample has a mean redshift of 0.026. The galaxies presented by Bernard-Salas et al. (2009) have instead an average redshift of 0.0074, while the sample by Goulding & Alexander (2009) has an average redshift of 0.0044. For the other lines in the 10\u2005\u2212\u200536\u2006\u03bcm interval, Goulding & Alexander (2009) did not report a detection, and we used the data presented in Bernard-Salas et al. (2009) for a total of 15 objects. Both Bernard-Salas et al. (2009) and Goulding & Alexander (2009) reported data from the high-resolution Spitzer-IRS spectra. Data in the 50\u2005\u2212\u2005205\u2006\u03bcm interval were taken from D\u00edaz-Santos et al. (2017). For the GOALS sample, 20 starburst galaxies were taken from Fern\u00e1ndez-Ontiveros et al. (2016), and 23 objects were taken from the ISO-LWS observations of Negishi et al. (2001). As a result, we obtained a total sample of 193 objects. Lastly, the PAH features\u2019 fluxes were measured from the low-resolution Spitzer-IRS spectra by Brandl et al. (2006), including 12 objects from the sample of Bernard-Salas et al. (2009) and 179 objects from Stierwalt et al. (2014).","Citation Text":["Goulding & Alexander (2009)"],"Functions Text":["It is worth noting that the original samples from","and Bernard-Salas et al. (2009) have spectra solely covering the central region of the galaxies."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1359,1386]],"Functions Start End":[[1309,1358],[1387,1483]]}
{"Identifier":"2018AandA...613A...7Y__Kostogryz_et_al._(2016)_Instance_1","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)"],"Functions Text":["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",", according to the selected wavelength, surface gravity, and temperature of a star."],"Functions Label":["Uses","Uses"],"Citation Start End":[[581,604]],"Functions Start End":[[401,580],[604,687]]}
{"Identifier":"2019MNRAS.484..814G__T\u00e1pai,_Zolt\u00e1n_&_L\u00e1szl\u00f3_2015_Instance_1","Paragraph":"There are many methods, beyond those already employed here, that can be used to identify, and probe the masses of, black holes. This includes reverberation mappings of AGNs (e.g. Bahcall, Kozlovsky & Salpeter 1972; Blandford & McKee 1982; Netzer & Peterson 1997), the \u2018fundamental plane of black hole activity\u2019 (Merloni et al. 2003; Falcke, K\u00f6rding & Markoff 2004), spectral modelling of the high-energy X-ray photon coming from the hot accretion discs around IMBHs (Pringle & Rees 1972; Narayan & Yi 1995), high-ionization optical emission lines (Baldwin et al. 1981; Kewley et al. 2001), and high-spatial-resolution observations of maser emission using radio and millimetre\/submillimetre interferometry (e.g. Miyoshi et al. 1995; Greenhill et al. 2003; Humphreys et al. 2016; Asada et al. 2017). In addition, the merging of black holes is now quite famously known to produce gravitational radiation during their orbital decay (Abbott et al. 2016). The merging of galaxies containing their own central IMBH is similarly expected to result in the eventual merging of these black holes. The Kamioka Gravitational Wave Detector (KAGRA; Aso et al. 2013) will be a 3-km-long underground interferometer in Japan capable of detecting the gravitational radiation emanating from collisions involving black holes with masses up to 200\u2009M\u2299 (T\u00e1pai, Zolt\u00e1n & L\u00e1szl\u00f3 2015). The planned Deci-Hertz Interferometer Gravitational wave Observatory (DECIGO; Kawamura et al. 2011) and the European, Laser Interferometer Space Antenna (LISA) Pathfinder mission15 (Anza et al. 2005; McNamara 2013), with their greater separation of mirrors, will be able to detect longer wavelength gravitational waves and thus better reach into the domain of intermediate-mass and supermassive black hole mergers, the latter of which are currently being searched for via \u2018pulsar timing arrays\u2019 (PTAs) (e.g. Hobbs et al. 2010; Kramer & Champion 2013; Shannon et al. 2015). A key constraint to the expected detection threshold of such signals from PTAs \u2013 in particular the background of cosmic ripples from the merger of massive black holes (themselves arising from the merger of galaxies) \u2013 is the (black hole)-to-(host galaxy\/bulge) mass ratio (see equation 4 for spiral galaxies). An additional source of long-wavelength gravitational radiation will arise from the inspiral of compact stellar-mass objects, such as neutron stars and black holes, around these IMBHs (Mapelli et al. 2012). It is reasonable to expect that the densely packed nuclear star clusters, which coexist with low-mass SMBHs (e.g. Gonz\u00e1lez Delgado et al. 2008; Seth et al. 2008; Graham & Spitler 2009), will similarly surround many IMBHs. Gravitational radiation and the gravitational tidal disruption of ill-fated stars that venture too close to these black holes (Komossa et al. 2009, Komossa 2013, and references therein; Zhong, Berczik & Spurzem 2015; Stone & Metzger 2016; Lin et al. 2018) are therefore expected from these astrophysical entities. There is, therefore, an array of future observations that could yield further confidence and insight into the realm of IMBHs.","Citation Text":["T\u00e1pai, Zolt\u00e1n & L\u00e1szl\u00f3 2015"],"Functions Text":["The Kamioka Gravitational Wave Detector","will be a 3-km-long underground interferometer in Japan capable of detecting the gravitational radiation emanating from collisions involving black holes with masses up to 200\u2009M\u2299"],"Functions Label":["Background","Background"],"Citation Start End":[[1330,1357]],"Functions Start End":[[1086,1125],[1151,1328]]}
{"Identifier":"2016ApJ...831...11S__Honeycutt_1992_Instance_1","Paragraph":"The secondary eclipses of the binary are 0.03\u20130.04 mag deep, and so relatively small data reduction issues can have large effects on the fidelity of the light curve. Our brightness measurements were derived from aperture photometry using DAOPHOT (Stetson 1987), although we took several additional steps to increase the precision of the results. We conducted a curve-of-growth analysis of 12 apertures photometered per star using DAOGROW (Stetson 1990) in order to correct all measured stars to a uniform large aperture. We then attempted to unify the data for each filter to a consistent zero point by using ensemble photometry (Sandquist et al. 2003; Honeycutt 1992). This essentially uses all measured stars on the frame to determine magnitude offsets resulting from differences in exposure time, airmass, atmospheric transparency, and the like. Our implementation iteratively fits for position-dependent corrections that result from variations from point-spread function across the frame. These steps each brought noticeable reductions in the amount of scatter in the light curves. However, we still found that the shape of the ground-based light curves of the secondary eclipse did not match what was expected from the Kepler observations. After some investigation, we found that features in the light curve were correlated with those for the brightest star near KIC 9777062 in our images. (The star was about 70 pixels or 28\u2033 distant on the sky, so there was no significant overlap of the point-spread function when the seeing was generally 4\u20138 pixels FWHM.) We found that after subtracting the light curve of this star the out-of-eclipse light levels and the shapes of the primary and secondary eclipse light curves were much more consistent from night to night. When more distant stars on the images were tested, we found a much poorer degree of correlation. We conclude that short length-scale variations were not being corrected for (and could not be corrected for, due to lack of star sampling on the image) with whole-image zero-point corrections or image-scale position-dependent corrections. The variations were of around 0.01 mag size, but could significantly affect the secondary eclipse light curves. We could not identify any features on our flat-field images that could explain the light-curve variations, and they occurred whether dome flats, twilight flats, or hybrid flats (combining the smoothed large-scale variations from the twilight flats and small-scale variations from dome flats) were used. The ground-based eclipse observations are shown in the bottom rows of Figure 1.","Citation Text":["Honeycutt 1992"],"Functions Text":["We then attempted to unify the data for each filter to a consistent zero point by using ensemble photometry","This essentially uses all measured stars on the frame to determine magnitude offsets resulting from differences in exposure time, airmass, atmospheric transparency, and the like. Our implementation iteratively fits for position-dependent corrections that result from variations from point-spread function across the frame."],"Functions Label":["Uses","Background"],"Citation Start End":[[653,667]],"Functions Start End":[[521,628],[670,992]]}
{"Identifier":"2018AandA...615A.161M__Hartmann_1904_Instance_1","Paragraph":"\u03b4 Ori Aa+Ab = Mintaka Aa+Ab = 34 Ori Aa+Ab = HD 36486 A+B = BD \u221200 983 A+B = ALS 14 779 A+B. \u03b4 Ori is Orion\u2019s Belt westernmost star and a multiple system at the center of the Mintaka cluster (Caballero & Solano 2008). It has two close bright visual components (Aa and Ab) and two distant dim ones (B and C) that will not be considered here. Aa is itself composed of two spectroscopic components (Aa1+Aa2) in an eclipsing orbit with a5.7325 d period, while Ab is two magnitudes fainter than the two Aa components and located 0.\u2032\u2032 267 away in 1993 and moving away from the primary (Hartmann 1904; Harvin et al. 2002). We observed the system on two consecutive nights and in both cases we were able to spatially resolve Ab from Aa, this being one of the two systems in this paper (\u03c3 Ori is the other one) where the power of Lucky Spectroscopy becomes more apparent. The two spectra for Aa yield the same spectral type as for the combined Aa+Ab value from GOSSS I, O9.5 IINwk, but the lines are slightly narrower. When analyzing the data for GOSSS I, the combined spectrum was close to being (n), something that does not happen for the spatially resolved spectra (this has the additional advantage of giving us a more purespectrum for a classical O9.5 II classification standard). There are very small variations between the two epochs for Aa, likely due to the motion of Aa1 and Aa2 but the signature of the latter in the combined spectrum is very weak (Shenar et al. 2015). The two epochs for Ab yield the same spectral type, O9.7 III:(n), and in both cases there are indications of only a slight residual contamination from Aa, with the second epoch being noisier. The (n) qualifier indicates that Ab is a fast rotator, as previously noted by Harvin et al. (2002) and Shenar et al. (2015). The differences between the spectral types of Aa (Aa1+Aa2) and Ab are consistent with the Teff and log g differences measured by Shenar et al. (2015) using spectral disentangling (as opposed to the spatial deconvolution used here). Richardson et al. (2015) also spatially resolved Aa and Ab in the UV using HST\/STIS and obtained very similar values of Teff around 31 kK for Aa1 and Ab, with error bars close to 2000 K, which is also consistent with our spectral classifications. That paper using UV data is the only one we have found where \u03b4 Ori Aa+Ab was spatially deconvolved, making our result the first time it has been done in the optical.","Citation Text":["Hartmann 1904"],"Functions Text":["Aa is itself composed of two spectroscopic components (Aa1+Aa2) in an eclipsing orbit with a5.7325 d period, while Ab is two magnitudes fainter than the two Aa components and located 0.\u2032\u2032 267 away in 1993 and moving away from the primary"],"Functions Label":["Background"],"Citation Start End":[[580,593]],"Functions Start End":[[341,578]]}
{"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"],"Functions Text":["We also plot the SFR in the outer regions of spiral galaxies","square symbols)."],"Functions Label":["Uses","Uses"],"Citation Start End":[[916,934]],"Functions Start End":[[854,914],[936,952]]}
{"Identifier":"2021MNRAS.508.2019B__G\u00fcrkan_et_al._2018_Instance_1","Paragraph":"Assuming that non-jetted H\u2009ii galaxies are powered by SF, the relation between radio luminosities in such galaxies and MBH can be deciphered as a link between thermal\/non-thermal SF luminosity and BH mass, in the form $L_{\\rm SF} \\propto M_{\\rm BH}^{0.61}$, which has the largest scatter of the observed linear regressions, due to the large uncertainties of MBH estimates at low values and numerous radio non-detections. This unprecedented relation is likely the result of the observed link found between SF indicators such as radio, far-infrared, optical or line luminosities, and galaxy mass (e.g. Mauch & Sadler 2007; G\u00fcrkan et al. 2018). Moreover, the SF contribution in the radio band with respect to the AGN activity is expected to increase with galaxy mass (Aird, Coil & Georgakakis 2019). In fact, the observed correlation suggests that the radio production due to SF broadly increases with BH mass, as the latter quantity is approximately a constant proportion of the galaxy mass (from a third down to a fifth in the lowest MBH \u223c 105 M\u2299, H\u00e4ring & Rix 2004; Reines & Volonteri 2015; Mart\u00edn-Navarro & Mezcua 2018), i.e. the stellar content: more stars likely produce more radio emission. This is consistent with the idea that SF, although mostly included in the disc (Bluck et al. 2020) and known to scale with galaxy mass (Speagle et al. 2014), would possibly also scale with MBH. More precisely, the nuclear starburst is plausibly set by the available gas mass present in the NSC (MNSC, Fern\u00e1ndez-Ontiveros, Prieto & Acosta-Pulido 2009; Neumayer, Seth & B\u00f6ker 2020), which scales with the galaxy dynamical mass ($M_{\\rm NSC} \\propto M _{\\mathrm{ dyn}}^{0.55}$, Scott & Graham 2013), which in turn respond to the SMBH gravitational well (Cen 2015; Pitchford et al. 2016). Such a sequence of links eventually enacts the radio-MBH and radio-optical relations observed for star-forming galaxies, e.g setting the fraction of radio emission with respect to the photoionizing energy from young stellar populations and SN products.","Citation Text":["G\u00fcrkan et al. 2018"],"Functions Text":["This unprecedented relation is likely the result of the observed link found between SF indicators such as radio, far-infrared, optical or line luminosities, and galaxy mass"],"Functions Label":["Uses"],"Citation Start End":[[621,639]],"Functions Start End":[[421,593]]}
{"Identifier":"2020ApJ...896..169P__Dullemond_&_Dominik_2004_Instance_1","Paragraph":"Pre-main-sequence stars are typically surrounded by protoplanetary disks, and since such disks are generally highly optically thick at optical\/near-infrared (NIR) wavelengths, they may cast shadows on their surroundings. Such disk shadows have been observed on a wide range of angular scales; recently, high-contrast imaging has revealed shadows cast on outer disks by misaligned inner disks on subarcsecond scales (e.g., Marino et al. 2015; Benisty et al. 2018; Casassus et al. 2018). Such shadows reveal that angular distortions are common in inner disks. This is consistent with the existence of a class of self-shadowed disks in which the inner disk scale height is higher than that of the outer disk, leading to a general shadowing and related cooling of the outer disk (Dullemond & Dominik 2004; Dong 2015). A related type of disk shadow, but on a vastly larger angular scale of arcminutes, are shadows cast on large-scale reflection nebulosity. These great disk shadows can occur if a young, usually spatially unresolved, star-disk system is illuminating a reflection nebula, and are especially apparent for systems viewed close to edge-on (Hodapp et al. 2004; Pontoppidan & Dullemond 2005). Neither of these two types of disk shadow should be confused with silhouette disks that obscure background nebulosity, such as the Orion proplyds (O\u2019dell 1998), or dark dust lanes in isolated edge-on disks (Burrows et al. 1996; Stapelfeldt et al. 1998; Duch\u00eane et al. 2010). The projection of the disk onto a large reflection nebula can greatly magnify a small structure in the obscuring disk. Indeed, the apparent angular size of great disk shadows is only limited by the size of the reflection nebula illuminated by the central star, and may be orders of magnitude larger than the protoplanetary disk itself. Great disk shadows therefore present a unique opportunity to explore the geometry of disks on scales otherwise not resolved by direct imaging, primarily the disk scale height, inclination, and position angle.","Citation Text":["Dullemond & Dominik 2004"],"Functions Text":["Such shadows reveal that angular distortions are common in inner disks. This is consistent with the existence of a class of self-shadowed disks in which the inner disk scale height is higher than that of the outer disk, leading to a general shadowing and related cooling of the outer disk"],"Functions Label":["Similarities"],"Citation Start End":[[776,800]],"Functions Start End":[[486,774]]}
{"Identifier":"2021MNRAS.501.3781R___2017_Instance_2","Paragraph":"While spatially extended optical jets and bipolar CO molecular outflows have been observed in numerous Class 0\/I protostars (e.g. Reipurth & Bally 2001; Bally 2016, and references therein), near-infrared high-resolution spectroscopy and spectroimaging observations in the past two decades have made it possible to study the kinematics of the outflowing gas and physical properties at the base of the jet within a few hundred au of the driving source in Class 0\/I protostars (e.g. Davis et al. 2001, 2003, 2011; Nisini et al. 2005, 2016; Caratti o Garatti et al. 2006; Takami et al. 2006; Antoniucci et al. 2008, 2011, 2017; Garcia Lopez et al. 2008, 2013). These microjets are bright in [Fe\u2009ii] forbidden and H2 rovibrational emission lines, hence showing the presence of forbidden emission-line (FEL) regions and molecular hydrogen emission-line (MHEL) regions in low-mass Class 0\/I protostars. While multiple low- and high-velocity components are observed in both MHELs and FELs, the higher velocity gas is slightly further offset from the driving source than the slower gas, and the kinematics of the H2 emission differs from [Fe\u2009ii] emission, revealing complicated kinematic structures. Evidence of H2 emission from cavity walls is also seen in some protostars, suggesting the presence of a wide-angled wind. Strong emission in the well-known accretion diagnostics of Paschen and Brackett hydrogen recombination lines is observed in protostars, with the ratio of the accretion luminosity to bolometric luminosity spanning from \u223c0.1 to \u223c1. The mass accretion and loss rates for Class 0\/I low-mass protostars span the range of 10\u22126\u201310\u22128 M\u2299 yr\u22121, and the derived jet efficiencies (ratio between mass ejection and accretion rates) range between \u223c1 per\u2009cent and 10 per\u2009cent (e.g. Davis et al. 2001, 2003, 2011; Nisini et al. 2005, 2016; Caratti o Garatti et al. 2006; Takami et al. 2006; Antoniucci et al. 2008, 2011, 2017; Garcia Lopez et al. 2008, 2013). These measurements are within the range predicted by the magnetohydrodynamic jet launching models (e.g. Frank et al. 2014).","Citation Text":["Antoniucci et al.","2017"],"Functions Text":["The mass accretion and loss rates for Class 0\/I low-mass protostars span the range of 10\u22126\u201310\u22128 M\u2299 yr\u22121, and the derived jet efficiencies (ratio between mass ejection and accretion rates) range between \u223c1 per\u2009cent and 10 per\u2009cent (e.g.","These measurements are within the range predicted by the magnetohydrodynamic jet launching models"],"Functions Label":["Background","Similarities"],"Citation Start End":[[1887,1904],[1917,1921]],"Functions Start End":[[1543,1778],[1956,2053]]}
{"Identifier":"2021ApJ...916...57G__Gedalin_2017_Instance_1","Paragraph":"Collisionless shocks are one of the most ubiquitous phenomena in space plasmas. The directed flow energy is converted into other forms at the shock front: ion and electron heating, particle acceleration, and magnetic field enhancement. One of the most important problems of the collisionless shock physics is prediction of the post-shock (downstream) state of the plasma and fields given the corresponding plasma and magnetic field state before the shock (upstream). In the absence of dissipative processes the relations between the upstream and downstream parameters (Rankine\u2013Hugoniot relations, boundary conditions, and jump conditions) are just conservation laws of particle number, momentum, and energy. There is a complete hierarchy of scales in collisionless shocks and, accordingly, the boundary conditions depend on the scale at which the boundaries are placed. The standard Rankine\u2013Hugoniot relations (RH) are formulated on the largest, magnetohydrodynamic (MHD) scale, at which the distributions thermalize (Kennel 1994). In many space plasma environments the ambient conditions change on spatial or temporal scales, which are too short to achieve such thermalization. Observational comparison of the upstream and downstream plasma is often done in various regions that do not satisfy the conditions for establishing the MHD RH. In the close vicinity of the shock front, ion distributions are significantly gyrophase dependent, and the jump conditions at the very transition should take this into account (Gedalin & Balikhin 2008; Gedalin 2016a). On larger scales or for measurements invoking substantial temporal and\/or spatial averaging, the distributions become gyrotropic but can still remain anisotropic, since isotropization may be slow. The corresponding RH should take into account this anisotropy (Lyu & Kan 1986; Gedalin 2017; Gedalin et al. 2020). Additional complications arise when there are different populations of ions that undergo gyrotropization and isotropization on different scales. Eventual equilibration of temperatures may never happen. Even for different populations of ions of the same kind eventual merging into one thermal population may never be observed. The latter situation is expected to be typical when pickup ions (PUIs) are important constituents, as at the termination shock (Zank et al. 1996; Li et al. 2008; Richardson et al. 2008; Burrows et al. 2010; Matsukiyo & Scholer 2011, 2014; Ariad & Gedalin 2013; Jokipii 2013; Mostafavi et al. 2017, 2018; Kumar et al. 2018). In this case, MHD has to be replaced with a multispecies model, where each population obeys the corresponding conservation laws separately, while the magnetic field jump is obtained by combining all species together. In this approach one has to know the equations of state for each species. Typically, some form of thermodynamics motivated state equations are assumed (Florinski et al. 2009; Borovikov et al. 2011; Pogorelov et al. 2013; Zank et al. 2014; Mostafavi et al. 2018). This approach implicitly assumes that the population is in a kind of local thermodynamic equilibrium. However, observations show that this is not the case. In fact, the downstream ion distribution is determined primarily by ion dynamics in the macroscopic fields of the shock front. Interaction with the electromagnetic fluctuations plays a secondary role causing slow relaxation toward equilibrium, while binary collisions are negligible. The resolution of particle observations in the vicinity of the heliospheric termination shock corresponds to the scale on which the solar wind protons are possibly (but not certainly) isotropic but PUI are still anisotropic. Yet most studies of RH at the termination shock aim at the determination of jump conditions on the MHD scale, assuming isotropy of the distributions. Some previous analyses within a two-fluid model (solar wind and PUI) introduced heat flux and collisionless viscosity of PUI (Zank et al. 2014; Mostafavi et al. 2016, 2017, 2018; Zank 2016). This approach explicitly assumes that ion scattering within the shock transition is substantial so that the focused transport equation is valid throughout the shock (Zank et al. 2014). However, the shock transition is scatter-free (Toptyghin 1980; Drury 1983), so that the approach loses the physics of relation between the kinetic and MHD scales. Accordingly, the obtained magnetic profiles are monotonic and cannot explain observed overshoots (Mostafavi et al. 2018). Previous attempts to take into account the ion kinetics at a shock front assumed magnetic moment conservation (Fahr & Chalov 2008; Fahr et al. 2012; Fahr & Siewert 2013, 2015). However, it has been shown that magnetic moment is, in general, not conserved (Terasawa 1979; Gedalin 2020), and such approximation may be satisfactory only in the perpendicular regime. Proper establishment of the boundary conditions on MHD scales requires establishing a relation of the ultimately isotropic distributions to the collisionless ion dynamics within a shock front. This objective requires proper determination of downstream gyrophase averaged distributions, which are gyrotropic but anisotropic, with further relating these to the isotropic pressure (Gedalin et al. 2020). In the present paper, we numerically determine the parameters of these distributions, which will be used for modifying the Rankine\u2013Hugoniot relation.","Citation Text":["Gedalin 2017"],"Functions Text":["On larger scales or for measurements invoking substantial temporal and\/or spatial averaging, the distributions become gyrotropic but can still remain anisotropic, since isotropization may be slow. The corresponding RH should take into account this anisotropy"],"Functions Label":["Background"],"Citation Start End":[[1833,1845]],"Functions Start End":[[1557,1815]]}
{"Identifier":"2020ApJ...899....4X__Zhou_et_al._2019_Instance_1","Paragraph":"It is also very encouraging to note that recent analyses of heavy-ion reaction experiments in terrestrial laboratories and properties of NSs from multiple messengers have led to some new progress in constraining the Esym(\u03c1) up to about twice the saturation density. For example, shown in Figure 1 are the values of symmetry energy at 2\u03c10, that is, Esym(2\u03c10), from (1) the FOPI-LAND (Russotto et al. 2011) and (2) the ASY-EOS (Russotto et al. 2016) Collaborations by analyzing the relative flows and yields of light mirror nuclei as well as neutrons and protons in heavy-ion collisions at beam energies of 400 MeV\/nucleon, (3) (Chen) an extrapolation of the systematics of low-density symmetry energy (Chen 2015), (4) (Zhang & Li) direct inversions of observed NS radii, tidal deformability, and maximum mass in the high-density EOS space (Zhang et al. 2018; Zhang & Li 2019a, 2020a), (5) (Xie & Li) a Bayesian inference from the radii of canonical NSs observed by using X-rays and gravitational waves from GW170817 (Xie & Li 2019), (6) (Zhou et al.) analyses of NS radii, tidal deformability, and maximum mass within an extended Skyrme Hartree\u2013Fock approach (eSHF) (Zhou & Chen 2019; Zhou et al. 2019), (7) (Nakazato & Suzuki) analyzing cooling timescales of protoneutron stars as well as the radius and tidal deformability of GW170817 (Nakazato & Suzuki 2019), and (8) a Bayesian inference directly from the X-ray data of seven quiescent low-mass X-ray binaries in globular clusters (Baillot d\u2019Etivaux et al. 2019). Despite the rather different assumptions and methods used in analyzing the different types of laboratory and observational data, it is very interesting to see that they all together are consistent with a fiducial value of Esym(2\u03c10) = 47 MeV within the still relatively large error bars of the individual analyses. Moreover, several recent theoretical studies also predicted values of Esym(2\u03c10) consistent with its fiducial value of 47 MeV. For example, an upper bound of Esym(2\u03c10) \u2264 53.2 MeV was derived recently by Tong et al. (2020) by studying the radii of neutron drops using the state-of-the-art nuclear energy density functional theories. Quantum Monte Carlo calculations using local interactions derived from chiral EFT up to next-to-next-to-leading order predicted a value of Esym(2\u03c10) \u2248 46 \u00b1 4 MeV (Lonardoni et al. 2020), while the latest many-body perturbation theory calculations with consistent nucleon-nucleon and three-nucleon interactions up to fourth order in the EFT expansion predicted a value of Esym(2\u03c10) \u2248 45 \u00b1 3 MeV (Drischler et al. 2020). They are both consistent with the fiducial value of \n\n\n\n\n\n and have much smaller error bars. It is worth noting that the chiral EFT is currently applicable to a maximum density of about 2\u03c10.","Citation Text":["Zhou et al. 2019"],"Functions Text":["For example, shown in Figure 1 are the values of symmetry energy at 2\u03c10, that is, Esym(2\u03c10), from","(6)","analyses of NS radii, tidal deformability, and maximum mass within an extended Skyrme Hartree\u2013Fock approach (eSHF)","Despite the rather different assumptions and methods used in analyzing the different types of laboratory and observational data, it is very interesting to see that they all together are consistent with a fiducial value of Esym(2\u03c10) = 47 MeV within the still relatively large error bars of the individual analyses."],"Functions Label":["Uses","Uses","Uses","Similarities"],"Citation Start End":[[1184,1200]],"Functions Start End":[[266,363],[1032,1035],[1050,1164],[1517,1830]]}
{"Identifier":"2017ApJ...834..135T__Zolotov_et_al._2015_Instance_1","Paragraph":"To explain the morphological transformation, two main evolutionary paths have been proposed in the literature. A slow cosmological path naturally follows from the strong redshift evolution of galaxy sizes, \n\n\n\n\n\n (Mosleh et al. 2012; Newman et al. 2012; van der Wel et al. 2014b; Shibuya et al. 2015). Star-forming galaxies quench star formation and add to the passive population with approximately the same size in a later epoch (van Dokkum et al. 2015; Lilly & Carollo 2016). A second, fast path involves a downward transition in the mass\u2013size plane, at approximately constant redshift (Barro et al. 2013, 2014; Dekel & Burkert 2014; Zolotov et al. 2015). This process requires a substantial \u201ccompaction\u201d of the formally extended star-forming galaxies. One possible mechanism would be a major merger, which is known from observations and simulations to lead to substantial angular momentum redistribution, orbit reconfiguration, and mixing (Mihos & Hernquist 1996; Wuyts et al. 2010). Another possibility is an internal angular momentum redistribution within the star-forming disk. This process has been considered to be effective at high redshift (Noguchi 1999; Immeli et al. 2004a, 2004b; Elmegreen et al. 2008; Genzel et al. 2008; Bournaud et al. 2011), when galaxies are gas rich (Tacconi et al. 2013) and effective viscous dissipation leads to radial inward transport of gas and stars with a timescale of a few hundred megayears (Dekel et al. 2009) and buildup of a central dense core (bulge component) through circumnuclear concentration of gas. Nelson et al. (2016b) find in massive galaxies at z \u223c 1.4 that central 1 kpc regions are highly attenuated by dust and are responsible for half of the total star formation rate (SFR). In conjunction with morphological quenching (Martig et al. 2009; Genzel et al. 2014a), and powerful AGN outflows (Bower et al. 2006; Croton et al. 2006; F\u00f6rster Schreiber et al. 2014; Genzel et al. 2014b), the compaction process may then lead to an inside-out quenching near the Schechter mass (Tacchella et al. 2015, 2016).","Citation Text":["Zolotov et al. 2015"],"Functions Text":["A second, fast path involves a downward transition in the mass\u2013size plane, at approximately constant redshift","This process requires a substantial \u201ccompaction\u201d of the formally extended star-forming galaxies."],"Functions Label":["Background","Background"],"Citation Start End":[[636,655]],"Functions Start End":[[478,587],[658,754]]}
{"Identifier":"2018ApJ...865...22C__Criscuoli_et_al._2013_Instance_1","Paragraph":"The strong linearity between the two indices shown in Figure 5 must be ascribed to their strong temperature dependence and to the proximity of the Mg ii core and FUV formation heights on one side, and Mg ii continuum and MUV on the other. Most of the radiation contributing to the FUV and MUV originates in the higher photosphere and chromosphere (Thuillier et al. 2012), where atmosphere models of quiet and magnetic structures present similar (and rather small) temperature and density gradients (c.f.r. Figure 1 in Fontenla et al. 2011). This means that, at least at first order, an increase of the activity corresponds to an increase of equal amount of the temperature and density from which both MUV and FUV, as well as Mg ii wings and core, originate. Under the assumption that the lambda-integrated fluxes can be described by a Plank function, it is easy to show that the color index must scale linearly with the temperature, at least for small temperature variations. Therefore, since the spectral brightness variations over a magnetic cycle do not exceed a few tens of K (Fontenla et al. 2011), the [FUV\u2013MUV] color index must also scale linearly with the activity. In LTE, a linear dependence of the core\/wing ratio is expected (e.g., Criscuoli et al. 2013), but such dependence is probably less obvious for the Mg ii h and k line, whose cores form in NLTE conditions. Nevertheless, collisions play an important role in the formation of the line (Linsky & Avrett 1970), making the core of this doublet strongly sensitive to temperature (e.g., Carlsson et al. 2015). Indeed, numerical magnetohydrodynamic simulations indicate that the brightness temperature derived from the Mg ii h and k cores is a good approximation for the plasma temperature at the corresponding formation heights (Leenaarts et al. 2013b). Moreover, the Bremen Mg ii index is derived from data of relatively modest spectral resolution, so that the cores are not well resolved. This means that the core intensities used to construct the Mg ii index must have a non-negligible contribution from several layers of the atmosphere, from the lower photosphere, where the response function of the spectral region between the two h and k peaks reaches its maximum (Uitenbroek 1997), to the base of the transition region (Thuillier et al. 2012; Leenaarts et al. 2013b; Carlsson et al. 2015), where the cores form. Employing again a Plank function to describe both the core and the wing intensities of the Mg ii doublet it is easy to show that the Mg ii index must scale linearly with temperature, at least for small temperature perturbations.","Citation Text":["Criscuoli et al. 2013"],"Functions Text":["In LTE, a linear dependence of the core\/wing ratio is expected (e.g.,","but such dependence is probably less obvious for the Mg ii h and k line, whose cores form in NLTE conditions."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1244,1265]],"Functions Start End":[[1174,1243],[1268,1377]]}
{"Identifier":"2021ApJ...923..105L__Owens_et_al._2013_Instance_2","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"],"Functions Text":["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"],"Functions Label":["Motivation"],"Citation Start End":[[2877,2894]],"Functions Start End":[[2667,2857]]}
{"Identifier":"2015ApJ...799...55G__Rosenvinge_et_al._2008_Instance_1","Paragraph":"We use data from instruments on board STEREO-A, STEREO-B, ACE, SOHO, Wind, the Geostationary Operational Environmental Satellites (GOES), MESSENGER and the Solar Dynamics Observatory (SDO). Remote-sensing observations of the CME and the activity phenomena over the solar surface were provided by the Sun Earth Connection Coronal and Heliospheric Investigation (SECCHI) instrument suite on board STEREO (Howard et al. 2008), the Large Angle and Spectrometric Coronagraph experiment (LASCO) on board SOHO (Brueckner et al. 1995), the Atmospheric Imaging Assembly (AIA; Lemen et al. 2012) on board SDO, and the X-ray telescopes on board GOES. Synoptic maps including PFSS model results provided by the Global Oscillation Network Group (GONG, http:\/\/gong.nso.edu\/) and by Solarsoft PFSS package were also examined. Radio observations were provided by the S\/WAVES (Bougeret et al. 2008) investigation on board STEREO and by the WAVES (Bougeret et al. 1995) experiment on board Wind. In situ energetic particle observations were provided by the Solar Electron and Proton Telescope (SEPT; M\u00c3\u00bcller-Mellin et al. 2008), the Low Energy Telescope (LET; Mewaldt et al. 2008), and the High Energy Telescope (HET; von Rosenvinge et al. 2008) on board STEREO (all of them part of the IMPACT instrument suite; Luhmann et al. 2008), the Comprehensive Suprathermal and Energetic Particle Analyzer (COSTEP; M\u00c3\u00bcller-Mellin et al. 1995), and the Energetic and Relativistic Nuclei and Electron (ERNE; Torsti et al. 1995) instrument on board SOHO, the Electron, Proton, and Alpha Monitor (EPAM; Gold et al. 1998) on board ACE, the 3D Plasma and Energetic Particle Investigation (3DP) on board Wind, and the Energetic Particle and Plasma Spectrometer (EPPS; Andrews et al. 2007) on board MESSENGER. Additionally, Fe\/O ratios at ACE were obtained using data from the Solar Isotope Spectrometer (SIS; Stone et al. 1998). Finally, solar wind plasma and magnetic field data were obtained from the Plasma and Suprathermal Ion Composition (PLASTIC; Galvin et al. 2008) investigation on board STEREO, the STEREO Magnetic Field Experiment (Acu\u00c3\u00b1a et al. 2008), the ACE Magnetic Field Experiment (Smith et al. 1998), and the ACE\/Solar Wind Electron Proton Alpha Monitor (SWEPAM; McComas et al. 1998). Due to gaps in ACE\/SWEPAM data, plasma data from the Charge, Element, and Isotope Analysis System (CELIAS; Hovestadt et al. 1995) instrument on board SOHO were also used. Magnetic field data from MESSENGER were obtained from the Magnetometer Instrument on board this s\/c (Anderson et al. 2007). Interplanetary disturbance identifications were cross-checked using the STEREO and ACE level 3 event list maintained by L. Jian (http:\/\/www-ssc.igpp.ucla.edu\/forms\/stereo\/stereo_level_3.html, http:\/\/www.srl.caltech.edu\/ACE\/ASC\/DATA\/level3\/) and the near-Earth ICME list maintained by I. Richardson and H. Cane (http:\/\/www.srl.caltech.edu\/ACE\/ASC\/DATA\/level3\/icmetable2.htm; see also Cane & Richardson 2003).","Citation Text":["von Rosenvinge et al. 2008"],"Functions Text":["Radio observations were provided by","and the High Energy Telescope"],"Functions Label":["Uses","Uses"],"Citation Start End":[[1200,1226]],"Functions Start End":[[811,846],[1164,1193]]}
{"Identifier":"2021MNRAS.507.1623C__Lesgourgues_2011_Instance_1","Paragraph":"To describe the H\u2009i bispectrum in redshift space, we apply the standard redshift space kernels (see Heavens et al. 1998; Scoccimarro et al. 1999 for derivations) such that\n(10)$$\\begin{eqnarray*}\r\nB_\\rm {H\\,\\small {I}}(\\boldsymbol{k}_{1},\\boldsymbol{k}_{2},\\boldsymbol{k}_{3}) &amp;=&amp; 2\\, \\overline{T}_\\rm {H\\,\\small {I}}^2\\left[Z_1(\\boldsymbol{k}_{1}) Z_1(\\boldsymbol{k}_{2}) Z_2(\\boldsymbol{k}_{1}, \\boldsymbol{k}_{2}) P_\\text{lin}(k_{1}) P_\\text{lin}(k_{2})\\right.\\nonumber\\\\\r\n&amp;&amp;\\left. +\\, \\text{cycl.}\\right] D_\\text{FoG}(\\boldsymbol{k}_{1},\\boldsymbol{k}_{2},\\boldsymbol{k}_{3},\\sigma _\\text{B}),\r\n\\end{eqnarray*}$$where cycl. represents cyclic permutations which run over all possible pairs of $\\boldsymbol{k}_{1},\\boldsymbol{k}_{2}$ and $\\boldsymbol{k}_{3}$. Plin represents the real-space, linear matter power spectrum for which we use the CLASS Boltzmann solver (Blas, Lesgourgues & Tram 2011; Lesgourgues 2011). Z1 is given in equation (7) and Z2 denotes the second-order kernel and is given by\n(11)$$\\begin{eqnarray*}\r\nZ_2(\\boldsymbol{k}_{i}, \\boldsymbol{k}_{j}) &amp;=&amp; b_1 F_2(\\boldsymbol{k}_{i}, \\boldsymbol{k}_{j})+f \\mu _{ij}^2 G_2(\\boldsymbol{k}_{i}, \\boldsymbol{k}_{j})+\\frac{f \\mu _{ij} k_{ij}}{2}\\nonumber\\\\\r\n&amp;&amp;\\times \\,\\left[\\frac{\\mu _i}{k_i} Z_1(\\boldsymbol{k}_{j})+\\frac{\\mu _j}{k_{j}} Z_1(\\boldsymbol{k}_{i})\\right] + \\frac{b_2}{2},\r\n\\end{eqnarray*}$$where $\\boldsymbol{k}_{ij} = \\boldsymbol{k}_{i} + \\boldsymbol{k}_{j}$ and $\\mu _{ij} = \\boldsymbol{k}_{ij}\\cdot \\hat{{\\bf z}}\/k_{ij}$. F2 and G2 denote the second-order kernels for the real-space density and velocity fields and are given by\n(12)$$\\begin{eqnarray*}\r\nF_2(\\boldsymbol{k}_{i}, \\boldsymbol{k}_{j}) = \\frac{5}{7} + \\frac{m_{ij}}{2}\\left(\\frac{k_i}{k_{j}} + \\frac{k_{j}}{k_i}\\right) + \\frac{2}{7} m_{ij}^2,\r\n\\end{eqnarray*}$$(13)$$\\begin{eqnarray*}\r\nG_2(\\boldsymbol{k}_{i}, \\boldsymbol{k}_{j}) = \\frac{3}{7} + \\frac{m_{ij}}{2}\\left(\\frac{k_i}{k_{j}} + \\frac{k_{j}}{k_i}\\right)+\\frac{4}{7} m_{ij}^2 ,\r\n\\end{eqnarray*}$$where $m_{ij} = (\\boldsymbol{k}_{i}\\cdot \\boldsymbol{k}_{j})\/(k_i k_{j})$. The final term in equation (10), DFoG, is a phenomenological factor to address some non-linear RSD effects not sufficiently modelled by the redshift kernels alone. On smaller scales, internal motion inside virialized structures produces a radial smearing to the density field in redshift space, known as the fingers-of-god (FoG) effect (Jackson 1972). It is common to include a term which describes the FoG (Taruya, Nishimichi & Saito 2010), even when including higher order perturbation theory terms and should be seen as a phenomenological damping required to correct for non-linear effects (Verde et al. 1998; Gil-Mar\u00edn et al. 2014). For our choice of model, this factor is given by (Gil-Mar\u00edn et al. 2015)\n(14)$$\\begin{eqnarray*}\r\nD_\\text{FoG}(\\boldsymbol{k}_{1},\\boldsymbol{k}_{2},\\boldsymbol{k}_{3},\\sigma _\\text{B}) = \\left[1 + \\frac{1}{2}\\left(k_{1}^2 \\mu _1^2 + k_{2}^2\\mu _2^2 + k_{3}^2\\mu _3^2\\right)^2 \\sigma _\\text{B}^2\\right]^{-2} . \\nonumber\\\\\r\n\\end{eqnarray*}$$It is worth noting that it is necessary for galaxy surveys to also include modelling of shot noise caused by discreteness effects in their bispectra analyses. However, for H\u2009i IM, where unresolved signal is integrated over, this should not be a limiting factor (Spinelli et al. 2020). We therefore do not consider any treatment of shot noise in our analysis, making the assumption that this should be very low in H\u2009i IM observations.","Citation Text":["Lesgourgues 2011"],"Functions Text":["Plin represents the real-space, linear matter power spectrum for which we use the CLASS Boltzmann solver"],"Functions Label":["Uses"],"Citation Start End":[[915,931]],"Functions Start End":[[778,882]]}
{"Identifier":"2017ApJ...840...91Y__Shatsky_&_Tokovinin_2002_Instance_1","Paragraph":"Another binary parameter is the stellar mass ratio distribution. The choice of the mass ratio can be called the pairing function, which is combining stars into binary systems. When random pairing is used (i.e., binary companions are randomly chosen from a given IMF), we find similar mass segregation results to those for star clusters containing only single stars. However, random pairing is ruled out observationally, and there is also a lack of theoretical backing for a random pairing function (Shatsky & Tokovinin 2002; Kouwenhoven et al. 2005, 2007a, 2007b). Indeed, massive stars preferentially choose other massive stars as their binary companions (Kouwenhoven et al. 2010). This implies that both the geometric average and the median MST method may be biased by the short MST edges from binary systems, although the LnND method is less influenced. An alternative binary pairing function is the so-called ordered pairing (i.e., mass ratio of the binary star components \u223c1, in particular for massive stars; see e.g., Kouwenhoven et al. 2009; Oh & Kroupa 2012), which matches stars in order of their mass distribution. We re-pair our test model using ordered pairing and find some extremely high values of the mass segregation degree when using the MST method. This is illustrated in Figure 5, where we see that the mass segregation degree for ordered pairing functions is biased toward very high values for relatively small massive-star samples. These extremely high degrees of mass segregation are not physical, because they are mainly produced by the short MST edges in the binary systems. This is illustrated by the sharp steps in the solid curves, which greatly deviates from the non-mass-segregation distribution equal to unity. Therefore, the gmMST and mMST methods should not be considered as a global compactness measure of massive stars in cases where binary fractions are significant, i.e., >10%. Massive stars may have much shorter MST edges than low-mass stars, due to the preferential presence of binaries among massive stars (e.g., Sana et al. 2014).","Citation Text":["Shatsky & Tokovinin 2002"],"Functions Text":["When random pairing is used (i.e., binary companions are randomly chosen from a given IMF), we find similar mass segregation results to those for star clusters containing only single stars. However, random pairing is ruled out observationally, and there is also a lack of theoretical backing for a random pairing function"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[499,523]],"Functions Start End":[[176,497]]}
{"Identifier":"2020MNRAS.492..708M__Davidson_&_Fesen_1985_Instance_1","Paragraph":"Due to the high luminosity and seemingly long-term flux stability, the Crab pulsar and its pulsar wind nebula (PWN) is one of the most studied sources in the very high energy (VHE, $\\rm {E} \\gt 100$\u2009GeV) regime. For many years, Crab has been used as a standard candle in X- and gamma-ray astronomy (Hester 2008). The Crab nebula was the first TeV gamma-ray source discovered (in 1989 by the Whipple 10-m telescope; Weekes et al. 1989), and soon after detected by numerous facilities above 100\u2009GeV (Smith et al. 2000; Aharonian et al. 2006; Abdo et al. 2012; Aleksi\u0107 et al. 2015; Meagher & VERITAS Collaboration 2015; Abeysekara et al. 2017). It has a (energy-dependent) angular size of \u223c0.1\u00b0 and its distance has been estimated to be \u22482.2\u2009kpc, corresponding to a physical size of \u22483.8\u2009pc (Trimble 1973; Davidson & Fesen 1985; Kaplan et al. 2008). The nebula non-thermal spectrum can be described by two components, a synchrotron component extending from radio to high energy gamma-rays and a second component emerging above 1\u2009GeV (Atoyan & Aharonian 1996). The latter is interpreted as inverse Compton scattering (IC) of the same particles against soft background photons: cosmic microwave background (CMB), far-infrared (FIR), and near-infrared (NIR) background, and the synchrotron photons of the nebula itself or Synchrotron Self-Compton (SSC). The pulsar has a spin period of $\\rm {P} = 33$\u2009ms, a spin-down rate of $\\rm {\\dot{P}} = 4.21\\times 10^{-13}$ and a spin-down luminosity of $\\rm {L_{spin}} = 3.8\\times 10^{38}\\ \\rm {erg}\\ \\rm {s^{-1}}$. The pulsed emission between 0.1\u2009 and 100\u2009GeV is believed to be due to synchrotron-curvature radiation (Abdo et al. 2010; Ansoldi et al. 2016) and its spectrum is well parametrized by a power law with a sub-exponential cut-off function of spectral index \u03b3P = 1.59, the break located at an energy of about 500\u2009MeV and curvature index of \u03ba = 0.43. In addition, a power-law component emerges above the cut-off extending above 100\u2009GeV (Aliu et al. 2011; Aleksi\u0107 et al. 2012; Ansoldi et al. 2016).","Citation Text":["Davidson & Fesen 1985"],"Functions Text":["It has a (energy-dependent) angular size of \u223c0.1\u00b0 and its distance has been estimated to be \u22482.2\u2009kpc, corresponding to a physical size of \u22483.8\u2009pc"],"Functions Label":["Background"],"Citation Start End":[[803,824]],"Functions Start End":[[642,787]]}
{"Identifier":"2022ApJ...935..135B__Thomas_et_al._2019_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":["Thomas et al. 2019"],"Functions Text":["At large galactocentric radii (>10 kpc) this includes, among others,","and streams of stars kicked up from the disk that undergo phase mixing, sometimes referred to as \u201cfeathers\u201d (e.g.,"],"Functions Label":["Background","Background"],"Citation Start End":[[1044,1062]],"Functions Start End":[[761,829],[903,1017]]}
{"Identifier":"2015AandA...574A..50J__Takeda_et_al._(2008)_Instance_1","Paragraph":"Several studies that analyzed the metallic content in the atmospheres of evolved stars hosting planets have been published in the last decade. The first results were based on small samples and\/or the abundances were not obtained with a homogeneous technique. Schuler et al. (2005) derived the metallicity for one GWP and gathered abundances from the literature for another seven. They reported that, on average, GWP were metal-poor compared with planet hosting dwarfs. Similar results were found by Sadakane et al. (2005) analyzing 4 GWP. In 2007, Pasquini et al. studying 14 GWP (4 from the literature), concluded that in contrast to the distribution of main-sequence stars with planets, the GWP distribution does not favor high metallicity objects. Conversely, Hekker & Mel\u00e9ndez (2007), analyzing a sample of 380 GK giant stars including 20 with planets (15 from the literature), found an enhancement for GWP of 0.13 dex compared with stars without planets. Takeda et al. (2008), with a sample of 322 giants, including 10 planet-hosts, did not find any metallicity offset. Ghezzi et al. (2010a) found that the metallicity distribution of 16 GWP displays an average that is 0.17 dex more metal-poor than the sample of 117 planet-hosting dwarfs and, furthermore, that the subgiant sample is more metal-rich by 0.12 dex. Johnson et al. (2010) ruled out a flat metallicity relationship among their sample of 246 subgiants (36 with planets, including unpublished candidates). More recently, with larger samples, Mortier et al. (2013) did not find any metallicity enhancement in 71 evolved stars with planets in comparison with 733 evolved stars without planets, with metallicity values gathered from the literature. Finally, Maldonado et al. (2013) found a metallicity enhancement in 16 subgiants with planets relative to 55 without planets (50 from literature). These authors did not find evidence of a metallicity offset between giants with and without planets, analyzing 43 GWP and 67 GWOP. However, for stars with masses above 1.5 M\u2299, they reported a slight metallicity enhancement for giants with planets relative to the control sample. ","Citation Text":["Takeda et al. (2008)"],"Functions Text":["with a sample of 322 giants, including 10 planet-hosts, did not find any metallicity offset."],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[960,980]],"Functions Start End":[[982,1074]]}
{"Identifier":"2022MNRAS.509..567S__Spinoso_et_al._2017_Instance_1","Paragraph":"When a galaxy is isolated, star formation quenching is mostly driven by internal processes (Larson 1974; Dekel & Silk 1986; Vulcani et al. 2021). Active galatic nuclei (AGNs) feedback create an outflow of gas preventing further hot gas accretion from the circumgalactic medium (Dalla Vecchia & Schaye 2008; Bongiorno et al. 2016; Trussler et al. 2020). Star formation depends on gravitational instabilities, which may be prevented due to a transition from a stellar disc to a spheroid (\u2018morphological quenching\u2019; Martig et al. 2009). Bars in spiral galaxies may drive gas inflows, which enhance central star formation (bar-driven evolution; Spinoso et al. 2017). However, in the local Universe most of the galaxies live in groups\/clusters (Geller & Huchra 1983). Even before crossing the virial radius, infalling galaxies can stop accreting new gas (\u2018Starvation\u2019; Larson, Tinsley & Caldwell 1980; Balogh, Navarro & Morris 2000; van de Voort et al. 2017). For instance, Trussler et al. (2020) suggest galaxy quenching has an extended phase (\u223c5\u2009Gyr) of starvation. In addition, infalling galaxies can lose gas, stars and even dark matter via gravitational tides (\u2018tidal mass-loss\u2019 - TML; Johnston, Sigurdsson & Hernquist 1999; Read et al. 2006). Within the virial radius, the hot gas in the intracluster medium (ICM) exerts pressure on galaxies moving within the cluster and may remove gas via Ram Pressure Stripping (RPS; Gunn & Gott 1972; Abadi, Moore & Bower 1999). Clusters provide a suitable environment for both direct and indirect galaxy interactions, especially in its core. Direct encounters may lead to galaxy mergers and cause a starburst event over a short time-scale and quickly exhaust a galaxy\u2019s gas supply (Springel & Hernquist 2005; Cox et al. 2008; Teyssier, Chapon & Bournaud 2010), whereas repeated indirect encounters may leave interacting galaxies with disturbed morphologies. At last, it is important to note that clusters are built up by the accretion of galaxy groups. Cluster galaxies may therefore be affected by \u2018pre-processing\u2019, in which galaxy properties were altered even before entering the cluster (Fujita 2004; Mahajan 2013; Sarron et al. 2019).","Citation Text":["Spinoso et al. 2017"],"Functions Text":["Bars in spiral galaxies may drive gas inflows, which enhance central star formation (bar-driven evolution;"],"Functions Label":["Background"],"Citation Start End":[[641,660]],"Functions Start End":[[534,640]]}
{"Identifier":"2017AandA...599A..13Y__Bu\u010d\u00edk_et_al._(2009)_Instance_1","Paragraph":"Figure 2 shows solar wind plasma and magnetic field measurements for a CIR that occurred between July 26 and 27, 2003 (days of year 207\u2013208). Following Chotoo et al. (2000), Richardson et al. (1993), we marked four regions in the plot: the slow wind region (S), the compressed slow wind region (S\u2032), the compressed fast wind region (F\u2032), and the fast wind itself (F). Throughout four regions, the mean charge states of iron measured by ACE\/SWICS lies around 11+, consistent with typical values in the solar wind (Lepri et al. 2001). The stream interface (S\u2032-F\u2032) is indicated by the vertical line in Fig. 2 and is characterized by a drop of the O7+\/O6+ abundance ratio measured with SWICS in the bulk solar wind (Wimmer-Schweingruber et al. 1997, 1999). The leading (S-S\u2032) and trailing edge (F\u2032-F) of the CIR were determined by the total pressure (Jian et al. 2006). Bu\u010d\u00edk et al. (2009) found that CIR boundaries can be well defined when the total pressure exceeds 50 pPa (indicated by the horizontal dashed line in Fig. 2), which is slightly higher than that in the background solar wind, which typically is 20\u221230 pPa, according to Jian et al. (2006). The total pressure P was obtained from the sum of plasma and magnetic field pressure, that is, \\hbox{$P=n_{\\rm p}v^{2}_{\\rm th}m+B^2\/2\\mu_0$}P=npvth2m+B2\/2\u03bc0, where np and vth are the proton density and thermal speed, respectively, and B is the magnitude of the magnetic field. Because SOHO has no magnetometer, we used magnetic field data from ACE\/MAG (which is also around L1). Comparing plasma parameters (bulk speed, thermal speed, and proton density) measured by PM with those of SWEPAM, we see that the physical conditions at SOHO and ACE were almost the same, and that the time difference between passages of the CIR boundaries is less than ten minutes. The CIR shown in Fig. 2 was bounded by a reverse shock (vertical line separating F\u2032 from F). We clearly see that the suprathermal He++ intensity peaks inside the decelerated and compressed fast-wind region (F\u2032), close to the reverse shock. In contrast, suprathermal particles are very rare in the S and S\u2032 regions. After passage of the reverse shock, suprathermal particles continue to be observable for more than one day. They are commonly believed to be the sunward particles accelerated by the reverse shock far beyond the Earth orbit. In other words, the observer saw the duration of the CIR particle event, which was longer than that of the CIR compression region itself. The background level shown in green was estimated using the method described above. The signal-to-noise ratio (S\/N) in the F and F\u2032 regions is higher than 100, confirming that our observations are due to real He++ particles. ","Citation Text":["Bu\u010d\u00edk et al. (2009)"],"Functions Text":["found that CIR boundaries can be well defined when the total pressure exceeds 50 pPa (indicated by the horizontal dashed line in Fig. 2), which is slightly higher than that in the background solar wind, which typically is 20\u221230 pPa, according to Jian et al. (2006)."],"Functions Label":["Uses"],"Citation Start End":[[866,885]],"Functions Start End":[[886,1151]]}
{"Identifier":"2021MNRAS.506.5468Z__Begeman,_Broeils_&_Sanders_1991_Instance_1","Paragraph":"Since the free-fall time-scale tff depends on the gravity law, we study how changing the gravity law affects the value of n in the KS law. Our main aim in this paper is to derive for the first time the KS law from a basic description in the framework of Milgromian dynamics (MOND; Milgrom 1983; Famaey & McGaugh 2012). MOND is an alternative approach to a cold dark matter-dominated universe deduced from the flattening of observed rotation curves of spiral galaxies under the assumption of Newtonian dynamics. In MOND, these dynamical discrepancies are addressed by a generalization of Newtonian gravity (for a thorough review, see e.g. Famaey & McGaugh 2012). Within the classical MOND framework, the Newtonian gravitational acceleration gN is replaced in the spherically symmetric case by $g=\\sqrt{g_{\\rm {N}} a_0}$ when the gravitational acceleration is far smaller than the critical acceleration a0 = 1.2 \u00d7 10\u221210 m s\u22122 (Begeman, Broeils & Sanders 1991; Gentile, Famaey & de Blok 2011). In less symmetric configurations, the equations of motion are derived from a Lagrangian, yielding standard equations of motion but with a generalized Poisson equation for the gravitational field (Bekenstein & Milgrom 1984; Milgrom 2010). MOND predicted the very tight radial acceleration relation (RAR) between the gravity g implied by disc galaxy rotation curves and the Newtonian gravity gN, resulting from their baryonic distribution (Lelli et al. 2017; Li et al. 2018). The RAR is also evident in stacked galaxy\u2013galaxy weak lensing measurements that probe out to larger radii (Milgrom 2013; Brouwer et al. 2021). The external field effect (EFE) predicted by MOND (Bekenstein & Milgrom 1984) and required for consistency with data on Solar neighbourhood wide binaries (Banik & Zhao 2018c; Pittordis & Sutherland 2019) has recently been confirmed at high significance by comparing galaxies in isolated and more crowded environments (Haghi et al. 2016; Chae et al. 2020). Detailed numerical simulations of disc galaxy secular evolution in MOND have recently been conducted for the case of M33 (Banik et al. 2020) and for a Milky Way or M31-like surface density (Roshan et al. 2021), while star formation has also been explored with high-resolution simulations (Renaud, Famaey & Kroupa 2016). The possible cosmological context has been explored in e.g. Haslbauer, Banik & Kroupa (2020) and Asencio, Banik & Kroupa (2021). The MOND corrections to Newtonian gravity might be capturing effects of the quantum vacuum (Milgrom 1999; Smolin 2017; Verlinde 2017; Senay, Mohammadi Sabet & Moradpour 2021).","Citation Text":["Begeman, Broeils & Sanders 1991"],"Functions Text":["Within the classical MOND framework, the Newtonian gravitational acceleration gN is replaced in the spherically symmetric case by $g=\\sqrt{g_{\\rm {N}} a_0}$ when the gravitational acceleration is far smaller than the critical acceleration a0 = 1.2 \u00d7 10\u221210 m s\u22122"],"Functions Label":["Background"],"Citation Start End":[[925,956]],"Functions Start End":[[662,923]]}
{"Identifier":"2020AandA...644A.108V__Birrer_&_Amara_2018_Instance_1","Paragraph":"The most commonly used technique to create a mock lens system from simulated galaxies is to extract a mass map from a particle-based simulation and use it to calculate lensing quantities (i.e. lensing potential and its first and second derivatives) needed to emulate the gravitationally lensed images. For this purpose, galaxies from high resolution hydrodynamical simulations, including, for example, EAGLE (Evolution and Assembly of GaLaxies and their Environments, Schaye et al. 2015; Crain et al. 2015) or Illustris (Vogelsberger et al. 2014b,a), have been widely used. Different types of software, such as lenstronomy (Birrer & Amara 2018) and GLAMER (Gravitational Lensing Simulations with Adaptive Mesh Refinement, Metcalf & Petkova 2014), can handle the inference of lensing quantities from mass maps using fast Fourier transform convolution. Fast Fourier is a commonly used technique to speed up the calculation of lensing quantities, which imply computationally expansive numerical integration but it remains a demanding procedure (Metcalf & Petkova 2014; Plazas 2020). One could wonder what mass map resolution should be used and what size of map is relevant to be sufficiently precise in the mock creation while minimising the computational time. One generally considers that a strongly lensed system is determined by the projected mass inner to the lensed images. This would suggest that a region extending over a few Einstein radii (\u03b8E) is sufficient for the simulations. However, this consists in effectively ignoring any source of shear, and\/or perturbations caused by substructures and\/or anisotropy in the mass distributions. Moreover, depending on the symmetry of the problem, cutting the mass distribution at a given radius not only automatically removes the mass beyond that radius, but it may also introduce numerical artefacts that could wrongly be attributed to properties of the examined lens mass distribution. In this paper we focus on this latter point and quantify the impact of the shape (and size) of the integration domain on the lensing quantities inference.","Citation Text":["Birrer & Amara 2018"],"Functions Text":["Different types of software, such as lenstronomy","can handle the inference of lensing quantities from mass maps using fast Fourier transform convolution."],"Functions Label":["Background","Background"],"Citation Start End":[[624,643]],"Functions Start End":[[574,622],[747,850]]}
{"Identifier":"2021AandA...648A.120R__Shimajiri_et_al._2014_Instance_1","Paragraph":"We derived the mean values of the 13CO\/C18O abundance ratio for the three star-forming regions, obtaining values of 13CO\/C18O = 10.9 \u00b1 7.5 for Taurus, 13CO\/C18O = 17.3 \u00b1 11.5 for Perseus, and 13CO\/C18O = 23.7 \u00b1 10.4 for Orion. The expected value is 13CO\/C18O = 7.5\u22129.8, assuming 12C\/13C = 57\u221267 and 16O\/18O = 500\u2212600 (Gerin et al. 2015; Langer & Penzias 1990; Wilson & Rood 1994). As mentioned in Sect. 5.7, the 13CO\/C18O abundance ratio increases in regions of enhanced UV field (Shimajiri et al. 2014; Ishii et al. 2019; Areal et al. 2018). This variation is interpreted in terms of the selective photodissociation and isotopic fractionation (see Bron et al. 2018; Fuente et al. 2019). This effect is produced because the more abundant CO isotopolog shields itself from the effect of UV photons more efficiently than less abundant isotopologs (Stark et al. 2014; Visser et al. 2009). Our data confirm this trend in Orion, although the observed positions are located at a distance >1 pc from the ionized nebula M 42, which shows that the whole cloud is illuminated by a strong UV field emitted by young massive stars in the Trapezium cluster (Pabst et al. 2019). In contrast to X(C18O), which presents uniform abundance in the three regions, the mean 13CO abundance is a factor of approximately two higher in Orion than in Taurus. One interesting question pertains to whether the observed variation in the 13CO abundance is related to a variation in the 12CO\/13CO ratio (Roueff et al. 2015; Colzi et al. 2020) or is revealing variations for a higher 12CO abundance in Orion. Because of the higher gas temperature and incident UV field in Orion, selective photodissociation would work in the direction of increasing N(12CO)\/N(13CO) in Orion with respect to Taurus. Therefore, a higher 12CO abundance in Orion stands as the most likely explanation. Our results are based on significant approximations (one single phase, gas-dust thermalization), and therefore a detailed multi-transition study of these compounds is required to confirm this result.","Citation Text":["Shimajiri et al. 2014"],"Functions Text":["As mentioned in Sect. 5.7, the 13CO\/C18O abundance ratio increases in regions of enhanced UV field"],"Functions Label":["Uses"],"Citation Start End":[[481,502]],"Functions Start End":[[381,479]]}
{"Identifier":"2021AandA...653A.111R__Jones_et_al._(2021)_Instance_2","Paragraph":"As done by Le F\u00e8vre et al. (2020), we visually inspect the ancillary data, the intensity maps, the velocity and velocity dispersion fields presented in Sect. 3.1 to search for the presence of multiple components or disturbed morphology near the position of the targets. The channel maps, the spectra and the PVDs are checked together searching for consistent emission features. By taking into account the results of the initial qualitative classification by Le F\u00e8vre et al. (2020) and of the more recent quantitative analysis of a subsample of the ALPINE targets by Jones et al. (2021), we proceed with a more in-depth characterization of the [CII]-detected galaxies aimed at obtaining a robust merger fraction at z\u2004\u223c\u20045. Adopting the same criteria described in Sect. 2 to differentiate the targets and considering the S\/N of the minor merger component as described in Sect. 3.1, we find a slightly lower fraction (\u223c31%, 23 out of 75 [CII]-detected sources) of mergers7 if compared to the 40% found by Le F\u00e8vre et al. (2020), with 12, 20 and 7% of the sample made by rotating, extended and compact dispersion dominated sources, respectively. To be more conservative in the classification of the galaxies (especially for obtaining a more robust merger statistics), we define the remaining 30% of the sample as \u2018uncertain\u2019, a new category that includes the weak galaxies (as described in Le F\u00e8vre et al. 2020) and also objects that, by visual inspection, present features that are intermediate to those of various classes. This category is similar to the \u2018uncertain\u2019 (UNC) class introduced in Jones et al. (2021) that, based on the results of the 3DBarolo fits, contains sources they are unable to classify because of the low S\/N and\/or spectral resolution, or contrasting evidence in their classification criteria. Although the uncertain category is populated by a significantly larger fraction of sources with respect to the weak class (\u223c16%) in Le F\u00e8vre et al. (2020), we recover the same qualitative morpho-kinematic distribution of the previous analysis, confirming the high fraction of rotators and mergers at these early epochs.","Citation Text":["Jones et al. (2021)"],"Functions Text":["This category is similar to the \u2018uncertain\u2019 (UNC) class introduced in","that, based on the results of the 3DBarolo fits, contains sources they are unable to classify because of the low S\/N and\/or spectral resolution, or contrasting evidence in their classification criteria."],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[1590,1609]],"Functions Start End":[[1520,1589],[1610,1812]]}
{"Identifier":"2020MNRAS.499.4394M__Bate_et_al._2014_Instance_1","Paragraph":"For these four remaining FHSC candidates (L1451-mm, MC35-mm, SM1N, and Oph A-N6) that have been observed at intermediate scales (few 100 au to few 1000 au) a final confirmation of their true evolutionary state requires higher resolution observations. For L1451-mm, the compact outflow needs to be resolved to investigate its morphology and kinematics, as a higher velocity component (an indication of protostellar nature) could be revealed by observations with a beam smaller than 100 au, similar to the case of B1b-N (Hirano 2019). An additional goal of high-resolution observations for L1451-mm and the remaining youngest candidates should be to investigate the temperature and density profiles of the envelope at scales from few au to 100 au. This is because simulations show that the temperature remains lower than \u223c30 K even at several tens of au up to 100 au from the centre (Bate et al. 2014; Tomida et al. 2015; Hincelin et al. 2016; Young et al. 2019) during the FHSC stage. On the other hand, Class 0 sources show temperatures of 20\u201330 K or higher at scales of several 100 au (sufficient for thermal evaporation of CO) that results in the inner envelope and disc being easily detected using C18O observations (Yen et al. 2015, 2017; Stephens et al. 2018). This holds even in very low luminosity objects, for which the unexpected large extent of C18O is interpreted as evidence of a previous burst of accretion (Frimann et al. 2017; Hsieh et al. 2018). As for the density profile, simulations of the FHSC stage show a flat inner region, corresponding to the FHSC structure and extending up to \u223c10 au (Tomida et al. 2013; Bate et al. 2014). For a protostar, on the other hand, the density profile should increase towards the central \u223c1 au region (Young et al. 2019). Observations of the continuum emission with a resolution better than a few tens of au are likely required to model the emission and provide a density and temperature profile that can probe the relevant scales. Additional line observations with a similar resolution can also help to further distinguish between the different models. We note that, as pointed out in Young et al. (2019), distinguishing a dense core with only an FHSC and one that has recently formed a protostar but in which the FHSC structure is still present is likely not possible, even with high-resolution observations. Given the optically thick nature of the FHSC core, it is difficult to probe the physical properties within the FHSC structure. Despite this, finding a source with density and temperature profiles as well as with outflow properties consistent with the theoretical predictions will provide convincing evidence in support of a bona fide FHSC.","Citation Text":["Bate et al. 2014"],"Functions Text":["An additional goal of high-resolution observations for L1451-mm and the remaining youngest candidates should be to investigate the temperature and density profiles of the envelope at scales from few au to 100 au. This is because simulations show that the temperature remains lower than \u223c30 K even at several tens of au up to 100 au from the centre"],"Functions Label":["Future Work"],"Citation Start End":[[882,898]],"Functions Start End":[[533,880]]}
{"Identifier":"2015MNRAS.451.4290S__Torrey_et_al._2014_Instance_1","Paragraph":"Hydrodynamical simulations of evolving galaxies allow us to calibrate these diagnostics by measuring their observability given a set of formation scenarios and physical processes (e.g. Jonsson et al. 2006; Rocha et al. 2007; Lotz et al. 2008a; Bush et al. 2010; Narayanan et al. 2010; Hayward et al. 2013; Snyder et al. 2013; Lanz et al. 2014). The quality and breadth of these experiments are limited by the availability of computational resources and the fidelity of models for galaxy physics such as star formation, supernovae, and the interstellar medium (ISM). It has only recently become widespread to model the formation of galaxies ab initio (e.g. Governato et al. 2004; Agertz, Teyssier & Moore 2011; Guedes et al. 2011; Marinacci, Pakmor & Springel 2013; Ceverino et al. 2014), and the realism continues to improve (Stinson et al. 2012; Hopkins et al. 2014; Torrey et al. 2014), albeit with still widely varying physics models (e.g. Scannapieco et al. 2012; Kim et al. 2014). Prior to these advances, studies were limited to small numbers of isolated galaxies or mergers to inform common diagnostics of galaxy evolution, an approach with a significant limitation: they do not fully account for cosmological context, such as gas accretion and the breadth of assembly histories. In addition to mergers, models of high-redshift galaxy formation (e.g. Dekel, Sari & Ceverino 2009; Dekel et al. 2013) have recently appreciated the tight coupling between gas accretion and disc evolution (e.g. Cacciato, Dekel & Genel 2012; Danovich et al. 2012; Dekel & Krumholz 2013), as well as bulge and super-massive black hole (SMBH) growth mediated by turbulent motions or violent disc instability (e.g. Bournaud et al. 2011; Porter et al. 2014) and the evolution of giant clumps (Dekel & Burkert 2013). These important processes likely complicate interpretation of a given observation, and recent studies of galaxy morphology have begun to exploit simulations including them (e.g. Scannapieco et al. 2010; Pedrosa, Tissera & De Rossi 2014).","Citation Text":["Torrey et al. 2014"],"Functions Text":["The quality and breadth of these experiments are limited by the availability of computational resources and the fidelity of models for galaxy physics such as","It has only recently become widespread to model the formation of galaxies ab initio","and the realism continues to improve"],"Functions Label":["Motivation","Motivation","Motivation"],"Citation Start End":[[868,886]],"Functions Start End":[[345,502],[566,649],[788,824]]}
{"Identifier":"2021MNRAS.508.2583Z___2016_Instance_1","Paragraph":"Located in the star-forming region \u03c1-Ophiuchi, inside the dark cloud L1689N and at a distance of 141 pc (Dzib et al. 2018), IRAS16293\u22122422 is a well-studied Young Stellar Object (YSO) classified as a Class 0 source with less than 104 yr (Andre, Ward-Thompson & Barsony 1993), and represents one of the very early stages of low-mass star formation. It was the first source identified as a hot corino (Blake et al. 1994; van Dishoeck et al. 1995) based on the detection of Complex Organic Molecules (COMs) in the source, which was later supported by follow-up studies (Ceccarelli et al. 1998, 2000; Sch\u00f6ier et al. 2002; Crimier et al. 2010; J\u00f8rgensen et al. 2011, 2016; Pineda et al. 2012; Oya et al. 2016; Jacobsen et al. 2018; van der Wiel et al. 2019). Higher resolution observations revealed that IRAS16293\u22122422 is in fact a triple system, composed of sources A1 and A2, separated by 54 au from each other (Maureira et al. 2020) and source B, 738 au (5 arcse; Wootten 1989) away from source A. Due to this larger separation, tidal truncation between the three protostars is discarded and therefore source B is considered to have evolved as an isolated source (Rodr\u00edguez et al. 2005). It was initially proposed to be either an evolved T Tauri star (Stark et al. 2004; Takakuwa et al. 2007) or a very young object (Chandler et al. 2005), however, Chandler et al. (2005) suggested that source B has large-scale infalls based on SO line emission. Pineda et al. (2012) confirmed the infall of an inner envelope, with mass accretion rates of 4.5 \u00d7 10\u22125 M\u2299yr\u22121, based on ALMA detections of inverse P-Cygni profiles in CH3OCHO-E, CH3OCHO-E-A and H2CCO, ruling out the possibility of it being a T Tauri star. The interpretations of infall from these profiles was also suggested by J\u00f8rgensen et al. (2012) and Zapata et al. (2013). Unlike the A1 and A2 protostars, source B has not shown clear signs of outflow launching, explained by the lack of free\u2013free emission at low frequencies (Chandler et al. 2005; Rodr\u00edguez et al. 2005; Loinard et al. 2007; Rao et al. 2009; Liu et al. 2018; Hern\u00e1ndez-G\u00f3mez et al. 2019b) and also based on molecular lines (Loinard et al. 2002; van der Wiel et al. 2019).","Citation Text":["J\u00f8rgensen et al.","2016"],"Functions Text":["It was the first source identified as a hot corino","based on the detection of Complex Organic Molecules (COMs) in the source, which was later supported by follow-up studies"],"Functions Label":["Background","Background"],"Citation Start End":[[639,655],[662,666]],"Functions Start End":[[348,398],[445,565]]}
{"Identifier":"2021AandA...647A.140C__Gianninas_et_al._2016_Instance_2","Paragraph":"In recent years, numerous low-mass and ELM WDs have been detected in the context of relevant surveys, such as the SDSS, ELM, SPY and WASP (see, e.g., Koester et al. 2009; Brown et al. 2010, 2016, 2020; Kilic et al. 2011, 2012; Gianninas et al. 2015; Kosakowski et al. 2020). The discovery of their probable precursors, namely, the so-called low-mass pre-WDs, has triggered an interest in these types of objects because of the possibility of studying the evolution of the progenitors that lead to the WD phase. Even more interestingly, the detection of multi-periodic brightness variations in low-mass WDs (Hermes et al. 2012, 2013a,b; Kilic et al. 2015, 2018; Bell et al. 2017, 2018; Pelisoli et al. 2018), and low-mass pre-WDs (Maxted et al. 2013, 2014; Gianninas et al. 2016; Wang et al. 2020) has brought about new classes of pulsating stars known as ELMVs and pre-ELMVs, respectively (ELM and pre-ELM variables, respectively). It has allowed for the study of their stellar interiors using the tools of asteroseismology, similarly to the case of other pulsating WDs such as ZZ Ceti stars or DAVs \u2013pulsating WDs with H-rich atmospheres \u2013 and V777 Her or DBVs \u2013 pulsating WDs with He-rich atmospheres (Winget & Kepler 2008; Fontaine & Brassard 2008; Althaus et al. 2010; C\u00f3rsico et al. 2019). The pulsations observed in ELMVs are compatible with global gravity (g)-mode pulsations. In the case of pulsating ELM WDs, the pulsations have large amplitudes mainly at the core regions (Steinfadt et al. 2010; C\u00f3rsico et al. 2012; C\u00f3rsico & Althaus 2014a), allowing for the study of their core chemical structure. According to nonadiabatic computations (C\u00f3rsico et al. 2012; Van Grootel et al. 2013; C\u00f3rsico & Althaus 2016), these modes are probably excited by the \u03ba\u2005\u2212\u2005\u03b3 (Unno et al. 1989) mechanism acting at the H-ionization zone. In the case of pre-ELMVs, the nonadiabatic stability computations for radial (Jeffery & Saio 2013) and nonradial p- and g-mode pulsations (C\u00f3rsico et al. 2016; Gianninas et al. 2016; Istrate et al. 2016b) revealed that the excitation is probably due to the \u03ba\u2005\u2212\u2005\u03b3 mechanism, acting mainly in the zone of the second partial ionization of He, with a weaker contribution from the region of the first partial ionization of He and the partial ionization of H. The presence of He in the driving zone is crucial to having the modes destabilized by the \u03ba\u2005\u2212\u2005\u03b3 mechanism (C\u00f3rsico & Althaus 2016; Istrate et al. 2016b).","Citation Text":["Gianninas et al. 2016"],"Functions Text":["In the case of pre-ELMVs,","and nonradial p- and g-mode pulsations","revealed that the excitation is probably due to the \u03ba\u2005\u2212\u2005\u03b3 mechanism, acting mainly in the zone of the second partial ionization of He, with a weaker contribution from the region of the first partial ionization of He and the partial ionization of H."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[1988,2009]],"Functions Start End":[[1828,1853],[1927,1965],[2033,2281]]}
{"Identifier":"2018MNRAS.477.3520L__Abolfathi_et_al._2018_Instance_1","Paragraph":"Over time, the data releases have treated the Balmer line regions in different ways. The presence of the artificial curvature was first reported by Busca et al. (2013) in the context of the DR9 data release. To minimize this effect, a different scheme was used in DR12 (Alam et al. 2015, see their table 2) by using a linear function (instead of an iterative b-spline procedure) to interpolate the flux over the masked regions. Surprisingly, we observe that this data reduction change was only applied to the Balmer \u03b2, \u03b3, and \u03b4 lines but not applied to the Balmer \u03b1 line. For the latter one, the original iterative spline interpolation scheme has been used all along. As a result, an absorption-like feature at the location of Balmer \u03b1 is found in SDSS data releases 9 up to now, i.e. the latest data release 14 (Abolfathi et al. 2018). To illustrate this, we show examples of calibration vectors for SDSS BOSS DR9 (Ahn et al. 2012; Dawson et al. 2013), DR12 (Alam et al. 2015), eBOSS DR14 (Dawson et al. 2016; Abolfathi et al. 2018) data release as well as calibration vectors for the MaNGA survey (Bundy et al. 2015) DR14 data release (Abolfathi et al. 2018) in Fig. 6. In the upper panel, we show small-scale features in calibration vectors from the same plate (number 3647) processed by the DR9 (blue), DR12 (green), and DR14 (red) pipelines. The spectra are obtained by normalizing the large-scale features in calibration vectors with a cubic b-spline with break points separated by 50 \u00c5. The three spectra are mostly identical while the smooth curving features at the wavelengths of Balmer series in DR9 disappear in DR12 and DR14. From DR12, the pipeline corrects the features with linear interpolation across wavelength regions with Balmer \u03b2, \u03b3, and \u03b4 lines (Alam et al. 2015). One can also observe the change in the median residual spectra of DR9 and DR12 in Fig. 1. However, the H\u03b1 feature remains uncorrected from DR9 to DR14. We also show a DR7 calibration vector (black) in which most of the wiggles are absent. As pointed previously, this is due to the fact that the DR7 pipeline interpolates the calibration vectors using an effective scale larger than that used in subsequent data releases.","Citation Text":["Abolfathi et al. 2018"],"Functions Text":["Surprisingly, we observe that this data reduction change was only applied to the Balmer \u03b2, \u03b3, and \u03b4 lines but not applied to the Balmer \u03b1 line. For the latter one, the original iterative spline interpolation scheme has been used all along. As a result, an absorption-like feature at the location of Balmer \u03b1 is found in SDSS data releases 9 up to now, i.e. the latest data release 14"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[813,834]],"Functions Start End":[[428,811]]}
{"Identifier":"2018ApJ...869..121M__Manfroid_et_al._2009_Instance_1","Paragraph":"An example of a low [14N\/15N] value is potentially interesting for understanding the origins of material in our own solar system. The observed value in N2 of \u223c200 is lower than values measured for both atmospheres of Earth and Venus (\u223c270; Junk & Svec 1958; Hoffman et al. 1979) and inferred for the protosolar nebula (\u223c440; Owen et al. 2001; Fouchet et al. 2004; Meibom et al. 2007; Marty et al. 2010). It is more similar to the extremely low values of [14N\/15N] found in primitive solar system material, like cometary CN and HCN (\u223c140; Arpigny et al. 2003; Bockel\u00e9e-Morvan et al. 2008; Manfroid et al. 2009) and in organic material in chondritic meteorites (\u223c50\u2013150; Briani et al. 2009; Bonal et al. 2010), and that recently observed in HCN in a sample of T Tauri and Herbig disks (Guzm\u00e1n et al. 2017). Low [14N\/15N] values are seen to be correlated with high D abundances in carbonaceous chondrites, \u201ccluster\u201d interplanetary dust particles, and comets Wild2 and Hale-Bopp (Messenger 2000; Busemann et al. 2006; Floss et al. 2006; Al\u00e9on 2010), an enrichment that occurred either in the protosolar disk or in the precursor cloud core. Recent ISM modeling suggests that this correlation may not occur in the core stage, as a correlation between D and 15N is not expected (due to, e.g., sensitive dependence of chemistry on ortho-to-para variations in collisional partners like H2; Wirstr\u00f6m et al. 2012; De Simone et al. 2018). This is supported by observations showing a hint of anticorrelation between the abundances of D and 15N using N2H+ in a sample of massive prestellar cores (Fontani et al. 2015b). However, the potential presence of simultaneously high 15N and D in N2 suggests that it could still be possible for this to occur in an individual core, allowing for an interstellar origin for this solar system feature. As the solar system is suggested to have formed in a rich stellar cluster experiencing a nearby supernova (Adams 2010; Dukes & Krumholz 2012; Pfalzner 2013), the environment of a massive protocluster like Sgr B2 may actually be quite relevant for understanding the neighborhood of the protosolar nebula.","Citation Text":["Manfroid et al. 2009"],"Functions Text":["It is more similar to the extremely low values of [14N\/15N] found in primitive solar system material, like cometary CN and HCN"],"Functions Label":["Similarities"],"Citation Start End":[[588,608]],"Functions Start End":[[404,530]]}
{"Identifier":"2020AandA...639A..88C__Chatzistergos_et_al._2019b_Instance_1","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":["Chatzistergos et al. 2019b"],"Functions Text":["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","Paper II, hereafter)."],"Functions Label":["Background","Background"],"Citation Start End":[[907,933]],"Functions Start End":[[544,905],[935,956]]}
{"Identifier":"2022ApJ...931...70B__Gabrielse_et_al._2012_Instance_1","Paragraph":"RFs can propagate from the magnetotail to Earth over a long distance more than 10 R\nE together with BBFs behind them (Runov et al. 2009; Cao et al. 2010). Studies have suggested that RFs are crucial regions for particle acceleration, pitch-angle evolution, wave\u2013particle interactions, and electromagnetic energy conversion during their Earthward propagation. For instance, rapid increases in energy fluxes of electrons and ions from tens to hundreds of keV are a typical feature of RF events (Khotyaintsev et al. 2011; Liu et al. 2013, 2018c, 2021a, 2022b; Zhou et al. 2018; Liu & Fu 2019; Gabrielse et al. 2021), pitch-angle distribution of suprathermal electrons can evolve dramatically around RFs (Runov et al. 2013; Liu et al. 2020), strong particle and wave activity can occur in the vicinity of RFs (Ono et al. 2009; Zhou et al. 2009, 2014; Fu et al. 2014; Breuillard et al. 2016; Greco et al. 2017; Yang et al. 2017), and RFs are associated with energy conversion from electromagnetic fields to particles (Sitnov et al. 2009; Huang et al. 2015; Khotyaintsev et al. 2017; Liu et al. 2018a, 2022a). The energetic plasma in the vicinity of RFs plays a key role in connecting the magnetotail with the inner magnetosphere because they carry a large amount of energy and can be injected into the inner magnetosphere to affect the ring current and radiation belt (Gabrielse et al. 2012; Duan et al. 2014; Turner et al. 2014). Possible mechanisms responsible for the energization of particles around RFs have been widely investigated based on both spacecraft observations and numerical simulations during the past decade. The strong convection electric field induced by the strong magnetic field gradient of RFs provides significant adiabatic acceleration of the ambient particles (Birn et al. 2004, 2013, 2015; Gabrielse et al. 2012, 2014, 2016; Ganushkina et al. 2013; Liu et al. 2016; Turner et al. 2016). Nonadiabatic effects, caused by particle reflection ahead of the RFs (Zhou et al. 2018), resonance with RFs (Ukhorskiy et al. 2013, 2017), and scattering by wave emissions (Zhou et al. 2009; Greco et al. 2017), are also significant for particle energization. These above studies usually assumed that the RF surface has a planar boundary at a typical thickness comparable to the ion gyroradius and below (Nakamura et al. 2002; Sergeev et al. 2009; Zhou et al. 2009; Schmid et al. 2011; Liu et al. 2013; Vapirev et al. 2013). Divin et al. (2015b) revealed that the RF surface is unstable to instabilities ranging from electron scales to ion scales. Simulation studies found that RFs can be unstable to interchange instability and that finger-like structures on ion\u2013electron hybrid scales can develop at the RF (Vapirev et al. 2013). Such finger-like structures are found to play a role in modulating the electron acceleration process (Wu et al. 2018). Bai et al. (2022) also reported significant ion trapping acceleration at the RF with ion-scale ripples. Unlike these surface structures with ion or ion\u2013electron hybrid scales, Liu et al. (2018b) recently reported that the RF layer has electron-scale density gradients, currents, and electric fields, based on the MMS mission, which consists of four spacecraft separated by 30 km. Such electron-scale ripple structure can be generated by lower hybrid drift instability (Divin et al. 2015b; Pan et al. 2018). Liu et al. (2021c) presented a detailed investigation of energy flux densities at two RFs with\/without the electron-scale surface ripples and indicated that surface ripples may play an important role in the particle dynamics. But how such electron-scale RF structure impacts the electron energization and transport still remains unknown. In this paper, with the aid of observation-based test-particle simulation, we aim to investigate in detail the effect of the front surface ripples on the local electron dynamics.","Citation Text":["Gabrielse et al. 2012"],"Functions Text":["The energetic plasma in the vicinity of RFs plays a key role in connecting the magnetotail with the inner magnetosphere because they carry a large amount of energy and can be injected into the inner magnetosphere to affect the ring current and radiation belt"],"Functions Label":["Motivation"],"Citation Start End":[[1364,1385]],"Functions Start End":[[1104,1362]]}
{"Identifier":"2022MNRAS.511.1714T__Pe\u00f1arrubia_et_al._2005_Instance_1","Paragraph":"While these similar characteristics suggest that the OA and C could have a relationship, the nature of this relationship is not clear. Kawata, Thom & Gibson (2003) used numerical simulations to investigate whether Complex C could have been produced by the passage of a satellite galaxy through the Milky Way disc, and while they found that such an event could produce a structure with the general properties of Complex C and the OA, they did not favour this hypothesis because they could not identify compelling evidence of the putative satellite. However, subsequently several stellar streams and structures have been discovered in this region of the outer Galaxy, and the progenitor(s) of these stellar structures could potentially also explain the origin of Complex C and the OA in a scenario like the one investigated by Kawata et al. (2003). The most well-known stellar stream in this region is the Monoceros Ring (e.g. Newberg et al. 2002; Pe\u00f1arrubia et al. 2005), but surveys such as the Sloan Digital Sky Survey (SDSS; York et al. 2000) and the Gaia Survey (Gaia Collaboration 2016) have revealed additional stellar streams in this part of the Galaxy. Table 1 summarizes some properties of stellar streams in this direction including Monoceros \u2018North\u2019, the Anticentre Stream, Hr\u00edd, Gaia 9, and GD1. It is intriguing to note that some of the metallicities and locations of these stellar streams are similar to the published metallicities of Complex C and the OA, but as we will see in this paper, the metallicities of these HVCs are not securely measured \u2013 there are degeneracies in the ionization models used to derive metallicities, and it is possible that these HVCs have significantly higher metallicities than the published abundances. It is also worth noting that some of these stellar streams, including the northern Monoceros Ring, Hr\u00edd, and GD1, have similar velocities to the HVCs. The Anticentre Stream, on the other hand, exhibits significantly different kinematics with mean velocities of $\\it {v}_{\\rm LSR} = +48$ and +78 km\u2009s\u22121 in two directions studied by Grillmair, Carlin & Majewski (2008). However, spatially overlapping streams with different kinematics could lead to a roiled, turbulent region, which in turn could make precipitation more likely (Voit 2021). To more visually show the spatial correspondence of the stellar structures and the gas clouds, Fig. 1 shows a schematic map of the projected locations of the streams and clouds from Table 1. Note that Fig. 1 is centred and zoomed in on a set of active galactic nuclei employed in this paper (see below) and does not show the full extent of the HVCs and streams.","Citation Text":["Pe\u00f1arrubia et al. 2005"],"Functions Text":["The most well-known stellar stream in this region is the Monoceros Ring (e.g."],"Functions Label":["Background"],"Citation Start End":[[946,968]],"Functions Start End":[[847,924]]}
{"Identifier":"2021MNRAS.506.1045M__Marshall_et_al._2013_Instance_1","Paragraph":"Discovered in 1977 from its bright H \u03b1 emission (Stephenson & Sanduleak 1977), SS433\u2019s defining characteristics are undoubtedly the helical motion of highly collimated jets of plasma launched from its innermost regions, and mass-loaded, non-polar outflows (Fabian & Rees 1979; Margon et al. 1979) which together inflate the surrounding W50 supernova remnant. Knots in SS433\u2019s jet can be resolved at radio frequencies using very long baseline interferometry (VLBI) and indicate the presence of highly relativistic electrons (Vermeulen et al. 1987), while the baryon content is revealed by emission lines ranging from H and He lines in the optical through to highly ionized Fe lines in the X-rays (Kotani et al. 1994; Marshall et al. 2013). The Doppler shifts of the lines indicate precession of the accreting system with a period of \u2248 162 d, also seen in optical (He ii) emission lines originating from the non-polar wind (Fabrika 1997). Both the jets and winds carry a large kinetic luminosity (>1038 erg s\u22121, e.g. Marshall et al. 2002), which requires extraction of energy via accretion on to a compact object. While the nature of the compact object in SS433 remains somewhat unknown (although dynamical arguments suggest the presence of a black hole \u2013 Blundell, Bowler & Schmidtobreick 2008), the rate of mass transfer from the companion star, as inferred from the IR excess (Shkovskii 1981; Fuchs et al. 2006), is thought to be \u223c1 \u00d7 10\u22124 M\u2299 yr\u22121, orders of magnitude in excess of the Eddington limit for any plausible stellar remnant (>300 times the Eddington mass accretion rate for a typical stellar mass black hole of around 10 M\u2299). Classical theory and radiation magnetohydrodynamic (RMHD) simulations agree that such \u2018super-critical\u2019 rates of accretion will lead to a radiatively supported, large scale height (H\/R \u2248 1, where H is the height of the disc at distance R from the compact object) accretion disc with powerful winds launched from the surface at mildly relativistic speeds (Shakura & Sunyaev 1973; Poutanen et al. 2007; Ohsuga & Mineshige. 2011; Takeuchi et al. 2013; Jiang et al. 2014; Sadowski et al. 2014).","Citation Text":["Marshall et al. 2013"],"Functions Text":["while the baryon content is revealed by emission lines ranging from H and He lines in the optical through to highly ionized Fe lines in the X-rays"],"Functions Label":["Background"],"Citation Start End":[[716,736]],"Functions Start End":[[548,694]]}
{"Identifier":"2020MNRAS.498..385J__Chomiuk_&_Povich_2011_Instance_1","Paragraph":"In Fig. 9, we show the total galactic star formation rate as a function of simulation time t for each isolated disc galaxy. Following the initial vertical collapse of the disc and the subsequent star formation \u2018burst\u2019 from t \u223c 30\u2009Myr to t \u223c 250\u2009Myr, the SFR settles down to a rate of \u223c2\u20134\u2009M\u2299\u2009yr\u22121. We make absolutely sure to consider each isolated disc in its equilibrium state by examining the cloud population during a later time interval, between t = 600\u2009Myr and t = 1\u2009Gyr (grey-shaded region). Over this period, the SFR declines only gradually, by a total of around 0.5\u2009M\u2299\u2009yr\u22121. These values are consistent with the current observed SFR in the Milky Way (Murray & Rahman 2010; Robitaille & Whitney 2010; Chomiuk & Povich 2011; Licquia & Newman 2015). We may also consider the resolved star-forming behaviour on scales of 750\u2009pc, as studied in nearby galaxies by Bigiel et al. (2008). In Fig. 10, we display the star formation rate surface density as a function of gas surface density for each of our simulated galaxies, at a simulation time of 600\u2009Myr. The top row shows the 2D projection maps of the CO-bright molecular gas column density $\\Sigma _{{\\rm H}_2}$. These are computed via the total gas column density in equation (35) and degraded using a Gaussian filter of FWHM = 750\u2009pc. The corresponding projections of the star formation rate surface densities \u03a3SFR are displayed in the central row. Details for the production of all maps are given in Appendix A. In the bottom row, the values of $\\Sigma _{{\\rm H}_2}$ and \u03a3SFR in each pixel of the spatially degraded projection maps are compiled to produce a single histogram. The loci of our simulation data fall close to the observed star formation relations obtained by Bigiel et al. (2008), denoted by the orange contours, though with a population of points at lower densities and star formation rates than are reached by the observations. These points arise because we consider all CO-emitting gas down to a molecular hydrogen surface density of $\\Sigma _{{\\rm H}_2} = 10^{-3.5} \\: {\\rm M}_\\odot {\\rm pc}^{-2}$ (see Fig. 4). This avoids taking an arbitrary cut on $\\Sigma _{\\rm H_2}$, but also captures much lower levels of CO emission than could be detected by current observatories.","Citation Text":["Chomiuk & Povich 2011"],"Functions Text":["These values are consistent with the current observed SFR in the Milky Way"],"Functions Label":["Similarities"],"Citation Start End":[[708,729]],"Functions Start End":[[583,657]]}
{"Identifier":"2022ApJ...928..167Z__Kr\u00fcger_&_Foucart_2020_Instance_1","Paragraph":"Compared with BNS mergers, which would definitely eject a certain amount of materials to produce EM signals, some NSBH binaries may not tidally disrupt the NS component and, hence, would not make bright EM counterparts such as sGRBs and kilonovae.\n9\n\n\n9\nDuring the final merger phase for plunging NSBH binaries, some weak EM signals may be produced because of the charge and magnetic field carried by the NS (e.g., Dai 2019; Pan & Yang 2019; Zhang 2019; D\u2019Orazio et al. 2022; Sridhar et al. 2021). The tidal disruption probability of NSBH mergers and the brightness of NSBH EM signals are determined by the BH mass, BH spin, NS mass, and NS equation of state (EoS; e.g., Belczynski et al. 2008; Kyutoku et al. 2011, 2013, 2015; Fern\u00e1ndez et al. 2015; Kawaguchi et al. 2015, 2016; Foucart 2012; Foucart et al. 2018; Barbieri et al. 2019; Kr\u00fcger & Foucart 2020; Fragione & Loeb 2021; Fragione 2021; Zhu et al. 2020, 2021c, 2021e; Raaijmakers et al. 2021; Li & Shen 2021; Tiwari et al. 2021). An NSBH merger tends to be a disrupted event and produces bright EM signals if it has a low-mass BH with a high projected aligned spin and a low-mass NS with a stiff EoS. The parameter space in which an NSBH merger can undergo tidal disruption may be very limited. Recently, LIGO\u2013Virgo\u2013KAGRA (LVK) Collaboration reported three high-confidence GWs from NSBH candidates, i.e., GW190814, GW200105_162426, and GW200115_042309 (Abbott et al. 2020, 2021a; Nitz et al. 2021). In spite of many efforts for follow-up observations of these three events, no confirmed EM counterpart candidate has been identified (e.g., Coughlin et al. 2020; Gompertz et al. 2020; Kasliwal et al. 2020;Page et al. 2020; Thakur et al. 2020; Alexander et al. 2021; Anand et al. 2021; Dobie et al. 2022; Kilpatrick et al. 2021). Abbott et al. (2021a), Zhu et al. (2021c), and Fragione (2021) showed that the parameter space of these GW candidates mostly lies outside the disrupted parameter region, so these candidates are likely plunging events with a high probability. There are many mysteries surrounding NSBH binaries, such as the proportion of disrupted events in cosmological NSBH mergers, their cosmological contribution to elements heavier than iron, the formation channel of NSBH binaries, and so on. Systemic research on the population properties of NSBH binaries can help us address these mysteries and unveil the nature of cosmological NSBH binaries.","Citation Text":["Kr\u00fcger & Foucart 2020"],"Functions Text":["The tidal disruption probability of NSBH mergers and the brightness of NSBH EM signals are determined by the BH mass, BH spin, NS mass, and NS equation of state (EoS; e.g."],"Functions Label":["Uses"],"Citation Start End":[[837,858]],"Functions Start End":[[498,669]]}
{"Identifier":"2019ApJ...877...33Z__Dunn_et_al._2010_Instance_1","Paragraph":"To investigate the emission-line properties in the blue system, we employed the photoionization code Cloudy (Version 17.01, Ferland et al. 2017) and applied the measured emission-line ratios to these models to simulate the possible physical conditions and processes in the medium. The simple model is a slab-shaped gas with a unique density and homogeneous chemical composition of solar values, irradiated directly by the central ionization continuum source. For the BLR around one of the hypothesized binary black holes, such primordial models can be simply described using parameters such as density nH, hydrogen column density NH, ionization parameter U, and spectral energy distribution (SED) of the incident radiation. A typical AGN ionization continuum is applied as the incident SED, which is a combination of a blackbody \u201cBig Bump\u201d and power laws.6\n\n6\nSee details in Hazy, a brief introduction to Cloudy; http:\/\/www.nublado.org.\n The component peaks at \u22481 Ryd and is parameterized by TBB = 1.5 \u00d7 105 K. The slope of the X-ray component, the X-ray to UV ratio, and the low-energy slope are set as \u03b1x = \u22122 (beyond 100 keV) and \u22121 (between 1.36 eV and 100 keV), \u03b1ox = \u22121.4, and \u03b1UV = \u22120.5, respectively. This UV-soft SED is regarded as more realistic for radio-quiet quasars than the other available SEDs provided by Cloudy (see the detailed discussion in Section 4.2 of Dunn et al. 2010). We calculated a series of photoionization models with different ionization parameters, densities, and hydrogen column densities. The ranges of these parameters are \u22124.0 \u2264 log10 U \u2264 1.0, 3.0 \u2264 log10 nH (cm\u22123) \u2264 11.0, and 19.0 \u2264 log10 NH (cm\u22122) \u2264 24.0, with a step of 0.2 dex, which could well cover the possible parameter space of the broad-line and narrow-line regions. In Figure 3, the flux ratios of He i\u03bb10830\/H\u03b1 are presented in the log U\u2013log NH space for the five densities (log nH (cm\u22123) = 3, 5, 7, 9, and 11). The extensive parameter space is almost enough to cover all the routine possibilities of the AGN \u201cnormal\u201d emission-line regions. However, the typical BLR gases with nH \u223c 109\u20131010 cm\u22123, NH \u223c 1022 cm\u22122 and U \u223c 10\u22122\u201310\u22121 present the relatively strong He i \u03bb10830 emission. If the blue system of SDSS J1536+0441 is also compared, the observed ratio of He i10830\/H\u03b16564 is at least larger than 0.1, which is two times the 3\u03c3 upper limit. Further extensive calculation in a larger parameter space suggests that a parameter combination range of UnH \u2248 1011.5 cm\u22123 with an exceedingly high density of nH \u2265 1012 cm\u22123 would reproduce the observed flux ratios of the hydrogen and helium lines. If the blue system emits from such a high-density medium, the size would be 50 times less than that of the red system BLR (based on the sub-parsec binary hypothesis, in which both BLRs are illuminated by the same ionizing flux), and the number ratios of ions in the blue system and red system would be only approximately five per thousand. Even accounting for the higher emission efficiency (within an order of magnitude), the high-density medium is not sufficient to emit blue system lines comparable to those of the red system.","Citation Text":["Dunn et al. 2010"],"Functions Text":["This UV-soft SED is regarded as more realistic for radio-quiet quasars than the other available SEDs provided by Cloudy (see the detailed discussion in Section 4.2 of"],"Functions Label":["Uses"],"Citation Start End":[[1376,1392]],"Functions Start End":[[1209,1375]]}
{"Identifier":"2021MNRAS.500.3438O__Vidotto_et_al._2014b_Instance_1","Paragraph":"In addition to the wind models of \u03bb And, we also present here the first full surface magnetic field observations of this star, finding a strong magnetic field for such an evolved star. These observations, carried out with the NARVAL spectropolarimeter, allow us to constrain the surface magnetic field of \u03bb And. These derived surface magnetic fields can constrain the lower boundary of the 3D magnetohydrodynamic wind simulations that we run. Usually, we see a decay in magnetic field strength as solar-type stars evolve, as their activity decreases along with their rotation (Skumanich 1972; Vidotto et al. 2014b; Booth et al. 2020). However, this subgiant star seems to have a relatively strong large-scale magnetic field compared to the Sun. The exact process through which this star would reach this stage in its evolution with such a magnetic field is yet unknown. Potential reasons are that it began with a much stronger dynamo in its past than anticipated, or perhaps the secondary companion had some effect on the primary star at a point in the past. \u03bb And differs from the Sun as it is an RS Canum Venaticorum (RS CVn) variable, meaning it is a variable binary system. The variability on this star is likely due to magnetic spots coming in and out of view due to stellar rotation (Baliunas & Dupree 1979, 1982; Donati, Henry & Hall 1995; Henry et al. 1995; O\u2019Neal et al. 2001; Sanz-Forcada, Brickhouse & Dupree 2001; Frasca et al. 2008; Drake et al. 2011). RS CVn systems, in particular, can present observed levels of chromospheric and coronal activity that are orders of magnitude higher than in single stars with similar spectral types (Ayres & Linsky 1980; Walter & Bowyer 1981). This is likely caused by the increase in activity when the two stars interact with each other, which can lead to rotational synchronization of the system (e.g. Lanza & Rodon\u00f2 2004; an analogous process has been inferred to take place in close-in planet\u2013star systems, Cuntz, Saar & Musielak 2000). For the purposes of this work, we assume the binarity of this system does not affect our wind models. Compared to the Sun, \u03bb And is metal-poor ([Fe\/H] = \u22120.46 \u00b1 0.04 dex, Maldonado & Villaver 2016). We do not include the effects of different metal abundances on the stellar wind and stellar evolution, but the effects of which have been examined in other works (Suzuki 2018).","Citation Text":["Vidotto et al. 2014b"],"Functions Text":["Usually, we see a decay in magnetic field strength as solar-type stars evolve, as their activity decreases along with their rotation","However, this subgiant star seems to have a relatively strong large-scale magnetic field compared to the Sun. The exact process through which this star would reach this stage in its evolution with such a magnetic field is yet unknown."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[593,613]],"Functions Start End":[[443,575],[635,869]]}
{"Identifier":"2022MNRAS.509.3599T__Matteo,_Springel_&_Hernquist_2005_Instance_1","Paragraph":"The typical value of the mass outflow rate for sources accreting below or close to the Eddington limit is $\\dot{M}_{\\rm out} \\gtrsim 5\\!-\\!10{{\\ \\rm per\\ cent}} \\,\\dot{M}_{\\rm acc}$, for both UFOs and non-UFOs (Tombesi et al. 2012). In this scenario, even for the UFOs with the lowest allowed velocity, the mechanical power is enough to exercise a significant feedback impact on the surrounding environment. Looking at the comparison between mass accretion rate and mass outflow rate for our source (see Table 3), the upper limit on the mass outflow rate for Wind 1 is extremely high, but the values for Wind 2 and Wind 3 are still comparable with the values of quasars and Seyfert galaxies (Tombesi et al. 2012). If we consider the lower limits instead, their values are well below the average. Theoretical works (Di Matteo, Springel & Hernquist 2005; King 2010; Ostriker et al. 2010; Debuhr, Quataert & Ma 2011) showed that, in order to have a significant feedback impact in the environment surrounding an AGN, it is required a minimum ratio between the mechanical power of the outflow and the bolometric luminosity of ${\\sim}0.5{{\\ \\rm per\\ cent}}$. Tombesi et al. (2012) showed that actually the lower limit of this value for UFOs is ${\\sim}0.3{{\\ \\rm per\\ cent}}$ and for non-UFOs is ${\\sim}0.02 \\!-\\! 0.8{{\\ \\rm per\\ cent}}$. According to what found for the ratio between the mass outflow rate and the mass accretion rate, looking at the upper limits on $\\dot{K}\/L_{\\rm b,out}$ of our source, i.e. multiphase and multiscale X-ray winds, IRAS 04416+1215 fits well in this scenario in which the outflowing winds can impress a feedback. Indeed, the upper limits on $\\dot{K}\/L_{\\rm b,out}$ is comparable with the kinetic coupling efficiency, defined as the ratio of the kinetic luminosity of outflows to the AGN radiative luminosity (Eout\/Lrad), calculated with the feedback model for hyper-Eddington accretion by Takeo, Inayoshi & Mineshige (2020), using the outflow velocities and the $\\dot{M}\/\\dot{M}_{\\rm acc}$ values of IRAS 04416+1215. Instead, the lower limits are below the minimum value required to generate at least a weak feedback. Considering only the values derived from the lower limits on the distance, i.e. the situation in which the X-ray winds are cospatial, we would be in a scenario in which the source loses much luminosity due to advection inside the disc, resulting in a much lower efficiency for wind production as most of the radiation remains trapped inside the disc. This deduction is supported also by the results of the lower limits on the ratio between the momentum rate of the outflows and the momentum of the radiation. Outflows accelerated through the continuum radiation pressure are expected to have a $\\dot{p}_{\\rm out}\/\\dot{p}_{\\rm rad}\\sim 1$ (King & Pounds 2015). The median value of this ratio for UFOs is \u223c0.96 after the relativistic correction and \u223c0.64 without the relativistic corrections (Luminari et al. 2020). The values we found for the lower limits on $\\dot{p}_{\\rm out}\/\\dot{p}_{\\rm rad}$ of IRAS 04416+1215 are again well below the median. Thus, the outcoming luminosity of the source is not enough to accelerate the material to the escape velocity, which is required for a wind to leave the system, suggesting that likely in the scenario of the cospatial winds the outflows observed in IRAS 04416+1215 could be accelerated by other mechanisms such as magnetohydrodynamic processes.","Citation Text":["Di Matteo, Springel & Hernquist 2005"],"Functions Text":["Theoretical works","showed that, in order to have a significant feedback impact in the environment surrounding an AGN, it is required a minimum ratio between the mechanical power of the outflow and the bolometric luminosity of ${\\sim}0.5{{\\ \\rm per\\ cent}}$.","Instead, the lower limits are below the minimum value required to generate at least a weak feedback. Considering only the values derived from the lower limits on the distance, i.e. the situation in which the X-ray winds are cospatial, we would be in a scenario in which the source loses much luminosity due to advection inside the disc, resulting in a much lower efficiency for wind production as most of the radiation remains trapped inside the disc."],"Functions Label":["Uses","Uses","Uses"],"Citation Start End":[[815,851]],"Functions Start End":[[796,813],[914,1152],[2044,2495]]}
{"Identifier":"2020ApJ...892...68N__Laurikainen_et_al._2004_Instance_1","Paragraph":"We also examine our \n\n\n\n\n\n measurement of NGC 3504 in the context of the empirical compilations of \n\n\n\n\n\n and \n\n\n\n\n\n scaling relations (Kormendy & Ho 2013; McConnell et al. 2013; Scott et al. 2013; Saglia et al. 2016) in Figure 16. The stellar-bulge velocity dispersion of NGC 3504 is determined from Ho et al. (2009), while the bulge mass of \n\n\n\n\n\nM\u2299 was derived from Section 2.2. For a sanity check, we estimate the bulge mass of NGC 3504 using our H-band MGE model and adapting the effective radius of the bulge derived from the bulge-disk-bar decomposition model (Laurikainen et al. 2004). Accounting for the dynamical M\/LH (Section 5.2), we obtain \n\n\n\n\n\nM\u2299. Our result shows that the best-fit \n\n\n\n\n\n of NGC 3504 is fully consistent with the empirical \n\n\n\n\n\n and \n\n\n\n\n\n relations of Kormendy & Ho (2013), McConnell et al. (2013), and Saglia et al. (2016), but outside \n\n\n\n\n\n uncertainty of the Scott et al. (2013) and Savorgnan et al. (2016) empirical \n\n\n\n\n\n relations for \u201cS\u00e9rsic\u201d galaxies (those without central cores). Compared with the theoretical predictions of the bimodality in the BH accretion efficiency model (e.g., Pacucci et al. 2015, 2018), our measurement is consistent with the \n\n\n\n\n\n correlation, but a positive outlier is the \n\n\n\n\n\n relation up to 1\u03c3. At the mass of \n\n\n\n\n\n M\u2299, the SMBH of NGC 3504 lies within the same mass regime as the BHs derived in the Combes et al. (2019) sample and lies between the samples of lower- and higher-mass galaxies previously and currently studied with ALMA (Davis et al. 2013, 2017, 2018; Onishi et al. 2015, 2017; Barth et al. 2016a, 2016b; Boizelle et al. 2019; Nagai et al. 2019; North et al. 2019; Smith et al. 2019, T. Davis et al. 2020, in preparation, D. Nguyen et al. 2020, in preparation), respectively. All of these works prove that the cold gas-dynamical method observed with ALMA at high spatial resolution now can work well in a wide range of BH masses covering six orders of magnitude from 105 to \n\n\n\n\n\n.","Citation Text":["Laurikainen et al. 2004"],"Functions Text":["For a sanity check, we estimate the bulge mass of NGC 3504 using our H-band MGE model and adapting the effective radius of the bulge derived from the bulge-disk-bar decomposition model"],"Functions Label":["Uses"],"Citation Start End":[[568,591]],"Functions Start End":[[382,566]]}
{"Identifier":"2018ApJ...861...28S__Temmer_et_al._2011_Instance_1","Paragraph":"In addition, the graduated cylindrical shell (GCS) model, which was developed by Thernisien et al. (2006, 2009) and Thernisien (2011), is applied to obtain the three-dimensional parameters of these CMEs. Figure 5 shows the fitting results of these CMEs. Seen from these images, the GCS model can well represent the topology of these CMEs. The last three columns in Table 1 show the fitting results of these CMEs, including the propagation directions, velocities and face-on angular widths. Assuming a constant velocity and considering the influence of the propagation direction and angular width on the prediction of the arrival time suggested by Shen et al. (2014), CME-1 would arrive at the Earth around the time of September 6 22:27 UT. However, previous results show that fast CME would decelerate during their propagation in interplanetary space (e.g., Gopalswamy et al. 2001a, 2005; Vr\u0161nak 2001; Vr\u0161nak & \u017dic 2007; Temmer et al. 2011; Lugaz & Kintner 2012, and reference therein). Such deceleration may make the CME-1 arrive at the Earth later than September 6 22:27 UT. Thus, CME-1 is more likely to be the solar source of ICME-1. Seen from the Figure 3, the front edge of the CME-2 is lower than the front edge of CME-1 indicating that CME-2 would arrive at the Earth later than CME-1. Thus, CME-2 was the solar source of the ICME-2. It should be noted that, based on the fitting results of GCS model, CME-2 is faster than CME-1 and their propagation directions are close to each other. Thus, these two CMEs are expected to interact in interplanetary space. Seen from the in situ observations, possible interaction region signatures were detected between these two ICMEs with lower magnetic field, higher velocity, high density, and higher plasma beta. Furthermore, similar analysis shows that CME-3 might arrive at Earth after September 7 19:21 UT. Considering the long duration of Ejecta-4 and the larger angular width of CME-3, we verify that CME-3 is the solar source of ICME-4 and the driver of the second shock. It should be noted that, no obvious Earth-directed CME could be identified as the solar source of Ejecta-3. A possible explanation is that this ejecta structure is formed in the sheath region of ICME-4 during its propagation outward (e.g., Zheng & Hu 2018, and reference therein).","Citation Text":["Temmer et al. 2011"],"Functions Text":["However, previous results show that fast CME would decelerate during their propagation in interplanetary space (e.g.,"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[921,939]],"Functions Start End":[[740,857]]}
{"Identifier":"2018AandA...613A...3Q__Kelly_et_al._2017_Instance_1","Paragraph":"As a prototypical Seyfert 2 galaxy with starburst at a distance of 14.4 Mpc (1\u2033 = 72 pc, Bland-Hawthorn et al. 1997), NGC 1068 was observed at radio (Greenhill et al. 1996), millimeter (Schinnerer et al. 2000), infrared (Jaffe et al. 2004), optical (Antonucci & Miller 1985), UV (Antonucci et al. 1994), and X-ray (Kinkhabwala et al. 2002). High spatial resolution CO (1\u20130) observations show two molecular spiral arms with a diameter of ~40\u2033 and a northern half-bar, while a CO (2\u20131) map reveals a nuclear ring with two bright knots in the CND region (Schinnerer et al. 2000). The dense gas fraction as traced by HCN (1\u20130) (Tacconi et al. 1994; Helfer & Blitz 1995) and CS (2\u20131) (Tacconi et al. 1997; Takano et al. 2014) in the nuclear region is higher than the two arms. Observations of CO (3\u20132) (Krips et al. 2011; Tsai et al. 2012; Garc\u00eda-Burillo et al. 2014) showed that the difference of molecular gas temperatures between the nuclear region and the two arms was not as large as that of densities. Dozens of molecular lines at millimeter wavelength were detected at CND with single-dish observations (Usero et al. 2004; Nakajima et al. 2011, 2013; Aladro et al. 2013). Moreover, several molecules were detected and resolved toward NGC 1068 with interferometers in the past few years (Tosaki et al. 2017; Kelly et al. 2017; Furuya & Taniguchi 2016; Izumi et al. 2016; Imanishi et al. 2016; Nakajima et al. 2015; Viti et al. 2014; Takano et al. 2014; Garc\u00eda-Burillo et al. 2014, 2016). The molecular gas in the CND region was denser and hotter than that in the starburst ring, while chemical properties in the two regions were also different (Viti et al. 2014). The highest molecular gas temperature was higher than 150 K, and the gas density was above 105 cm\u22123 in the CND region (Viti et al. 2014). The distribution of different species of molecules were also different: CO isotopic species, for instance, were enhanced in the starburst ring, while the shock\/dust related molecules were enhanced in the CND region (Nakajima et al. 2015). The spatially resolved observations showed that the CND region was a complex dynamical system. For instance, the east and west dots were dominated by a fast shock and a slower shock (Kelly et al. 2017), while the dust torus also showed complex kinematics (Garc\u00eda-Burillo et al. 2016). Gas inflow was driven by a past minor merger (Furuya & Taniguchi 2016), while the outflow was AGN driven (Garc\u00eda-Burillo et al. 2014). We conducted adeeper survey of millimeter lines toward the CND region of NGC 1068 with the IRAM 30 m telescope, with the goal to quantify the gas properties in the CND. Compared to previous single-dish observations, our data probe weaker transition lines, which could place more constraints on the physical and chemistry properties of the CND.","Citation Text":["Kelly et al. 2017"],"Functions Text":["Moreover, several molecules were detected and resolved toward NGC 1068 with interferometers in the past few years"],"Functions Label":["Background"],"Citation Start End":[[1309,1326]],"Functions Start End":[[1174,1287]]}
{"Identifier":"2016MNRAS.461..666K__Joshi_et_al._2011_Instance_2","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"],"Functions Text":["The F\u03ba-test is preferred when the magnitude difference between the object and comparison stars is large"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1147,1164]],"Functions Start End":[[1042,1145]]}
{"Identifier":"2019AandA...630A...2H__Kofman_et_al._(2015)_Instance_1","Paragraph":"Several laboratory experiments have been performed to determine the mechanical properties of possible building blocks of comets. In particular, G\u00fcttler et al. (2009), Schr\u00e4pler et al. (2015), Lorek et al. (2016), and Katsuragi & Blum (2017) investigated how the compressibility of ice layers and different dust structures created by random ballistic deposition as well as pebble sub-structures (with scales of centimeters, millimeters, and 0.1 mm) depend on the volume filling factor. The respective results are summarized in Fig. 6. Remote observations lead to multiple independent estimates for the volume filling factor. Kofman et al. (2015) derived values of 0.15\u20130.25, P\u00e4tzold et al. (2016) estimated 0.25\u20130.30, and a study by Fulle et al. (2016) resulted in 0.21\u20130.37. The corresponding range is depicted in Fig. 6 by the horizontal bar shaded in blue. As thevolume filling factor is known, the laboratory measurements can be used to determine the most likely material of the surface. For the observed volume filling factor, dust and ice layers have a compressive strength above 103 Pa. A surface made up of such dust or ice layers is therefore inconsistent with the results of Groussin et al. (2015) or the compressive strength derived as part of this work. In contrast, a surface consisting of layers of dust and ice aggregates could explain all observations, except for the MPa range derived by Spohn et al. (2015). This is strong evidence for the presence of aggregate layers on the surface of 67P, that is, a surface composed of sub-decimeter-sized aggregates, as inferred from the gravitational collapse scenario. This conclusion can be drawn without precise knowledge of the actual value of the compressive strength, as the upper limit derived as part of this work and by Groussin et al. (2015), even considering all possible errors, is well below 103 Pa for different parts of 67P. The presence of such aggregate layers can only be explained by a formation dominated by relatively low-velocity collisions. Hierarchical agglomeration would cause a significantly higher compressive strength as a result of impact-compaction of the pebbles and possible break-up of aggregates (Blum 2018). The derived mechanical properties combined with the available laboratory models therefore suggest that the gentle collapse of aggregate ensembles played a major role in the formation ofcomets like 67P.","Citation Text":["Kofman et al. (2015)"],"Functions Text":["derived values of 0.15\u20130.25","The corresponding range is depicted in Fig. 6 by the horizontal bar shaded in blue. As thevolume filling factor is known, the laboratory measurements can be used to determine the most likely material of the surface."],"Functions Label":["Uses","Uses"],"Citation Start End":[[624,644]],"Functions Start End":[[645,672],[775,990]]}
{"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)"],"Functions Text":["are at the levels of >4\u03c3 for the analysis of"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1345,1357]],"Functions Start End":[[1300,1344]]}
{"Identifier":"2020MNRAS.491..903W__Rosotti_et_al._2014_Instance_1","Paragraph":"The process of planet formation is strongly dependent on the stellar birth environment. The majority of stars exist in clusters or associations within their first few Myr of evolution (Lada & Lada 2003; Longmore et al. 2014; Krumholz, McKee & Bland-Hawthorn 2019), during which time they also host protoplanetary discs (PPDs \u2013 e.g. Haisch, Lada & Lada 2001; Ribas et al. 2014). Multiple feedback mechanisms influence disc evolution. In sufficiently dense environments, star\u2013disc encounters can truncate the disc and induce increased accretion rates (Clarke & Pringle 1993; Ostriker 1994; Hall, Clarke & Pringle 1996; Pfalzner et al. 2005a; Olczak, Pfalzner & Spurzem 2006; Pfalzner, Olczak & Eckart 2006; de Juan Ovelar et al. 2012; Breslau et al. 2014; Rosotti et al. 2014; Winter et al. 2018a). Recent studies indicate that in the solar neighbourhood such interactions only have a significant effect in the early stages of cluster evolution due to enhanced stellar multiplicity and substructure, and therefore set initial conditions rather than destruction time-scales (Bate 2018; Winter et al. 2018b; Winter, Booth & Clarke 2018c). However, in regions with massive stars, external photoevaporation by far-ultraviolet (FUV) and extreme-ultraviolet (EUV) photons can rapidly disperse PPDs (Johnstone, Hollenbach & Bally 1998; St\u00f6rzer & Hollenbach 1999; Armitage 2000; Clarke 2007; Fatuzzo & Adams 2008; Adams 2010; Facchini, Clarke & Bisbas 2016; Ansdell et al. 2017; Haworth et al. 2018b; Winter et al. 2018b). Additionally, before the dispersal of the parent giant molecular cloud (GMC), ram pressure stripping can truncate PPDs (Wijnen et al. 2017a) or additional material can be accreted (Moeckel & Throop 2009; Scicluna et al. 2014), leading to the destruction and reforming of discs during the embedded phase (Bate 2018). If a PPD is destroyed quickly by feedback in dense stellar environments, planets may be unable to form, depending on the efficiency of the formation mechanisms (Youdin & Goodman 2005; Johansen & Lambrechts 2017; Ormel, Liu & Schoonenberg 2017; Haworth et al. 2018a). Given the apparent ubiquity of grouped star formation, quantifying the destruction time-scales for PPDs due to neighbour feedback is of great relevance for understanding the demographics of PPDs and exoplanetary systems.","Citation Text":["Rosotti et al. 2014"],"Functions Text":["Multiple feedback mechanisms influence disc evolution. In sufficiently dense environments, star\u2013disc encounters can truncate the disc and induce increased accretion rates"],"Functions Label":["Background"],"Citation Start End":[[754,773]],"Functions Start End":[[378,548]]}
{"Identifier":"2021AandA...647A.132K__Kouloumvakos_et_al._2015_Instance_1","Paragraph":"Studies of SEP events are important for different reasons. On one hand, solar eruptive events are well-observed processes of energetic-particle acceleration (Vainio & Afanasiev 2018), which can be studied in detail using a multi-messenger approach, complementing particle data with observations in different wavelengths (e.g., Plainaki et al. 2014; Cliver 2016; Kocharov et al. 2017). For this purpose, the peak flux intensity and detailed temporal variability of the particle flux are important as signatures of the acceleration process in the solar corona and the interplanetary medium (e.g., Desai & Giacalone 2016; Kong et al. 2017). Accordingly, numerous studies were focused on peak fluxes of SEPs and corresponding acceleration and transport processes (e.g., Kouloumvakos et al. 2015; Kocharov et al. 2017). On the other hand, enhanced fluxes of energetic particles affect the radiation environment near the Earth (e.g., Webber et al. 2007; Mishev et al. 2015), making not only the peak fluxes but also the fluence (event-integrated flux) and its spectral shape of significant importance, especially for extreme events (e.g., Cliver et al. 2020). We emphasise that SEP fluences can not be used for the detailed study of SEP acceleration processes, because (i) the observations at 1 AU are also modified by transport, and (ii) different energies in the fluence spectrum can be dominated by different acceleration mechanisms or by the same mechanism operating under different conditions. It is evident from the proton time-intensity profiles alone that the fluence at MeV and 10 MeV energies is often dominated by acceleration at interplanetary shocks (e.g., Reames 1999). However, the question is more open at 100 MeV and GeV energies peaking much earlier, with possible contributions from flares and\/or coronal shocks as the main candidates to account for the acceleration (see Cliver 2016, and references therein). Even if the same CME-driven shock were responsible for the acceleration of 10 MeV and 1 GeV protons, the former would typically be accelerated mainly in the solar wind and the latter in the corona, and there is no reason to suggest that the spectral form of the fluence would reveal something common about the accelerator properties.","Citation Text":["Kouloumvakos et al. 2015"],"Functions Text":["Accordingly, numerous studies were focused on peak fluxes of SEPs and corresponding acceleration and transport processes (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[766,790]],"Functions Start End":[[638,765]]}
{"Identifier":"2015MNRAS.450.4364N__Wu_et_al._2004_Instance_1","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"],"Functions Text":["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"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[874,888]],"Functions Start End":[[498,785]]}
{"Identifier":"2016AandA...589A..44G__within_2000_Instance_1","Paragraph":"W51e2 is the strongest and best-studied HC HII region in the W51 Main cluster, and it is believed to be powered by an O8-type young star (e.g., Shi et al. 2010a). A number of interferometric studies conducted with varying angular resolutions, at centimetre (cm) and (sub)millimetre (mm) bands, identified molecular and ionized gas undergoing infall and rotation toward W51e2. VLA observations of the NH3 inversion lines (1,\u20091) and (2,\u20092) seen in absorption (1\\hbox{$\\farcs$}.\u030b1 beamsize) revealed radial infall on scales larger than 5000 AU toward the W51e2 core (Zhang & Ho 1997). Higher angular resolution observations of the (3,\u20093) NH3\u2009absorption line (0\\hbox{$\\farcs$}.\u030b3 beamsize) showed signatures of rotation within 2000 AU based on a position-velocity (pv) diagram (Zhang & Ho 1997). Zhang et al. (1998) identified a velocity gradient in a CH3CN transition at 2 mm, deriving a position angle (PA) of 20 \u00b1 20\u00b0. Keto & Klaassen (2008) imaged the H53\u03b1 radio recombination line (RL) with the VLA (0\\hbox{$\\farcs$}.\u030b45 beamsize) and they claimed rotation in the ionized gas along the axis of a molecular bipolar outflow (oriented NW-SE) imaged with the SMA in the CO (2\u22121) line (1\u2032\u2032 beamsize), suggesting a simple inflow\/outflow picture in a single high-mass young stellar object (YSO). However, higher resolution observations, using the SMA at the wavelengths of 0.85 mm (0\\hbox{$\\farcs$}.\u030b3 beamsize) and 1.3 mm (0\\hbox{$\\farcs$}.\u030b7 beamsize), revealed a more complex picture, by resolving W51e2 into three subcores (Shi et al. 2010a): W51e2-W, corresponding to the HC HII region, W51e2-E, located about 1\u2032\u2032\u2009east of the HC HII region and corresponding to the brightest dust continuum source, and W51e2-NW, the weakest continuum component, located about 1\u2032\u2032\u2009 NW of the HC HII region. Shi et al. (2010b) imaged the CO (3\u22122) line (with a 0\\hbox{$\\farcs$}.\u030b7 beamsize) and established that the driving source of the powerful molecular outflow in this region is the protostellar core W51e2-E, and not the HC HII region W51e2-W, challenging the scenario proposed by Keto & Klaassen (2008). Etoka et al. (2012) used MERLIN to image the Class II 6.7 GHz CH3OH\u2009masers (typical signpost of HMSF), and found that the bulk of maser emission is indeed concentrated toward W51e2-E, and not the HC HII region W51e2-W. This further supports the scenario proposed by Shi et al. (2010a), where the ongoing star formation activity in the region is not concentrated on the HC HII region but toward its companion 1\u2032\u2032 to the east. ","Citation Text":["Zhang & Ho 1997"],"Functions Text":["VLA observations of the NH3 inversion lines (1,\u20091) and (2,\u20092) seen in absorption (1\\hbox{$\\farcs$}.\u030b1 beamsize) revealed radial infall on scales larger than 5000 AU toward the W51e2 core"],"Functions Label":["Background"],"Citation Start End":[[564,579]],"Functions Start End":[[376,562]]}
{"Identifier":"2021ApJ...920..139M__in_2002_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":["McBride et al. 2007"],"Functions Text":["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"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[2379,2398]],"Functions Start End":[[2060,2377]]}
{"Identifier":"2022MNRAS.512..439C__Lian_et_al._2021_Instance_2","Paragraph":"It is still unclear whether this incompatibility is evidence against the spatially flat \u039bCDM model or is caused by unidentified systematic errors in one of the established cosmological probes or by evolution of the parameters themselves with the redshift (Dainotti et al. 2021b, 2022). Newer, alternate cosmological probes could help alleviate this issue. Recent examples of such probes include reverberation-mapped quasar (QSO) measurements that reach to redshift z \u223c 1.9 (Czerny et al. 2021; Khadka et al. 2021a,b; Yu et al. 2021; Zaja\u010dek et al. 2021), H\u2009ii starburst galaxy measurements that reach to z \u223c 2.4 (Mania & Ratra 2012; Ch\u00e1vez et al. 2014; Gonz\u00e1lez-Mor\u00e1n et al. 2019, 2021; Cao, Ryan & Ratra 2020, 2022a; Cao et al. 2021a; Johnson, Sangwan & Shankaranarayanan 2022; Mehrabi et al. 2022), QSO angular size measurements that reach to z \u223c 2.7 (Cao et al. 2017, 2020, 2021a; Ryan, Chen & Ratra 2019; Lian et al. 2021; Zheng et al. 2021), QSO flux measurements that reach to z \u223c 7.5 (Risaliti & Lusso 2015, 2019; Khadka & Ratra 2020a,b, 2021, 2022; Lusso et al. 2020; Yang, Banerjee & \u00d3 Colg\u00e1in 2020; Li et al. 2021; Lian et al. 2021; Luongo et al. 2021; Rezaei, Sol\u00e0 Peracaula & Malekjani 2021; Zhao & Xia 2021),1 and the main subject of this paper, gamma-ray burst (GRB) measurements that reach to z \u223c 8.2 (Amati et al. 2008, 2019; Cardone, Capozziello & Dainotti 2009; Cardone et al. 2010; Samushia & Ratra 2010; Dainotti et al. 2011, 2013a,b; Postnikov et al. 2014; Wang, Dai & Liang 2015; Wang et al. 2016, 2022; Fana Dirirsa et al. 2019; Khadka & Ratra 2020c; Hu, Wang & Dai 2021; Dai et al. 2021; Demianski et al. 2021; Khadka et al. 2021c; Luongo et al. 2021; Luongo & Muccino 2021; Cao et al. 2021a). Some of these probes might eventually allow for a reliable extension of the Hubble diagram to z \u223c 3\u20134, well beyond the reach of Type Ia supernovae. GRBs have been detected to z \u223c 9.4 (Cucchiara et al. 2011), and might be detectable to z = 20 (Lamb & Reichart 2000), so in principle GRBs could act as a cosmological probe to higher redshifts than 8.2.","Citation Text":["Lian et al. 2021"],"Functions Text":["QSO flux measurements that reach to z \u223c 7.5"],"Functions Label":["Background"],"Citation Start End":[[1125,1141]],"Functions Start End":[[947,990]]}
{"Identifier":"2021AandA...647A..35B__Laffon_et_al._(2010)_Instance_2","Paragraph":"When we now compare the photodesorption yields at 541 eV between fluences  5 \u00d7 1015 photon cm\u22122 and at 3 \u00d7 1017 photon cm\u22122 in Fig. 2, the CO2 photodesorption first increases from 2.6 \u00d7 10\u22122 molecule\/photon to 7.3 \u00d7 10\u22122 molecule\/photon. We also observed this phenomenon for CO photodesorption yield (the data are not shown for more clarity), which increased from 1.9 \u00d7 10\u22122 molecule\/photon to 3.5 \u00d7 10\u22122 molecule\/photon. Second, the estimated yield for the X-ray photodesorption of CH3OH from pure methanol ice decreased by almost one order of magnitude from 9.0 \u00d7 10\u22123 molecule\/photon to 1.3 \u00d7 10\u22123 molecule\/photon. This indi- cates that the photodesorption of CH3OH is higher for a lower fluence received by the ice when more intact methanol molecules are present in the ice. This aging process favors the photodesorption of simpler molecules such as CO2 or CO. Laffon et al. (2010) estimated with NEXAFS spectroscopy (at the C K-edge) that X-ray irradiation at 150 eV of pure methanol ice at 20 K leads to a survival rate of 50% for methanol after an absorbed dose of 1.1 MGy. In our fixed-energy experiments, we irradiated pure methanol ice with fluences between 5 \u00d7 1015 photon cm\u22122 and 2 \u00d7 1016 photon cm\u22122. Because we irradiated a volume of 0.1 cm2 \u00d7 100 ML, with a meanenergy of ~550 eV, and when we consider a volumic mass of condensed methanol of ~ 0.64 g cm\u22123 (at 20 K; Luna et al. 2018) and an X-ray absorption cross section of ~ 0.6 Mbarn (Ishii & Hitchcook 1988), the absorbed doses used in our fixed energy experiments change from ~ 2 to ~ 15 MGy, whichis quite similar to the absorbed doses in Laffon et al. (2010). This indicates that we could expect a methanol destruction rate of about 50% for our low-fluence experiments. In similar experiments, when irradiating a H2 O:CH4:NH3 (2:1:1) ice mixture covered by a layer of CO:CH3OH (3:1) with 250\u20131250 eV X-rays during 120 min with a flux of 7.6 \u00d7 1014 photon s\u22121, higher by almost two order of magnitudes than our experiments, Ciaravella et al. (2020) did not detect a desorption signal on the mass channel 31 (attributed to methanol desorption) and estimated that only ~ 20% of methanol molecules remained intact in the first minutes of the irradiation. The irradiation flux therefore appears to be critical for detecting methanol desorption in X-ray irradiation experiments of methanol-containing ices. A lower X-ray flux appears to favor methanol desorption because the methanol destruction rate is lower. This destruction of methanol molecules could also have a significant effect on the formation and desorption of more complex molecules.","Citation Text":["Laffon et al. (2010)"],"Functions Text":["the absorbed doses used in our fixed energy experiments change from ~ 2 to ~ 15 MGy, whichis quite similar to the absorbed doses in"],"Functions Label":["Similarities"],"Citation Start End":[[1611,1631]],"Functions Start End":[[1479,1610]]}
{"Identifier":"2022MNRAS.515.5267B__Landau_&_Lifshitz_1959_Instance_1","Paragraph":"Recent studies have shown that one can use energy balance arguments that include the large-scale magnetic field, ${\\mathrm{{\\boldsymbol {\\mathit {B}}}}_0}$, to derive scaling laws between the Alfv\u00e9nic and kinetic fluid quantities (Federrath 2016; Beattie et al. 2020; Skalidis & Tassis 2021; S+2021). The dimensionless magnetic energy density, by which we mean the magnetic energy density normalized to the mean thermal pressure3$\\rho _0 c_{\\rm s}^2$, is\n(5)$$\\begin{eqnarray}\r\n{e_{\\mathrm{ mag}}}= \\frac{B^2}{8\\pi c_{\\rm s}^2 \\rho _0} = \\frac{1}{8\\pi c_{\\rm s}^2 \\rho _0}\\Bigg (B_0^2+ \\underbrace{2\\delta \\mathrm{{{\\mathit {B}}}}\\cdot \\mathrm{{{\\mathit {B}}}}_0}_{{\\substack{\\text{coupling} \\\\\\text{term}}}} {} + \\delta B^2 \\Bigg),\r\n\\end{eqnarray}$$where $B_0^2$ is the large-scale field contribution to the total energy, \u03b4B2 is the turbulent field contribution, and $2\\delta \\mathrm{{\\boldsymbol {\\mathit {B}}}}\\cdot \\mathrm{{\\boldsymbol {\\mathit {B}}}}_0$ is the coupling term between the two field components. In the linear perturbation theory limit of the MHD equations, \u03b4B2 includes contributions from shear Alfv\u00e9n, fast and slow magnetosonic compressive eigenmodes (e.g. Landau & Lifshitz 1959). Because $\\delta \\mathrm{{\\boldsymbol {\\mathit {B}}}}\\cdot \\mathrm{{\\boldsymbol {\\mathit {B}}}}_0 = \\delta B_{\\parallel }B_0$, the coupling term only contains the component of magnetic field fluctuations that are parallel to the large-scale field. In linear theory, both fast and slow magnetosonic compressible modes are able to perturb the field variables parallel to ${\\mathrm{{\\boldsymbol {\\mathit {B}}}}_0}$, so under the lens of linear theory, the coupling term is the fluctuation contribution from the compressible modes in the turbulence scaled by B0 (Bhattacharjee, Ng & Spangler 1998). Furthermore, for sub-Alfv\u00e9nic turbulence Beattie et al. (2021b) showed that converging, shocked flows along magnetic field lines excite strong \u03b4B\u2225 fluctuations, which travel roughly at the theoretical fast Alfv\u00e9n mode speed. Therefore, it is likely, assuming that \u03b4B\u2225\/B0 \u226a 1 (this is indeed the case for ${\\mathcal {M}_{\\text{A0}}}\\lt 1$ plasmas; see left panel of fig. 5 in Beattie et al. 2022) where a linear theory may be valid for the magnetic field, the coupling term contains significant energy contributions from fast magnetosonic modes excited by shocked gas that converges and forms dense filaments perpendicular to magnetic field lines.","Citation Text":["Landau & Lifshitz 1959"],"Functions Text":["In the linear perturbation theory limit of the MHD equations, \u03b4B2 includes contributions from shear Alfv\u00e9n, fast and slow magnetosonic compressive eigenmodes (e.g."],"Functions Label":["Uses"],"Citation Start End":[[1178,1200]],"Functions Start End":[[1014,1177]]}
{"Identifier":"2020AandA...641A.155V__Puglisi_et_al._2019_Instance_2","Paragraph":"The scenario presented above has been formulated in various flavors to individually explain several of the properties reported here. The main addition of this work, namely the excitation of CO in distant main-sequence and starburst galaxies, fits in the general picture that we sketched. The ensemble of properties and correlations that we reported here can be also used to revisit the definition of what a starburst is. A standard operational classification is based on the distance from the observed empirical M\u22c6-SFR correlation, the main sequence. This proved to be a useful distinction and an excellent predictor of several trends (e.g., Sargent et al. 2014), but recent results, including our present and previous analysis (Puglisi et al. 2019), show that the demarcation between starburst and main-sequence galaxies is more blurred that we previously considered. We do detect starburst-like behaviors in galaxies on the main sequence (Elbaz et al. 2018), likely linked to the existence of transitional objects (Popping et al. 2017; Barro et al. 2017b; G\u00f3mez-Guijarro et al. 2019; Puglisi et al. 2019, and in prep. to limit the references to recent works based on submillimeter observations). Such transition might well imply an imminent increase of the SFR, driving the object in the realm of starbursts (e.g., Barro et al. 2017b), or its cessation, bringing the system back onto or even below the main sequence (G\u00f3mez-Guijarro et al. 2019; Puglisi et al. 2019), with the CO properties potentially able to distinguish between these two scenarios. Regardless of these transitional objects, a definition of starburst based on \u03a3SFR, rather than \u0394MS, would naturally better account for the observed molecular gas excitation properties, dust temperatures and opacities, or SFE (see also Elbaz et al. 2011; Rujopakarn et al. 2011; Jim\u00e9nez-Andrade et al. 2018; Tacconi et al. 2020). As an example, in Fig. 8 we show the mean SLED of the subsample of galaxies with both CO (2\u2005\u2212\u20051) and CO (5\u2005\u2212\u20054) coverage, split at its median \u03a3SFR. While only tentative at this stage, this suggests a trend of increasing CO excitation with \u03a3SFR, consistently with Fig. 7 and what mentioned above.","Citation Text":["Puglisi et al. 2019"],"Functions Text":["We do detect starburst-like behaviors in galaxies on the main sequence","likely linked to the existence of transitional objects","to limit the references to recent works based on submillimeter observations"],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1086,1105]],"Functions Start End":[[869,939],[961,1015],[1120,1195]]}
{"Identifier":"2018AandA...615A..61C__Gladders_et_al._2013_Instance_1","Paragraph":"To tackle this problem, one must study the recent star formation history (SFH) of galaxies. This information is embeddedin their spectral energy distribution (SED). However, recovering it through SED modeling is complex and subject to many uncertainties and degeneracies. Although an average SFH of galaxies can be derived assuming that they follow the MS (Heinis et al. 2014; Ciesla et al. 2017), galaxies are expected to exhibit complex SFHs, with short-term fluctuations, requiring sophisticated SFH parametrizations to model them (e.g., Lee et al. 2010; Pacifici et al. 2013; Behroozi et al. 2013; Pacifici et al. 2016). However, implementation of these models is complex and a large library is needed to model all galaxies properties. Instead, numerous studies have used simple analytical forms to model galaxies SFH (e.g., Papovich et al. 2001; Maraston et al. 2010; Pforr et al. 2012; Gladders et al. 2013; Simha et al. 2014; Buat et al. 2014; Boquien et al. 2014; Ciesla et al. 2015; Abramson et al. 2016; Ciesla et al. 2016, 2017). Furthermore, SFH parameters are known to be difficult to constrain from broadband SED modeling with a general agreement on the difficulty to constrain the age of the galaxy, here defined as the age of the oldest star, from broad-band SED fitting (e.g., Maraston et al. 2010; Pforr et al. 2012; Buat et al. 2014; Ciesla et al. 2015, 2017). To understand the origin of the scatter of the MS, we need to use an analytical SFH that is able to recover recent variations of the SFR with a precision better than the scatter of the MS itself, that is, 0.3 dex. Recently, Ciesla et al. (2017) showed that a delayed SFH to which we add a flexibility on the recent SFH provides SFRs that are more accurate than those estimated by other typical analytical SFHs (\u03c4 models, delayed, etc.). This SFH was tested in Ciesla et al. (2016). Studying a sample of local galaxies from the Virgo cluster, that is, galaxies known to have undergone a fast drop of star formation activity due to ram pressure stripping, we showed that the amplitude of the flexibility can be constrained by broadband SED modeling as long as ultraviolet (UV) rest frame and near-IR data are available.","Citation Text":["Gladders et al. 2013"],"Functions Text":["However, implementation of these models is complex and a large library is needed to model all galaxies properties. Instead, numerous studies have used simple analytical forms to model galaxies SFH (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[892,912]],"Functions Start End":[[625,828]]}
{"Identifier":"2022ApJ...938...92B__Zikanov_&_Thess_1998_Instance_1","Paragraph":"Flows with \n\n\n\nRem\u226a1\n\n and \n\n\n\nN\u223c\ue23b(1)\n\n have the distinct property that the induced magnetic field is quickly diffused away, yet the Lorentz force is not negligible. This limit is referred to as the quasi-static approximation to MHD (which we call \u201cQMHD\u201d henceforth; Moffatt 1967; Sommeria & Moreau 1982; Davidson 1995; Knaepen & Moreau 2008; Davidson 2013), and has been studied mainly in metallurgy and in MHD experiments due to the typically low conductivity of liquid metals (Alemany et al. 1979; Sommeria 1988; Gallet et al. 2009; Klein & Poth\u00e9rat 2010; Poth\u00e9rat & Klein 2014; Baker et al. 2018), although recent numerical studies on its turbulent properties and anisotropy have been done as well (Zikanov & Thess 1998; Burattini et al. 2008; Favier et al. 2010, 2011; Reddy & Verma 2014; Verma 2017). After nondimensionalizing the equations of MHD using the uniform density \u03c1, \u2113, and u, and taking the limits above, one is left with a single dynamical equation for the velocity\n2\n\n\n\n\u2202v\u2202t+v\u00b7\u2207v=\u2212\u2207p*\u2212Ro\u22121x\u02c6\u2225\u03a9\u00d7v\u2212N\u2207\u22122(x\u02c6\u2225B0\u00b7\u2207)2v+F,\n\nwhere p\n* is the total pressure modified by rotation and magnetic pressure, Ro\n\u22121 \u2261 2\u03a9\u2113\/u is the inverse Rossby number (quantifying the relative strength of the Coriolis force), \n\n\n\nx\u02c6\u2225\u03a9\n\n and \n\n\n\nx\u02c6\u2225B0\n\n are unit vectors in the direction of rotation and the background magnetic field, respectively, and \nF\n is a generic forcing term that can include dissipation such as viscosity and a body force (to be specified in Section 3). The background magnetic field is fixed in time and is uniform in space, such that \n\n\n\n\u2207\u00d7B0=B0\u2207\u00d7x\u02c6\u2225B0=0\n\n. Care must be taken if considering a spatially dependent background magnetic field \nB\n\n\n0\n(\nx\n), as the resulting equation will not be the same. See the discussion in Section 5. This equation is accompanied with the incompressibility condition \u2207 \u00b7 \nv\n = 0. The induced magnetic field can be found using a diagnostic relation\n3\n\n\n\nb=\u2212\u2207\u22122x\u02c6\u2225B0\u00b7\u2207v,\n\nwhich would be \n\n\n\nb=\u2212\u2207\u22122B0\u00b7\u2207v\/\u03b7\n\n in dimensional variables.","Citation Text":["Zikanov & Thess 1998"],"Functions Text":["This limit is referred to as the quasi-static approximation to MHD","although recent numerical studies on its turbulent properties and anisotropy have been done as well"],"Functions Label":["Background","Background"],"Citation Start End":[[703,723]],"Functions Start End":[[166,232],[602,701]]}
{"Identifier":"2016MNRAS.461.2328M__Smith_et_al._2005_Instance_1","Paragraph":"This work, which focuses on mass reconstruction from gravitational lensing, is only the first in a series, which will exploit our mesh-free numerical techniques. Two regimes are typically distinguished in lensing mass reconstruction. Strong lensing is usually confined to the inner-most core of the gravitational lens and produces spectacular observational constraints such as multiple images of the same source, gravitational arcs or even rings. The domain of weak lensing is further away from the centre of the lens but spans large areas and manifests itself by the weak distortion in the shape of background galaxies behind the lens. Reconstruction techniques are divided into two classes, although this distinction is by no means unique or even consistent in some cases. Parametric techniques (e.g. Kneib et al. 1996; Broadhurst et al. 2005; Smith et al. 2005; Halkola, Seitz & Pannella 2006; Jullo et al. 2007; Zitrin et al. 2009; Oguri 2010; Newman et al. 2013; Jullo et al. 2014; Monna et al. 2014; Johnson et al. 2014, for some recent examples) assume a parametric form of the underlying mass density distribution for the lens and typically make the assumption that light traces mass in the positioning of these parametric forms. On the other hand, free-form1 methods (see e.g. Broadhurst, Taylor & Peacock 1995; Bartelmann et al. 1996; Abdelsalam, Saha & Williams 1998; Bridle et al. 1998; Seitz, Schneider & Bartelmann 1998; Brada\u010d et al. 2005a; Cacciato et al. 2006; Liesenborgs, De Rijcke & Dejonghe 2006; Diego et al. 2007; Jee et al. 2007; Coe et al. 2008; Brada\u010d et al. 2009; Merten et al. 2009; Williams & Saha 2011; Merten et al. 2011, 2015; Diego et al. 2015, for some recent examples) usually do not make this assumption and purely rely on the input data either based on weak lensing, strong lensing or a combination of the two. This is possible while using a reconstruction mesh and directly inverting the underlying equations describing lensing on this mesh. In the following, we introduce a free-form method combining weak and strong lensing, which uses our new mesh-free numerical framework. This method translates original ideas by Bartelmann et al. (1996), Seitz et al. (1998), Brada\u010d et al. (2005a), Cacciato et al. (2006) and Merten et al. (2009) into the flexible and efficient mesh-free numerical domain. Alternative implementations of such ideas can e.g. be found in Brada\u010d et al. (2009).","Citation Text":["Smith et al. 2005"],"Functions Text":["Reconstruction techniques are divided into two classes, although this distinction is by no means unique or even consistent in some cases. Parametric techniques (e.g.","assume a parametric form of the underlying mass density distribution for the lens and typically make the assumption that light traces mass in the positioning of these parametric forms."],"Functions Label":["Background","Background"],"Citation Start End":[[846,863]],"Functions Start End":[[637,802],[1053,1237]]}
{"Identifier":"2020ApJ...903L..12H__Zhong_et_al._2020_Instance_1","Paragraph":"Magnetic reconnection (MR) may occur in various space and astrophysical plasma environments, among which the planetary magnetopause boundaries separating the solar wind and magnetospheric origins of plasmas and magnetic field are some of the most likely sites for the occurrence of MR. Due to the easy access to the in situ spacecraft observations the Earth\u2019s magnetopause is the most widely studied space plasma environment for MR (Paschmann et al. 1979; Vaivads et al. 2004; Graham et al. 2014). In particular, the Magnetospheric Multiscale (MMS) mission has contributed greatly to the kinetic physics of magnetopause reconnection (Burch et al. 2016; Hasegawa et al. 2017; Zhong et al. 2020). Many studies have shown that an initial Harris type equilibrium profile with constant total pressures \n\n\n\n\n\n and antiparallel magnetic field with or without a guide field (Harris 1962) may tend to develop MR geometry. In particular, two major categories of MR have been proposed: the steady state model with a single X line and the outflow approaching the Alfv\u00e9n speed (Petschek 1964), and the tearing mode instability with a series of X and O lines and mild plasma velocity (Furth et al. 1963). Numerous fluid and kinetic simulations have been carried out to examine the various aspects of MR processes for the past 50 yr (Hau & Chiou 2001; Guo et al. 2015; Landi et al. 2015). In particular, the effects of pressure or temperature anisotropy on MR have been examined by a number of authors (Chen & Palmadesso 1984; Shi et al. 1987; Birn & Hesse 2001; Chiou & Hau 2002, 2003; Hung et al. 2011). In the MHD models the double-polytropic (DP) laws are widely adopted as the energy closures to study the effects of temperature anisotropy and energy closures on MR and tearing mode instability (Chiou & Hau 2002, 2003; Hung et al. 2011). It is shown that the mirror type temperature anisotropy of \n\n\n\n\n\n may greatly enhance the growth rate of tearing mode instability and the merging rate of single X-line reconnection. In particular, the coupling of tearing and mirror instabilities may lead to relatively larger magnetic islands as compared to the cases with isotropic pressure and the mirror waves with anticorrelated density and magnetic field may be present in the vicinity of X lines.","Citation Text":["Zhong et al. 2020"],"Functions Text":["In particular, the Magnetospheric Multiscale (MMS) mission has contributed greatly to the kinetic physics of magnetopause reconnection"],"Functions Label":["Background"],"Citation Start End":[[675,692]],"Functions Start End":[[498,632]]}
{"Identifier":"2015MNRAS.450.3458C__Cichowolski_et_al._2001_Instance_7","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"],"Functions Text":["see fig. 1 of"],"Functions Label":["Background"],"Citation Start End":[[2785,2808]],"Functions Start End":[[2771,2784]]}
{"Identifier":"2022MNRAS.509.1504M__Komissarov_2006_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":["Komissarov 2006"],"Functions Text":["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."],"Functions Label":["Background"],"Citation Start End":[[1055,1070]],"Functions Start End":[[843,1054]]}
{"Identifier":"2018MNRAS.480.1081K__Hunter_et_al._2008_Instance_1","Paragraph":"Work in recent years has shown that rotation is a key ingredient in shaping the evolution of massive stars with very low metallicities (Z  SMC metallicity down to Population III stars; see Meynet & Maeder 2017, and references therein). Fast rotators (with initial vrot \u2273300\u2009km\u2009s\u22121) are expected to lead to chemically homogeneous evolution (CHE) in which the star becomes brighter and hotter, and thus more ionizing photons, specially in the extreme UV, are emitted than in the corresponding non-rotating case (e.g. Brott et al. 2011; Levesque et al. 2012; Yoon, Dierks & Langer 2012). Presently, although we still know little about rotation velocities of massive stars and their variation with environment, observations seem to favor fast rotators at low Z (e.g. Martayan et al. 2007; Hunter et al. 2008). Model predictions suggest that the effects of rotation, like CHE, should be enhanced at lower metallicities. There is an increase of theoretical and observational evidence which supports the significant role of rotation among the generations of first stars, with Z=0 or extremely low metallicities, and fast rotating massive stars were likely common phenomena in the early Universe (e.g. Leitherer 2008; Chiappini et al. 2008; Maeder & Meynet 2012; Yoon et al. 2012). Modern models for low-metallicity massive fast-rotating single stars which undergo CHE have been published by Sz\u00e9csi et al. (2015); the authors called them transparent wind UV intense stars (or TWUIN stars). Considering all this and the extremely low-Z of SBS\u20090335 \u2212 052E, we compare our observations with the TWUINs predictions,6 following the same approach that we used previously for non-rotating WCE and massive O stellar models. Taking the computed Q(He\u2009ii)=7.37 \u00d7 1048\u2009photon\u2009s\u22121 for the most massive 294\u2009M\u2299 TWUIN, we require that \u223c430 (\u223c300) such stars are necessary to explain the Q(He\u2009ii)Int (Q(He\u2009ii)$_{\\rm He\\, \\small {II}-MB}$). Again, hundreds of these super-massive TWUINs do not match the M\u22c6, SSCs of SBS\u20090335 \u2212 052E. If we instead, apply the same approach using lower mass TWUIN models, it makes even harder to account for the observations. Besides the Q(He\u2009ii) budget, we have measured values of He\u2009ii \u03bb4686\/H\u03b2 as high as 0.06 within the He\u2009ii \u03bb4686 main body region. Under ionization-bounded conditions, even the most massive TWUIN models cannot reproduce these values of He\u2009ii \u03bb4686\/H\u03b2 (see table B.1 from Sz\u00e9csi et al. 2015).","Citation Text":["Hunter et al. 2008"],"Functions Text":["Presently, although we still know little about rotation velocities of massive stars and their variation with environment, observations seem to favor fast rotators at low Z (e.g."],"Functions Label":["Background"],"Citation Start End":[[785,803]],"Functions Start End":[[585,762]]}
{"Identifier":"2019MNRAS.488.1066C__Link_et_al._1992_Instance_1","Paragraph":"A pulsar rotates at an angular velocity \u03a9 with inertial moment I and magnetic moment $\\boldsymbol {\\mathcal {M}}$. It is described by the canonical rotation-down equation in the magnetic dipole model:\n(4)\r\n\\begin{eqnarray*}\r\n\\dot{\\Omega } = - K \\Omega ^3,\r\n\\end{eqnarray*}\r\nwhere $K = \\frac{\\mu _0 | \\boldsymbol{\\mathcal {M}} |^2 \\sin ^2 \\theta }{6 \\pi c^3 I}$, \u03bc0 is vacuum permeability, and c is speed of light. \u03b8 is inclination angle, namely, the angle between $\\boldsymbol {\\mathcal {M}}$ and $\\boldsymbol {\\Omega }$. The possible change of macroscopic magnetic moment of pulsars at the moment of a glitch occurring was discussed by Link and Epstein (Link et al. 1992; Link & Epstein 1997). Namely, K and \u03a9 changed while the rotation-down equation remains the same at glitches, \n(5)\r\n\\begin{eqnarray*}\r\n\\frac{\\Delta \\dot{\\Omega }}{\\dot{\\Omega }} \\approx \\frac{\\Delta K}{K} = 2 \\frac{\\Delta (| \\boldsymbol{\\mathcal {M}} | \\sin \\theta)}{| \\boldsymbol{\\mathcal {M}} | \\sin \\theta } - \\frac{\\Delta I}{I} .\r\n\\end{eqnarray*}\r\nFrom observations, we know $\\frac{\\Delta \\dot{\\Omega }}{\\dot{\\Omega }} \\sim 10^{-3}$. Merely variation of orbit angular momentum cannot describe the magnetic moment changing, as $- \\frac{\\Delta I}{I} = \\frac{\\Delta \\Omega }{\\Omega } \\lesssim 10^{-6} \\ll \\frac{\\Delta \\dot{\\Omega }}{\\dot{\\Omega }}$. It motivates us to consider the origin of the equation (5) related with conservation of total angular momentum. As Einstein-de Haas effect suggests that $\\boldsymbol {S}$ and $\\boldsymbol {\\Omega }$ can transfer into each other caused by the total angular momentum conservation, \n(6)\r\n\\begin{eqnarray*}\r\n\\Delta \\boldsymbol {J} = \\Delta (\\boldsymbol {S} + I \\boldsymbol {\\Omega }) = \\boldsymbol {T} \\Delta t \\approx 0,\r\n\\end{eqnarray*}\r\nwhere $\\boldsymbol {S}=(1\/\\gamma)\\boldsymbol {\\mathcal {M}}$, \u0394t is the duration of a glitch, and $\\boldsymbol {T}$ is the torque. The total angular momentum conservation would turn to be reasonable, when glitches regarding as instantaneous processes, namely \u0394t \u2192 0, as the traditional glitch model (Baym et al. 1969) does. In principal, $\\boldsymbol {T} \\Delta t \\ne 0$ is permitted in our model. For the first step, we wish to obtain approximated results with the simplest case of equation (6).","Citation Text":["Link et al. 1992"],"Functions Text":["The possible change of macroscopic magnetic moment of pulsars at the moment of a glitch occurring was discussed by Link and Epstein"],"Functions Label":["Background"],"Citation Start End":[[655,671]],"Functions Start End":[[522,653]]}
{"Identifier":"2016MNRAS.459.1422E__Ebrahimi_&_Bhattacharjee_2014_Instance_1","Paragraph":"We begin with our results from global DNS MHD simulations of the MRI in cylindrical (r, \u03d5, z) geometry using the DEBS (Schnack et al. 1987; Ebrahimi et al. 2009) initial-value code to solve the non-linear, viscous and resistive MHD equations\n\n(1)\n\n\\begin{eqnarray}\n\\frac{\\mathrm{\\partial} \\boldsymbol A }{ \\mathrm{\\partial} t } &amp;=&amp; -{\\boldsymbol E} = S\\boldsymbol V\\times \\boldsymbol B - \\eta \\boldsymbol J\n\\end{eqnarray}\n\n\n(2)\n\n\\begin{eqnarray}\n\\rho \\frac{\\mathrm{\\partial} \\boldsymbol V }{ \\mathrm{\\partial} t } &amp;=&amp; -S \\rho \\boldsymbol V\\cdot \\ \\nabla \\boldsymbol V + S\\boldsymbol J \\times \\boldsymbol B +P_{\\rm m} \\nabla ^2 \\boldsymbol V -S \\frac{\\beta _0}{2}\\nabla P\n\\end{eqnarray}\n\n\n(3)\n\n\\begin{eqnarray}\n\\frac{\\mathrm{\\partial} P }{ \\mathrm{\\partial} t } &amp;=&amp; -S\\nabla \\cdot (P \\boldsymbol V) - S (\\Gamma -1) P \\nabla \\cdot \\boldsymbol V\n\\end{eqnarray}\n\n\n(4)\n\n\\begin{eqnarray}\n\\frac{\\mathrm{\\partial} \\rho }{ \\mathrm{\\partial} t } &amp;=&amp; -S\\nabla \\cdot (\\rho \\boldsymbol V)\n\\end{eqnarray}\n\n\n(5)\n\n\\begin{eqnarray}\n\\boldsymbol B &amp;=&amp; \\nabla \\times \\boldsymbol A\n\\end{eqnarray}\n\n\n(6)\n\n\\begin{eqnarray}\n\\boldsymbol J &amp;=&amp; \\nabla \\times \\boldsymbol B,\n\\end{eqnarray}\n\nwhere the variables, \u03c1, P, V, B, J, and \u0393 are the density, pressure, velocity, magnetic field, current, and ratio of the specific heats, respectively. We use the same normalization (Schnack et al. 1987; Ebrahimi et al. 2009; Ebrahimi & Bhattacharjee 2014), where time, radius and velocity are normalized to the outer radius a, the resistive diffusion time \u03c4R = a2\/\u03bc0\u03b7, and the Alfv\u00e9n velocity $V_{\\rm A} = B_0\/\\sqrt{\\mu _0 \\rho _0}$, respectively. The dimensionless parameters, S = \u03c4RVA\/a and Pm, are the Lundquist number and the magnetic Prandtl number (the ratio of viscosity to resistivity), respectively. the initial state satisfies the equilibrium force balance condition $\\frac{\\beta _0}{2}\\nabla p = \\rho V_{\\phi }^2\/r$, where $\\beta _0 \\equiv 2 \\mu _0 P_0\/B_0^2$ is normalized to the axis value, and the initial pressure and density profiles are assumed to be radially uniform and unstratified. Pressure and density are evolved, however, they remain fairly uniform during the computations. A no-slip boundary condition is used for the poloidal flow and flow fluctuations. The inner and outer radial boundaries are perfectly conducting so that the tangential electric field, the normal component of the magnetic field, and the normal component of the velocity vanish. The tangential component of the velocity is the rotational velocity of the wall. The azimuthal (\u03d5) and axial (z) boundaries are periodic. We assume a radial pressure gradient balances the centrifugal force in equilibrium, but radial gravity and a radial pressure force are interchangeable for our incompressible, unstratified circumstance. The pressure gradient, rather than gravity, is what balances the centrifugal force in cylindrical laboratory experiments designed to test the MRI (Goodman & Ji (2002)).","Citation Text":["Ebrahimi & Bhattacharjee 2014"],"Functions Text":["We use the same normalization","where time, radius and velocity are normalized to the outer radius a, the resistive diffusion time \u03c4R = a2\/\u03bc0\u03b7, and the Alfv\u00e9n velocity $V_{\\rm A} = B_0\/\\sqrt{\\mu _0 \\rho _0}$, respectively."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1436,1465]],"Functions Start End":[[1362,1391],[1468,1658]]}
{"Identifier":"2021MNRAS.500.1772N__Siegel_&_Ciolfi_2015_Instance_1","Paragraph":"While these early studies demonstrated the viability of neutron star mergers as a major r-process site, they identified only one ejection channel: \u2018dynamical ejecta\u2019 that are tidally flung out by gravitational torques. Since they are never substantially heated, these ejecta carry their original \u03b2 \u2212equilibrium electron fraction from the original neutron star, Ye \u2248 0.05, and this enormous neutron-richness allows them to undergo a \u2018fission cycling\u2019 process (Goriely, Bauswein & Janka 2011; Korobkin et al. 2012), which produces a very robust r-process abundance distribution close to the solar pattern for A \u2265 130, but hardly any lighter r-process elements. Oechslin, Janka & Marek (2007) pointed out that there is a second channel of mass ejection that also happens on a dynamical time-scale: shock-heated matter from the interface where the stars come into contact. As of today, many more mass ejection channels have been discussed: matter that becomes unbound on secular time-scales (\u223c1 s) from the post-merger accretion torus (Beloborodov 2008; Metzger, Piro & Quataert 2008; Fernandez & Metzger 2013; Fernandez et al. 2015; Just et al. 2015; Siegel & Metzger 2017, 2018; Fernandez et al. 2019; Miller et al. 2019a), as MHD-driven winds (Siegel & Ciolfi 2015) and by viscous effects (Shibata, Kiuchi & Sekiguchi 2017; Radice et al. 2018a; Shibata & Hotokezaka 2019) from a long-lived neutron star merger remnant. Similar to the case of proto-neutron stars, the enormous neutrino luminosities (>1053 erg s\u22121) after a neutron star merger can also drive substantial matter outflows (Ruffert et al. 1997; Rosswog & Ramirez-Ruiz 2002; Dessart et al. 2009; Perego et al. 2014; Martin et al. 2015; Radice et al. 2018b). The secular torus ejecta contain approximately 40 per cent of the initial torus mass and, although the latter may vary substantially from case to case, they likely contribute the lion\u2019s share to the total ejecta mass. Due to their different thermal histories and exposure times to neutrinos, the ejecta channels can have different electron fractions Ye and therefore different nucleosynthesis yields.1 For electron fractions below a critical value, $Y_{\\rm e}^{\\rm crit}\\approx 0.25$ (Korobkin et al. 2012; Lippuner & Roberts 2015), lanthanides and actinides are efficiently produced, which, due to their open f-shells, have particularly high bound\u2013bound opacities (Barnes & Kasen 2013; Kasen, Badnell & Barnes 2013; Tanaka & Hotokezaka 2013; Tanaka et al. 2020) and therefore lead to red transients that peak days after the merger. Ejecta with electron fractions above $Y_{\\rm e}^{\\rm crit}$, in contrast, only produce \u2018lighter\u2019 elements with lower opacities and thus result in bluer transients that peak after about 1 d. Opaque, low-Ye ejecta blocking the view on high-Ye ejecta can lead to a \u2018lanthanide curtaining\u2019 effect (Kasen, Fern\u00e1ndez & Metzger 2015; Wollaeger et al. 2018), which will efficiently block blue light. Therefore, it is important to understand the layering, dynamics, interaction and potential mixing of different ejecta channels.","Citation Text":["Siegel & Ciolfi 2015"],"Functions Text":["As of today, many more mass ejection channels have been discussed:","as MHD-driven winds"],"Functions Label":["Background","Background"],"Citation Start End":[[1243,1263]],"Functions Start End":[[869,935],[1222,1241]]}
{"Identifier":"2016AandA...586A..80O__Fornasier_et_al._2015_Instance_3","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"],"Functions Text":["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","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."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[2053,2074]],"Functions Start End":[[1893,2051],[2076,2328]]}
{"Identifier":"2020MNRAS.492.3241V__Frebel_et_al._2006_Instance_1","Paragraph":"Progress in this field will require large statistical samples of metal-poor stars in a variety of environments within the Local Group. Unfortunately, metal-poor stars are exceedingly rare and difficult to find, being overwhelmed by the more numerous metal-rich populations in the Galaxy. Examination of the Besan\u00e7on model of the Galaxy (Robin et al. 2003), which is guided by a theoretical framework for the formation and evolution of the main stellar populations, suggests that a typical halo field has only one in \u223c2000 stars with [Fe\/H]  \u22123 between 14  V  18 (Youakim et al. 2017). Enormous effort has gone into the discovery and study of extremely, ultra, and hyper metal-poor stars with [Fe\/H]  \u22123.0, \u22124.0, and \u22125.0, respectively. Most of the known metal-poor stars have been found in dedicated surveys, such as objective prism surveys (the HK survey and Hamburg-ESO survey, Beers, Preston & Shectman 1992; Christlieb et al. 2002, 2008; Beers & Christlieb 2005; Frebel et al. 2006; Sch\u00f6rck et al. 2009), wide-band photometric surveys (Schlaufman & Casey 2014), and blind spectroscopic surveys, such as the Sloan Digital Sky Survey (SDSS) SEGUE and BOSS surveys (Yanny et al. 2009; Eisenstein et al. 2011; Dawson et al. 2013), and from the Large Sky Area Multi-Object Fibre Spectroscopic Telescope (LAMOST; Cui et al. 2012). According to the SAGA data base (see Suda et al. 2017, and references therein), there are \u223c500 stars with [Fe\/H]  \u22123.0, though fewer than half have detailed chemical abundances. Recently, narrow-band photometric surveys have shown higher success rates for finding metal-poor stars, particularly SkyMapper (Keller et al. 2007; DaCosta et al. 2019) and the Pristine survey (Starkenburg et al. 2017b; Youakim et al. 2017; Aguado et al. 2019a). Pristine photometry with follow-up Keck II\/DEIMOS spectroscopy has also been used to increase sample sizes and improve the chemodynamical studies of faint satellites (Draco II and Sgr II, Longeard et al. 2018, 2019). At the same time, Simon (2018) has shown that Gaia DR2 proper motion cleaning may also be a promising way to find new metal-poor members of ultra-faint dwarf galaxies.","Citation Text":["Frebel et al. 2006"],"Functions Text":["Most of the known metal-poor stars have been found in dedicated surveys, such as objective prism surveys (the HK survey and Hamburg-ESO survey"],"Functions Label":["Background"],"Citation Start End":[[967,985]],"Functions Start End":[[736,878]]}
{"Identifier":"2018ApJ...860...88B__Paradijs_1978_Instance_1","Paragraph":"Thermonuclear (type-I) X-ray bursts are intermittently observed from many neutron star low-mass X-ray binaries (LMXBs; Strohmayer & Bildsten (2006) and references therein). Such a burst originates from an unstable thermonuclear burning of the accreted matter accumulated on the neutron star surface (Joss 1977; Lamb & Lamb 1978; Strohmayer & Bildsten 2006). For most bursts, the observed X-ray intensity rises in \u22480.5\u20135 s, decays in \u223c10\u2013100 s as the neutron star surface cools down after the nuclear burning, and the typical recurrence time is a few hours to days (Galloway et al. 2008). The burst spectrum is traditionally described with a blackbody model, and the best-fit burst blackbody normalization, which can be identified as the burst emission area, usually matches well with the expected surface area of a neutron star (Hoffman et al. 1977; Swank et al. 1977). These motivated an effort to measure the neutron star radius using the burst continuum spectrum, where the normalization of the burst blackbody is expected to be proportional to the square of the stellar radius (e.g., van Paradijs 1978; Goldman 1979; van Paradijs 1979; van Paradijs & Lewin 1986). Note that such a radius measurement is extremely important to probe the superdense and degenerate core matter of neutron stars, which is a fundamental problem of physics (e.g., Lattimer & Prakash 2007; Bhattacharyya et al. 2017). However, a reliable radius measurement using this method has so far not been possible due to a number of systematic uncertainties, which include (1) burst emission from and the visibility of an unknown fraction of the neutron star surface and (2) plausible deviation of the burst spectrum from a blackbody, etc. (Bhattacharyya 2010; Bhattacharyya et al. 2010; Kajava et al. 2017a). Nevertheless, the use of continuum burst spectrum remains a promising method to measure the neutron star radius for the following reasons: (1) neutron star LMXBs provide many complementary methods to measure neutron star parameters, the joint application of which has a potential to significantly reduce the systematics (Bhattacharyya 2010) and (2) it was generally believed that the much more intense burst emission could be reliably distinguished from the persistent (i.e., non-burst) emission during the burst.","Citation Text":["van Paradijs 1978"],"Functions Text":["These motivated an effort to measure the neutron star radius using the burst continuum spectrum, where the normalization of the burst blackbody is expected to be proportional to the square of the stellar radius (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[1088,1105]],"Functions Start End":[[870,1087]]}
{"Identifier":"2019AandA...632A.104G__Hirabayashi_et_al._2016_Instance_1","Paragraph":"Finally, our observations are consistent with the bilobate shape of the nucleus of comet 8P\/Tuttle. As noted in Sect. 1, this shape is likely common among comets because it was found for four out of the six comets for which we have spatially resolved images. This is also the case of the trans-Neptunian object 2014 MU69 (Ultima Thule) observed by the New Horizon spacecraft (Stern et al. 2019). This binary configuration has some implications for the formation and evolution of 8P\/Tuttle. A contact binary could result from (i) the accretion at low velocity of two primordial objects (Massironi et al. 2015; Davidsson et al. 2016), (ii) the disruption of a monolithic object due to excessive spin-up resulting from non-gravitational forces or YORP5 effect followed by a reaccretion (Boehnhardt 2004; \u0106uk 2007; Hirabayashi et al. 2016), or (iii) the catastrophic disruption of a monolithic object by a collision followed by a re-accretion (Jutzi & Benz 2017; Schwartz et al. 2018). On the one hand, with a low thermal inertia compared with NEAs, the YORP effect is low for comets, in particular for NIC, which have an elongated orbit and spend most of their time far from the Sun, and it may not be sufficient to increase the spin rate of the nucleus to the point where centrifugal exceed gravitational forces. On the other hand, comet 8P\/Tuttle has been on a very stable orbit for centuries, and it is likely an evolved comet, as suggested by its low activity, so that it could have been much more active in the past. For cometary nuclei, the primary cause for spin-up is torques caused by outgassing, therefore it is possible that 8P\/Tuttle formed as a monolithic body and became a contact binary after its injection into the inner Solar System as a result of excessive spin-up resulting from non-gravitational forces. This scenario has been proposed for comet 67P\/Churyumov-Gerasimenko by Hirabayashi et al. (2016). Alternatively, if the binary nature of comet 8P\/Tuttle is the result of a primordial accretion or a catastrophic collision in the early Solar Sytem, it could have persisted until now. Similar examples are offered by some binary asteroids that can be stable over the age of the Solar System (Chauvineau et al. 1991), or as proposed by Davidsson et al. (2016) for comet 67P\/Churyumov-Gerasimenko. For comet 8P\/Tuttle, it is however not possible to distinguish the solution of a binary nucleus that formed in the first billion years of our Solar System (e.g., Matonti et al. 2019) from a more recent origin following its injection into the inner Solar System (e.g., Hirabayashi et al. 2016).","Citation Text":["Hirabayashi et al. 2016"],"Functions Text":["A contact binary could result from","(ii) the disruption of a monolithic object due to excessive spin-up resulting from non-gravitational forces or YORP5 effect followed by a reaccretion"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[811,834]],"Functions Start End":[[490,524],[633,782]]}
{"Identifier":"2018ApJ...867..120H__Takahashi_et_al._1990_Instance_1","Paragraph":"Although the magnetic (i.e., one-photon) pair production is also taken into account, most electron\u2013positron pairs are found to be produced via photon\u2013photon (i.e., two-photon) collisions, which take place via two paths. One path is through the collisions of two MeV photons, both of which were emitted from the equatorial ADAF. Another path is through the collisions of TeV and eV photons; the former photons were emitted by the gap-accelerated leptons via inverse Compton process, while the latter were emitted from the ADAF via the synchrotron process. There is, indeed, a third path, in which the gap-emitted GeV curvature photons collide with the ADAF-emitted keV inverse Compton photons; however, this path is negligible, particularly when \n\n\n\n\n\n. If the pairs are produced via TeV\u2013eV collisions (i.e., via the second path) outside the gap outer boundary, they have outward ultrarelativistic momenta to easily \u201cclimb up the hill\u201d of the potential k0 (Takahashi et al. 1990; see also Figure 2 of HP16) and propagate to large distances without turning back. However, if the pairs are produced via MeV\u2013MeV collisions (i.e., via the first path), they are produced with subrelativistic outward momenta; thus, they eventually return to fall onto the horizon owing to the strong gravitational pull inside the separation surface (Figure 2 of HP16). When the returned pairs arrive at the gap outer boundary, only positrons can penetrate into the gap because of E\u2225  0. Accordingly, electrons accumulate at the boundary, whose surface charge leads to the jump of the normal derivative of E\u2225. Thus, although the stationary gap solutions show that the \u03b3-ray spectrum little depends on the injected current density (Section 5.5), the gap solution inevitably becomes time dependent owing to the increasing discontinuity of \n\n\n\n\n\n with an accumulated surface charge (in this case, electrons) at the outer boundary. If the injected current is much small compared to the GJ current, the time dependence will be mild. However, if the injected current becomes a good fraction of the GJ current, the assumption of the stationarity becomes invalid, as pointed out by Levinson & Segev (2017). In this sense, a caution should be made in the applicability of the stationary solutions presented in this paper, when the injected current is non-negligible compared to the current created within the gap.","Citation Text":["Takahashi et al. 1990"],"Functions Text":["If the pairs are produced via TeV\u2013eV collisions (i.e., via the second path) outside the gap outer boundary, they have outward ultrarelativistic momenta to easily \u201cclimb up the hill\u201d of the potential k0","see also Figure 2 of HP16) and propagate to large distances without turning back."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[956,977]],"Functions Start End":[[753,954],[979,1060]]}
{"Identifier":"2021AandA...655A..12T__Tang_et_al._2017b_Instance_3","Paragraph":"Using the RADEX3 non local thermodynamic equilibrium (LTE) modeling program (van der Tak et al. 2007) with collisional rate coefficients from Wiesenfeld & Faure (2013), we modeled the relation between the gas kinetic temperature and the measured average of para-H2CO 0.5 \u00d7 [(322\u2013221 + 321\u2013220)\/303\u2013202] ratios, adopting a 2.73 K background temperature, an average observational linewidth of 4.0 km s\u22121, and column densities N(para-H2CO) = 2.7 \u00d7 1012 and 3.7 \u00d7 1012 cm\u22122 for N113 and N159W, respectively. The results are shown in Fig. 5. The values of the para-H2CO column density were obtained with APEX data (beam size ~30\u2033; Tang et al. 2017b), which cover similar regions. Different column densities of para-H2CO only weakly affect derived kinetic temperatures (see Fig. 3 in Tang et al. 2017b or Fig. 4 in Tang et al. 2018a; this was also shown in Fig. 13 and discussed in Sect. 4.3.1 of Mangum & Wootten 1993) as long as all lines are optically thin. Considering that the relation between the gas temperature and the para-H2CO line ratio may vary at different spatial densities (see Fig. 2 in Tang et al. 2017b), we modeled it at spatial densities 104, 105, and 106 cm\u22123 in Fig. 5. It appears that Tkin at n(H2) = 105 cm\u22123 is consistently lower than values at 104 and 106 cm\u22123 by \u227223% and \u227234%, respectively, for Tkin \u2272 100 K. Local thermodynamic equilibrium (LTE) is a good approximation for the H2CO level populations under optically thin and high-density conditions (Mangum & Wootten 1993; Tang et al. 2017a,b, 2018b). Following the method applied by Tang et al. (2017b) in their Eq. (2), we plot the relation between the LTE kinetic temperature, TLTE, and the para-H2CO (3\u20132) line ratio in Fig. 5. Apparently, TLTE agrees well with Tnon-LTE at volume densities n(H2) ~ 105 cm\u22123 as long as Tkin \u2272 100 K. Previous observations show that para-H2CO (3\u20132) is sensitive to gas temperature at density 105 cm\u22123 (Ginsburg et al. 2016; Immer et al. 2016; Tang et al. 2017b). The spatial density measured with para-H2CO (303\u2013202) and C18O (2\u20131) in N113 and N159W is n(H2) ~ 105 cm\u22123 on a size of ~30\u2033 (Tang et al. 2017b). Therefore, here we adopt 105 cm\u22123 as an averaged spatial gas density in the N113 and N159W regions.","Citation Text":["Tang et al. 2017b"],"Functions Text":["Considering that the relation between the gas temperature and the para-H2CO line ratio may vary at different spatial densities (see Fig. 2 in","we modeled it at spatial densities 104, 105, and 106 cm\u22123 in Fig. 5."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1097,1114]],"Functions Start End":[[955,1096],[1117,1185]]}
{"Identifier":"2020MNRAS.492.5675C__P\u00e9rez-Montero_et_al._2019_Instance_2","Paragraph":"In regarding AGNs, the Te method tends to underestimate the oxygen abundance by an average value of about 0.6 dex in comparison to estimations based on strong-line methods and it produces subsolar O\/H values for most of these objects (Dors et al. 2015; Dors, Freitas-Lemes & \u00c2mores 2020). An alternative method to derive the metallicity or abundances in the nuclear regions of spiral galaxies is the extrapolation of the radial oxygen abundance. Along decades, results based on this indirect method have indicated Z near or slightly above the solar value in nuclear regions (Vila-Costas & Edmunds 1992; Zaritsky, Kennicutt & Huchra 1994; van Zee et al. 1998; Pilyugin, V\u00edlchez & Contini 2004; Gusev et al. 2012; Dors et al. 2015; Zinchenko et al. 2019), in consonance with predictions of chemical evolution models (e.g. M\u00f3lla & D\u00edaz 2005) and with the use of strong-line methods (e.g. Groves, Dopita & Sutherland 2004; Groves, Heckman & Kauffmann 2006; Feltre, Charlot & Gutkin 2016; P\u00e9rez-Montero et al. 2019; Thomas, Kewley & Dopita 2019; Dors et al. 2020). Therefore, Temethod does not seem to work for AGNs. The origin of the discrepancy between Z values calculated via Te method and via strong-line methods, the so-called Teproblem, could be attributed, in part, to the presence of heating\/ionization by gas shock in the narrow-line region (NLR) of AGNs. In fact, Contini (2017) carried out detailed modelling of AGN optical emission lines by using the SUMA code (Contini & Aldrovandi 1983) and suggested the presence of gas shock with low velocity ($v \\: \\lesssim \\: 400 \\: \\rm km \\:s^{-1}$) in a sample of Seyfert 2 nuclei. This result is supported by recent spatially resolved observational studies of Seyfert 2 nuclei, in which the presence of gas outflows with velocity of the order of 100\u2013300 $\\rm km \\: s^{-1}$ have been found (e.g. Riffel, Storchi-Bergmann & Riffel 2017; Riffel, Hekatelyne & Freitas 2018). Moreover, the Te problem can also be originated due to the use of an unappropriate calculation of the ionization correction factor (ICF) for oxygen in AGNs (P\u00e9rez-Montero et al. 2019; Dors et al. 2020).","Citation Text":["P\u00e9rez-Montero et al. 2019"],"Functions Text":["Moreover, the Te problem can also be originated due to the use of an unappropriate calculation of the ionization correction factor (ICF) for oxygen in AGNs"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[2078,2103]],"Functions Start End":[[1921,2076]]}
{"Identifier":"2021MNRAS.500.3957E__Springel_&_Hernquist_2003_Instance_1","Paragraph":"The results presented in this paper are based on data from IllustrisTNG,1the next-generation suite of state-of-the-art magnetohydrodynamical cosmological simulations of galaxy formation (Marinacci et al. 2018; Naiman et al. 2018; Nelson et al. 2018; Pillepich et al. 2018b; Springel et al. 2018). Building on the success of its predecessor Illustris (Genel et al. 2014; Vogelsberger et al. 2014a,b; Nelson et al. 2015; Sijacki et al. 2015), IllustrisTNG follows the same fundamental approach but includes improved aspects and novel features in its galaxy formation model and expands its scope to several simulated volumes and improved resolution. The models for galaxy formation include physical processes such as gas heating by a spatially uniform and time-dependent UV background, primordial and metal-line gas cooling, a subgrid model for star formation, and the unresolved structure of the interstellar medium (Springel & Hernquist 2003), as well as models for the evolution and chemical enrichment of stellar populations, which track nine elements (H, He, C, N, O, Ne, Mg, Si, and Fe) in addition to europium and include yields from supernovae Ia, II, and AGB stars (Vogelsberger et al. 2013; Torrey et al. 2014). Furthermore, IllustrisTNG incorporates improved feedback implementations for galactic winds caused by supernovae as well as accretion and feedback from black holes. In particular, depending on accretion, black hole feedback occurs in two modes: low accretion rates result in purely kinetic feedback while high accretion rates invoke thermal feedback (Weinberger et al. 2017). Galactic winds are injected isotropically and the wind particles\u2019 initial speed scales with the one-dimensional dark matter velocity dispersion (Pillepich et al. 2018a). Magnetic fields are amplified self-consistently from a primordial seed field and follow ideal magnetohydrodynamics (Pakmor & Springel 2013). The TNG simulations were run using the moving mesh code arepo (Springel 2010). Here, concepts from adaptive mesh refinement and smooth particle hydrodynamics are combined to create an unstructured, moving Voronoi tessellation. IllustrisTNG follows the \u039bCDM framework, adopting cosmological parameters according to recent constraints from Planck data: matter density $\\Omega _\\rm {m} = 0.3089$, baryonic density $\\Omega _\\rm {b} = 0.0486$, cosmological constant $\\Omega _\\Lambda = 0.6911$, Hubble constant h = 0.6774, normalization \u03c38 = 0.8159, and spectral index $n_\\rm {s} = 0.9667$ (Planck Collaboration I 2016).","Citation Text":["Springel & Hernquist 2003"],"Functions Text":["The models for galaxy formation include physical processes such as gas heating by a spatially uniform and time-dependent UV background, primordial and metal-line gas cooling, a subgrid model for star formation, and the unresolved structure of the interstellar medium"],"Functions Label":["Uses"],"Citation Start End":[[915,940]],"Functions Start End":[[647,913]]}
{"Identifier":"2017AandA...607A.126Y__Gunawardhana_et_al._2011_Instance_1","Paragraph":"Observed high mass end power-law index of the galaxy-wide IMF resulting from the calculated IGIMF, \\hbox{$\\alpha_3^{\\mathrm{gal}}$}\u03b13gal (i.e., \u03b1gal in Eq. (15) for m> 1 M\u2299), for a constant SFR over \u03b4t = 10 Myr in dependence of the galaxy-wide SFR. In Fig. 5, \\hbox{$\\alpha_3^{\\mathrm{gal}}$}\u03b13gal values diverge for different SFRs and also vary for different m at the high mass end. As at each m value there exists a different \\hbox{$\\alpha_3^{\\mathrm{gal}}$}\u03b13gal\u2013SFR relation, we plot solid lines for log\u200910(m\/M\u2299) = 0.2, 0.4, ..., 2, i.e., 1.58, 2.51, ..., 100 M\u2299 from black to gray (top to bottom) for the fiducial model and dotted lines for the corresponding SolarMetal model defined in Sect. 2.1. The blue squares are data from the GAMA galaxy survey (Gunawardhana et al. 2011). The red triangles and red dash-dotted line are data from Weidner et al. (2013b); the left triangle is for the MW field, the middle three triangles are galaxy studies, the right triangle is for the bulges of the MW and M\u200931, and the dash-dotted line is their IGIMF model assuming \u03b2 = 2. A recent study has suggested that the 2 M\u2299\/yr SFR for MW is overestimated (Chomiuk & Povich 2011), but we leave this data point the same as in Weidner et al. (2013b). Gargiulo et al. (2015) report consistency between their IGIMF model assuming \u03b2 = 2 (thick yellow dashed line) and the [\u03b1\/Fe] abundance ratios of elliptical galaxies. The purple diamond is an individual analysis for the dwarf galaxy NGC 2915 (Bruzzese et al. 2015). Green stars are based on the Lee et al. (2009) 11HUGS observations of dwarf galaxies. The black circle is an observation for the solar neighborhood from Rybizki & Just (2015) with adopted MW SFR from Robitaille & Whitney (2010) as an upper limit of the solar neighborhood SFR because the Sun is located in an inter-arm region where the relevant SFR is significantly smaller (toward the direction indicated by the arrow; see Sect. 4.6 for further details). The thin horizontal dashed line represents the canonical IMF index \u03b12 = \u03b13 = 2.3.","Citation Text":["Gunawardhana et al. 2011"],"Functions Text":["The blue squares are data from the GAMA galaxy survey"],"Functions Label":["Uses"],"Citation Start End":[[758,782]],"Functions Start End":[[703,756]]}
{"Identifier":"2015AandA...576L..16P__Li_et_al._(2013)_Instance_1","Paragraph":"To constrain the size of outflows that we could have missed, we performed simple simulations. We placed artificial, unipolar secondary sources next to a primary point source model representing Sgr A* and compared the closure phases obtained from the resulting artificial visibility data with the observations. We considered two geometries: a single-point source and a jet composed of ten equally spaced point sources (knots) with equal fluxes. We probed four orientations for the simulated outflows (see Fig. 2): along the major axis, along the minor axis of the beam, the jet direction claimed by Li et al. (2013), and the jet direction claimed by Yusef-Zadeh et al. (2012). We used total fluxes of 0.2 Jy and 0.55 Jy for the artificial sources; these values ensure that our simulated outflows are sufficiently faint to not violate the constraints given by the known recent brightness evolution of Sgr A* (0.2 Jy from the mean variability of \u224815% from June 2013 to February 2014 at 41 GHz, with 0.55 Jy corresponding to the strongest variation in the same period, Chandler & Sjouwerman 2014). For each simulation setup, we measured the average of absolute values of the closure phases for each triangle. We varied the distances of the model sources (for the jet model: the largest distance) from Sgr A* until we found a critical distance at which the absolute values of the simulated closure phases exceeded those of the observations by more than the 1\u03c3 error at all triangles. We summarize our results in Table 1. As expected, the critical distances are smaller for brighter outflows. Jet-like structures lead to larger critical distances than equally luminous single, compact sources. As a consequence of the very elongated beam, the critical distances for sources located along the major axis of the beam are larger by a factor of \u22487 than for those located along the minor axis. In a few cases (denoted \u201cN\/A\u201d in Table 1), the absolute values of the simulated closure phases were similar to those of the observations for all distances of the model sources, meaning that we were unable to identify a critical distance. Overall, our observations limit the extension of asymmetric (in the observer frame) jet-like outflows from Sgr A* to projected distances of \u22482.5 mas along the major axis and \u22480.4 mas along the minor axis. ","Citation Text":["Li et al. (2013)"],"Functions Text":["We probed four orientations for the simulated outflows","the jet direction claimed by"],"Functions Label":["Uses","Uses"],"Citation Start End":[[598,614]],"Functions Start End":[[444,498],[569,597]]}
{"Identifier":"2020MNRAS.491.5881Y__McConnell_et_al._2011_Instance_1","Paragraph":"There is good observational and theoretical evidence that supermassive black holes (SMBHs) exist in nearly every galaxy in universe. Understanding the properties of these SMBHs will clarify their roles in galaxy formation and evolution across the cosmology history (e.g. Kormendy & Ho 2013). There are mainly two parameters for an SMBH, i.e. mass (MBH) and spin, which need to be determined. For a few very nearby (100 Mpc) quiescent galaxies, including our Galaxy, SMBH masses can be measured through the stellar or gaseous dynamics method (e.g. Tremaine et al. 2002; McConnell et al. 2011). It has been found that nearby quiescent galaxies follow a tight correlation between the central SMBH mass and the bulge or spheroid stellar velocity dispersion (\u03c3*), which is called MBH\u2013\u03c3* relation (e.g. Kormendy & Ho 2013). Active galactic nuclei (AGNs) can be classified into type 1 or type 2 AGNs, depending on whether the broad-line regions (BLRs) can be viewed directly. For type 1 AGNs, the BLR can be used as a probe of the gravitational potential of the SMBHs. The SMBH mass can be weighed through the BLR clouds for type I AGNs across cosmos time. The SMBH masses in type I AGNs can be calculated as follows (e.g. Kaspi et al. 2000; Bian & Zhao 2002; Peterson et al. 2004; Collin et al. 2006; Du et al. 2016a; Yu et al. 2019):\n(1)$$\\begin{eqnarray*}\r\nM_{\\rm BH} =f_{\\rm BLR}\\frac{R_{\\rm BLR}~(\\Delta V)^2}{G},\r\n\\end{eqnarray*}$$where G is the gravitational constant. RBLR is the distance from black hole to the BLRs, and can be estimated from the reverberation mapping (RM) method (e.g. Blandford & McKee 1982; Peterson 1993). \u0394V is the velocity of the BLR clouds, and usually traced by the full width at half-maximum (FWHM) or the line dispersion (\u03c3H\u2009\u03b2) of the broad H\u2009\u03b2 emission line. fBLR is a virial factor to characterize the kinematics, geometry, and inclination of the BLR clouds. Using the MBH\u2013\u03c3* relation, we recently did the calibration of fBLR and found $f_{\\rm BLR} \\propto \\rm FWHM^{-1.11}$ when FWHM(H\u2009\u03b2) is used as the tracer of \u0394V in equation (1) (Mejia-Restrepo et al. 2018; Yu et al. 2019). It is consistent with the results by the BLR dynamical model to fit simultaneously the AGNs continuum\/H\u2009\u03b2 light curves and H\u2009\u03b2 line profiles (e.g. Li et al. 2018; Pancoast et al. 2018; Williams et al. 2018).","Citation Text":["McConnell et al. 2011"],"Functions Text":["For a few very nearby (100 Mpc) quiescent galaxies, including our Galaxy, SMBH masses can be measured through the stellar or gaseous dynamics method (e.g."],"Functions Label":["Background"],"Citation Start End":[[569,590]],"Functions Start End":[[392,546]]}
{"Identifier":"2016ApJ...831...37K__Dadina_2008_Instance_1","Paragraph":"Based on previous works (Kawamuro et al. 2013; Tazaki et al. 2013), we start with a base-line model,\n\n\n\n\n\nin the XSPEC terminology. This model includes absorbed primary X-ray emission (i.e., a cut-off power law), a scattered component, and a reflection continuum from distant cold matter accompanied by a narrow iron-K\u03b1 line. Optically thin thermal emission from the host galaxy (apec in XSPEC) and other emission\/absorption lines (zgauss) are also added if they are significantly required with a confidence level above 90% in terms of \u0394\u03c72. Because it is difficult to determine the cut-off energy from our data, we fix it at 300 keV, which is a typical value measured in nearby AGNs (Dadina 2008). The first constant factor, NXIS, is applied to the primary power-law component in the Suzaku spectra to absorb possible time variability between the Suzaku (one epoch) and Swift\/BAT (averaged for 70 months) observations. The second constant term represents the scattered fraction, fscatt. As a reflection component from the torus, we employ the pexrav model, which calculates a reflected spectrum from an optically thick slab with a solid angle of \u03a9 irradiated by a point source (Magdziarz & Zdziarski 1995). We set the reflection strength, R = \u03a9\/2\u03c0, as a free parameter, and fix the inclination angle at 60\u00b0. It is confirmed that even if 30\u00b0 is adopted, best-fit parameters do not significantly change. The shape of the incident spectrum is assumed to be the same as the power-law component. We basically assume that the reflection and scattered components did not vary between the Suzaku and Swift\/BAT observations, considering that the size of the reflector most likely has a parsec scale. In the three low-mass LLAGNs (NGC 4395, NGC 5273, and NGC 5643), however, we assume that the reflection component varied in accordance with the primary emission because of a smaller size of the emitting regions. Thus, in these targets, R is defined with respect to the primary component in the Suzaku data. The zgauss component represents an iron-K\u03b1 fluorescence line. The line width is fixed at 20 eV, which corresponds to a typical velocity dispersion of \u223c2000 km s\u22121 measured with Chandra\/HETGS in local Seyfert galaxies (Shu et al. 2010). We always consider the Galactic absorption \n\n\n\n\n\n, which is calculated with the nh command (Kalberla et al. 2005) in FTOOLS.","Citation Text":["Dadina 2008"],"Functions Text":["Because it is difficult to determine the cut-off energy from our data, we fix it at 300 keV, which is a typical value measured in nearby AGNs"],"Functions Label":["Uses"],"Citation Start End":[[684,695]],"Functions Start End":[[541,682]]}
{"Identifier":"2019ApJ...875...61M__Marco_2006_Instance_1","Paragraph":"A substantial fraction of metal-poor stars that have recently evolved off the MS, e.g., giants and planetary nebulae (PNe), have been influenced by binary interactions. The IMF is significantly weighted toward low-mass stars (Bastian et al. 2010; Kroupa et al. 2013), and the MW star formation rate was \u22483 times larger \u224810 Gyr ago than it is now (Governato et al. 2007; De Lucia et al. 2014). Based on the measured IMF and modeled galactic star formation history, we estimate that \u224855% of MW giants and PNe have old, solar-type progenitors (\u03c4* > 7 Gyr, M \u2248 0.8\u20131.2 \n\n\n\n\n\n). Such old, low-mass giants tend to be metal-poor (Ratnatunga & Yoss 1991; Carollo et al. 2010; Mackereth et al. 2017). The metallicity trend therefore dramatically affects the properties of low-mass evolved stars. For example, the enhanced close binary fraction of metal-poor solar-type stars substantially strengthens the conclusion that the shaping of PN morphologies is the result of binary interactions (Moe & De Marco 2006; De Marco 2009; Jones & Boffin 2017). Providing further corroboration, Badenes et al. (2015) measured the delay-time distribution of bright PNe in the LMC and discovered two distinct populations of PN progenitors: an old channel (\u03c4* = 5\u20138 Gyr) deriving from solar-type stars (M \u2248 1.0\u20131.2 \n\n\n\n\n\n) and a young channel (35\u2013800 Myr) evolving from late-B\/early-A stars (\u22482\u20138 \n\n\n\n\n\n). According to the measured age\u2013metallicity relation of the LMC (Olszewski et al. 1991; Pagel & Tautvaisiene 1998; Cole et al. 2005; Carrera et al. 2011; Piatti & Geisler 2013), the old, solar-type progenitors are metal-poor ([Fe\/H] \u2272 \u22121.0) and hence have a large close binary fraction of Fclose = 40%\u201350%. The young progenitors have a higher metallicity of [Fe\/H] \u2248 \u22120.4 but are sufficiently massive so that they also have a large close binary fraction of Fclose = 40%\u201360%. Meanwhile, evolved stars with intermediate masses (M \u2248 1.2\u20132.0 \n\n\n\n\n\n) in the LMC have intermediate metallicities and therefore a smaller close binary fraction of Fclose \u2248 30%. If PNe derive from interactions in close binaries, then the variations in Fclose with respect to mass and metallicity can explain the observed bimodal mass\/age distribution of PN progenitors in the LMC.","Citation Text":["Moe & De Marco 2006"],"Functions Text":["For example, the enhanced close binary fraction of metal-poor solar-type stars substantially strengthens the conclusion that the shaping of PN morphologies is the result of binary interactions"],"Functions Label":["Similarities"],"Citation Start End":[[981,1000]],"Functions Start End":[[787,979]]}
{"Identifier":"2020MNRAS.492.2510L__Galsgaard_et_al._2007_Instance_1","Paragraph":"We suggest that the jet\u2013facula collision does not cause a change in the field-line connectivity and only leads to the redistribution of jet material. There are two main reasons supporting this viewpoint. The first is the 3D magnetic configuration at the collision region. The penumbra and facula are both located at negative-polarity magnetic fields and the relative orientation of the two flux systems is nearly parallel (Fig. 5). When the magnetic field lines are parallel it is not easy to have reconnection (Feynman & Martin 1995; Galsgaard et al. 2007). As suggested by Gopalswamy et al. (2009), there are two possibilities during the CME\u2013CH interaction. When the magnetic fields in the CME and the CH are parallel, deflection of the CME takes place. When the CME field lines are antiparallel to the CH field lines, magnetic reconnection is possible. In our event, the jet\u2013facula collision is similar to the first situation of CME\u2013CH interaction. The second reason is the absence of reconnection signatures. The jet\u2013facula collision does not produce any impulsive brightenings in the collision region in the H\u2009\u03b1 and EUV images. Although RFJ was heated after the collision, the heating of RFJ is less intense and concentrated than the reconnection heating, and is probably caused by collision-induced conversion of kinetic energy to thermal energy. Due to the interaction of the primary FJs with the facula structure, the material of primary FJs is redistributed, with partial material propagating along the QSL fan plane and forming the AFJs. It is the apparent material end that differentiates the primary FJs and AFJs observationally. In the aspect of magnetic fields, the primary FJs and AFJs are along the same magnetic field lines from the sunspot. Direct observations of the jet\u2013facula collision process and apparent fan-shaped jets are very rare. Thus it is difficult to give an exact physical explanation based on the present observations. If more related cases could be observed in future, we expect to be able to interpret in depth the physical mechanism of collision-induced apparent jets.","Citation Text":["Galsgaard et al. 2007"],"Functions Text":["When the magnetic field lines are parallel it is not easy to have reconnection"],"Functions Label":["Uses"],"Citation Start End":[[535,556]],"Functions Start End":[[432,510]]}
{"Identifier":"2022ApJ...933..243F__Woosley_&_Bloom_2006b_Instance_2","Paragraph":"Gamma-ray bursts (GRBs) are among the most powerful gamma-ray sources in the universe. They could be generated from the merger of binary compact objects (BCOs; Duncan & Thompson 1992; Usov 1992; Thompson 1994; Metzger et al. 2011) or the death of massive stars (Woosley 1993; Paczy\u0144ski 1998; Woosley & Bloom 2006a). The merger of BCOs; a black hole (BH)\u2013a neutron star (NS) or NS\u2013NS, leading to kilonovae (KNe), is correlated with short-duration gamma-ray bursts (sGRBs; T\n90\n\n10\n\n\n10\n\nT\n90 is defined as the time during which the cumulative number of collected counts above background rises from 5% to 95%. \u2272 2 s; Li & Paczy\u0144ski 1998; Rosswog 2005; Metzger et al. 2010; Kasen et al. 2013; Metzger 2017). On the other hand, long-duration gamma-ray bursts (lGRBs; T\n90 \u2273 2 s; Kouveliotou et al. 1993) are associated with the core collapse (CC) of dying massive stars (Woosley 1993; Galama et al. 1998) leading to supernovae (SNe; Bloom et al. 1999; Woosley & Bloom 2006b). It is believed that in both scenarios large quantities of materials with a wide range of velocities are ejected. In the framework of CC-SNe (depending on the type of SN association), several materials ejected with sub-relativistic velocities less than \u03b2 \u2272 0.4\n11\n\n\n11\nHereafter, we adopt natural units c = \u210f = 1. have been reported (see, e.g., Kulkarni et al. 1998; Bloom et al. 1999; Woosley & Bloom 2006b; Valenti et al. 2008; Gal-Yam 2017; Izzo et al. 2019, 2020; Modjaz et al. 2020; Nicholl et al. 2020). Regarding the merger of two NSs, sub-relativistic materials such as the cocoon, the shock breakout, and the dynamical and wind ejecta are launched with velocities in the range 0.03 \u2272 \u03b2 \u2272 0.8\n12\n\n\n12\nSome authors have considered the shock breakout material in the sub-, trans-, and ultra-relativistic regimes (see, e.g., Kyutoku et al. 2014; Metzger et al. 2015; Fraija et al. 2019c). (see, e.g., Dessart et al. 2009; Metzger & Fern\u00e1ndez 2014; Fern\u00e1ndez et al. 2015; Kyutoku et al. 2014; Metzger et al. 2015; Nagakura et al. 2014; Murguia-Berthier et al. 2014; Lazzati et al. 2017, 2018; Goriely et al. 2011; Hotokezaka et al. 2013; Bauswein et al. 2013; Wanajo et al. 2014). While the mass and velocity inferred for the first GRB\/KN association\n13\n\n\n13\nGRB 170817A\/AT 2017gfo. were M\nej \u2248 (10\u22124\u221210\u22122)M\n\u2299 and \u03b2 \u2248 (0.1\u22120.3), respectively (Coulter et al. 2017; Arcavi et al. 2017; Cowperthwaite et al. 2017; Nicholl et al. 2017; Metzger 2019), the mass and velocity inferred for the first GRB\/SN association\n14\n\n\n14\nGRB 980425\/SN1998bw. was M\nej \u2248 10\u22125\nM\n\u2299 and \u03b2 \u2248 (0.2\u20130.3), respectively (Kulkarni et al. 1998).","Citation Text":["Woosley & Bloom 2006b"],"Functions Text":["Hereafter, we adopt natural units c = \u210f = 1. have been reported (see, e.g.,"],"Functions Label":["Uses"],"Citation Start End":[[1357,1378]],"Functions Start End":[[1240,1315]]}
{"Identifier":"2016ApJ...831..200W__Lor\u00e9n-Aguilar_&_Bate_2015_Instance_1","Paragraph":"HD 97048 is yet another disk for which the dust emission at (sub-)mm wavelengths shows evidence of axisymmetric ring-like structures, here on spatial scales of around tens of au (Walsh et al. 2014; ALMA Partnership et al. 2015; Andrews et al. 2016; Nomura et al. 2016; Zhang et al. 2016). We predict that this substructure will be clearly evident in images of HD 97048 at higher spatial resolution (\u224810\u201320 au, see Figure 13). There remains much debate in the literature on the origin of such axisymmetric substructure in protoplanetary disks including gaps and dust traps carved by forming planets (see, e.g., Dipierro et al. 2015; Pinilla et al. 2015; Rosotti et al. 2016), a change in dust opacity properties at the positions of snow lines (e.g., Banzatti et al. 2015; Zhang et al. 2015; Guidi et al. 2016; Okuzumi et al. 2016), and toroidal dust traps created by hydrodynamic or magnetohydrodynamic effects (see, e.g., Pinilla et al. 2012b; Lor\u00e9n-Aguilar & Bate 2015; Ruge et al. 2016). To distinguish between each of the scenarios requires observations of dust emission at multiple and well-separated frequencies to determine the radial dust size and density distribution (and dust opacity index) along with emission from optically thin gas tracers to determine the gas surface density. Planets will create deep gaps in the gas surface density as well as influencing the dust (note that this is dependent on the planet mass, see e.g., Rosotti et al. 2016), toroidal instabilities will create much shallower features in the gas surface density, and opacity changes at snow lines will affect only the dust emission and will have no effect on the gas. We note that the ringed substructure seen here has very recently been confirmed in scattered light images of HD 97048 taken with VLT\/SPHERE (Ginski et al. 2016). An initial (and shallow) comparison of the data sets shows remarkable coincidence between the positions of the (sub-)mm peaks and gaps and those seen in scattered light. That such structure is seen in both small (\u2248\u03bcm-sized) dust grains in the disk atmosphere and large (\u2248mm-sized) dust grains in the disk midplane points toward a (proto)planetary system origin; however, further data, particularly to better constrain the gas structure, are needed for confirmation. Since this paper has been accepted for publication, ALMA Cycle 2 data of HD 97048, for which longer baseline data were available and imaged with a uv clip (>160 k\u03bb), resulted in a beam of 048 \u00d7 026 (18\u00b0) and resolved the inner dust cavity (40 \u2013 46 au) and the bright dust ring at \u2248150 au (van der Plas et al. 2016).","Citation Text":["Lor\u00e9n-Aguilar & Bate 2015"],"Functions Text":["There remains much debate in the literature on the origin of such axisymmetric substructure in protoplanetary disks including","and toroidal dust traps created by hydrodynamic or magnetohydrodynamic effects (see, e.g.,"],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[944,969]],"Functions Start End":[[426,551],[831,921]]}
{"Identifier":"2022AandA...665L...1F__Peck_et_al._2001_Instance_1","Paragraph":"In our model, we compute the tidal response of the oceans and the solid-Earth to luni-solar semi-diurnal forcing, both combined with mimetic continental drift driven by plate tectonics. We focus on the dependence of dissipation on the Earth\u2019s spin rate. We combine two analytical approaches that describe long-wavelength barotropic tidal flows over shallow spherical and hemispherical shells. The spherical shell describes a global ocean that we assume had existed in the earliest eons of the lifetime of the Earth (Motoyama et al. 2020). The existence of an early ocean is supported by evidence from the analysis of detrital zircon around 4.4 Ga (Wilde et al. 2001), from the interaction between the ocean and continental crust 4 billion years ago (Mojzsis et al. 1996), and from records of the oxygen isotope composition of seawater (Peck et al. 2001; Johnson & Wing 2020). The \u201cglobality\u201d of this ocean is justified by the analysis of continental crust growth curves based on geochemical evidence in zircon crystallization ages (Dhuime et al. 2012; Hawkesworth et al. 2020). In compliance with these curves, we consider that a hemispherical oceanic shell has taken over in the most recent times. In our model, the center of this hemispheric continental cap follows the evolution of the paleogeographic center. In doing so, we emphasize on the role of \u201ccontinentality\u201d in the tidal response, while avoiding the under-sampling of geometric scenarios due to theoretical limitations (Hansen 1982; Tyler 2021) or due to uncertainties in plate tectonic models (Matthews et al. 2016; Daher et al. 2021). To compute this evolution, we adopt the recently developed paleogeographic reconstructions that cover the past billion years (Merdith et al. 2021). A postprocessing of these reconstructions allows us to produce the latitudinal evolution of the center of the continental cap captured in Fig. 1. The tidal frequencies at which oceanic resonances are excited and the amplitudes of these resonances vary with the surface position of the hemispherical ocean (Fig. B.1). Super-continental formations and breakups thus have their mark on the predicted lunar recession rate.","Citation Text":["Peck et al. 2001"],"Functions Text":["The existence of an early ocean is supported by","and from records of the oxygen isotope composition of seawater"],"Functions Label":["Background","Background"],"Citation Start End":[[836,852]],"Functions Start End":[[539,586],[772,834]]}
{"Identifier":"2015AandA...574A..62S__Burkepile_et_al._(2004)_Instance_1","Paragraph":"Quiescent prominences are objects formed by relatively cool plasma in the hot corona. Their cool material occurs mostly in the dipped magnetic field lines. In quiescent prominences, which can persist from several hours to several days, magnetic dips form quasi-vertical structures called threads. Quiescent prominences are often observed as part of magnetic structures composed of three coronal patterns: the prominence itself, surrounded by a low-density cavity and a dense helmet streamer overlying the cavity (Engvold 1989). The disruption of the helmet streamer often signifies the beginning of a coronal mass ejection (CME) that can reflect the three-part structure of the helmet streamer (when observed in the white light): the bright leading shell surrounding a dark cavity in which bright prominence material occurs (Crifo et al. 1983; Hundhausen 1999). As the masses of CMEs are most often inferred from the visible-light observations, the masses given in literature are usually only the masses of the leading shell. However, Burkepile et al. (2004) also saw highly structured material most probably originating from a prominence in 63% of CMEs associated with eruptive prominences at the limb observed in visible light using the broadband filter (\u03bb\u03bb 5000\u22125350\u2009\u00c5). Although such observations are called as white light in literature, it would be better to consider them as observations in visible light, because naturally integral intensities in the visible spectral range are observed including both continuum and absorption lines, while white light is just a continuum caused by Thomson scattering on free electrons without the lines. Thus, the presence of absorption spectral lines in visible light complicates very much the estimation of prominence mass, but even if the influence of absorption lines was eliminated, continuum alone only allows us to estimate the mass of ionised material, thus mass of the prominence would be underestimated. Low (1996) stated that the contribution of an erupting prominence to the total mass of CME is usually one order of magnitude less than that of the shell, but in some cases these contributions can be comparable. Low et al. (2003) proposed theoretically an importance of the prominence mass for the deposit of magnetic energy for driving a CME. Thus, developing methods for the estimation of the total mass of prominences can help to explain the connection between prominences and CMEs and provide an important quantity for more accurate estimation of total mass ejected by a CME. ","Citation Text":["Burkepile et al. (2004)"],"Functions Text":["However,","also saw highly structured material most probably originating from a prominence in 63% of CMEs associated with eruptive prominences at the limb observed in visible light using the broadband filter (\u03bb\u03bb 5000\u22125350\u2009\u00c5).","Although such observations are called as white light in literature, it would be better to consider them as observations in visible light, because naturally integral intensities in the visible spectral range are observed including both continuum and absorption lines, while white light is just a continuum caused by Thomson scattering on free electrons without the lines."],"Functions Label":["Background","Background","Compare\/Contrast"],"Citation Start End":[[1035,1058]],"Functions Start End":[[1026,1034],[1059,1273],[1274,1644]]}
{"Identifier":"2016MNRAS.463..696L__Randich_et_al._2006_Instance_1","Paragraph":"Another possible explanation for the peculiar solar composition is that some of the dust in the pre-solar nebula was radiatively cleansed by luminous hot stars in the solar neighbourhood before the formation of the Sun and its planets. This dust-cleansing scenario is supported by the finding that the open cluster M67 seems to have a chemical composition closer to the solar composition than most solar twins (\u00d6nehag et al. 2011, hereafter O11 and \u00d6nehag, Gustafsson & Korn 2014, hereafter O14). They suggested that the proto-solar nebula was dust-cleansed by massive stars, similar to what happened for the proto-cluster cloud of M67, while the majority of solar twins in the field would presumably have formed in less massive clusters where no nearby high-mass star (\u226515 M\u2299) was formed. M67 offers the possibility for studying the solar-type stars in a dense cluster environment. This cluster has about solar metallicity ([Fe\/H] in the range \u22120.04 to +0.03, e.g. Hobbs & Thorburn 1991; Yong, Carney & Teixera de Almeida 2005; Randich et al. 2006; Pasquini et al. 2008). The age of M67 is also comparable with that of the Sun: 3.5\u20134.8 Gyr (Yadav et al. 2008). Pasquini et al. (2008) listed 10 solar-twin candidates in M67, including M67-1194 which hosts a hot Jupiter with a period of 6.9 d and a minimum mass of 0.34 MJup (Brucalassi et al. 2014). M67-1194 was studied by O11 and re-visited by O14. They found that unlike nearby solar twins, which were systematically analysed by Mel\u00e9ndez et al. (2009) and Ram\u00edrez, Mel\u00e9ndez & Asplund (2009), the chemical composition of M67-1194 is more solar-like and therefore suggested that the Sun may have been formed in a similar cluster environment, perhaps even in M67. While this scenario is plausible when considering chemical abundances and ages, the dynamics are problematic: the Sun has an orbit close to the Galactic plane, while M67 is presently some 450 pc above the plane. The probability that the Sun, if formed in M67 with an orbit similar to the present cluster orbit, was scattered or diffused out of the cluster into the present solar orbit, is found to be quite low, if the existence of the outer planetary system is taken into account (Pichardo et al. 2012). Gustafsson et al. (2016) suggested that the high-altitude metal-rich clusters (such as M67) were formed in orbits close to the Galactic plane and later scattered to higher orbits by interaction with giant molecular clouds and spiral arms. Thus, it is possible, though not very probable, that the Sun formed in such a cluster before scattering occurred. Currently, only one solar twin in M67 (M67-1194) was spectroscopically analysed with high precision (\u223c0.02 dex). It is thus crucial to analyse the chemical composition of additional solar-type stars in M67 of very high precision.","Citation Text":["Randich et al. 2006"],"Functions Text":["This cluster has about solar metallicity ([Fe\/H] in the range \u22120.04 to +0.03, e.g."],"Functions Label":["Motivation"],"Citation Start End":[[1029,1048]],"Functions Start End":[[883,965]]}
{"Identifier":"2016MNRAS.455..449H__Oman_et_al._2015_Instance_1","Paragraph":"With only six free parameters, the standard \u039b cold dark matter (\u039bCDM) cosmological model fits no less than 2500 multipoles in the cosmic microwave background (CMB) angular power spectrum (Planck Collaboration XVI 2014), the Hubble diagram of Type Ia supernovae, the large-scale structure matter power spectrum, and even the detailed scale of baryonic acoustic oscillations. It thus provides the current basis for simulations of structure formation, and is extremely successful down to the scale of galaxy clusters and groups. Nevertheless, it still faces numerous challenges on galaxy scales. Among these, the most important ones are the too-big-to-fail problem (Boylan-Kolchin, Bullock & Kaplinghat 2011) and the satellite-plane problem (e.g. Pawlowski, Pflamm-Altenburg & Kroupa 2012; Ibata et al. 2014) for dwarf galaxies, the tightness of the baryonic Tully\u2013Fisher relation (McGaugh 2012; Vogelsberger et al. 2014), or the unexpected diversity of rotation curve shapes at a given mass scale (Oman et al. 2015). The latter problem is actually a subset of a more general problem, i.e. that the shapes of rotation curves indeed do not depend on the Dark Matter (DM) halo mass, contrary to what would be expected in \u039bCDM, but rather on the baryonic surface density, as has long been noted (e.g. Zwaan et al. 1995). This makes the problem even worse, since the rotation curve shapes are not only diverse at a given mass scale, but uniform at a given baryonic surface density scale, implying a completely ununderstood fine-tuning of putative feedback mechanisms. On the other hand, this behaviour of rotation curves is an a priori prediction of the formula proposed by Milgrom more than 30 yr ago (Milgrom 1983a,b), relating the total gravitational field to the Newtonian field generated by baryons alone, and which can be interpreted as a modification of Newtonian dynamics on galaxy scales below a characteristic acceleration (Modified Newtonian Dynamics (MOND), for a review see Famaey & McGaugh 2012; Milgrom 2014). With this simple formula, high surface brightness (HSB) galaxies are predicted to have rotation curves that rise steeply before becoming essentially flat, or even falling somewhat to the not-yet-reached asymptotic circular velocity, while low surface brightness (LSB) galaxies are predicted to have rotation curves that rise slowly to the asymptotic velocity. This is precisely what is observed, and was predicted by Milgrom long before LSB galaxies were even known to exist. The formula also predicts the tightness of the baryonic Tully\u2013Fisher relation.","Citation Text":["Oman et al. 2015"],"Functions Text":["Nevertheless, it still faces numerous challenges on galaxy scales. Among these, the most important ones are","or the unexpected diversity of rotation curve shapes at a given mass scale"],"Functions Label":["Background","Background"],"Citation Start End":[[996,1012]],"Functions Start End":[[526,633],[920,994]]}
{"Identifier":"2018ApJ...864...49P__Tacconi_et_al._2013_Instance_1","Paragraph":"Figure 9 summarizes our constraints on the molecular gas mass fraction in the analyzed samples of galaxies. In order to convert CO luminosity to molecular gas mass, we consider two different assumptions for \u03b1CO. First, we assume a constant value of \u03b1CO = 3.6 M\u2299 (K km s\u22121 pc2)\u22121 adopted by some previous studies (e.g., Daddi et al. 2010b; Decarli et al. 2014; Walter et al. 2016). We then also consider a metallicity-dependent conversion factor, evaluated by assuming a redshift and stellar mass\u2013metallicity relation (Genzel et al. 2015). Previous studies investigating optically and FIR-selected galaxy samples have estimated the relationship between gas mass fraction, stellar mass, redshift, and SFR offset from the main sequence (e.g., Genzel et al. 2015; Scoville et al. 2017). While the PHIBSS project estimated molecular gas masses by measuring the CO(3\u20132) line emission (Tacconi et al. 2013; Genzel et al. 2015), Scoville et al. (2016, 2017) used the flux on the Rayleigh\u2013Jeans tail of the dust continuum emission to estimate the total gas masses. We here assume that the samples of galaxies plotted, although not complete to any degree due to their preselection for having a spectroscopic redshift, may be somewhat representative of the star-forming main sequence (Figure 9). Their star formation rates are consistent with scatter around the main sequence and appear to include as many galaxies above and below the main sequence estimated by Speagle et al. (2014). Although our CO detections (both blind and with previous spectroscopic redshifts) are indicative of gas fractions compatible with or above (GN19) expectations for main-sequence galaxies, the individual CO nondetections and the stacked signal appear to be systematically lower than the predicted averages, suggesting lower gas mass fractions than might be expected (Figure 9). We can quantify the apparent deficit in stacked signal relative to expectations for the \n\n\n\n\n\n sample by calculating expected gas masses for the individual stacked galaxies predicted as a function of their redshift, stellar masses, and star formation rates. The expected sample average molecular gas mass is \n\n\n\n\n\n adopting the best-fit relation by Genzel et al. (2015) and \n\n\n\n\n\n according to the relation by Scoville et al. (2017). The constant CO luminosity conversion factor above would imply a ratio between expected and observed stacked CO luminosity of 4.8 \u00b1 2.4 and 6.3 \u00b1 3.1, according to the relations by Genzel et al. (2015) and Scoville et al. (2017), respectively. Applying instead the metallicity-dependent CO conversion factor suggested by Genzel et al. (2015) to individual galaxies would somewhat reduce the tension, implying ratios of 3.0 \u00b1 1.7 and 3.8 \u00b1 2.1, according to the relations by Genzel et al. (2015) and Scoville et al. (2017), respectively. While the constraints for low stellar mass galaxies may be compatible with an evolving CO conversion factor due to low metallicity, this is unlikely to resolve the apparent conflict at the high-mass end and may point to lower-than-expected gas masses.","Citation Text":["Tacconi et al. 2013"],"Functions Text":["While the PHIBSS project estimated molecular gas masses by measuring the CO(3\u20132) line emission","Scoville et al. (2016, 2017) used the flux on the Rayleigh\u2013Jeans tail of the dust continuum emission to estimate the total gas masses."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[879,898]],"Functions Start End":[[783,877],[921,1055]]}
{"Identifier":"2021ApJ...910...86R__Stark_et_al._2017_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":["Stark et al. 2017"],"Functions Text":["Complementing these observations, the spectroscopic confirmation","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."],"Functions Label":["Background","Background"],"Citation Start End":[[1131,1148]],"Functions Start End":[[948,1012],[1264,1484]]}
{"Identifier":"2018MNRAS.481.3573L__Robotham_et_al._2017_Instance_1","Paragraph":"Another important test for galaxy formation models, is whether they place the right amount of stellar mass into discs and bulges. Moffett et al. (2016) measured the SMF separating galaxies into different morphological types and also into discs\/bulges. With this, they derived the fractional contribution from bulges\/discs to the total stellar mass in bins of stellar mass. This is quite a difficult measurement to do in observations as light profile fitting is required, which can be robustly done in very disc- and bulge-dominated galaxies, but when both components contribute similarly, the measurement is less robust (Robotham et al. 2017). In Fig. 11, we compare Shark with these measurements for two measurements of bulge mass.9 The first one is considering all the mass in the central concentration (regardless of whether it was formed due to mergers or disc instabilities; solid line), and the second one assumes that the bulge mass formed through disc instabilities (either through the starburst triggered by the gas being fueled to the centre or the stars that are transferred from the disc to the bulge) is part of the disc (dashed line). The bulge mass formed via mergers include both the stars formed via merger-driven starbursts and stars that were accreted by the bulge as a result of a merger (but that were formed in the disc of the primary and\/or in the secondary galaxies). The latter is done as pseudobulges in Moffett et al. (2016) were added up to the disc rather than the bulge, and those are thought to form through secular processes taking place in the discs of galaxies (Kormendy & Kennicutt 2004). The effect of assigning the bulge mass formed via disc instabilities to the disc has the effect of shifting the transition from disc- to bulge-dominated stellar budget to higher stellar masses, much closer to the Moffett et al. (2016) observations. In Shark, we find that the formation of elliptical galaxies (i.e. spheroid dominated, massive galaxies) is dominated by galaxy mergers rather than disc instabilities, as disc instabilities only increase the bulge contribution by \u224810 per cent at $10^{11}\\, \\rm M_{\\odot }$. This is qualitatively similar to the finding in large cosmological hydrodynamical simulations that elliptical galaxies form primarily via galaxy mergers (Wellons et al. 2016; Clauwens et al. 2018; Lagos et al. 2018a).","Citation Text":["Robotham et al. 2017"],"Functions Text":["This is quite a difficult measurement to do in observations as light profile fitting is required, which can be robustly done in very disc- and bulge-dominated galaxies, but when both components contribute similarly, the measurement is less robust"],"Functions Label":["Motivation"],"Citation Start End":[[621,641]],"Functions Start End":[[373,619]]}
{"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)"],"Functions Text":["Set II: we now focus on the observational constraints on the model parameters after the inclusion of the local measurement of H0 by","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."],"Functions Label":["Uses","Uses"],"Citation Start End":[[133,152]],"Functions Start End":[[1,132],[153,296]]}
{"Identifier":"2021ApJ...907...47L__Lee_et_al._2019_Instance_1","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"],"Functions Text":["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;"],"Functions Label":["Similarities"],"Citation Start End":[[889,904]],"Functions Start End":[[576,872]]}
{"Identifier":"2018AandA...616A..11G__Westerhout_1957_Instance_1","Paragraph":"Gaia DR2 contains unprecedented information about the Galaxy, which should allow us to infer its current structure, its equilibrium state, its evolution, modes of mass growth over time, dark matter distribution (and perhaps nature), to cite a few of the questions of modern Galactic astrophysics. As an example, it has been known for several decades that the Galactic disc contains large-scale non-axisymmetric features, including a central boxy\/peanut-shaped bar (Okuda et al. 1977; Maihara et al. 1978; Weiland et al. 1994; Dwek et al. 1995; Binney et al. 1997; Babusiaux & Gilmore 2005; L\u00f3pez-Corredoira et al. 2005; Rattenbury et al. 2007; Cao et al. 2013) and its possible in-plane extension (Hammersley et al. 2000; Benjamin et al. 2005; Cabrera-Lavers et al. 2007; Wegg et al. 2015), a warp (Burke 1957; Kerr 1957; Westerhout 1957; Weaver 1974; Djorgovski & Sosin 1989; Evans et al. 1998; Gyuk et al. 1999; Drimmel & Spergel 2001; L\u00f3pez-Corredoira et al. 2002; Momany et al. 2006; Robin et al. 2008; Reyl\u00e9 et al. 2009; Am\u00f4res et al. 2017), and spiral arms (Georgelin & Georgelin 1976; Taylor & Cordes 1993; Drimmel 2000; Bissantz & Gerhard 2002; Churchwell et al. 2009; Vall\u00e9e 2014; Reid et al. 2014; Hachisuka et al. 2015; Hou & Han 2015). However, full knowledge of these asymmetric structures, that is, of their spatial extent, pattern speeds, and number (in case of spiral arms) is still lacking. Since asymmetries constitute the driver of the secular evolution in galaxy discs (see e.g. Minchev et al. 2012; Fouvry et al. 2015; Halle et al. 2015; Aumer et al. 2017 and Kormendy 2013, for a review) by redistributing angular momentum between the inner and outer disc and between its baryonic and dark matter content (Debattista & Sellwood 2000; Bournaud & Combes 2002; Athanassoula 2003; Martinez-Valpuesta et al. 2006; Combes 2011), quantifying their characteristics is fundamental for understanding to what extent the Milky Way has \u201csimply\u201d evolved secularly in the last ~9 Gyr (Hammer et al. 2007; Martig et al. 2014), or whether some more complex evolutionary scenarios need to be invoked.","Citation Text":["Westerhout 1957"],"Functions Text":["As an example, it has been known for several decades that the Galactic disc contains large-scale non-axisymmetric features, including","a warp"],"Functions Label":["Background","Background"],"Citation Start End":[[822,837]],"Functions Start End":[[297,430],[791,797]]}
{"Identifier":"2020ApJ...891...59R__Feruglio_et_al._2015_Instance_1","Paragraph":"In order to understand the origin of the termination of the powerful jet in Mrk 231, consider the very low accretion rates in these FRI NLRGs, three to four orders of magnitude less than a quasar (Chiaberge et al. 1999, 2002; Hardcastle et al. 2009). However, there are the occasional FRI broad-line galaxies such as the Seyfert 1 galaxy, 3C 120, which has a very prominent FRI morphology on a 400 kpc scale (Walker et al. 1987). Thus, the low accretion rate of FRI NLRGs is not the full explanation of the discrepancy with Mrk 231, but is likely related. One difference between 3C 120 and Mrk 231 is that Mrk 231 has a low-ionization BAL wind (Lipari et al. 1994; Smith et al. 1995). In low radio states, it has also displayed evidence of a high-ionization X-ray-absorbing wind (Feruglio et al. 2015; Reynolds et al. 2017). There are also extreme amounts of intrinsic optical absorption in the galaxy itself from dusty gas (Lipari et al. 1994; Smith et al. 1995). All three circumstances point to a very dense nuclear environment through which the jet must propagate. Furthermore, in Reynolds et al. (2009), it was argued that the density of the free\u2013free absorbing screen at K1 is consistent with the jet being stopped by the BAL wind at K1. By contrast, the extremely low, undetectable accretion of gas in FRI NLRGs is consistent with a very low-density nuclear environment. The launching of the jet in Mrk 231 does not seem to be stopped by the dense nuclear environment, but its propagation does seem to be thwarted. The case of 3C 120 does not seem to be accommodated by this discussion. It is noted that there is no evidence of either a BAL wind or a high-ionization X-ray-absorbing wind in 3C 120 (Oke & Zimmerman 1979; Ballentyne et al. 2004). Thus, it might be the case that there is not an extremely dense nuclear environment near the source of the jet in 3C 120; hence, jet propagation is not hindered. We cannot rule out the possibility that the BAL wind in Mrk 231 also diminishes the power of the jet-launching mechanism as well as providing a drag on its propagation.","Citation Text":["Feruglio et al. 2015"],"Functions Text":["In low radio states, it has also displayed evidence of a high-ionization X-ray-absorbing wind"],"Functions Label":["Background"],"Citation Start End":[[780,800]],"Functions Start End":[[685,778]]}
{"Identifier":"2021AandA...647A..35B__Laffon_et_al._(2010)_Instance_1","Paragraph":"When we now compare the photodesorption yields at 541 eV between fluences  5 \u00d7 1015 photon cm\u22122 and at 3 \u00d7 1017 photon cm\u22122 in Fig. 2, the CO2 photodesorption first increases from 2.6 \u00d7 10\u22122 molecule\/photon to 7.3 \u00d7 10\u22122 molecule\/photon. We also observed this phenomenon for CO photodesorption yield (the data are not shown for more clarity), which increased from 1.9 \u00d7 10\u22122 molecule\/photon to 3.5 \u00d7 10\u22122 molecule\/photon. Second, the estimated yield for the X-ray photodesorption of CH3OH from pure methanol ice decreased by almost one order of magnitude from 9.0 \u00d7 10\u22123 molecule\/photon to 1.3 \u00d7 10\u22123 molecule\/photon. This indi- cates that the photodesorption of CH3OH is higher for a lower fluence received by the ice when more intact methanol molecules are present in the ice. This aging process favors the photodesorption of simpler molecules such as CO2 or CO. Laffon et al. (2010) estimated with NEXAFS spectroscopy (at the C K-edge) that X-ray irradiation at 150 eV of pure methanol ice at 20 K leads to a survival rate of 50% for methanol after an absorbed dose of 1.1 MGy. In our fixed-energy experiments, we irradiated pure methanol ice with fluences between 5 \u00d7 1015 photon cm\u22122 and 2 \u00d7 1016 photon cm\u22122. Because we irradiated a volume of 0.1 cm2 \u00d7 100 ML, with a meanenergy of ~550 eV, and when we consider a volumic mass of condensed methanol of ~ 0.64 g cm\u22123 (at 20 K; Luna et al. 2018) and an X-ray absorption cross section of ~ 0.6 Mbarn (Ishii & Hitchcook 1988), the absorbed doses used in our fixed energy experiments change from ~ 2 to ~ 15 MGy, whichis quite similar to the absorbed doses in Laffon et al. (2010). This indicates that we could expect a methanol destruction rate of about 50% for our low-fluence experiments. In similar experiments, when irradiating a H2 O:CH4:NH3 (2:1:1) ice mixture covered by a layer of CO:CH3OH (3:1) with 250\u20131250 eV X-rays during 120 min with a flux of 7.6 \u00d7 1014 photon s\u22121, higher by almost two order of magnitudes than our experiments, Ciaravella et al. (2020) did not detect a desorption signal on the mass channel 31 (attributed to methanol desorption) and estimated that only ~ 20% of methanol molecules remained intact in the first minutes of the irradiation. The irradiation flux therefore appears to be critical for detecting methanol desorption in X-ray irradiation experiments of methanol-containing ices. A lower X-ray flux appears to favor methanol desorption because the methanol destruction rate is lower. This destruction of methanol molecules could also have a significant effect on the formation and desorption of more complex molecules.","Citation Text":["Laffon et al. (2010)"],"Functions Text":["estimated with NEXAFS spectroscopy (at the C K-edge) that X-ray irradiation at 150 eV of pure methanol ice at 20 K leads to a survival rate of 50% for methanol after an absorbed dose of 1.1 MGy."],"Functions Label":["Uses"],"Citation Start End":[[865,885]],"Functions Start End":[[886,1080]]}
{"Identifier":"2016MNRAS.461.1745E__Goerdt_et_al._2010_Instance_1","Paragraph":"Given that the CDM paradigm only begins to face significant problems at precisely such scales when complex baryonic physics begins to play an important role, it is natural to inquire whether it is the central culprit behind erroneous theoretical predictions. It was for example realized early on that energy from supernovae may be sufficient for driving gas out of the potential wells of dwarf galaxies, the associated mass deficit resulting in the expansion of the central halo region and the flattening of the density profile. More generally, many hydrodynamical simulations implementing stellar and active galactic nuclei (AGN) baryonic feedback processes in a cosmological context are able to reproduce cores (e.g. Governato et al. 2010, 2012; Macci\u00f2 et al. 2012b; Martizzi et al. 2012; Di Cintio et al. 2014; Chan et al. 2015). However, the complexity of such simulations obscures the physical mechanisms through which these processes affect the dark matter distribution. These mechanisms normally invoke \u2018heating\u2019 of the cold central density cusp through an irreversible process, such as dynamical friction from infalling clumps (El-Zant, Shlosman & Hoffman 2001; El-Zant et al. 2004; Tonini, Lapi & Salucci 2006; Romano-D\u00edaz et al. 2008; Goerdt et al. 2010; Cole, Dehnen & Wilkinson 2011; Del Popolo et al. 2014; Nipoti & Binney 2015). Alternatively, repeated gravitational potential fluctuations induced by stellar winds, supernova explosions and AGN could also dynamically heat the central halo (Read & Gilmore 2005; Mashchenko, Couchman & Wadsley 2006; Mashchenko, Wadsley & Couchman 2008; Peirani, Kay & Silk 2008; Governato et al. 2012; Pontzen & Governato 2012, 2014; Zolotov et al. 2012; Martizzi, Teyssier & Moore 2013; Teyssier et al. 2013; Madau, Shen & Governato 2014; Ogiya & Mori 2014). Although the last mechanism may seem most closely related to the supernovae driven wind outflows discussed above, it is in principle more closely connected to the dynamical friction proposal, in the sense that it involves irreversible stochastic dynamics: one may envisage the potential fluctuations leading to cusp-core transformation as originating from stochastic density variations; the relevant \u2018clumps\u2019 would be associated with fluctuation scales, as opposed to physically distinct objects dissipating orbital energy via dynamical friction; nevertheless, the basic physical mechanism through which the energy is transferred to the dark matter is similar. For, as is the case in general with processes involving fluctuation and dissipation, fluctuations in a gravitational system can be approximated as stochastic processes described by power spectra and correlation functions, and they can be accompanied by dissipation in the form of dynamical friction (Chandrasekhar 1943; Nelson & Tremaine 1999).","Citation Text":["Goerdt et al. 2010"],"Functions Text":["However, the complexity of such simulations obscures the physical mechanisms through which these processes affect the dark matter distribution. These mechanisms normally invoke \u2018heating\u2019 of the cold central density cusp through an irreversible process, such as dynamical friction from infalling clumps"],"Functions Label":["Background"],"Citation Start End":[[1245,1263]],"Functions Start End":[[833,1134]]}
{"Identifier":"2019ApJ...876L..28D__Lamb_et_al._2018_Instance_2","Paragraph":"In Figures 1 and 2 we show that the X-ray (1.7 keV5\n\n5\nThis value corresponds to the geometric mean of the XRT energy band, at which the error of the estimated flux can be reasonably suppressed.\n), optical (R-band), and radio (6 GHz) fluxes varied with the time of observation applied to the proper corrections if observed at a distance of 200 Mpc, motivated by the fact that the averaged sensitive range of the Advanced LIGO\/Virgo detectors in their full-sensitivity run is about 210 Mpc, for the current samples. Due to the faintness of the SGRB afterglow emission, there are gaps of the data between the previous more distant events and GW170817\/GRB 170817A (please note that for the latter we only consider the quick decline phase as the early part is significantly influenced by the beam effect of the off-axis outflow). Therefore we extrapolate the very late (t > 200 day) X-ray and optical afterglow data of GRB 170817A to t \u223c 2 day after the burst and then compare them to other events. The radio to X-ray spectrum of the forward shock afterglow emission of GW170817\/GRB 170817A is f\u03bd \u221d \u03bd\u22120.6, which yields a p = 2.2 in the slow-cooling synchrotron radiation scenario (Lamb et al. 2018; Troja et al. 2018). In the jet model, such a p can also reasonably account for the very late flux decline of \n\n\n\n\n\n (Lamb et al. 2018). The extrapolation function of the forward shock emission of GRB 170817A to early times is thus taken as f \u221d t\u22122.2. Surprisingly, the forward shock afterglow emission of GW170817\/GRB 170817A, the first neutron star merger event detected by Advanced LIGO\/Virgo, are among the brightest ones for all SGRBs detected so far. Just a few events have X-ray afterglow emission brighter than that of GRB 170817A, as demonstrated in the right panels of Figure 1. The same conclusion holds for the optical and radio afterglow data, as shown in Figure 2, though these two samples are rather limited. We have also compared the distribution of the isotropic gamma-ray energy Eiso, calculated in the rest-frame energy band of 1\u2013104 keV, for the SGRBs with well-measured spectra, and found no significant difference for the SGRBs with and without \u201clong-lasting\u201d afterglow emission (see Figure 3; where the number of events for the X-ray sample are smaller than that presented in Figure 1 because some bursts lack reliable spectral measurements). GRB170817A and GRB 150101B (Troja et al. 2018), two short events with the weakest detected prompt emission, have \u201cbright\u201d late-time afterglow emission because of their off-axis nature.","Citation Text":["Lamb et al. 2018"],"Functions Text":["In the jet model, such a p can also reasonably account for the very late flux decline of"],"Functions Label":["Uses"],"Citation Start End":[[1312,1328]],"Functions Start End":[[1215,1303]]}
{"Identifier":"2017MNRAS.464.3385W__Tsiganis_et_al._2005_Instance_1","Paragraph":"Oort Cloud. The conclusions that can be reached from this figure about the formation of the Oort Cloud are well known. For example, T93 showed that the parameter space in which an Oort Cloud forms is quite restricted, and that those Oort Clouds that do form have a narrow range of semimajor axes \u223c10\u2009000 au. While this parameter space is inhabited by Uranus and Neptune in the Solar system, which should thus readily supply objects to the Oort Cloud (even accounting for the possibility that these planets may have started closer to the Sun; Tsiganis et al. 2005), it could be that Oort Clouds are relatively rare. Many simulations have confirmed these predictions regarding the ability of planets to implant objects in the Oort Cloud (e.g. Dones et al. 2004) while also showing further subtleties such as the ability of Jupiter and Saturn to place a small fraction of the objects they scatter into the Oort Cloud even if ejection is the predominant outcome in such encounters (Brasser, Duncan & Levison 2008). The time-scale predicted for the scattering process to occur is also borne out in numerical simulations. For example, compare the prediction of Fig. 1 that it should take 0.1\u20131 Gyr for Uranus and Neptune to implant material in the Oort Cloud with fig. 13 of Dones et al. (2015). The radius at which the Oort Cloud forms in the simulations also agrees with that predicted of \u223c10\u2009000 au, with some studies includingthe differentiation between inner and outer Oort Clouds (e.g. Lewis, Quinn & Kaib 2013; Brasser & Schwamb 2015). Inspection of Fig. 1 shows that Oort Clouds could form at smaller orbital radii, but that such an outcome requires both low-mass planets and a large system age; for example, if Neptune and Uranus were each of Earth mass, then the Oort Cloud would be at \u223c1000 au, but would take \u223c20 Gyr to form. Another way to achieve a small Oort Cloud on a shorter time-scale is to place the planetary system in a dense stellar environment (as discussed in Section 3.2).","Citation Text":["Tsiganis et al. 2005"],"Functions Text":["While this parameter space is inhabited by Uranus and Neptune in the Solar system, which should thus readily supply objects to the Oort Cloud (even accounting for the possibility that these planets may have started closer to the Sun;","it could be that Oort Clouds are relatively rare."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[542,562]],"Functions Start End":[[308,541],[565,614]]}
{"Identifier":"2022ApJ...926..206Z__Schwab_et_al._2016_Instance_1","Paragraph":"The first possible scenario is that all FRBs are generated from magnetars, the population of non-repeating FRBs are from magnetars born in the delay formation channels (Margalit et al. 2019) such as neutron star\u2212neutron star mergers associated with short gamma-ray bursts (Rosswog et al. 2003; Price & Rosswog 2006; Giacomazzo & Perna 2013), neutron star\u2212white dwarf mergers possibly associated with rapidly evolving supernovae (SNe) (Toonen et al. 2018; Zhong & Dai 2020), or white dwarf\u2212white dwarf mergers (Yoon et al.2007; Schwab et al. 2016) possibly associated with SNe Ia, and accretion-induced collapse of white dwarfs (Nomoto & Kondo 1991; Tauris et al. 2013; Schwab et al. 2015); the population of repeating FRBs are from magnetars born in the prompt formation channels such as type I superluminous SNe, long gamma-ray bursts, or core-collapse SNe (Duncan & Thompson 1992; Kouveliotou et al. 1998; Kasen & Bildsten 2010). Nonetheless, it seems to be untrue based on recent observations in host galaxies and offsets of the repeaters FRBs 20180916B (Marcote et al. 2020; Tendulkar et al. 2021) and 20200120E (Bhardwaj et al. 2021a; Kirsten et al. 2021). So it may be that the observed non-repeating and repeating FRBs are respectively from less active and more active magnetars due to different magnetic field strengths. If this is the case, however, it is unclear whether the distribution discrepancies of properties between the one-off events and repeaters can arise from different magnetic field strengths. Moreover, if the one-off events and repeaters are produced by an identical mechanism like that generating FRB 20200428 and its associated X-ray burst, it is conceivable that the more active magnetars yielding repeaters could produce more frequent more violent outbreaks resulting in brighter bursts than the less active magnetars yielding one-off events. This appears to be inconsistent with the statistical results shown in Figure 3, though it is still a complicated issue.","Citation Text":["Schwab et al. 2016"],"Functions Text":["The first possible scenario is that all FRBs are generated from magnetars, the population of non-repeating FRBs are from magnetars born in the delay formation channels","or white dwarf\u2212white dwarf mergers"],"Functions Label":["Background","Background"],"Citation Start End":[[527,545]],"Functions Start End":[[0,167],[474,508]]}
{"Identifier":"2018AandA...613A..35K__Leloudas_et_al._2011_Instance_1","Paragraph":"As shown in Fig. 3, the differences in metallicity between different SN subclasses are not significant. This is in contradiction with what is expected from single-star evolution theory, where metallicity-driven winds are crucial: type Ic SNe, which are the most highly stripped, would show the highest metallicity, followed by type Ib, and finally the H-rich type II SNe. The observations, on the other hand, reveals that this is not the case. Some SNe Ic are even located in the low-metallicity part of the distribution in the current sample. This result strengthens the notion that metallicity may not play an important role in deciding theresulting SN type, in agreement with other works based on SN environments (Anderson et al. 2010, 2015; Leloudas et al. 2011; Galbany et al. 2016a). The environments of broad-lined SNe IcBL are found to be relatively metal poor compared to the normal CCSNe, in agreement with previous studies (Modjaz et al. 2011; Galbany et al. 2016a). However, we note that there are only two such SNe in the current sample. The explosion site of SN 1998bw (the first SN to be associated with a GRB: 980425; Galama et al. 1998; Kr\u00fchler et al. 2017) in this study shows a lower metallicity of 12 + log(O\/H) = 8.30 dex compared to the GRB-less SN 2009bb (Pignata et al. 2011), 12 + log(O\/H) = 8.49 dex. Levesque et al. (2010a), using slit spectroscopy of the explosion site, concluded that the high metallicity of SN 2009bb site is consistent with typical GRB-less SNe IcBL and not with GRB hosts. Their metallicity value recalculated on the Marino et al. (2013) N2 scale is 12 + log(O\/H) = 8.52 dex. These two different cases illustrate the importance of metallicity in deciding whether an SN IcBL progenitor would also produce GRB or not (Modjaz et al. 2008; Levesque et al. 2010b). Progenitors with higher metallicity are not able to spin fast enough and thus produce high angular momentum essential for GRB jet production, eventually producing a GRB-less SN IcBL (Woosley & Bloom 2006).","Citation Text":["Leloudas et al. 2011"],"Functions Text":["This result strengthens the notion that metallicity may not play an important role in deciding theresulting SN type, in agreement with other works based on SN environments"],"Functions Label":["Similarities"],"Citation Start End":[[745,765]],"Functions Start End":[[544,715]]}
{"Identifier":"2022MNRAS.517.5744G__Caro_et_al._2016_Instance_3","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"],"Functions Text":["Absorption band shifts of CO ice in the UV and IR ranges only occurred at deposition temperatures above 20\u2009K","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."],"Functions Label":["Background","Compare\/Contrast"],"Citation Start End":[[912,934]],"Functions Start End":[[783,891],[937,1286]]}
{"Identifier":"2022AandA...658A..35K__Boquien_et_al._2019_Instance_1","Paragraph":"It is important in our analysis to only keep sources for which the SFR and stellar mass estimates are reliable. To this end, we first excluded sources whose fit had reduced \u03c72, \n\n\n\n\n\u03c7\n\nred\n2\n\n>5\n\n$ \\chi ^2_{\\rm red}>5 $\n\n\n\n. This criterion is based on visual inspection of the SED fits and has been adopted in previous studies (e.g., Masoura et al. 2018; Mountrichas et al. 2019; Buat et al. 2021). Increasing the limit to \n\n\n\n\n\u03c7\n\nred\n2\n\n>6\n\n$ \\chi ^2_{\\rm red}>6 $\n\n\n\n adds more sources to our sample, but most of them have bad fits and thus unreliable host galaxy measurements. Reducing the threshold to \n\n\n\n\n\u03c7\n\nred\n2\n\n>4\n\n$ \\chi ^2_{\\rm red}>4 $\n\n\n\n would exclude additional sources, the vast majority of which have reliable fits. This criterion eliminates 133 AGN (9% of the sample). For each parameter calculated by the SED fitting process, X-CIGALE estimates two values. One is evaluated from the best-fit model, and one weights all models allowed by the parametric grid, with the best-fit model having the heaviest weight (Boquien et al. 2019). This weight is based on the likelihood, exp (\u2212\u03c72\/2), associated with each model. A large difference between these two values for a specific parameter indicates that the fitting process did not result in a reliable estimation for this parameter (Mountrichas et al. 2021a,c; Buat et al. 2021). Thus, when we compared the stellar masses of type 1 and 2 AGN, we only included in our analysis sources with \n\n\n\n\n1\n5\n\n\u2264\n\n\nM\n\n\n\u2217\n,\nbest\n\n\n\nM\n\n\n\u2217\n,\nbayes\n\n\n\n\u2264\n5\n\n\n$ \\frac{1}{5}\\leq \\frac{M_{*,\\rm best}}{M_{*, \\rm bayes}} \\leq 5 $\n\n\n, where M\u2217,best and M\u2217,bayes are the best-fit and Bayesian-fit values of M*, respectively. This method to exclude unreliable estimations has been applied in recent studies (Mountrichas et al. 2021a,c; Buat et al. 2021). Using different values for the boundaries of the criterion (i.e. 0.1\u20130.33 for the lower limit and 3\u201310 for the upper limit) does not affect the results of our analysis. This criterion reduces the sample to 944 X-ray AGN (93% of the initial dataset). Of these systems, 729 are type 1 and 215 are type 2 (Table 2). Similarly, for the comparison of the SFR of the two AGN types, we require \n\n\n\n\n1\n5\n\n\u2264\n\n\n\nSFR\nbest\n\n\n\n\nSFR\nbayes\n\n\n\n\u2264\n5\n\n\n$ \\frac{1}{5}\\leq \\frac{\\mathrm{SFR_{best}}}{\\mathrm{SFR_{bayes}}} \\leq 5 $\n\n\n, where SFRbest and SFRbayes are the best and Bayesian values of the SFR, respectively, estimated by X-CIGALE. This reduces our sample to 860 X-ray AGN (85% of the initial dataset). Of these sources, 673 are type 1 and 187 are type 2. We note that throughout our analysis we used the Bayes calculations of X-CIGALE for the various parameters.","Citation Text":["Boquien et al. 2019"],"Functions Text":["One is evaluated from the best-fit model, and one weights all models allowed by the parametric grid, with the best-fit model having the heaviest weight"],"Functions Label":["Uses"],"Citation Start End":[[1030,1049]],"Functions Start End":[[877,1028]]}
{"Identifier":"2015AandA...584A.103S__Baym_et_al._1971b_Instance_1","Paragraph":"Also plotted in Fig. 2 is the pressure in the outer crust from some popular EoSs that model the complete structure of the NS. The figure is drawn up to nb = 3  \u00d7  10-4 fm-3, thus comprising the change from the outer crust to the inner crust in order to allow comparison of the EoSs also in this region (notice, however, that inner crust results are not available for the FRDM). We show in Fig. 2 the EoS from the recent BSk21 Skyrme nuclear effective force (Pearson et al. 2012; Fantina et al. 2013; Potekhin et al. 2013; Goriely et al. 2010) tabulated in (Potekhin et al. 2013). The parameters of this force were fitted (Goriely et al. 2010) to reproduce with high accuracy almost all known nuclear masses, and to various physical conditions including the neutron matter EoS from microscopic calculations. We see in Fig. 2 that after the experimentally constrained region, the BSk21 pressure is similar to the BCPM and FRDM results, with just some displacement around the densities where the composition changes from a nucleus to the next one. In the seminal work of BPS (Baym et al. 1971b) the nuclear masses for the outer crust were provided by an early semi-empirical mass table. The corresponding EoS is seen to display a similar pattern with the BCPM, FRDM, and BSk21 results in Fig. 2. The EoS by Lattimer & Swesty (1991), taken here in its Ska version (Lattimer 2015; LS-Ska), and the EoS by (Shen et al. 1998b,a; Sumiyoshi 2015; Shen-TM1) were computed with, respectively, the Skyrme force Ska and the relativistic mean-field model TM1. In the two cases the calculations of masses are of semiclassical type and A and Z vary in a continuous way. Therefore, these EoSs do not present jumps at the densities associated with a change of nucleus in the crust. Beyond this feature, the influence of shell effects in the EoS is rather moderate because to the extent that the pressure at the densities of interest is largely determined by the electrons, small changes of the atomic number Z compared with its semiclassical estimate modify only marginally the electron density and, consequently, the pressure. The LS-Ska EoS shows good agreement with the previously discussed models, with some departure from them in the transition to the inner crust. The largest discrepancies in Fig. 2 are observed with the Shen-TM1 EoS that in this region predicts a softer crustal pressure with density than the other models. ","Citation Text":["Baym et al. 1971b"],"Functions Text":["In the seminal work of BPS","the nuclear masses for the outer crust were provided by an early semi-empirical mass table.","The corresponding EoS is seen to display a similar pattern with the BCPM, FRDM, and BSk21 results in Fig. 2."],"Functions Label":["Background","Background","Similarities"],"Citation Start End":[[1073,1090]],"Functions Start End":[[1045,1071],[1092,1183],[1184,1292]]}
{"Identifier":"2019ApJ...881...38E__Pinto_et_al._2016_Instance_1","Paragraph":"We fitted the XMM-Newton spectra of ULX-1 and ULX-2 from observation 0794581201 with an absorbed power-law model, using two tbabs absorption components, one frozen to the Galactic value of NH = 1.84 \u00d7 1021 cm\u22122 and the other allowed to vary. In the case of ULX-1, a power-law model was not sufficient to produce a good fit, with \u03c72\/dof = 131.9\/90 and the fit showing significant soft residuals at \u223c1 keV. These soft residuals are known to be a common feature in the spectra of ULXs and bright X-ray binaries (e.g., Bauer & Brandt 2004; Carpano et al. 2007; Middleton et al. 2015b), and found in other ULXs to be a combination of emission and absorption features related to powerful outflowing winds (including in NGC 6946 ULX-3; Pinto et al. 2016). We used an additional Gaussian component to empirically fit these soft residuals, which resulted in a very large statistical improvement and provided an acceptable fit (\u03c72\/dof = 84.5\/87). Using the best-fitting model for each source, we calculated the absorbed 0.3\u201310 keV flux of both objects. We do not correct for absorption since extending steep power laws, which are empirical rather than physical models, to low energies is likely to overestimate the true flux of the object beneath the absorption\u2014unabsorbed fluxes using such models can be a factor of 2 or 3 higher than those for more physically motivated models, and the amount of absorption is itself model-dependent, so we make fewer assumptions by just considering the observed, absorbed flux. For comparison, we also fitted both spectra with an absorbed multicolor disk model. Although this was able to produce a statistically acceptable fit, the Gaussian representing the soft residuals in ULX-1 contributes to a physically unreasonable portion of the spectrum, broadening to fit the majority of the soft emission, and the residuals for ULX-2 are even less well characterized than for a power-law fit. Therefore, a steep power law seems to be the best empirical model for both spectra. We present the spectral fit results for these two sources in Table 3 and the best-fitting model plots in Figure 2.","Citation Text":["Pinto et al. 2016"],"Functions Text":["These soft residuals are known to be a common feature in the spectra of ULXs and bright X-ray binaries","and found in other ULXs to be a combination of emission and absorption features related to powerful outflowing winds (including in NGC 6946 ULX-3;"],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[729,746]],"Functions Start End":[[405,507],[582,728]]}
{"Identifier":"2019MNRAS.487.5666S__Wolfe_et_al._1986_Instance_1","Paragraph":"The observations of the characteristic double-horned line profiles (Stewart et al. 2014) are all restricted to very low redshifts (z  0.5). Our understanding of the kinematics of H\u2009i at high redshifts (z > 0.5) majorly depends on the study of the QSO absorption spectra where the absorption features are mainly dominated by the Lyman\u2212\u03b1 systems (Zafar et al. 2013). The systems with the largest column densities, the DLAs, are proposed to be the progenitors of the present-day spiral galaxies (Wolfe et al. 2005). Modelling of DLAs observations suggests that DLAs resemble rotating disc galaxies with circular velocities typically of the order of $100\\!-\\!200\\, {\\rm km\\, s}^{-1}$ (Wolfe et al. 1986; Kauffmann & Charlot 1994; Klypin et al. 1995; Lanzetta, Wolfe & Turnshek 1995; Wolfe 1995; Jedamzik & Prochaska 1998). A number of theoretical studies indicate that the DLAs can have circular velocities as low as ${\\sim } 50 \\, {\\rm km\\, s}^{-1}$ (e.g. Kauffmann 1996) which is supported by a number of numerical simulations (Pontzen et al. 2008; Cen 2012; Bird et al. 2014, 2015). QSO absorption spectra also suggest that the gas inside the DLAs has a velocity dispersion of $\\sigma _v \\approx 5\\!-\\!10\\, {\\rm km\\, s}^{-1}$ (Wolfe et al. 2005) which is comparable with the nearby galaxies. Other observations also reveal that the high-redshift galaxies show rotational motion (e.g. see Pettini 2009 and references therein). Based on these considerations we have assumed that across the entire range z \u2264 6 the H\u2009i resides in rotating disc galaxies which exhibit a double-horned line profile similar to those seen for local galaxies. The value of the parameters \u03b1, hf, \u03c3v are found to vary from galaxy to galaxy in the local Universe. The statistics of these parameters and their redshift evolution are largely unknown. In our analysis each simulation corresponds to fixed values of these parameters which are held constant over the entire z range. We have carried out simulations covering the entire range of parameter values in order to estimate how this variation affects the RSD. Considering vcirc which determines the overall width of the line profile, it may be noted that the values are determined by the halo mass distribution which evolves with z.","Citation Text":["Wolfe et al. 1986"],"Functions Text":["Modelling of DLAs observations suggests that DLAs resemble rotating disc galaxies with circular velocities typically of the order of $100\\!-\\!200\\, {\\rm km\\, s}^{-1}$"],"Functions Label":["Background"],"Citation Start End":[[681,698]],"Functions Start End":[[513,679]]}
{"Identifier":"2020AandA...639A..46B__\u0160tver\u00e1k_et_al._(2009)_Instance_2","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)"],"Functions Text":["Findings by",", 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."],"Functions Label":["Differences","Differences"],"Citation Start End":[[783,804]],"Functions Start End":[[771,782],[804,1010]]}
{"Identifier":"2017ApJ...834..178Y__Tachihara_et_al._2007_Instance_2","Paragraph":"In order to investigate the gas kinematics at an early evolutionary stage and the formation of Keplerian disks, we conduct ALMA observations toward three candidate young protostars, Lupus 3 MMS, IRAS 15398\u22123559, and IRAS 16253\u22122429. They are selected from our SMA sample (Yen et al. 2015a). These three protostars all have relatively low protostellar masses (0.1 M\u2299), inferred from the infalling motions in their protostellar envelopes, and they do not show clear signs of a spin-up rotation on a 1000 au scale; i.e., no signatures of Keplerian disks are seen in our SMA observations (Yen et al. 2015a). Lupus 3 MMS is a Class 0 protostar with a bolometric luminosity (Lbol) of 0.41 L\u2299 and a bolometric temperature (Tbol) of 39 K in the Lupus 3 cloud at a distance of 200 pc (Tachihara et al. 2007; Comer\u00f3n 2008; Dunham et al. 2013). Our SMA results suggest that the protostellar mass in Lupus 3 MMS can be as low as 0.1 M\u2299 (Yen et al. 2015a). IRAS 15398\u22123559 is a Class 0\/I protostar with an Lbol of 1.2 L\u2299 and a Tbol of 61 K in the Lupus 1 cloud at a distance of 150 pc (Froebrich 2005; Comer\u00f3n 2008). Early single-dish observations of its CO outflow suggest that IRAS 15398\u22123559 is close to face on (van Kempen et al. 2009). Recent SMA and ALMA observations show that it is actually closer to edge on (Oya et al. 2014; Bjerkeli et al. 2016). With this new estimated inclination angle (\u223c70\u00b0), our SMA data suggest a low protostellar mass (0.1 M\u2299) and a low specific angular momentum in the protostellar envelope (\u223c1 \u00d7 10\u22124 km s\u22121 pc; Yen et al. 2015a). IRAS 16253\u22122429 is a Class 0 protostar with an Lbol of 0.24 L\u2299 and a Tbol of 36 K in the \u03c1 Ophiuchus star-forming region at a distance of 125 pc (Dunham et al. 2013). Both CARMA and our SMA results suggest that its protostellar mass is 0.1 M\u2299 (Tobin et al. 2012a; Yen et al. 2015a). These three protostars are all embedded in dense cores with masses \u22730.5 M\u2299 (Froebrich 2005; Tachihara et al. 2007; Enoch et al. 2009). Therefore, they are excellent targets by which to study the gas motions on a 100 au scale at an early evolutionary stage.","Citation Text":["Tachihara et al. 2007"],"Functions Text":["These three protostars are all embedded in dense cores with masses \u22730.5 M\u2299","Therefore, they are excellent targets by which to study the gas motions on a 100 au scale at an early evolutionary stage."],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[1930,1951]],"Functions Start End":[[1838,1912],[1973,2094]]}
{"Identifier":"2017AandA...601A.130R__Hern\u00e1ndez_et_al._2013_Instance_1","Paragraph":"Various strategies have been devised to identify modes. One can, for instance, search for frequency patterns appropriate for rapid rotation. The background for this search is the discovery of asymptotically uniform frequency spacings in the numerically computed spectra of uniformly rotating polytropic models (Ligni\u00e8res et al. 2006; Reese et al. 2008) and differentially rotating realistic self-consistent field (SCF) models (Reese et al. 2009a). These uniform spacings have also been modelled through asymptotic semi-analytical formulas (Pasek et al. 2012). In observed spectra, recurrent frequency spacings that may correspond to the large separation or half its value have been found in some stars (Garc\u00eda Hern\u00e1ndez et al. 2009; Garc\u00eda Hern\u00e1ndez et al. 2013; Papar\u00f3 et al. 2016). Moreover, Garc\u00eda Hern\u00e1ndez et al. (2015) show that mean density estimates based on this type of spacings (obtained via a scaling relation similar to the one in Reese et al. 2008, but based on SCF models) are compatible with independent mass and radii measurements obtained for \u03b4 Scuti stars in binary systems. Nonetheless, it is expected that various effects may contribute to hide these regular frequency patterns. First, as mentioned before, the full spectrum is a superposition of sub-spectra corresponding to different classes of modes and some of the uniform spacings only concern one class. This complicates their detection in the full spectrum. Also, owing to their asymptotic nature, these spacings might not be relevant to analyse the low to moderate (up to radial order n ~ 10) frequency domain, typical of most rapidly rotating pulsators. A third effect that may come into play is the presence of mixed modes in evolved stars and\/or sharp sound speed gradients since they can potentially modify the regular spacings. Finally, mode selection effects that are due to the non-linearly determined intrinsic mode amplitudes could affect the detectability of the regular patterns. As a first attempt, Reese et al. (2009b) developed a strategy to find these frequency spacings but ran into difficulties when including chaotic modes, which come from another class of modes. Ligni\u00e8res et al. (2010) addressed the same question with encouraging results but their analysis was restricted to the asymptotic regime and relied on simplifying assumptions regarding the spectrum of chaotic modes and the mode visibilities. In this paper, our first goal is to search for regular frequency spacings in the most realistic synthetic spectra available, using relevant frequency ranges and accurate visibility calculations. While they can provide guidance to a similar search in real data, we already know that these results must be taken with caution since the intrinsic mode amplitudes used in this paper are not realistic, but based on ad-hoc prescriptions. ","Citation Text":["Garc\u00eda Hern\u00e1ndez et al. 2013"],"Functions Text":["In observed spectra, recurrent frequency spacings that may correspond to the large separation or half its value have been found in some stars"],"Functions Label":["Background"],"Citation Start End":[[733,761]],"Functions Start End":[[560,701]]}
{"Identifier":"2018MNRAS.473.3810Y__Mitrushchenkov_et_al._2017_Instance_2","Paragraph":"The lack of data on inelastic processes due to collisions with neutral hydrogen atoms has been a major limitation on modelling of F-, G- and K-star spectra in statistical equilibrium, and thus to reliably proceeding beyond the assumption of local thermodynamic equilibrium (LTE) in analysis of stellar spectra and the determination of elemental abundances. This problem has been well documented, e.g. see Lambert (1993); Barklem (2016a) and references therein. Significant progress has been made in recent times through detailed full-quantum scattering calculations, based on quantum chemical data, for the cases of simple atoms such as Li, Na, Mg and Ca (Belyaev & Barklem 2003; Barklem, Belyaev & Asplund 2003; Belyaev et al. 2010; Barklem et al. 2010; Belyaev et al. 2012; Barklem et al. 2012; Mitrushchenkov et al. 2017). These calculations have demonstrated the importance of the ionic-covalent curve crossing mechanism leading naturally to charge transfer processes (mutual neutralization and ion-pair production), in addition to excitation and de-excitation processes. The importance of this mechanism has allowed various simplified model approaches to be developed, which may be used in cases where suitable quantum chemistry data are not been available. In particular a semi-empirical model has been employed for Al, Si, Be and Ca (Belyaev 2013a,b; Belyaev, Yakovleva & Barklem 2014b; Yakovleva, Voronov & Belyaev 2016; Belyaev et al. 2016), and a theoretical model based on a two-electron asymptotic linear combinations of atomic orbitals (LCAO) approach, has also been employed for Ca (Barklem 2016b, 2017). Comparisons of the two methods show quite good agreement and reasonable agreement with the full quantum calculations is found, particularly for the most important processes with the largest rates (Barklem 2016b, 2017; Mashonkina, Sitnova & Belyaev 2017; Mitrushchenkov et al. 2017). Thus, the model approaches provide a useful route for obtaining estimates of the rates for these processes for many elements of astrophysical interest.","Citation Text":["Mitrushchenkov et al. 2017"],"Functions Text":["Comparisons of the two methods show quite good agreement and reasonable agreement with the full quantum calculations is found, particularly for the most important processes with the largest rates"],"Functions Label":["Similarities"],"Citation Start End":[[1873,1899]],"Functions Start End":[[1619,1814]]}
{"Identifier":"2018ApJ...863..162M__Liu_et_al._2013_Instance_2","Paragraph":"NLFFF extrapolation provides the reconstructed coronal magnetic field for AR 11158 from 2011 February 13 \u2212 2011 February 15 (Figures 1(d)\u2013(f)). The field lines (yellow lines) within the core of the AR have arcade-like structure with a relatively strong twist mainly near the PIL. These figures show that the magnetic field evolved during this period. Although we did not quantitatively compare the field lines with the observation, in general, the reconstructed coronal field morphologies match with the observations in Figures 1(a)\u2013(c). The general morphologies and the locations of the high-twist fields are also in agreement with many previous studies (Jing et al. 2012; Sun et al. 2012; Dalmasse et al. 2013; Inoue et al. 2013, 2014a; Liu et al. 2013; Wang et al. 2013; Aschwanden et al. 2014; Malanushenko et al. 2014; Zhao et al. 2014). Unlike Zhao et al. (2014) who could identify the twisted flux rope from the topology of the reconstructed coronal field, we could not find an obvious topological signature of a flux rope existing in our NLFFF during our analysis time window. It might be due to the fact that there was little magnetic flux with twist higher than one turn in our NLFFF and it is difficult to topologically define it as a flux rope. However, our result is consistent with other NLFFF results (Jing et al. 2012; Sun et al. 2012; Liu et al. 2013; Wang et al. 2013; Inoue et al. 2014a; Malanushenko et al. 2014). The high-twist region in our result is also in agreement with the region with high helicity flux (Dalmasse et al. 2013) and the location of the flare ribbons (Bamba et al. 2013; Liu et al. 2013), as well as the high current density region (Janvier et al. 2014). Figures 1(g)\u2013(i) show the evolution of the twist distribution map, with the magnetic twist of the field lines plotted at the footpoints of field lines according to a color scale. This shows that the high-twist (strongly right-handed twist corresponding to Tw > 0.5) areas are concentrated in only a limited part of the AR. The high-twist area grew and became even more twisted just before the X2.2 flare (Figure 1(i)). Most parts of the AR have twist values less than 0.25, but near the PIL the twist can reach more than 0.5, even up to about a full turn. This is consistent with the results of Sun et al. (2012) and Inoue et al. (2014a). A high-twist (strong negative\/left-handed twist) area also developed in the eastern part of the AR, which did not exist initially on February 13. Both of these high-twist areas produced several flares. However, here we focus on the flares that resulted from the high-twist core region near the center of the AR, where the M6.6 and X2.2 flares occurred.","Citation Text":["Liu et al. 2013"],"Functions Text":["However, our result is consistent with other NLFFF results"],"Functions Label":["Similarities"],"Citation Start End":[[1352,1367]],"Functions Start End":[[1257,1315]]}
{"Identifier":"2015AandA...582A..42K__Cardelli_et_al._(1989)_Instance_1","Paragraph":"Spectral energy distributions in accordance with the physical parameters of the best-fit spectra were generated with FASTWIND, to provide the flux density per surface unit through our studied bands. The composite flux measured at Earth from a binary at a distance d, at a wavelength \u03bb, reddened to extinction A(\u03bb), is given by \\begin{eqnarray*} f_{\\lambda}=\\dfrac{1}{d^2}\\left(R_{1}^{2}F_{1,\\lambda}+R_{2}^{2}F_{2,\\lambda}\\right) \\times 10^{-0.4A(\\lambda)} \\end{eqnarray*}f\u03bb=1d2(R12F1,\u03bb+R22F2,\u03bb)\u00d710\u22120.4A(\u03bb)where R1,2 and F1,2 are the radii and the surface fluxes of the components, respectively. The composite SED was reddened according to the new family of optical and near-infrared extinction laws for O-type stars provided by Ma\u00edz Apell\u00e1niz et al. (2014), which constitute an improvement of the widely used extinction laws by Cardelli et al. (1989). We ran a fitting algorithm over a wide range of distances with a step of 0.05 kpc, setting free the monochromatic parameters R5495 and E(4405 \u2212 5495) for the type and amount of extinction, respectively, and the best-fit values were considered to be those that minimized the weighted-\u03c72. To estimate the uncertainty of our measurements, we used a Monte Carlo approach. In particular, we ran the fitting procedure 1000 times using sets of randomly selected parameters (photometry and radii) within their uncertainties assuming a Gaussian distribution is used. The corresponding values of distance are shown in Table 6 for every set of radii and temperatures determined from the four different fit models. Our adopted radii measured with ELC yielded d = 3.52 \u00b1 0.08 kpc, E(4405 \u2212 5495) = 3.66 \u00b1 0.06 mag, R5495 = 3.26 \u00b1 0.04 and A5495 = 11.9 \u00b1 0.1 mag, based on a photometric Teff2 ~ 36 kK. Assuming a spectroscopic Teff2 = 37 kK, the distance changed slightly to d = 3.55 \u00b1 0.08 kpc. Having three degrees of freedom, the reduced-\u03c72 of our resulting fits was measured to be \\hbox{$\\chi^2_\\textrm{red}\\sim11$}\u03c7red2~11. ","Citation Text":["Cardelli et al. (1989)"],"Functions Text":["The composite SED was reddened according to the new family of optical and near-infrared extinction laws for O-type stars provided by Ma\u00edz Apell\u00e1niz et al. (2014), which constitute an improvement of the widely used extinction laws by"],"Functions Label":["Uses"],"Citation Start End":[[829,851]],"Functions Start End":[[596,828]]}
{"Identifier":"2022AandA...661A.129S__Rodr\u00edguez-Almeida_et_al._2021_Instance_3","Paragraph":"Radio astronomy is recognized as one of the most effective techniques to search for interstellar molecules. By comparing the spectra of candidate molecules in the laboratory with the spectra observed in astronomical surveys, we can determine whether these molecules exist in interstellar space. Therefore, it is necessary to provide rotational spectra of candidates for astronomical detection. Radio astronomy has helped to detect several sulfur-containing molecules in the ISM in recent years: in particular, thiols, the sulfur analogs of alcohols. Methanethiol (or methyl mercaptan, CH3SH) was detected in the Sagittarius B2 (Sgr B2) region of the center of our Galaxy (Linke et al. 1979; Gibb et al. 2000; M\u00fcller et al. 2016; Rodr\u00edguez-Almeida et al. 2021) and in the protostar IRAS 16293-2422 (Majumdar et al. 2016). Two groups reported to have detected several signs of ethanethiol (C2H5SH) in Sgr B2 (M\u00fcller et al. 2016) and Orion (Kolesnikov\u00e1 et al. 2014). Most recently, Rodriguez-Almeida reported the first unambiguous detection of ethanethiol in the ISM, toward the G+0.603-0.027 molecular cloud (Rodr\u00edguez-Almeida et al. 2021). Moreover, several sulfur-containing species have been observed in comets (Altwegg et al. 2017). Some recent efforts, both from spectroscopy and astronomical searches, to detect S-stitutes of other classes of compounds have also been reported. For instance, thioformic acid (HC(O)SH) was very recently detected in G+0.693\u20130.027. Its trans-isomer has an abundance of ~1 \u00d7 10\u201310 (Rodr\u00edguez-Almeida et al. 2021). Conversely, thioformamide (NH2CHS), the counterpart of for-mamide (NH2CHO), was characterized in the laboratory up to 660 GHz, and its transitions were searched for toward the hot cores Sgr B2(N1S) and Sgr B2(N2), but it was not detected (Motiyenko et al. 2020). The rotational spectrum of thioac-etamide was recently analyzed in the 59.6\u2013110.0 GHz frequency region (5.03\u20132.72 mm). Its emission was searched for in regions associated with star formation using the IRAM 30 m ASAI observations toward the prestellar core L1544 and the outflow shock L1157\u2013B1. The molecule was not detected, but the study allowed placing constraints on the thioacetamide abundances (Maris et al. 2019; Remijan et al. 2022).","Citation Text":["Rodr\u00edguez-Almeida et al. 2021"],"Functions Text":["For instance, thioformic acid (HC(O)SH) was very recently detected in G+0.693\u20130.027. Its trans-isomer has an abundance of ~1 \u00d7 10\u201310"],"Functions Label":["Background"],"Citation Start End":[[1516,1545]],"Functions Start End":[[1382,1514]]}
{"Identifier":"2015MNRAS.454.1468K__Winckel_2003_Instance_2","Paragraph":"Owing to their dusty circumstellar environments, a large mid-infrared (mid-IR) excess is a characteristic feature of post-AGB stars and a detection of cold circumstellar material using mid-IR photometry can be used to identify these objects. The first extensive search for these objects was initiated in the mid-80's using results from the Infrared Astronomical Satellite (Neugebauer et al. 1984) which enabled the identification of post-AGB stars in our Galaxy (Kwok 1993). The Toru$\\acute{\\rm n}$ catalogue (Szczerba et al. 2007) for Galactic post-AGB stars lists around 391 very likely post-AGB objects. The Galactic sample of post-AGB stars have been found to be a very diverse group of objects (Van Winckel 2003). Studies showed that the majority of the optically visible Galactic post-AGB stars could be classified based on their spectral energy distributions (SEDs) into two groups: shell-sources and disc-sources (Van Winckel 2003). The shell-sources show a double-peaked SED with the hot central star peaking at shorter wavelengths while the cold, detached, expanding dust shell peaks at longer wavelengths. This type of SED is considered to be characteristic of objects that follow the single-star evolution scenario mentioned above. The disc-sources do not show two distinct flux peaks in the mid-IR but they do display a clear near-infrared (near-IR) excess indicating that circumstellar dust must be close to the central star, near sublimation temperature. It is now well established that this feature in the SED indicates the presence of a stable compact circumbinary disc, and therefore these sources are referred to as disc-sources (de Ruyter et al. 2006; Deroo et al. 2007; Gielen et al. 2011a; Hillen et al. 2013). The rotation of the disc was resolved with the ALMA array (Bujarrabal et al. 2013a) in one object and using single dish observations Bujarrabal et al. (2013b) confirmed that disc rotation is indeed widespread. Moreover, these disc-sources are confirmed to be binaries and show orbital periods between 100 and 2000 d (Van Winckel et al. 2009; Gorlova et al. 2014). In contrast, for the Galactic shell-sources long-term radial velocity monitoring efforts have not yet resulted in any clear detected binary orbit (Hrivnak et al. 2011), which either confirms the single-star nature of these objects or introduces a possibility that these systems can have companions on very wide orbits.","Citation Text":["Van Winckel 2003"],"Functions Text":["Studies showed that the majority of the optically visible Galactic post-AGB stars could be classified based on their spectral energy distributions (SEDs) into two groups: shell-sources and disc-sources"],"Functions Label":["Background"],"Citation Start End":[[922,938]],"Functions Start End":[[719,920]]}
{"Identifier":"2019MNRAS.487...24G__Rogers_2015_Instance_2","Paragraph":"NASA\u2019s Kepler mission has unveiled a wealth of new planetary systems (e.g. Borucki et al. 2010). These systems offer new insights into the process of planet formation and evolution. One of Kepler\u2019s key findings is that the most common planets in our Galaxy, observed to date, are between 1 and 4 R\u2295, i.e. larger than Earth but smaller than Neptune (Fressin et al. 2013; Petigura, Marcy & Howard 2013). Further observations revealed a transition in average densities at planet sizes \u223c1.5 R\u2295 (Marcy et al. 2014; Rogers 2015), with smaller planets having densities consistent with rocky compositions while larger planets having lower densities indicating significant H\/He envelopes. In addition, Owen & Wu (2013) noticed a bimodal distribution of observed planet radii. Since then, refined measurements have provided strong observational evidence for the sparseness of short-period planets in the size range of \u223c1.5\u20132.0 R\u2295 relative to the smaller and larger planets, yielding a valley in the small exoplanet radius distribution (e.g. Fulton et al. 2017; Fulton & Petigura 2018). For example, the California-Kepler Survey reported measurements from a large sample of 2025 planets, detecting a factor of \u223c2 deficit in the relative occurrence of planets with sizes \u223c1.5\u20132.0 R\u2295 (Fulton et al. 2017). Studies suggest that this valley likely marks the transition from the smaller rocky planets: \u2018super-Earths\u2019, to planets with significant H\/He envelopes typically containing a few per cent of the planet\u2019s total mass: \u2018sub-Neptunes\u2019 (e.g. Lopez & Fortney 2013, 2014; Owen & Wu 2013; Rogers 2015; Ginzburg, Schlichting & Sari 2016). Furthermore, the location of this valley is observed to decrease to smaller planet radii, Rp, with increasing orbital period, P. In a recent study involving asteroseismology-based high precision stellar parameter measurements for a sample of 117 planets, a slope $\\text{d log} R_\\mathrm{ p}\/ \\text{d log} P = -0.09^{+0.02}_{-0.04}$ was reported for the radius valley by Van Eylen et al. (2018). A similar value for the slope of $-0.11^{+0.03}_{-0.03}$ was reported by Martinez et al. (2019).","Citation Text":["Rogers 2015"],"Functions Text":["Studies suggest that this valley likely marks the transition from the smaller rocky planets: \u2018super-Earths\u2019, to planets with significant H\/He envelopes typically containing a few per cent of the planet\u2019s total mass: \u2018sub-Neptunes\u2019 (e.g."],"Functions Label":["Background"],"Citation Start End":[[1574,1585]],"Functions Start End":[[1293,1529]]}
{"Identifier":"2022ApJ...936..102A__Williams_et_al._2006_Instance_1","Paragraph":"Following the original BGK kinetic-theoretical scheme (Bernstein et al. 1957), we transform the above equations to the energy frame, defined as\n5\n\n\n\nw=12v2+\u03d5.\n\nIn the energy frame, the ion distribution transforms into\n6\n\n\n\nfi(w)=\u0393(\u03bai)\u03c0\u03bai\u22123\/2\u0393(\u03bai\u22121\/2)1+2w\u03bai\u22123\/2\u2212\u03bai,\n\nwhere\n7\n\n\n\nf(x,v)dv=f(w)dw\/2w\u2212\u03d5.\n\nHere, the total energy (w) is normalized with \n\n\n\nmvth,i2\n\n, i.e., 2k\n\nB\n\nT\n\ni\n. As the ions encounter a negative potential well (a pulse), depending on their respective velocities, some of them will become trapped and some of them will pass through. Hence two types of population exist: a trapped population, and a passing population. We assume the passing populations to follow the initial distribution function - a kappa distribution function. In addition, we also assume the form of the potential in which particles are trapped to be prescribed. Spacecraft observations show that wave potential structures of Gaussian form are common in space and astrophysical plasmas (Matsumoto et al. 1994; Williams et al. 2006). A negative wave potential well acts as a perturbation capable of trapping ions in it. We assume this potential well to have Gaussian form, given by\n8\n\n\n\n\u03d5(x)=\u2212\u03c8exp\u2212x22\u03b42,\n\nwhere \u03c8 (>0) denotes the amplitude, and \u03b4 is the width of the perturbation, respectively. More precisely, \u03b4 is actually the distance where the potential decreases to 0.6065 times the maximum amplitude of \u03c8. The FWHM of the perturbation is actually given by \u0394 = 2.35\u03b4. The net charge density can thus be expressed as\n9\n\n\n\nd2\u03d5dx2=1\u2212Hne\u2212ni,p\u2212ni,tr+H,\n\nwhere n\n\ni,p\n is the passing ion density, and \n\n\n\nni,tr\n\n is the trapped ion density. Particles with suitable velocities falling in the potential range \n\n\n\n\u2212\u2212\u03d5,+\u2212\u03d5\n\n will become trapped, and the rest will pass through. Thus, the range of integration for both passing and trapped ion distributions is given by\n10\n\n\n\nni,p=\u222b\u2212\u221e\u2212\u2212\u03d5fp(x,v)dv+\u222b+\u2212\u03d5\u221efp(x,v)dv.\n\nAs the passing ions follow the kappa distribution, we obtain the passing ion density as\n11\n\n\n\nni,p=1\u22122AB\u2212\u03d52F1[\u03bai,1\/2,3\/2;\u03d5\/B],\n\nwhere 2\nF\n1 is the hypergeometric function of the first kind. Now that we have obtained the passing ion density, we move on to the derivation of the trapped ion density. In terms of distribution function, the trapped ion density is given by\n12\n\n\n\nni,tr=\u222b\u2212\u2212\u03d5+\u2212\u03d5ftr(x,v)dv.\n\nWe may now derive the trapped ion density by rearranging Equation (9) as\n13\n\n\n\nni,tr=1\u2212Hne\u2212ni,p\u2212d2\u03d5dx2+H.\n\nSubstituting from Equation (11) and differentiating the potential in Equation (8) twice, in Equation (13), we obtain the trapped electron density as\n14\n\n\n\nni,tr=1\u2212H1\u2212Tr\u03d5(\u03bae\u22123\/2)0.5\u2212\u03bae+2\u03d5log\u2212\u03d5\u03c8\u03b42+\u03d5\u03b42\u22121+2AB\u2212\u03d52F1[\u03bai,1\/2,3\/2;+\u03d5\/B]+H.\n\nFor convenience, we make a transformation \u2212 \u03d5 = \u2223\u03d5\u2223 \u2192 \u03c1, (>0), and we obtain the trapped ion density as\n15\n\n\n\nni,tr=1\u2212H\u03c1Tr\u03bae\u221232+10.5\u2212\u03bae\u22122\u03c1log\u03c1\u03c8\u03b42\u2212\u03c1\u03b42\u22121+2AB\u03c12F1[\u03bai,1\/2,3\/2;\u2212\u03c1\/B]+H.\n\n\n","Citation Text":["Williams et al. 2006"],"Functions Text":["Spacecraft observations show that wave potential structures of Gaussian form are common in space and astrophysical plasmas","We assume this potential well to have Gaussian form"],"Functions Label":["Uses","Uses"],"Citation Start End":[[997,1017]],"Functions Start End":[[850,972],[1106,1157]]}
{"Identifier":"2016ApJ...821...74J__D\u00edaz_et_al._2016_Instance_1","Paragraph":"Recent theoretical work has suggested that the presence, or lack thereof, of long-period giant planets could affect the formation of such systems. Batygin & Laughlin (2015) argued that the migration of Jupiter within our own solar system might have disrupted a massive primordial inner protoplanetary disk that could have formed multiple short-period super-Earths; they predicted that systems like the Kepler short-period multiple systems should typically lack long-period giant planets. A related question is, how common are planetary systems broadly similar in architecture to our solar system, with small close-in planets and more distant giant planets? We can begin to answer these questions in the near future through the combination of searches for short-period super-Earths and data from the long-term RV programs that have been monitoring many bright FGK stars for well over a decade. Super-Earths can be found with either high-precision RV observations or space-based transit searches. Such high-precision RV surveys include those being undertaken currently with HARPS (e.g., D\u00edaz et al. 2016), HARPS-N (M15), APF (Vogt et al. 2014), and CHIRON (Tokovinin et al. 2013), and in the near future with MINERVA (Swift et al. 2015), CARMENES (Quirrenbach et al. 2014), ESPRESSO (M\u00e9gevand et al. 2014), and SPIRou (Artigau et al. 2014). The major upcoming space-based transit survey is that of TESS (Ricker et al. 2015). Long-term RV programs include the McDonald Observatory Planet Search (e.g., Endl et al. 2016), the Anglo-Australian Planet Search (e.g., Jones et al. 2010), the Lick-Carnegie Exoplanet Survey (e.g., Rowan et al. 2016), the CORALIE planet search (Marmier et al. 2013), and the planet search at ESO (e.g., Zechmeister et al. 2013). Long-period giant planets will also be found by Gaia, which will produce a huge sample of astrometrically detected planets (Perryman et al. 2014). While most of the Kepler sample is too faint to have been observed previously by long-term RV surveys (e.g., Coughlin et al. 2015), Gaia will be able to astrometrically detect long-period planets around many of these stars. Our own McDonald Observatory Planet Search program now has a baseline of 12\u201315 years for \u223c200 FGKM stars, and a handful of stars also have lower precision observations dating back more than 25 years. HD 219134 is one of these stars, and here we present an analysis of our RV observations of this star, as well as our data on the stellar activity.","Citation Text":["D\u00edaz et al. 2016"],"Functions Text":["Super-Earths can be found with either high-precision RV observations or space-based transit searches. Such high-precision RV surveys include those being undertaken currently with HARPS (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[1085,1101]],"Functions Start End":[[893,1084]]}
{"Identifier":"2021ApJ...921..179L__Hayes_et_al._2019_Instance_1","Paragraph":"Quasi-periodic pulsations (QPPs) often refer to the quasi-periodic intensity variations during solar\/stellar flares (see Zimovets et al. 2021, for a recent review). In many observations, the flare QPPs were found to show a nonstationary property in the time series integrated over the whole Sun\/star or over the oscillation region, for instance, each pulsation has an anharmonic and symmetric triangular profile shape (e.g., Kolotkov et al. 2015; Nakariakov et al. 2019). The signature of flare QPPs can be detected in flare light curves across a broad band of the electromagnetic spectrum, i.e., radio\/microwave emissions (Ning et al. 2005; Reznikova & Shibasaki 2011; Nakariakov et al. 2018; Yu & Chen 2019), UV\/EUV wavelengths (Shen et al. 2018; Hayes et al. 2019; Reeves et al. 2020; Miao et al. 2021), SXR\/HXR and \u03b3-ray channels (Nakariakov et al. 2010; Ning 2017; Hayes et al. 2020; Li et al. 2020c), and the H\u03b1 (Srivastava et al. 2008; Kashapova et al. 2020; Li et al. 2020b) or Ly\u03b1 (Van Doorsselaere et al. 2011; Milligan et al. 2017; Li 2021) emissions. The quasi-periods of these QPPs were reported from subseconds to tens of minutes (e.g., Tan et al. 2010; Shen et al. 2013, 2019; Kolotkov et al. 2018; Karlick\u00fd & Ryb\u00e1k 2020; Clarke et al. 2021). It should be stated that the observed periods are generally related to the specific channels or flare phases (Tian et al. 2016; Dennis et al. 2017; Pugh et al. 2019), suggesting that the various classes of QPPs could be produced by different generation mechanisms (e.g., Kupriyanova et al. 2020). In the literature, the flare-related QPPs were most often explained by magnetohydrodynamic (MHD) waves, more specifically sausage waves, kink waves, and slow waves (Li et al. 2020a; Nakariakov & Kolotkov 2020; Wang et al. 2021), or by a repetitive regime of magnetic reconnection that could be spontaneous (i.e., self-oscillatory process) or triggered owing to external MHD oscillations (Thurgood et al. 2017; Yuan et al. 2019; Clarke et al. 2021). They can also be interpreted in terms of the LRC-circuit oscillation in current-carrying loops (Tan et al. 2016; Li et al. 2020b) or caused by the interaction between supra-arcade downflows and flare loops (Xue et al. 2020; Samanta et al. 2021).","Citation Text":["Hayes et al. 2019"],"Functions Text":["The signature of flare QPPs can be detected in flare light curves across a broad band of the electromagnetic spectrum, i.e., radio\/microwave emissions"],"Functions Label":["Background"],"Citation Start End":[[749,766]],"Functions Start End":[[472,622]]}
{"Identifier":"2019AandA...632A.104G__Hirabayashi_et_al._2016_Instance_2","Paragraph":"Finally, our observations are consistent with the bilobate shape of the nucleus of comet 8P\/Tuttle. As noted in Sect. 1, this shape is likely common among comets because it was found for four out of the six comets for which we have spatially resolved images. This is also the case of the trans-Neptunian object 2014 MU69 (Ultima Thule) observed by the New Horizon spacecraft (Stern et al. 2019). This binary configuration has some implications for the formation and evolution of 8P\/Tuttle. A contact binary could result from (i) the accretion at low velocity of two primordial objects (Massironi et al. 2015; Davidsson et al. 2016), (ii) the disruption of a monolithic object due to excessive spin-up resulting from non-gravitational forces or YORP5 effect followed by a reaccretion (Boehnhardt 2004; \u0106uk 2007; Hirabayashi et al. 2016), or (iii) the catastrophic disruption of a monolithic object by a collision followed by a re-accretion (Jutzi & Benz 2017; Schwartz et al. 2018). On the one hand, with a low thermal inertia compared with NEAs, the YORP effect is low for comets, in particular for NIC, which have an elongated orbit and spend most of their time far from the Sun, and it may not be sufficient to increase the spin rate of the nucleus to the point where centrifugal exceed gravitational forces. On the other hand, comet 8P\/Tuttle has been on a very stable orbit for centuries, and it is likely an evolved comet, as suggested by its low activity, so that it could have been much more active in the past. For cometary nuclei, the primary cause for spin-up is torques caused by outgassing, therefore it is possible that 8P\/Tuttle formed as a monolithic body and became a contact binary after its injection into the inner Solar System as a result of excessive spin-up resulting from non-gravitational forces. This scenario has been proposed for comet 67P\/Churyumov-Gerasimenko by Hirabayashi et al. (2016). Alternatively, if the binary nature of comet 8P\/Tuttle is the result of a primordial accretion or a catastrophic collision in the early Solar Sytem, it could have persisted until now. Similar examples are offered by some binary asteroids that can be stable over the age of the Solar System (Chauvineau et al. 1991), or as proposed by Davidsson et al. (2016) for comet 67P\/Churyumov-Gerasimenko. For comet 8P\/Tuttle, it is however not possible to distinguish the solution of a binary nucleus that formed in the first billion years of our Solar System (e.g., Matonti et al. 2019) from a more recent origin following its injection into the inner Solar System (e.g., Hirabayashi et al. 2016).","Citation Text":["Hirabayashi et al. (2016)"],"Functions Text":["For cometary nuclei, the primary cause for spin-up is torques caused by outgassing, therefore it is possible that 8P\/Tuttle formed as a monolithic body and became a contact binary after its injection into the inner Solar System as a result of excessive spin-up resulting from non-gravitational forces. This scenario has been proposed for comet 67P\/Churyumov-Gerasimenko by"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1892,1917]],"Functions Start End":[[1519,1891]]}
{"Identifier":"2020AandA...641A.155V__G\u00f3mez-Guijarro_et_al._2019_Instance_2","Paragraph":"The scenario presented above has been formulated in various flavors to individually explain several of the properties reported here. The main addition of this work, namely the excitation of CO in distant main-sequence and starburst galaxies, fits in the general picture that we sketched. The ensemble of properties and correlations that we reported here can be also used to revisit the definition of what a starburst is. A standard operational classification is based on the distance from the observed empirical M\u22c6-SFR correlation, the main sequence. This proved to be a useful distinction and an excellent predictor of several trends (e.g., Sargent et al. 2014), but recent results, including our present and previous analysis (Puglisi et al. 2019), show that the demarcation between starburst and main-sequence galaxies is more blurred that we previously considered. We do detect starburst-like behaviors in galaxies on the main sequence (Elbaz et al. 2018), likely linked to the existence of transitional objects (Popping et al. 2017; Barro et al. 2017b; G\u00f3mez-Guijarro et al. 2019; Puglisi et al. 2019, and in prep. to limit the references to recent works based on submillimeter observations). Such transition might well imply an imminent increase of the SFR, driving the object in the realm of starbursts (e.g., Barro et al. 2017b), or its cessation, bringing the system back onto or even below the main sequence (G\u00f3mez-Guijarro et al. 2019; Puglisi et al. 2019), with the CO properties potentially able to distinguish between these two scenarios. Regardless of these transitional objects, a definition of starburst based on \u03a3SFR, rather than \u0394MS, would naturally better account for the observed molecular gas excitation properties, dust temperatures and opacities, or SFE (see also Elbaz et al. 2011; Rujopakarn et al. 2011; Jim\u00e9nez-Andrade et al. 2018; Tacconi et al. 2020). As an example, in Fig. 8 we show the mean SLED of the subsample of galaxies with both CO (2\u2005\u2212\u20051) and CO (5\u2005\u2212\u20054) coverage, split at its median \u03a3SFR. While only tentative at this stage, this suggests a trend of increasing CO excitation with \u03a3SFR, consistently with Fig. 7 and what mentioned above.","Citation Text":["G\u00f3mez-Guijarro et al. 2019"],"Functions Text":["Such transition might well imply an imminent increase of the SFR, driving the object in the realm of starbursts","or its cessation, bringing the system back onto or even below the main sequence","with the CO properties potentially able to distinguish between these two scenarios."],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Future Work"],"Citation Start End":[[1419,1445]],"Functions Start End":[[1198,1309],[1338,1417],[1469,1552]]}
{"Identifier":"2018ApJ...866L...1S__Pecharrom\u00e1n_et_al._1999_Instance_3","Paragraph":"It was found that the complex dielectric function from Pecharrom\u00e1n et al. (1999) for the sample obtained by heating bayerite at 1273 K, assuming a spheroid with depolarization parameters of (0.35, 0.003), produced an opacity with 11, 20, 28, and 32 \u03bcm features, so this component was included in the models. However, with only this component, the observed 20 \u03bcm features in the residual spectra were found to be wider than those in the models. By adding the opacity of the sample obtained by heating boehmite at 1173 K, the width of the 20 \u03bcm feature could be matched. This was done using the complex dielectric function for the sample obtained by heating boehmite at 1173 K from Pecharrom\u00e1n et al. (1999), assuming a spheroid with depolarization parameters of (0.35, 0.035). The complex dielectric functions of the samples obtained by heating bayerite and boehmite to various temperatures (Pecharrom\u00e1n et al. 1999) were derived by modeling the reflectance spectra of pellets obtained by pressing powders of these materials under great pressure. This method required Pecharrom\u00e1n et al. (1999) to assume an effective medium theory, such that a pellet is a mixture of one of their samples with a matrix of air. Pecharrom\u00e1n et al. (1999) noted that heating bayerite at 500\u00b0C eliminates the XRD pattern of bayerite, and they note that at 700\u00b0C, the infrared reflectance spectrum of the boehmite sample no longer shows OH\u2212 stretching bands. This must mean that the samples obtained from heating bayerite at 1273 K and from heating boehmite at 1173 K are no longer bayerite or boehmite, respectively. XRD performed by Pecharrom\u00e1n et al. (1999) of the sample of bayerite prepared at 1273 K suggests only \u03b8-alumina was present, and their infrared and NMR spectroscopy confirms this. XRD of their sample obtained from heating boehmite to 1173 K (Pecharrom\u00e1n et al. 1999) suggests \u03b4-alumina to be present, though some amounts of \u03b8-alumina and \u03b1-alumina are present, as they deduce from XRD and infrared and NMR spectroscopy.","Citation Text":["Pecharrom\u00e1n et al. 1999"],"Functions Text":["The complex dielectric functions of the samples obtained by heating bayerite and boehmite to various temperatures","were derived by modeling the reflectance spectra of pellets obtained by pressing powders of these materials under great pressure"],"Functions Label":["Uses","Uses"],"Citation Start End":[[891,914]],"Functions Start End":[[776,889],[916,1044]]}
{"Identifier":"2017MNRAS.471.4256V__King_et_al._2000_Instance_1","Paragraph":"Our ideas about which systems may survive spiral-in and produce WR X-ray binaries were triggered by the realization that the peculiar X-ray binary SS433 has avoided going into CE evolution, and that the donor star in this system is transferring mass to the compact star by stable Roche lobe overflow (RLOF; King & Begelman 1999; King, Taam & Begelman 2000). By analysing the properties of this system we realized that it is the high mass of the compact star in this system [4.3(\u00b10.8)\u2009\u2009M\u2299] in combination with the relatively low mass of its donor star [12.3(\u00b13.3)\u2009M\u2299] (Hillwig & Gies 2008), which allowed it to avoid CE evolution and enables it to gently spiral in without ever coalescing with its donor. As we consider SS433 to be a \u2018keystone\u2019 for understanding the formation of the WR X-ray binaries, we give in Section 2 a brief overview of its properties and evolutionary state. The avoidance of its going into CE evolution is \u2013 as we will argue \u2013 a consequence of the donor star having a radiative envelope (King et al. 2000), in combination with the donor star and accretor having a mass ratio less than 3.5. In Section 3, we then examine for which donor masses, mass ratios and orbital periods HMXBs will, when they start Roche lobe overflow, avoid going into CE evolution and may survive as WR X-ray binaries with short orbital periods. We also examine under which conditions they may still survive after having gone into CE evolution. In this section, some examples are given on how a number of well-known observed WR+O binaries with relatively short orbital periods are expected to evolve in the future, and are expected to produce WR X-ray binaries and, as a final evolutionary state, close double BHs. This model for producing double BHs is different from the ones proposed by Belczynski et al. (2016), Marchant et al. (2016) and de Mink & Mandel (2016). In Section 4, we attempt to estimate the birth rate of WR X-ray binaries in the Galaxy on the basis of our model, and find it to be still higher than observed and discuss possible ways to minimise this discrepancy. In Section 5, we discuss the results and estimate the possible birth rate of double BHs based on our model.","Citation Text":["King et al. 2000"],"Functions Text":["The avoidance of its going into CE evolution is \u2013 as we will argue \u2013 a consequence of the donor star having a radiative envelope"],"Functions Label":["Uses"],"Citation Start End":[[1012,1028]],"Functions Start End":[[882,1010]]}
{"Identifier":"2016AandA...586A..92P__Maciesiak_et_al._2012_Instance_1","Paragraph":"We calculated the dependence of the width of the profiles on the pulse period, considering the different frequencies separately. We note that the pulse width is not a direct reflection of the beam size or diameter (i.e. 2\u03c1, where \u03c1 is beam radius). For a visual representation of the geometry see for instance Maciesiak et al. (2011) and Bilous et al. (2014). In fact, only if the observer\u2019s line of sight cuts the emission centrally for magnetic inclination angles, \u03b1, that are not too small (i.e. \u03b1>  ~ 60\u00b0), w \u2248 2\u03c1. In such a case, when the emission beam is confined by dipolar open field lines, we would expect a P\u2212 1 \/ 2 dependence, which has indeed been observed when correcting for geometrical effects by transforming the pulse width into a beam radius measurement (see Rankin 1993; Gil & Krawczyk 1996; Maciesiak et al. 2012). For circular beams, profile width and beam radius are related by the relation first derived by Gil et al. (1984): (5)\\begin{equation} {\\rho_{10}} = 2\\sin^{-1} \\left[\\sin{\\alpha}\\sin{(\\alpha+\\beta)}\\sin^2\\left(\\frac{w_{10}}{4}\\right)+\\sin^2\\left(\\frac{\\beta}{2}\\right)\\right]^{1\/2}\\cdot \\end{equation}\u03c110=2sin-1sin\u03b1sin(\u03b1+\u03b2)sin2w104+sin2\u03b221\/2\u00b7The angle \u03b2 is the impact angle, measured at the fiducial phase, \u03c6, which describes the closest approach of our line of sight to the magnetic axis. This equation is derived under the assumption that the beam is symmetric relative to the fiducial phase. Typically, widths are measured at a certain intensity level (e.g. 50% or 10%, as here), and \u03c1 values are derived accordingly. In many cases, profiles are indeed often asymmetric relative to the chosen midpoint, or become so as they evolve with frequency. We note that for a central cut of the beam (\u03b2 = 0) and for an orthogonal rotator (\u03b1 ~ 90\u00b0) the equation reduces to \u03c1 = w\/ 2 as expected, while in a more general case, where \u03b2 = 0 and \u03b1 \u226b \u03c1 the relation reduces to \u03c1 = (w\/ 2)sin\u03b1. In principle, it is possible to determine \u03b1 and \u03b2 with polarisation measurements. However, in reality the duty cycle of the pulse is often too small to obtain reliable estimates (see Lorimer & Kramer 2004). Alternatively, at least for \u03b1, the relation reported by Rankin (1993) can be used: (6)\\begin{equation} \\label{eqn:w50} w_{50,{ \\rm core}}(1~{\\rm GHz}) = 2.45^\\circ \\cdot P^{-0.5\\pm 0.2}\/\\sin(\\alpha), \\end{equation}w50,core(1GHz)=2.45\u25e6\u00b7P\u22120.5\u00b10.2\/sin(\u03b1),calculated from the observed width dependency on period for the core components of pulsars (see Sect. 5.1), which is intrinsically related to the polar cap geometry. Equation (6) is valid at 1 GHz, but can be applied at LOFAR frequencies, maintaining the same dependence, if the impact angle \u03b2 \u226a \u03c1core; sin\u03b1 should be ignored for orthogonal rotators. Additionally, Rankin (1993), Gil et al. (1993), Kramer et al. (1994), Gould & Lyne (1998) suggested that \u201cparallel\u201d \u03c1 \u2212 P relations are found if the radio emission of the pulsar can be classified and separated into emission from \u201cinner\u201d and \u201couter\u201d cones, which seem to show different spectral properties (see Sect. 5.1 for details). ","Citation Text":["Maciesiak et al. 2012"],"Functions Text":["In such a case, when the emission beam is confined by dipolar open field lines, we would expect a P\u2212 1 \/ 2 dependence, which has indeed been observed when correcting for geometrical effects by transforming the pulse width into a beam radius measurement (see"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[811,832]],"Functions Start End":[[519,776]]}
{"Identifier":"2020ApJ...900..143S__Potgieter_et_al._2014_Instance_1","Paragraph":"It is well known that there exists an anticorrelation between the cyclic variations of GCR intensity and SSN (Forbush 1954), and the time lag in this relationship has also been reported by Forbush (1958). The time lag in solar modulation of GCRs should be closely related to the transport of solar wind plasma with its embedded HMF through the modulation region. This effect has been considered in numerical models of GCR modulation to study the long-term modulation mechanism (e.g., Le Roux & Potgieter 1990; Ferreira & Potgieter 2004; Bobik et al. 2012; Potgieter et al. 2014; Qin & Shen 2017). Moreover, the delay time in odd solar cycles was found to be longer than that in even ones (e.g., Nagashima & Morishita 1980; Mavromichalaki & Petropoulos 1984; Usoskin et al. 1998; Mavromichalaki et al. 2007; Ross & Chaplin 2019). Drift theory has predicted that positively charged particles mainly drift inward through the polar region of the heliosphere during the A > 0 solar magnetic cycle, while they predominantly participate in the heliosphere along the wavy heliospheric current sheet when A  0 (Jokipii et al. 1977; Jokipii & Thomas 1981). Thus, positively charged particles will take a longer time to reach Earth when A  0, and such modulation mechanism may cause a longer delay time in odd cycles (Van Allen 2000; Cliver & Ling 2001). In addition, the time lag varies with the energy\/rigidity of the particle, which is easy to understand: the higher the particle\u2019s energy is, the shorter its propagation time in the modulation region. Such phenomenon is implied from several theoretical (O\u2019Gallagher 1975; Strauss et al. 2011) and observational (Burger & Swanenburg 1973; Nymmik & Suslov 1995) studies. Meanwhile, the difference in the energy dependence of the time lag in odd and even solar cycles was also shown by Nymmik & Suslov (1995). However, in the correlation studies of Ross & Chaplin (2019), an obvious rigidity dependence on the delay time between SSN and GCR intensity recorded by NMs was not found. Thus, we focus on the GCR data from spacecraft rather than those from NMs in this work.","Citation Text":["Potgieter et al. 2014"],"Functions Text":["This effect has been considered in numerical models of GCR modulation to study the long-term modulation mechanism (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[556,577]],"Functions Start End":[[363,483]]}
{"Identifier":"2018AandA...619A..59S__Sander_et_al._2015_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":["Sander et al. 2015"],"Functions Text":["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","POWR"],"Functions Label":["Background","Background"],"Citation Start End":[[1988,2006]],"Functions Start End":[[1706,1927],[1960,1964]]}
{"Identifier":"2022MNRAS.516..167B__Steinborn_et_al._2015_Instance_1","Paragraph":"Otherwise (i.e. if there is insufficient mass left in the sub-grid gas reservoir), the mass deficit $m_\\mathrm{BH} + \\Delta m - m_\\mathrm{BH}^\\mathrm{dynamical}$ (the last term denotes the dynamical mass of the SMBH) is drawn from the surrounding gas particles. In eagle, this was done by stochastically swallowing individual gas neighbours. This is not an ideal approach: the momentum imparted on the SMBH from the swallowed gas particle may artificially dislodge it from its position, particularly without instantaneous repositioning (as also discussed by Steinborn et al. 2015). The mass of a gas particle is also typically much greater than \u0394m, so that the dynamical mass of SMBH particles remains systematically above its sub-grid mass. Both issues become more severe when individual gas particles have been enriched to masses well above their initial value due to stellar outflows, which is particularly common in massive, gas-poor galaxies. Instead of swallowing entire particles, we therefore transfer a (typically very) small fraction of mass from all gas neighbours to the SMBH simultaneously, with the mass \u03b4mi \u2018nibbled\u2019 from each neighbour i weighted in analogy to their contribution to the gas density at the location of the SMBH,\n(4)$$\\begin{eqnarray}\r\n\\delta m_i = (1-\\epsilon _\\mathrm{r}) \\Delta m \\left[\\frac{w_i m_i}{\\sum _j (w_j m_j)}\\right],\r\n\\end{eqnarray}$$where wi is the kernel weight of particle i, mi its mass, and the sum is over all neighbours; as above, the factor of (1 \u2212 \u03f5r) accounts for the mass converted to energy. In addition to mass, a fraction \u03b4mi\/mi of the momentum of neighbour i is also transferred to the SMBH. To prevent individual gas particles from becoming too light, we exclude any neighbour that would be reduced to less than half its initial mass and accept that the dynamical mass of the SMBH grows slightly less than desired in this case. In practice, we have found that this limit is never reached in the simulations presented here, because stars typically inject far more mass into gas particles than is drained by SMBHs.","Citation Text":["Steinborn et al. 2015"],"Functions Text":["This is not an ideal approach: the momentum imparted on the SMBH from the swallowed gas particle may artificially dislodge it from its position, particularly without instantaneous repositioning (as also discussed by"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[558,579]],"Functions Start End":[[342,557]]}
{"Identifier":"2017ApJ...834...20A__Temi_et_al._2007a_Instance_1","Paragraph":"Lenticular galaxies seems to have a wider range of properties compared to ellipticals that resemble more the old definition of ETGs. However, even in ellipticals, large differences prevail. Recent observations of elliptical galaxies with Spitzer and Herschel (Temi et al. 2005, 2007a, 2007b, 2009; Smith et al. 2012; Agius et al. 2013; Mathews et al. 2013) have revealed that the far-infrared (FIR) luminosity LFIR from these galaxies can vary by \u223c100 among galaxies with similar optical luminosity. The 70 \u03bcm band luminosities (from Temi et al. 2007a, 2009), is a good example of such a huge scatter in the FIR luminosity of elliptical galaxies. Some of the high L70 galaxies are members of a small subset of ellipticals having radio detections of neutral and molecular gas. A few others may be S0 galaxies which, because of their rotationally supported disks, often contain large masses of cold gas and dust. Ellipticals containing large excess masses of dust and cold gas probably result from significant galaxy mergers in the past. However a fraction of elliptical galaxies appear to be completely normal but L70 in these galaxies still ranges over a factor of \u223c30, far larger than can be explained by uncertainties in the estimate of the FIR spectral energy distribution (SED) due to local stellar mass loss. While a significant fraction of the cold gas mass in low- to intermediate-mass ETGs is thought to have an external, merger-related origin (e.g., Davis et al. 2011), in the most massive ETGs the cold gas phases are presumably generated internally (Davis et al. 2011; David et al. 2014; Werner et al. 2014). Mergers with gas- and dust-rich galaxies have often been suggested for the origin of dust in all elliptical galaxies (e.g., Forbes 1991). Although the merger explanation is almost certainly correct in some cases, mergers cannot explain most of the observed scatter in L70. A crucial element in our understanding of the evolution of galaxies toward ETGs is the mutual role played by the major merging of galaxies and the secular star formation quenching. Neither of these scenarios yet accounts for all the observational evidence, and one could assume both contributing to some extent.","Citation Text":["Temi et al.","2007a"],"Functions Text":["However, even in ellipticals, large differences prevail. Recent observations of elliptical galaxies with Spitzer and Herschel","have revealed that the far-infrared (FIR) luminosity LFIR from these galaxies can vary by \u223c100 among galaxies with similar optical luminosity."],"Functions Label":["Background","Background"],"Citation Start End":[[260,271],[278,283]],"Functions Start End":[[133,258],[357,499]]}
{"Identifier":"2020ApJ...905..111Z__Jiri\u010dka_et_al._2001_Instance_2","Paragraph":"Surveys of radio bursts in decimetric wavelengths is presented in papers by Isliker & Benz (1994) and Jiri\u010dka et al. (2001), within 1\u20133 GHz and 0.8\u20132.0 GHz frequency ranges, respectively. Some of these bursts are still not well understood. This is a case of the slowly positively drifting bursts (SPDBs). They appear in groups or as single bursts, with a duration of an individual burst from 1 to several seconds and their frequency drift is lower than about 100 MHz s\u22121 (Jiri\u010dka et al. 2001). The SPDBs seem to be similar to the reverse type III bursts (Aschwanden 2002) but their frequency drift is much smaller. The majority of observed SPDBs are connected to solar flares (Jiri\u010dka et al. 2001), and they appear many times at the very beginning of the flares (Benz & Simnett 1986; Kotr\u010d et al. 1999; Kaltman et al. 2000; Karlick\u00fd et al. 2018). Kaltman et al. (2000) reported on several SPDBs observed during three solar flares in the 0.8\u20132 GHz frequency range. They found frequency drifts of the observed SPDBs to be within the 20\u2013180 MHz s\u22121 range. Kotr\u010d et al. (1999) studied one of those flares. By combining the radio and spectral plus imaging H\u03b1 observations, they explained the observed SPDBs as radio emission generated by downwards propagating shock waves. Based on numerical simulations of the formation of thermal fronts in solar flares, Karlick\u00fd (2015) proposed that SPDBs observed in the 1\u20132 GHz range could be a signature of a thermal front. Furthermore, Karlick\u00fd et al. (2018) reported the observation of an SPDB (1.3\u20132.0 GHz) observed during the impulsive phase of an eruptive flare. They found time coincidence between the SPDB occurrence, an appearance of an ultraviolet (UV)\/EUV multithermal plasma blob moving down along the dark H\u03b1 loop at approximately 280 km s\u22121, and the observed change of H\u03b1 profile at the footpoint of that dark loop. Combining these observations they concluded that observed SPDB was likely generated by the thermal front formed in front of the falling EUV blob.","Citation Text":["Jiri\u010dka et al. 2001"],"Functions Text":["Some of these bursts are still not well understood. This is a case of the slowly positively drifting bursts (SPDBs). They appear in groups or as single bursts, with a duration of an individual burst from 1 to several seconds and their frequency drift is lower than about 100 MHz s\u22121"],"Functions Label":["Background"],"Citation Start End":[[472,491]],"Functions Start End":[[188,470]]}
{"Identifier":"2022MNRAS.513.4556Z__Swartz,_Wheeler_&_Harkness_1991_Instance_1","Paragraph":"Type II supernovae (SNe II) are believed to originate from the core collapse of massive stars with initial masses larger than 8\u2009M\u2299 (Heger et al. 2003). They are characterized by prominent P-Cygni profiles of Balmer series in their spectra (Filippenko 1997). Photometrically, they are classified as SNe IIP if their light curves show extended plateau features (\u223c100\u2009d), and SNe IIL if their light curves display post-peak linear declines (Barbon, Ciatti & Rosino 1979). Such a two-type classification is favoured by some statistical studies of SNe II (e.g. Faran et al. 2014a, b). However, Anderson et al. (2014) pointed out that if the sample is sufficiently large, one will find that SNe IIP and SNe IIL belong to a continuous distribution. Such a distribution is believed to be related to the mass of the hydrogen envelope maintained by the progenitor stars, as evidenced by the results that SNe having larger envelope masses produce light curves with shallower plateau slopes and longer plateau durations (Barbon et al. 1979; Swartz, Wheeler & Harkness 1991; Blinnikov & Bartunov 1993). In spite of the variety of envelope masses, the initial masses of the progenitor stars of SNe II are restricted to a certain range. Stellar evolution theory predicts an upper limit of 25\u2009M\u2299 for SN-II progenitors (Heger et al. 2003). Based on analyses of pre-explosion images (e.g. Smartt 2009), the progenitors of most SNe II are found to be red supergiants (RSGs), but their masses lie in a narrow range (i.e. \u223c9\u201317\u2009M\u2299). This inconsistency is known as the \u2018RSG problem\u2019 and can be partially explained by the \u2018failed SNe\u2019 theory (Lovegrove & Woosley 2013). As SNe II can also be used to determine distances, thus they are intriguing to the cosmology community. The measurement methods include the Expanding Photosphere Method (EPM; Kirshner & Kwan 1974; Hamuy 2001) and the Standard Candle Method (SCM; Hamuy & Pinto 2002). The basic idea of EPM is to derive the intrinsic luminosity from the photospheric radius and the temperature, and then compare the intrinsic luminosity with the apparent value to obtain the distance, while the SCM is based on the correlation between the expansion velocity and the luminosity at a specific epoch.","Citation Text":["Swartz, Wheeler & Harkness 1991"],"Functions Text":["Such a distribution is believed to be related to the mass of the hydrogen envelope maintained by the progenitor stars, as evidenced by the results that SNe having larger envelope masses produce light curves with shallower plateau slopes and longer plateau durations"],"Functions Label":["Motivation"],"Citation Start End":[[1029,1060]],"Functions Start End":[[742,1007]]}
{"Identifier":"2021ApJ...911...75M__Yamauchi_et_al._2004_Instance_1","Paragraph":"With the advent of the Parker Solar Probe (PSP) we are able to directly measure properties of the solar wind closer to the Sun than ever before (Fox et al. 2016). One of the early highlight discoveries of PSP is the omnipresence of strong local deflections of the magnetic field in the solar wind, mostly referred to as switchbacks (Kasper et al. 2019; Horbury et al. 2020). These structures, also called folds, jets, or spikes, are not necessarily full reversals of the local magnetic field, unlike some names suggest. In fact, the measured distribution of deflections resembles a power law, with most of the deflections at small angles relative to the Parker spiral (Dudok de Wit et al. 2020). There is evidence that switchbacks are localized kinks in the magnetic field and not polarity reversals or closed loops (McManus et al. 2020; Whittlesey et al. 2020). However, switchbacks are not a new finding. They have already been observed for several decades, by, e.g., Helios (e.g., Horbury et al. 2018), out to 1.3 au by Ulysses (e.g., Balogh et al. 1999; Neugebauer & Goldstein 2013). The novelty in the PSP observations is their sharpness and omnipresence (among other features; see Dudok de Wit et al. 2020), indicating that switchbacks are a more frequent feature closer to the Sun. This observation lends the possibility that switchbacks may originate lower in the solar atmosphere. The formation mechanism(s) of switchbacks, and whether they represent large-amplitude Alfv\u00e9n waves or structures advected by the solar wind, is as of yet not known. Several explanations about their origin have been put forth. Among these theories, interchange reconnection is the most focused upon (Yamauchi et al. 2004; Fisk 2005; Fisk & Kasper 2020). In this scenario, switchbacks are generated in the solar corona, and form at the reconnection sites between open and closed magnetic fluxes. In some studies, it is argued that the interchange reconnection results in magnetic flux ropes, which are ejected by the reconnection outflow, and are advected by the wind (Drake et al. 2020). In other studies, reconnection is thought to generate either Alfv\u00e9nic (He et al. 2020) or fast magnetosonic (Zank et al. 2020) wave pulses. Alternatively, as the distribution of deflections appears to be featureless and monotone in switchbacks (Dudok de Wit et al. 2020), they may not be a distinct feature but a manifestation of the ensuing turbulent dynamics in the solar wind (Squire et al. 2020). Thus, it is still not clear whether switchbacks originate in the lower solar atmosphere, or represent dynamic features of solar wind turbulence. Additionally, it is not clear whether they are wavelike perturbations, propagating at the Alfv\u00e9n or some other characteristic speed, or structures that are advected by the wind, such as flux ropes. However, some observations offer good constraints on the nature of switchbacks. Given that switchbacks are characterized by a strong Alfv\u00e9nic correlation of their velocity and magnetic field perturbations, a nearly constant magnetic pressure, and velocity enhancements along the propagation direction, an interpretation in terms of propagating nonlinear Alfv\u00e9n waves is a plausible scenario (Matteini et al. 2014). Alternatively, their localization in the perpendicular direction and the kink-like geometry may indicate their kink fast magnetoacoustic nature (e.g., Van Doorsselaere et al. 2008). An additional option is the association of the switchbacks with kink solitons, which keep their shape because of the balance between nonlinear and dispersive effects (see, e.g., Ruderman et al. 2010), i.e., a stationary nonlinear kink wave pulse. Other observations, such as a sharp rise in ion temperature at the boundaries of switchbacks are more compatible with an origin by reconnection (Farrell et al. 2020; Mozer et al. 2020).","Citation Text":["Yamauchi et al. 2004"],"Functions Text":["Several explanations about their origin have been put forth. Among these theories, interchange reconnection is the most focused upon","In this scenario, switchbacks are generated in the solar corona, and form at the reconnection sites between open and closed magnetic fluxes."],"Functions Label":["Background","Background"],"Citation Start End":[[1689,1709]],"Functions Start End":[[1555,1687],[1743,1883]]}
{"Identifier":"2021ApJ...917...59C__Schiller_et_al._2020_Instance_1","Paragraph":"With regard to Ni isotopes in the metal phase of iron meteorites, multiple studies have analyzed samples belonging to the same iron groups (e.g., IIAB, IIIAB, IVA, IVB) using broadly similar techniques (i.e., wet chemical separation of Ni followed by mass spectrometry). The initial studies reported no nucleosynthetic anomalies (Cook et al. 2006; Quitt\u00e9 et al. 2006). As analytical precision improved, anomalies were reported in some subsequent studies (e.g., Bizzarro et al. 2007; Regelous et al. 2008; Steele et al. 2011) but not in others (e.g., Dauphas et al. 2008; Chen et al. 2009), and no clear picture has emerged. Moreover, the interpretation of Ni isotopic data is hampered by the omission of data for 64Ni, the least abundant Ni isotope, from multiple studies. For Fe isotopes, initial studies pointed to a high degree of homogenization among various meteorite classes (e.g., Dauphas et al. 2008; Tang & Dauphas 2012). However, more recent studies (Cook & Sch\u00f6nb\u00e4chler 2017; Schiller et al. 2020) have revealed small non-mass-dependent variations in Fe isotopes. A related issue, also without consensus, concerns the magnitude of the initial abundance of the radionuclide 60Fe in the early solar system. A wide variety of sample types have been analyzed in an attempt to constrain the solar system initial 60Fe\/56Fe ratio including CAIs (e.g., Quitt\u00e9 et al. 2007), chondrules (e.g., Tachibana et al. 2006; Telus et al. 2018), chondritic metal (Cook et al. 2006), and differentiated meteorites, such as angrites and eucrites (e.g., Tang & Dauphas 2012, 2015). These studies have broadly used one of two analytical approaches: wet chemical separation of Ni followed by mass spectrometry or in situ analyses using ion microprobe (aka SIMS). Among these studies, estimates of 60Fe\/56Fe span more than a factor of 100 (e.g., Quitt\u00e9 et al. 2007; Tang & Dauphas 2012, 2015) from \u22482 \u00d7 10\u22126 to \u22481 \u00d7 10\u22128. Determining the initial 60Fe\/56Fe ratio in the protoplanetary disk is important for studies of early solar system chronology. Furthermore, radioactive decay of 60Fe may have been an important heat source on early-formed bodies, but this depends on its abundance (Kohman & Robison 1980; Yoshino et al. 2003).","Citation Text":["Schiller et al. 2020"],"Functions Text":["However, more recent studies","have revealed small non-mass-dependent variations in Fe isotopes."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[987,1007]],"Functions Start End":[[931,959],[1009,1074]]}
{"Identifier":"2017ApJ...837...89W__Henry_et_al._2013_Instance_1","Paragraph":"For over a decade, a tight correlation between metallicity and galaxy stellar mass (\n\n\n\n\n\n), i.e., the mass\u2013metallicity relation (MZR), has been quantitatively established, from the vast database of local galaxies observed by the Sloan Digital Sky Survey (SDSS; Tremonti et al. 2004; Zahid et al. 2012; Andrews & Martini 2013). This relation has been further extended to high redshifts, using deep near-infrared (IR) spectroscopy facilitated by large ground-based and space-based telescopes (Erb et al. 2006; Maiolino et al. 2008; Zahid et al. 2011; Henry et al. 2013; Steidel et al. 2014; Sanders et al. 2015; Guo et al. 2016). The measurements of the MZR as a function of redshift can cast useful constraints on various galaxy evolution models, since the slope of the MZR is sensitive to the properties of outflows, such as the mass loading factor and the outflow speed (see, e.g., Dav\u00e9 et al. 2012; Lu et al. 2015a). This slope can also be explained by variations of star-formation efficiency and gas mass fraction in galaxies with different stellar masses (see, e.g., Baldry et al. 2008; Zahid et al. 2014). The normalization of the MZR can shed light upon the stellar chemical yield across cosmic time (Finlator & Dav\u00e9 2008). Mannucci et al. (2010) first suggested that there exists a so-called fundamental metallicity relation (FMR) in the 3D parameter space spanned by \n\n\n\n\n\n, star-formation rate (SFR), and metallicity, such that the MZR is merely a 2D projection of this more fundamental 3D manifold (see also Hunt et al. 2016). This 3D scaling relation shows a tight scatter (\u223c0.05 dex) in metallicity and is speculated to not evolve with z. In this context, the apparent redshift evolution of the MZR normalization originates primarily from sampling the FMR in terms of galaxies with different SFR. This concept of the FMR is in accord with the gas regulator model proposed by Lilly et al. (2013), even though mergers can also play a subtle role in shaping the form of the FMR by increasing the scatter (Michel-Dansac et al. 2008). However, at high redshifts, the validity of the FMR is still under investigation (see, e.g., Wuyts et al. 2014; Sanders et al. 2015).","Citation Text":["Henry et al. 2013"],"Functions Text":["This relation has been further extended to high redshifts, using deep near-infrared (IR) spectroscopy facilitated by large ground-based and space-based telescopes"],"Functions Label":["Background"],"Citation Start End":[[550,567]],"Functions Start End":[[328,490]]}
{"Identifier":"2016ApJ...817..152X__Brien_et_al._2006_Instance_1","Paragraph":"Afterglows of gamma-ray bursts (GRBs) are generally believed to be produced by a relativistic jet interacting with the surrounding medium (e.g., M\u00e9sz\u00e1ros & Rees 1997; Sari et al. 1998; Kumar & Zhang 2015). Very early multi-wavelength afterglows are critical to reveal the properties of the radiating fireball and its environment as well as the central engine of GRBs. Using the prompt slewing and precise locating capacity of the X-ray telescope (XRT) on board the Swift mission, very early X-ray and optical afterglows were obtained with the XRT and ground-based optical telescopes. The early X-ray afterglows are usually dominated by erratic X-ray flares and the tail emission of prompt gamma-rays are due to the arrival time delay of photons in the high latitude of the radiating fireball (Liang et al. 2006; O\u2019Brien et al. 2006; Zhang et al. 2006, 2007a). The flares and prompt tail emission are usually not seen in the early optical afterglow data (Li et al. 2012; Wang et al. 2013). About one-third of well-sampled optical afterglow light curves show a clear smooth bump (Li et al. 2012). It may be attributed to deceleration of the fireball by the ambient medium (see Rees & M\u00e9sz\u00e1ros 1992; M\u00e9sz\u00e1ros & Rees 1993 and Sari & Piran 1999 for the thin shell case; Kobayashi et al. 1999 and Kobayashi & Zhang 2007 for the thick shell case). This may give the most robust estimate for the initial Lorentz factor (\n\n\n\n\n\n) of the fireball since the deceleration time weakly depends on other model parameters than \n\n\n\n\n\n (e.g., Molinari et al. 2007; Melandri et al. 2010; Xin et al. 2015). With a sample of GRBs with a detected optical afterglow onset bump, Liang et al. (2010) derived their \n\n\n\n\n\n and found a tight \n\n\n\n\n\n relation (see also a \n\n\n\n\n\n relation in Lu et al. 2012). Furthermore, by incorporating the peak energy of the \n\n\n\n\n\n spectrum in the burst frame (\n\n\n\n\n\n) into the \n\n\n\n\n\n relation, a tighter \n\n\n\n\n\n relation is found (Liang et al. 2015). This relation places strong constraints on the composition of GRB fireballs.","Citation Text":["O\u2019Brien et al. 2006"],"Functions Text":["The early X-ray afterglows are usually dominated by erratic X-ray flares and the tail emission of prompt gamma-rays are due to the arrival time delay of photons in the high latitude of the radiating fireball"],"Functions Label":["Background"],"Citation Start End":[[812,831]],"Functions Start End":[[584,791]]}
{"Identifier":"2022ApJ...926..151Z__Saxena_et_al._2020_Instance_1","Paragraph":"Unlike the CMB, the 21 cm signal is highly non-Gaussian, because patchy, bubble-like structures of ionized hydrogen (H ii) regions are produced surrounding the ionizing sources. Thus, there is potentially a wealth of information in the 21 cm signal that is not contained in the 21 cm power spectrum, a two-point statistics of 21 cm brightness temperature fluctuations that is traditionally well studied in the literature. It is therefore essential to develop new methods that maximally exploit the full information in the 3D 21 cm images obtained by the SKA. Toward this goal, conventionally, new summary statistics that can only be measured with imaging have been proposed. These include the three-point correlation function (Hoffmann et al. 2019; Jennings et al. 2020), bispectrum (Yoshiura et al. 2015; Shimabukuro et al. 2016, 2017; Majumdar et al. 2018, 2020; Hutter et al. 2020; Saxena et al. 2020; Kamran et al. 2021), one-point statistics (Harker et al. 2009; Shimabukuro et al. 2015; Gorce et al. 2021), topological quantities such as the Minkowski functionals (Gleser et al. 2006; Chen et al. 2019; Kapahtia et al. 2021) and Betti numbers (Giri & Mellema 2021), the cross correlation between the 21 cm line and other probes, such as the CO line (Gong et al. 2011; Lidz et al. 2011), the C ii line (Gong et al. 2012; Beane & Lidz 2018), the kinetic Sunyaev\u2013Zel\u2019dovich (kSZ) effect (Ma et al. 2018; La Plante et al. 2020), and novel techniques such as the antisymmetric cross correlation between the 21 cm line and CO line (Zhou et al. 2021). Since those summary statistics are fully determined by the parameters in the reionization models (hereafter \u201creionization parameters\u201d), in principle, Monte Carlo Markov Chain (MCMC) methods can be employed to constrain the reionization parameters from measurements of those statistics with futuristic 21 cm experiments (see, e.g., Watkinson et al. 2022), just as the MCMC analysis with the 21 cm power spectrum (Greig & Mesinger 2015, 2017, 2018).","Citation Text":["Saxena et al. 2020"],"Functions Text":["Toward this goal, conventionally, new summary statistics that can only be measured with imaging have been proposed. These include","bispectrum"],"Functions Label":["Background","Background"],"Citation Start End":[[885,903]],"Functions Start End":[[559,688],[772,782]]}
{"Identifier":"2020MNRAS.499.1788W__Wolfire_et_al._2003_Instance_1","Paragraph":"Owing to their brightness at rest-frame FIR wavelengths, the ionized and neutral species of Carbon, Nitrogen, and Oxygen are powerful diagnostic lines for tracing the ISM of nearby and distant galaxies. When combined with photodissociation region (PDR) models (Tielens & Hollenbach 1985), measurement of the emission from different lines provide a means to constrain quantities such as the ionization rate and metallicity of the ISM of galaxies. Multilevel FIR transition lines have now been widely surveyed in local galaxies, originally by ISO (e.g. Luhman et al. 1998; Malhotra et al. 2001) and more recently using the PACS spectrometer (Poglitsch et al. 2010) on the ESA Herschel Space Observatory (Pilbratt et al. 2010). The [C\u2009ii]158\u2009\u03bcm line is typically the brightest in star-forming galaxies, arising from ionized, and even neutral gas where it is the main coolant (Wolfire et al. 2003). Another commonly observed FIR line tracer of the ionized ISM is [N\u2009ii]122\u2009\u03bcm. It has the advantage that it can be found associated with lower excitation gas, close to that observed in our own Galaxy (e.g. Goldsmith et al. 2015; Herrera-Camus et al. 2016). Another major coolant of the ISM is [O\u2009i]63\u2009\u03bcm (Wolfire et al. 2003). Owing to its high excitation temperature and critical density, it can dominate the cooling in regions of starburst activity (Kaufman et al. 1999; Kaufman, Wolfire & Hollenbach 2006; Brauher, Dale & Helou 2008; Croxall et al. 2012; Narayanan & Krumholz 2017). When combined with measurements of the [C\u2009ii]158\u2009\u03bcm line intensity and FIR luminosity, the [O\u2009i]63\u2009\u03bcm line intensity can constrain the FUV field, G, and the gas density using PDR models. The luminosity in these FIR lines generally exhibits a deficit in the most FIR luminous galaxies compared to the trend expected from lower luminosity galaxies (e.g. Malhotra et al. 2001; Graci\u00e1-Carpio et al. 2011; D\u00edaz-Santos et al. 2017). This has made the emission from lines like [O\u2009i]63\u2009\u03bcm more challenging to detect at high-redshifts.","Citation Text":["Wolfire et al. 2003"],"Functions Text":["The [C\u2009ii]158\u2009\u03bcm line is typically the brightest in star-forming galaxies, arising from ionized, and even neutral gas where it is the main coolant"],"Functions Label":["Background"],"Citation Start End":[[873,892]],"Functions Start End":[[725,871]]}
{"Identifier":"2021MNRAS.508.2823C__Baumgardt,_Makino_&_Ebisuzaki_2004_Instance_1","Paragraph":"Globular clusters are of particular dynamical interest because of the phenomena of core collapse and binary burning. Stellar interactions in dense globular clusters tend to drive the central core towards a collapse to a state of extremely small radius and high density. Primordial binary populations may delay the onset of this core collapse by serving as dynamical energy sources, but eventually become depleted by various destruction mechanisms (e.g. Verbunt & Freire 2014), allowing core collapse to proceed (e.g. Fregeau et al. 2003, and references therein). Several studies suggest that black hole binaries, rather than main-sequence (MS) binaries, are critical for delaying core collapse (e.g. Breen & Heggie 2013; Wang et al. 2016; Kremer et al. 2019). In addition, an intermediate-mass black hole (IMBH) in a cluster core will act as a strong central energy source able to delay, or even prevent, core collapse (e.g. Baumgardt, Makino & Ebisuzaki 2004; Gill et al. 2008; L\u00fctzgendorf, Baumgardt & Kruijssen 2013). In any case, after an initial collapse, the cluster core will undergo gravothermal oscillations of expansion and contraction, as first demonstrated by Sugimoto & Bettwieser (1983), during which the core radius remains quite small. About 20\u201325 globular clusters, including NGC 6397 and NGC 6752, have very compact (rc \u2272 10\u2009arcsec), high-density cores that appear to be post-collapse. The post-collapse oscillations of these cores should produce episodic bursts of strongly enhanced dynamical binary formation and ejection during the densest phases, as noted by Lugger et al. (2007). Their reasoning is based on the scaling of the encounter rate \u0393, which is given by the integral of \u03c12\/v over the cluster volume, where \u03c1 is spatial mass density and v is the velocity dispersion (Verbunt & Hut 1987; Bahramian et al. 2013). The encounter rate can be approximated by the simplified expression $\\Gamma \\propto \\rho _0^2 r_c^3\/v_0$ (Pooley et al. 2003), where \u03c10 is central density, rc is the core radius, and v0 is the central velocity dispersion. This results in $\\Gamma \\propto r_c^{-1.4}$ for a simple homologous model for core collapse, in which $\\rho _0 \\propto r_c^{-2.2}$ and $v_0 \\propto r_c^{-0.05}$ (Cohn 1980).","Citation Text":["Baumgardt, Makino & Ebisuzaki 2004"],"Functions Text":["In addition, an intermediate-mass black hole (IMBH) in a cluster core will act as a strong central energy source able to delay, or even prevent, core collapse (e.g."],"Functions Label":["Background"],"Citation Start End":[[925,959]],"Functions Start End":[[760,924]]}
{"Identifier":"2022MNRAS.512.3137Z__Katz_et_al._1999_Instance_3","Paragraph":"However, it is not straightforward to explain H2 formation in astronomical sources even when the catalytic roles of dust grains are introduced into models. Interstellar species are believed to be formed on cold grain surfaces via the so called Langmuir\u2013Hinshelwood mechanism (Watson & Salpeter 1972; Pickles & Williams 1977; Hasegawa, Herbst & Leung 1992). To form H2, H atoms accrete on dust grains and then bind weakly with surfaces, which is known as physisorption. They can overcome the diffusion barrier and move on the grain surfaces via quantum tunnelling or thermal hopping. However, laboratory and theorectical studies showed that quantum tunnelling does not contribute much to the mobility of H atoms on silicate, carbonaceous or solid amorphous water (ASW) surfaces (Pirronello et al. 1997, 1999; Katz et al. 1999; Nyman 2021). If H atoms encounter other H atoms, then H2 molecules are formed. But H atoms can also desorb and leave grain surfaces. A hydrogen atom must reside on a grain long enough to find a partner H atom to form H2. As the dust temperature increases, the H atom desorption and diffusion rates also increase. So the temperature of grain surfaces must be sufficiently low so that an H atom can encounter another one before it desorbs. On the other hand, the temperature of grain surfaces must be high enough so that H atoms can diffuse on the grain surface. The parameter that measures how strongly species are to bound to grain surfaces is called desorption energy. It was found that if we assume a single H desorption energy in models, the dust temperature range over which efficient H2 formation occurs is narrow (6\u201310 K for olivine grains) (Katz et al. 1999). Moreover, it was found that the highest dust temperature at which H2 can be formed efficiently is less than 17\u2009K (Katz et al. 1999). However, the grain surface temperature in the unshielded diffuse clouds, where hydrogen molecules are believed to be efficiently formed, is around 20 K (Li & Draine 2001).","Citation Text":["Katz et al. 1999"],"Functions Text":["Moreover, it was found that the highest dust temperature at which H2 can be formed efficiently is less than 17\u2009K"],"Functions Label":["Background"],"Citation Start End":[[1807,1823]],"Functions Start End":[[1693,1805]]}
{"Identifier":"2016ApJ...817...12P__Chamandy_et_al._2014_Instance_2","Paragraph":"Large-scale magnetic fields with strength of the order of 1\u201310 \u03bcG have been observed in disk galaxies (e.g., Beck et al. 1996; Fletcher 2010; Beck 2012; Beck & Wielebinski 2013; Van Eck et al. 2015). The origin of these fields can be explained through mean-field dynamo theory (Ruzmaikin et al. 1988; Beck et al. 1996; Brandenburg & Subramanian 2005a; Kulsrud & Zweibel 2008). The conservation of magnetic helicity is one of the key constraints in these models, and also leads to the suppression of the \u03b1-effect. The operation of the mean-field dynamo automatically leads to the growth of magnetic helicity of opposite signs between the large-scale and small-scale magnetic fields (Pouquet et al. 1976; Gruzinov & Diamond 1994; Blackman & Field 2002). To avoid catastrophic suppression of the dynamo action (\u03b1-quenching), the magnetic helicity due to the small-scale magnetic field should be removed from the system (Blackman & Field 2000, 2001; Kleeorin et al. 2000). Mechanisms suggested to produce these small-scale magnetic helicity fluxes are: advection of magnetic fields by an outflow from the disk through the galactic fountain or wind (Shukurov et al. 2006; Sur et al. 2007; Chamandy et al. 2014), magnetic helicity flux from anisotropy of the turbulence produced by differential rotation (Vishniac & Cho 2001; Subramanian & Brandenburg 2004, 2006; Sur et al. 2007; Vishniac & Shapovalov 2014), and through diffusive flux (Kleeorin et al. 2000, 2002; Brandenburg et al. 2009; Mitra et al. 2010; Chamandy et al. 2014). The outflow of magnetic helicity from the disk through dynamo operation leads to the formation of a corona (Blackman & Field 2000). According to Taylor's hypothesis, an infinitely conducting corona would resistively relax to force-free field configurations under the constraint of global magnetic helicity conservation (Woltjer 1960; Taylor 1974; Finn & Antonsen 1983; Berger & Field 1984; Mangalam & Krishan 2000). In this paper, we include advective and diffusive fluxes in a simple semi-analytic model of a galactic dynamo that transfers magnetic helicity outside the disk and consequently builds up a force-free corona in course of time. We first solve the time-dependent dynamo equations by expressing them as separable in variables r and z. The radial part of the dynamo equation is solved using an eigenvector expansion constructed using the steady-state solutions of the dynamo equation. The eigenvalues of the z part of the solution are obtained by solving a fourth-order algebraic equation, which primarily depends upon the turbulence parameters and the magnetic helicity fluxes. Once the dynamo solutions are written out as parametric functions of these parameters, the evolution of the mean magnetic field is computed numerically by simultaneously solving the dynamical equations for \u03b1-quenching and the growth of large-scale coronal magnetic helicity. Since the large-scale magnetic field lines cross the boundary between the galactic disk and the corona, the magnetic helicity of the large-scale magnetic field in the disk volume is not well defined. Hence we use the concept of gauge-invariant relative helicity (Finn & Antonsen 1983; Berger & Field 1984; Berger 1985) to estimate the large-scale magnetic helicity in the disk and the corona. Here the gauge-invariant relative helicity for the cylindrical geometry is calculated using the prescription given in Low (2006, 2011). We then investigate the dependence of the saturated mean magnetic field strength and its geometry on the magnetic helicity fluxes within the disk and the corresponding evolution of the force-free field in the corona.","Citation Text":["Chamandy et al. 2014"],"Functions Text":["and through diffusive flux"],"Functions Label":["Background"],"Citation Start End":[[1504,1524]],"Functions Start End":[[1404,1430]]}
{"Identifier":"2017AandA...605A.102C__Konstantinova-Antova_et_al._2013_Instance_1","Paragraph":"Magnetic fields are actively searched for at the surface of all kinds of stars throughout the Hertzsprung-Russell diagram (HRD), as they probably impact stellar evolution from birth to death in various ways (see e.g., Donati & Landstreet 2009 for a general review; and Wade et al. 2016; Alecian et al. 2013; Vidotto et al. 2014; Folsom et al. 2016 for recent spectropolarimetric surveys on massive, intermediate-mass, and low-mass stars). Recently, magnetic fields have been unambiguously detected via Zeeman signatures in a large sample of single G-K giants observed with the spectropolatimeters TBL\/Narval and CFHT\/ESPaDOnS (Konstantinova-Antova et al. 2013; Auri\u00e8re et al. 2015; Borisova et al. 2016; Tsvetkova et al. 2017). Interestingly, the cool intermediate-mass evolved stars with surface magnetic fields are found to cluster in certain regions of the HRD that correspond to precise moments of their evolution where convective envelopes make up a significant fraction of the stellar mass. These observations are of particular importance for stellar-evolution modeling, since magnetic fields may play a crucial role in the angular momentum evolution of different types of stars. For instance, it is well established that magnetic fields affect the rotation rate of solar-type stars through the torque applied to stellar surfaces by magnetically coupled stellar winds (e.g., Schatzman 1962; Matt et al. 2015; R\u00e9ville et al. 2015; Amard et al. 2016). However, the efficiency and the impact of magnetic braking for evolved stars are not well studied yet. In addition, because of the angular momentum they transport (e.g., Spruit 1999; Mathis & Zahn 2005), magnetic fields may play an important role in explaining the properties of core and surface rotation in giants as seen by asteroseismology (e.g., Beck et al. 2012; Mosser et al. 2012b,a; Deheuvels et al. 2012, 2014; Cantiello et al. 2014; Di Mauro et al. 2016), as well as the rotation rate of white dwarf remnants (Suijs et al. 2008). By comparing the observed rotational and magnetic properties of stars at different evolutionary phases with the predictions of rotating stellar models, one may thus obtain strong constraints on the input physics of the stellar models. Conversely, it is important to know whether the possible presence and the global properties of the magnetic field of a given star can be anticipated from its position in the Hertzsprung-Russell diagram (see Gregory et al. 2012 for the case of pre-main sequence stars). ","Citation Text":["Konstantinova-Antova et al. 2013"],"Functions Text":["Recently, magnetic fields have been unambiguously detected via Zeeman signatures in a large sample of single G-K giants observed with the spectropolatimeters TBL\/Narval and CFHT\/ESPaDOnS"],"Functions Label":["Background"],"Citation Start End":[[627,659]],"Functions Start End":[[439,625]]}
{"Identifier":"2019MNRAS.488.5748W__Mao_&_Schneider_1998_Instance_1","Paragraph":"The search for observational evidence of dark matter substructure in galaxies is on-going, as traditional methods for in-directly detecting dark matter sub-haloes (e.g. modelling tidal streams and gravitational lensing) have yet to agree on either the Milky Way\u2019s or an external galaxies\u2019 current substructure composition. Gravitational lensing allows for constraints to be placed on the dark matter substructure content of external galaxies, as orbiting sub-haloes will result in anomalies in strong gravitational lenses (Mao & Schneider 1998; Dalal & Kochanek 2002; Vegetti et al. 2012). The disruption of stellar streams by dark matter sub-haloes in the Milky Way has also been well studied (e.g. Johnston, Spergel & Haydn 2002; Ibata et al. 2002; Carlberg 2009; Erkal et al. 2016; Sanders, Bovy & Erkal 2016; Bovy, Erkal & Sanders 2017; Carlberg 2017), with gaps in stellar streams believed to be a tell-tale sign that a stream has recently encountered a sub-halo. Hence the presence of gaps in a tidal stream, or lack thereof, can be used to constrain the dark-matter substructure properties of the Milky Way (Yoon, Johnston & Hogg 2011; Carlberg 2012; Erkal & Belokurov 2015a,b; Bovy 2016; Carlberg 2016; Banik et al. 2018; Bonaca et al. 2019). However gaps in tidal streams, as well as overdensities and asymmetry, can also be produced as stars are tidally stripped from a star cluster along its orbit (K\u00fcpper et al. 2010), disc shocking (Odenkirchen et al. 2003), spiral arms (Dehnen et al. 2004), a tri-axial halo (K\u00fcpper et al. 2015), the Galactic bar (Pearson, Price-Whelan & Johnston 2017), interactions with giant molecular clouds (Amorisco et al. 2016), the stream\u2019s progenitor cluster (Webb & Bovy 2018), and complex dynamical histories (i.e a time-dependent tidal field due to galaxy growth via mergers or accretion) (Carlberg 2018). Banik & Bovy (2018) recently studied these effects in detail for the Pal 5 stream and found that the bar and molecular clouds can each individually explain the observed structure of the Pal 5 stream. Hence new methods are required in order to help search for dark matter substructure in the Milky Way and constrain its properties.","Citation Text":["Mao & Schneider 1998"],"Functions Text":["Gravitational lensing allows for constraints to be placed on the dark matter substructure content of external galaxies, as orbiting sub-haloes will result in anomalies in strong gravitational lenses"],"Functions Label":["Motivation"],"Citation Start End":[[523,543]],"Functions Start End":[[323,521]]}
{"Identifier":"2021MNRAS.506.1978L__Krivov_et_al._2018_Instance_1","Paragraph":"This dynamical limit on planetesimal sizes could suggest that these may exist in the belt a factor of 103 larger than the ${\\sim }1\\,$km sized lower limit predicted by collisional replenishment. However, if the assumed \u03b1 = 3.5 Dohnanyi (1969) size distribution continued up to planetesimals of this size, then the disc mass would be $M_{\\rm {disc}}{\\sim }220\\, M_{\\oplus }$, a factor of \u223c3 higher than the prediction of Sch\u00fcppler et al. (2016) and over an order of magnitude larger than the lower limit derived from the collisional lifetime and age of the system. This would not violate our earlier calculation since these \u226b1\u2009km planetesimals could be abundant in the disc without having collided within the age of the system. Whilst this higher total disc mass is still consistent with the dust mass measurements of protoplanetary discs (for example, see Andrews & Williams 2005; Ansdell et al. 2016; Cieza et al. 2019), high debris disc masses, i.e. those in the range of $100\\!-\\!1000\\, M_{\\oplus }$, become problematic since these would require a very high efficiency of primordial dust being incorporated into these larger planetesimals (an example of the so-called \u2018disc mass problem\u2019, see Krivov et al. 2018; Krivov & Wyatt 2021). Nevertheless, it might still be possible to explain the observed level of stirring by embedded bodies while retaining a lower disc mass (i.e. if the largest bodies are less frequent than the \u03b1 = 3.5 size distribution would predict). If instead the size distribution had a slope of \u03b1 = 3.7, then even with these 1200\u2009km bodies, the total disc mass estimate would be reduced from ${\\sim }220\\, M_\\oplus$ to ${\\sim }6\\, M_\\oplus$, since the total number of these would be greatly reduced. On the other hand, the size distribution could be truncated, being much steeper for planetesimals larger than a few kilometres, and shallower in other regions. We estimate the effect that this can have on the derived disc mass using equation (9) of Krivov & Wyatt (2021), and find that the total mass can be reduced by an order of magnitude from ${\\sim }220\\, M_\\oplus$ to ${\\sim }17\\, M_\\oplus$, if based on a triple power law size distribution, with qmed = 4, qbig = 3 and q = 3.5. Alternatively the measured vertical scale height could be due to other dynamical interactions (e.g. stirring by a planet internal or external to the belt, or a recent stellar fly-by) or even be a remnant of the primordial disc (e.g. if this disc was born stirred, Booth & Clarke 2016). In summary, whilst q1 Eri at an age of ${\\sim }1.4\\, \\rm {Gyr}$ is an outlier in terms of its brightness, it need not be an outlier in terms of its disc mass unless many much larger planetesimals are present.","Citation Text":["Krivov et al. 2018"],"Functions Text":["Whilst this higher total disc mass is still consistent with the dust mass measurements of protoplanetary discs","high debris disc masses, i.e. those in the range of $100\\!-\\!1000\\, M_{\\oplus }$, become problematic since these would require a very high efficiency of primordial dust being incorporated into these larger planetesimals (an example of the so-called \u2018disc mass problem\u2019, see"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1196,1214]],"Functions Start End":[[727,837],[922,1195]]}
{"Identifier":"2015AandA...576A...5C__J\u00f8rgensen_et_al._2012_Instance_1","Paragraph":"The relative abundances of the three species are derived from the column densities in Table 2 and are compared with other star-forming regions and comets in Table 3. The (CH2OH)2\/CH2OHCHO abundance ratio of ~0.3\u20130.5 previously derived in IRAS 16293 by J\u00f8rgensen et al. (2012) was revised. Indeed, the assignment in J\u00f8rgensen et al. (2012) was based on only one line of the gGg\u2032 conformer of ethylene glycol about 200 cm-1 (~290 K, M\u00fcller & Christen 2004) above the lowest-energy aGg\u2032 conformer \u2013 and thus tentative. An analysis from observations of six transitions of the lower energy conformer from ALMA Cycle 1 observations at 3 mm (four spectral windows at 89.48\u201389.73, 92.77\u201393.03, 102.48\u2013102.73, and 103.18\u2013103.42 GHz; J\u00f8rgensen et al., in prep.) results in a higher ethylene glycol-to-glycolaldehyde abundance ratio of 1.0\u2009\u00b1\u20090.3. This new estimate is consistent with the ratio expected between the aGg\u2032 and gGg\u2032 conformers under thermal equilibrium conditions at 300\u2009K, the excitation temperature of glycolaldehyde derived in IRAS 16293 (J\u00f8rgensen et al. 2012). The (CH2OH)2\/CH2OHCHO abundance ratio in IRAS2A is estimated at 5.5 \u00b1 1.0 if we consider the column densities derived from the rotational diagrams. It is slightly lower (4.6), however, if we use the column density of ethylene glycol of 1.1 \u00d7 1016 cm-2 that does not overproduce the peak intensities of a few lines (see Fig. 3). The (CH2OH)2\/CH2OHCHO abundance ratio consequently is a factor ~5 higher than in the Class 0 protostar IRAS 16293. It is also higher than in the other star-forming regions (see Table 3), but similar to the lower limits derived in comets (\u22733\u20136). This indicates that the glycolaldehyde chemistry may in general vary among hot corinos. It is possible that like IRAS2A, other very young low-mass protostars show high (CH2OH)2\/CH2OHCHO abundance ratios, in agreement with the cometary values. The CH3OCHO\/CH2OHCHO column density ratio found in IRAS2A (~20) ranges between the values derived in the molecular clouds from the Galactic center (~3.3\u20135.2) and the high-mass star-forming regions (~40\u201352). A lower limit of 2 was derived for comet Hale-Bopp. ","Citation Text":["J\u00f8rgensen et al. (2012)"],"Functions Text":["The (CH2OH)2\/CH2OHCHO abundance ratio of ~0.3\u20130.5 previously derived in IRAS 16293 by","was revised."],"Functions Label":["Differences","Differences"],"Citation Start End":[[252,275]],"Functions Start End":[[166,251],[276,288]]}
{"Identifier":"2015ApJ...799...99K__Coelho_&_Gadotti_2011_Instance_2","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"],"Functions Text":["Some studies find weak statistical links among AGN activity and the presence of bars (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[1590,1611]],"Functions Start End":[[1414,1505]]}
{"Identifier":"2018MNRAS.479.4073L__Chatterjee_et_al._2008_Instance_1","Paragraph":"As can be seen from Fig. 1, in both the optical and the NIR band, several predominant flares appearing in the whole period considered are usually surrounded by other smaller subflares. In addition, it seems that the NIR light curves follow a similar flaring activity, lasting for several days or weeks, to the optical bands, although the NIR flux is more variable compared than the optical one, increasing by about 7 times in the J band, from 9.68 to 67.85 mJy (for an example, see Table 1). An exponential rise\/decay model can appropriately fit the blazar light curve, which generally consists of superpositions of a steady background flux density and a number of flare components caused by events in the jet or the accretion disc\/corona region (e.g. Valtaoja et al. 1999; Chatterjee et al. 2008; Abdo et al. 2010b; Chatterjee et al. 2012; Guo et al. 2016; Li et al. 2017, 2018, and references therein). Following Abdo et al. (2010b), we define the multipeaked model to analyse the temporal flare profiles during the interval of interest:\n(3)\r\n\\begin{eqnarray*}\r\nF(t) = {F_{\\rm c}} + \\sum _{i}\\left[{F_i}{\\left({{{\\rm e}^{\\frac{{{T_{0,i}} - t}}{{{T_{r,i}}}}}} + {{\\rm e}^{\\frac{{{\\rm t} - {T_{0,i}}}}{{{T_{{\\rm d},i}}}}}}} \\right)^{ - 1}}\\right],\r\n\\end{eqnarray*}\r\nwhere Fc is the assumed constant background flux level underlying the flare components and is constrained to be less than or equal to the lowest value of flux, and the sum runs over all of the individual flares being fitted. For any one of the individual flares, Fi is the variability amplitude, T0, i is the epoch of the peak, and Tr, i and Td, i are the time-scales of rise and decay, respectively. Equation (3) represents flares within a limited time interval, because the observational gaps exist regularly or occasionally. In order to ensure the accuracy of the fitting model, the divided observational segments in different bands mentioned in Section 2 were used to analyse the temporal flare profiles. Each segment of the light curves was first interpolated by means of cubic spline interpolations through the 1-d binned light curves, then smoothed by a five-point moving-average filter to reduce the impact of the short-term variabilities.","Citation Text":["Chatterjee et al. 2008"],"Functions Text":["An exponential rise\/decay model can appropriately fit the blazar light curve, which generally consists of superpositions of a steady background flux density and a number of flare components caused by events in the jet or the accretion disc\/corona region (e.g."],"Functions Label":["Uses"],"Citation Start End":[[774,796]],"Functions Start End":[[492,751]]}
{"Identifier":"2022MNRAS.511.4946N__Vivek,_Srianand_&_Gupta_2016_Instance_1","Paragraph":"Outflows from the central regions of active galactic nuclei (AGNs) are thought to be the main agents that regulate the evolution of both the central supermassive black holes as well as the host galaxies (Silk & Rees 1998; Di Matteo, Springel & Hernquist 2005). The presence of high-velocity outflows from AGNs can be established from the evidence provided by the blueshifted broad absorption lines (BALs) seen in the spectra of 10\u201320 per\u2009cent of quasi-stellar objects (QSOs)1 (Weymann et al. 1991). BALQSOs are classified into three sub-classes based on the ionization state of the absorbing gas: (a) high-ionization BALQSOs (HiBALs) consist of absorption from high-ionization lines such as C iv, Si iv, and N v (Gibson et al. 2008; Filiz Ak et al. 2013; Vivek, Srianand & Gupta 2016; Vivek 2019; Mishra et al. 2021) (b) low-ionization BALQSOs (LoBALs) show absorption from low-ionization lines such as Mg ii and Al iii along with the high-ionization lines (Voit, Weymann & Korista 1993; Vivek et al. 2014; Vivek, Srianand & Dawson 2018; Yi et al. 2019), and (c) iron-LoBALs (FeLoBALs) are LoBALs with excited fine-structure Fe ii and\/or Fe iii absorption lines (Vivek et al. 2012; McGraw et al. 2015). The observed BALQSO fraction is explained either by an orientation model, where the line of sight intersects with the BAL absorbing clouds in 10\u201320 per\u2009cent QSOs (Weymann et al. 1991; Elvis 2000), or an evolutionary model, where the QSO spends 10\u201320 per\u2009cent of its lifetime in the BALQSO phase (Farrah et al. 2007; Lipari et al. 2009). Although the similarity in the optical\/NIR properties of BAL and non-BALQSOs, higher reddening and polarization features in BALQSOs support the orientation model, the observed radio spectral indices of BALQSOs and the discovery of polar winds in BALQSOs challenge this scenario (Reichard et al. 2003; Zhou et al. 2006; Wang, Wang & Wang 2007; Montenegro-Montes et al. 2008). On the other hand, the anticorrelation between the radio loudness parameter and balnicity index, Compact Steep Spectrum (CSS) nature of BALQSOs, and the redshift dependence of BALQSO fraction favour the evolutionary model (Gregg, Becker & de Vries 2006; Farrah et al. 2007; Allen et al. 2011).","Citation Text":["Vivek, Srianand & Gupta 2016"],"Functions Text":["BALQSOs are classified into three sub-classes based on the ionization state of the absorbing gas: (a) high-ionization BALQSOs (HiBALs) consist of absorption from high-ionization lines such as C iv, Si iv, and N v"],"Functions Label":["Background"],"Citation Start End":[[755,783]],"Functions Start End":[[499,711]]}
{"Identifier":"2021MNRAS.503.1319G__Bolton_et_al._2008_Instance_1","Paragraph":"In this paper, following the method proposed by Ofek et al. (2003), we use the differential optical depth to lensing with respect to the lens redshift zl as the probability density. For a statistical sample that contains Nl strong lensing systems, the log-likelihood of observing the lens at redshift zl is given by\n(9)$$\\begin{eqnarray*}\r\n{\\rm ln}\\mathcal {L}({\\bf p})=\\sum _{i=1}^{N_l} {\\rm ln} \\delta p_i({\\bf p}),\r\n\\end{eqnarray*}$$where p represents the set of the VDF parameters (\u03b1, \u03b2) and the galaxy evolution parameters (\u03bdn, \u03bdv). Now one can perform Monte Carlo simulations of the posterior likelihood ${\\cal L} \\sim \\exp {(- \\chi ^2 \/ 2)}$, where the \u03c72 function is defined as\n(10)$$\\begin{eqnarray*}\r\n\\chi ^2=-2{\\rm ln}\\mathcal {L}.\r\n\\end{eqnarray*}$$in our statistical analysis of lens redshift distribution. The sample used in this paper is primarily drawn from Sloan Lens ACS Survey (SLACS) and recent large-scale observations of galaxies, which is compiled and summarized in Cao et al. (2015) and Shu et al. (2017). The combined sample includes 91 lenses from SLACS (Bolton et al. 2008; Auger et al. 2009; Shu et al. 2017) and an extension of the SLACS survey known as \u2018SLACS for the Masses\u2019 (S4TM) (Shu et al. 2015, 2017), 35 lenses from the BOSS emission-line lens survey (BELLS) (Brownstein et al. 2012) and BELLS for GALaxy-Ly\u03b1 EmitteR sYstemsGALLERY (BELLS GALLERY) (Shu et al. 2016a,b), 26 lenses from the Strong Lensing Legacy Survey (SL2S) (Sonnenfeld et al. 2013a,b), and five lenses from Lenses Structure and Dynamic (LSD) (Treu & Koopmans 2002; Koopmans & Treu 2003; Treu & Koopmans 2004). The advantage of this recently assembled lens sample, the detailed information of which is described and listed in Chen et al. (2019), lies in its well-defined observational selection criteria satisfying the assumption of spherical lens mass model. Fig. 1 shows the redshift distributions of the lensing systems used in our analysis. However, a statistical analysis requires a sample that is complete and has well-characterized, homogeneous selection criteria. Note that the lensing systems collected in this analysis are selected in very different manners. For instance, the SLACS, S4TM, and BELLS surveys, respectively, selected candidates from the spectroscopic observations of ETGs and look for the presence of higher redshift emission lines in Sloan Digital Sky Survey I (Eisenstein et al. 2001) and Sloan Digital Sky Survey-III (Eisenstein et al. 2011). These lens candidates were followed up with HST ACS snapshot imaging and after image processing. Therefore, in order to verify the completeness of the full early-type lens sample (hereafter Sample A), one additional subsample will also be applied to discuss its utility for the redshift test: 126 deflector-selected lenses from SLACS, S4TM, BELLS, and BELLS GALLERY (hereafter Sample B). Such choice is also motivated by the fact that the SLACS and BELLS lenses could be moderately suffered from the finite Sloan fibre size (Brownstein et al. 2012).","Citation Text":["Bolton et al. 2008"],"Functions Text":["The sample used in this paper is primarily drawn from Sloan Lens ACS Survey (SLACS)","The combined sample includes 91 lenses from SLACS"],"Functions Label":["Uses","Uses"],"Citation Start End":[[1081,1099]],"Functions Start End":[[820,903],[1030,1079]]}
{"Identifier":"2018AandA...611A..74R__Grady_et_al._2013_Instance_3","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)"],"Functions Text":["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,","detected two spiral arms and polarized light down to 0.\u2032\u2032 1 (15 au) from the star."],"Functions Label":["Background","Background"],"Citation Start End":[[1583,1602]],"Functions Start End":[[1414,1582],[1603,1685]]}
{"Identifier":"2021AandA...650A.205V__Fontaine_et_al._2012_Instance_1","Paragraph":"The formation of the approximately 40% of sdB stars that appear to be single has been a mystery for decades. In the absence of a companion, it is hard to explain how the star can expel most of its envelope on the RGB and still achieve core-He burning ignition. Recently, Pelisoli et al. (2020) suggested that all sdB stars might originate from binary evolution. Merger scenarios involving two low-mass white dwarfs have also been investigated (Webbink 1984; Han et al. 2002, 2003; Zhang & Jeffery 2012), but several facts challenge this hypothesis. First, compact low-mass white dwarf binaries are quite rare, even though some candidates are identified (Ratzloff et al. 2019). Second, the mass distributions of single and binary sdB stars are indistinguishable (Fontaine et al. 2012, Table 3 in particular). This mass distribution is mainly obtained from asteroseismology (some sdB stars exhibit oscillations, which allow the precise and accurate determination of the stellar parameters, including total mass; Van Grootel et al. 2013) and also from binary light-curve modeling for hot subdwarfs in eclipsing binary systems. Single and binary mass distributions peak at ~ 0.47 M\u2299, which is the minimum core mass required to ignite He through a He-flash at the tip of RGB (stars of \u2273 2.3 M\u2299 are able to ignite He quietly at lower core masses, down to ~0.33 M\u2299, but the more massive the stars, the rarer they are). A mass distribution of single sdB stars from mergers, in contrast, would be much broader (0.4\u22120.7 M\u2299; Han et al. 2002). With the DR2 release of Gaia (Gaia Collaboration 2018) and precise distances for many hot subdwarfs (Geier 2020), it is now also possible to build a spectrophotometric mass distribution for a much larger sample than what was achieved with the hot subdwarf pulsators or those in eclipsing binaries. Individual masses are much less precise than those obtained by asteroseismology or binary light-curve modeling (Schneider et al. 2019). However, single and binary spectrophotometric mass distributions share the same properties here as well, which tends to disprove the hypothesis of different origins for single and binary sdB stars. The third piece of evidence against merger scenarios (which would most likely result in fast-rotating objects) is the very slow rotation of almost all single sdB stars, as obtained through v sin i measurements (Geier & Heber 2012) or from asteroseismology (Charpinet et al. 2018). Moreover, their rotation rates are in direct line with the core rotation rates observed in RC stars (Mosser et al. 2012), which is another strong indication that these stars and the single sdB stars do share a same origin, that is, that they are post-RGB stars.","Citation Text":["Fontaine et al. 2012"],"Functions Text":["Second, the mass distributions of single and binary sdB stars are indistinguishable (",", Table 3 in particular)."],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[762,782]],"Functions Start End":[[677,762],[782,807]]}
{"Identifier":"2016AandA...595A.106C__Korista_&_Goad_2000_Instance_1","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"],"Functions Text":["The choice of SED is very important in the BLR modeling, as different lines respond to the continuum variations on different time scales"],"Functions Label":["Motivation"],"Citation Start End":[[573,592]],"Functions Start End":[[435,571]]}
{"Identifier":"2022AandA...666L...5G__Esquej_et_al._2014_Instance_1","Paragraph":"More recently, Garc\u00eda-Bernete et al. (2022) found that the PAH molecules responsible for the 11.3 \u03bcm PAH emission band are more resilient in the hard environments often present in AGN. In particular, the authors found larger 11.3\/7.7 \u03bcm and 11.3\/6.2 \u03bcm PAH ratios in AGN-dominated systems compared to SF galaxies, indicating a larger fraction of neutral PAH molecules (as noted by Smith et al. 2007 using a sample of relatively weak AGN). However, these studies were limited by the spatial resolution (\u223c4\u2033) and the low spectral resolution (R\u2004\u223c\u200460\u2013130) of Spitzer\/InfraRed Spectrograph (IRS). Previous sub-arcsecond angular resolution N-band (\u223c8\u201313 \u03bcm) ground-based spectroscopic studies investigated the 11.3 \u03bcm PAH feature in the nuclear and circumnuclear regions of AGN (e.g., H\u00f6nig et al. 2010; Gonz\u00e1lez-Mart\u00edn et al. 2013; Alonso-Herrero et al. 2014, 2016; Ramos et al. 2014; Esquej et al. 2014; Garc\u00eda-Bernete et al. 2015; Jensen et al. 2017; Esparza-Arredondo et al. 2018). However, these works were unable to provide definitive details regarding the effect of the AGN on the PAH molecules due to limited wavelength coverage and sensitivity. The changes in the PAH properties due to the presence of the AGN might be more prominent in their innermost regions of galaxies. Therefore, the unprecedented combination of high angular and spectral resolution (R\u2004\u223c\u20041500\u2005\u2212\u20053500) in the entire mid-IR range (4.9\u201328.1 \u03bcm) afforded by the James Webb Space Telescope (JWST)\/Mid-Infrared Instrument (MIRI; Rieke et al. 2015; Wells et al. 2015; Wright et al. 2015) is key to investigating PAH properties. In this Letter we report on the first investigation of PAH emission in the nuclear regions of three luminous Seyfert (Sy) galaxies and compare them with emission from SF regions using JWST\/MIRI Medium Resolution Spectrograph (MRS) data. This enables us, for the first time, to characterise the PAH properties of local luminous Sy galaxies (log (Lbol)> 44.46 erg s\u22121)1 at sub-arcsecond scales (\u223c0.45\u2033, \u223c142\u2013245 pc).","Citation Text":["Esquej et al. 2014"],"Functions Text":["Previous sub-arcsecond angular resolution N-band (\u223c8\u201313 \u03bcm) ground-based spectroscopic studies investigated the 11.3 \u03bcm PAH feature in the nuclear and circumnuclear regions of AGN (e.g.,","However, these works were unable to provide definitive details regarding the effect of the AGN on the PAH molecules due to limited wavelength coverage and sensitivity."],"Functions Label":["Background","Motivation"],"Citation Start End":[[880,898]],"Functions Start End":[[592,778],[980,1147]]}
{"Identifier":"2022AandA...661A..10B__Ghirardini_et_al._2021a_Instance_2","Paragraph":"It is also possible that these clusters have a smaller extent and can just be missed by our extent selection as our detection algorithm sets the extent to zero if it is smaller than 6 (Brunner et al. 2022). Following the method presented in Ghirardini et al. (2021a), we estimated several dynamical properties of the clusters in the point source sample and compared them with the extent-selected sample presented in Ghirardini et al. (2021a). In Fig. 5 we compare the distributions of the core radii (Rcore) constrained by the V06 model and the concentration parameter (cSB) between these two samples. The concentration parameter is defined as the ratio of the surface brightness within 0.1 R500 to the surface brightness within R500 (Ghirardini et al. 2021a; Santos et al. 2008; Maughan et al. 2012). Intuitively, the expectation is that the smaller the core radius, the more compact the cluster. The left panel of Fig. 5 clearly shows that the clusters in the point source sample have relatively smaller core radii, hence the emission is more concentrated in a smaller area. Consistently, the concentration of the point source sample shows a clear excess in higher values than the extent-selected sample, indicating that a significantly larger fraction of cool-core clusters and clusters host a central AGN. We performed the same experiment by applying cuts in flux 1.5 \u00d7 10\u221214 ergs s\u22121 cm\u22122 and in detection likelihood to test whether the clusters are missed by the extent selection because they are fainter and\/or more compact than the extent-selected clusters. The distribution of number density, core radius, and concentration parameter remains the same, indicating that the population of clusters in the point source catalog is more compact than the extent-selected sample. The extent-selected sample does not show a clear bias toward cool-core clusters or clusters with a central AGN, but contains the fraction of cool-cores is similar to that of SZ surveys (Ghirardini et al. 2021a). In this sample, we observe the opposite trend.","Citation Text":["Ghirardini et al. (2021a)"],"Functions Text":["we estimated several dynamical properties of the clusters in the point source sample and compared them with the extent-selected sample presented in"],"Functions Label":["Uses"],"Citation Start End":[[416,441]],"Functions Start End":[[268,415]]}
{"Identifier":"2016MNRAS.463.2348S__Narlikar_&_Padmanabhan_1991_Instance_1","Paragraph":"We appear to live in a flat, homogeneous and isotropic expanding Universe (at scales >100 Mpc) well described by the Friedmann\u2013Lema\u00eetre\u2013Robertson\u2013Walker (FLRW) metric where the scalefactor, a(t), describes the time dependence of the Universe geometry: a(t)\u2009\u221d\u2009t1\/2 for radiation domination, a(t)\u2009\u221d\u2009t2\/3 for matter domination, and $a(t)\\propto \\exp (\\scriptstyle\\sqrt{\\frac{\\Lambda }{3}}ct)$ for dark energy domination, where \u039b = 3H0\u03a9\u039bc\u22122 is the cosmological constant (e.g. Beringer et al. 2012) \u2013 see Table 1. A common normalization of the FLRW metric defines the scalefactor equal to unity at the present time (a(t0) = 1; e.g. Liddle & Lyth 1993), yielding, for a Universe dominated by a positive cosmological constant (cf. Table 1)\n\n(5)\n\n\\begin{eqnarray}\na_{\\Lambda }(t)=\\exp \\left(c\\sqrt{\\frac{\\Lambda }{3}}(t-t_{0})\\right),&amp; t_{\\Lambda }\\le t\\le t_{0} \\:.\n\\end{eqnarray}\n\nFor the matter-dominated Universe, we have\n\n(6)\n\n\\begin{eqnarray}\na_{\\text{m}}(t)=a_{\\Lambda }(t_{\\Lambda })\\left(\\frac{t}{t_{\\Lambda }}\\right)^{2\/3}, \\:\\: t_{\\text{eq}}\\le t\\le t_{\\Lambda } \\:.\n\\end{eqnarray}\n\nBetween the end of inflation at t = te \u223c 10\u221233\u2009s; (e.g. Narlikar & Padmanabhan 1991) and t = teq = 2.37 \u00d7 1012 s, the Universe was radiation-dominated since at the latter instant a radiation-matter equality is reached (Table 1). Exceptions to radiation-domination in the te \u2264 t \u2264 teq time interval might have taken place during cosmological phase transitions such as the QCD (when the Universe might have been dust-like). We, then, split the radiation-domination epoch into three parts. For the interval between the end of the QCD phase transition (t+) and teq, we define\n\n(7)\n\n\\begin{equation}\na_{\\text{rl}}(t)=a_{\\text{m}}(t_{\\text{eq}})\\left(\\frac{t}{t_{\\text{eq}}}\\right)^{1\/2}, \\:\\: t_{+}\\le t\\le t_{\\text{eq}} \\:.\n\\end{equation}\n\nDuring the QCD phase transition, we define\n\n(8)\n\n\\begin{eqnarray}\na_{\\text{QCD}}(t)=a_{\\text{rl}}(t_{+})\\left(\\frac{t}{t_{+}}\\right)^{n_{\\text{QCD}}}, \\:\\: t_{-}\\le t\\le t_{+},\n\\end{eqnarray}\n\nwhere nQCD = 2\/3 if the Universe experiences a QCD dust-like phase or nQCD = 1\/2 if the Universe continues to be radiation-dominated during that epoch. Finally, between the end of inflation (te) and the beginning of the QCD phase transition (t\u2212), we define\n\n(9)\n\n\\begin{equation}\na_{\\text{rm}}(t)=a_{\\text{QCD}}(t_{-})\\left(\\frac{t}{t_{-}}\\right)^{1\/2}, \\:\\: t_{\\text{e}}\\le t\\le t_{-} \\:.\n\\end{equation}\n\n","Citation Text":["Narlikar & Padmanabhan 1991"],"Functions Text":["Between the end of inflation at t = te \u223c 10\u221233\u2009s; (e.g.","and t = teq = 2.37 \u00d7 1012 s, the Universe was radiation-dominated since at the latter instant a radiation-matter equality is reached (Table 1)."],"Functions Label":["Background","Background"],"Citation Start End":[[1146,1173]],"Functions Start End":[[1090,1145],[1175,1318]]}
{"Identifier":"2021MNRAS.505..435S__Goodman_2009_Instance_1","Paragraph":"Detections of ionic, atomic, and molecular species in exoplanetary atmospheres serve as a unique and strong diagnostic of those chemical and dynamical processes driving their formation and evolution. Their detection and abundance measurements could act as indicators of planetary formation scenarios and reveal connections to the primordial protoplanetary disc and the host star (Williams & Cieza 2011; Mordasini et al. 2016; Madhusudhan et al. 2017). Furthermore, discoveries of atmospheric chemical species allow us to better understand various thermodynamical processes and chemistry, winds in the upper atmosphere (Goodman 2009; Snellen et al. 2010; Brogi et al. 2016; Madhusudhan et al. 2016; Wyttenbach et al. 2020), and to probe planetary interiors and various bulk properties through their abundances (Kite et al. 2016; Thorngren & Fortney 2019; Madhusudhan et al. 2020). A whole host of ions, atoms, and molecules have been detected through a variety of, often complementary, techniques, such as differential spectrophotometry using low-to-mid resolution spectroscopy (e.g. Gibson et al. 2012, 2017; Deming et al. 2013; Kreidberg et al. 2014; Kirk et al. 2016; Nortmann et al. 2016), and high resolution spectroscopic techniques (e.g. Redfield et al. 2008; Snellen et al. 2008; Rodler, Lopez-Morales & Ribas 2012; Birkby et al. 2013; Hoeijmakers et al. 2015, 2018, 2020; Brogi et al. 2016; Birkby et al. 2017; \u017d\u00e1k et al. 2019; Ehrenreich et al. 2020). To date, ionic species such as Fe\u2009ii and Ti\u2009ii (Hoeijmakers et al. 2019), atomic absorption from Na, K, H \u03b1, and He (e.g. Redfield et al. 2008; Sedaghati et al. 2016; Casasayas-Barris et al. 2017; Spake et al. 2018; Chen et al. 2020; Seidel et al. 2020), and molecules such as H2O, CH4, and CO (e.g. Konopacky et al. 2013; Brogi et al. 2014; Fraine et al. 2014; Barman et al. 2015; Sing et al. 2016) have been detected through the aforementioned techniques. Needless to say that this list of detected constituents is by no means exhaustive, nor that of methods employed to detect exoplanetary atmospheres. For instance, high-resolution imaging instruments such as SPHERE (Beuzit et al. 2019) and GRAVITY (Gravity Collaboration et al. 2017), both at the VLT (ESO\u2019s Very Large Telescope), through combination with low-dispersion spectroscopy, have facilitated direct measurements of exoplanetary atmospheres (Samland et al. 2017; Gravity Collaboration et al. 2020).","Citation Text":["Goodman 2009"],"Functions Text":["Furthermore, discoveries of atmospheric chemical species allow us to better understand various thermodynamical processes and chemistry, winds in the upper atmosphere"],"Functions Label":["Background"],"Citation Start End":[[619,631]],"Functions Start End":[[452,617]]}
{"Identifier":"2018AandA...616A.139G__Rao_et_al._(2009)_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":["Rao et al. (2009)"],"Functions Text":["The observations of NGC 1333 IRAS4A and IRAS16293 are presented in Girart et al. (2006) and","respectively."],"Functions Label":["Background","Background"],"Citation Start End":[[575,592]],"Functions Start End":[[483,574],[594,607]]}
{"Identifier":"2020MNRAS.494.2465B__Lee,_Sode-Yome_&_Park_1991_Instance_1","Paragraph":"Here we demonstrate that, over a fixed time interval, the planar three-body problem can be solved by means of a multilayered deep artificial neural network (ANN; e.g. see LeCun, Bengio & Hinto 2015). These networks are designed for high-quality pattern recognition by mirroring the function of our brains (McCulloch & Pitts 1943; Rosenblatt 1985) and have been successfully applied to a wide variety of pattern recognition problems in science and industry, even mastering the game of Go (Silver et al. 2016). The abundance of real-world applications of ANNs is largely a consequence of two properties: (i) an ANN is capable of closely approximating any continuous function that describes the relationship between an outcome and a set of covariates, known as the universal approximation theorem (Cybenko 1989; Hornik 1991); and (ii) once trained, an ANN has a predictable and a fixed computational burden. Together, these properties lead to the result that an ANN can be trained to provide accurate and practical solutions to Newton\u2019s laws of motion, resulting in major improvements in computational economy (Lee, Sode-Yome & Park 1991) relative to modern technologies. Our proof-of-principle method shows that an ANN can accurately match the results of converged solutions found using the arbitrary precision numerical integrator that, for computationally challenging scenarios, e.g. during multiple close encounters, can offer numerical solutions at a fraction of the time cost and CO2 expense. We demonstrate the importance of training an ANN on converged solutions. This enables the trained ANN to accurately predict particle locations even when a conventional \u2018double-precision\u2019 numerical integrator fails dramatically. By training an ANN that can accurately compute particle trajectories during close encounters, our work extends previous work training neural networks on an n-body-type problem (e.g. Quito, Monterola & Saloma 2001; Battaglia et al. 2016). Our findings also add to the growing body of literature that supports machine learning technologies being developed to enrich the assessment of chaotic systems (Pathak et al. 2018; Stinis 2019) and providing alternative approaches to classical numerical solvers more broadly (Hennig, Osborne & Girolami 2015).","Citation Text":["Lee, Sode-Yome & Park 1991"],"Functions Text":["Together, these properties lead to the result that an ANN can be trained to provide accurate and practical solutions to Newton\u2019s laws of motion, resulting in major improvements in computational economy","relative to modern technologies."],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[1108,1134]],"Functions Start End":[[905,1106],[1136,1168]]}
{"Identifier":"2022ApJ...926..208Y__Zheng_&_Hu_2018_Instance_1","Paragraph":"Magnetic flux ropes are frequently observed magnetic structures with helical magnetic field lines and a strong core field. In spacecraft observations, flux ropes are often identified from bipolar variations in one magnetic field component and the enhancement of magnetic strength at the center. There are various observations of flux ropes in the magnetosphere (e.g., Khurana et al. 1995; Slavin 2003; Yang et al. 2013; Poh et al. 2019; Sun et al. 2019), boundary of magnetosphere (e.g., Rijnbeek et al. 1984; Kawano & Russell 1997; Fear et al. 2008, 2009; Akhavan-Tafti et al. 2018; Hwang et al. 2018; Yao et al. 2020), and solar wind (e.g., Zheng & Hu 2018; Blanco-Cano et al. 2019; Bai et al. 2020). Flux ropes are also observed in the magnetospheres of Mercury (Zhong et al. 2020a), Mars (Briggs et al. 2011), Jupiter (Sarkango et al. 2021), and Saturn (Jasinski et al. 2016). Flux ropes are widely believed to result from magnetic reconnection. In earlier studies, some flux ropes near the magnetopause are referred to as flux transfer events (FTEs; Russell & Elphic 1978), which allow plasma transportation across the magnetopause. Various mechanisms for FTE generation have been proposed, such as single X-line reconnection (Scholer 1988; Southwood et al. 1988) and multiple X-line reconnection (Lee & Fu 1985). Recent simulations show that reconnection can generate not only fluid-scale flux ropes but also those with ion or even electron scales (Daughton et al. 2011; Hoilijoki et al. 2019; Lu et al. 2020). These kinetic-scale flux ropes may interact with one another to form larger-scale flux ropes (Daughton et al. 2011), which could be manifested in spacecraft observations as entangled flux ropes (Wang et al. 2017; \u00d8ieroset et al. 2019; Qi et al. 2020). There are also observations of ion-scale flux ropes near the ion\/electron diffusion region (Wang et al. 2016; Hwang et al. 2018; Poh et al. 2019; Dong et al. 2020). The flux ropes can change in size while moving through convection, often expanding (Dong et al. 2017; Akhavan-Tafti et al. 2018, 2019) and sometimes contracting (Hasegawa et al. 2016). Eastwood et al. (2012) investigated that the flux ropes in the day side and in the distant-tail magnetopause have similar orientations and comparable magnetic flux content, indicating that the flux ropes may be in quasi-equilibrium as they are convected tailwards along the magnetopause. Moreover, ion-scale flux ropes are believed to be responsible for exciting waves (Huang et al. 2016; Wang et al. 2019) or accelerating particles (Zhu et al. 2019; Zhong et al. 2020b), implying that they are actively involved in magnetospheric physics.","Citation Text":["Zheng & Hu 2018"],"Functions Text":["There are various observations of flux ropes in the","and solar wind (e.g.,"],"Functions Label":["Background","Background"],"Citation Start End":[[643,658]],"Functions Start End":[[295,346],[621,642]]}
{"Identifier":"2022ApJ...934..145I__from_1995_Instance_1","Paragraph":"The projected position angle (PA) of the jet gradually rotates counterclockwise with increasing distance from the jet base. In the 2.3 GHz VLBA image the PA is 50\u00b0 \u00b1 5\u00b0 east of north at about 20 mas from the core and the jet bends toward the north\u2013south direction closer to the core, consistently with the jet orientation in the nonsimultaneous 1.6 GHz image at 5\u201310 mas from the core (Shen et al. 1999). At 8.7 GHz we find a PA of 25\u00b0 \u00b1 5\u00b0 at about 5 mas from the core, consistent with the archival VLBA monitoring results at 15 GHz, showing a persistent jet orientation at a PA of about 30\u00b0 in the epochs from 1995 to 2013 at angular scales \u223c5 mas (Kellermann et al. 1998; Pushkarev et al. 2017). The same MOJAVE observations hint at more variability of the jet PA on the smallest resolved scales \u223c1 mas. Observations by Shen et al. (2002) with the VLBA between 1994 and 2000 across four frequencies (5, 12, 15, and 43 GHz) showed a consistent PA orientation of 30\u00b0 at 5 GHz with a clockwise shift of about 51\u00b0\u201367\u00b0 at 43 GHz. At 86 GHz a possible bent jet structure is seen, with the inner jet oriented with a PA of about \u221240\u00b0 less than 0.3 mas from the core, and an apparent transition to a northeast direction further out. The \u221240\u00b0 PA with respect to the core is consistent with the PA of the single jet component located \u223c400 \u03bcas from the core, imaged from 2018 April GMVA observations (see Figure 4 of Issaoun et al. 2021). 2018 images at 86 GHz do not indicate jet bending. At 230 GHz, the component C1 is located at a PA of \u221245\u00b0 and C2 at a PA of \u221235\u00b0 with respect to the core C0. This morphology can potentially be explained by a helical structure in the jet (Conway & Murphy 1993; Steffen et al. 1995). Such a helical jet structure can be caused by an orbiting lower-mass secondary black hole around a stable primary central black hole, as also proposed for 4C 73.18 (Roos et al. 1993) and OJ 287 (e.g., Dey et al. 2021; G\u00f3mez et al. 2022), precession caused by a wobbling disk (Britzen et al. 2018), or a large-scale accretion flow that is tilted with respect to the black hole spin axis (e.g., as in M81; Mart\u00ed-Vidal et al. 2011).","Citation Text":["Kellermann et al. 1998"],"Functions Text":["At 8.7 GHz we find a PA of 25\u00b0 \u00b1 5\u00b0 at about 5 mas from the core, consistent with the archival VLBA monitoring results at 15 GHz, showing a persistent jet orientation at a PA of about 30\u00b0 in the epochs from 1995 to 2013 at angular scales \u223c5 mas"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[651,673]],"Functions Start End":[[405,649]]}
{"Identifier":"2021AandA...656A..16C__Bruno_&_Carbone_2013_Instance_2","Paragraph":"Investigations of the turbulent nature of solar wind fluctuations have been ongoing for more than half a century (see, e.g., Bruno & Carbone 2016). Advances have been made consistently thanks to the increasingly accurate measurements of several dedicated space mission as well as to the enormous improvement of numerical calculation, new detailed models and theoretical frameworks, and the development of specific data analysis techniques. Nevertheless, the extremely complex nature of the system and the coexistence of multiple actors, scales, and dynamical regimes have led to a number of questions that remain open (Viall & Borovsky 2020). Among these, the very nature of the turbulent cascade of the solar wind flow and its relationship with the small-scale processes still need to be described in full (Tu & Marsch 1995; Bruno & Carbone 2013; Matthaeus & Velli 2011; Chen 2016). Magnetic field fluctuations have been characterized with great detail at magnetohydrodynamic and kinetic scales, for example, through spectral and high-order moments analysis (Tu & Marsch 1995; Bruno & Carbone 2013). The anisotropic nature of magnetic turbulence has also been addressed, and is still being debated, due to the limited access to three-dimensional measurements in space (see, e.g., Horbury et al. 2008, 2012; Sorriso-Valvo et al. 2010; Yordanova et al. 2015; Verdini et al. 2018; Telloni et al. 2019a; Oughton & Matthaeus 2020). Velocity fluctuations have been studied thoroughly (see, e.g., Sorriso-Valvo et al. 1999; Bruno & Carbone 2013), although the kinetic scales still remain quite unexplored for instrumental limitations, most notably in the sampling time resolution. Both the velocity and magnetic field show highly variable turbulence properties, with well developed spectra, strong intermittency (Sorriso-Valvo et al. 1999), anisotropy, and linear third-order moments scaling (Sorriso-Valvo et al. 2007; Carbone et al. 2011). The level of Alfv\u00e9nic fluctuations (mostly but not exclusively found in fast streams, see e.g., D\u2019Amicis et al. 2011; Bruno et al. 2019) are believed to be associated with the state of the turbulence. In particular, solar wind samples containing more Alfv\u00e9nic fluctuations are typically associated with less developed turbulence, as inferred from both shallower spectra and reduced intermittency (see Bruno & Carbone 2013, and references therein). This is consistent with the expectation that uncorrelated Alfv\u00e9nic fluctuations contribute to reduce the nonlinear cascade by sweeping away the interacting structures (Dobrowolny et al. 1980), as also confirmed by the observed anticorrelation between the turbulent energy cascade rate and the cross-helicity (Smith et al. 2009; Marino et al. 2011a,b).","Citation Text":["Bruno & Carbone 2013"],"Functions Text":["Magnetic field fluctuations have been characterized with great detail at magnetohydrodynamic and kinetic scales, for example, through spectral and high-order moments analysis"],"Functions Label":["Background"],"Citation Start End":[[1078,1098]],"Functions Start End":[[884,1058]]}
{"Identifier":"2021AandA...654A.124W__Gao_et_al._2015_Instance_1","Paragraph":"In view of the ultrahigh mass of the merger product of around 2.5\u2006M\u2299, it is undoubtedly necessary and important to further test the existence of post-merger NSs (Gao et al. 2016; Li et al. 2016, 2017; Ai et al. 2018; Zhu et al. 2018; Sarin et al. 2020; Beniamini & Lu 2021), which can provide a robust constraint on the equation of state of the NS matter and then improve our understanding of the low-energy feature of strong interaction. Besides the GRB 170817A\/AT 2017gfo event, searches for possible kilonova emission have already been carried out in the afterglows of many SGRBs since 2013 (Berger et al. 2013; Tanvir et al. 2013; Yang et al. 2015; Jin et al. 2015, 2016, 2018, 2020; Gao et al. 2015, 2017; Kasliwal et al. 2017). Among the SGRBs with a kilonova candidate, GRB 160821B has one of the lowest redshifts, of namely z\u2004=\u20040.162. From its optical\/near-infrared(NIR) afterglow, an obvious excess was found. Because of its close distance, the kilonova emission associated with GRB 160821B is in principle detectable and can provide a natural explanation for the observed optical\/NIR excess (Lamb et al. 2019; Troja et al. 2019). In view of its luminosity, which is lower than that of AT 2017gfo, the kilonova after GRB 160821B can in principle be modeled with a pure radioactive power. However, it could still be necessary to mention that a significant internal plateau appeared in the early X-ray afterglow during the first few hundred seconds (see the insert in Fig. 1), which indicated that a post-merger NS also exists in this event. According to these observations, Ma et al. (2021) suggested that the post-merger NS could collapse into a black hole and then the subsequent kilonova could be powered by the accretion onto the black hole. Nevertheless, alternatively, as suggested by Yu et al. (2018), the steep decay after the internal plateau may not represent the collapse of the NS, but may simply be caused by the suppression of the magnetic dipole radiation of the NS. In this case, the spin-down of the NS of relatively low luminosity can still power the kilonova emission, which can be generally called a mergernovae (Yu et al. 2013). This scenario can provide a natural explanation for the AT 2017gfo emission. Therefore, in our opinion, this situation could also appear in the case of GRB 160821B. This paper is devoted to testing whether or not there is a nonthermal emission component arising from the interaction between the NS wind and the merger ejecta, as mentioned above for AT 2017gfo.","Citation Text":["Gao et al. 2015"],"Functions Text":["Besides the GRB 170817A\/AT 2017gfo event, searches for possible kilonova emission have already been carried out in the afterglows of many SGRBs since 2013"],"Functions Label":["Background"],"Citation Start End":[[688,703]],"Functions Start End":[[439,593]]}
{"Identifier":"2020ApJ...894..107I__Kastner_et_al._1994_Instance_1","Paragraph":"AFGL 2136 IRS 1 (also referred to as CRL 2136, G17.64+0.16, and IRAS 18196\u22121331) is a luminous (1.0 \u00d7 105 L; Lumsden et al. 2013), high-mass (45 \u00b1 10 M; Maud et al. 2019) protostar that is believed to be in the latter stages of its evolution due to a variety of observed characteristics (Boonman & van Dishoeck 2003; Maud et al. 2018 and references therein). It is located at a distance of 2.2 kpc away from the Sun (Urquhart et al. 2014), and has been extensively observed from centimeter to micron wavelengths, at low and high angular resolution, and low and high spectral resolution. The myriad observations paint a picture where a single, isolated massive protostar is driving a wide-angle bipolar outflow through its natal cloud. The large scale outflow is observed in CO emission at millimeter wavelengths, with both the red and blue lobes being about 100\u2033 in extent (Kastner et al. 1994; Maud et al. 2018). Closer to the central source (2\u2033\u201310\u2033), the outflow cavity walls are seen in scattered light at near-infrared wavelengths (Kastner et al. 1992; Murakawa et al. 2008; Maud et al. 2018). The cool molecular envelope exhibits ice and dust absorption bands (Willner et al. 1982; Keane et al. 2001b; Dartois et al. 2002; Gibb et al. 2004), as well as molecular emission at millimeter wavelengths (van der Tak et al. 2000a, 2000b), but a much warmer component is also inferred from several different molecules seen in absorption in the near- to mid-infrared (Mitchell et al. 1990; Lahuis & van Dishoeck 2000; Keane et al. 2001a; Boonman et al. 2003; Boonman & van Dishoeck 2003; Goto et al. 2013, 2019; Indriolo et al. 2013a). The presence of a dust disk on small spatial scales was suggested by near-infrared polarization imaging (Minchin et al. 1991; Murakawa et al. 2008) and by mid-infrared interferometric observations (de Wit et al. 2011; Boley et al. 2013). A compact source was marginally resolved at centimeter wavelengths, along with a cluster of nearby 22 GHz H2O masers (Menten & van der Tak 2004), but only with the recent ALMA 1.3 mm continuum observations has the 93 \u00d7 71 mas dust disk been fully resolved (Maud et al. 2019). Thermal line emission at 232.687 GHz from the H2O \u03bd2 = 1\u20131, 55,0\u201364,3 transition has the same spatial extent as the dust emission, and the H2O gas velocities indicate Keplerian rotation within the disk (Maud et al. 2019). It is ideal that the reader has a clear picture of the AFGL 2136 region in mind to best understand the discussion throughout this paper. In particular, Figure 10 of Maud et al. (2018) provides an up-to-date schematic diagram of the AFGL 2136 region, and Figures 1 and 2 of Maud et al. (2019) present the compact disk observed in dust and gas emission, respectively.","Citation Text":["Kastner et al. 1994"],"Functions Text":["The large scale outflow is observed in CO emission at millimeter wavelengths, with both the red and blue lobes being about 100\u2033 in extent"],"Functions Label":["Background"],"Citation Start End":[[874,893]],"Functions Start End":[[735,872]]}
{"Identifier":"2021AandA...650L...4B__DeForest_et_al._2016_Instance_1","Paragraph":"As Parker Solar Probe (PSP) descends deeper into the solar corona on its succeeding orbits, its measurements reveal features of the heliospheric plasma that will significantly increase our fundamental understanding of the workings of the solar corona, the origins of the solar wind, and the behavior of heliospheric energetic particle populations (Fox et al. 2016). In the first several orbits, PSP approached progressively nearer to the range of altitudes at which the accelerating solar wind speed exceeds the Alfv\u00e9n speed. This critical zone is not a simple surface, but a more irregularly defined region above which (incompressive) magnetohydrodynamic signals, such as Alfv\u00e9n waves, can no longer return to the lower altitude corona. In this very region, the analysis of heliospheric imaging (HI; DeForest et al. 2016) has described a transition from striated images that appear to be highly collimated due to a structured magnetic field and lower plasma beta to higher altitude images that appeared to be more disordered and isotropic, a condition described as \u201cflocculation.\u201d A number of interesting features have been reported from observations in the region, including periods of near-corotation and, notably, the appearance of magnetic reversals or \u201cswitchbacks\u201d (Bale et al. 2019) and accompanying plasma jets (Kasper et al. 2019). Switchbacks have received considerable attention, with a focus on understanding how various plasma properties respond in and near them (Bale et al. 2019; Kasper et al. 2019; Dudok de Wit et al. 2020; McManus et al. 2020; Mozer et al. 2020; Whittlesey et al. 2019). Their origin is also a subject of active discussion, with ideas ranging from distant generation in the lower corona via interchange reconnection (Axford et al. 1999; Fisk & Kasper 2020; Zank et al. 2020) to a roll-up of the magnetic field by vortices in shear-driven turbulence (Ruffolo et al. 2020). Here we examine another important set of PSP observations and their behavior near switchbacks, namely the behavior of energetic particles as measured on PSP by the integrated Science Investigation of the Sun (IS\u2299IS) instrument suite (McComas et al. 2016, 2019).","Citation Text":["DeForest et al. 2016"],"Functions Text":["In this very region, the analysis of heliospheric imaging (HI;","has described a transition from striated images that appear to be highly collimated due to a structured magnetic field and lower plasma beta to higher altitude images that appeared to be more disordered and isotropic, a condition described as \u201cflocculation.\u201d"],"Functions Label":["Background","Background"],"Citation Start End":[[801,821]],"Functions Start End":[[738,800],[823,1081]]}
{"Identifier":"2018AandA...609A..28B__Hansen_&_Milosavljevi\u0107_2003_Instance_1","Paragraph":"A number of works have tried to explain this conundrum with different ideas, which follow one of three possibilities. When Genzel et al. (1996) first discovered missing RGB stars, they proposed that this might be due to stellar collisions depleting giant stars in the innermost parts through the high stellar densities that are reached near the SMBH. Later, Davies et al. (1998), Alexander (1999), Bailey & Davies (1999), and Dale et al. (2009) addressed this idea in detail and came to the conclusion that it can only explain the absence of the brightest and most extended giant stars. A different suggested possibility to explain the missing stars is that our GC does not only have one, but a binary of two massive black holes. This hypothesised binary could indeed carve a core into the stellar distribution through three-body interactions, as shown by a number of authors (Baumgardt et al. 2006; Portegies Zwart et al. 2006; Matsubayashi et al. 2007; L\u00f6ckmann & Baumgardt 2008; Gualandris & Merritt 2012). Nonetheless, the mass of the secondary needs to be of the order of \n\\hbox{${\\sim}10^5~M_{\\odot}$}\n~105M\u2299\n\n to explain the observed core. Such a massive secondary black hole would require the Milky Way to have experienced a major merger relatively recently, which is excluded by observations (see Hansen & Milosavljevi\u0107 2003; Yu & Tremaine 2003; Chen & Liu 2013). Moreover, the existence of such a massive secondary black hole is largely ruled out from a number of other considerations, for instance, constraints on the proper motion of Sgr A* from radio interferometry (see Gualandris & Merritt 2009). A number of inspiraling smaller-mass black holes can also create a shallow stellar density profile in the centre, which would relax the major merger requirement, as has been shown by Mastrobuono-Battisti et al. (2014). It has also been put forward that a star cluster falling towards the GC could increase the density profile outside of 10\u2032\u2032 , so that within this distance the profile would be like a core (Kim & Morris 2003; Ernst et al. 2009; Antonini et al. 2012; Antonini 2014). However, mass segregation would rebuild a steeper profile in as fast as a quarter of the relaxation time (as shown by Preto & Amaro-Seoane 2010; Amaro-Seoane & Preto 2011). This requirement would hence need a steady inflow of clusters to maintain a weak cusp profile in the centre. Finally, Merritt (2010) and Antonini (2014) found that if the nuclear cluster in the GC formed with an extended enough initial core profile, the current stellar distribution would still not be dynamically relaxed. While this solution is possible, it requires fine-tuning in the initial conditions to produce the density distribution seen in the GC. Amaro-Seoane & Chen (2014) proposed that the discs of young stars observed at the GC (Paumard et al. 2006) are connected to the missing bright giants: the precursor gaseous discmust have gone through a fragmentation phase that produced dense enough clumps to ensure an efficient removal of the outer layers of the giants through collisions, rendering them invisible to observations. Their degenerate cores would nonetheless populate the same area of phase space where the missing bright giants should be. Kieffer & Bogdanovi\u0107 (2016) recently showed that in order to be viable, this scenario requires the total mass of the fragmenting disc to have been several orders of magnitude higher than that of the early-type stars in the stellar discs in the GC.","Citation Text":["Hansen & Milosavljevi\u0107 2003"],"Functions Text":["Such a massive secondary black hole would require the Milky Way to have experienced a major merger relatively recently, which is excluded by observations (see"],"Functions Label":["Background"],"Citation Start End":[[1306,1333]],"Functions Start End":[[1147,1305]]}
{"Identifier":"2021MNRAS.504..444C__Steiner_&_McClintock_2012_Instance_1","Paragraph":"After a first part of ballistic, high-speed motion, the deceleration of RK1 was rather abrupt, which is something not observed in the majority of discrete ejecta from BH XRBs (e.g Mirabel & Rodr\u00edguez 1994; Fender et al. 1999; Miller-Jones et al. 2012). This scenario is consistent with a jet that travels first at constant speed in a low-density region of the ISM, which constitutes a large scale cavity around the system, before hitting the higher density wall of the cavity itself, as already proposed in Hao & Zhang (2009). Radio emission at late times would be produced by the external shocks between the plasma blob and the ISM cavity wall, in analogy with GRB afterglows (Wang et al. 2003). Those cavities have been suggested to exist at \u223cpc scales at least for XTE J1550\u2013564 and H1743\u2013322 (Hao & Zhang 2009; Steiner & McClintock 2012; Steiner, McClintock & Reid 2012; Migliori et al. 2017). We obtained the best modelling for the RK1 motion by a combination of a linear motion and a Sedov phase (see Section 3.3), achieved at late times due to the jet sweeping up ISM material on its path, in a similar way as Miller-Jones et al. (2011), which derived it from Wang et al. (2003). We therefore suggest that MAXI J1348\u2013630 is located in a similar cavity, possibly carved by previous jet activity, or resulting from the action of accretion disc winds. However, so far such winds have not been detected for MAXI J1348\u2013630. If we take our estimate of \u03b10 \u223c 25 arcsec, the upper limit on the inclination angle of the discrete ejection of \u03b81 \u223c 44\u25cb and the lowest acceptable distance of 1.6 kpc, we obtain that RK1 travelled at least \u223c0.3 pc before reaching the angular distance \u03b10, which is consistent with what has been observed for other discrete ejections (Corbel et al. 2000; Gallo et al. 2004) and could be taken as a rough lower limit on the size of the cavity.7 As argued by Hao & Zhang (2009), the presence of an underdense cavity could be a common characteristic of BH XRBs environments (Heinz 2002), and its existence could be strictly required for the jet to travel such a long distance (Hao & Zhang 2009).","Citation Text":["Steiner & McClintock 2012"],"Functions Text":["Those cavities have been suggested to exist at \u223cpc scales at least for XTE J1550\u2013564 and H1743\u2013322"],"Functions Label":["Uses"],"Citation Start End":[[815,840]],"Functions Start End":[[697,795]]}
{"Identifier":"2018ApJ...852L...1Z__Metzger_2017_Instance_1","Paragraph":"The short gamma-ray burst (GRB) 170817A was observed by the Fermi- Gamma-ray Burst Monitor (GBM; Goldstein et al. 2017). The fact that it was not detected by the Insight Hard X-ray Modulation Telescope (Li et al. 2017) suggests that the burst had a very weak fluence and a soft spectrum. The burst was highly noticeable because of its connection to gravitational-wave (GW) event GW170817, which was detected by the Laser Interferometer Gravitational-wave Observatory, about 1.7 s before the GBM was triggered (Abbott et al. 2017a). To understand the event, it is crucial that we obtain all of the object\u2019s physical properties. Compared with other short GRBs, the luminosity of GRB 170817A was extremely weak, thus suggesting that the jet was off-axis to the line of sight (Albert et al. 2017; Burgess et al. 2017; Fong et al. 2017; Granot et al. 2017a; He et al. 2017; Jin et al. 2017; Margutti et al. 2017; Metzger 2017; Troja et al. 2017; Wang et al. 2017; Xiao et al. 2017). The GW fitting parameters suggest an angle of less than 28\u00b0 (Abbott et al. 2017a). Given that no prompt X-rays were detected, a number of predictions have been made for the off-axis angle. From observations by the \n\n\n\n\n\n and \n\n\n\n\n\n telescopes, Evans et al. (2017) suggested that the viewing angle was \n\n\n\n\n\n. Constraints from deep Chandra observations suggest that it was greater than 23\u00b0 (Haggard et al. 2017), whereas constraints from radio-frequency observations of a relativistic jet (Alexander et al. 2017) suggested it was \n\n\n\n\n\n. On the basis of the upper limit from the Atacama Large Millimeter\/submillimeter Array (ALMA) and Giant Metrewave Radio Telescope (GMRT) at radio bands, Kim et al. (2017) found the angle to be 41\u00b0 (or 17\u00b0) but did not rule out other values. The modeling for several different bands suggested that the viewing angle was \n\n\n\n\n\n (Ioka & Nakamura 2017), \n\n\n\n\n\n (Granot et al. 2017a), or \n\n\n\n\n\n (Guidorzi et al. 2017). The \n\n\n\n\n\n versus \n\n\n\n\n\n plot in Pozanenko et al. (2017) shows that GRB 170817A belongs to neither the long nor short GRB groups, thus supporting the idea that it was triggered by an off-axis jet. However, the exact off-axis angle is still hard to predict. If it was derived from the relativistic beaming effect, it would have degenerated with the jet\u2019s Lorentz factor; if it was derived from the parameter fitting of the GW signal, it would have degenerated with the distance.","Citation Text":["Metzger 2017"],"Functions Text":["Compared with other short GRBs, the luminosity of GRB 170817A was extremely weak, thus suggesting that the jet was off-axis to the line of sight"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[908,920]],"Functions Start End":[[627,771]]}
{"Identifier":"2017ApJ...845...86E__Soler_&_Terradas_2015_Instance_3","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"],"Functions Text":["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.,"],"Functions Label":["Background"],"Citation Start End":[[1923,1944]],"Functions Start End":[[1576,1735]]}
{"Identifier":"2018ApJ...855...48Q__Falgarone_&_Passot_2003_Instance_1","Paragraph":"The dust condensations could not arise from a thermal Jeans fragmentation process. If that is the case, with a density of 104\u2013105 cm\u22123 for the surrounding medium, one may expect the mass of the condensations on the order of 20\u201350 M\u2299 (the Jeans mass at 105\u2013104 cm\u22123) and the nearest separations between the condensations around 0.2\u20130.5 pc (the Jeans length at 105\u2013104 cm\u22123); both are clearly inconsistent with the observations. Alternatively, small dense structures can be temporary density fluctuations frequently created and destroyed by supersonic turbulence (e.g., Elmegreen 1999; Biskamp 2003; Falgarone & Passot 2003). However, \u03c3nth is found to be subsonic or at most transonic. Goicoechea et al. (2016) also found that there is only a gentle level of turbulence in the Bar. So the turbulence does not seem to be strong enough to produce the condensations. Another force that could potentially compress the cloud and produce high-density structures is a high-pressure wave from the expansion of the H ii region. Goicoechea et al. (2016) detected a fragmented ridge of high-density substructures at the molecular cloud surface and three periodic emission maxima in HCO+ (4\u20133) from the cloud edge to the interior of the Bar, providing evidence that a high-pressure wave has compressed the cloud surface and moved into the cloud to a distance of \u223c15\u2033 from the dissociation front. The dust condensations are also located within a distance of 15\u2033 from the dissociation front, and thus are very likely over-dense structures created as the compressive wave passed by. The complex clumpy appearance of the condensations is probably related to the front instability of the compressive wave (Goicoechea et al. 2016), or an instability developed across different layers of the Bar (e.g., the thin-shell instability, Garc\u00eda-Segura & Franco 1996). The velocity structure of dense gas around the dust condensations may provide insights into the feasibility of this scenario. Figure 10 shows the intensity-weighted velocity map of the H2CS (71,7\u201361,6) emission. A velocity gradient in a northwest\u2013southeast direction (i.e., along the direction of the propagation of the compressive wave) is seen in the map, and is consistent with the scenario that the gas is being compressed by a high-pressure wave.","Citation Text":["Falgarone & Passot 2003"],"Functions Text":["Alternatively, small dense structures can be temporary density fluctuations frequently created and destroyed by supersonic turbulence (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[598,621]],"Functions Start End":[[427,567]]}
{"Identifier":"2021AandA...652A.124N__Wi\u015bniewska_et_al._(2016)_Instance_1","Paragraph":"Figure 9 presents Fourier power spectrum for wave period P versus height. The initial pulse has a Gaussian spectrum of wave number k which results in a spectrum of period P. The steepening of the magnetoacoustic waves results from the growing wave amplitude with height. Hence, waves with shorter wavelengths and wave periods are present for higher y-values in their Fourier spectra. For the pulse launched from y = y0 = 0 Mm (top), the main wave period of the downward-propagating waves becomes approximately the same for all values of y  0 Mm and is equal to about 250 s. Higher up, however, the period P decays with increasing y, and attains values close to 200 s. As a result of the cut-off only short-period waves can propagate upwards while long-period waves become evanescent. Hence, the relative contribution of long P waves weakens with increasing y. We note that some of our data fits the observational findings of Wi\u015bniewska et al. (2016), represented by diamonds over-plotted on the power spectra, and Kayshap et al. (2018), denoted by dots. The agreement of the theory with the observational data indicates that the results can be used to determine the background structure of the solar atmosphere and confirms that wave generation by the solar granulation in the partially ionized plasma dominates the behavior of the waves. For a pulse at the bottom of the photosphere (y = y0 = 0 Mm, Fig. 9 (top)), a jump in the dominant wave period is observed close to the height y = 1.5 Mm, which reaches a magnitude of 300 s. The signal launched from the bottom of the photosphere with the main period P = 250 s thus reaches the corona with P = 300 s. The wave periods in the photosphere are lower than Pac in this region, because max Pac = 240 s at y = 0.5 Mm, which means that magnetoacoustic waves are evanescent. However, in the corona above y = 2.3 Mm the acoustic cut-off period reaches values larger than P = 300 s. However, this is not the case for a pulse launched at a somewhat greater height, in the middle of the photosphere, namely for y = y0 = 0.25 Mm (see Fig. 9 (bottom)). Moreover, in this case the layers below the photosphere (y\u2004\u20040) oscillate with a dominant wave period of 225 s. In the photosphere, this wave period is also dominant and this time it increases slightly with height in the photosphere. Moreover, in this case there is a second dominant period, namely 250 s. In fact, this wave period is also the dominant one in the upper chromosphere and low corona in this case.","Citation Text":["Wi\u015bniewska et al. (2016)"],"Functions Text":["We note that some of our data fits the observational findings of","represented by diamonds over-plotted on the power spectra","The agreement of the theory with the observational data indicates that the results can be used to determine the background structure of the solar atmosphere and confirms that wave generation by the solar granulation in the partially ionized plasma dominates the behavior of the waves."],"Functions Label":["Similarities","Uses","Compare\/Contrast"],"Citation Start End":[[925,949]],"Functions Start End":[[860,924],[951,1008],[1054,1338]]}
{"Identifier":"2017MNRAS.465..492M__Taverna_et_al._2015_Instance_1","Paragraph":"Given the quite strong surface magnetic field of the M7, thermal radiation is expected to be polarized, either if emission is from a bare surface or from an atmosphere (see Turolla et al. 2004; Potekhin 2014). The polarization properties are quite different in the two cases, although there are still uncertainties, especially at optical\/ultraviolet (UV) wavelengths. One of the first predictions of quantum electrodynamics (QED), even before it was properly formulated, was vacuum birefringence, and, in particular, that a strong magnetic field affects the propagation of light through it (Heisenberg & Euler 1936; Weisskopf 1936). In thermally emitting INSs, radiation comes from a region comparable with the entire star surface, over which the magnetic field direction changes substantially. In the absence of QED vacuum polarization effects, this would produce a drastic depolarization of the radiation collected at infinity (Heyl, Shaviv & Lloyd 2003, see also Taverna et al. 2015; Gonz\u00e1lez Caniulef et al. 2016 and references therein). Vacuum birefringence dramatically increases the linear polarization of the observed radiation, from a level of a few \u2009per\u2009cent up to even \u223c100\u2009per\u2009cent, depending on the viewing geometry and the surface emission mechanism (Heyl & Shaviv 2000, 2002; Heyl et al. 2003; Taverna et al. 2015; Gonz\u00e1lez Caniulef et al. 2016). Detecting polarization in the thermal emission from the surface of an INS will be therefore extremely valuable. First, and independently on the physical conditions of the emitting region, the detection of a large degree of linear polarization in the signal would constitute the observational evidence of QED vacuum polarization effects in the strong-field regime. Secondly, the polarization observables can be compared with emission models and help to uncover the physical conditions of INS surfaces and atmospheres, ideally complementing spectral observations (Taverna et al. 2014; Gonz\u00e1lez Caniulef et al. 2016).","Citation Text":["Taverna et al. 2015"],"Functions Text":["In the absence of QED vacuum polarization effects, this would produce a drastic depolarization of the radiation collected at infinity","see also"],"Functions Label":["Background","Background"],"Citation Start End":[[966,985]],"Functions Start End":[[795,928],[957,965]]}
{"Identifier":"2019AandA...622A.146M__Arribas_et_al._(2014)_Instance_2","Paragraph":"Previous works (e.g. Holt et al. 2011; Arribas et al. 2014; Villar Mart\u00edn et al. 2014, 2015) have found very high reddening and densities associated with ionised outflows in local objects (e.g. H\u03b1\/H\u03b2\u2004\u223c\u20044.91 and ne\u2004\u2273\u20041000 cm\u22123, Villar Mart\u00edn et al. 2014). Concerning the reddening, although we find that the outflowing gas is generally less affected by dust extinction than the disc, the median value of the total distribution is significantly affected by dust (H\u03b1\/H\u03b2\u2004\u223c\u20044.16), with tails up to H\u03b1\/H\u03b2\u2004\u2273\u20046. Similarly, the outflow density of MAGNUM galaxies is higher than the values in the disc gas, but appears to be far lower than the values found by these authors. This could stem from the fact that the galaxies studied by Holt et al. (2011), Arribas et al. (2014) are local luminous or ultra-luminous infrared galaxies (U\/LIRGs), and those of Villar Mart\u00edn et al. (2014, 2015) are highly obscured Seyfert 2, thus sampling sources that are more gas and dust rich compared to our sample. However, our values are also lower than the outflow densities found in Perna et al. (2017) (ne\u2004\u223c\u20041200 cm\u22123), who targeted optically selected AGNs from the SDSS, and F\u00f6rster Schreiber et al. (2018b), who presented a census of ionised gas outflows in high-z AGN with the KMOS3D survey (ne\u2004\u223c\u20041000 cm\u22123). A possible explanation could be related to the high quality of our MUSE data, which also allows us to detect the faint [S\u202fII] emission associated with lower density regions. If we calculate the median densities of the disc and outflow components, weighting for the [S\u202fII] line flux, we obtain higher values (ne\u2004\u223c\u2004170 cm\u22123 and ne\u2004\u223c\u2004815 cm\u22123, for disc and outflow, respectively). This shows that previous outflow density values from the literature could be biased towards higher ne because they are based only on the most luminous outflowing regions, characterised by a higher S\/N. This could also mean that outflows at high-z could be far more extended than the values we observe.","Citation Text":["Arribas et al. (2014)"],"Functions Text":["This could stem from the fact that the galaxies studied by Holt et al. (2011),","are local luminous or ultra-luminous infrared galaxies (U\/LIRGs), and those of Villar Mart\u00edn et al. (2014, 2015) are highly obscured Seyfert 2, thus sampling sources that are more gas and dust rich compared to our sample."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[744,765]],"Functions Start End":[[665,743],[766,987]]}
{"Identifier":"2020AandA...641A.155V__G\u00f3mez-Guijarro_et_al._2019_Instance_1","Paragraph":"The scenario presented above has been formulated in various flavors to individually explain several of the properties reported here. The main addition of this work, namely the excitation of CO in distant main-sequence and starburst galaxies, fits in the general picture that we sketched. The ensemble of properties and correlations that we reported here can be also used to revisit the definition of what a starburst is. A standard operational classification is based on the distance from the observed empirical M\u22c6-SFR correlation, the main sequence. This proved to be a useful distinction and an excellent predictor of several trends (e.g., Sargent et al. 2014), but recent results, including our present and previous analysis (Puglisi et al. 2019), show that the demarcation between starburst and main-sequence galaxies is more blurred that we previously considered. We do detect starburst-like behaviors in galaxies on the main sequence (Elbaz et al. 2018), likely linked to the existence of transitional objects (Popping et al. 2017; Barro et al. 2017b; G\u00f3mez-Guijarro et al. 2019; Puglisi et al. 2019, and in prep. to limit the references to recent works based on submillimeter observations). Such transition might well imply an imminent increase of the SFR, driving the object in the realm of starbursts (e.g., Barro et al. 2017b), or its cessation, bringing the system back onto or even below the main sequence (G\u00f3mez-Guijarro et al. 2019; Puglisi et al. 2019), with the CO properties potentially able to distinguish between these two scenarios. Regardless of these transitional objects, a definition of starburst based on \u03a3SFR, rather than \u0394MS, would naturally better account for the observed molecular gas excitation properties, dust temperatures and opacities, or SFE (see also Elbaz et al. 2011; Rujopakarn et al. 2011; Jim\u00e9nez-Andrade et al. 2018; Tacconi et al. 2020). As an example, in Fig. 8 we show the mean SLED of the subsample of galaxies with both CO (2\u2005\u2212\u20051) and CO (5\u2005\u2212\u20054) coverage, split at its median \u03a3SFR. While only tentative at this stage, this suggests a trend of increasing CO excitation with \u03a3SFR, consistently with Fig. 7 and what mentioned above.","Citation Text":["G\u00f3mez-Guijarro et al. 2019"],"Functions Text":["We do detect starburst-like behaviors in galaxies on the main sequence","likely linked to the existence of transitional objects"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1058,1084]],"Functions Start End":[[869,939],[961,1015]]}
{"Identifier":"2022MNRAS.512.2854P__P\u00e9tri,_Heyvaerts_&_Bonazzola_2002_Instance_1","Paragraph":"The aforementioned fluid description offers a good starting point to understand the global electric circuit made of charge and current densities. However, it neglects some fundamental kinetic aspects required to self-consistently include single particle acceleration as well as radiation feedback. As kinetic simulations are much more demanding than fluid models, this approach was only scarcely investigated in the last century. Let us mention Krause-Polstorff & Michel (1985) who computed axisymmetric dead pulsar magnetospheres called electrospheres. Due to the axisymmetry of the problem they used rings of charges instead of point particles. Later with the advent on computational power, Smith, Michel & Thacker (2001) showed with slightly more sophisticated simulations that a fully field magnetosphere is unstable and collapse to an electrosphere. The first full three-dimensional electrosphere was constructed by McDonald & Shearer (2009), using an electromagnetic Particle in Cell (PIC) code. They neglect pair creation and therefore did not add any particle injection process. Eventually Philippov & Spitkovsky (2014) computed the first two-dimensional axisymmetric pulsar magnetosphere for an aligned rotator by permanently injecting particle supposed to be released from the surface, avoiding to end to an electrosphere configuration (P\u00e9tri, Heyvaerts & Bonazzola 2002). Depending on the volume injection rate, they were able to find any equilibrium between the force-free and the fully charge separated state. Chen & Beloborodov (2014) improved this model by adding a prescription for the pair creation, putting a threshold on the lepton Lorentz factor. Following the same lines, Cerutti et al. (2015) assumed particle injection only from the vicinity of the stellar surface. Belyaev (2015) injected particles from regions where a parallel electric field exists. The first full three-dimensional PIC simulations of a pulsar magnetosphere were performed by Philippov, Spitkovsky & Cerutti (2015b). For an aligned rotator Philippov et al. (2015a) also included general-relativistic corrections with frame-dragging. Soon after some observational signature predictions were added to compute light curves and spectra emanating from curvature and or synchrotron radiation like for instance Cerutti, Philippov & Spitkovsky (2016b) who then included polarization (Cerutti, Mortier & Philippov 2016a). This PIC simulations were then extended to the striped wind well outside the light-cylinder to study its dissipation (Cerutti & Philippov 2017; Cerutti, Philippov & Dubus 2020). The oblique magnetosphere with radiation and general-relativistic correction was eventually computed by Philippov & Spitkovsky (2018). Several other groups performed similar simulations like Brambilla et al. (2018) or Kalapotharakos et al. (2017, 2018) who tried to explicitly connect their simulation results to gamma-ray observations. Alternatively, more simply test particle trajectories can be explored within a fluid code (see for instance Brambilla et al. 2015).","Citation Text":["P\u00e9tri, Heyvaerts & Bonazzola 2002"],"Functions Text":["Eventually Philippov & Spitkovsky (2014) computed the first two-dimensional axisymmetric pulsar magnetosphere for an aligned rotator by permanently injecting particle supposed to be released from the surface, avoiding to end to an electrosphere configuration"],"Functions Label":["Background"],"Citation Start End":[[1347,1380]],"Functions Start End":[[1087,1345]]}
{"Identifier":"2021AandA...646L...9M__Quan_&_Herbst_(2007)_Instance_1","Paragraph":"The chemistry of C4H3N isomers in cold molecular clouds was discussed by Balucani et al. (2000), and more specifically by Balucani et al. (2002), based on crossed molecular beam experiments and ab initio calculations. In these studies, it was pointed out that reactions of the CN radical with methyl acetylene and allene are barrier-less and exothermic when producing CH3C3N and CH2CCHCN in the methyl acetylene reaction, and CH2CCHCN and HCCCH2CN in the reaction involving allene. Indeed, the reactions of CN with CH3CCH and CH2CCH2 were measured to be rapid at low temperatures (Carty et al. 2001). This chemical scheme was implemented in a chemical model by Quan & Herbst (2007) to explain the abundance of cyanoallene in TMC-1. Later on, Abeysekera et al. (2015) measured the product branching ratios of the reaction between CN and methyl acetylene at low temperatures using a chirped-pulse uniform flow and found that HC3N is the major product, while CH3C3N accounts for 22% of the products and CH2CCHCN is not formed. These results are in contrast with those obtained from crossed molecular beam experiments (Huang et al. 1999; Balucani et al. 2000, 2002), where CH2CCHCN is observed as a product of the CN + CH3CCH reaction. Therefore, the most stable isomer, CH3C3N, can be formed in the reaction of CN and methyl acetylene, the second most stable isomer, CH2CCHCN, can be formed when CN reacts with CH2CCH2 \u2013 and perhaps also with CH3CCH, depending on whether one gives credit to the chirped-pulse uniform flow experiment or to the crossed molecular beam ones \u2013 and the least stable isomer, HCCCH2CN, can only be formed in the reaction between CN and allene. These neutral-neutral reactions involving CN are therefore likely routes to the three C4H3N isomers in cold interstellar clouds such as TMC-1, where abundant CN, CH3CCH, and probably CH2CCH2 (non-polar and thus cannot be detected at radio wavelengths) are present. Moreover, the presence of HCCCH2CN (and perhaps also CH2CCHCN) can be used as a proxy of the non-polar C3H4 isomer allene since this isomer is only formed from CH2CCH2 in the aforementioned reactions of CN.","Citation Text":["Quan & Herbst (2007)"],"Functions Text":["This chemical scheme was implemented in a chemical model by","to explain the abundance of cyanoallene in TMC-1."],"Functions Label":["Background","Background"],"Citation Start End":[[661,681]],"Functions Start End":[[601,660],[682,731]]}
{"Identifier":"2022MNRAS.509..619I__R\u00f3\u017ca\u0144ska_et_al._2011_Instance_1","Paragraph":"A number of sources of systematic error disproportionately affect the soft X-rays. One important example is absorption by partially ionized material around the AGN. Although the response of the absorbing gas to changes in the irradiating flux from the AGN does contribute its own time lag, this should only influence the lag-energy spectrum for Fourier frequencies lower than those of interest for our analysis (Silva, Uttley & Costantini 2016). The influence of absorption on the soft X-ray region of the spectrum (e.g. Miller, Turner & Reeves 2008; Miller et al. 2010), however, is potentially much more important and can influence the shape of the lag-energy spectrum (I19). It will therefore be prudent to select the AGN with the clearest view of the inner regions. The soft X-ray region of the reflection spectrum is also the most sensitive to modelling assumptions such as vertical disc structure (Nayakshin, Kazanas & Kallman 2000; Done & Nayakshin 2007; R\u00f3\u017ca\u0144ska et al. 2011; Vincent et al. 2016). Moreover, X-ray reverberation models currently always assume that the time taken for photons to be reprocessed and reemitted in the disc atmosphere is very small compared with the light-crossing delays. This is a very good assumption for the fluorescence and scattering processes dominant for E \u2273 3 keV. However, the time taken for soft excess photons to be reprocessed will be longer, since these photons undergo enough interactions to approximately thermalize. This thermalization time is still expected to be small, but it is not yet clear if it is small enough to be neglected entirely. It is also important to note that the model we have explored here does not include a warm (kTe \u223c 0.1 keV), optically thick (\u03c4 \u223c 10\u201340) corona in addition to the hot corona, as is often invoked to explain the observed soft excess in AGNs (Czerny & Elvis 1987; Middleton et al. 2009; Done et al. 2012; Petrucci et al. 2018; Ursini et al. 2020). Models 1, 2, and 4 therefore effectively assume that the observed soft excess is dominated by reflection (e.g. Jiang et al. 2018; Garc\u00eda et al. 2019). If in reality there is indeed a warm corona, its presence would make it more difficult to constrain the soft X-ray shape of the reflection spectrum (Xu et al. 2021). Since we believe our models to be most robust for E \u2273 3 keV, we ran an alternative fit ignoring E  3 keV (Model 3), and found that the 3\u03c3 error on H0 almost doubles (dH0 \u2248 60 km s\u22121 Mpc\u22121) compared to the other fits. It is therefore desirable to also include soft X-rays, but it is encouraging that constraints can even be achieved without. Ignoring soft X-rays entirely is an extreme measure. An alternative approach would be to keep the soft X-rays but trial a variety of different models, for instance with and without a warm corona. This approach would likely return uncertainties somewhere between the two extreme scenarios explored here.","Citation Text":["R\u00f3\u017ca\u0144ska et al. 2011"],"Functions Text":["The soft X-ray region of the reflection spectrum is also the most sensitive to modelling assumptions such as vertical disc structure"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[962,982]],"Functions Start End":[[770,902]]}
{"Identifier":"2019MNRAS.489.2792Z__Zahid_et_al._2014_Instance_1","Paragraph":"The compact blobs instead likely have a different origin. The fact that compact blobs are unresolved even at the HST resolution, that they are found at \u223c1 kpc distance from the galaxy barycentre, that they have relatively small stellar masses (\u227215 per\u2009cent of the underlying disc), but are actively forming stars suggests that they are star-forming regions likely originated due to disc instability and fragmentation of the galaxy disc (Bournaud et al. 2014; Mandelker et al. 2017). The in situ formation of the compact blobs is further supported by their metallicity. In fact, while the disc properties are consistent with the stellar mass\u2013metallicity relation of z \u223c 2 star-forming galaxies (e.g. Maiolino et al. 2008; Zahid et al. 2014), compact blobs instead show metallicities inconsistent with the mass\u2013metallicity relation (Fig. 7, bottom right panel). They have a comparable metallicity to the discs, but \u223c1.5 dex lower stellar masses, so they are \u223c1 dex above the mass\u2013metallicity relation. Metallicity measurements for statistical sample of galaxies with M\u22c6 \u223c 108 M\u2299 at z \u223c 2 are still lacking and therefore at the low-mass end we are showing an extrapolation of the mass\u2013metallicity relation derived for galaxies with M\u22c6 \u2273 109 M\u2299. In Fig. 7, we also show the average location of dwarf galaxies with M\u22c6 \u223c 108 M\u2299 at z \u2272 1 (Kirby et al. 2013; Calabr\u00f2 et al. 2017; Hidalgo 2017). Despite some of them may have gas-phase metallicities up to 12 + log\u2009(O\/H) \u223c 8.5 (S\u00e1nchez Almeida et al. 2018), on average our compact blobs seem to be \u223c0.5 dex more metal-rich than dwarf galaxies (Fig. 7). The high metallicity of compact clumps further suggests that they formed in situ, due to the gravitational collapse of pre-enriched gas in unstable regions of the galaxy disc. The young ages of the blobs reported in Fig. 7 support these conclusions. In fact, the metallicity of star-forming regions is altered in about one galactic dynamical time (\u2273100 Myr), increasing due to the active star formation and internal production of metals. The fact that our sample clumps with metallicity measurements have ages \u2272 50 Myr points towards the conclusion that they formed in situ from metal-rich gas, since this time-scale is too short for the gas to be self-enriched due to internal star formation (Bournaud 2016). Finally, the metallicity of our sample of clumps does not clearly correlate with their star formation rate, as indeed expected if they formed from pre-enriched material, although larger statistical samples are needed to confirm this finding.","Citation Text":["Zahid et al. 2014"],"Functions Text":["The in situ formation of the compact blobs is further supported by their metallicity. In fact, while the disc properties are consistent with the stellar mass\u2013metallicity relation of z \u223c 2 star-forming galaxies (e.g.","compact blobs instead show metallicities inconsistent with the mass\u2013metallicity relation (Fig. 7, bottom right panel)."],"Functions Label":["Similarities","Differences"],"Citation Start End":[[721,738]],"Functions Start End":[[483,698],[741,859]]}
{"Identifier":"2018AandA...616A..99K__Narang_et_al._2016_Instance_2","Paragraph":"The high-resolution imaging observations of TR from IRIS reveal the ubiquitous presence of network jets. We have used three different IRIS observations of the quiet sun, which are located near the disk center. On the basis of careful inspection, 51 network jets are identified from three QS observations and used for further analysis. These 51 network jets are very well resolved and are not affected by the dynamics of other jets. The study is focused on the rotating motion of network jets along with the estimation of their other properties (speed, height, and lifetime). The mean speed, as predicted by statistical distributions of the speed, is 140.16 km s\u22121 with a standard deviation of 39.41 km s\u22121. The mean speed of network jets is very similar, as reported in previous works (e.g., Tian et al. 2014; Narang et al. 2016). However, in case of their lifetimes, we found a value that is almost double (105.49 s) that of the previously reported mean lifetime of the network jets (49.6 s; Tian et al. 2014). As mentioned above, we took only those network jets that are very well resolved in space and in the time; these criteria exclude short lifetime network jets. Therefore, our statistical distribution of the lifetime predicts a higher mean lifetime. The mean length of the network jets is 3.16 Mm with a standard deviation of 1.18 Mm. In the case of CH network jets, Tian et al. (2014) have reported that most of the network jets have lengths from 4.0 to 10.0 Mm. However, the mean length for QS network jets is smaller (3.53 Mm; Narang et al. 2016). So, the mean length for QS network jets from the present work is in good agreement with Narang et al. (2016). In addition, the apparent speed and length of these network jets are positively correlated, which is very similar to what has already been reported in previous works (Narang et al. 2016). Finally, we can say that these networks jets are very dynamic features of the solar TR, as revealed by their estimated properties.","Citation Text":["Narang et al. 2016"],"Functions Text":["However, the mean length for QS network jets is smaller (3.53 Mm;"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1539,1557]],"Functions Start End":[[1473,1538]]}
{"Identifier":"2016ApJ...819...23S__Kendall_et_al._1992_Instance_1","Paragraph":"Single-photon VUV excitation from the N2(X1\n\n\n\n\n\n) ground state can only populate the ungerade states. Many of these ungerade states can be coupled by spin\u2013orbit or vibronic interactions. Hence, only the ungerade states of N2 are calculated and presented. Because of the potential diffuse nature of the excited states, their wavefunctions were calculated using a large basis set composed by the aug-cc-pVQZ quality, which is augmented by 3s and 2p diffuse Gaussian-type orbitals (GTOs; Dunning 1989; Kendall et al. 1992; Woon & Dunning 1995). In addition, the expected high density of the electronic states in this energy range should lead to the mixing of their electronic wavefunctions. Therefore, a multiconfigurational approach is adopted where the electronic computations were carried out using the Complete Active Space Self Consistent Field (CASSCF; Knowles & Werner 1985; Werner & Knowles 1985) approach followed by the internally contracted Multi Reference Configuration Interaction (MRCI; Knowles & Werner 1988; Werner & Knowles 1988) technique as implemented in MOLPRO (Werner et al. 2012). We followed the procedure established in Spelsberg & Meyer (2001), Ndome et al. (2008), and Hochlaf et al. (2010a, 2010b). Briefly, the CASSCF active space is constructed by the valence molecular orbitals (MOs) of N2 increased by one sg and one pg MOs. This allows better relaxation of the wavefunctions of the N2 electronic states whose configurations differ in their s and p orbital occupations as those of interest at present. In MRCI, all CASSCF configurations were taken into account as reference. Finally, the CASSCF wavefunctions were used to evaluate the spin\u2013orbit matrix elements. These spin\u2013orbit computations were performed in Cartesian coordinates, where the CASSCF wavefunctions were used as the multielectron basis for the two-step spin\u2013orbit coupling calculation (Llusar et al. 1996; Zeng et al. 2011) through the Breit\u2013Pauli Hamiltonian (Berning et al. 2000) as implemented in MOLPRO.","Citation Text":["Kendall et al. 1992"],"Functions Text":["Because of the potential diffuse nature of the excited states, their wavefunctions were calculated using a large basis set composed by the aug-cc-pVQZ quality, which is augmented by 3s and 2p diffuse Gaussian-type orbitals (GTOs;"],"Functions Label":["Uses"],"Citation Start End":[[500,519]],"Functions Start End":[[256,485]]}
{"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)"],"Functions Text":["Figure 5 shows the comparison between our EWs measurements and the line strengths measured by","(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,"],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Similarities"],"Citation Start End":[[1530,1551]],"Functions Start End":[[1436,1529],[1552,1637],[1638,1721]]}
{"Identifier":"2020MNRAS.496.3448D__Joy,_Sahni_&_Starobinsky_2008_Instance_1","Paragraph":"The aforementioned form of power-law primordial power spectrum is a prediction of inflation where the scalar field (inflaton) slowly rolls down to the bottom of the flat inflationary potential. With the constraints on the tilt and an upper bound on the amplitude of tensor perturbation with respect to scalar perturbation, various surveys have ruled out a wide class of models. However, fundamental questions such as the energy scale of inflation and the detailed shape of the potential remain unanswered. Note that any changes in the nearly flat potential will eventually lead to certain features in the spectrum. Local glitches in the potential including rapid change of its amplitude, or the break in its first or second derivatives (Starobinsky 1992; Starobinsky 1998; Adams, Cresswell & Easther 2001; Covi et al. 2006; Joy, Sahni & Starobinsky 2008; Joy et al. 2009; Hazra et al. 2010; Miranda, Hu & Adshead 2012; Benetti 2013; Cadavid & Romano 2015; Chluba, Hamann & Patil 2015), false vacuum decay (leading to open inflation, in particular) (Linde 1999; Linde, Sasaki & Tanaka 1999; Bousso, Harlow & Senatore 2014), or an inflection point in the potential (Allahverdi & Mazumdar 2006; Jain et al. 2009), or oscillations in the potential (Ashoorioon & Krause 2006; Pahud, Kamionkowski & Liddle 2009; Biswas, Mazumdar & Shafieloo 2010; Flauger et al. 2010; Aich et al. 2013; Hazra 2013; Peiris, Easther & Flauger 2013; Easther & Flauger 2014; Meerburg, Spergel & Wandelt 2014; Motohashi & Hu 2015; Miranda et al. 2016) all lead to local and non-local oscillations in the spectrum. Direct reconstruction of the primordial spectrum from the Planck data (Hazra et al. 2014c) hints at large-scale oscillations, an intermediate-scale burst of oscillations, and persistent high-frequency oscillations within intermediate to small scales. While these types of features can be obtained by different potentials, in this work we will be using the WWI, which is known the provide these local and non-local features in a unified framework.","Citation Text":["Joy, Sahni & Starobinsky 2008"],"Functions Text":["Note that any changes in the nearly flat potential will eventually lead to certain features in the spectrum. Local glitches in the potential including rapid change of its amplitude, or the break in its first or second derivatives","all lead to local and non-local oscillations in the spectrum."],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[824,853]],"Functions Start End":[[506,735],[1524,1585]]}
{"Identifier":"2021ApJ...908...95H__Goto_et_al._2011_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":["Goto et al. 2011"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[621,637]],"Functions Start End":[[307,550]]}
{"Identifier":"2015MNRAS.450...53H__Mu\u00f1oz_et_al._2014_Instance_2","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"],"Functions Text":["In a shearing disc, if the cells adapt in a truly Lagrangian manner, then they are inevitably deformed into a highly sheared\/irregular shape"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1198,1215]],"Functions Start End":[[1056,1196]]}
{"Identifier":"2015MNRAS.450..630S__Moore_et_al._1996_Instance_1","Paragraph":"While there are increasing efforts to try to explain the SFR dependence on the environment, by conducting surveys at high redshift (e.g. Hayashi et al. 2010; Matsuda et al. 2011; Sobral et al. 2011; Muzzin et al. 2012; Koyama et al. 2013; Darvish et al. 2014; Tal et al. 2014), so far such studies have not been able to fully reveal the physical processes leading to the ultimate quenching of (satellite) star-forming galaxies (e.g. Peng et al. 2010; Muzzin et al. 2012, 2014). Several strong processes have been proposed and observed, such as harassment (e.g. Moore et al. 1996), strangulation (e.g. Larson, Tinsley & Caldwell 1980) and ram pressure stripping (e.g. Gunn & Gott 1972; Fumagalli et al. 2014). Observations are also showing a variety of blueshifted rest-frame ultraviolet (UV) absorption lines which indicate that most star-forming galaxies at least at z \u223c 1\u20132 are able to drive powerful gas outflows (e.g. Shapley et al. 2003; Weiner et al. 2009; Kornei et al. 2012) which may play a significant role in quenching, particularly if those happen in high-density environments. Evidence of such galactic winds have also been seen in e.g. F\u00f6rster Schreiber et al. (2009) through broad components in the rest-frame optical H\u03b1 and [N\u2009ii] emission line profiles (e.g. Genzel et al. 2011). Spatially resolved observations allow for constraints on the origin of the winds within galaxies, and on the spatial extent of the outflowing gas, which are essential to derive mass outflow rates. In field environments, it is expected that such outflows will not be able to escape the halo (as long as it is massive enough and it is not a satellite), and in many conditions would likely come back and further fuel star formation (e.g. Hopkins et al. 2014). However, in the most massive clusters, such strong outflows will likely result in significant amounts of gas being driven out of the subhaloes that host star-forming galaxies, enriching the intracluster medium (ICM) and quickly quenching star-forming galaxies with the highest SFRs\/highest outflow rates.","Citation Text":["Moore et al. 1996"],"Functions Text":["Several strong processes have been proposed and observed, such as harassment (e.g."],"Functions Label":["Background"],"Citation Start End":[[561,578]],"Functions Start End":[[478,560]]}
{"Identifier":"2018AandA...618A..67C__Moriguchi_et_al._(2002)_Instance_2","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)","Moriguchi et al. (2002)"],"Functions Text":["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","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."],"Functions Label":["Background","Compare\/Contrast"],"Citation Start End":[[492,515],[1670,1693]],"Functions Start End":[[335,491],[1694,2049]]}
{"Identifier":"2017AandA...606A..17M__Kennicutt_(1998)_Instance_2","Paragraph":"The SFR reported in Table C.1 refers to a stellar mass range from Mlow = 0.1M\u2299 to Mup = 100M\u2299, is averaged over the past \u0394t = 100 Myr, and was calculated using the standard SFR(LIR) relationship from Kennicutt (1998; here scaled to a Chabrier 2003, IMF) (1)\\begin{equation} \\label{eq:sfr} \\textit{SFR}=10^{-10}\\times L_{\\rm IR}[{L}_{\\sun}]\\, {M}_{\\sun}~{\\rm yr}^{-1}. \\end{equation}SFR=10-10\u00d7LIR[L\u2299]\u2009M\u2299yr-1.This calibration relies on the starburst synthesis models of Leitherer & Heckman (1995), and it is based on the assumption of solar metallicity, and an optically thick (\u03c4dust \u226b 1) starburst region, in which case LIR is a good proxy of the system\u2019s bolometric luminosity (LIR \u2243 Lbol), and hence a sound, calorimetric probe of the obscured, current stellar birth rate. A possible caveat is that the contribution to the dust heating by more evolved stellar populations (the cirrus component; e.g. Helou 1986; Lonsdale Persson & Helou 1987; Walterbos & Greenawalt 1996) is not taken into account. If the cirrus ISM component heated by the more general galactic UV radiation field contributes to LIR, then the Kennicutt (1998) relationship overestimates the SFR. Another issue is the fact that some percentage of the UV photons can escape the starburst region without being absorbed, and hence are not reprocessed into IR photons (indeed, some of our SMGs are visible in the rest-frame UV images; Miettinen et al. 2017b). The MAGPHYS code also gives the SFR as an output, and contrary to the aforementioned LIR diagnostic, the model permits for the heating of the dust by older and longer-lasting stellar populations. We found that the SFR(LIR) is somewhat higher on average than SFRMAGPHYS: the SFR(LIR) \/SFRMAGPHYS ratio was found to range from 0.47 to 6.92 with a median of \\hbox{$1.31^{+0.83}_{-0.17}$}1.31-0.17+0.83, where the \u00b1 errors represent the 16th\u201384th percentile range (see the corresponding panel in Fig. 2). If, instead of \u0394t = 100 Myr, the aforementioned comparison is done by using the SFRMAGPHYS values averaged over the past \u0394t = 10 Myr, the median SFR(LIR) \/SFRMAGPHYS ratio is found to be \\hbox{$1.15^{+0.38}_{-0.27}$}1.15-0.27+0.38, which is consistent with the results obtained by da Cunha et al. (2015). Unless otherwise stated, in our subsequent analysis we use the SFR averaged over the past 100 Myr as calculated using Eq. (1). ","Citation Text":["Kennicutt (1998)"],"Functions Text":["If the cirrus ISM component heated by the more general galactic UV radiation field contributes to LIR, then the","relationship overestimates the SFR."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1112,1128]],"Functions Start End":[[1000,1111],[1129,1164]]}
{"Identifier":"2018MNRAS.480.4931V__Dokkum_2001_Instance_1","Paragraph":"We made use of the SALT product data generated by the in-house pipeline called PySALT (Crawford et al. 2010), which mosaics the individual CCD data to a single FITS file, corrects for cross-talk effects, and performs bias and gain corrections. We then carried out further reduction steps using our own custom tools, written in Python\/PyRAF, which consist of several modules designed to handle specific steps of processing. First, we trim the regions of the CCD not containing any usable data and then fill the CCD gaps using a gradient fill. To do this, we calculate the gradient across the gap by measuring the flux values of the pixels to the right and the left of the gap, computing the difference and normalizing it to the number of pixels belonging to the CCD gap (the gradient). We then use the gradient to determine the flux values for pixels in each row of the image and replace the zero-value pixels by these flux values. This step is important to prevent discontinuities in the brightness distribution which affect the flat-fielding process which relies on the modelling of the large-scale illumination structure. Another advantage of this gradient fill is seen in the subsequent step of background subtraction. A CCD gap which appears straight in the original images becomes curved as a result of the coordinate transformation step. This causes background subtraction to result in discontinuous patches. By filling the CCD gap, these patches can be avoided. We then remove the cosmic rays using the LACosmic algorithm (van Dokkum 2001). Calibration in wavelength is achieved using arc lamps and has typical error of \u00b10.35 \u00c5. We use the arc lamp to determine a coordinate transformation to \u2018straighten\u2019 frames so that every column corresponds to a unique wavelength making the background sky shape fitting and subtraction possible. Using the spectrophotometric standard data, we determine the relative flux calibration. We then check individual frames for alignment and co-add them. During this step, a standard deviation frame is constructed from which an error frame is obtained. Before the co-addition, to have consistent noise characteristics, we equalize the effective exposures of the frames \u2013 this is needed because the pupil of SALT changes by design during the observation. If Oi(x, y) are the individual observed frames, the final spectrum F(x, y) and error spectrum E(x, y) can be written as, \n(1)\r\n\\begin{equation*}\r\nF(x,y) = \\left\\langle O(x,y) \\right\\rangle = \\frac{\\sum _i^N O_i(x, y)}{N}, \r\n\\end{equation*}\r\n(2)\r\n\\begin{equation*}\r\nE(x,y) = \\sqrt{\\frac{\\sum _i^N (O_i(x,y) - \\left\\langle O(x,y)\\right\\rangle)^2}{N-1}}. \r\n\\end{equation*}\r\nAn alternate way of obtaining the error frames is to start from the raw CCD images whose noise properties are relatively well understood to try to model the effects of every single reduction step on the errors. This is, however, very complicated and so we adopt the above technique of determining our final error frame. Since we are combining only six to eight frames, we might overestimate our errors, but this overestimation does not affect the findings presented in the paper. Our final step is to correct for foreground extinction and this is done using the deredden task in IRAF.","Citation Text":["van Dokkum 2001"],"Functions Text":["We then remove the cosmic rays using the LACosmic algorithm"],"Functions Label":["Uses"],"Citation Start End":[[1530,1545]],"Functions Start End":[[1469,1528]]}
{"Identifier":"2021MNRAS.500.2336Y__Lin_et_al._2020_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":["Lin et al. 2020"],"Functions Text":["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)"],"Functions Label":["Background"],"Citation Start End":[[900,915]],"Functions Start End":[[502,719]]}
{"Identifier":"2018ApJ...856..136P__Burkhart_et_al._2010_Instance_2","Paragraph":"Depending on the specific driver, the characteristics of turbulence will then be imprinted within the ISM mainly as three-dimensional density and velocity fluctuations, and these fluctuations have been traditionally studied via correlation functions such as the spatial power spectrum (SPS) (e.g., Crovisier & Dickey 1983), \u0394-variance (e.g., Stutzki et al. 1998), and structure function (e.g., Padoan et al. 2002; Burkhart et al. 2015b). In particular, the SPS approach has been applied to observations of various Galactic and extragalactic environments (e.g., Plume et al. 2000; Dickey et al. 2001; Elmegreen et al. 2001; Burkhart et al. 2010; Combes et al. 2012; Zhang et al. 2012; Pingel et al. 2013), showing power spectral slopes \u03b2 roughly ranging from \u22122.7 to \u22123.7 depending on the tracers used (e.g., H i, carbon monoxide (CO), and dust). These slopes essentially provide information on the relative amount of structure as a function of spatial scale and can be compared with theoretical models of turbulence (mainly numerical simulations) to characterize turbulence cascade (e.g., Burkhart et al. 2010), to determine the influence of shocks (e.g., Beresnyak et al. 2005), to reveal the injection and dissipation scales of turbulent energy (e.g., Kowal & Lazarian 2007; Federrath & Klessen 2013; Chen et al. 2015), and to trace the evolution of MCs (e.g., Burkhart et al. 2015a). The proximity and abundance of multi-wavelength observations make MCs in the solar neighborhood an ideal laboratory for probing the impact of turbulence on their formation and evolution. In this paper, we focus on the Perseus MC, which is a nearby (\u223c300 pc; e.g., Herbig & Jones 1983; \u010cernis 1990), low-mass (\u223c2 \u00d7 104 M\u2299; e.g., Sancisi et al. 1974; Lada et al. 2010) cloud. Its star formation activities, as well as atomic and molecular gas content, have been extensively examined over the past decade (e.g., Ridge et al. 2006; J\u00f8rgensen et al. 2007; Pineda et al. 2008; Lee et al. 2012, 2014, 2015; Mercimek et al. 2017), revealing that the cloud consists of several individual dark and star-forming regions (e.g., B5, B1, B1E, IC 348, and NGC 1333) and is actively forming low- to intermediate-mass stars (see Bally et al. 2008 for a review).","Citation Text":["Burkhart et al. 2010"],"Functions Text":["These slopes essentially provide information on the relative amount of structure as a function of spatial scale and can be compared with theoretical models of turbulence (mainly numerical simulations) to characterize turbulence cascade (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[1089,1109]],"Functions Start End":[[846,1088]]}
{"Identifier":"2018AandA...617A..86L__Tian_2017_Instance_2","Paragraph":"The IRIS spectra measure the flare in a \u201csit-and-stare\u201d mode with a roll angle of 45\u2218. The spectral scale is \u223c25.6 m\u00c5 per pixel in the far-ultraviolet (FUV) wavelengths. The IRIS slit crosses the flaring loop and one ribbon (Fig. 1). Two red bars enclose the flaring loop region used to study the quasi-periodic oscillations in this work. IRIS spectrum was pre-processed with the SSW routines of \u201ciris_orbitval_corr_l2.pro\u201d (Tian et al. 2014; Cheng et al. 2015) and \u201ciris_prep_despike.pro\u201d (De Pontieu et al. 2014). To improve the signal-to-noise ratio, we apply a running average over five pixels to the IRIS spectra along the slit (Tian et al. 2012, 2016). We also manually perform the absolute wavelength calibration using a relatively strong neutral line, O\u202fI 1355.60 \u00c5 (see De Pontieu et al. 2014; Tian et al. 2015; Tian 2017). IRIS observations show that Fe\u202fXXI 1354.08 \u00c5 is a hot (\u223c11 MK) and broad emission line and is always blended with many narrow chromospheric lines at the flaring ribbons (Li et al. 2015b, 2016b; Tian et al. 2015; Young et al. 2015; Brosius et al. 2016; Polito et al. 2016). However, the Fe\u202fXXI 1354.08 \u00c5 line is much stronger than those blended emission lines at the flaring loops (Tian et al. 2016). Figure 2a gives the time evolution of the line profiles of Fe\u202fXXI 1354.08 \u00c5, averaged over the slit positions between \u223c18.3\u2033 and 21.6\u2033. Figure 2 panels b\u2212f show the spectral line profiles at the time indicated by the yellow lines in panel a. We can see that only the cool line of C\u202fI 1354.29 \u00c5 is blended with the hot line of Fe\u202fXXI 1354.08 \u00c5, but its contribution is negligible. Therefore, double Gaussian functions superimposed on a linear background are used to fit the IRIS spectra at \u201cO\u202f\u202fI\u201d window (Tian et al. 2016). Next, we can extract the hot line of Fe\u202fXXI 1354.08 \u00c5, as shown by the turquoise profile. The purple profile is the cool line of C\u202fI 1354.29 \u00c5. Two orange peaks represent the cool lines of O\u202fI 1354.60 \u00c5 and C\u202fI 1354.84 \u00c5 (Tian 2017), which are far away from the flaring line of Fe\u202fXXI 1354.08 \u00c5. Finally, the line properties of Fe\u202fXXI 1354.08 \u00c5 are extracted from the fitting results, that is, Doppler velocity, peak intensity, and line width (Li et al. 2016b; Tian et al. 2016; Tian & Chen 2018).","Citation Text":["Tian 2017"],"Functions Text":["Two orange peaks represent the cool lines of O\u202fI 1354.60 \u00c5 and C\u202fI 1354.84 \u00c5","which are far away from the flaring line of Fe\u202fXXI 1354.08 \u00c5."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1978,1987]],"Functions Start End":[[1900,1976],[1990,2051]]}
{"Identifier":"2016MNRAS.461.1719C__Harris_et_al._2012_Instance_1","Paragraph":"HATLAS12-00 had already been identified as a candidate gravitationally lensed galaxy as a result of its high submm flux (i.e. F500 > 100 mJy), red Herschel colours and the lack of a bright optical or radio counterpart (see e.g. Negrello et al. 2010 for a discussion of the selection of lens candidates in H-ATLAS and other Herschel surveys). This source was therefore observed spectroscopically in the submm. A CO spectroscopic redshift of 3.26 was first suggested by Z-spec (Bradford et al. 2004) observations, then subsequently confirmed by observations by the CARMA interferometer (Leeuw et al., in preparation) and the Zpectrometer instrument (Harris et al. 2007) on the Greenbank Telescope (Harris et al. 2012; see also Fu et al. 2012). Additional followup observations in the optical, near-IR, submm and other wavelengths were targeted at the lensed z = 3.26 source and the foreground objects responsible for the lensing, resulting in detailed analyses of this lensing system by Fu et al. (2012) and Bussmann et al. (2013). Their conclusions are that the z = 3.26 source HATLAS12-00 is subject to gravitational lensing, with a magnification of 9.6 \u00b1 0.5 in both the submm continuum and CO, and 16.7 \u00b1 0.8 in the K\u2032 band, by two foreground galaxies, one at a spectroscopically determined redshift of 1.22, and another with photometry suggesting that it lies at a similar redshift. The submm photometry of HATLAS12-00 at 890\u2009\u03bcm acquired with the Submillimeter Array (SMA) as part of this programme (Fu et al. 2012; Bussmann et al. 2013) is fully consistent with the 870\u2009\u03bcm and 850\u2009\u03bcm fluxes derived for this source from the LABOCA and SCUBA2 data to be presented here. The spectral energy distribution (SED) of the lensed source, after correcting for the lensing amplification, is well matched by the optically thick SED model for Arp220 from Rangwala et al. (2011), with a lensing-corrected far-IR luminosity of 1.2 \u00b1 0.2 \u00d7 1013\u2009L\u2299, and an implied star formation rate of 1400 \u00b1 300 \u2009M\u2299 yr\u22121. In many ways the unlensed properties of this object match those of the broader population of bright submm selected galaxies first discovered by the SCUBA submm imager (see e.g. Chapman et al. 2005; Clements et al. 2008; Micha\u0142owski, Hjorth & Watson 2010). The unlensed 870\u2009\u03bcm flux of this object would be \u223c7.7 mJy.","Citation Text":["Harris et al. 2012"],"Functions Text":["A CO spectroscopic redshift of 3.26 was first suggested by Z-spec","observations, then subsequently confirmed by observations","and the Zpectrometer instrument","on the Greenbank Telescope"],"Functions Label":["Similarities","Similarities","Similarities","Similarities"],"Citation Start End":[[696,714]],"Functions Start End":[[409,474],[498,555],[615,646],[668,694]]}
{"Identifier":"2022MNRAS.513.3458B__Robertson_et_al._2019_Instance_1","Paragraph":"Among the most viable mechanisms of cusp-core transformation that require changes to the assumed cosmogony is one that was proposed specifically as a possible solution to the cusp-core problem. It proposes that the DM is in fact not collisionless but self-interacting (SIDM; Spergel & Steinhardt 2000; Yoshida et al. 2000; Dav\u00e9 et al. 2001; Col\u00edn et al. 2002; Vogelsberger, Zavala & Loeb 2012; Rocha et al. 2013; see Tulin & Yu 2018 for a review). In SIDM, particles can exchange energy and momentum through elastic scattering, causing an outside-in energy redistribution within the centre of DM haloes, resulting in the formation of an isothermal core. The time-scale on which an initially cuspy SIDM halo forms a flat and isothermal core is roughly given by the time it takes for each DM particle in the inner halo to scatter at least once (Vogelsberger et al. 2012; Rocha et al. 2013). The strength of the self-interaction in SIDM models is parametrized in terms of the momentum transfer cross-section per unit mass, \u03c3T\/m\u03c7. Depending on the specific SIDM model, \u03c3T\/m\u03c7 can either be constant or dependent on the relative velocity between the two scattering DM particles. SIDM is an efficient mechanism of cusp-core transformation in dwarf-size haloes for $\\sigma _T\/m_\\chi \\gtrsim 1\\, {\\rm cm^2g^{-1}}$, whereas SIDM haloes are virtually indistinguishable from CDM haloes if $\\sigma _T\/m_\\chi \\lesssim 0.1\\, {\\rm cm^2g^{-1}}$ (Zavala, Vogelsberger & Walker 2013). The most stringent and precise constraints on the self-interaction cross-section have been put on the scales of galaxy clusters (e.g. Robertson, Massey & Eke 2017; Robertson et al. 2019) and large elliptical galaxies (Peter et al. 2013), where observations require that $\\sigma _T\/m_chi \\lesssim 1\\, {\\rm cm^2g^{-1}}$. On smaller scales, Read, Walker & Steger (2018) concluded that $\\sigma _T\/m_\\chi \\lesssim 0.6\\, {\\rm cm^2g^{-1}}$, based on their findings that the central density profile of the MW dwarf spheroidal galaxy Draco is cuspy (see also the SIDM results of Valli & Yu 2018). Moreover, based on a DM only analysis of the updated too-big-to-fail problem, Zavala et al. (2019) concluded that SIDM models with a constant cross-section of $\\sigma _T\/m_\\chi \\sim 1\\, {\\rm cm^2g^{-1}}$ fail to explain the apparently large central densities of the host haloes of the ultra-faint satellites of the MW (Errani, Pe\u00f1arrubia & Walker 2018). It should be pointed out that the constraints on \u03c3T\/m\u03c7 on the scale of dwarf galaxies are affected by significantly larger systematic uncertainties than on the scales of galaxy clusters or elliptical galaxies. Moreover, Zavala et al. (2019) demonstrate that SIDM with a strongly velocity-dependent self-interaction cross-section may provide a natural explanation for the observed diversity in the rotation curves of the MW dwarf spheroidals (see also Correa 2021). The strong dependence of the self-interaction cross-section on the typical DM velocities would create a bimodal distribution of rotation curves in the MW satellites in which the heavier haloes have constant density cores while the lighter haloes have undergone gravothermal collapse and have very steep central cusps as a consequence. The same mechanism of gravothermal collapse might be accelerated by tidal interactions in the environment of the MW leading to an agreement between constant cross-section SIDM models with $\\sigma _T\/m_\\chi \\sim 3\\, {\\rm cm^2g^{-1}}$ and the internal kinematics of MW satellites (e.g. Kahlhoefer et al. 2019; Sameie et al. 2020).","Citation Text":["Robertson et al. 2019"],"Functions Text":["The most stringent and precise constraints on the self-interaction cross-section have been put on the scales of galaxy clusters (e.g.","where observations require that $\\sigma _T\/m_chi \\lesssim 1\\, {\\rm cm^2g^{-1}}$."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1630,1651]],"Functions Start End":[[1466,1599],[1704,1784]]}
{"Identifier":"2022AandA...666A.107S__BeyondPlanck_2022_Instance_1","Paragraph":"While new data will refine our understanding of the interplanetary medium, it is necessary for future models to be consistent with both new and archival data. It is precisely this need that motivates the COSMOGLOBE project1, which aims to create a framework that will allow for the refinement of astrophysical models jointly with the raw data from complementary experiments. This form of joint analysis is already being explored within the framework of WMAP (Watts et al. 2022) and LiteBIRD (Aurlien et al., in prep.), in combination with Planck Low Frequency Instrument data (BeyondPlanck 2022). ZE is an especially promising direction for joint analysis, in part because HFI made observations of ZE at complementary wavelengths to DIRBE. While the DIRBE model was modified during the HFI analysis, no attempt was made to improve upon the geometrical representation of the model components using the larger effective dataset. Our understanding of the interplanetary medium has improved since the development of the DIRBE IPD model (see, e.g., Reach et al. 1997; Reach 2010a,b). The AKARI satellite, which observed in the infrared at wavelengths between 6 and 180 \u00b5m, detected small-scale structures in the ZE which were not well-characterized by the DIRBE model (Ootsubo et al. 2016). While the DIRBE model has been successful in describing the large-scale diffuse ZE, modern high-resolution high-frequency and infrared experiments will require IPD models that can more effectively resolve the small-scale structures in the ZE. As such, an update of the community state-of-art ZE model is long overdue. In this COSMOGLOBE framework, such a model refinement is a natural byproduct of a joint analysis of DIRBE, HFI, and other data. To improve upon this model, it is essential to have the model agree with the data at all wavelengths. The modeling of ZE is in the process of being implemented in the Commander framework, which will ultimately be used for the joint processing of HFI and DIRBE data. However, a stand-alone code is useful for agile model building and data analysis. As such, ZodiPy will function as the first step toward this goal in the COSMOGLOBE framework, in addition to making ZE corrections on arbitrary data more accessible to the community.","Citation Text":["BeyondPlanck 2022"],"Functions Text":["This form of joint analysis is already being explored within the framework of WMAP","and LiteBIRD","in combination with Planck Low Frequency Instrument data"],"Functions Label":["Background","Background","Background"],"Citation Start End":[[577,594]],"Functions Start End":[[375,457],[478,490],[519,575]]}
{"Identifier":"2017AandA...599A..97H__Carretta_et_al._2009b_Instance_1","Paragraph":"Amongst the oldest stellar systems known to exist in the Milky Way (MW) are metal-poor globular clusters (GCs). These accumulations of stars do not seem to have undergone substantial star formation for extended periods. Given the limited quality of the available data, for a long time color-magnitude diagrams (CMDs) of GCs appeared to be narrow and could be readily described by a single isochrone. These observations have justified the establishment of the long-lasting paradigm that considers CGs as prime examples of simple stellar populations (SSPs), that is, the results of very short bursts of star formation in their natal clouds. However, improved photometric precision indicates the presence of sub-populations in the cluster CMDs that are inconsistent with the SSP assumption, for a number of luminous GCs in a variety of bandpasses. Thus, early detections of chemical abundance variations (e.g., Cohen 1978) could be more easily explained in a scenario involving several populations. Moreover, in recent years evidence has grown supporting the statement that GCs are generally composed of two or three chemically distinct populations. These subpopulations are separated by a few tens to hundreds of Myr in age and show vastly varying abundances of light elements such as C, N, O, Na, Mg, and Al (see, e.g., Carretta et al. 2009b; Gratton et al. 2012, and references therein). Theoretical considerations (see, e.g., D\u2019Ercole et al. 2008, 2011) imply that GCs could have lost the majority of the initial stellar content of the first population, which consequently should have ended up in the Galactic halo. In fact, numerous studies found metal-poor GCs to be consistent with the abundance trends of the MW halo at equally low metal content (e.g., Pritzl et al. 2005; Koch et al. 2009; Koch & McWilliam 2014; Villanova et al. 2016). We address this scenario by adding NGC 6426 to the short list of metal-poor clusters with available information on detailed chemical abundances. There are only two GCs in the Harris catalog (Harris 1996, 2010 edition) more metal poor than NGC 6426. At 12.9 \u00b1 1.0 Gyr, the cluster is the oldest in the age compilation by Salaris & Weiss (2002). At a galactocentric distance of Rgc = 14.4 kpc and a galactic latitude of 16.23\u00b0 it is located in the transition region between inner and outer halo. Previous studies found consistent [Fe\/H]1 values: \u22122.20 \u00b1 0.17 dex (Zinn & West 1984), \u22122.33 \u00b1 0.15 (Hatzidimitriou et al. 1999), and \u22122.39 \u00b1 0.04 dex (Dias et al. 2015). The latter value originates from the very first spectroscopic analysis of NGC 6426 at low resolution, which also stated [Mg\/Fe] = 0.38 \u00b1 0.06 dex. To date, there has been no study further addressing the detailed metal content of this cluster. ","Citation Text":["Carretta et al. 2009b"],"Functions Text":["Moreover, in recent years evidence has grown supporting the statement that GCs are generally composed of two or three chemically distinct populations. These subpopulations are separated by a few tens to hundreds of Myr in age and show vastly varying abundances of light elements such as C, N, O, Na, Mg, and Al (see, e.g.,"],"Functions Label":["Background"],"Citation Start End":[[1319,1340]],"Functions Start End":[[996,1318]]}
{"Identifier":"2021MNRAS.500.2336Y___2018_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":["Vu\u010deti\u0107 et al.","2018"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[1324,1338],[1349,1353]],"Functions Start End":[[1041,1171]]}
{"Identifier":"2020ApJ...891...10L__Song_et_al._2009_Instance_1","Paragraph":"We need to build data sets to train and evaluate the model. Then the data sets are prepared in the following way: (1) if the AR does not flare within 24 hr after the observation time, the No-flare (weaker than C1.0) label is assigned to the magnetogram sample in the same AR. (2) If the C\/M\/X-level flare occurs within 24 hr after the observation time, the corresponding flare label (i.e., C, M, or X) is assigned to the magnetogram sample. Note that there are a number of ARs producing recurring flares with different flare levels within 24 hr. For the first flare of one AR, the corresponding flare label is assigned to the magnetogram sample within 24 hr after the observation time. Then for the following flares of the same AR, the corresponding labels are assigned to the magnetogram sample in the period from the end of its prior flare to the end of this flare. (3) We adopt a four-level AR classification scheme based on the maximum GOES-level flare an AR ever yields (Song et al. 2009; Yuan et al. 2010; Liu et al. 2017). In other words, ARs are further categorized into four levels (i.e., No-flare, C, M, and X) if they yield at least one flare with such GOES-level criterion but no flares with higher GOES-level criterion: \u201cLevel = X\u201d indicates that an AR yields at least one X-level flare; \u201cLevel = M\u201d indicates that an AR yields at least one M-level flare but no X-level flares; \u201cLevel = C\u201d indicates that an AR yields at least one C-level flare but no M\/X-level flares; \u201cLevel = No-flare\u201d indicates that an AR only yields microflares (weaker than C1.0 flares). Finally, we gather 870 ARs and 136134 magnetogram samples in total, including 443 X-level, 6534 M-level, 72412 C-level, and 56745 No-flare level samples. Note that the magnetogram samples with multiple ARs (Bobra et al. 2014) are not included in our work. For the \u2265M class, magnetogram samples of M\/X-level flare in an AR are defined as positive class, while all the others are defined as negative class. For the \u2265C class, magnetogram samples of C\/M\/X-level flare in an AR are defined as positive class, while all the others are defined as negative class.","Citation Text":["Song et al. 2009"],"Functions Text":["We adopt a four-level AR classification scheme based on the maximum GOES-level flare an AR ever yields"],"Functions Label":["Uses"],"Citation Start End":[[976,992]],"Functions Start End":[[872,974]]}
{"Identifier":"2021MNRAS.506.1962S__Nordlander_&_Lind_2017_Instance_1","Paragraph":"Among the odd Z elements, we derived the abundances of Na, Al, K, and Sc using the high-resolution spectra and the details of the lines are given in Table A1. The Na D lines were not used as the lines are too strong for deriving the abundances. So, we depended on the weaker lines 5682 \u00c5 and 5688 \u00c5 to derive the Na abundance. The NLTE corrections for these Na i lines were performed from Lind et al. (2011) whenever available and the average values are listed in Table 3. Al i lines at 6696 and 6698 are used for deriving the Al abundance and these lines have negligible contributions from NLTE effects (Baumueller & Gehren 1997; Nordlander & Lind 2017). For the case of K, whenever the feature at 7698.98 \u00c5 has contributions from telluric features, the abundance from 7664.87 \u00c5 is quoted, otherwise the average of the abundances from both the lines is quoted in Table 3. Both these lines are sensitive to NLTE effects (Takeda et al. 2002; Kobayashi et al. 2006; Andrievsky et al. 2010; Prantzos et al. 2018; Reggiani et al. 2019) and the NLTE corrections depend on the effective temperature and surface gravity of the model. Using the NLTE grid provided in Reggiani et al. (2019), we have estimated the NLTE corrections for the program stars and the NLTE corrected values are given in Table 3. The NLTE grid does not cover very low log g values, so we could not estimate the NLTE corrections for HE 1152\u22120355 and HE 0314\u22120143. We could not measure the K abundance in HD 5223 due to the contamination from telluric features. For the case of Sc abundance, we could not find any study of Sc line formation in NLTE except in Zhang, Gehren & Zhao (2008) where they have studied the NLTE effects only for the sun and identified negligible NLTE effects for Sc ii lines. So, we report only LTE results for this element in Table 3. The resolution of the NIR spectra was too low to resolve the HF feature from the nearby features, so we could not derive F abundance in these stars.","Citation Text":["Nordlander & Lind 2017"],"Functions Text":["Al i lines at 6696 and 6698 are used for deriving the Al abundance and these lines have negligible contributions from NLTE effects"],"Functions Label":["Uses"],"Citation Start End":[[631,653]],"Functions Start End":[[473,603]]}
{"Identifier":"2015AandA...584A..75V__Essen_et_al._(2014)_Instance_4","Paragraph":"The data presented here comprise quasi-simultaneous observations during secondary eclipse of WASP-33 b around the V and Y bands. The predicted planet-star flux ratio in the V-band is 0.2 ppt, four times lower than the accuracy of our measurements. Therefore, we can neglect the planet imprint and use this band to measure the stellar pulsations, and most specifically to tune their current phases (see phase shifts in von Essen et al. 2014). Particularly, our model for the light contribution of the stellar pulsations consists of eight sinusoidal pulsation frequencies with corresponding amplitudes and phases. Hence, to reduce the number of 24 free parameters and the values they can take, we use prior knowledge about the pulsation spectrum of the star that was acquired during von Essen et al. (2014). As the frequency resolution is 1\/\u0394T (Kurtz 1983), 3.5 h of data are not sufficient to determine the pulsations frequencies. Therefore, during our fitting procedure we use the frequencies determined in von Essen et al. (2014) as starting values plus their derived errors as Gaussian priors. As pointed out in von Essen et al. (2014), we found clear evidences of pulsation phase variability with a maximum change of 2 \u00d7 10-3 c\/d. In other words, as an example after one year time a phase-constant model would appear to have the correct shape with respect to the pulsation pattern of the star, but shifted several minutes in time. To account for this, the eight phases were considered as fitting parameters. The von Essen et al. (2014) photometric follow-up started in August, 2010, and ended in October, 2012, coinciding with these LBT data. We then used the phases determined in von Essen et al. (2014) during our last observing season as starting values, and we restricted them to the limiting cases derived in Sect. 3.5 of von Essen et al. (2014), rather than allowing them to take arbitrary values. The pulsation amplitudes in \u03b4 Scuti stars are expected to be wavelength-dependent (see e.g. Daszy\u0144ska-Daszkiewicz 2008). Our follow-up campaign and these data were acquired in the blue wavelength range. Therefore the amplitudes estimated in von Essen et al. (2014), listed in Table 1, are used in this work as fixed values. This approach would be incorrect if the pulsation amplitudes would be significantly variable (see e.g., Breger et al. 2005). Nonetheless, the short time span of LBT data, and the achieved photometric precision compared to the intrinsically low values of WASP-33\u2019s amplitudes, make the detection of any amplitude variability impossible. ","Citation Text":["von Essen et al. (2014)"],"Functions Text":["As pointed out in","we found clear evidences of pulsation phase variability with a maximum change of 2 \u00d7 10-3 c\/d. In other words, as an example after one year time a phase-constant model would appear to have the correct shape with respect to the pulsation pattern of the star, but shifted several minutes in time."],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[1114,1137]],"Functions Start End":[[1096,1113],[1139,1433]]}
{"Identifier":"2016MNRAS.460.3472E__Ercolano_et_al._2008a_Instance_1","Paragraph":"We use the set of wind solutions (density and velocity distribution of gas in the wind) for primordial discs (i.e. gas-rich, optically thick discs, which do not have an evacuated inner cavity) calculated by Owen et al. (2010, 2011) and EO10 for a 0.7 M\u2299 star and X-ray luminosities (0.1 \u2264 h\u03bd \u2264 10 keV) of LX = 2 \u00d7 1028, 2 \u00d7 1029 and 2 \u00d7 1030 erg s\u22121. These were obtained by means of two-dimensional hydrodynamic calculations using the zeus code (Stone & Norman 1992; Stone, Mihalas & Norman 1992; Hayes et al. 2006), modified to include the effects of X-ray irradiation with a parametrization of the gas temperature as a function of the local ionization parameter. The dust radiative transfer and photosionisation code mocassin (Ercolano et al. 2003; Ercolano, Barlow & Storey 2005; Ercolano et al. 2008a), modified according to Ercolano et al. (2008b), was used produce the temperature parametrization. The atomic data base of the mocassin code included opacity data from Verner et al. (1993) and Verner & Yakovlev (1995), energy levels, collision strengths and transition probabilities from Version 5.2 of the CHIANTI data base (Landi et al. 2006, and references therein) and hydrogen and helium free\u2013bound continuous emission data of Ercolano & Storey (2006). The ionizing spectrum used to calculate the temperature parametrization was calculated by Ercolano et al. (2009), using the plasma code of Kashyap & Drake (2000) from an emission measure distribution based on that derived for RS CVn type binaries by Sanz-Forcada, Brickhouse & Dupree (2002), which peaks at 104 K and fits to Chandra spectra of T-Tauri stars by Maggio et al. (2007), which peaks at around 107.5 K. This spectrum has a significant EUV component (13.6\u2009eV \u2264 h\u03bd \u2264 0.1\u2009keV), with roughly LEUV = LX. Solar abundances (Asplund, Grevesse & Sauval 2005), depleted according to Savage & Sembach (1996) were assumed, namely (number density, with respect to hydrogen): He\/H = 0.1,\u2009C\/H = 1.4 \u00d7 104,\u2009N\/H = 8.32 \u00d7 105,\u2009O\/H = 3.2 \u00d7 104,\u2009Ne\/H = 1.2 \u00d7 104,\u2009Mg\/H = 1.1 \u00d7 106,\u2009Si\/H = 1.7 \u00d7 106,\u2009S\/H = 2.8 \u00d7 105. More details about the codes and setup of the models can be found in Ercolano et al. (2008b, 2009) and Owen et al. (2010).","Citation Text":["Ercolano et al. 2008a"],"Functions Text":["The dust radiative transfer and photosionisation code mocassin","was used produce the temperature parametrization. The atomic data base of the mocassin code included opacity data from Verner et al. (1993) and Verner & Yakovlev (1995), energy levels, collision strengths and transition probabilities from Version 5.2 of the CHIANTI data base (Landi et al. 2006, and references therein) and hydrogen and helium free\u2013bound continuous emission data of Ercolano & Storey (2006)."],"Functions Label":["Uses","Uses"],"Citation Start End":[[783,804]],"Functions Start End":[[665,727],[854,1262]]}
{"Identifier":"2019MNRAS.487.1626Q__Merloni,_Fabian_&_Ross_2000_Instance_1","Paragraph":"Low-mass X-ray binaries (LMXBs) which either contain a black hole (BH) or a neutron star (NS), accreting matter from its low-mass companion star (\u22721M\u2299) are ideal natural laboratories for studying the physics of accretion and jet around a BH or an NS (e.g. Migliari & Fender 2006). According to the timing and spectral features in the X-ray band, LMXBs are generally divided into two main spectral states, i.e. the high\/soft state and the low\/hard state (Gilfanov 2010, for review). For BH-LMXBs, when they are in the high\/soft state, the accretion flow is widely believed to be dominated by the optically thick, geometrically thin, cool accretion disc (Shakura & Sunyaev 1973), and the X-ray spectrum can be well described by a multicolour blackbody spectrum (e.g. Mitsuda et al. 1984; Makishima et al. 1986; Merloni, Fabian & Ross 2000). Whereas, for NS-LMXBs, besides the emission from the disc, there is a significant thermal emission from the boundary layer between the accretion disc and the surface of the NS (Popham & Narayan 1992; Inogamov & Sunyaev 1999; Popham & Sunyaev 2001; Gilfanov & Sunyaev 2014, for review). Observationally, for both BH-LMXBs and NS-LMXBs, when they are in the low\/hard state, generally, the accretion flow is suggested to be dominated by the optically thin, geometrically thick, hot, advection-dominated accretion flow (ADAF) (Done, Gierli\u0144ski & Kubota 2007, for review). Theoretically, the ADAF solution has been studied in detail by several researchers since it was discovered in 1970\u2019s (Ichimaru 1977; Rees et al. 1982; Narayan & Yi 1994, 1995a,b; Abramowicz et al. 1995; Chen et al. 1995; Yuan & Narayan 2014, for review). In the BH case, the ADAF solution is a kind of radiatively inefficient accretion flow, in which a fraction of the viscously dissipated energy will be advected into the event horizon of the BH. While in the NS case, the viscously dissipated energy advected onto the surface of the NS will eventually be radiated out, so the ADAF solution is radiatively efficient (Narayan & Yi 1995b). Qiao & Liu (2018b) calculated the structure and the corresponding emergent spectrum of the ADAF around a weakly magnetized NS within the framework of the self-similar solution of the ADAF. The authors compared the electron temperature of the ADAF around a NS and a BH, it is found that the electron temperature of the ADAF around an NS is systemically lower than that of a BH, which is consistent with observations (Burke, Gilfanov & Sunyaev 2017; Qiao & Liu 2018b). Meanwhile, the authors compared the Compton y-parameter (defined as $y={{4kT_{\\rm e}}\\over {m_{\\rm e}c^2}} \\rm {Max}(\\tau _{\\rm es}, \\tau ^2_{\\rm es})$, with Te being the election temperature, me being the electron mass, c being the speed of light, and \u03c4es being the Compton scattering optical depth) of the ADAF around an NS and a BH, it is found that the Compton y-parameter of the ADAF around an NS is systemically lower than that of a BH, producing a softer X-ray spectrum, which is also consistent with observations (Wijnands et al. 2015; Parikh et al. 2017; Sonbas, Dhuga & G\u00f6\u011f\u00fc\u015f 2018; Qiao & Liu 2018b).","Citation Text":["Merloni, Fabian & Ross 2000"],"Functions Text":["For BH-LMXBs","and the X-ray spectrum can be well described by a multicolour blackbody spectrum (e.g."],"Functions Label":["Background","Background"],"Citation Start End":[[809,836]],"Functions Start End":[[482,494],[678,764]]}
{"Identifier":"2019AandA...625A.121M__Beaug\u00e9_&_Nesvorn\u00fd_2012_Instance_1","Paragraph":"The final location of close-in giant planets in our models reflects the strength of the tides that we include in our modeling, which play a very important role in the decay of planetary orbits. These are dynamical tides (e.g., Lai 1997; Ivanov & Papaloizou 2004, 2007, 2011) and in our simulations we used a formulation given by Ivanov & Papaloizou (2007) as described inSect. 2. However, the impulse approximation used in the evaluation of dynamical tides becomes a poor approximation when the circularization proceeds (e.g., Mardling 1995a,b) and the eccentricity becomes low. Equilibrium tides become then effective (e.g., Beaug\u00e9 & Nesvorn\u00fd 2012 and references therein) and the tidal evolution may occur on a longer timescale. In short, at the beginning of the orbital evolution that leads to the formation of hot\/warm Jupiters, dynamical tides are important in forcing the decay of the orbit. In the last part of the dynamical evolution when the eccentricity has become low, equilibrium tides are more important in determining the location where the planet stops. Unfortunately, at present it is not known when and how the two tides switch. When we change the magnitude of two tides, the final location of the planets can be adjusted (Beaug\u00e9 & Nesvorn\u00fd 2012). However, rather than introducing artificial effects, we continue to use dynamical tides in our simulations even for low eccentricities but we stop our simulations when the energy, decreasing from the tide at the pericenter, overcomes the orbital energy leading to a clustering of tidally circularized planets around 0.02 au. However, the final distribution of the inclination of the planets does not depend on this choice and highly misaligned planets would be produced anyway. Since the tidal evolution of planets with arbitrary inclinations is still not well known, we assume that planetary inclination is not significantly changed during tidal evolution (Barker & Ogilvie 2009). Thus, the planets maintain the inclination they have when the circularization begins.","Citation Text":["Beaug\u00e9 & Nesvorn\u00fd 2012"],"Functions Text":["Equilibrium tides become then effective (e.g.,","and references therein) and the tidal evolution may occur on a longer timescale."],"Functions Label":["Uses","Uses"],"Citation Start End":[[626,648]],"Functions Start End":[[579,625],[649,729]]}
{"Identifier":"2021MNRAS.508.4332M__Sickafoose_et_al._2001_Instance_1","Paragraph":"As a first step, we need to specify the photoelectron sheath features. In this course, we first evaluate the steady state potential over the lunar surface (equation 4), and then after, we use this as a boundary condition to solve the Poisson equation (equation 2) and estimate the photoelectron sheath profile. In calculations, Lyman \u03b1 (\u03bb \u223c 121.57 nm, 10.29 eV, \u039b \u223c 3 \u00d7 1011 cm\u20132 s\u20131) spike of solar photon radiation (Bauer 1973) is considered as the dominant source for the generation of photoelectrons from the lunar surface. The work function of the regolith material is taken from Grobman & Blank (1969), where it is suggested to vary in the range \u03d5r \u223c (4\u20136) V for the region across the subsolar point and limb. Moreover, Draine\u2019s formulation is accounted to determine the lunar surface\u2019s photoelectric efficiency (Draine 1978; Draine & Salpeter 1979) \u2013 its spectral dependence can be represented as ${\\chi _{\\nu r}} = {\\chi _o}[1 - ({\\phi _r}\/{E_\\nu })]$. For instance, for the Lyman \u03b1 radiation \u03c7\u03bdr = 0.042 for optimum efficiency \u03c7o = 0.1 (Sickafoose et al. 2001) and \u03d5r = 6 V (Grobman & Blank 1969). Another significant parameter is the surface temperature which describes the electron population within the lattice available for the photoemission. Lunar Reconnaissance Orbiter based measurements (Williams et al. 2017) suggest that the surface temperature may vary from the equator (\u223c400 K) to the terminator (poles, \u223c150 K). In order to take this account, we use the latitude (\u03b8) dependent empirical relation ${T_\\theta } = {T_0}[1 - (5\/4\\pi )\\theta ]$; for instance, at \u03b8 = 70\u00b0, and To \u2248 205 K. These three parameters, viz., \u03d5r, To, \u03c7\u03bd, and \u039b drive the photoemission current from the lunar regolith. The nominal solar wind plasma parameters are considered for calculating collection current over lunar regolith; the constituents are considered as40-41nes \u2248 nis = 8.7 cm\u20133 and Tes \u2248 Tis = 1.4 \u00d7 105 K (Mann et al. 2011; Kureshi et al. 2020). These solar radiation and wind plasma parameters might vary widely during active solar events and alter surface charging and sheath features. Popel et al. (2018) suggest the dust number density may also vary in a wide range depending on lunar altitude and particle size; for instance, nd \u223c 800 cm\u20133 for the particles of size 100 nm \u2264 ao \u2264 200 nm and \u03b8 = 77\u00b0. Note that the secondary electron emission (Seitz 1940; Misra, Mishra & Sodha 2013) from the lunar regolith (and floating dust) is ignored, as it minimally contributes to the charging of sunlit surfaces (Mishra & Bhardwaj 2020). These parameters, along with equation (4), yield steady state potential over the sunlit locations. This estimate of the surface potential is used as a boundary condition (i.e. at l = 0, \u03c5 = \u03c5o) alongwith \u03c5\u2019 = \u03c5 = 0 as $l \\to \\infty $ to solve the Poisson equation (equation 2) numerically \u2013 using this framework, the sheath structure is derived in terms of electric potential (\u03c5), electric field (Es), and photoelectron population density (npe).","Citation Text":["Sickafoose et al. 2001"],"Functions Text":["For instance, for the Lyman \u03b1 radiation \u03c7\u03bdr = 0.042 for optimum efficiency \u03c7o = 0.1"],"Functions Label":["Uses"],"Citation Start End":[[1046,1068]],"Functions Start End":[[961,1044]]}
{"Identifier":"2019ApJ...883..130L__Hansen_2009_Instance_1","Paragraph":"(3) What is the likelihood of obtaining Mercury and Mars analogs in systems with Venus\u2013Earth pair analogs? Recent terrestrial planet formation studies managed to statistically produce low-mass Mars analogs. In particular, there are currently five main competing models. (1) Grand Tack: the protoplanetary disk is truncated at \u223c1 au by perturbations of an inward-then-outward gas-driven migrating Jupiter within the first 1\u20133 Myr of the solar system history, leaving a disk with mass concentrated within that distance after the disk gas dispersal (Walsh et al. 2011; Jacobson & Morbidelli 2014; Walsh & Levison 2016; Brasser et al. 2016a); (2) Empty Asteroid Belt: embryos and planetesimals in the protoplanetary disk formed concentrated within a narrow belt at \u223c0.7\u20131 au (Hansen 2009; Drazkowska et al. 2016; Raymond & Izidoro 2017; Ogihara et al. 2018); (3) Early Instability: the protoplanetary disk is perturbed by the giant planets\u2019 instability that occurred within \u223c10 Myr after the disk gas dispersal, strongly depleting the disk mass beyond \u223c1.3 au (Clement et al. 2018, 2019b); (4) Pebble Accretion: embryos and planetesimals form preferentially in the inner and outer regions of the protoplanetary disk, respectively, and the disk mass is concentrated within \u223c1.5 au (Levison et al. 2015; Chambers 2016); (5) Sweeping Secular Resonance: the growth of embryos and planetesimals is inhibited beyond \u223c1\u20131.5 au by perturbations of secular resonances that swept the disk during the disk gas dispersal (Bromley & Kenyon 2017). In all these models, Mars analogs may form as a result of the disk mass depletion or absence of mass beyond \u223c1\u20131.5 au. Additionally, from the standpoint of the initial conditions of the disk after gas dispersal, models 1 and 2 are very similar, while models 3\u20135 probably also share similar properties. However, in the majority of those studies, it remains unclear whether these Mars analogs were obtained in systems that also contained Venus and Earth analogs. Moreover, even if they were, it is also unclear what fraction these three-planet analog systems would represent compared to all the systems obtained in those studies, because their results are often presented with all planets mixed (e.g., in plots of distance vs. mass). Furthermore, Mercury and Mars analogs in those studies are often defined as planets from the mixed population that satisfy an arbitrary distance range: e.g., 0.5 au and 1.2\u20132.0 au, respectively. However, this approach can lead to incomplete classifications in a given system (e.g., by failing to identify more massive Mars analogs at 1.2 au or failing to properly identify planet analogs). Indeed, it is difficult to discriminate Venus from Earth analogs, unambiguously identify Mercury\/Mars analogs, and avoid misclassifications in past studies. If the five models described above are successful in reproducing a low-mass Mars, how can the best among them be discriminated without a proper system classification?","Citation Text":["Hansen 2009"],"Functions Text":["In particular, there are currently five main competing models.","(2) Empty Asteroid Belt: embryos and planetesimals in the protoplanetary disk formed concentrated within a narrow belt at \u223c0.7\u20131 au"],"Functions Label":["Background","Background"],"Citation Start End":[[772,783]],"Functions Start End":[[207,269],[639,770]]}
{"Identifier":"2017AandA...605A.121M__Thompson_et_al._2017_Instance_1","Paragraph":"The effects of the wet air, the variable water vapour cells, are the main cause of the refraction at submillimetre\/millimetre wavelengths. The dipole moment of water makes water vapour, the wet component in the troposphere, a strong absorber at submillimetre\/millimetre wavelengths and significantly increases the refractive index of the air. Because the water vapour is not well mixed there are localised pockets of air with different refractive indices. In what is called the \u201cfrozen-screen\u201d hypothesis (Taylor 1938), these pockets, or turbulent eddies, are assumed to be fixed in the atmospheric layer that advects over an interferometric array (Thompson et al. 2017). Thus, these cause various delays in the path length (variable in time and position) along the line of sight to each antenna. Interferometers are sensitive to the variations in path length, the interferometric phase difference, between pairs of antennas that form a baseline. For a given baseline (distance and orientation) the line-of-sight path to each component of an astronomical source has an intrinsic phase that relates the measured intensities to their location in an image. Thus any additional variable atmospheric delays that cause anomalous phase changes on many baselines making up an array have the effect of blurring the interferometric image; this is analogous to the effect of seeing at optical and infrared wavelengths. The introduced delays scale linearly with the difference in precipitable water vapour (\u0394PWV) between an antenna pair (excluding dispersive effects) and linearly with frequency. The correlated signals between pairs of antennas (the visibilities V = V0ei\u03c6) become partly decorrelated as a result of the phase noise. The reduced coherence for the visibilities is given by (1)\\begin{equation} \\label{eqn0} \\langle V \\rangle = V_0 \\times \\langle {\\rm e}^{\\rm i\\phi} \\rangle = V_0 \\times {\\rm e}^{-\\phi^2_{\\rm rms}\/2} , \\end{equation}\u27e8V\u27e9=V0\u00d7\u27e8ei\u03c6\u27e9=V0\u00d7e\u2212\u03c6rms2\/2,","Citation Text":["Thompson et al. 2017"],"Functions Text":["In what is called the \u201cfrozen-screen\u201d hypothesis","these pockets, or turbulent eddies, are assumed to be fixed in the atmospheric layer that advects over an interferometric array"],"Functions Label":["Background","Background"],"Citation Start End":[[649,669]],"Functions Start End":[[456,504],[520,647]]}
{"Identifier":"2021AandA...653A.154T__VI_2020_Instance_1","Paragraph":"The global volume-weighted neutral fraction of hydrogen QH\u202fI in the high-resolution volume is presented as the thick blue line in Fig. 10. The reionization process starts when the first stars are born, and by z50\u2004\u2243\u20047.63, half of the volume is reionized. We identify the redshifts at which 1%, 10%, 50%, 90% and 99% of the volume is ionized as z01\u2004=\u200411.13, z10\u2004=\u20048.68, z50\u2004=\u20047.63, z90\u2004=\u20046.58, and z99\u2004=\u20045.92 respectively; corresponding to a reionization duration \u0394z\u2004=\u2004z99\u2005\u2212\u2005z10\u2004=\u20042.8 (\u0394t\u2004\u2243\u2004385\u2006Myr), broadly consistent with the estimates of Robertson et al. (2015) The dark and light shaded areas in Fig. 10 correspond to the 1\u03c3 and 2\u03c3 constraints on the redshift of reionization from the cosmic microwave background measurements of the Planck mission (Planck Collaboration VI 2020), with a reionization midpoint zre\u2004=\u20047.67\u2005\u00b1\u20050.73. We also compare the OBELISK reionization history to a selection of observational constraints: Black hexagons correspond to the measurements of the Lyman-\u03b1 forest transmission (Ly\u03b1 forest, Fan et al. 2006b), the green circles show constraints on the IGM opacity from the fraction of Lyman-\u03b1 emitters in Lyman-break galaxy samples (Schenker et al. 2014; Ono et al. 2012; Pentericci et al. 2014; Robertson et al. 2013; Tilvi et al. 2014), the purple diamonds show measurements from quasar damping wings by Mortlock et al. (2011), Schroeder et al. (2013), Ba\u00f1ados et al. (2018), \u010eurov\u010d\u00edkov\u00e1 et al. (2020), the red diamonds show similar measurements on gamma-ray bursts (GRB, Totani et al. 2006, 2016), and the black squares from Ouchi et al. (2010), Ota et al. (2008) represent constraints derived from the evolution of the Lyman-\u03b1 luminosity function. Some of these data points come from the compilations of Bouwens et al. (2015). Overall, we find that the simulation agrees with most observations in terms of reionization history, despite the fact that we focus on an overdense region. Interestingly, the simulation manages to capture residual neutral fraction after reionization is complete at z\u2004\u20046 similar to what is observed. We discuss this point further below.","Citation Text":["Planck Collaboration VI 2020"],"Functions Text":["The dark and light shaded areas in Fig. 10 correspond to the 1\u03c3 and 2\u03c3 constraints on the redshift of reionization from the cosmic microwave background measurements of the Planck mission","with a reionization midpoint zre\u2004=\u20047.67\u2005\u00b1\u20050.73."],"Functions Label":["Uses","Uses"],"Citation Start End":[[752,780]],"Functions Start End":[[564,750],[783,830]]}
{"Identifier":"2022AandA...666L...5G__Esparza-Arredondo_et_al._2018_Instance_1","Paragraph":"More recently, Garc\u00eda-Bernete et al. (2022) found that the PAH molecules responsible for the 11.3 \u03bcm PAH emission band are more resilient in the hard environments often present in AGN. In particular, the authors found larger 11.3\/7.7 \u03bcm and 11.3\/6.2 \u03bcm PAH ratios in AGN-dominated systems compared to SF galaxies, indicating a larger fraction of neutral PAH molecules (as noted by Smith et al. 2007 using a sample of relatively weak AGN). However, these studies were limited by the spatial resolution (\u223c4\u2033) and the low spectral resolution (R\u2004\u223c\u200460\u2013130) of Spitzer\/InfraRed Spectrograph (IRS). Previous sub-arcsecond angular resolution N-band (\u223c8\u201313 \u03bcm) ground-based spectroscopic studies investigated the 11.3 \u03bcm PAH feature in the nuclear and circumnuclear regions of AGN (e.g., H\u00f6nig et al. 2010; Gonz\u00e1lez-Mart\u00edn et al. 2013; Alonso-Herrero et al. 2014, 2016; Ramos et al. 2014; Esquej et al. 2014; Garc\u00eda-Bernete et al. 2015; Jensen et al. 2017; Esparza-Arredondo et al. 2018). However, these works were unable to provide definitive details regarding the effect of the AGN on the PAH molecules due to limited wavelength coverage and sensitivity. The changes in the PAH properties due to the presence of the AGN might be more prominent in their innermost regions of galaxies. Therefore, the unprecedented combination of high angular and spectral resolution (R\u2004\u223c\u20041500\u2005\u2212\u20053500) in the entire mid-IR range (4.9\u201328.1 \u03bcm) afforded by the James Webb Space Telescope (JWST)\/Mid-Infrared Instrument (MIRI; Rieke et al. 2015; Wells et al. 2015; Wright et al. 2015) is key to investigating PAH properties. In this Letter we report on the first investigation of PAH emission in the nuclear regions of three luminous Seyfert (Sy) galaxies and compare them with emission from SF regions using JWST\/MIRI Medium Resolution Spectrograph (MRS) data. This enables us, for the first time, to characterise the PAH properties of local luminous Sy galaxies (log (Lbol)> 44.46 erg s\u22121)1 at sub-arcsecond scales (\u223c0.45\u2033, \u223c142\u2013245 pc).","Citation Text":["Esparza-Arredondo et al. 2018"],"Functions Text":["Previous sub-arcsecond angular resolution N-band (\u223c8\u201313 \u03bcm) ground-based spectroscopic studies investigated the 11.3 \u03bcm PAH feature in the nuclear and circumnuclear regions of AGN","However, these works were unable to provide definitive details regarding the effect of the AGN on the PAH molecules due to limited wavelength coverage and sensitivity."],"Functions Label":["Background","Motivation"],"Citation Start End":[[948,977]],"Functions Start End":[[592,771],[980,1147]]}
{"Identifier":"2021AandA...654A..80S__Urrutia_et_al._2019_Instance_1","Paragraph":"In this study we are focusing on rest-frame UV emission red-wards of the strongest UV emission line, Ly\u03b1, partially motivated by the challenges of observing this line at high redshift (z\u2004\u2273\u20046) where the significantly neutral CGM and IGM absorbs the Ly\u03b1 photons escaping the galaxy along the line of sight (e.g., Dijkstra et al. 2011; Laursen et al. 2011, 2019; Dijkstra 2017). Nevertheless, the Ly\u03b1 line itself has improved our understanding of star-forming galaxies in the (early) Universe. In particular, the asymmetric Ly\u03b1 line profile has enabled redshift confirmations of large samples of sources at both 2\u2004\u2004z\u2004\u20046 (e.g., Steidel et al. 2014; Le Fevre et al. 2015; Herenz et al. 2017; Inami et al. 2017; Urrutia et al. 2019) and high redshift at z\u2004>\u20046 (e.g., Finkelstein et al. 2013; Oesch et al. 2015; Schmidt et al. 2016; Tilvi et al. 2016; Huang et al. 2016; Pentericci et al. 2018; Fuller et al. 2020). The resonant scattering of the photons and the resulting (occasional) multipeaked emission has been shown to relate closely to the column density and dynamics of the neutral hydrogen in the ISM and the CGM (Verhamme et al. 2015; Gazagnes et al. 2018, 2020). The fraction of galaxies with confirmed Ly\u03b1 emission has been used to probe the evolution (or lack thereof) of the fraction of LAEs among Lyman-break galaxies from low redshift to the EoR (e.g., Treu et al. 2013; Pentericci et al. 2014; Tilvi et al. 2014; de Barros et al. 2017; Caruana et al. 2018; Kusakabe et al. 2020). Together with the observed velocity offset of the Ly\u03b1 line resulting from resonant scattering (Schenker et al. 2013; Erb et al. 2014; Hashimoto et al. 2015; Stark et al. 2017; Verhamme et al. 2018), this has probed the amount of neutral gas in the IGM and has constrained the neutral fraction of the Universe during the EoR (Ouchi et al. 2010; Greig et al. 2017; Mason et al. 2018a,b; Banados et al. 2018; Hoag et al. 2019). Furthermore, comparisons between Ly\u03b1 and H\u03b1 or UV emission line strengths have been used to study the production efficiency and escape of ionizing photons from LAEs (Nakajima et al. 2016; Matthee et al. 2017; Harikane et al. 2018; Lam et al. 2019; Maseda et al. 2020). It is therefore of interest to relate and compare the measured rest-frame UV emission lines red-wards of Ly\u03b1 studied here, with the characteristics of the Ly\u03b1 line itself and the properties of the LAEs in our sample.","Citation Text":["Urrutia et al. 2019"],"Functions Text":["In particular, the asymmetric Ly\u03b1 line profile has enabled redshift confirmations of large samples of sources at both 2\u2004\u2004z\u2004\u20046 (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[706,725]],"Functions Start End":[[491,623]]}
{"Identifier":"2022ApJ...938..124Z__Tu_&_Marsch_1995_Instance_1","Paragraph":"In the context of solar wind turbulence, one crucial question concerns turbulence evolution and heating (Parashar et al. 2015; Chen 2016; Viall & Borovsky 2020). Numerous correlations exist between solar wind parameters and magnetic fluctuations, which can provide important indications and constraints on the physics of turbulence evolution and heating processes. Early observations showed that faster solar wind usually corresponds to higher proton temperature (e.g., Marsch et al. 1982; Elliott et al. 2012). The faster solar wind tends to be more imbalanced with greater wave energy flux antisunward than sunward (Tu & Marsch 1995), and can be characterized by larger inertial-scale magnetic energy density and steeper proton-scale magnetic spectrum (Bruno et al. 2014). A study by Wind observations revealed that higher proton temperature is associated with a steeper proton-scale magnetic spectrum (Leamon et al. 1998b). Statistical studies by Wind and Advanced Composition Explorer measurements later revealed that higher proton temperature is related to larger inertial-scale magnetic energy density (Smith et al. 2006; Vech et al. 2018). Recent statistical studies by Wind and Parker Solar Probe (PSP) observations further revealed that higher proton temperature is linked to larger proton-scale magnetic energy density (Zhao et al. 2020, 2022). On the other hand, cross and magnetic helicities at inertial and kinetic scales, respectively, are often employed to describe the imbalance and handedness of solar wind turbulence. A correlation between cross and magnetic helicities was reported, which was interpreted as the signature of cyclotron-resonant dissipation in solar wind turbulence (Leamon et al. 1998a). A strong correlation (with a correlation coefficient (CC) up to 0.8 occasionally) between the proton-scale spectral index and magnetic helicity was found (Pine et al. 2020; Zhao et al. 2021, 2022). Via PSP observations, a mild correlation (with a CC of 0.36) between the spectral index and cross helicity was displayed in the most recent literature (Huang et al. 2021). However, less attention has been paid to the relations between these correlations despite the large body of research.","Citation Text":["Tu & Marsch 1995"],"Functions Text":["The faster solar wind tends to be more imbalanced with greater wave energy flux antisunward than sunward"],"Functions Label":["Background"],"Citation Start End":[[618,634]],"Functions Start End":[[512,616]]}
{"Identifier":"2022ApJ...940....5D__B\u00e1lazs_et_al._2003_Instance_1","Paragraph":"Figure 1 shows the best fits of a multiple-Gaussian model to the T\n90 distributions of distinct Swift\/BAT GRB samples. Interestingly, we find from Figures 1(a)\u2013(d) that the lognormal T\n90 durations of all Swift GRBs except those with good spectra in sample V are triply distributed. It is confirmed in Figure 1(e) that the Swift\/BAT GRBs with a well-measured peak energy are bimodally distributed at a boundary of \u223c1 s, which is in good agreement with Zhang et al. (2020). Note that the third component remains controversial. On the other hand, Hakkila et al. (2000) suggested that the third subgroup proved by statistics was only a deviation caused by complex instrumental effects, which could reduce the duration of some weak long pulses. In addition, another classification scheme uses a scatter plot of flux and duration fitted with two-dimensional Gaussian functions (B\u00e1lazs et al. 2003). Some authors pointed out that there are more than two clusters (Mukherjee et al. 1998; Horv\u00e1th 1998). Unfortunately, the physical origin of the extra components cannot be reasonably interpreted. The two-Gaussian fit to data in Figure 1(e) demonstrates that the duration distribution peaks at 0.21 \u00b1 0.38 s with a spread of 0.93 dex for short bursts and at 43.97 \u00b1 1.05 s with a spread of 1.12 dex for long GRBs. The best fit returns a good reduced Chi-square of \n\n\n\n\u03c7\u03bd2\n\n \u2248 0.82, indicating that two classes are evidently reconfirmed and separated at T\n90 \u2248 1.06 s instead of at T\n90 = 2 s shown by CGRO\/BATSE data. The dividing line of 1 s is consistent with some previous results of Swift\/BAT GRBs (Berger et al. 2013; Gruber et al. 2014; Zhang et al. 2020). Additionally, the number of components in each T\n90 distribution of Figure 1 has been accordingly tested by the traditional Bayesian information criterion (BIC) as we previously did (e.g., Zhang et al. 2016; Li et al. 2021b). Figure 1 demonstrates that the S\/N level and the EE component are two important contributors to the third class of GRBs. In other words, one of them can separately confuse the classification of GRBs in terms of T\n90 only.","Citation Text":["B\u00e1lazs et al. 2003"],"Functions Text":["In addition, another classification scheme uses a scatter plot of flux and duration fitted with two-dimensional Gaussian functions"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[873,891]],"Functions Start End":[[741,871]]}
{"Identifier":"2020MNRAS.495.4508E__Heinke_et_al._2014_Instance_3","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"],"Functions Text":["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"],"Functions Label":["Uses"],"Citation Start End":[[1994,2012]],"Functions Start End":[[1785,1969]]}
{"Identifier":"2022ApJ...924...56S__Schreiber_et_al._2015_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":["Schreiber et al. 2015"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[712,733]],"Functions Start End":[[401,593]]}
{"Identifier":"2018ApJ...853..148C__Shibuya_et_al._2014_Instance_2","Paragraph":"LAE galaxies are defined by a high equivalent width (EW > 20 \u212b) Ly\u03b1 line and are believed to be composed of extremely large regions of active star formation. Many efforts have been made to detect and characterize LAE galaxies (e.g., Conselice et al. 2003; Conselice 2004; Ravindranath et al. 2006; Shimasaku et al. 2006; Bournaud et al. 2007; Ouchi et al. 2008, 2017; Elmegreen et al. 2009a, 2009b; Tacconi et al. 2010; Gronwall et al. 2011; Kashikawa et al. 2011; Mandelker et al. 2014; Moody et al. 2014; Guo et al. 2015). In general, these galaxies appear as clusters of bright clumps, sometimes with a background of continuum emission. Evidence suggests that these clumps are larger and brighter than most star-forming regions in nearby low-redshift galaxies (Elmegreen et al. 2009a). Efforts have been made in quantifying mass, star formation rates, gas composition, and kinematics, as well as other LAE properties (e.g., Nilsson et al. 2009; Ono et al. 2010a, 2010b; Swinbank et al. 2010; Tacconi et al. 2010; Shibuya et al. 2014; Livermore et al. 2015; Nakajima et al. 2016; Hashimoto et al. 2017). These have revealed a wealth of information about the early universe, but they are ultimately limited by LAE surface brightnesses. Most studies rely upon stacks of galaxies and can draw only limited inferences about individual LAEs. Other studies show that LAE dust content, particularly clumpy dust, in the interstellar medium (ISM) can have an impact on most LAE observables (Kobayashi et al. 2007, 2010; Verhamme et al. 2008; Duval et al. 2014). Finkelstein et al. (2009) showed that clumpy dust models can provide a good fit to a set of z \u223c 4.5 LAEs, although they invoked a multiphase ISM that may be unlikely to form in nature (Laursen et al. 2013). Nevertheless, dust in LAE galaxy ISM could cause some of the irregularity in LAE surface-brightness profiles (Buck et al. 2017). With limited resolution, however, it is difficult to make this distinction. A further challenge to morphological studies is that the clump sizes are near the resolution limit of instrumental point spread functions (PSFs) and often cannot be distinguished from point sources (Guo et al. 2015). As a result, direct imaging studies cannot decisively determine whether the clumps are different in nature from star-forming regions in our local universe or if the larger apparent size is merely an artifact of insufficient resolution (Shibuya et al. 2014; Kobayashi et al. 2016; Tamburello et al. 2017; Fisher et al. 2017).","Citation Text":["Shibuya et al. 2014"],"Functions Text":["As a result, direct imaging studies cannot decisively determine whether the clumps are different in nature from star-forming regions in our local universe or if the larger apparent size is merely an artifact of insufficient resolution"],"Functions Label":["Background"],"Citation Start End":[[2420,2439]],"Functions Start End":[[2184,2418]]}
{"Identifier":"2020ApJ...895L...8R__Kapferer_et_al._2009_Instance_1","Paragraph":"As clusters of galaxies assemble, they dynamically transform the physical properties of in-falling cluster members. Galaxies falling into a cluster experience ram pressure from dense intracluster medium (ICM) gas that can potentially unbind their individual gas reservoirs (Gunn et al. 1972). This process referred to as ram pressure stripping (RPS) can eventually remove a galaxy\u2019s entire gas supply, making it an important quenching pathway for satellite galaxies (Vollmer et al. 2001; Tonnesen et al. 2007). Observationally, RPS results in disturbed galaxy morphologies and trailing tails of stripped gas (e.g., Kenney et al. 2004; van Gorkom 2004; Cramer et al. 2019). The most extreme examples have been dubbed \u201cjellyfish\u201d galaxies, due to the evocative morphologies of their star-forming tails (Ebeling et al. 2014; Boselli et al. 2016; Poggianti et al. 2016). Prior to complete gas removal, moderate values of ram pressure have also been shown to increase the star formation rate in galaxies both observed (Crowl & Kenney 2006; Merluzzi et al. 2013; Vulcani et al. 2018) and simulated (Kronberger et al. 2008; Kapferer et al. 2009; Tonnesen & Bryan 2009; Bekki 2014). In this picture, the increased pressure initially helps compress the gas and triggers increased star formation. Other proposed processes for accretion of gas to the centers of galaxies include gravitational instabilities and the inspiral of preferentially low angular momentum clumps that lose angular momentum to a wind: drag from the nonrotating ICM operating on dense clumps (Schulz & Struck 2001; Tonnesen & Bryan 2009; Ramos-Mart\u00ednez et al. 2018). Over time, the interstellar medium (ISM) is fully stripped from the galaxy and star formation ceases. Recently, a very high incidence of AGN (5\/7) has been observed in a sample of jellyfish galaxies (Poggianti et al. 2017), and comprehensive follow-up of this sample has led to the identification of AGN-driven outflows (Radovich et al. 2019) and a compelling case for AGN feedback in action (George et al. 2019). It is plausible that the same mechanisms that initially promote star formation can also fuel active galactic nuclei (AGN) during the ram pressure stripping process. Indeed, outside of cluster environments, ram pressure induced shocks are known to produce nuclear inflows that can fuel AGN in the context of galaxy merger simulations (Barnes 2002; Capelo & Dotti 2017; Blumenthal & Barnes 2018).","Citation Text":["Kapferer et al. 2009"],"Functions Text":["Prior to complete gas removal, moderate values of ram pressure have also been shown to increase the star formation rate in galaxies both observed","and simulated"],"Functions Label":["Background","Background"],"Citation Start End":[[1117,1137]],"Functions Start End":[[867,1012],[1078,1091]]}
{"Identifier":"2016MNRAS.461.3982B__Walsh_&_Richardson_2008_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":["Walsh & Richardson 2008"],"Functions Text":["However, close encounters with the planets proved not to be enough for creation of the current population of binary systems"],"Functions Label":["Background"],"Citation Start End":[[609,632]],"Functions Start End":[[464,587]]}
{"Identifier":"2020MNRAS.496.5528M__Greif_et_al._2012_Instance_1","Paragraph":"\nTheoretically predicted magnetic field in the formation of the first stars. The evolution of a dynamo in a collapsing minihalo depends on a large number of parameters: the initial density, nH,0, the turbulent velocity, vt (which we parametrize in terms of the virial velocity, vt = \u03d5tvvir), the temperature, T, the mass of the collapsing cloud, M0, the rate of collapse (parametrized by \u03d5ff), and the rate at which these quantities vary with density (denoted by qx for quantity x). (The initial value of the field, B0, enters only logarithmically, and is important only if it is many orders of magnitude less than our estimate of \u223c10\u221216 G.) Choosing values of these parameters that are consistent with simulations (e.g. those of Greif et al. 2012), we find that the time for the field to grow from its initial amplitude of \u223c10\u221216 G to equipartition at the viscous scale, B\u03bd \u223c 10\u22128 G, is less than the virial time in the minihalo; hence, the exponential growth of the field occurs at approximately constant gas density. This rapid growth of the field is consistent with that found in previous work (e.g. Schleicher et al. 2010; Schober et al. 2012b). The subsequent non-linear dynamo amplification is sufficient to bring the field energy to within about an order of magnitude of equipartition; none the less, the overall amplification of the field is generally dominated by compression. We estimate that the field first reaches equipartition with turbulent velocities of the order of 2 km s\u22121 (taken from simulations) at a value of \u223c10\u22124 G; the field subsequently grows as $n_{\\rm H}^{1\/2}$. The field reaches equipartition with the central 5 per\u2009cent of the mass of the gas. Our conclusion that the field reaches equipartition in a minihalo at z \u223c 25 differs from that of Xu & Lazarian (2016), who found that equipartition was not reached until a time of about 6 \u00d7 108 yr (the age of the Universe at z \u2243 8) since they did not consider the increase in density that occurs in star formation.","Citation Text":["Greif et al. 2012"],"Functions Text":["Choosing values of these parameters that are consistent with simulations (e.g. those of"],"Functions Label":["Uses"],"Citation Start End":[[730,747]],"Functions Start End":[[642,729]]}
{"Identifier":"2021ApJ...913L..14H__Priest_&_Schrijver_2000_Instance_1","Paragraph":"Thanks to the observations with high time resolution and high optical sensitivity from state-of-the-art facilities, various dynamic phenomena, and processes (e.g., magnetic reconnections, jet flows, and oscillatory waves) have been observed to be omnipresent in the multi-layers of solar atmosphere (Shibata et al. 2007; Tomczyk et al. 2007; He et al. 2010; Tian et al. 2014; Shen et al. 2018). In the past, magnetic reconnection and wave dissipation are viewed as two seemingly distinct mechanisms opposite to one another when dealing with the problems of coronal heating and solar wind origin (e.g., Cranmer & Van Ballegooijen 2010). Magnetic reconnection rapidly converts magnetic energy to particle energy, causing emission flare in multiple wave bands and even triggering coronal mass ejections (Priest & Schrijver 2000). Nano-jets are observed in closed loop system and have been suggested as being caused by the slingshot effect of newly reconnected magnetic field lines with a small angle of shear before reconnection (Antolin et al. 2021). Oscillatory wave signatures (e.g., oscillations in radiation intensity and Doppler velocities) have been considered as propagating or standing in both open strands or closed loops throughout the solar atmosphere (Wang et al. 2009). Flares are found to be often accompanied by quasi-periodic pulsations in multiband radiation intensities (Nakariakov et al. 2010; Van Doorsselaere et al. 2016; McLaughlin et al. 2018). The causal relation between reconnection and Alfv\u00e9n waves was explored in 3D magnetohydrodynamic (MHD) simulations using the Wal\u00e9n test of the field-aligned current and parallel vorticity as emitting from the reconnection site (Ma et al. 1995). In the chromosphere, excitation of kink or Alfv\u00e9nic waves, as well as compressive slow-mode waves, are also discovered to be associated with the formation of type-II spicules, which are launched by magnetic reconnection (He et al. 2009a, 2009b; Liu et al. 2014).","Citation Text":["Priest & Schrijver 2000"],"Functions Text":["Magnetic reconnection rapidly converts magnetic energy to particle energy, causing emission flare in multiple wave bands and even triggering coronal mass ejections"],"Functions Label":["Background"],"Citation Start End":[[801,824]],"Functions Start End":[[636,799]]}
{"Identifier":"2018AandA...616A..99K__Narang_et_al._2016_Instance_3","Paragraph":"The high-resolution imaging observations of TR from IRIS reveal the ubiquitous presence of network jets. We have used three different IRIS observations of the quiet sun, which are located near the disk center. On the basis of careful inspection, 51 network jets are identified from three QS observations and used for further analysis. These 51 network jets are very well resolved and are not affected by the dynamics of other jets. The study is focused on the rotating motion of network jets along with the estimation of their other properties (speed, height, and lifetime). The mean speed, as predicted by statistical distributions of the speed, is 140.16 km s\u22121 with a standard deviation of 39.41 km s\u22121. The mean speed of network jets is very similar, as reported in previous works (e.g., Tian et al. 2014; Narang et al. 2016). However, in case of their lifetimes, we found a value that is almost double (105.49 s) that of the previously reported mean lifetime of the network jets (49.6 s; Tian et al. 2014). As mentioned above, we took only those network jets that are very well resolved in space and in the time; these criteria exclude short lifetime network jets. Therefore, our statistical distribution of the lifetime predicts a higher mean lifetime. The mean length of the network jets is 3.16 Mm with a standard deviation of 1.18 Mm. In the case of CH network jets, Tian et al. (2014) have reported that most of the network jets have lengths from 4.0 to 10.0 Mm. However, the mean length for QS network jets is smaller (3.53 Mm; Narang et al. 2016). So, the mean length for QS network jets from the present work is in good agreement with Narang et al. (2016). In addition, the apparent speed and length of these network jets are positively correlated, which is very similar to what has already been reported in previous works (Narang et al. 2016). Finally, we can say that these networks jets are very dynamic features of the solar TR, as revealed by their estimated properties.","Citation Text":["Narang et al. (2016)"],"Functions Text":["So, the mean length for QS network jets from the present work is in good agreement with"],"Functions Label":["Similarities"],"Citation Start End":[[1648,1668]],"Functions Start End":[[1560,1647]]}
{"Identifier":"2022AandA...666A..51M__Djura\u0161evi\u0107_1993_Instance_1","Paragraph":"RX Cas is part of a group of interacting binaries showing light curves similar to \u03b2 Lyrae and showing a long photometric cycle of unknown origin (Mennickent 2017). Few of these interacting binaries have been studied in terms of their physical changes during the long cycle. Here we provide a brief comparison with some of them. In RX Cas, we find that the long cycle of around 516 days can be explained in terms of changes of the physical parameters of the accretion disk, which was mentioned along with other possible causes in the past (K\u0159\u00ed\u017e et al. 1980). The disk turns out to be hotter and thicker at long-cycle maximum, but it does not change appreciably in radial extension. We note that a hotter disk during maximum was predicted by a simpler study in the past (Djura\u0161evi\u0107 1993). Contrarily, large changes in disk radius and thickness have been inferred in OGLE-BLG-ECL-157529, a binary characterized by orbital and long periods of 24\n\n\n\n.\n\nd\n\n\n\n$ \\overset{\\text{d}}{.} $\n\n\n8 and \u223c850 days, while its long cycle can be reproduced by occultation of the gainer by a variable disk thickness (Mennickent & Djura\u0161evi\u0107 2021). On the other hand, the system OGLE-LMC-DPV-097 is characterized by a long-cycle length of \u223c306 days and an amplitude of \u223c0.8 mag in the I band. This binary has an orbital period of 7\n\n\n\n.\n\nd\n\n\n\n$ \\overset{\\text{d}}{.} $\n\n\n75 and relatively low stellar masses of 5.5 and 1.1 M\u2299. In this system the orbital light curve changes in a systematic way during the long cycle. At the minimum of the long cycle the secondary eclipse practically disappears and during the ascending branch the system is brighter in the first quadrant than during the second quadrant. The disk radius is 7.5 R\u2299 at the minimum and 15.3 R\u2299 at the ascending branch of the long cycle; its temperature at the outer edge changes from 6870 to 4030 K in these two stages (Garc\u00e9s et al. 2018). OGLE-LMC-DPV-065 is another interesting binary with a large amplitude of the long cycle. The long cycle is characterized by a double hump light curve in the I and V bands whose general shape is nearly constant with only minor variations. After a continuous decrease in the long-period from 350 to 218 d lasting about 13 yr, the long cycle remained almost constant for about 10 yr. However, the orbital light curve is fairly constant in shape, and no clear link between the disk structure and long cycle has been observed for this system, in a clear contrast to OGLE-LMC-DPV-097 and RX Cas (Mennickent et al. 2019). From the above we conclude that we cannot draw general conclusions about disk changes and photometric long cycles in these few systems as they show different behaviors that cannot be easily adapted to a global single picture.","Citation Text":["Djura\u0161evi\u0107 1993"],"Functions Text":["We note that a hotter disk during maximum was predicted by a simpler study in the past"],"Functions Label":["Similarities"],"Citation Start End":[[769,784]],"Functions Start End":[[681,767]]}
{"Identifier":"2018MNRAS.476.2591V__Nikolic,_Cullen_&_Alexander_2004_Instance_1","Paragraph":"Galaxy interactions represent a fundamental component of our current view of hierarchical galaxy evolution. Studies based on both observations and simulations have shown that galaxy collisions and mergers can dramatically affect the galaxies undergoing the interaction, by, e.g. triggering nuclear activity (e.g. Kennicutt 1984; Kennicutt et al. 1987; Ellison et al. 2011, 2013a; Silvermann et al. 2011; Satyapal et al. 2014), producing colour changes (e.g. Larson & Tinsley 1978; Darg et al. 2010; Patton et al. 2011), disrupting morphologies (e.g. Kaviraj et al. 2011; Patton et al. 2016; Lofthouse et al. 2017), and altering the metallicities (e.g. Rupke et al. 2010; Perez, Michel-Dansac & Tissera 2011; Scudder et al. 2012; Torrey et al. 2012). The most evident effect driven by galaxy encounters is probably the triggering of new episodes of star formation, which can occur both in the pre-merger regime between first pericentre and coalescence (e.g. Nikolic, Cullen & Alexander 2004; Ellison et al. 2008, 2013b; Patton et al. 2011; Scudder et al. 2012), and in the post-merger phase, where the two nuclei of the interacting galaxies have merged together (e.g. Kaviraj et al. 2012; Kaviraj 2014; Ellison et al. 2013a). The idea that galaxy mergers have a strong impact on the star formation activity is supported by studies of Ultra-Luminous InfraRed Galaxies (ULIRGs), i.e. galaxies with IR luminosities exceeding 1012 L\u2299 and characterized by star formation rates (SFRs) up to \u223c1000 M\u2299 yr\u22121 (e.g. Barnes & Hernquist 1991; Mihos & Hernquist 1994; Daddi et al. 2010; Scoville et al. 2015). Observations have revealed that the majority of ULIRGs reside in interacting systems (e.g. Sanders & Mirabel 1996; Veilleux, Kim & Sanders 2002; Kartaltepe et al. 2010, 2012; Haan et al. 2011). Nevertheless, ULIRGs are rare and extreme examples of highly star-forming galaxies. Most galaxy\u2013galaxy interactions result in SFR increases of at most a factor of a few, as shown in both numerical simulations (e.g. Di Matteo et al. 2008) and observations of galaxy pairs and post-mergers (Ellison et al. 2008; Martig & Bournaud 2008; Jogee et al. 2009; Robaina et al. 2009; Scudder et al. 2012).","Citation Text":["Nikolic, Cullen & Alexander 2004"],"Functions Text":["The most evident effect driven by galaxy encounters is probably the triggering of new episodes of star formation, which can occur both in the pre-merger regime between first pericentre and coalescence (e.g."],"Functions Label":["Background"],"Citation Start End":[[957,989]],"Functions Start End":[[750,956]]}
{"Identifier":"2022MNRAS.516.3532M__Johnson_et_al._2017_Instance_1","Paragraph":"The size\u2013SFR relation has been already well-defined in the local Universe investigating H\u2009ii regions in nearby spiral and irregular galaxies by Kennicutt (1988). On the other hand, outliers are mainly hosted by interacting systems as shown in the case of the Antennae galaxy (Bastian et al. 2006), and have been thoroughly investigated on a larger samples with DYNAMO (Green et al. 2014) and LARS (Messa et al. 2019). In Fig. 10, we show the relation between the size and star formation rate of clumps, including samples across different redshifts (Swinbank et al. 2009, 2012; Jones et al. 2010; Livermore et al. 2012, 2015; Johnson et al. 2017). We consider samples where SFRs was estimated through SED fitting (as for our sample) or from H\u2009\u03b1, using the prescription by Kennicutt (1988). Our sample and the Sunburst clumps (Vanzella et al. 2021a), despite being at z \u223c 2\u20136, have sizes comparable to local H\u2009ii regions and GCs (Swinbank et al. 2009; Jones et al. 2010; Livermore et al. 2012; Johnson et al. 2017), but they have SFRs \u223c300 times higher. They also have SFRs \u223c100 times higher than clumps from Wuyts et al. (2014) and Livermore et al. (2015), which are part of a sample of lensed systems at z > 2. In particular, we compare our sample with two well-known H\u2009ii regions from the local Universe: 30 Doradus and II Zw40 (Vanzi et al. 2008). One possible explanation for the high SFR measured in our sample is enhanced interactions, or larger gas reservoirs at high redshift. Similar cases have been observed in the local Universe, for example, the Antennae galaxy which is a local merging system hosting six star-forming complexes whose SFR ranges from 0.2 to 1.4 M\u2299 yr\u22121, significantly higher than other local star-forming regions. Three of those complexes show signatures of Wolf\u2013Rayet stars, implying young ages of $\\sim 5\\, \\rm Myr$ (Bastian et al. 2006). Similar properties are observed in the Sunburst 5.1 knot presented in detail in Vanzella et al. (2021a) and in some of our clumps shown in fig. D.1 of Vanzella et al. (2021b). Other scenarios, such as the fragmentation of gas-rich discs, are potentially responsible for the observed high SFR among high redshift clumpy structures (Noguchi 1999; Dekel et al. 2009). During this process, cold gas cools becoming unstable and the galactic disc starts to fragments and forms clumpy structures. Further on, such newly formed structures lead to the increased SFR as discussed in i.e. Immeli et al. (2004).","Citation Text":["Johnson et al. 2017","Johnson et al. 2017"],"Functions Text":["In Fig. 10, we show the relation between the size and star formation rate of clumps, including samples across different redshifts","Our sample","despite being at z \u223c 2\u20136, have sizes comparable to local H\u2009ii regions and GCs","but they have SFRs \u223c300 times higher."],"Functions Label":["Uses","Similarities","Similarities","Differences"],"Citation Start End":[[625,644],[992,1011]],"Functions Start End":[[418,547],[789,799],[849,926],[1014,1051]]}
{"Identifier":"2019AandA...627A.172R__Rozitis_&_Green_(2013)_Instance_3","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)"],"Functions Text":["However, a study of non-convex shape models for fast two to four hour rotators in","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."],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[1004,1026]],"Functions Start End":[[922,1003],[1027,1193]]}
{"Identifier":"2017AandA...604A..53C__Bolatto_et_al._(2013)_Instance_1","Paragraph":"The reasons for such a tight \u2013 and linear \u2013 correlation between \\hbox{$L^{\\prime}_{\\rm CO(1{-}0)}$}LCO(1\u22120)\u2032 and M\u2217 (or LK) observed consistently across different galaxy samples except in early type galaxies have been little explored in the literature. The interpretation favoured by Leroy et al. (2005) is that CO emission and stars are linked through the hydrostatic pressure in the disk, which depends mainly on the stellar surface density (\u03a3\u2217) and sets the rate at which Hi is converted into H2. An alternative explanation that we propose here goes back to the nature of the optically thick 12CO emission and to the approximately linear relation between 12CO luminosity and virial mass found for GMCs (see e.g. Scoville et al. (1987), Solomon et al. (1987), and Bolatto et al. (2013) for a more recent compilation). By extrapolating from the relation shown by Scoville et al. (1987), our observed \\hbox{$L^{\\prime}_{\\rm CO(1{-}0)}-M_*$}LCO(1\u22120)\u2032\u2212M\u2217 correlation may be so tight and close to linear simply because the global 12CO luminosity is a very good tracer of the dynamical mass in star-forming galaxies, assuming that in this class of objects the bulk of the CO emission traces molecular gas clouds in virial motions (for example in a rotating disk). This explanation of course assumes that in most star-forming galaxies the stellar mass is also a good tracer of the dynamical mass. Following on from our hypothesis, a possible explanation for the break down of the \\hbox{$L^{\\prime}_{\\rm CO(1{-}0)}-M_*$}LCO(1\u22120)\u2032\u2212M\u2217 relation in early types is that in these sources, even when CO is detected, the motions of the CO-emitting clouds are poor tracers of the total dynamical mass of the system. Interferometric CO observations of large samples of early type galaxies have indeed shown that the CO emission in these objects tends to be rather compact (on average extending over ~ 1 kpc) compared to the optical extent of the galaxy (Alatalo et al. 2013; Davis et al. 2013). ","Citation Text":["Bolatto et al. (2013)"],"Functions Text":["An alternative explanation that we propose here goes back to the nature of the optically thick 12CO emission and to the approximately linear relation between 12CO luminosity and virial mass found for GMCs (see e.g.","for a more recent compilation"],"Functions Label":["Uses","Uses"],"Citation Start End":[[766,787]],"Functions Start End":[[500,714],[788,817]]}
{"Identifier":"2020AandA...638A..44B__Jim\u00e9nez-Serra_et_al._2010_Instance_1","Paragraph":"Combining these different velocity signatures, the data show strong signatures of two gas components at different velocities (around 6 km s\u22121 apart) that converge to a common intermediate velocity at the location of the infrared dark cloud and active star-formingregion, similar to filament formation via gravitationally driven, converging gas flows (e.g., G\u00f3mez & V\u00e1zquez-Semadeni 2014). We interpret these signatures as indicators of converging gas streams that may trigger the star formation event at its center (e.g., V\u00e1zquez-Semadeni et al. 2006; Heitsch et al. 2008; Banerjee et al. 2009; G\u00f3mez & V\u00e1zquez-Semadeni 2014). Position-velocity diagrams based on simulations of cloud-cloud collisions sometimes show a characteristic pattern of lower-level emission between two main velocity components, a so-called bridging feature (e.g., Haworth et al. 2015, 2018). Similar signatures were also reported in observations (e.g., Jim\u00e9nez-Serra et al. 2010; Henshaw et al. 2013; Dobashi et al. 2019; Fujita et al. 2019). The pv-diagrams of the G28.3 region presented here (Fig. 9) show different signatures in the sense that there is not a lower-intensity bridge between the two well-defined components, but that the two velocity components converge at the center of the cloud toward a central, high-intensity velocity component. However, the absences of a bridging feature does not necessarily rule out the formation of G28.3 in a cloud-cloud collision, as this feature is not always visible in simulations. For example, simulations by Bisbas et al. (2017) show that the low-intensity bridge feature may merge into a centrally peaked pv-diagram during the evolution of the cloud-cloud collisions, while the colliding-cloud simulations of Clark et al. (2019) yield only a single central velocity component in CO or [CI], with multiple components only becoming apparent when the [CII] emission from the cloud is considered. In addition, the multiple components and bridging feature may not be visible in cases where our line of sight is oriented at a large angle to the direction of motion of the clouds. We return to the interpretation in Sects. 4.4 and 4.5.","Citation Text":["Jim\u00e9nez-Serra et al. 2010"],"Functions Text":["Similar signatures were also reported in observations (e.g.,"],"Functions Label":["Similarities"],"Citation Start End":[[928,953]],"Functions Start End":[[867,927]]}
{"Identifier":"2021MNRAS.508..637S__Johnston_et_al._2019_Instance_1","Paragraph":"We also present the correlation functions of centrals and satellites in Fig. 17. This is a worthwhile exercise for a variety of reasons, not least that there is information on the small-scale IA signal missed in the large-scale fits. While we cannot, at the present time, fit the IA signal on scales \u223c1\u2009h\u22121\u2009Mpc, the qualitative comparison can be instructive. Given that there is some evidence that they behave differently, we consider blue and red satellites\/centrals separately here. We also drop the full sample density tracer, and instead use one of the four (red\/blue, satellite\/central) subsamples. The motivation here is that, while on large scales, the density tracer is effectively just that: a probe of the large-scale matter distribution multiplied by a linear galaxy bias, on scales approaching the one halo regime this no longer holds. Fig. 17 shows these new data vectors. As shown we measure both wg + and w++, and recompute the covariance matrices with the appropriate densities. For reference, the dark red crosses also show the equivalent satellite\/central red galaxy wg + correlations from KiDS\u00d7GAMA here (c.f. Johnston et al. 2019 fig. 7, red points\/band). On large scales at least, our illustristng red sample is consistent with their measurements. There are a few interesting features here to note, however. First, we see a relatively strong red galaxy 1h contribution on scales 1 h\u22121 Mpc. Although the general trends match the real data, with ss and to a lesser extent cs exhibiting strong scale-dependent IAs in this regime, the magnitude is somewhat higher in our sample. This is particularly interesting, given that our sample characteristics are similar (\u2329L\u232a\/L0 = 0.91 and 0.34 for our red and blue samples, respectively, compared with their \u223c0.99 and 0.50). As discussed briefly in Section 6.2.1, we observe a persistent non-zero IA signal in blue galaxies on large scales; here we can see it is dominated by the cc correlation, with a smaller contribution from sc. Also notable is that, in line with some previous direct IA measurements (e.g. Singh et al. 2015), the large-scale satellite correlations do not appear to vanish on large scales. Focusing on the right-hand panels, the purple and pink points are consistently positive and non-zero. While small compared with the red central terms, and consistent with the dark red points from GAMA, there appears to be a detectable signal at the precision allowed by illustristng.","Citation Text":["Johnston et al. 2019"],"Functions Text":["For reference, the dark red crosses also show the equivalent satellite\/central red galaxy wg + correlations from KiDS\u00d7GAMA here (c.f.","fig. 7, red points\/band).","On large scales at least, our illustristng red sample is consistent with their measurements.","There are a few interesting features here to note, however.","Although the general trends match the real data,","with ss and to a lesser extent cs exhibiting strong scale-dependent IAs in this regime, the magnitude is somewhat higher in our sample.","This is particularly interesting, given that our sample characteristics are similar (\u2329L\u232a\/L0 = 0.91 and 0.34 for our red and blue samples, respectively, compared with their \u223c0.99 and 0.50)."],"Functions Label":["Uses","Uses","Similarities","Compare\/Contrast","Similarities","Differences","Compare\/Contrast"],"Citation Start End":[[1129,1149]],"Functions Start End":[[995,1128],[1150,1175],[1176,1268],[1269,1328],[1411,1459],[1460,1595],[1596,1784]]}
{"Identifier":"2021MNRAS.508..637S__Chang_et_al._2019_Instance_1","Paragraph":"It is now well established that the weak lensing of distant galaxies by foreground mass provides a relatively clear window on to the large-scale structure of the Universe. This is true whether that foreground mass is in the form of discrete matter concentrations, as traced by galaxies (i.e. galaxy\u2013galaxy lensing; Mandelbaum et al. 2013; Leauthaud et al. 2017; Joudaki et al. 2018; Prat et al. 2018; Blake et al. 2020), massive dark matter haloes (cluster lensing; Melchior et al. 2017; Dark Energy Survey Collaboration 2020), or the continuous large-scale matter distribution (cosmic shear; Heymans et al. 2013; Dark Energy Survey Collaboration 2016; Troxel et al. 2018; Chang et al. 2019; Hamana et al. 2020; Hildebrandt et al. 2020; Amon et al. 2021; Asgari et al. 2021; Secco, Samuroff et al. 2021). Though the measurement method and the exact form of the theory predictions differ slightly in the three cases, they are all fundamentally probes of the growth of structure at low redshift. Similarly cross-correlations between galaxy lensing and other observables can be powerful probes in their own right; recent examples include galaxy lensing \u00d7 CMB lensing (Schaan et al. 2017), voids correlated with CMB lensing (Vielzeuf et al. 2021), and galaxy weak lensing crossed with gamma-ray emission (Ammazzalorso et al. 2020), each of which provide probes of dark matter with slightly different sensitivities. A measurement of cosmological weak lensing, however, is subject to a range of systematic effects; that is, observational effects that mimic a cosmological lensing signal, and so bias cosmological inference if one neglects them. Depending on the systematic in question, the most effective mitigation strategy may be quite different. In broad terms, however, the standard approach is to either (i) mitigate systematics where possible, either by applying a calibration to the data, or discarding the data points most strongly affected or (ii) marginalize over them with a parametric model. Often a combination of the two is appropriate, and the prior used in (ii) is informed by additional data or simulations, and detailed testing of the calibration step in (i).","Citation Text":["Chang et al. 2019"],"Functions Text":["It is now well established that the weak lensing of distant galaxies by foreground mass provides a relatively clear window on to the large-scale structure of the Universe. This is true whether that foreground mass is in the form","or the continuous large-scale matter distribution (cosmic shear;"],"Functions Label":["Background","Background"],"Citation Start End":[[673,690]],"Functions Start End":[[0,228],[528,592]]}
{"Identifier":"2019MNRAS.488.5029H__Stacey_et_al._2010_Instance_2","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"],"Functions Text":["The [C\u2009ii] deficit persists when including high-z galaxies (e.g."],"Functions Label":["Motivation"],"Citation Start End":[[1105,1123]],"Functions Start End":[[1040,1104]]}
{"Identifier":"2022MNRAS.510.5302I__Fields_et_al._2014_Instance_1","Paragraph":"Lithium is the only metal element produced during the Big Bang nucleosynthesis (BBN) due to the lack of stable nuclei with mass number eight (Fields, Molaro & Sarkar 2014). The element abundances predicted by the standard BBN theory for the baryonic density coming from the Planck mission agree well with those observed, except for 7Li (Fields 2011; Coc, Uzan & Vangioni 2014). Indeed, the abundance of lithium measured in the low-metallicity Galactic halo stars is A(7Li)  = log[N(7Li)\/N(H)] + 12 = 2.25 (Spite & Spite 1982; Sbordone et al. 2010; Bonifacio et al. 2015), which is \u223c3 times below the estimate of the standard cosmological model A(7Li) = 2.72 \u00b1 0.06 (Cyburt et al. 2016). The latter value depends on the baryon-to-photons ratio $\\eta = \\frac{N_b}{N_{\\gamma }} \\propto \\Omega _B h^2$, with \u03a9b is the cosmological baryon density and h is the dimensionless hubble parameter (Planck Collaboration XIII 2016). This problem is also known as the Cosmological Lithium problem (Fields et al. 2014). A possible solution can be ascribed to convective diffusion in the pre-main-sequence phase as well as during the lifetime of these halo stars (Fu et al. 2015) or to new physics beyond the standard model. On the other hand, the young stellar populations in our Galaxy show Li-abundances four times greater than the SBBN estimate and more than one order of magnitude greater than the halo stars (Spite 1990; Lambert & Reddy 2004; Lodders, Palme & Gail 2009; Ram\u00edrez et al. 2012; Fu et al. 2018). The evidence of a growth requires the existence of additional lithium factories. In the last decades several astrophysical Li sources have been proposed, like Galactic cosmic-rays, AGB stars, low-mass Carbon stars, type II supernovae, and classical novae (D\u2019Antona & Matteucci 1991; Romano et al. 1999; Prantzos 2012; Matteucci 2021). The recent detection in the outburst spectra of classical novae of 7Li and 7Be\u2009ii, an isotope whose unique decay channel is into lithium through electron capture, have confirmed these objects as Li producers. The corresponding yields inferred have placed nova explosions as the main lithium factories in the Galaxy. The time-scales involved also match, as shown by detailed Galactic chemical evolution (Izzo et al. 2015; Tajitsu et al. 2015; Molaro et al. 2016; Izzo et al. 2018; Cescutti & Molaro 2019; Grisoni et al. 2019; Molaro et al. 2020a; Matteucci 2021).","Citation Text":["Fields et al. 2014"],"Functions Text":["This problem is also known as the Cosmological Lithium problem"],"Functions Label":["Background"],"Citation Start End":[[984,1002]],"Functions Start End":[[920,982]]}
{"Identifier":"2016ApJ...824...92H__Douglas_&_Ballai_2007_Instance_1","Paragraph":"A couple of years after the closure of the solar flare myth debate, Dere et al. (1997) and Thompson et al. (1998) reported on the presence of wavefronts observed on the solar disk that propagated away from the site of the flare. They were originally termed \u201cEIT waves\u201d (Thompson et al. 1999) after the instrument that first detected them (EIT on board the Solar and Heliospheric Observatory (SOHO) Delaboudini\u00e9re et al. 1995) and have now been intensely studied for over 15 years. Over that time they have been assigned different terms, such as flare waves (e.g., Thompson et al. 2000; Warmuth et al. 2003), extreme-UV (EUV) waves (e.g., Patsourakos & Vourlidas 2012; Muhr et al. 2014), coronal waves (e.g., Douglas & Ballai 2007; Webb & Howard 2012), and coronal bright fronts (e.g., Gallagher & Long 2011, and references therein). While there are differences in the precise definitions of these terms, broadly speaking they describe the same phenomenon: a bright front propagating away from a point-like origin on the solar disk. In many cases the origin appears to be from the site of a solar flare associated with a CME (e.g., Thompson 1999; Thompson & Myers 2009; Wills-Davey & Attrill 2009; Nitta et al. 2013), but some workers have reported a stronger correlation of wave occurrence with CMEs rather than flares (e.g., Wills-Davey & Attrill 2009, and references therein). For simplicity, henceforth in this paper we refer to them as EIT waves although we do not use the EIT instrument for any measurements in our study. We define EIT waves with the criteria described in Section 4. Nitta et al. (2013) provide a recent review of the various physical narratives that have been assigned to EIT waves, such as the coronal manifestations of Moreton waves (e.g., Thompson 1999; Warmuth et al. 2001, 2002), the base of expanding overhead CME magnetic structures (e.g., Delan\u00e9e & Aulanier 1999), and solitons (Wills-Davey et al. 2007). Other reviews on EIT waves can be found by Wills-Davey & Attrill (2009), Gallagher & Long (2011), and Patsourakos & Vourlidas (2012).","Citation Text":["Douglas & Ballai 2007"],"Functions Text":["Over that time they have been assigned different terms, such as","coronal waves (e.g.,"],"Functions Label":["Background","Background"],"Citation Start End":[[708,729]],"Functions Start End":[[481,544],[687,707]]}
{"Identifier":"2015ApJ...815..129S__Shen_et_al._2011_Instance_1","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"],"Functions Text":["For comparison, we show published observations in the same redshift range from the literature in the right panel of Figure 5"],"Functions Label":["Uses"],"Citation Start End":[[1003,1019]],"Functions Start End":[[833,957]]}
{"Identifier":"2018AandA...615A.148D__Weidner_et_al._(2010)_Instance_3","Paragraph":"The last column in Table 1 reports the number of OB stars minus the \u201cdiffuse\u201d population estimated from their density in the Reference field (22.5 stars per square degree): as is immediately seen, the M-star statistics is much larger than the OB star statistics. This can hardly be considered surprising, if an ordinary IMF (e.g., that from Weidner et al. 2010) is assumed for the star-formation region. In Fig. 19 we show the density ratio between M and OB stars, which provides a consistency test between our results and a plausible IMF: this ratio varies however by a large factor, close to 20, among our subregions. This might reflect differences in the respective IMFs, but also differences in completeness among the stellar samples considered for the various regions. We first note that the ratio between M and OB stars in NGC 6231 is dramatically lower than anywhere else in Sco OB1. We can indeed expect that M stars are detected less efficiently in the inner parts of NGC 6231, where the density of bright stars is very large, and their diffuse glare raises the limiting magnitude locally. As already discussed above in Sect. 4.1, this causes our sample of M stars in NGC 6231 to be highly incomplete. Moreover, we determined above that NGC 6231 is significantly more extincted, by almost half a magnitude in V, than Tr 24, and this implies a higher minimum detectable mass among NGC 6231 M stars compared to Tr 24 (see the MDA diagrams in Fig. 5); this effect reduces the completeness of the M-star sample in NGC 6231 more than in Tr 24. If Tr 24 is also slightly younger than NGC 6231, as we argue below, our M-stars in Tr 24 will reach down to lower masses than in NGC 6231, with a steep increase in the detected M-star population: adopting the IMF from Weidner et al. (2010), the predicted number of cluster M stars doubles considering the mass interval 0.25\u20130.5 M\u2299 rather than 0.35\u20130.5 M\u2299. If Tr 24 is younger than NGC 6231, moreover, its stars in the mass range 2.5\u20133 M\u2299 might not have yet reached their ZAMS position as B stars, and therefore would not be counted among OB stars; this would further raisethe M\/OB star ratio there by up to 30%. Therefore, the proportions of both M and OB stars that are detected in a young cluster will depend on their age and extinction, in accordance with the MDA diagrams, even for a fixed, spatially uniform photometric sensitivity. We estimated using the Weidner et al. (2010) IMF the expected range for the observed M\/OB number ratio. Siess et al. (2000) predict that the latest-type B stars have a mass of ~ 3.5 M\u2299 at 2 Myr, and ~ 2.2 M\u2299 at 10 Myr, that is, in the range of ages expected for Sco OB1 clusters. The MDA diagrams of Fig. 5 predict that the lowest-mass stars we are able to detect using the available Sco-OB1 data have ~ 0.2 M\u2299, even assuming the most favorable (and unlikely) circumstances of an age less than 2 Myr and negligible reddening. The extreme values found for the M\/OB ratio are then ~ 3.8 for a minimum M-star mass as high as 0.35 M\u2299 and an old age of 10 Myr, and ~ 20 for a minimum M-star massas low as 0.2 M\u2299 and age of 2 Myr. These extremes are also shown as horizontal lines in Fig. 19. We note that the M\/OB ratio in NGC 6231 falls well within this range; however, both Tr 24 regions are significantly richer of M stars than expected, by more than a factor of two and well above (statistical) errors. If true, then paradoxically this part of the OB association would form preferentially lower-mass stars. Of course, more detailed studies are needed to confirm this result. In the G345.45+1.50 region the M\/OB ratio is highest, and far above predictions from the IMF: we may tentatively explain this since this region is very young, and some of its most massive members, like IRAS 16562-3959, are still in formation, thus decreasing the number of optically revealed OB stars. The lowest M\/OB ratio in NGC 6231 is unlikely to be real, since as discussed above our M-star sample in this densest subregion is likely incomplete.","Citation Text":["Weidner et al. (2010)"],"Functions Text":["We estimated using the","IMF the expected range for the observed M\/OB number ratio."],"Functions Label":["Uses","Uses"],"Citation Start End":[[2409,2430]],"Functions Start End":[[2386,2408],[2431,2489]]}
{"Identifier":"2021MNRAS.503..815V__Schmalzing,_Buchert_&_Kerscher_1995_Instance_1","Paragraph":"Level crossing statistics is a pioneering approach for characterizing stochastic processes introduced by S. O. Rice (Rice 1944, 1945). Up-, down-, and conditional crossing statistics are modifications to the primary definition of level crossing (Bardeen et al. 1986; Bond & Efstathiou 1987; Ryden 1988; Ryden et al. 1989; Matsubara 1996, 2003; Brill 2000; Shahbazi et al. 2003; Movahed & Khosravi 2011; Ghasemi Nezhadhaghighi et al. 2017). Minkowski functionals, which are also closely related to the crossing statistics, provide 1 + D functionals to quantify the morphology in D dimensions (Hadwiger 2013) and have been utilized for cosmological random fields (Mecke, Bucheri & Wagner 1994; Schmalzing, Buchert & Kerscher 1995; Schmalzing & G\u00f3rski 1998; Matsubara 2003, 2010; Hikage, Komatsu & Matsubara 2006; Codis et al. 2013; Ling et al. 2015; Fang, Li & Zhao 2017). A number of critical sets including peaks (hills), troughs (lakes), saddles, voids, skeleton, genus, and Euler characteristics, are more popular in cosmology for different purposes and they have been fully explored for Gaussian stochastic fields. Some extensions for non-Gaussian and anisotropic conditions have been done in some research (Matsubara 2003; Pogosyan et al. 2009; Pogosyan, Pichon & Gay 2011; Gay et al. 2012; Codis et al. 2013). More recently, Betti numbers, Euler characteristics, and Minkowski functionals for a set of cosmological 3D fields have been examined extensively (Pranav et al. 2019). The scaling approach for investigating cosmological stochastic fields has been discussed by Borgani (1995), Movahed et al. (2011). Standard estimators like three- and four-point functions in real space, bispectra and trispectra in harmonic space, multiscaling methods such as wavelets (Planck Collaboration 2014b, 2016d; and references therein), and regeneration of stochastic processes based on the Fokker\u2013Planck equation (Ghasemi et al. 2006) have also been considered.","Citation Text":["Schmalzing, Buchert & Kerscher 1995"],"Functions Text":["Minkowski functionals, which are also closely related to the crossing statistics","and have been utilized for cosmological random fields"],"Functions Label":["Background","Background"],"Citation Start End":[[692,727]],"Functions Start End":[[440,520],[607,660]]}
{"Identifier":"2017AandA...602A..75R__Kaneko_&_Yokoyama_(2015)_Instance_2","Paragraph":"We can see how this is consistent with the development of small scales via phase mixing by introducing local wavenumbers for the variation with \u03b1 and \u03b2, (3)\\begin{eqnarray} \\xi \\propto \\exp {\\rm i}\\left[ \\int \\kappa_\\alpha {\\rm d}\\alpha + \\int \\kappa_\\beta {\\rm d}\\beta \\right]\\!. \\label{number3} \\end{eqnarray}\u03be\u221dexpi\u222b\u03ba\u03b1d\u03b1+\u222b\u03ba\u03b2d\u03b2.Here \u03ba\u03b1 and \u03ba\u03b2 are the wavenumbers in \u03b1 and \u03b2 and have units that are the inverse of the units of their respective coordinates. These wavenumbers should be distinguished from the perpendicular components of the usual wave vector k, which has units of 1\/length. The different wavenumbers may be related through the scale factors (h) that relate elemental coordinate increments to physical distances: dr = e\u03b1h\u03b1d\u03b1 + e\u03b2h\u03b2d\u03b2 + e\u03b3h\u03b3d\u03b3, where e\u03b1 is a unit vector in the \u03b1 direction, etc. In this notation \u2207\u22a5 = (e\u03b1\/h\u03b1)\u2202\/\u2202\u03b1 + (e\u03b2\/h\u03b2)\u2202\/\u2202\u03b2, and noting that \u2207\u22a5\u03be \u2248 ik\u22a5\u03be, Eq. (3) yields \\begin{eqnarray} \\bdelp\\xi &amp;=&amp; {\\rm i} \\vec{k}_\\perp \\xi = {\\rm i} \\left( \\frac{\\vec{e}_\\alpha}{h_\\alpha}\\frac{ \\partial}{\\partial\\alpha} +\\frac{\\vec{e}_\\beta}{h_ \\beta}\\frac{ \\partial}{\\partial \\beta}\\right)\\xi \\\\ &amp;\\equiv&amp; {\\rm i}\\left( \\frac{\\vec{e}_\\alpha}{h_\\alpha}\\kappa_\\alpha +\\frac{\\vec{e}_\\beta}{h_ \\beta}\\kappa_\\beta\\right)\\xi. \\label{number4} \\end{eqnarray}\u2207\u22a5\u03be=ik\u22a5\u03be=ie\u03b1h\u03b1\u2202\u2202\u03b1+e\u03b2h\u03b2\u2202\u2202\u03b2\u03be\u2261Equating components of the second and fourth expressions in Eq. (5) gives the expected relations between the various wavenumbers, (6)\\begin{eqnarray} k_\\alpha = \\kappa_\\alpha \/h_\\alpha, \\qquad k_\\beta = \\kappa_ \\beta \/h_ \\beta. \\label{number5} \\end{eqnarray}k\u03b1=\u03ba\u03b1\/h\u03b1,\u2001k\u03b2=\u03ba\u03b2\/h\u03b2.Equations (2) and (5) give a direct and elegant expression for the perpendicular wave vector as (7)\\begin{eqnarray} \\vec{k}_\\perp \\approx -(\\bdel \\omega_{\\rm c}) t, \\label{number6} \\end{eqnarray}k\u22a5\u2248\u2212(\u2207\u03c9c)t,which is a generalisation to three dimensions of the results of Mann et al. (1995), (Wright et al. 1999) and Kaneko & Yokoyama (2015) for lower dimensional systems, which developed phase mixing in only one perpendicular coordinate. The above expression allows phase mixing in both perpendicular directions, giving physical phase mixing lengths (or wavelengths) in the \u03b1 and \u03b2 directions of (8)\\begin{eqnarray} L_{{\\rm ph}\\alpha} = \\frac{2\\pi}{|k_\\alpha |} \\equiv \\frac{2\\pi h_\\alpha}{|\\partial \\omega_{\\rm c}\/\\partial\\alpha |t}, \\qquad L_{{\\rm ph}\\beta} = \\frac{2\\pi}{|k_ \\beta |} \\equiv \\frac{2\\pi h_\\beta}{|\\partial \\omega_{\\rm c}\/\\partial \\beta |t}\\cdot \\label{number7} \\end{eqnarray}Lph\u03b1=2\u03c0|k\u03b1|\u22612\u03c0h\u03b1|\u2202\u03c9c\/\u2202\u03b1|t,\u2001Lph\u03b2=2\u03c0|k\u03b2|\u22612\u03c0h\u03b2|\u2202\u03c9c\/\u2202\u03b2|t\u00b7If the phase mixing lengths are expressed in the same units as \u03b1 and \u03b2, rather than physical length as in Eq. (8), slightly simpler expressions are found, i.e. (9)\\begin{eqnarray} \\ell_{{\\rm ph}\\alpha} = \\frac{2\\pi}{|\\kappa_\\alpha |} \\equiv \\frac{2\\pi}{|\\partial \\omega_{\\rm c}\/\\partial\\alpha |t}, \\qquad \\ell_{{\\rm ph}\\beta} = \\frac{2\\pi}{|\\kappa_ \\beta |} \\equiv \\frac{2\\pi}{|\\partial \\omega_{\\rm c}\/\\partial \\beta |t}\\cdot \\label{number8} \\end{eqnarray}\u2113ph\u03b1=2\u03c0|\u03ba\u03b1|\u22612\u03c0|\u2202\u03c9c\/\u2202\u03b1|t,\u2001\u2113ph\u03b2=2\u03c0|\u03ba\u03b2|\u22612\u03c0|\u2202\u03c9c\/\u2202\u03b2|t\u00b7The development of the phase mixing length can be pictured simply as the tendency for each field line to oscillate with its own natural frequency. Even if all the field lines start to oscillate with the same phase, they soon drift out of phase with one another as time passes. Not only does the phase mixing process generate perpendicular scales, but points of constant phase can be seen to move across field lines. This phase motion has been seen in magnetospheric data of Alfv\u00e9n waves (see the review by Wright & Mann 2006) and the simulations of coronal oscillations by Kaneko & Yokoyama (2015). These studies note that the direction of motion is related to the spatial variation of \u03c9c. The results of these papers for the perpendicular phase velocity in physical space generalise to Vph = \u03c9c\/k\u22a5, giving the components (10)\\begin{eqnarray} V_{{\\rm ph}\\alpha} = \\frac{-\\omega_{\\rm c} h_\\alpha}{(\\partial \\omega_{\\rm c}\/\\partial\\alpha )t}, \\qquad V_{{\\rm ph}\\beta} = \\frac{-\\omega_{\\rm c} h_ \\beta}{(\\partial \\omega_{\\rm c}\/\\partial \\beta )t}, \\qquad \\label{number9} \\end{eqnarray}Vph\u03b1=\u2212\u03c9ch\u03b1(\u2202\u03c9c\/\u2202\u03b1)t,\u2001Vph\u03b2=\u2212\u03c9ch\u03b2(\u2202\u03c9c\/\u2202\u03b2)t,\u2001If the excitation occurred at a time ti, the subsequent properties are found by replacing t with t\u2212ti in the above formulae. ","Citation Text":["Kaneko & Yokoyama (2015)"],"Functions Text":["This phase motion has been seen in magnetospheric data of Alfv\u00e9n waves","and the simulations of coronal oscillations by","These studies note that the direction of motion is related to the spatial variation of \u03c9c."],"Functions Label":["Uses","Uses","Uses"],"Citation Start End":[[3614,3638]],"Functions Start End":[[3457,3527],[3567,3613],[3640,3730]]}
{"Identifier":"2021AandA...646A.142R__Nelson_&_Melrose_1985_Instance_1","Paragraph":"CME-driven shocks accelerate not just protons, but also the electrons in the solar corona (Holman & Pesses 1983; Schlickeiser 1984; Kirk 1994; Mann et al. 1995, 2001; Mann & Klassen 2005). These accelerated electron beams can be observed as type II bursts in the solar radio radiation in the metric wave range (Wild & McCready 1950; Uchida 1960). Type II bursts require electrons escaping from the shock front, and the lack of these bursts implicates the absence of accelerated electrons as type II bursts occur when 0.2\u201310 KeV electrons are accelerated in the shock front (see, e.g., Bale et al. 1999; Knock et al. 2001; Mann & Klassen 2005). Energetic electrons are unstable to Langmuir waves, thus they are converted into radio emission at the local plasma frequency and its harmonic (see Nelson & Melrose 1985). Therefore, type II radio bursts hold crucial information of both the shock and the surrounding ambient medium in which the CME-driven shock propagates (Gopalswamy et al. 2008a). Although almost every large SEP event is accompanied by a type II radio burst (Gopalswamy 2003; Cliver et al. 2004) that indicates CME-driven particle acceleration (Gosling 1993; Reames 1999), we have ten events in our sample that lack a type II burst: 14 August 2010, 03 August 2011, 04 March 2012, 26 May 2012, 27 May 2012, 14 June 2012, 08 September 2012, 14 December 2012, 21 April 2013, and 06 November 2013. Of the ten events, two originate in the west, three in the east, and five at the disk center. On average, these events have a maximum velocity of about 1000 km s\u22121 and maximum Mach number of about 0.94. The protons accelerated by these events have peak fluxes of about 23.2 cm\u22122 s\u22121 sr\u22121 in the > 10 MeV band, making these events the slowest and weakest SEPs in the sample. As the propagation of radio bursts does not depend on the magnetic connectivity, a possible explanation for their absence could be that the path of the radio burst did not coincide with the instrument on board the satellite or that the detection of the waves was below the range of the radio instrument, hence missing the signature. Detailed investigation is required to understand the absence of DH type II radio bursts in these events.","Citation Text":["Nelson & Melrose 1985"],"Functions Text":["Energetic electrons are unstable to Langmuir waves, thus they are converted into radio emission at the local plasma frequency and its harmonic (see"],"Functions Label":["Background"],"Citation Start End":[[792,813]],"Functions Start End":[[644,791]]}
{"Identifier":"2019AandA...626A..49P__Springel_et_al._(2005)_Instance_1","Paragraph":"To provide enough galaxies to adequately train a neural network, EAGLE galaxies from the simulation snapshots with a redshift of less than 1.0 were used. Objects with M\u22c6 greater than 1010\u2006M\u2299 were selected while the merging partner of the merging systems must be larger than 109\u2006M\u2299. The merging partner must also be more than 10% of the M\u22c6 of the primary galaxy. Galaxies were deemed to have merged when they are tracked as two galaxies in one simulation snapshot and then tracked as one galaxy in the following snapshot in the EAGLE merger trees catalogue (Qu et al. 2017). This prevents the inclusion of chance flybys that may be selected as mergers if the EAGLE galaxies were selected based on proximity. Systems that are projected to merge, using a closing velocity extrapolation, within the next 0.3 Gyr (pre-merger) or are projected to have merged, again using a closing velocity extrapolation of the progenitors, within the last 0.25 Gyr (post-merger) were selected, along with a number of non-merging systems, and gri band images were created of these systems. Springel et al. (2005) have shown that the effects of a merger are visible for approximately 0.25 Gyr after the merger event while the pre-merger stage is much longer. However, we chose to have the pre and post merger period approximately equal as tests conducted with longer pre-merger times showed no improvement, see discussion in Sect. 4.2. We note, however, that the merger timing may suffer from imprecision as a result of the coarse time resolution of the EAGLE simulation, that is the time between snapshots, which becomes coarser at lower simulation redshift. Each galaxy was imaged at an assumed distance of 10 Mpc and each image contains all material within 100 kpc of the centre of the target galaxy and is 256\u2005\u00d7\u2005256 pixels, where 256 pixels corresponds to a physical size of 60 kpc. There are 537 pre-merger, 339 post-merger and 335 non-merging systems, each with six random projections to increase the size of the training set. Each of the six projections are treated as individual galaxies resulting in 3222 pre-merger, 2034 post-merger and 2010 non-merging galaxy images for training. The pre-mergers and post-mergers were combined to form the merger class, keeping the pre-merger image if the same galaxy appears in both sets.","Citation Text":["Springel et al. (2005)"],"Functions Text":["have shown that the effects of a merger are visible for approximately 0.25 Gyr after the merger event while the pre-merger stage is much longer. However, we chose to have the pre and post merger period approximately equal as tests conducted with longer pre-merger times showed no improvement, see discussion in Sect. 4.2."],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1068,1090]],"Functions Start End":[[1091,1412]]}
{"Identifier":"2017MNRAS.469.4620B___2017_Instance_1","Paragraph":"In Table 3, we present a detail comparison of our theoretical energy values from the layer n = 7. For this purpose, we have taken four data sets available in the literature. First, we compare with the ab initio values of Fritzsche (1995) because they have used a larger configuration space than previous ab initio computations (Vajed-Samii & MacDonald 1981; Huang et al. 1983). Secondly, theoretical values from the CHIANTI V.8 (Landi et al. 1999; Del Zanna et al. 2015) because these are considered the most appropriate data for astrophysical applications (Tr\u00e4bert 2014). Thirdly, with the most recent semi-empirically predicted and tentative experimental data values of Del Zanna & Badnell (2016). Fourthly, with the NIST data base (NIST Atomic Spectra Database (ver. 5.3) 2017), which is considered as the most critically analysed data base. Comparison with observed energies (Del Zanna & Badnell 2016) and with the NIST values shows that the present ab initio values are not fully converged to experimental values. In order to achieve better accuracy relative to experimental values, we suggest more balanced and extended reference configuration sets as a starting point of the orbital expansion. This, however, results in very large expansions of wave functions, which is computationally very expensive. Nevertheless, with the present set of calculations for CV correlations, results for the lowest 31 levels are better than obtained before in any ab initio investigation. A very good agreement with experiment is achieved to an accuracy better than 0.5 per cent for low lying levels and 0.8 per cent for higher lying levels, whereas this accuracy is 1\u20132\u2009per\u2009cent for the mcdf calculations of Fritzsche (1995) and for superstructure calculations available at CHIANTI V.8 (Landi et al. 1999; Del Zanna et al. 2015). In some more detail, there is no experimental connection in NIST data base to the 5 levels with J = 7\/2 and two levels with J = 9\/2 for the 3s23p43d configuration with other levels. The present results of energies for these seven levels with less than 0.45 per cent error with new tentative experimental assignments (Del Zanna & Badnell 2016) testify them. However, we suggest different level ordering than provided by semi-empirical formalism of Del Zanna & Badnell (2016). We briefly discuss later the shift in level ordering and reasons for it. We also note that there is 2100 cm\u22121 difference between the latest experimental energy (Del Zanna & Badnell 2016) and NIST data base for the level no. 31 (2D3\/2). The present value of 683 070 cm\u22121 confirms the experimental energy (Del Zanna & Badnell 2016), and we suggest a replacement in the NIST data base with new experimental value.","Citation Text":["NIST Atomic Spectra Database (ver. 5.3) 2017"],"Functions Text":["Fourthly, with the NIST data base","which is considered as the most critically analysed data base.","Comparison with observed energies","and with the NIST values shows that the present ab initio values are not fully converged to experimental values."],"Functions Label":["Uses","Uses","Similarities","Similarities"],"Citation Start End":[[735,779]],"Functions Start End":[[700,733],[782,844],[845,878],[906,1018]]}
{"Identifier":"2022MNRAS.516.2500C__Lin_et_al._2009_Instance_1","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"],"Functions Text":["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"],"Functions Label":["Motivation"],"Citation Start End":[[1492,1507]],"Functions Start End":[[1272,1465]]}
{"Identifier":"2021MNRAS.504.3316B__than_2000_Instance_3","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)"],"Functions Text":["As for the planet\u2019s phase offset,","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)."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[646,669]],"Functions Start End":[[612,645],[670,818]]}
{"Identifier":"2015MNRAS.446.1140T__Murray_et_al._2010_Instance_3","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"],"Functions Text":["There are at least three reasons why radiative feedback is an essential ingredient of the galaxy formation process.","Thirdly, it is difficult to explain the large gas turbulence values observed in star-forming regions without including the momentum input by radiation"],"Functions Label":["Background","Background"],"Citation Start End":[[2072,2090]],"Functions Start End":[[1353,1468],[1920,2070]]}
{"Identifier":"2018MNRAS.473.4279D__Markevitch_et_al._2002_Instance_1","Paragraph":"The Hubble Frontier Fields program1 (or HFF hereafter, Lotz et al. 2017) provides the most remarkably detailed examples of gravitational lensing by galaxy clusters, registering hundreds of multiply lensed galaxies for charting galaxy formation to unprecedented depths (see e.g. Jauzac et al. 2014, 2015a,b; Lam et al. 2014; Zitrin et al. 2014; Diego et al. 2015b,c, 2016; Kawamata et al. 2016; Limousin et al. 2016; Mahler et al. 2017). Furthermore, most of these HFF clusters are in a state of collision, enhancing their value for assessing the collisionality of dark matter (DM), a basic assumption of the standard particle interpretation of DM (Markevitch et al. 2002, 2004; Springel & Farrar 2007; Randall et al. 2008). Many clusters exhibit significant, but modest, offsets between the peak of the DM distribution and the centroid of the X-ray emission (Markevitch et al. 2004; Clowe et al. 2006; Mahdavi et al. 2007; Dawson et al. 2012; Menanteau et al. 2012), which is expected if DM is collisionless. These observations can provide a constraint on the DM cross-section (Markevitch et al. 2004; Randall et al. 2008). It is important that these differences are evaluated with the guidance of hydrodynamical models, as complex multibody interactions may also separate the DM from the plasma that can be explained without new physics, as is clearly evident in extreme cases of the bullet cluster (Mastropietro & Burkert 2008), and like the El Gordo cluster (Molnar & Broadhurst 2015), where high-speed collisions between pairs of clusters are ongoing. More direct evidence for collisional DM would be concluded from differences between the stellar and DM distributions as the stars behave like collisionless particles and we should expect the collisionless DM to follow the gravitational potential in the same way. Offsets between the position of the DM central peak and the peak of the luminous matter are difficult to explain with a standard \u039bcold dark matter but are naturally produced for reasonable values of the DM cross-section (Rocha et al. 2013). A difference of this nature has been claimed recently based on detailed lensing data in the centre of a cluster that contains four bright member galaxies (Massey et al. 2015). In case of the HFF clusters, it is interesting that our free-form analysis of MACS0146 also indicates a possible offset between the lensing-based centroids of the brightest galaxies and their luminous stellar centroids. These differences are subtle and it will be important to look at a larger sample and the model dependencies, and systematic uncertainties, in detail to support any claim of new physics.","Citation Text":["Markevitch et al. 2002"],"Functions Text":["Furthermore, most of these HFF clusters are in a state of collision, enhancing their value for assessing the collisionality of dark matter (DM), a basic assumption of the standard particle interpretation of DM"],"Functions Label":["Motivation"],"Citation Start End":[[648,670]],"Functions Start End":[[437,646]]}
{"Identifier":"2021ApJ...916...61F__Hutsem\u00e9kers_et_al._2019_Instance_1","Paragraph":"There are several popular scenarios employed to explain the observational features of CL AGNs. One scenario is that the broad emission lines are obscured by the torus or moving clouds over the observer\u2019s line of sight (Goodrich & Miller 1989), while only several CL AGNs can be explained by this scenario (Wang et al. 2019; Jaffarian & Gaskell 2020; Kokubo & Minezaki 2020). The features observed in most CL AGNs, e.g., the complex multiband spectral variabilities, and strong changes seen in the infrared or low level of polarization, strongly argue against the scattering (or obscuration) scenario (e.g., Sheng et al. 2017; Mathur et al. 2018; Stern et al. 2018; Hutsem\u00e9kers et al. 2019; Kynoch et al. 2019, and the references therein). Another attractive scenario is that the CL AGNs are indeed tidal disruption events (Merloni et al. 2015; Kawamuro et al. 2016; Yang et al. 2019; Padmanabhan & Loeb 2020; Ricci et al. 2020; Zhang 2021), while it is not a general mechanism to explain all CL AGNs, such as repeating CL AGNs. The changing look phenomenon triggered by the change of the mass accretion rate of the accretion disk is a rather straightforward model (Husemann et al. 2016; Liu et al. 2019; Ruan et al. 2019; Ai et al. 2020; Sniegowska et al. 2020); however, it has a fatal problem that propagation timescale of the gas in a thin accretion disk is much longer than the observed timescale in CL AGNs, unless the viscous thin disk model is somewhat revised (Lawrence 2018). Dexter & Begelman (2019) proposed that the magnetically elevated disk model could help to explain the changing look timescale. Based on this scenario, Scepi et al. (2021) suggested that a magnetic flux inversion in a magnetically arrested disk is able to explain the CL event in 1ES 1927+654. Sniegowska et al. (2020) suggested that a narrow unstable zone between the outer thin disk and the inner ADAF could cause the periodic outbursts in repeating CL AGNs. Recently, the effects of large-scale magnetic fields on this scenario for repeating CL AGNs have been studied in detail by Pan et al. (2021), and Lyu et al. (2021) suggested that the change of broad emission lines in Mrk 1018 might be regulated by the evolution of accretion disk.","Citation Text":["Hutsem\u00e9kers et al. 2019"],"Functions Text":["The features observed in most CL AGNs, e.g., the complex multiband spectral variabilities, and strong changes seen in the infrared or low level of polarization, strongly argue against the scattering (or obscuration) scenario (e.g.,"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[665,688]],"Functions Start End":[[375,606]]}
{"Identifier":"2018ApJ...866L...1S__Pecharrom\u00e1n_et_al._1999_Instance_5","Paragraph":"It was found that the complex dielectric function from Pecharrom\u00e1n et al. (1999) for the sample obtained by heating bayerite at 1273 K, assuming a spheroid with depolarization parameters of (0.35, 0.003), produced an opacity with 11, 20, 28, and 32 \u03bcm features, so this component was included in the models. However, with only this component, the observed 20 \u03bcm features in the residual spectra were found to be wider than those in the models. By adding the opacity of the sample obtained by heating boehmite at 1173 K, the width of the 20 \u03bcm feature could be matched. This was done using the complex dielectric function for the sample obtained by heating boehmite at 1173 K from Pecharrom\u00e1n et al. (1999), assuming a spheroid with depolarization parameters of (0.35, 0.035). The complex dielectric functions of the samples obtained by heating bayerite and boehmite to various temperatures (Pecharrom\u00e1n et al. 1999) were derived by modeling the reflectance spectra of pellets obtained by pressing powders of these materials under great pressure. This method required Pecharrom\u00e1n et al. (1999) to assume an effective medium theory, such that a pellet is a mixture of one of their samples with a matrix of air. Pecharrom\u00e1n et al. (1999) noted that heating bayerite at 500\u00b0C eliminates the XRD pattern of bayerite, and they note that at 700\u00b0C, the infrared reflectance spectrum of the boehmite sample no longer shows OH\u2212 stretching bands. This must mean that the samples obtained from heating bayerite at 1273 K and from heating boehmite at 1173 K are no longer bayerite or boehmite, respectively. XRD performed by Pecharrom\u00e1n et al. (1999) of the sample of bayerite prepared at 1273 K suggests only \u03b8-alumina was present, and their infrared and NMR spectroscopy confirms this. XRD of their sample obtained from heating boehmite to 1173 K (Pecharrom\u00e1n et al. 1999) suggests \u03b4-alumina to be present, though some amounts of \u03b8-alumina and \u03b1-alumina are present, as they deduce from XRD and infrared and NMR spectroscopy.","Citation Text":["Pecharrom\u00e1n et al. (1999)"],"Functions Text":["noted that heating bayerite at 500\u00b0C eliminates the XRD pattern of bayerite, and they note that at 700\u00b0C, the infrared reflectance spectrum of the boehmite sample no longer shows OH\u2212 stretching bands."],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1209,1234]],"Functions Start End":[[1235,1435]]}
{"Identifier":"2017AandA...605A..88L__Bernstein_et_al._2002_Instance_1","Paragraph":"Altogether, the approximately thirty molecules recently detected have confirmed the chemical complexity in the nebula, and generated our interest for the present study. Of these species, we will focus our attention on the seventeen species listed by molecular families in Table 1. As can be seen in this table, the WHISPER survey allowed the detection of some organic molecules in the Horsehead nebula, such as formaldehyde (H2CO) and methanol (CH3OH), which constitute key species in the likely synthesis of more complex organic molecules such as some prebiotic molecules (Bernstein et al. 2002; Mu\u00f1oz Caro et al. 2002; Garrod et al. 2008). Because they are detected in a wide variety of interstellar sources \u2013 in hot cores (Sutton et al. 1995; Ceccarelli et al. 2000), dark clouds (Bergman et al. 2011), shocked regions (e.g. Sakai et al. 2012; Codella et al. 2012; Tafalla et al. 2010) and even in comets (Mumma & Charnley 2011; Cordiner et al. 2015) \u2013 it is of prime importance to understand well how these precursor molecules form. H2CO is commonly thought to form both in the gas-phase and on grain surfaces, while CH3OH is believed to be only formed on grain surfaces (Garrod et al. 2006; Geppert et al. 2006). Guzman et al. (2013) reported the observations of these two molecules toward the Horsehead nebula in both the PDR and Core positions. Unable to reproduce the observed abundances of either H2CO or CH3OH at the PDR position with only pure gas-phase models, they concluded that, for this region, both species are formed on grain surfaces and then photodesorbed into the gas phase. On the other hand, at the Core position, a pure gas-phase model can reproduce the observed H2CO abundance, while photodesorption of ices is still needed to explain the observed abundance of CH3OH. Other organic molecules were reported in the Horsehead nebula as first detections in a PDR environment, including HCOOH (formic acid), CH2CO (ketene), CH3CHO (acetaldehyde), and CH3CCH (propyne) (Guzman et al. 2014). Their abundances were found to be higher at the PDR position than at the Core, revealing that complex organic chemistry is also occurring in UV-illuminated neutral gas (Guzman et al. 2014). Of these molecules, some \u2013 HCOOH, CH2CO, and CH3CHO \u2013 have now also been detected in the Orion bar PDR (Cuadrado et al. 2016, 2017). ","Citation Text":["Bernstein et al. 2002"],"Functions Text":["As can be seen in this table, the WHISPER survey allowed the detection of some organic molecules in the Horsehead nebula, such as formaldehyde (H2CO) and methanol (CH3OH), which constitute key species in the likely synthesis of more complex organic molecules such as some prebiotic molecules"],"Functions Label":["Uses"],"Citation Start End":[[574,595]],"Functions Start End":[[281,572]]}
{"Identifier":"2015AandA...579A.102B__Boselli_et_al._2009_Instance_2","Paragraph":"Once corrected for dust attenuation, H\u03b1 luminosities can be transformed into star formation rates (SFR, in\u2009M\u2299 yr-1) using a factor that depends on the assumed IMF and stellar model7: (10)\\begin{equation} {SFR = k({\\rm H}\\alpha) \\times L({\\rm H}\\alpha)_{\\rm cor}} . \\end{equation}SFR=k(H\u03b1)\u00d7L(H\u03b1)cor.We recall that this relation is valid only under the assumption that the mean star formation activity of the emitting galaxies is constant on a timescale of a few Myr, roughly comparable to the typical time spent by the stellar population responsible for the ionisation of the gas on the main sequence (Boselli et al. 2009; Boissier 2013; Boquien et al. 2014). The ionising stars are O and early-B stars, whose typical age is \u2272107 yr. The stationarity condition is generally satisfied in massive, normal, star-forming galaxies undergoing secular evolution. In these objects, the total number of OB associations is significantly larger than the number of HII regions under formation and of OB stars reaching the final stage of their evolution, thus their total H\u03b1 luminosity is fairly constant with time. This might not be the case in strongly perturbed systems or in dwarf galaxies, where the total star formation activity can be dominated by individual giant HII regions (Boselli et al. 2009; Weisz et al. 2012), and the IMF is only stochastically sampled (Lee et al. 2009; Fumagalli et al. 2011; da Silva et al. 2014). The HRS sample is dominated by relatively massive galaxies undergoing secular evolution. For these objects, Eq. (10) can thus be applied. The sample, however, also includes galaxies in the Virgo cluster region, where the perturbation induced by the cluster environment might have affected their star formation rate (e.g. Boselli & Gavazzi 2006, 2014). Models and simulations have shown that in these objects the suppression of star formation occurs on timescales of a few hundred Myr (Boselli et al. 2006, 2008a,b, 2014d). These timescales are relatively long compared to the typical age of O-B stars. The recent work of Boquien et al. (2014) has clearly shown that the Lyman continuum emission tightly follows the rapid variations in the star formation activity of simulated galaxies down to timescales of a few Myrs. We can thus safely consider that the linear relation between the H\u03b1 luminosity and the star formation rate given in Eq. (10) is satisfied in the HRS sample. ","Citation Text":["Boselli et al. 2009"],"Functions Text":["In these objects, the total number of OB associations is significantly larger than the number of HII regions under formation and of OB stars reaching the final stage of their evolution, thus their total H\u03b1 luminosity is fairly constant with time. This might not be the case in strongly perturbed systems or in dwarf galaxies, where the total star formation activity can be dominated by individual giant HII regions"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1271,1290]],"Functions Start End":[[855,1269]]}
{"Identifier":"2020ApJ...888...46C__Singh_et_al._1995_Instance_1","Paragraph":"The theoretical models mentioned above have made assumptions. The validity of these assumptions needs to be examined. Besides, all these models involve parameters to be determined. These problems could be better understood through numerical simulations. Hurlburt et al. (1994) performed a two-dimensional numerical simulation of compressible flow on downward overshooting. They investigated the dependence of subadiabatic extent \u03b4 on stability parameter S. They revealed a scaling law of \u03b4 \u221d S\u22121 in the penetrative layer and \u03b4 \u221d S\u22121\/4 in the overshooting layer, respectively. This result is in good agreement with Zahn\u2019s analytic model (Zahn 1991). Freytag et al. (1996) performed two-dimensional numerical simulations with a realistic description of radiation and ionization on A-type stars and DA white dwarf stars. They described the overshooting as a diffusive process, and derived an exponential decay parameter for the diffusion. Early attempts of low-resolution three-dimensional numerical simulations on overshooting were made by Singh et al. (1994; upward overshooting), Singh et al. (1995, 1998), and Saikia et al. (2000) (downward overshooting). The scaling laws derived from the numerical simulations of downward overshooting (Singh et al. 1995, 1998; Saikia et al. 2000) agree well with Zahn\u2019s analytical model. Only a qualitative result was given in the numerical simulations of upward overshooting (Singh et al. 1994). High-resolution numerical simulations of downward overshooting across a wide range of parameters were presented by Brummell et al. (2002). They confirmed the \u22121\/4 scaling law of the thermal adjustment overshooting layer, while the \u22121 scaling law of the nearly adiabatic penetrative layer was absent in the simulations. Based on a semianalytic model, Rempel (2004) argued that the absence of the nearly adiabatic penetrative layer is caused by the large energy flux specified in the numerical simulations. Numerical experiments on Boussinesq flow were performed by Korre et al. (2019). They reported steeper scaling laws of \u03b4 \u221d S\u22121\/3 or \u03b4 \u221d S\u22121\/2, depending on the steepness of the background radiative temperature gradient. Simulations with realistic physical variables on stellar core convection were performed by Browning et al. (2004) and Brun et al. (2005). The effects of rotation and magnetic field are considered. They found that the penetrative convection yields a prolate shape of a nearly adiabatic region. Kitiashvili et al. (2016) performed 3D radiative hydrodynamic simulations of the outer layer of a moderate-mass star (1.47 solar mass). Their result discovered a nearly adiabatic layer and a deeper subadiabatic layer. The recent work of Brun et al. (2017) simulated the differential rotation and overshooting in solar-like stars. Their result indicated that slow rotators favor a wider overshooting region near the poles and fast rotators at mid-to-low latitude. Hotta (2017) performed numerical simulations on the solar overshooting with very low energy fluxes F. He found that the overshooting distance obeys a scaling law of \u03b4 \u221d F0.31. K\u00e4pyl\u00e4 (2019) conducted numerical experiments on downward overshooting by considering the effect of the smoothness of the heat conduction profiles. He discovered that the power-law index of the overshooting distance on the energy flux is smaller in the smooth heat conduction profile than in the step profile. Efforts on prediction of 321D turbulent theory were made by Arnett et al. (2015) and Arnett & Moravveji (2017). They separated the overshooting region into three layers: a fully mixed layer, a partially mixed wave layer, and an extra diffusive mixing layer. With the scale analysis of turbulent plumes and eddies, Viallet et al. (2015) discussed the three possible regimes of turbulent overshooting: a diffusion-dominated regime (only mix composition), a penetrative regime (transition within the boundary layer), and an entrainment regime (mix both entropy and composition). The selection criterion of different regimes during a stellar evolution calculation is not well defined yet.","Citation Text":["Singh et al. (1995"],"Functions Text":["Early attempts of low-resolution three-dimensional numerical simulations on overshooting were made by","(downward overshooting)"],"Functions Label":["Background","Background"],"Citation Start End":[[1080,1098]],"Functions Start End":[[936,1037],[1132,1155]]}
{"Identifier":"2019AandA...622A.106M__Lanz_et_al._(2010)_Instance_2","Paragraph":"The standard single-frequency detection methods for point sources in the CMB and far IR are based on wavelet techniques (Vielva et al. 2003; Barnard et al. 2004; Gonz\u00e1lez-Nuevo et al. 2006) or on the matched filter (or MF hereafter, Tegmark & de Oliveira-Costa 1998; Herranz et al. 2002; Barreiro et al. 2003; L\u00f3pez-Caniego et al. 2006, see also Herranz & Vielva 2010 for a review.). Wavelets are well suited for the detection of compact sources due to their good position-scale determination properties, whereas the MF is the optimal linear detector-estimator because it provides the maximum signal-to-noise (S\/N) amplification for a source with a known shape (usually the point-spread function, or PSF hereafter, of the telescope) embedded in statistically homogeneous and spatially correlated noise. By default, these techniques are applicable only to single-frequency sky images: even for multiwavelength observatories such as the Herschel Space Observatory (Pilbratt et al. 2010) or Planck (Tauber et al. 2010), the standard detection pipelines have produced individual source catalogs for each frequency band (see e.g., Planck Collaboration VII 2011; Planck Collaboration XXVIII 2014; Planck Collaboration XXVI 2016; Maddox et al. 2018). The next logical step is to boost the signal of faint sources by combining the different bands into a single detection, that is, \u201cmultifrequency detection\u201d. Most of the blind component separation algorithms that are used for diffuse components in microwave and far IR astronomy can not deal with the high diversity of spectral behaviors associated to the different populations of extragalactic compact sources (see for example Leach et al. 2008). However, over the last few years a number of multifrequency compact source detection techniques have been proposed in the literature (Herranz & Sanz 2008; Herranz et al. 2009; Lanz et al. 2010, 2013; Planck Collaboration Int. LIV 2018). A review on the topic can be found in Herranz et al. (2012). In particular, if the spatial profile and the spectral energy distribution (SED) of the sources are known, and if the cross-power spectrum is known, or can be estimated from the data, the optimal linear detection method is the matched multifilter (or MMF hereafter, Herranz et al. 2002). Lanz et al. (2010) also showed that the MMF can be generalized for the case where the SED of the sources is not known. This generalization outperforms the single-frequency MF in terms of S\/N and can be used to infer the spectral index of synchrotron-dominated radio sources, as shown in Lanz et al. (2013). However, in this paper we will incorporate a specific SED to the MMF in order to derive a photometric redshift estimation of dusty galaxies and high-redshift star forming galaxies detected in the IR part of the spectrum1. We will do so by applying the multifrequency MMF filter to the first and second data releases of the Herschel Astrophysical Terahertz Large Area Survey (the Herschel-ATLAS or H-ATLAS, Eales et al. 2010), the largest single key project carried out in open time with the Herschel Space Observatory. We restrict our multifrequency analysis to the three wavelength bands covered by the SPIRE instrument aboard Herschel (Griffin et al. 2010), centered around 250, 350 and 500 \u03bcm. As discussed in Hopwood et al. (2010), Lapi et al. (2011), Gonz\u00e1lez-Nuevo et al. (2012), Pearson et al. (2013) and Donevski et al. (2018), the SPIRE bands are ideal for capturing the peak in the SED corresponding to dust emission of star-forming galaxies at z\u2004\u223c\u20042, that is redshifted from its rest-frame wavelength around 70\u2013100 \u03bcm to the SPIRE wavelengths: This is the redshift range where galaxies have formed most of their stars. At higher redshifts, dusty star-forming galaxies (DSFGs) occupy the most massive halos and are among the most luminous objects found at z\u2004\u2273\u20044 (Micha\u0142owski et al. 2014; Oteo et al. 2016; Ikarashi et al. 2017). These high-redshift DSFGs have markedly red colors as seen by SPIRE, with rising flux densities from 250 to 500 \u03bcm (the so-called \u201c500 \u03bcm-risers\u201d), and have received a great deal of attention in the recent years (see for example Ivison et al. 2016; Negrello et al. 2017; Strandet et al. 2017). The DSFGs, and particularly the 500 \u03bcm risers uncovered by Herschel, are providing much insight into the early star forming history of the universe. However, sensitivity and limited angular resolution severely constrain the power of this type of objects as astrophysical probes. The sensitivity of SPIRE allows for the direct detection of only the brightest, and thus rarest objects, at the bright end of the luminosity function. By means of our multifrequency MMF technique, we intend to enhance the detectability and statistical significance of very faint red objects in the H-ATLAS source catalog and so expand the list of reliable 500 \u03bcm-riser candidates.","Citation Text":["Lanz et al. (2010)"],"Functions Text":["also showed that the MMF can be generalized for the case where the SED of the sources is not known.","This generalization outperforms the single-frequency MF in terms of S\/N and can be used to infer the spectral index of synchrotron-dominated radio sources, as shown in Lanz et al. (2013)."],"Functions Label":["Background","Compare\/Contrast"],"Citation Start End":[[2277,2295]],"Functions Start End":[[2296,2395],[2396,2583]]}
{"Identifier":"2015ApJ...808...56M__Deming_et_al._2015_Instance_1","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"],"Functions Text":["The 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"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[748,766]],"Functions Start End":[[583,746]]}
{"Identifier":"2022MNRAS.516.5289M__Thompson_et_al._2015_Instance_2","Paragraph":"Given the number densities within the mass-dissociation index plane of Fig. 8, we now ask ourselves whether known dissociated clusters, such as the Bullet cluster, are expected in L210N1024NR? The Bullet Cluster has a mass of $\\sim 1.5 \\times 10^{15} \\, {\\rm M}_{\\odot }$ (e.g. Clowe et al. 2004; Brada\u010d et al. 2006; Clowe et al. 2006) and we estimated a dissociation index of SBullet \u223c 0.335 \u00b1 0.06. As seen in Fig. 8 there are no Bullet cluster analogues (structures of approximate mass and dissociation) in L210N1024NR, this is unsurprising as a simulation requires a significantly larger volume than that of L210N1024NR ((210cMpc\u2009h\u22121)3) to expect such an object (e.g. Lee & Komatsu 2010; Thompson & Nagamine 2012; Bouillot et al. 2015; Kraljic & Sarkar 2015; Thompson et al. 2015). From the distribution presented in Fig. 8, it is trivial to estimate the required cosmological volume (the effective volume, Veff) to expect structures of a given mass and dissociation index. By separating the 2D distribution on the mass-dissociation index planes into the component 1D distributions of mass and dissociation the effective volume is computed as\n(12)$$\\begin{eqnarray}\r\nV_\\text{eff}~^{-1} &amp;=&amp;\\int \\int \\,{\\rm{ d}} S \\, {\\rm{ d}} M \\phi (S, M) \\\\\r\n&amp;=&amp; \\int _{S_\\text{a}}^{S_\\text{b}} \\, {\\rm{ d}} S \\phi _S(S) \\int _{M_\\text{a}}^{M_\\text{b}} \\, {\\rm{ d}} M \\phi _M(M)~,\r\n\\end{eqnarray}$$where \u03d5S(S) is the number density function associated with S and $\\phi _\\mathit {M}(\\mathit {M})$ is the mass function presented in Fig. 7. Assuming a probable range of S = 0.335 \u00b1 0.06 and $1 \\lt M \\lt 2 \\times 10^{15} \\, {\\rm M}_{\\odot }$ we estimate a number density \u223c4.92 \u00d7 10\u221210 Mpc\u22123 or that an effective volume of \u223c2.03 Gpc3 would be required to observe a single Bullet-like cluster. This result is inline with the number density estimate of the order of \u223c10\u221210 Mpc\u22123 by Thompson et al. (2015), which improves on previous estimates (e.g. Lee & Komatsu 2010; Thompson & Nagamine 2012; Bouillot et al. 2015) due to more sophisticated halo finding methods (e.g. Behroozi, Wechsler & Wu 2013). Conversely, it was estimated by Kraljic & Sarkar (2015) (utilizing the same halo finder as Thompson et al. 2015) that given an effective volume of \u223c14.6 Gpc3, no Bullet cluster analogues are expected, however as indicated by a pairwise velocity distribution it would be expected that present binary halo\u2013halo orbits have the potential to form a Bullet-like object.","Citation Text":["Thompson et al. (2015)"],"Functions Text":["This result is inline with the number density estimate of the order of \u223c10\u221210 Mpc\u22123 by"],"Functions Label":["Similarities"],"Citation Start End":[[1881,1903]],"Functions Start End":[[1794,1880]]}
{"Identifier":"2022ApJ...935..135B__Mathur_1990_Instance_1","Paragraph":"All responses calculated in this paper only account for the direct response to a perturbing potential. In general, though, the response also has an indirect component that arises from the fact that neighboring regions in the disk interact with each other gravitationally. This self-gravity of the response, which we have ignored, triggers long-lived normal-mode oscillations of the slab that are not accounted for in our treatment. Several simulation-based studies have argued that including self-gravity is important for a realistic treatment of phase spirals (e.g., Darling & Widrow 2019a; Bennett & Bovy 2021). Using the Kalnajs matrix method (Kalnajs 1977; Binney & Tremaine 2008), we have made some initial attempts to include the self-gravity of the response in our perturbative analysis, along the lines of Weinberg (1991). Our preliminary analysis shows that the self-gravitating response is a linear superposition of two terms: (i) a continuum of modes given in Equation (12), dressed by self-gravity, that undergo phase mixing and give rise to the phase spiral; and (ii) a discrete set of modes called point modes or normal modes (see Mathur 1990; Weinberg 1991) that follow a dispersion relation. The continuum response can be amplified by self-gravity when the continuum frequencies, n\u03a9\nz\n + kv\n\nx\n, are close to the point-mode frequencies, \u03bd. Depending on the value of k, the normal modes can be either stable or unstable. Araki (1985) finds that in an isothermal slab the bending normal mode undergoes fire hose instability below a certain critical wavelength if \u03c3\n\nz\n\/\u03c3 \u2272 0.3, while the breathing normal mode becomes unstable above the Jeans scale. In the regime of stability, the normal modes are undamped oscillation modes in absence of lateral streaming (Mathur 1990) but are Landau damped otherwise (Weinberg 1991). For an isothermal slab with typical MW-like parameter values, the point modes are strongly damped since their damping timescale (inverse of the imaginary part of \u03bd) is of order their oscillation period (inverse of the real part of \u03bd), which turns out to be of order the vertical dynamical time, h\n\nz\n\/\u03c3\n\nz\n. Moreover, the normal-mode oscillations are coherent oscillations of the entire system, independent of the vertical actions of the stars, and are decoupled from the phase spiral in linear theory since the full response is a linear superposition of the two. Based on the above arguments, we conclude that self-gravity has little impact on the evolution of phase spirals in the isothermal slab, at least in the linear regime. We emphasize that Darling & Widrow (2019a), who found their phase spirals to be significantly affected by the inclusion of self-gravity, assumed a perturber-induced velocity impulse with magnitude comparable to the local velocity dispersion in the solar neighborhood; hence, their results are likely to have been impacted by nonlinear effects. Moreover, the self-gravitating response of an inhomogeneous disk embedded in a DM halo, as in the simulations of Darling & Widrow (2019a), can be substantially different from that of the isothermal slab. We intend to include a formal treatment of self-gravity along the lines of Weinberg (1991) in future work.","Citation Text":["Mathur 1990"],"Functions Text":["Our preliminary analysis shows that the self-gravitating response is a linear superposition of two terms: (i) a continuum of modes given in Equation (12), dressed by self-gravity, that undergo phase mixing and give rise to the phase spiral; and (ii) a discrete set of modes called point modes or normal modes (see","that follow a dispersion relation."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1145,1156]],"Functions Start End":[[831,1144],[1173,1207]]}
{"Identifier":"2018AandA...620A..80M__observations,_Gerin_et_al._2017_Instance_1","Paragraph":"Methyl formate is the species for which the largest number of transitions were identified. In total 24 transitions from both A and E species were observed (18 unblended), with energy levels between 100 and 500 K, allowing to constrain the properties of the emitting region from the LTE model. We ran a set of models using one source component of 0.6\u2033 in size (the angularresolution of the ALMA band 6 observations), and we modified the temperature and column density to fit the observed spectra. However, it was difficult to fit all the lines with only one component. The problem became more evident when using the deuterated substitutions, CH3OCOD and CH2DOCOH, and the vibrationally excited states \u03bdt = 1 and \u03bdt = 2. We have detected 12 transitions from CH3OCOD and four transitions from CH2DOCOH (three and two, respectively, are blended with other species), with upper energy levels between 90 and 170 K. In both cases the model gave better results when decreasing the temperature with respect to the best CH3OCOH model. On the other hand, the vibrationally excited lines (13 unblended transitions from \u03bdt = 1, and one transition from \u03bdt = 2), supported the high temperature model. Therefore, we built a model of the source with two components. One warm and compact component at 200 K and a size of 0.35\u2033 (that of the continuum compact emission obtained by ALMA Band 7 observations, Gerin et al. 2017), and another larger (0.6\u2033, approximately the angular resolution of the Band 6 observations) and colder component at 60 K. Figure B.4 shows two examples of the best one component models obtained using the main CH3OCOH (200 K), and the deuterated methyl formate lines (60 K), together with the final two component model. While some species give reasonable agreement with observations using either model (mostly those which observed transitions have similar energy levels), there are others (besides the ones discussed above) that fail in one of the models (e.g. HNCO and t-CH3CH2OCOH). We have checked that using other temperatures (between 50 and 250 K) in the one component models does not reproduce satisfactorily the observations. For the vibrationally excited states, we only used the compact component at 200 K, since the low temperature one does not contribute much to the line intensities. Indeed, energy levels of the observed transitions range between 300 and 500 K, supporting this scenario. We have also observed five transitions of 13CH3OCOH, with energies \u2264200 K and intensities of the order of 500 mK or below. Since the lines are weak, any of the models give reasonable agreement with the observations. Acetaldehyde is the other species for which many lines are observed. We have detected 15 transitions (10 unblended) from the main species, with energies between 100 and 300 K. We have also detected one transition arising from the deuterated substitution CH3CDO, and four transitions (two of them blended) from the vibrational state \u03bdt = 1, with energy levels of ~300\u2013400 K. We used the same model as methyl formate, except for CH3CHO \u03bdt = 1, for which only the compact and hot component was considered. The model allows to reproduce all transitions reasonably well, except for a few lines that are overestimated: two CH3OCOH transitions (Fig. B.1, second panel), three CH3OCOH \u03bdt = 1 lines (top panel in Fig. B.1, and second panel in Fig. B.3), and two lines from CH3CHO (see Fig. B.3, second and third panels). These are transitions with the lowest energies (~100 K) and the highest line strengths and Einstein coefficients, and may consequently have high opacities. We estimate the error in the fit to be ~50%.","Citation Text":["Gerin et al. 2017"],"Functions Text":["One warm and compact component at 200 K and a size of 0.35\u2033 (that of the continuum compact emission obtained by ALMA Band 7 observations,"],"Functions Label":["Uses"],"Citation Start End":[[1387,1404]],"Functions Start End":[[1249,1386]]}
{"Identifier":"2021MNRAS.504..146V__Vink_&_Gr\u00e4fener_2012_Instance_1","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"],"Functions Text":["this mass is significantly diminished via stellar winds already during core hydrogen (H) burning"],"Functions Label":["Background"],"Citation Start End":[[955,975]],"Functions Start End":[[857,953]]}
{"Identifier":"2021AandA...648A..14R__G\u00fcrkan_et_al._(2018)_Instance_1","Paragraph":"To obtain a more complete picture of the physical processes that shape star formation in the early Universe, very deep radio surveys overwide areas of sky are required to complement deep submillimetre surveys. Previous work has used high-resolution radio observations, typically at 1.4 GHz, as a method of pinpointing the position of submillimetre sources detected in single-dish surveys (Ivison et al. 2002; Chapman et al. 2005); however, such work has been limited by the depth of available radio sky surveys,with dedicated deep surveys over only small regions of sky, resulting in a view biased towards the brighter radio sources. Studies have largely included limited radio spectral coverage of SMGs, focusing on the nature of the FIRC and its relation to properties such as stellar mass and redshift (e.g. Yun et al. 2001; Ivison et al. 2010; Smith et al. 2014). The Low Frequency Array (LOFAR; van Haarlem et al. 2013) has opened up new ways of studying galaxies in the radio, and a number of studies have used LOFAR\u2019s capabilities to investigate this relationship between star formation and radio luminosity in the low-frequency regime \u2013 for example G\u00fcrkan et al. (2018), Read et al. (2018), Smith et al. (2021), and Wang et al. (2019a). However, these studies generally investigate the statistical properties of large samples of galaxies, in optically selected samples at low redshift (z \u22722), rather than probing the shapes of individual radio spectra. Thomson et al. (2019) conducted an in-depth study of high-frequency (>610 GHz) spectral curvature in SMGs, finding evidence of curved spectra that they attributed to spectral ageing of the synchrotron emission from star formation; their results implied estimated starburst ages consistent with expected SMG lifetimes. Studies at low frequencies, where we may be able to observe absorption processes affecting the shape of the spectrum, have been hampered by a lack of sufficiently deep, wide-area data. More comprehensive observations of the shape of the radio spectrum, extending to lower frequencies, can provide a probe of the physical conditions that give rise to extreme star formation in SMGs. Calistro Rivera et al. (2017) exploit LOFAR\u2019s frequency range to investigate the spectral shapes of star-forming galaxies and AGN, finding evidence of low-frequency spectral flattening in the star-forming sample. This sample is also constrained in redshift, focusing on local galaxies rather than the peak of star formation at z > 2, and so does not probe the bulk of the highly star-forming SMG population. Chy\u017cy et al. (2018) also find weak spectral flattening in local star-forming galaxies with LOFAR but largely attribute this slight effect to synchrotron losses, predicting stronger low-frequency spectral flattening due to free\u2013free absorption at high redshift, where galaxies with high SFRs are more common.","Citation Text":["G\u00fcrkan et al. (2018)"],"Functions Text":["The Low Frequency Array","has opened up new ways of studying galaxies in the radio, and a number of studies have used LOFAR\u2019s capabilities to investigate this relationship between star formation and radio luminosity in the low-frequency regime \u2013 for example","Read et al. (2018), Smith et al. (2021), and Wang et al. (2019a)."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[1157,1177]],"Functions Start End":[[868,891],[925,1156],[1179,1244]]}
{"Identifier":"2022MNRAS.516.4833J__Boffin_&_Jorissen_1988_Instance_1","Paragraph":"It is now clear that binary stars hold the key to understanding the formation of many planetary nebulae (PNe; Jones & Boffin 2017a). However, with much of the recent focus being placed on common envelope evolution (Boffin & Jones 2019), the role of wider binaries remains almost completely unconstrained (Tyndall et al. 2013). To date, only three wide binary central stars of PNe have confirmed periods (Jones et al. 2017) with other systems either being so wide that they are resolved (Ciardullo et al. 1999) or displaying composite spectra with giant companions (Tyndall et al. 2013). Four central star systems are known to be barium stars (Bond, Pollacco & Webbink 2003; Miszalski et al. 2012, 2013; L\u00f6bling, Boffin & Jones 2019) with only one having a known orbital period (LoTr 5; Jones et al. 2017; Aller et al. 2018). Barium stars are mostly \u2013 but not only: see e.g. Escorza et al. (2019) \u2013 giant stars of spectral type G-K that display enhanced abundances of carbon and s-process elements such as barium and strontium, and are now known to all be long-period binaries (100\u2009d \u2272 Porb \u2272 104\u2009d; McClure, Fletcher & Nemec 1980; Jorissen et al. 2019). The chemical contamination of barium stars is believed to be due to accretion of chemically enriched material from an evolved binary companion, likely through wind (Boffin & Jorissen 1988) or wind Roche lobe overflow (WRLOF; Theuns & Jorissen 1993; Nagae et al. 2004). WRLOF will occur when the acceleration radius of the asymptotic giant branch (now a white dwarf in these barium star systems) star\u2019s stellar wind is comparable to or greater than its Roche lobe radius (Mohamed & Podsiadlowski 2012), meaning that the wind itself is strongly influenced by the binary potential and can be accreted on to the companion via the first Lagrangian point. This accreted material chemically contaminates the surface of the companion as well as increases its spin rate due to the conservation of angular momentum (Theuns, Boffin & Jorissen 1996). Intriguingly, the amount of material required to account for the observed contamination is often in excess of that which would be expected to result in critical rotation rates in the companion, providing an indication that significant (and as yet not understood) angular momentum losses must be experienced in these systems (Matrozis, Abate & Stancliffe 2017). PNe with barium central stars offer an important window into this process as the presence of the short-lived nebula (\u03c4 \u2264 30\u2009000 yr; Jacob, Sch\u00f6nberner & Steffen 2013) means that the mass transfer has occurred too recently for significant changes in the companion\u2019s spin rate, and moreover the nebula itself traces the mass-loss history of the system as it is formed from the material that has escaped the binary potential (Jones 2018).","Citation Text":["Boffin & Jorissen 1988"],"Functions Text":["The chemical contamination of barium stars is believed to be due to accretion of chemically enriched material from an evolved binary companion, likely through wind"],"Functions Label":["Background"],"Citation Start End":[[1319,1341]],"Functions Start End":[[1154,1317]]}
{"Identifier":"2015ApJ...803...96S__T\u00f6r\u00f6k_et_al._2004_Instance_1","Paragraph":"Since their initial discovery with AIA on board the Solar Dynamics Observatory (SDO), HCs have been generally regarded as proxies for magnetic flux ropes (MFRs; volumetric plasma structures with magnetic field lines that wrap around a central axis). This is supported by the following observational studies: (1) Cheng et al. (2014a) observed an HC that showed helical threads winding around an axis. Simultaneously, cool filamentary materials descended spirally down to the chromosphere, providing direct observational evidence of an intrinsic helical structure for the HC. (2) Cheng et al. (2011) reported that an HC can grow during an eruption, similar to the MFR growth process according to the classic magnetic reconnection scenario in eruptive flares. Song et al. (2014a) presented the formation process of an HC during a coronal mass ejection (CME) and found that the HC was formed from coronal arcades through magnetic reconnection. These works further support the idea that an HC is an MFR structure based on the relation between the HC and magnetic reconnection. (3) Cheng et al. (2014b) found that an HC was initially cospatial with a prominence. Then a separation of the HC top from that of the prominence was observed during the eruption initiated by the ideal kink instability (T\u00f6r\u00f6k et al. 2004). It is widely accepted that a prominence\/filament can exist at the dip of a flux rope (Rust & Kumar 1994). Therefore, this observation offered further important support for the idea that an HC is an MFR; beside an HC, several lines of observations in the lower corona have also been proposed as MFRs, including sigmoid structures in an active region (Titov & D\u00e9moulin 1999; McKenzie & Canfield 2008) and coronal cavities in quiescent regions (Wang & Stenborg 2010). A sigmoid has either a forward or reverse S-shape with enhanced X-ray emissions (implying an entity of high temperature) with its center straddling along the polarity inversion line of the hosting active region. Zhang et al. (2012) showed that the HC initially appeared as a sigmoidal structure and then changed to a semi-circular shape. Therefore, a sigmoid and an HC might represent the same structure, and their different shapes are likely from different perspectives and evolution phases. Both structures feature a high temperature, a possible result of flare magnetic reconnection (e.g., Song et al. 2014a, 2014b). A coronal cavity, on the other hand, which is observed as a dark circular or oval structure above the solar limb in coronal images with temperatures close to the background corona (Fuller et al. 2008; Gibson et al. 2010; Kucera et al. 2012), is also interpreted as an MFR. As mentioned, the long-studied feature of solar filaments\/prominences shown best in H\u03b1 images has been interpreted as being situated along the dip in the MFR. Therefore, a prominence lying in the dip of a coronal cavity is not rare. The eruption of a coronal cavity (or filament) from a quiescent region does not show a high-temperature signature like an HC, which might be attributed to a lack of obvious heating acquired from the weak magnetic reconnection (e.g., Song et al. 2013).","Citation Text":["T\u00f6r\u00f6k et al. 2004"],"Functions Text":["Then a separation of the HC top from that of the prominence was observed during the eruption initiated by the ideal kink instability","Therefore, this observation offered further important support for the idea that an HC is an MFR"],"Functions Label":["Differences","Similarities"],"Citation Start End":[[1291,1308]],"Functions Start End":[[1157,1289],[1417,1512]]}
{"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"],"Functions Text":["The overall excesses within 805 and 1186 h of HESS exposures"],"Functions Label":["Uses"],"Citation Start End":[[1272,1282]],"Functions Start End":[[1210,1270]]}
{"Identifier":"2018AandA...612A..34D__Dexter_et_al._2010_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":["Dexter et al. 2010"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[825,843]],"Functions Start End":[[479,756]]}
{"Identifier":"2019ApJ...882..131M__Clements_et_al._2018_Instance_1","Paragraph":"The ice layer covering refractory grains in dense MCs consists mainly of amorphous water ice (Tielens & Allamandola 1987). A key property of the icy mantle is its porosity, which determines its ability to adsorb, desorb, and trap atoms and molecules. The actual degree of porosity of interstellar ices is still debated. There are indeed indications that the buildup of the ice in cold environments results in the formation of pores as shown by laboratory measurements (e.g., Dohn\u00e1lek et al. 2003; Bossa et al. 2014) and simulations (Cuppen & Herbst 2007; Clements et al. 2018). These two studies, based on Monte Carlo simulations, suggest a significant level of porosity at the surface\/subsurface, in particular in cold and dense environments. On the other hand, UV photons, exothermic reactions, and energetic ions tend to compact astrophysical ices. Estimates of the timescale for mantle compaction calculated from experiments range from a few up to 50 Myr (Raut et al. 2008; Palumbo et al. 2010; Accolla et al. 2011). It has also been suggested that the compaction of icy mantles may not be completed within the typical lifetime of MCs if they consist of ice mixtures, since compaction is slower for mixtures than for pure ices (Palumbo 2006). Transient events like the impact of energetic ions can also generate cavities in the ice surface and\/or subsurface. Such energetic events induce a strong local increase of the ice temperature, produce fragments, and are associated with sputtering and evaporation along the ion track. Such cavities are short-lived since the heated ice rearranges while relaxing (Mainitz et al. 2016). The missing O\u2013H dangling features near 3700 cm\u22121 in astronomical spectra have been taken as an indication of a low level of porosity of interstellar ices. However, laboratory data and simulations indicate that the absence of the O\u2013H dangling modes does not necessarily imply the complete absence of porosity (Raut et al. 2007; Isokoski et al. 2014).","Citation Text":["Clements et al. 2018"],"Functions Text":["There are indeed indications that the buildup of the ice in cold environments results in the formation of pores as shown b","and simulations","These two studies, based on Monte Carlo simulations, suggest a significant level of porosity at the surface\/subsurface, in particular in cold and dense environments."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[555,575]],"Functions Start End":[[320,442],[516,531],[578,743]]}
{"Identifier":"2022MNRAS.514.1169A__Fisher_et_al._1995_Instance_1","Paragraph":"SHT has some inherent geometrical advantages, especially for large-angle and deep surveys. A significant advantage is that the spherical coordinates apply to wide-angle surveys like BOSS without any flat-sky approximation. A traditional P(k) analysis relies on the flat-sky approximation to distinguish between transverse and line-of-sight modes, which is critical for accurately representing RSDs and distinguishing continuum foregrounds from line signal. By contrast, in the SHT formalism, both foregrounds and linear RSDs take exact, simple forms. This geometric advantage is shared with the related analysis technique of spherical Fourier-Bessel (SFB) decomposition (Fisher et al. 1995; Heavens & Taylor 1995; Leistedt et al. 2012; Rassat & Refregier 2012; Yoo & Desjacques 2013; Lanusse et al. 2015; Liu, Zhang & Parsons 2016; Grasshorn Gebhardt & Dor\u00e9 2021), which, in addtion to spherical harmonic decomposition, involves a further Fourier-Bessel transform and data-compression along the line-of-sight. However, SFB does not share another important feature of SHT, which is its ability to capture redshift-dependent change over cosmological time in deep surveys. For deep surveys, structure growth and changes in star formation rate break the assumption of translational invariance in the line-of-sight direction, rendering the P(k) or SFB statistic insufficient. However, since C\u2113(z, z\u2032) does not compress data along the line-of-sight direction, it describes redshift evolution. A study (Mondal et al. 2022) of simulated 21-cm data from the EoR found that an implementation of the SHT technique, MAPS (Mondal, Bharadwaj & Datta 2018; Mondal et al. 2019, 2020), obtains \u223c2 times more stringent error bars on model parameters than techniques that fail to capture redshift evolution due to data compression along the line-of-sight, such as P(k) or SFB. A final geometric advantage is that C\u2113(z, z\u2032) describes the data in observing coordinates of angle and frequency (or, equivalently, redshift) rather than re-gridding the data on to cosmological distances in an assumed cosmological model. An MCMC likelihood analysis that constrains cosmological parameters would therefore not need to recompute the data statistic at each step, which in principle would be needed for a P(k) or SFB analysis.","Citation Text":["Fisher et al. 1995"],"Functions Text":["This geometric advantage is shared with the related analysis technique of spherical Fourier-Bessel (SFB) decomposition","which, in addtion to spherical harmonic decomposition, involves a further Fourier-Bessel transform and data-compression along the line-of-sight.","However, SFB does not share another important feature of SHT, which is its ability to capture redshift-dependent change over cosmological time in deep surveys."],"Functions Label":["Similarities","Background","Differences"],"Citation Start End":[[671,689]],"Functions Start End":[[551,669],[865,1009],[1010,1169]]}
{"Identifier":"2022MNRAS.516.3900A__Cazaux_et_al._2022_Instance_1","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"],"Functions Text":["While Sn can also be formed from pure H2S ice by photo processing","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."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1018,1036]],"Functions Start End":[[951,1016],[1039,1222]]}
{"Identifier":"2018AandA...620A..31M__Palla_&_Stahler_1992_Instance_1","Paragraph":"Any massive stellar source from ~8 M\u2299 can burn hot enough to completely ionise the surrounding molecular material to form an H\u202fII region (Wood & Churchwell 1989; Churchwell 1990; Kurtz 2005) and destroy any complex chemical tracers that are traditionally used to understand the kinematics of their natal environments and their accretion discs. Initial arguments by Walmsley (1995) argued that high accretion rates will cause a very dense H\u202fII region that is optically thick to radio emission and thus the H\u202fII would not be seen. However, for heavily accreting massive YSOs the onset of the H\u202fII region could be delayed via stellar bloating (Palla & Stahler 1992; Hosokawa & Omukai 2009; Hosokawa et al. 2010; Kuiper & Yorke 2013) where the effective temperature of the star is much cooler than it would be considering a main sequence star of the same luminosity. In these models, a halt or considerable reduction in accretion (\u226a10\u22123 M\u2299 yr\u22121), or growth beyond ~30\u221240 M\u2299, will result in the YSO contracting to a \u201cmain-sequence\u201d configuration, heating significantly, and being able to create an H\u202fII region. The fine details of this transition are still somewhat unclear since they depend on the assumed accretion law and the initial conditions chosen for the stellar evolution calculations (Haemmerl\u00e9 & Peters 2016). Observations in the infra-red (IR) have been made in search of cool stellar atmospheres that point to bloated stars, however these studies remain inconclusive (Linz et al. 2009; Testi et al. 2010). Alternative scenarios are that H\u202fII regions can be gravitationally trapped at very early ionisation stages (see Keto 2003, 2007), or flicker due to chaotic shielding of the ionising radiation by an accretion flow, leading to a non-monotonous expansion (Peters et al. 2010b). Hyper Compact (HC) H\u202fII regions (0.03 pc) are thought to be the earliest ionisation stage and therefore could relate to the halt of accretion, be a marker of a transition phase.","Citation Text":["Palla & Stahler 1992"],"Functions Text":["However, for heavily accreting massive YSOs the onset of the H\u202fII region could be delayed via stellar bloating","where the effective temperature of the star is much cooler than it would be considering a main sequence star of the same luminosity. In these models, a halt or considerable reduction in accretion (\u226a10\u22123 M\u2299 yr\u22121), or growth beyond ~30\u221240 M\u2299, will result in the YSO contracting to a \u201cmain-sequence\u201d configuration, heating significantly, and being able to create an H\u202fII region."],"Functions Label":["Background","Background"],"Citation Start End":[[641,661]],"Functions Start End":[[529,639],[730,1105]]}
{"Identifier":"2019AandA...625A.121M__Beaug\u00e9_&_Nesvorn\u00fd_2012_Instance_2","Paragraph":"The final location of close-in giant planets in our models reflects the strength of the tides that we include in our modeling, which play a very important role in the decay of planetary orbits. These are dynamical tides (e.g., Lai 1997; Ivanov & Papaloizou 2004, 2007, 2011) and in our simulations we used a formulation given by Ivanov & Papaloizou (2007) as described inSect. 2. However, the impulse approximation used in the evaluation of dynamical tides becomes a poor approximation when the circularization proceeds (e.g., Mardling 1995a,b) and the eccentricity becomes low. Equilibrium tides become then effective (e.g., Beaug\u00e9 & Nesvorn\u00fd 2012 and references therein) and the tidal evolution may occur on a longer timescale. In short, at the beginning of the orbital evolution that leads to the formation of hot\/warm Jupiters, dynamical tides are important in forcing the decay of the orbit. In the last part of the dynamical evolution when the eccentricity has become low, equilibrium tides are more important in determining the location where the planet stops. Unfortunately, at present it is not known when and how the two tides switch. When we change the magnitude of two tides, the final location of the planets can be adjusted (Beaug\u00e9 & Nesvorn\u00fd 2012). However, rather than introducing artificial effects, we continue to use dynamical tides in our simulations even for low eccentricities but we stop our simulations when the energy, decreasing from the tide at the pericenter, overcomes the orbital energy leading to a clustering of tidally circularized planets around 0.02 au. However, the final distribution of the inclination of the planets does not depend on this choice and highly misaligned planets would be produced anyway. Since the tidal evolution of planets with arbitrary inclinations is still not well known, we assume that planetary inclination is not significantly changed during tidal evolution (Barker & Ogilvie 2009). Thus, the planets maintain the inclination they have when the circularization begins.","Citation Text":["Beaug\u00e9 & Nesvorn\u00fd 2012"],"Functions Text":["When we change the magnitude of two tides, the final location of the planets can be adjusted"],"Functions Label":["Uses"],"Citation Start End":[[1239,1261]],"Functions Start End":[[1145,1237]]}
{"Identifier":"2020MNRAS.491.5073P__Sutherland_&_Saunders_1992_Instance_1","Paragraph":"The catalogue used in this work is based on a far-IR sample selected in the \u223c2 deg2-wide COSMOS field and obtained within the Herschel-PEP survey (Lutz et al. 2011). We consider the latest released blind catalogue selected at 160-\u03bcm (DR1, 7047 sources) with >3\u03c3 flux density, corresponding to a flux limit of \u223c9.8 mJy. The choice of considering as parent sample a far-IR catalogue is guided by the necessity of having several detections at different wavelengths for each system to constrain the dust masses, and yet with a very simple selection function (see Section 3.1). From the original 160-\u03bcm selection, we built a multiband catalogue taking advantage of the extensive multiwavelength coverage in the COSMOS field. Concerning the other far-IR PACS band (100 \u03bcm) and the mid-IR 24-\u03bcm band, we use the association available in the DR1 release and based on the maximum likelihood technique (Sutherland & Saunders 1992; Ciliegi et al. 2001). For the cross-match with the SPIRE far-IR bands (250, 350, and 500 \u03bcm), we used the same catalogue considered in previous PACS-based works (i.e. Gruppioni et al. 2013; Delvecchio et al. 2015), the ones provided by the HerMES collaboration (Roseboom et al. 2010) using the Spitzer-MIPS 24-\u03bcm positions as priors to extract the SPIRE fluxes. Finally, the IRAC\/optical\/UV fluxes taken from the COSMOS2015 catalogue (Laigle et al. 2016) were merged to the PACS-160-\u03bcm sample by matching the 24-\u03bcm counterparts listed in both. The COSMOS2015 is NIR selected, where objects have been detected from the sum of the UltraVISTA-DR2 YJHK and z++ images. By construction, in comparison to the previous i-selected catalogue, this catalogue is missing a fraction of blue, faint, star-forming galaxies (Laigle et al. 2016). For this reason, we decided to cross-match the far-IR sources with no counterparts in the COSMOS2015 catalogue with the Ilbert et al. (2009) i-selected catalogue. Totally, among the 160-\u03bcm selected sources (7047), 6002 are with 24-\u03bcm counterparts (${\\sim }86{{\\ \\rm per\\ cent}}$), of which 5993 with available NIR or optical counterparts (${\\sim }99.9{{\\ \\rm per\\ cent}}$, 5783 in the COSMOS2015 and 210 in the Ilbert et al. 2009 catalogue). While the cross-matching with the optical\/NIR bands does not involve almost any source loss, a moderate (14${{\\ \\rm per\\ cent}}$) but not negligible fraction of the far-IR sources does not have 24-\u03bcm counterparts. A fraction of these sources are likely spurious sources, as shown by the simulations done for the DR1 release PEP catalogue (\u223c5 per\u2009cent at the 3\u03c3 flux level). In Section 3.1, we will describe our method to correct for incompleteness and for the presence of spurious systems.","Citation Text":["Sutherland & Saunders 1992"],"Functions Text":["Concerning the other far-IR PACS band (100 \u03bcm) and the mid-IR 24-\u03bcm band, we use the association available in the DR1 release and based on the maximum likelihood technique"],"Functions Label":["Uses"],"Citation Start End":[[893,919]],"Functions Start End":[[720,891]]}
{"Identifier":"2020AandA...638A..16T__Barnes_(2017)_Instance_1","Paragraph":"Figure 12 shows the results of our tidal evolution calculations. The left panel of Fig. 12 shows the planetary rotational evolution of GJ 1148 b due to star\u2013planet tides. After ~850 Myr, GJ 1148 b reaches a rotation period that is 2\u22153 of the orbital period, and remains there with Prot = 27.5 d. During the integration the planetary semi-major axis and eccentricity are mostly unaffected. An asymptotic rotation period that is shorter than synchronous and 2\/3 of the orbital period is expected for eb \u22730.24 in the constant Q tidal model (Goldreich & Peale 1966; Cheng et al. 2014). The time for GJ 1148 b to reach asymptotic rotation is inversely proportional to the initial Prot, as long as the initial Prot is much less than 27.5 d, and it depends on the other parameters of GJ 1148 b according to Eq. (3) of Barnes (2017) and Eq. (15) of Cheng et al. (2014). The rotational period of GJ 1148 b is thus very likely much longer than the orbital periods of the hypothetical exomoons, which could be dynamically stable only with orbital periods between 0.7 and 2 d. The right panel of Fig. 12 shows that the longer rotational period of GJ 1148 b (Prot = 27.5 d) leadsto strong orbital decay of the stable exomoon orbits due to tidal interactions with the planet. An exomoon eventually reaches the Roche limit where it is tidally disrupted by the gas giant. Not even one hypothetical \u201cstable\u201d exomoon in the context of Sect. 5.3.1 had survived this test. The maximum time a Mars-like exomoon could survive is ~55 M yr, while for Titan-like moons the maximum survival time is longer, ~255 M yr. The latter is longer by roughly the mass ratio of Mars to Titan, which can be understood from Eq. (2) of Barnes (2017) and Eq. (16) of Cheng et al. (2014). These timescales are optimistic since the orbital decay would start before the planet reaches the asymptotic spin state. In both cases the survival times are much shorter than the age of the system. Therefore, given the relatively fast orbital decay in the small stable region around the planet, we conclude that exomoons around GJ 1148 b are unlikely to exist.","Citation Text":["Barnes (2017)"],"Functions Text":["The time for GJ 1148 b to reach asymptotic rotation is inversely proportional to the initial Prot, as long as the initial Prot is much less than 27.5 d, and it depends on the other parameters of GJ 1148 b according to Eq. (3) of"],"Functions Label":["Uses"],"Citation Start End":[[811,824]],"Functions Start End":[[582,810]]}
{"Identifier":"2019ApJ...883...76R__Strateva_et_al._2005_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":["Strateva et al. 2005"],"Functions Text":["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.,"],"Functions Label":["Background"],"Citation Start End":[[533,553]],"Functions Start End":[[314,511]]}
{"Identifier":"2022MNRAS.513.5245A__Done_&_Jin_2016_Instance_1","Paragraph":"We assume the time-scales we observe here are generated in the corona itself (and note that longer time-scale changes will be driven by the disc outside of the corona) and are made visible by a changing electron temperature and density as a result of local turbulence and coupling to mass accretion rate propagations through the flow from rout to rin. We note that we can discount variations in the seed photon population as the driver for changes in the power spectrum, as the UV emission from the disc is established to be considerably less variable than the corona in NLS1s (Leighly 1999; Smith & Vaughan 2007; Ai et al. 2013; Alston, Vaughan & Uttley 2013; Done & Jin 2016). We assume that the variability generated locally at each radius (r\u03bd) is at the viscous frequency (see Churazov et al. 2001) such that\n(9)$$\\begin{eqnarray}\r\nr_{\\nu } = \\left[ \\frac{2\\pi \\nu }{\\alpha }\\left(\\frac{H}{R}\\right)^{-2}\\right]^{-2\/3} ,\r\n\\end{eqnarray}$$where the frequency is in units of c\/Rg (see e.g. Kato, Fukue & Mineshige 1998; Ar\u00e9valo & Uttley 2006). In the above, \u03b1 and $\\frac{H}{R}$ are the dimensionless viscosity parameter of Shakura & Sunyaev (1973) and scale height of the accretion disc, respectively. We assume that our frequency range of interest, 0.01\u22121 mHz (i.e. the range over which we can practically fit to the data) corresponds to radii between the ISCO (rin) and some radius within the true outer edge of the corona (i.e. rout \u2264 rcorona). The actual frequencies generated in our model therefore depend on the SMBH spin and the combination $\\left(\\frac{H}{R}\\right)^2 \\alpha$ (and somewhat on the SMBH mass \u2013 although here the range is small). Given the reported high spin values for these bright AGNs (Ogle et al. 2004; Fabian et al. 2013; Done & Jin 2016; Kara et al. 2017; Buisson et al. 2018b), we expect the ISCO to sit at \u223c1.25Rg. For the corona at the ISCO to produce variability above our upper frequency limit of 1 mHz requires $\\left(\\frac{H}{R}\\right)^2 \\alpha \\gtrsim 0.01$. We note that for the mass range subtended by our AGN sample (from 106.00 to 106.63\u2009M\u2299, see Table 2), should we instead assume zero spin (rin = 6Rg), the viscous frequency at rin is lower and we have strong curvature in our observed bandpass.","Citation Text":["Done & Jin 2016"],"Functions Text":["We note that we can discount variations in the seed photon population as the driver for changes in the power spectrum, as the UV emission from the disc is established to be considerably less variable than the corona in NLS1s"],"Functions Label":["Uses"],"Citation Start End":[[661,676]],"Functions Start End":[[352,576]]}
{"Identifier":"2018MNRAS.478...95K__Toal\u00e1_et_al._2012_Instance_1","Paragraph":"We subsequently considered the time evolution of a number of selected cores (based on the requirement that we have sufficient data to follow at least 200 kyr of evolution), finding a remarkably similar chemical evolution in all cores, including one that is strongly stabilized by turbulence and magnetic field and therefore not going through gravitational collapse. Assuming a more or less chemically homogeneous initial condition, the chemical evolution of the cores, particularly regarding deuteration effects, is thus very similar. Of course, environmental effects such as the local cosmic ray ionization rate could induce potential differences, by changing the ionization degree. Similarly, metallicity effects at lower densities may depend on the local conditions. At least within one filament, it is however plausible that the deuteration fraction is indeed indicative of chemical age. We find here that about two free-fall times (as defined for cylindrical systems, see Toal\u00e1 et al. 2012) are sufficient to reach core deuteration fractions of \u22730.1. We finally investigated also the radial structure of the core, finding overall similar properties as in the isolated collapsing cores studied by K\u00f6rtgen et al. (2017). The H2 ortho-to-para ratio appears to be approximately flat and only weakly dependent on radius. Both the deuteration fraction and the gas surface density show a peak on scales of about 1000\u20132000 au, which is particularly pronounced in the case of parallel magnetic fields. As found previously, this peak moves outward with increasing turbulent Mach number, indicating the amount of support against gravity. The difference in the visibility of the peak may result from the difference in the fragmentation mode in both cases, as previously described by Seifried & Walch (2015), thus potentially affecting the structure of the resulting cores. Overall, our results have shown that the observed high deuteration fraction in prestellar cores can be readily reproduced in simulations of turbulent magnetized filaments. We further found that deuteration fractions of order 0.1 can be produced independent of the specific history of the cores, both for high and low virial parameters. The latter suggests that deuteration is potentially very efficient.","Citation Text":["Toal\u00e1 et al. 2012"],"Functions Text":["We find here that about two free-fall times (as defined for cylindrical systems, see","are sufficient to reach core deuteration fractions of \u22730.1."],"Functions Label":["Uses","Uses"],"Citation Start End":[[977,994]],"Functions Start End":[[892,976],[996,1055]]}
{"Identifier":"2021MNRAS.504.5074S__Maraschi,_Ghisellini_&_Celotti_1992_Instance_1","Paragraph":"The broad-band spectral energy distribution (SED) of blazars is characterized by two broad humps, one at optical\/UV\/X-ray bands and the other in the HE \u03b3-ray band (see Padovani et al. 2017 for a recent review). It is believed that the first peak (low-energy component) is mostly due to synchrotron emission from relativistic electrons, whereas the origin of the second component is highly debatable. Within conventional leptonic scenarios, this component is produced when the synchrotron-emitting electrons inverse Compton up scatter the photons of internal (synchrotron self-Compton, SSC; Ghisellini, Maraschi & Treves 1985; Maraschi, Ghisellini & Celotti 1992; Bloom & Marscher 1996) or external (external inverse Compton, EIC; Sikora, Begelman & Rees 1994; B\u0142a\u017cejowski et al. 2000; Ghisellini & Tavecchio 2009) origin. The nature of the external photon fields depends on the distance of the emitting region from the central black hole (Sikora et al. 2009) and can be dominated either by the photons directly emitted from the accretion disc (Dermer, Schlickeiser & Mastichiadis 1992; Dermer & Schlickeiser 1993) or disc photons reflected from the broad-line region (BLR; Sikora et al. 1994) or IR photons emitted from the dusty torus (B\u0142a\u017cejowski et al. 2000). Recently, after associating TXS 0506+056 with the IceCube-170922A neutrino event (IceCube Collaboration et al. 2018a,b; Padovani et al. 2018), it is more evident that the HE component could be initiated by the interaction of energetic protons when they are effectively accelerated in the blazar jets. The HE component can be either from proton synchrotron emission (M\u00fccke & Protheroe 2001) or from secondary particles from pion decay (Mannheim & Biermann 1989; Mannheim 1993; M\u00fccke & Protheroe 2001; M\u00fccke et al. 2003; B\u00f6ttcher et al. 2013). In the latter case, blazars are also sources of very high energy neutrinos (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).","Citation Text":["Maraschi, Ghisellini & Celotti 1992"],"Functions Text":["Within conventional leptonic scenarios, this component is produced when the synchrotron-emitting electrons inverse Compton up scatter the photons of internal (synchrotron self-Compton, SSC;"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[626,661]],"Functions Start End":[[400,589]]}
{"Identifier":"2015MNRAS.446.2468E__Kauffmann_&_Haehnelt_2000_Instance_1","Paragraph":"In the following, all measurements and maps derived from snapshots prior to t = 800\u2009Myr are therefore using the original simulation (R+13), with the others extracted from the simulation at 1 pc resolution (for the gaseous component). Note that star formation was turned on at t \u223c 745\u2009Myr in the simulation as to avoid gas being prematurely consumed. Details on the implemented recipes for star formation, stellar feedback (photoionization, radiative pressure, supernova explosions) are described in R+13. Adding on AGN feedback would significantly impact on the distribution, kinematics and physical status of the gas, specifically for the close environment of the black hole (Ciotti & Ostriker 1997; Haehnelt, Natarajan & Rees 1998; Silk & Rees 1998; Kauffmann & Haehnelt 2000). For the present simulations, however, we do not include the potential feedback from an AGN, thus focusing on a time window (t \u223c 750\u2013830\u2009Myr) when we consider that the AGN itself is quiet (or in an \u2018off-state\u2019). This is partly justified by the assumption that AGN have low duty cycle at low redshift and for black holes of a few 106 M\u2299 (Haehnelt & Rees 1993; Wang et al. 2009; Shankar et al. 2010; Shankar, Weinberg & Miralda-Escud\u00e9 2013) and by the short time range we are considering. More importantly, it allows us to narrow down our study to probe the interplay between the dynamical evolution and the effect of star formation (similarly to e.g. Levine et al. 2008; Hopkins & Quataert 2010a). Turning on the AGN in such a simulation would be paramount to understand any potential fuelling cycle starting from the large-scale down to the vicinity of the black hole, and such an implementation has already been included in ramses by a direct calculation of the Bondi accretion rate (see Teyssier et al. 2011; Gabor & Bournaud 2014, and references therein). It would nevertheless require to probe various feedback schemes, and triggering mechanisms, which is beyond the scope of the present paper.","Citation Text":["Kauffmann & Haehnelt 2000"],"Functions Text":["Adding on AGN feedback would significantly impact on the distribution, kinematics and physical status of the gas, specifically for the close environment of the black hole"],"Functions Label":["Future Work"],"Citation Start End":[[752,777]],"Functions Start End":[[505,675]]}
{"Identifier":"2018MNRAS.476L...6R__Cyr_et_al._2000_Instance_1","Paragraph":"The coronal mass ejections (CMEs) are frequent discharge of huge energy and massive magnetized plasma from the solar corona into the heliosphere. They are of paramount importance in space physics for their key role in extreme space weather and geo-effectiveness, e.g. (Gosling 1993; Low 2001; Schrijver & Siscoe 2010; Cannon et al. 2013). In last few decades, the understanding of CMEs improved significantly because of space and ground-based observational data with the help of various modelling efforts. The studies are focused on the morphological and kinematic evolution of CMEs in the heliosphere, e.g. (Lindsay et al. 1999; St Cyr et al. 2000; Zurbuchen & Richardson 2006; Chen 2011; Webb & Howard 2012; Wang et al. 2016; Lugaz et al. 2017). By considering the number of CMEs emitted from the Sun during solar maximum and variations in their respective speeds, the interaction between multiple CMEs in the heliosphere is expected to be more frequent. The collision of multiple CMEs highly affect their dynamic evolution properties and contribute to enhanced geo-effectiveness, e.g. (Wang, Wang & Ye 2002; Farrugia & Berdichevsky 2004; Lugaz, Manchester IV & Gombosi 2005; Wang et al. 2005; Xiong et al. 2007; Shen et al. 2011, 2012; Lugaz et al. 2012; Temmer et al. 2012). To predict space weather effects near the Earth, an accurate estimation of arrival time of CMEs at the Earth is crucial (Mishra & Srivastava 2014). Besides this, the study of CME\u2013CME and CME\u2013solar wind interactions provide unique observational evidences to understand energy dissipation of large-scale magnetic clouds in interstellar medium and authenticate the physical processes predicted theoretically. Therefore, interaction of multiple CMEs needs to be examined in detail. The various results obtained from studies have justified CME\u2013CME collision as an inelastic\/elastic collision or superelastic collision, e.g. (Lugaz et al. 2012; Shen et al. 2012, 2016; Lugaz et al. 2017; Mishra et al. 2017). The magnetohydrodynamics (MHD) numerical simulations have striven to understand the physical mechanism involved in CME\u2013CME interaction, CME\u2013CME driven shock interactions and their consequences, e.g. (Niembro et al. 2015; Jin et al. 2016; Shen et al. 2016; Wu et al. 2016).","Citation Text":["St Cyr et al. 2000"],"Functions Text":["In last few decades, the understanding of CMEs improved significantly because of space and ground-based observational data with the help of various modelling efforts. The studies are focused on the morphological and kinematic evolution of CMEs in the heliosphere, e.g."],"Functions Label":["Background"],"Citation Start End":[[630,648]],"Functions Start End":[[339,607]]}
{"Identifier":"2020MNRAS.499.4666F__Micha\u0142owski_2015_Instance_1","Paragraph":"An example of these implications is the so-called \u2018dust budget crisis\u2019 introduced in Section 4.4: the dust masses currently estimated at z > 5 are not compatible with standard dust production channels and require an overhaul in our models of the initial mass function for star formation, of supernova production rates, or of dust growth in the ISM. Overall, the dust production rate would need to increase by one to two orders of magnitudes, as shown by Rowlands et al. (2014). The growth of dust grains through accretion in the ISM has been proposed as a solution (e.g. Mancini et al. 2015; Micha\u0142owski 2015; Popping et al. 2017), but there are doubts on the efficiency of accretion at high z, where high dust temperatures due to the CMB (see Section 3.3) keep the desorption time-scale for accreted materials short (Ferrara et al. 2016). The dust budget crisis is not only a problem at high redshift; it is observed, e.g. in the Magellanic Clouds (SMC, LMC). As explained in Srinivasan et al. (2016) using the dust mass fits by Gordon et al. (2014), the dust replenishment time-scale in the SMC from stellar sources alone is expected to be larger than the dust destruction time-scale and, in the worst-case scenario, longer than the lifetime of the Universe. Similarly, the ratio between the best LMC dust mass estimate by Gordon et al. (2014) and the dust injection estimates by Riebel et al. (2012) results in an LMC replenishment time-scale of 34 \u00b1 8 Gyr, exceeding the age of the Universe. Both the high redshift and the local Universe, therefore, show a dust budget crisis that could be alleviated \u2013 and, in the best case scenario, fully resolved \u2013 if the actual dust masses turned out to be lower than currently estimated, as our results suggest. More specifically, Rowlands et al. (2014) mention that dust opacity needs to be increased by just a factor of 7 to solve the high-redshift crisis (provided dust destruction by SNe is not efficient); in the LMC, the aforementioned replenishment time-scale would decrease to less than 2 Gyr if the dust mass were decreased by a factor of 20.","Citation Text":["Micha\u0142owski 2015"],"Functions Text":["The growth of dust grains through accretion in the ISM has been proposed as a solution (e.g."],"Functions Label":["Motivation"],"Citation Start End":[[592,608]],"Functions Start End":[[478,570]]}
{"Identifier":"2022ApJ...934..103H__Moore_et_al._2001_Instance_1","Paragraph":"Compared to various studies based on in situ measurements of MFR structures after eruption, the origination of CME-MFRs before and during eruptions remains elusive due to the complex environment in the solar source region and limited observations. At the present time, there are certain hypotheses on the formation process of MFRs. Some studies indicate that MFRs could exist prior to the eruption. For example, both Fan (2001) and Magara (2004) reported findings from magnetohydrodynamic (MHD) simulation that a twisted MFR initially formed below the photosphere can partially emerge into the low corona by magnetic buoyancy. Other studies suggest that the presence of preeruptive MFRs is not necessary and MFRs could be built up in the corona via magnetic reconnection processes associated with flares (Amari et al. 2003; Moore et al. 2001; Antiochos et al.1999; Jiang et al. 2021a, 2021b). To understand the physical processes more precisely for flare\u2013CME events, extensions of the standard 2D flare model have been proposed to account for much broader ranges of quantitative measurements with three-dimensional (3D) features intrinsic to realistic solar eruptions (Longcope et al. 2007; Aulanier et al. 2012; Priest & Longcope 2017; Aulanier & Dud\u00edk 2019). For example, quasi-3D models have been developed with a nonvanishing magnetic field component along the axis of the MFR and to illustrate the scenario that sequential reconnection along the magnetic polarity inversion line (PIL) forms the MFR in the first place (van Ballegooijen & Martens 1989; Longcope et al. 2007; Schmieder et al. 2015). This scenario has been widely applied to infer and interpret magnetic reconnection properties based on the observed flare ribbon morphology (Qiu et al. 2002, 2004, 2010; Hu et al. 2014; Kazachenko et al. 2017; Zhu et al. 2020). From such analyses, Qiu et al. (2004) illustrated that there is a temporal correlation between the magnetic reconnection rate and the acceleration of the CME (considered as the eruptive MFR) in the low corona. Such a correlation has been further established by Zhu et al. (2020) based on a statistical study of \u223c60 events. In addition, Qiu et al. (2007) and Hu et al. (2014) showed a correlation between the magnetic reconnection flux and the flux contents of the corresponding ICME\/MC flux ropes based on modeling results employing in situ spacecraft measurements. These results support the hypothesis that CME-MFRs can be formed by magnetic reconnection during the corresponding flare process. Recent simulation results also indicate clearly that the reconnection flux contributes to the axial (toroidal) flux of the CME-MFR in the early stage (Jiang et al. 2021a; Inoue et al. 2018).","Citation Text":["Moore et al. 2001"],"Functions Text":["Other studies suggest that the presence of preeruptive MFRs is not necessary and MFRs could be built up in the corona via magnetic reconnection processes associated with flares"],"Functions Label":["Motivation"],"Citation Start End":[[824,841]],"Functions Start End":[[627,803]]}
{"Identifier":"2021MNRAS.508.4767S__Chiaki,_Yoshida_&_Hirano_2016_Instance_1","Paragraph":"Although it is still unknown why the disc fragmentation for the primordial cases is well described by the simple relation such as equation (1), an important fact is that a barotropic EOS with \u03b3eff \u2243 1.1 approximately represents the gas thermal evolution during the cloud collapse for $n \\lesssim 10^{19}\\, \\mathrm{cm}^{-3}$ (Omukai & Nishi 1998). In fact, several simulations study the disc fragmentation assuming the same barotropic EOS with \u03b3eff = 1.1 for n \u2264 nth, resulting in the evolution described by equation (1) (Susa 2019). Since nth is only the characteristic quantity for this case, the simple scaling of equation (1) may be convincing. This suggests that the disc fragmentation with a different EOS should provide different evolution of Nc,b. For instance, it is well known that adding a tiny amount of heavy elements and dust grains alters the EOS of a collapsing cloud (e.g. Omukai 2000; Bromm et al. 2001; Omukai et al. 2005; Schneider et al. 2006, 2012a; Smith, Sigurdsson & Abel 2008; Jappsen et al. 2009; Safranek-Shrader, Milosavljevi\u0107 & Bromm 2014; Chiaki et al. 2015; Chiaki, Yoshida & Hirano 2016). While previous studies demonstrate that the dust cooling enhances the fragmentation during the cloud collapse (e.g. Meece, Smith & O\u2019Shea 2014; Smith et al. 2015; Chiaki & Wise 2019), its effect on the disc fragmentation remains to be further explored. Tanaka & Omukai (2014) investigate the evolution of the circumstellar disc in metal-poor environments developing 1D semi-analytical models. They predict that the discs with $Z \\sim 10^{-5}\\!-\\! 10^{-3}\\, \\mathrm{Z}_{\\odot }$ are subject to the efficient dust cooling and are more unstable than those for the primordial cases. Machida & Nakamura (2015) consider the disc fragmentation with various metallicities $0 \\le Z \\le 1\\, \\mathrm{Z}_{\\odot }$, performing a suite of 3D numerical simulations. They find qualitative differences between the cases with $Z \\lesssim 10^{-4}\\, \\mathrm{Z}_{\\odot }$ and $Z \\gtrsim 10^{-4}\\, \\mathrm{Z}_{\\odot }$; the vigorous disc fragmentation only occurs for the former metal-poor cases. Whereas Machida & Nakamura (2015) use the metallicity-dependent barotropic EOS, Chiaki & Yoshida (2020) recently report 3D simulations of the disc fragmentation solving the energy equation with relevant thermal processes coupled with a non-equilibrium chemical network. They find that for the metal-poor cases with $Z \\le 10^{-3}\\, \\mathrm{Z}_{\\odot }$, the disc fragmentation does not necessarily prevent the mass growth of the most massive protostar as many clumps are short-lived owing to the frequent merger or tidal disruption events. Although the above studies suggest the metallicity-dependence of the disc fragmentation, they both only follow the short-term evolution for $\\sim 100\\, \\mathrm{yr}$ since the first emergence of a protostar.","Citation Text":["Chiaki, Yoshida & Hirano 2016"],"Functions Text":["For instance, it is well known that adding a tiny amount of heavy elements and dust grains alters the EOS of a collapsing cloud (e.g."],"Functions Label":["Background"],"Citation Start End":[[1089,1118]],"Functions Start End":[[755,888]]}
{"Identifier":"2020AandA...641A.118F__Nakajima_et_al._(2018b)_Instance_1","Paragraph":"Comparison with z \u2248 2\u22124galaxies. High-ionization UV lines have been detected in the spectra of z\u2004\u2273\u20042 galaxies through gravitational lensing (e.g., Stark et al. 2014; Patr\u00edcio et al. 2016; Vanzella et al. 2016, 2017; Berg et al. 2018), spectral stacking (e.g., Nakajima et al. 2018b; Rigby et al. 2018; Saxena et al. 2020), and deep spectroscopic observations (e.g., Erb et al. 2010; Maseda et al. 2017; Amor\u00edn et al. 2017; Nanayakkara et al. 2019). The EWs of He\u202fII, O\u202fIII]\u03bb1666, and C\u202fIII] from these works are larger than those measured in the average spectra of our LAEs, with the exception of some of the MUSE sources at z\u2004\u2273\u20043 studied by Patr\u00edcio et al. (2016) and Nanayakkara et al. (2019), some of the stacks from Nakajima et al. (2018b), and some of the VANDELS He\u202fII emitters from Saxena et al. (2020). The EWs from Nakajima et al. (2018b) are computed from average spectra of a z\u2004\u2248\u20043 LAE population whose median UV luminosity is about two orders of magnitude brighter than ours. Most of the 2.4\u2004 \u2004z\u2004 \u20043.5 UV-selected, low-luminosity galaxies from Amor\u00edn et al. (2017) have higher He\u202fII and C\u202fIII] EWs, most likely because of their higher SFR (see Fig. 2). The line measurements from Saxena et al. (2020) are performed on the stacked spectra of UV continuum, bright He\u202fII emitters (\u221219\u2004 \u2004MUV\u2004 \u2004\u221222). The strong line emitters from Erb et al. (2010) and Berg et al. (2018) are brighter and more massive than the median value of our LAEs. The lensed galaxies from Vanzella et al. (2016, 2017) are among the least massive, most metal-poor, young, and faintest systems observed at z\u2004\u223c\u20043. With MUV\u2004> \u2004\u221216, the source ID14 from Vanzella et al. (2017) is roughly one order of magnitude fainter that the faintest LAE in our sample, while the source ID11 from Vanzella et al. (2016) has an remarkable blue UV slope (\u03b2\u2004=\u2004\u22122.95). Recently, Du et al. (2020) measured C\u202fIII] EWs in z\u2004\u223c\u20042 analogs of galaxies in the reionization era, obtaining values from 13.2 \u00c5 down to 1 \u00c5, reaching values as low as those of our LAEs. The authors also found differences in the C\u202fIII] EW depending on the selection criteria, with higher values of C\u202fIII] EW for emission lines rather than for continuum-selected galaxies. Mainali et al. (2020) detected two targets with C\u202fIII] emission reaching EW \u2248 17 \u2212 21 \u00c5 in the spectra of z\u2004\u223c\u20042 galaxies selected for their strong rest-optical line ([O\u202fIII]\u03bb5007+H\u03b2) EWs. These values are above the ones of our LAEs and are similar to those observed at z\u2004> \u20046 (e.g., Stark et al. 2015b, 2017; Hutchison et al. 2019).","Citation Text":["Nakajima et al. 2018b","Nakajima et al. (2018b)","Nakajima et al. (2018b)"],"Functions Text":["High-ionization UV lines have been detected in the spectra of z\u2004\u2273\u20042 galaxies through","spectral stacking","The EWs of He\u202fII, O\u202fIII]\u03bb1666, and C\u202fIII] from these works are larger than those measured in the average spectra of our LAEs, with the exception of","some of the stacks from","The EWs from","are computed from average spectra of a z\u2004\u2248\u20043 LAE population whose median UV luminosity is about two orders of magnitude brighter than ours."],"Functions Label":["Background","Background","Similarities","Similarities","Differences","Differences"],"Citation Start End":[[260,281],[720,743],[824,847]],"Functions Start End":[[33,117],[235,252],[449,596],[696,719],[811,823],[848,987]]}
{"Identifier":"2021ApJ...908...45B__Eyink_&_Sreenivasan_2006_Instance_1","Paragraph":"Here we shall set the potential \n\n\n\n\n\n and make the change of variable \n\n\n\n\n\n; hence, the vortical parts of the footpoint equations of motion are cast into a Hamiltonian form. The Hamiltonian describing the dynamics of magnetic footpoints driven by the vortical component of the turbulence on the photosphere is\n50\n\n\n\n\n\nThe footpoint Hamiltonian \n\n\n\n\n\n has the physical dimension of a frequency (s\u22121). From Figure 2 in Rincon et al. (2017), we further notice that the vortical component of the photospheric surface velocity field has a structure reminiscent of 2D Euler turbulence, which admits a vortex point representation (Eyink & Sreenivasan 2006) also on the sphere (Pavlov et al. 2001). There are three important parameters that can be extracted from the vortical component of photospheric turbulence. These parameters are the typical amplitude of the velocity fluctuations \n\n\n\n\n\n, the correlation length \n\n\n\n\n\n, and the correlation time \n\n\n\n\n\n of the turbulence, where \n\n\n\n\n\n is the rms amplitude of the fluctuating stream function \u03a6 and \n\n\n\n\n\n is the angular size of the turbulent eddies on the source surface. These three independent parameters \n\n\n\n\n\n, \n\n\n\n\n\n, and \n\n\n\n\n\n can all be derived from a more fundamental quantity, which is the covariance function \n\n\n\n\n\n, i.e., the autocorrelation function of the turbulent stream function (Creasey & Lang 2018). The covariance function is sufficient to characterize the vortical component of the surface flows provided that the stream function \u03a6 does not significantly deviate from a stationary isotropic Gaussian random field on the sphere. Hamilton equations with the Hamiltonian (50) determine a genuine problem of chaotic mixing (Ott 1993; Morrison 1998) on the sphere. We emphasize that mixing of footpoints on the photosphere relies on the time dependence of the Hamiltonian. In fact, the footpoint dynamics is regular (quasi-periodic) when the Hamiltonian is a constant of time, independently of the number of modes involved in its decomposition. The reason is that the dynamics generated by 1 degree of freedom Hamiltonians, here for the pair of canonically conjugate variables \u03bc and \u03d5, is regular. A simple mathematical proof of this statement is given by Morrison (1998). In one version of the protocol used by Giacalone & Jokipii (2004) to model the magnetic footpoint dynamics, the time dependence of the flow results from the (random) phase of the (complex) toroidal coefficients being renewed at random times, following the work of Simon et al. (1995).","Citation Text":["Eyink & Sreenivasan 2006"],"Functions Text":["From Figure 2 in Rincon et al. (2017), we further notice that the vortical component of the photospheric surface velocity field has a structure reminiscent of 2D Euler turbulence, which admits a vortex point representation","also on the sphere"],"Functions Label":["Uses","Uses"],"Citation Start End":[[626,650]],"Functions Start End":[[402,624],[652,670]]}
{"Identifier":"2016AandA...592L..11S__Vasyunin_&_Herbst_(2013)_Instance_1","Paragraph":"Methanol is believed to be formed on dust grains (Watanabe & Kouchi 2002) by subsequent hydrogenation of carbon monoxide, and its detection towards prestellar cores is already a challenge for current models given the absence of efficient desorption processes in these sources. Thermal desorption is out of question because of the low dust temperature. Recent laboratory studies showed that the photo-desorption of methanol from ices is likewise negligible (Bertin et al. 2016; Cruz-Diaz et al. 2016). The main desorption products when irradiating pure and mixed methanol ices are photo-fragments of methanol. An alternative route to explain the presence of methanol in the gas phase is the reactive\/chemical desorption that has been theoretically proposed by Garrod et al. (2007) and Vasyunin & Herbst (2013) and experimentally studied by Dulieu et al. (2013) and Minissale et al. (2016). On the other hand, c-C3H2 is mainly formed in the gas phase (e.g. Spezzano et al. 2013) and is expected to preferentially trace dense and chemically young gas, that is, gas where C atoms have not yet been mainly locked into CO. This gas rich in C atoms is expected in the outer envelope of an externally illuminated dense core (e.g. Aikawa et al. 2001). However, towards L1544, c-C3H2 only appears to trace one side of the core, the one closer to the sharp N(H2) edge and away from the CH3OH peak. This indicates that photo-chemistry is not uniformly active around L1544, most likely because the distribution of the envelope material (belonging to the filament within which L1544 is embedded) is not uniform, as is clearly shown by the Herschel map in Fig. 1. This figure shows that methanol traces a region farther away from the southern sharp edge of the N(H2) map, which is possibly more shielded from the ISRF and where most of the carbon is locked in CO. CH3OH is preferentially found at the northern edge of L1544 because here photochemistry does not play a major role (so C is locked in CO) and the density is low enough to maintain a higher fraction of CH3OH in the gas phase, but above the threshold value for CO freeze-out, a few \u00d7104 cm-3 (Caselli et al. 1999). Based on the Keto & Caselli (2010) model (that was updated by Keto et al. 2014), the volume density at the distance of the methanol peak is predicted to be 8 \u00d7 104 cm-3, which is just above the threshold value. In contrast, cyclopropenylidene has the most prominent peak towards the southern sharp edge of the H2 column density and extends along the semi-major axis of the core, almost parallel to the south-west edge of the N(H2) map. This behaviour is also clearly shown in Fig. 2, where the integrated intensities of both methanol and cyclopropenylidene are plotted against N(H2). c-C3H2 is also present at values of N(H2) that are lower than those of methanol, and it maintains a flat intensity profile, suggestive of a layer-like structure, with no significant increase towards the core centre. The CH3OH intensity instead shows a sharp rise up to column densities of about 1.6 \u00d7 1022 cm-2, and it declines at higher values. ","Citation Text":["Vasyunin & Herbst (2013)"],"Functions Text":["An alternative route to explain the presence of methanol in the gas phase is the reactive\/chemical desorption that has been theoretically proposed by Garrod et al. (2007) and","and experimentally studied by Dulieu et al. (2013) and Minissale et al. (2016).","On the other hand, c-C3H2 is mainly formed in the gas phase","and is expected to preferentially trace dense and chemically young gas, that is, gas where C atoms have not yet been mainly locked into CO.","This gas rich in C atoms is expected in the outer envelope of an externally illuminated dense core","However, towards L1544, c-C3H2 only appears to trace one side of the core, the one closer to the sharp N(H2) edge and away from the CH3OH peak. This indicates that photo-chemistry is not uniformly active around L1544, most likely because the distribution of the envelope material (belonging to the filament within which L1544 is embedded) is not uniform, as is clearly shown by the Herschel map in Fig. 1."],"Functions Label":["Background","Background","Background","Background","Background","Differences"],"Citation Start End":[[784,808]],"Functions Start End":[[609,783],[809,888],[889,948],[977,1116],[1117,1215],[1243,1648]]}
{"Identifier":"2016MNRAS.456..512C__Kronberg_et_al._2004_Instance_3","Paragraph":"Extended radio emission in galaxies is associated with both radio jets and lobes and with outflows, seen often as aligned radio sources in the opposite directions with respect to the central compact radio core. Giant radio galaxies (GRG) are extreme cases of this phenomenology with jets and lobes extending on \u223c Mpc scales suggesting that they are either very powerful or very old site for electron acceleration. In this respect, GRGs have a crucial role in the acceleration of cosmic rays over large cosmic scales (e.g. Kronberg et al. 2004), in the feedback mechanism of AGNs into the intergalactic and intracluster medium (e.g. Subrahmanyan et al. 2008) and in the seeding of large-scale magnetic fields in the universe (e.g. Kronberg et al. 2004) and they are excellent sites to determine the total jet\/lobe energetics in AGN-dominated structures (see e.g. Colafrancesco 2008, Colafrancesco & Marchegiani 2011). To date our knowledge of GRGs (see e.g. Ishwara-Chandra & Saikia 1999, 2002; Lara et al. 2001; Machalski, Jamrozy & Zola 2001; Schoenmakers et al. 2001; Kronberg et al. 2004; Saripalli et al. 2005; Malarecki et al. 2013; Butenko et al. 2014) is limited by their sparse numbers and by the difficulty of detecting them over large areas of the sky. Low-frequency radio observations have an enhanced capacity to detect the extended old electron population in these objects (see e.g. the recent Low Frequency Array \u2013 LOFAR \u2013 observation of the GRG UGC095551), but high-frequency radio observations are less efficient in this task due to the steep-spectra of giant radio lobes. In this context these sources will be ideal targets for the next coming deep, wide-field surveys like, e.g. the ATLAS survey of the Australia Telescope Network Facility (ATNF; see Norris et al. 2009) or the Square Kilometre Array (SKA) deep surveys that will have the potential to study their population evolution up to high redshifts and thus clarifying their role on the feedback for the evolution of non-thermal processes in large-scale structures.","Citation Text":["Kronberg et al. 2004"],"Functions Text":["To date our knowledge of GRGs (see e.g.","is limited by their sparse numbers and by the difficulty of detecting them over large areas of the sky. Low-frequency radio observations have an enhanced capacity to detect the extended old electron population in these objects (see e.g. the recent Low Frequency Array \u2013 LOFAR \u2013 observation of the GRG UGC095551), but high-frequency radio observations are less efficient in this task due to the steep-spectra of giant radio lobes.","In this context these sources will be ideal targets for the next coming deep, wide-field surveys like, e.g. the ATLAS survey of the Australia Telescope Network Facility (ATNF","or the Square Kilometre Array (SKA) deep surveys that will have the potential to study their population evolution up to high redshifts and thus clarifying their role on the feedback for the evolution of non-thermal processes in large-scale structures."],"Functions Label":["Background","Background","Future Work","Future Work"],"Citation Start End":[[1070,1090]],"Functions Start End":[[917,956],[1159,1588],[1589,1763],[1789,2040]]}
{"Identifier":"2015MNRAS.448.1847H__Kim_et_al._2003_Instance_1","Paragraph":"We assume that gas and dust are well mixed along each line of sight in which case the gas surface density (\u03a3gas) is proportional to the dust surface density (\u03a3dust) and the proportionality factor is the gas-to-dust mass ratio (rgd). The dust masses derived by SED fitting methods vary significantly and systematically (e.g. Galliano et al. 2011; Gordon et al. 2014) depending on the applied method and assumed dust optical properties. It is therefore important to use a value of rgd consistent with the used SED fitting methodology, i.e. the value of rgd that gives the correct gas mass given the derived dust mass. We have calibrated the rgd using the available gas tracers on a large scale. Within a 500 arcsec radius centred on NGC 346, we find an atomic gas mass of 2.7\u00d7106 M\u2299 using the Parkes+ATCA H\u2009i map (Kim et al. 2003) and a spin temperature of 60 K (Bernard et al. 2008). The molecular gas mass determined from CO(J = 1\u20130) within the same aperture is 7.4\u00d7104 M\u2299 when using a CO-to-H2 column density conversion factor (XCO) of 1021 [cm\u22122 (K\u2009km\u2009s\u22121)\u22121] (Bolatto, Wolfire & Leroy 2013). If we use an XCO of 1022 [cm\u22122 (K\u2009km\u2009s\u22121)\u22121], we find a molecular gas mass of 7.5\u00d7105 M\u2299. Such an elevated XCO may be more appropriate given the harsh radiation field that is prevalent in the region, which may induce enhanced photodissociation of the CO molecules. Thus, we derive a total gas mass in the range of 2.7\u00d7106\u20133.4\u00d7106 M\u2299. The dust mass derived using the dust continuum photometry in this aperture and applying the same SED fitting routine we use for the pixel-by-pixel analysis yields 2.7\u00d7103 M\u2299. The derived rgd ranges from \u223c1000 to \u223c1250. In the following, we use a fiducial value for rgd of 1250, which is much higher than the Galactic value (100; Draine & Li 2007) due to the lower metallicity of the SMC. Our value for rgd of 1250 is lower than the total SMC integrated value of 1740 from Gordon et al. (2014). The global value includes a significant amount of atomic gas at low column density in the outskirts of the SMC which is not detected in dust emission. This drives up the global rgd.","Citation Text":["Kim et al. 2003"],"Functions Text":["Within a 500 arcsec radius centred on NGC 346, we find an atomic gas mass of 2.7\u00d7106 M\u2299 using the Parkes+ATCA H\u2009i map"],"Functions Label":["Uses"],"Citation Start End":[[812,827]],"Functions Start End":[[693,810]]}
{"Identifier":"2021MNRAS.503.5179N__Blanton_et_al._2004_Instance_2","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"],"Functions Text":["Analysis of Chandra observations revealed a hole or bubble within the ICM, cospatial with the eastern lobe of the jet"],"Functions Label":["Background"],"Citation Start End":[[1235,1254]],"Functions Start End":[[1116,1233]]}
{"Identifier":"2020AandA...633A.163C__Aalto_et_al._2015_Instance_1","Paragraph":"By using the RADEX2 dense cloud models developed by Aalto et al. (2015) to reproduce the HCN(3\u20132)\/(1\u20130) line luminosity ratios in the outflow of Mrk 231, we can attempt to find a combination of XHCN, XCN, Tkin, and nH2 solutions that can also fit the CN\/HCN and CN spin doublet line ratios (Table 2). We assume that the HCN and CN line emissions arise from the same dense cloud population, while the low-J CO line emission is due to a different, more diffuse phase of the outflow. We recall that in these models (see also Aalto et al. 2015), the dense clouds can be either self-gravitating virialised clouds, which implies that their internal velocity dispersion (\u0394vsg) is locked to their mass (Mvir) and size (R) through \u0394vsg\u2004=\u2004(GMvir\/G)1\/2, or unbound clouds, for which \u0394v\u2004\u226b\u2004\u0394vsg. We explored CN and HCN abundances in the range between 10\u22128 and 10\u22126. We find that depending on whether the clouds are self-gravitating or unbound, the models produce very different values for the absolute CN and HCN abundances, hence XCN and XHCN remain quantitatively unconstrained for the outflow with current data. However, all possible solutions that fit the observed line ratios consistently require XCN\u2004> \u2004XHCN, with a CN abundance that is at least a factor of three higher than the HCN abundance. Gas densities for this outflow phase (traced by the CN and HCN emissions) are nH2\u2004\u223c\u2004105\u2005\u2212\u2005106 cm\u22123, with temperatures not much higher than Tkin\u2004\u223c\u200420 K. Because CN is a well-known PDR tracer (see also Sect. 1), these results strongly suggest that the whole dense cloud population in outflow is affected by UV radiation. We should mention that high CN abundances may also be due to cosmic rays (e.g. see work done on the Galactic centre by Harada et al. 2015), which are known to permeate the outflow of Mrk 231, as inferred by Gonz\u00e1lez-Alfonso et al. (2018) based on the high OH+ abundance. However, it is not clear whether a cosmic-ray chemistry would also explain XCN\u2004> \u2004XHCN.","Citation Text":["Aalto et al. (2015)"],"Functions Text":["By using the RADEX2 dense cloud models developed by","to reproduce the HCN(3\u20132)\/(1\u20130) line luminosity ratios in the outflow of Mrk 231, we can attempt to find a combination of XHCN, XCN, Tkin, and nH2 solutions that can also fit the CN\/HCN and CN spin doublet line ratios (Table 2)"],"Functions Label":["Uses","Uses"],"Citation Start End":[[52,71]],"Functions Start End":[[0,51],[72,299]]}
{"Identifier":"2022MNRAS.511.1121M__Reig_&_Nespoli_2013_Instance_4","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"],"Functions Text":["The HB pattern is generally observed in the subcritical regime and the DB pattern is observed in the supercritical regime"],"Functions Label":["Background"],"Citation Start End":[[2008,2027]],"Functions Start End":[[1885,2006]]}
{"Identifier":"2018ApJ...864...76Z__Plunkett_et_al._2013_Instance_1","Paragraph":"C18O (J = 2\u20131) observations revealed a rotational structure in the north\u2013south direction at the center of IRAS 4C, but whether the rotation is Keplerian is not known due to the low signal-to-noise ratio (S\/N) (Tobin et al. 2015). Spitzer IRAC observations show an outflow cavity structure highlighted by scattered light and shocked emission to the east side of the central source, with the west side being much fainter (Figure 19 of Tobin et al. 2015). The east side is therefore inferred to be the blueshifted side, because the blueshifted outflow cavity tends to be brighter than the redshifted cavity in NIR and MIR as it is less extincted. Despite the outflow cavity structure seen in infrared, there was no clear evidence of a molecular outflow from mm or submm observations. Previous 12CO observations either reported no detection of outflow emission toward this source (Plunkett et al. 2013; Tobin et al. 2015) or only weak compact blueshifted emission towards east of the continuum source (Stephens et al. 2018). One possible explanation for the weak 12CO outflow emission is that the outflow lies close to the plane of sky, so that the low-velocity outflow emission is easily mixed with the emission of the ambient gas, especially for abundant species like 12CO. Indeed, the inclination of the source is estimated to be nearly edge-on (Tobin et al. 2015) to about 25\u00b0 between the disk plane and the line of sight (Segura-Cox et al. 2016). On the other hand, the 13CO (J = 2\u20131) emission reveal a compact (\u22722\u2033) structure with the blueshifted emission slightly offset to the east of the redshifted emission, which was explained as a slow outflow (Koumpia et al. 2016). Here, we report that the outflow cavity structure is clearly detected in the CCH and CS emissions, with kinematics consistent with rotation with respect to the outflow axis. This allows us to measure the angular momentum distribution in the outflow, and further constrain its launching radii and launching mechanism.","Citation Text":["Plunkett et al. 2013"],"Functions Text":["Previous 12CO observations either reported no detection of outflow emission toward this source","or only weak compact blueshifted emission towards east of the continuum source"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[877,897]],"Functions Start End":[[781,875],[918,996]]}
{"Identifier":"2020AandA...633A..70P__Kocevski_et_al._2011_Instance_1","Paragraph":"Recent studies have used the spectral indices [OII], H\u03b4, and Dn4000 to probe the stellar population of galaxies at intermediate redshifts (0.5\u2004\u2272\u2004z\u2004\u2272\u20041.2) because they are available in the observed optical frame. All these indicators, when combined, can be used to distinguish actively star-forming, (post-) starburst, and old or passive galaxies because they are expected to occupy different regions of the possible parameter space (e.g. Couch & Sharples 1987; Balogh et al. 1999; Poggianti et al. 1999, 2009; Fritz et al. 2014). The [OII]\u03bb3737 emission traces on-going star formation (timescales of \u223c10 Myr, e.g. Couch & Sharples 1987; Poggianti et al. 1999, 2006; Kennicutt 1998; Kewley et al. 2004), but it also depends on the metallicity and can be a poor tracer for dusty galaxies (e.g. Kewley et al. 2004; Yan et al. 2006; Kocevski et al. 2011). By measuring the [OII] equivalent width (EW), we can also crudely trace the specific SFR (sSFR), which is found to be anti-correlated with stellar mass (e.g. Bridge et al. 2015; Cava et al. 2015; Darvish et al. 2015a), with more massive star-forming galaxies having lower [OII] EWs. Additionally, higher density environments are found to depress [OII] emission (e.g. Balogh et al. 1999; Darvish et al. 2015a). The H\u03b4 line (and other strong Balmer absorption lines) can be indicative of a post-starburst phase (\u2248100\u2005\u2212\u20051000 Myr after the burst, e.g. Couch & Sharples 1987; Balogh et al. 1999; Poggianti et al. 1999, 2009; Dressler et al. 2004; Vergani et al. 2010; Mansheim et al. 2017a) if a strong absorption (typical of A stars, where hydrogen absorption is strongest) is observed and no tracers of on-going star formation are found (Couch & Sharples 1987). Recently, Wu et al. (2018) found that the H\u03b4 EW correlates with stellar mass, with more massive galaxies having weaker H\u03b4 absorption lines, but they did not study the effect of the environment (see also e.g. Siudek et al. 2017, for a similar result on passive galaxies). Finally, a measure of the flux break at 4000 \u00c5 (D4000 and Dn4000, as defined by Bruzual 1983; Balogh et al. 1999, respectively) traces the age of the galaxy and also the stellar metallicity (especially for older systems) to a lesser degree. This break is produced by a combination of metal absorption on the atmosphere of old and cool stars and the lack of flux from young and hot OB stars (e.g. Poggianti & Barbaro 1997; Kauffmann et al. 2003), and so it is sensitive to the average age of the stellar population. The 4000 \u00c5 break is also found to be stronger for galaxies with higher stellar mass (e.g. Muzzin et al. 2012; Vergani et al. 2008; Hern\u00e1n-Caballero et al. 2013; Siudek et al. 2017; Wu et al. 2018), which indicates that their stellar populations might be older, in an average sense. In terms of local density, Muzzin et al. (2012) found that galaxies in cluster environments have stronger breaks on average than their field counterparts at similar stellar masses, which the authors argued can be explained by the different fractions of star-forming and quiescent galaxies in different environments.","Citation Text":["Kocevski et al. 2011"],"Functions Text":["The [OII]\u03bb3737 emission traces on-going star formation","but it also depends on the metallicity and can be a poor tracer for dusty galaxies (e.g."],"Functions Label":["Background","Background"],"Citation Start End":[[829,849]],"Functions Start End":[[530,584],[703,791]]}
{"Identifier":"2020AandA...638A.113C__Juneau_et_al._2013_Instance_1","Paragraph":"To date, large surveys focused on detecting AGN systems have been conducted at a range of different wavelengths and, in particular, for the IR and X-ray domains (Alexander et al. 2005, 2008; Ivison et al. 2004; Lutz et al. 2005; Men\u00e9ndez-Delmestre et al. 2007, 2009; Valiante et al. 2007; Pope et al. 2008; Bonzini et al. 2013; Smol\u010di\u0107 et al. 2017; Wang et al. 2013; Stach et al. 2019). Due to the dust extinction in the near-IR, optical, and UV, as well as gas absorption in X-ray bands, these surveys are often incomplete. In the mid- and far-IR, incompleteness of AGN surveys may arise from the fact that not all AGN have significant IR emission from a dusty torus and therefore may not be detected. It is postulated that up to a third of AGN are undetected in these surveys (Mateos et al. 2017). Moreover, other studies have compared AGN selected from various wavebands and find that their host galaxies tend to have different properties in terms of colour (Hickox et al. 2009) and star-formation rates (SFR; Juneau et al. 2013; Ellison et al. 2016). In particular, Hickox et al. (2009) illustrated that there is only very little overlap between their 122 radio-selected AGN and those selected by X-ray or IR. Therefore, dust-free radio surveys are needed to provided a more complete census of the AGN population. Traditional radio surveys are only sensitive to radio-loud (RL) AGN, which only represent a tiny fraction (10 \u223c 20%) of the whole AGN population; however, modern radio surveys can achieve a flux depth where radio-quiet AGN can be detected (see Prandoni et al. 2018 for a review). Recent work has focused on the radio as it is sensitive to AGN and star formation concordantly, thus providing a method of surveying AGN and star-formation activity across cosmic time (e.g. Smol\u010di\u0107 et al. 2017; Padovani et al. 2015). A lot of work has also been done to look for AGN-driven radio emission, which has been identified by an excess of radio emission compared to what is expected based on the radio-FIR correlation, holding for star-forming galaxies (e.g. Ivison et al. 2010; Condon et al. 2002; Thomson et al. 2014; Magnelli et al. 2015).","Citation Text":["Juneau et al. 2013"],"Functions Text":["Moreover, other studies have compared AGN selected from various wavebands and find that their host galaxies tend to have different properties","and star-formation rates (SFR;"],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[1013,1031]],"Functions Start End":[[800,941],[982,1012]]}
{"Identifier":"2016ApJ...819...59S__Castelli_&_Kurucz_2004_Instance_1","Paragraph":"HD 47366 (HIP 31674, BD\u221212 1566, HR 2437, TYC 5373-2001-1) is listed in the Hipparcos Catalogue (ESA 1997) as a K1 III star, with a visual magnitude of V = 6.11 and a color index of B \u2212 V = 0.994. The Hipparcos parallax \u03c0 = 12.5 \u00b1 0.42 mas (van Leeuwen 2007) corresponds to a distance of 80.0 \u00b1 2.7 pc. The reddening E(B \u2212 V) = 0.028 was obtained from the Galactic dust map of Schlegel et al. (1998), with the correction given by Bonifacio et al. (2000) and a scaling factor of \n\n\n\n\n\n, where d is the distance, b is the Galactic latitude, and h = 125 pc is the scale height of the reddening layer. The absolute magnitude MV = 1.51 was derived from the distance and the interstellar extinction AV = 3.1E(B \u2212 V). By adopting the broadband photometric color B \u2212 V and the estimated metallicity with the empirical calibration relation of Alonso et al. (1999, 2001), we derived the bolometric correction B.C. = \u22120.309 and the effective temperature Teff = 4866 \u00b1 100 K. We used a high signal-to-noise ratio (S\/N \u223c 200), iodine-free spectrum taken with HRS to measure the equivalent widths (EWs) of \u223c30 Fe i lines, to derive the iron abundance [Fe\/H]. The line lists as well as their oscillation strengths (\n\n\n\n\n\n) were mainly taken from Hekker & Mel\u00e9ndez (2007), in which iron lines were carefully selected to avoid any blend by atomic or CN lines. The model atmosphere used in this work was interpolated from the line-blanketed, local thermodynamic equilibrium (LTE) ATLAS9-ODFNEW grid (Castelli & Kurucz 2004). The microturbulent velocity vt was obtained by minimizing the trend between the abundances of different Fe i lines and their reduced equivalent widths (\n\n\n\n\n\n). The macroturbulent velocity was estimated with the empirical relations of vmacro versus Teff given by Hekker & Mel\u00e9ndez (2007), and the projected rotational velocity (\n\n\n\n\n\n) was determined with the method of Fekel (1997). The stellar mass, surface gravity (\n\n\n\n\n\n), radius, and age were derived using a Bayesian approach with the Geneva database (Lejeune & Schaerer 2001), which includes the post-helium flash phases for stars with M \u2265 1.7 M\u2299. First, we interpolated an extensive grid of evolutionary tracks, with \u0394M = 0.05 within 1.2 \u2264 M\/M\u2299 \u2264 3.6, \u0394[Fe\/H] = 0.02 within \u22120.4 \u2264 [Fe\/H] \u2264 +0.3, and 500 points in each track. Then, for each data point, we calculated the likelihood functions of \n\n\n\n\n\n, Teff and [Fe\/H] to match the observed values by assuming Gaussian errors. We adopted uniform prior probabilities of mass and [Fe\/H]. It is noted that stars evolving more slowly have a higher probability of being observed. Without correcting this evolution effect, the resulting parameters would bias toward the rapid evolution phases. We therefore weighted the probability of each point along its evolutionary track by the normalized time-step (ai+1,j \u2212 ai,j)\/(an,j \u2212 aa,j), where ai,j is the age of the ith interpolated point in the jth track, and n = 500 is the number of interpolated points in each track. Eventually, the probability distribution functions (PDFs) of the parameters yield M = 1.81 \u00b1 0.13 M\u2299, R = 7.30 \u00b1 0.33 R\u2299, \n\n\n\n\n\n, and age = 1.61 \u00b1 0.53 Gyr. The stellar mass is particularly important to derive the minimum masses of the orbiting planets detected with the Doppler technique. However, previous spectroscopic analyses gave discrepant results (1.87 M\u2299 by Liu et al. 2010; 1.2 M\u2299 by Mishenina et al. 2006) for HD 47366, which may be due to the different methods on finding Teff and \n\n\n\n\n\n. Our determinations were based on a similar method to that of Liu et al. (2010), but used the Geneva evolutionary tracks, instead of the Y2 model (Yi et al. 2003), which does not account for the evolutionary phases after the helium-core flash. We found that the probability that the star has passed through the RGB tip and in-core helium burning phase is \u223c88%. The stellar parameters of HD 47366 are listed in Table 1. In Figure 2, we plotted HD 47366 on the H-R diagram, together with the evolutionary tracks for stars with different masses and metal contents.","Citation Text":["Castelli & Kurucz 2004"],"Functions Text":["The model atmosphere used in this work was interpolated from the line-blanketed, local thermodynamic equilibrium (LTE) ATLAS9-ODFNEW grid"],"Functions Label":["Uses"],"Citation Start End":[[1482,1504]],"Functions Start End":[[1343,1480]]}
{"Identifier":"2020MNRAS.491.3860S__Schlickeiser_2003_Instance_1","Paragraph":"The main assumption is that the statistical properties of the interaction of the charged particles with the fields are dominated by Gaussian distributions. This is in correspondence with the Central Limit Theorem (CLT), which requires that all stochastic systems evolve asymptotically towards Gaussian statistics provided that (i) many interactions are involved, (ii) the change in state in individual interactions is always small, and (iii) subsequent changes of state are statistically independent of each other. Besides these necessary conditions for the applicability of the FP equation, also other simplifying assumption for a better tractability of the FP equation are being made, such as (1) the magnetic fluctuations are homogeneous in space, (2) the electromagnetic fields are quasi-static, (3) the interaction has a finite decorrelation time, etc. (see more details in the books Schlickeiser 2003; Zank 2014). Unfortunately, in astrophysical and laboratory plasmas, most of the above assumptions are not valid, yet the FP equation is used extensively, without a proof of its validity. This is especially true when the plasma particles are accelerated to high energies impulsively (e.g. in solar flares, coronal mass ejections, or the Earth\u2019s magnetotail). The acceleration volume is finite and the expected fluctuating electromagnetic fields are strong ($\\vert \\delta \\boldsymbol{B} \\vert \\ge \\vert \\boldsymbol{B}_0 \\vert$). In solar active regions, the complex magnetic topologies host many null magnetic points which are randomly distributed inside the erupting or flaring volume (Aulanier et al. 2000; Pontin 2011). In these cases, the interaction of the particles with the strong em disturbances is transient and has no time to lead to Gaussian statistics or to become homogeneous in space (Isliker, Archontis & Vlahos 2019). Before analysing the interaction of the particles with the em fluctuations, it is important to understand the evolution of the em waves. With the use of resistive MHD codes one can show that a spectrum of high amplitude electromagnetic fluctuations evolves rapidly and leads to a fragmented current system, where reconnecting current sheets and large amplitude magnetic fluctuations are present (Arzner & Vlahos 2004; Dmitruk, Matthaeus & Seenu 2004; Vlahos, Isliker & Lepreti 2004; Isliker, Vlahos & Constantinescu 2017a), and where the various statistics clearly are non-Gaussian, following largely the paradigm of the stable Levy distributions (Isliker et al. 2017a; Isliker et al. 2019). In strongly turbulent plasmas, the magnetic fluctuations are non-collective modes and cannot be described with a simple dispersion relation $\\omega =\\omega (\\boldsymbol{k}).$ The em environment generated from the evolution of large amplitude em fluctuations is well documented in the current literature and models much better many impulsive astrophysical and laboratory plasmas (see Dmitruk et al. 2004; Zhdankin et al. 2013; Isliker et al. 2017a and the reviews by Cargill et al. 2012; Karimabadi et al. 2014; Vlahos & Isliker 2019).","Citation Text":["Schlickeiser 2003"],"Functions Text":["Besides these necessary conditions for the applicability of the FP equation, also other simplifying assumption for a better tractability of the FP equation are being made, such as (1) the magnetic fluctuations are homogeneous in space, (2) the electromagnetic fields are quasi-static, (3) the interaction has a finite decorrelation time, etc. (see more details in the books","Unfortunately, in astrophysical and laboratory plasmas, most of the above assumptions are not valid, yet the FP equation is used extensively, without a proof of its validity."],"Functions Label":["Background","Background"],"Citation Start End":[[889,906]],"Functions Start End":[[515,888],[920,1094]]}
{"Identifier":"2022AandA...667A..15K__Yamashiki_et_al._2019_Instance_1","Paragraph":"In the search for new exoplanets, M dwarfs are ideal targets due to their high abundance in the Galaxy (Bochanski et al. 2010). However, M-type stars are prone to high levels of stellar activity (Walkowicz et al. 2011; Loyd et al. 2016, 2018b) that can impact the radial velocity and\/or transit signal of such systems through phenomena such as flaring (Tofflemire et al. 2012), star spots, plages and faculae (Boisse et al. 2011; Llama & Shkolnik 2015; Cauley et al. 2018; Roettenbacher et al. 2022; Bruno et al. 2022), and other activity-induced variability (Dumusque 2018; Rackham et al. 2019; Bellotti et al. 2022; Collier Cameron et al. 2021). Aside from the observational implications, the planet\u2019s physical and chemical state can be altered by stellar activity as well due to, for example, coronal mass ejections (CMEs) and stellar particle events (SPEs) (Yamashiki et al. 2019; Atri 2017, 2020; Segura et al. 2010), winds (Vidotto et al. 2015; Vidotto & Cleary 2020; Chebly et al. 2022; Colombo et al. 2022), and stellar flares (Segura et al. 2010; Venot et al. 2016; Chadney et al. 2017; Tilley et al. 2019; Chen et al. 2021; Louca et al. 2022), the latter being sudden releases of radiative energy triggered by magnetic reconnection (Benz & G\u00fcdel 2010). Stellar flares result in a temporary increase in incident flux on the planet\u2019s atmosphere, which in turn increases the photochemical reaction rates that can change the chemical composition. Photochemistry, and photolysis in particular, is a key driver of chemical disequilibrium in the atmospheres of close-orbiting, gaseous planets. Photolysis does not only deplete the upper layers from species such as CH4 and NH3, but it can enrich the middle regions with haze precursors such as HCN and C2H2 as well, particularly on cooler planets (Moses et al. 2011; Venot et al. 2012; Zahnle & Marley 2014; Ag\u00fandez et al. 2014; Moses 2014; Miguel & Kaltenegger 2014; Rimmer & Helling 2016; Drummond et al. 2016; Hobbs et al. 2019; Shulyak et al. 2020; Barth et al. 2021; Baeyens et al. 2022). Stellar flares thus have the potential to alter the chemical composition and, subsequently, alter the atmosphere\u2019s signature in transmission spectra.","Citation Text":["Yamashiki et al. 2019"],"Functions Text":["Aside from the observational implications, the planet\u2019s physical and chemical state can be altered by stellar activity as well due to, for example, coronal mass ejections (CMEs) and stellar particle events (SPEs)"],"Functions Label":["Background"],"Citation Start End":[[862,883]],"Functions Start End":[[648,860]]}
{"Identifier":"2020ApJ...901....8B__M\u00fcller-Mellin_et_al._1995_Instance_1","Paragraph":"Three semiannual galactic hydrogen spectra as a function of energy between 40 and 250 MeV have been obtained in three different consecutive time periods (from 2018 August 6 to 2020 January 5) very much inside the heliosphere (1 au); the energy profiles are shown as black circles in Figure 6. Each measured energy spectrum is compared to the theoretical prediction from the HelMod model (Boschini et al. 2019) in the same period (blue solid curve); the maximum and minimum uncertainties related to this prediction are also reported in the plots, as dashed and dotted lines, respectively. As a further comparison, data from the SOHO\/EPHIN spacecraft (red square marker) between 40 MeV and 53 MeV are also presented (M\u00fcller-Mellin et al. 1995). The agreement appears to be good in all the three examined periods, considering both statistical and systematic uncertainties. Ratio between HEPD data and models (displayed in the narrower bottom panels of Figure 6) gradually worsens at lower energies, particularly below 65 MeV, where the spectrum calculated by HEPD is systematically higher. Possible explanations for this discrepancy include a contamination from high-energy protons that is not fully removed using the simulation, and a possible contamination derived from nuclei fragmentation or from very inclined sub-cutoff protons that can enter the FoV of the instrument, even after the rigidity cutoff selection. However, although systematic uncertainties are higher than 10% in the lowest portion of the energy spectra, these results could help constrain theoretical models of particle transport from the border of the heliosphere, down to 1 au. From a comparison between the first spectrum (2018 August 6\u20132019 January 15) and the last one (2019 June 29\u20132020 January 5) an overall increase of \u223c9% is observed, in very good agreement with the variation observed in SOHO\/EPHIN (\u223c8.5%). This behavior is expected, because, as the solar activity continues to wind down (from 2018 to 2020), the effect of the Sun magnetic field diminishes, resulting in higher proton fluxes. On the other hand, HEPD data do not show a clear energy dependence in the modulation over time (typically lower energies should be more modulated with respect to higher ones); unfortunately, for HEPD the overall errors (statistical and systematic) in the first and last energy bins do not allow such a precise evaluation. Table 1 contains explicit values for the galactic hydrogen spectra in the three time periods and for each of the 16 energy bins allowed by the instrument resolution; statistical and systematic uncertainties are also reported.","Citation Text":["M\u00fcller-Mellin et al. 1995"],"Functions Text":["As a further comparison, data from the SOHO\/EPHIN spacecraft (red square marker) between 40 MeV and 53 MeV are also presented","The agreement appears to be good in all the three examined periods, considering both statistical and systematic uncertainties."],"Functions Label":["Uses","Similarities"],"Citation Start End":[[715,740]],"Functions Start End":[[588,713],[743,869]]}
{"Identifier":"2022ApJ...937...58I__Tabatabaei_et_al._2013_Instance_1","Paragraph":"For JW39 and JW100 the disks\u2019 slopes are consistent with linearity, that is the trend expected from Equation (4). This may suggests that the smoothing scale (\u223c10 kpc) is similar to the CRe transport scale, and, hence, the slopes resemble the standard SFR calibrators. Therefore, the spatial correlation slopes observed in the disks of JW39 and JW100 are consistent with the idea that the CRe transport scale in these galaxies is larger than observed in spiral galaxy disks (\u223c1\u20135 kpc). We investigate this result by comparing the spatial scales, L, of the two principal CRe transport processes in these systems: diffusion and advection. In the case of CRe diffusion, the spatial scale is:\n5\n\n\n\nL=4D\u03c4\u22431.1\u00d7D1028cm2s\u22121\u03c410Myrkpc,\n\nwhere \u03c4 is the timescale and D is the diffusion coefficient, which depends on the CRe energy and the local magnetic field power spectrum, and can vary between 1027 and 1029 cm2 s\u22121 (Strong et al. 2007). For advection, the typical spatial scale is:\n6\n\n\n\nL=V\u00b7\u03c4\u2243V100kms\u22121\u03c410Myrkpc,\n\nwhere V is the CRe velocity. Equation (5) shows that to reach L = 10 kpc, for a typical D = 1 \u00d7 1028 cm2 s\u22121, the timescale is of the order of \u223c900 Myr, which is longer than the typical CRe radiative time in galactic disks. Covering these spatial scales in \u2264108 yr requires a diffusion coefficient of D \u2265 9 \u00d7 1028 cm2 s\u22121, which is slightly higher than observed in spiral galaxies (e.g., Strong et al. 2007; Tabatabaei et al. 2013; Heesen et al. 2019). On the other hand, these scales could be consistent with the CRe advection (Equation (6)), as a typical velocity of 100 km s\u22121 would be able to cover 10 kpc in \u223c10 Myr, which is more reasonable for low-energy electrons emitting at 144 MHz. Therefore, our results might hint that CRe transport in the disks of these jellyfish galaxies, JW39 and JW100, is either dominated by advection, due to ram pressure which is stripping the nonthermal ISM (see Section 3.1), or that the CRe diffusion might be more efficient than in normal galaxies. The diffusion coefficient D depends on the local magnetic field configuration and turbulence spectrum (Strong et al. 2007), thus it may be possible that the RPS, by affecting the ISM\u2019s microphysics, may induce higher values of D and, hence, more efficient CRe diffusion (Equation (5)) than observed in normal spiral galaxies. However, we note that, for the rest of the sample, the spatial correlations in their stellar disks are not consistent with linearity (Figure 5). We argue that this might be due to projection effects that mix the disk and the extraplanar emissions, which, on the basis of what we observe for JW39, JO60, JW100, and JO206, follows a flat, almost uniform, distribution. Thus this blend may result in a flattening of the disks\u2019 slopes. Another possible explanation could be a discrepancy between the sampling resolution and the transport scale that was not solved by the smoothing.","Citation Text":["Tabatabaei et al. 2013"],"Functions Text":["Covering these spatial scales in \u2264108 yr requires a diffusion coefficient of D \u2265 9 \u00d7 1028 cm2 s\u22121, which is slightly higher than observed in spiral galaxies"],"Functions Label":["Differences"],"Citation Start End":[[1414,1436]],"Functions Start End":[[1230,1386]]}
{"Identifier":"2019MNRAS.482.5651M__Ruiz-Lapuente_et_al._2004_Instance_1","Paragraph":"To judge the origin of a star in an SNR, its kinetics characteristics may provide very important informations. Generally, except for being stripped-off a part of its envelope, the companion may receive a velocity kick from the supernova ejecta, but the kick velocity is usually much smaller than the orbital velocity (Marietta et al. 2000; Meng et al. 2007; Pakmor et al. 2008; Liu et al. 2012; Pan et al. 2012). Then, the orbital velocity of the companion at the moment of supernova explosion may represent its final space velocity to a great extant, especially for the sdB companions here, which almost do not receive any kick velocity for their large value of $A\/R_{\\rm 2}^{\\rm SN}$. If the spatial velocity of a star in an SNR is very different from the others in the SNR, the star is very possible to be the surviving companion in the remnant (Ruiz-Lapuente et al. 2004). In Fig. 7, we present the companion orbital velocity relative to binary centroid versus the companion radius at the moment of supernova explosion. From the figure, we can see that different companions may have very different orbital velocity. For MS companions, the orbital velocity is from 150 to 200 ${\\rm km\\, s^{\\rm -1}}$, while the RG companions have an orbital velocity of 50\u2013110 ${\\rm km\\, s^{\\rm -1}}$. For the sdB companions, the orbital velocity covers a large range, from 50 to 190 ${\\rm km\\, s^{\\rm -1}}$. Such a large range is mainly derived from the large initial orbital period range for the systems producing sdB companions, i.e. log\u2009(Pi\/d) is from \u223c0.4 to 1.2. In other words, although the mass transfer between a binary system must begin in HG for producing a sdB companion, it may occur at the stage very close to the MS end or at the end of the HG. The upper limit of the orbital velocity here is lower than that in Meng & Podsiadlowski (2017) by 60 ${\\rm km\\,s^{\\rm -1}}$, which originates from the fact that after MWD = 1.378\u2009M\u2299, the binary orbital period increases with mass transfer for a reversed mass ratio (see fig. 2 in Meng & Podsiadlowski 2017). For the same reason, the lower limit of the orbital velocity is also lower than that in Meng & Podsiadlowski (2017).","Citation Text":["Ruiz-Lapuente et al. 2004"],"Functions Text":["If the spatial velocity of a star in an SNR is very different from the others in the SNR, the star is very possible to be the surviving companion in the remnant"],"Functions Label":["Uses"],"Citation Start End":[[849,874]],"Functions Start End":[[687,847]]}
{"Identifier":"2022ApJ...925..123N__Frenklach_&_Feigelson_1989_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":["Frenklach & Feigelson 1989"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[1259,1285]],"Functions Start End":[[1007,1234]]}
{"Identifier":"2022MNRAS.509.6091H___2020a_Instance_1","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"],"Functions Text":["Current cosmological hydrodynamic simulations of galaxy formation employ a variety of subgrid models (e.g.","that artificially launch galactic winds,"],"Functions Label":["Background","Background"],"Citation Start End":[[731,749]],"Functions Start End":[[389,495],[751,791]]}
{"Identifier":"2021MNRAS.503..354G__Cantat-Gaudin_et_al._2020_Instance_2","Paragraph":"The spatial distribution of OB stars and associations, young long-period Cepheids and open clusters, star-forming regions, H\u2009ii regions, interstellar dust, and giant molecular and neutral gas clouds in the solar vicinity that have been in existence generally \u03c4 \u2272 108 yr is known to correlate with the location of the inner Sagittarius, the closest Orion, and outer Perseus spiral arm segments. (The distances for the vast majority of these spiral tracers have been determined in the literature with trigonometric or photometric methods.) The Sun is situated at the inner edge of the Orion arm (Levine et al. 2006; Hou & Han 2014; Nakanishi & Sofue 2016; Xu et al. 2018, 2021; Lallement et al. 2019; Reid et al. 2019; Skowron et al. 2019; Cantat-Gaudin et al. 2020; Fig. 2 above).3 These three spatial features nearby to the Sun appear to form part of the global spiral structure in the Galaxy. Contrary, the objects of older population with larger random velocities, for instance, main-sequence A\u2013K stars or the oldest Cepheids and open clusters, do not currently follow the exact location of those arms (e.g. Cantat-Gaudin et al. 2020, fig. 8 therein; Griv et al. 2020, fig. 7 therein). The latter can be explained by the difference in rotation velocity between the spiral density waves and the objects. Investigating the velocity field of Xu et al.\u2019s (2018) O and early B-type stars in the framework of the Lin\u2013Shu density-wave proposal, we also found that the Sun lies within the Orion arm, at the inner edge of this spiral. The radial distance from the Sun to the centre of the Orion arm is \u22480.2 kpc in the direction of the Galactic anticentre, the centre of the Sagittarius arm is \u22481.8 kpc from the Sun in the direction of the GC, and the width of the arms is \u22480.5 kpc. The radial distance between the centres of the Orion and Sagittarius arms near the Sun is \u03bbrad \u2248 2 kpc (cf. Hou & Han 2014; Wu et al. 2014; Bovy et al. 2015). As for us, the nearest Orion spiral arm forms part of the dominant density-wave structure of the system.","Citation Text":["Cantat-Gaudin et al. 2020"],"Functions Text":["These three spatial features nearby to the Sun appear to form part of the global spiral structure in the Galaxy. Contrary, the objects of older population with larger random velocities, for instance, main-sequence A\u2013K stars or the oldest Cepheids and open clusters, do not currently follow the exact location of those arms (e.g.","fig. 8 therein"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1110,1135]],"Functions Start End":[[781,1109],[1137,1151]]}
{"Identifier":"2017AandA...601A..87C__Falcke_(1996)_Instance_1","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} &amp;&amp;\\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}; \\\\ &amp;&amp;\\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)"],"Functions Text":["The above equation should reduce to the jet Lorentz factor profile used in","when \u03b6 = 1."],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[1077,1090]],"Functions Start End":[[1002,1076],[1114,1125]]}
{"Identifier":"2019ApJ...884..132K__Tanihata_et_al._2003_Instance_2","Paragraph":"First, we discuss the discrepancy of the distribution scale of the radio core positions based on the discussions of the internal shock model (Koyama et al. 2015; Niinuma et al. 2015). As is discussed there, the radio cores in Mrk 501 and Mrk 421 observed at 43 or 22 GHz can usually be considered as the internal shocked regions owing to the convex shape of the radio spectrum peaking around 10 GHz (Giroletti et al. 2008; Sokolovsky et al. 2010; Lico et al. 2012; Blasi et al. 2013). The standard internal shock model of blazars considers that the discrete ejecta with higher speeds (with Lorentz factor \u0393f) catch up with the preceding slower ejecta (with Lorentz factor \u0393s), and the collision leads to the nonthermal emission (e.g., Spada et al. 2001; Tanihata et al. 2003; Guetta et al. 2004; Kino et al. 2004; Ghisellini et al. 2005). Based on the model, the distribution scale of the internal shocks (\u0394DIS in Figure 7, defined as the difference between the largest distance between the internal shock and the central engine DIS,max and the closest one DIS,min) can be explained as the variation of the Lorentz factors of the ejecta (Koyama et al. 2015; Niinuma et al. 2015), by assuming the Lorentz factor ratio (\u0393f\/\u0393s) and the initial separation of the ejecta (IIS). The core stable within 200 \u03bcas constrained by the VERA can be explained by Lorentz factors within a factor of two variation for the slower ejecta, i.e., 8 \u2264 \u0393s \u2264 17, by assuming a minimum value of 8 (e.g., Kino et al. 2002), \u0393f\/\u0393s \u2264 1.01 (Tanihata et al. 2003), and IIS \u223c 1 Rs (Koyama et al. 2015). This time we refined the distribution scale of the radio core within 42 \u03bcas along its main jet axis, or 4.6 \u00d7 103 Rs deprojected (see Figure 7). Based on the same assumptions as in Koyama et al. (2015), to explain the further stable distribution scale of the internal shocks, the variation of Lorentz factors of the slower ejecta is constrained to be much smaller, within 30% or 8 \u2264 \u0393s \u2264 10. On the other hand, the radio core wandering of \u0394DIS \u223c 2.6 \u00d7 105 Rs in Mrk 421 can be explained by the maximum value as \u0393s \u223c 60 (with different assumptions; Niinuma et al. 2015). Even by applying the same assumptions to Mrk 421 as those for Mrk 501, the maximum of the slower Lorentz factor is estimated to be \u0393s \u223c 50, which is still a few times as large as that of Mrk 501. Therefore, even during the X-ray and VHE \u03b3-ray active states in 2012, the maximum Lorentz factors that explain the stability of Mrk 501's core are roughly a few times smaller than those for Mrk 421's wandering core, based on the internal shock model.","Citation Text":["Tanihata et al. 2003"],"Functions Text":["The core stable within 200 \u03bcas constrained by the VERA can be explained by Lorentz factors within a factor of two variation for the slower ejecta, i.e., 8 \u2264 \u0393s \u2264 17, by assuming","\u0393f\/\u0393s \u2264 1.01"],"Functions Label":["Background","Background"],"Citation Start End":[[1512,1532]],"Functions Start End":[[1273,1450],[1498,1510]]}
{"Identifier":"2019ApJ...886...34F__Sahijpal_&_Goswami_1998_Instance_2","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":["Sahijpal & Goswami 1998"],"Functions Text":["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"],"Functions Label":["Similarities"],"Citation Start End":[[1542,1565]],"Functions Start End":[[1376,1540]]}
{"Identifier":"2021MNRAS.501.2934C__Pinte_et_al._2019_Instance_1","Paragraph":"Understanding how the diverse populations of protoplanetary discs in young stellar regions results in the range of exoplanet types and architectures found in the Galaxy is one of the major goals of planet-formation theory. This is an extremely challenging task due in part to the limited observational constraints available. The Atacama Large Millimetre\/submillimetre Array (ALMA) is providing truly transformational images of protoplanetary discs with unprecedented sensitivity and resolution (Andrews 2020). However, millimetre wavelength images reveal the locations of small dust grains but provide little information on the presence of larger particles, beyond centimetre scales. Gas giant planets are mostly made of hydrogen and helium, which ALMA cannot directly observe; therefore, the information on the gas content relies on the observations of less abundant molecules, such as CO and its isotopologues, that are subjected to uncertain depletion processes in the gas-phase (e.g. Miotello et al. 2016). Planets might be detectable by ALMA, although indirectly, by the effects they have on the gas and\/or dust in the disc. When planets become massive enough, they can carve gaps (e.g. Rice et al. 2006; Pinilla, Benisty & Birnstiel 2012; Zhu et al. 2012) and disturb the dynamics of the gas (Teague et al. 2018; Casassus & P\u00e9rez 2019; Pinte et al. 2019). The minimum gap-opening mass depends on the viscosity and scale-height of the disc (Crida, Morbidelli & Masset 2006; Duffell & MacFadyen 2013), but mini-Neptune-mass (P\u00e9rez et al. 2019) or even Earth-mass planets (Rosotti et al. 2016; Dong & Fung 2017) could produce detectable gaps. Gaps consistent with fully formed planets have been imaged by ALMA in discs with estimated ages ranging from  1 Myr (HL Tau and Elias 2\u201324; ALMA Partnership et al. 2015; Cieza et al. 2017) to \u223c10 Myr (TW Hydra; Andrews et al. 2016). However, the origin of these gaps still remains to be established and several alternative explanations have been proposed, including the effect of snow-lines on the dust\/gas evolution of different volatiles (Zhang, Blake & Bergin 2015), magneto-hydrodynamic effects (Flock et al. 2015), secular gravitational instability (e.g. Youdin 2011; Takahashi & Inutsuka 2014), and viscous ring-instabilities (Dullemond & Penzlin 2018). Each one of the proposed mechanisms has their merits and shortcomings, and it is possible that different mechanisms operate together or dominate in different objects or in the same object at at different times. For a recent review on disc (sub)structures, see Andrews (2020). Substructures are also expected to be ubiquitous in protoplanetary discs from a theoretical point of view. Without substructures to halt the migration of mm-size grains at large radii, dust particles should migrate towards the innermost part of the disc in time-scales shorter than 0.1 Myr (e.g. Brauer et al. 2007), which is inconsistent with the observations showing significant mm emission at large radii (\u227310 au) at much older ages. Understanding the origin and evolution of substructures in protoplanetary discs and their implications for planet formation is currently one the main challenges in the field. To better understand the incidence and properties of disc substructures in any given molecular cloud, here we present 1.3 mm\/230 GHz continuum ALMA long-baseline observations at 3\u20135 au resolution of the 10 brightest targets of the \u2018Ophiuchus DIsc Survey Employing ALMA\u2019 (ODISEA) project (Cieza et al. 2019) that were not included in \u2018The disc Substructures at High Angular Resolution Project\u2019 (DSHARP) ALMA Cycle-4 Large Program (Andrews et al. 2018). Our new observations result in the largest sample of disc images at \u223c3\u20135 au resolution in any star-forming region observed so far at mm wavelengths (15 objects when combined with the brightest Ophiuchus objects in DSHARP). In Section 2, we discuss the sample selection, the long-baseline observations, and the data reduction. In Section 3, we characterize the observed substructures, including gaps, rings, inner discs, and cavities. In Section 4, we discuss individual objects and use the full sample of 15 bright Ophiuchus discs observed at high-resolution to construct a tentative evolutionary sequence in which the observed substructures are mostly driven by dust evolution and the formation of giant planets. We also discuss possible connections between the substructures observed in primordial discs and those seen in more evolved debris disc systems. A summary of our results and conclusions is presented in Section 5.","Citation Text":["Pinte et al. 2019"],"Functions Text":["When planets become massive enough, they can carve gap","and disturb the dynamics of the gas"],"Functions Label":["Background","Background"],"Citation Start End":[[1342,1359]],"Functions Start End":[[1130,1184],[1262,1297]]}
{"Identifier":"2016ApJ...829...29T__Haiman_et_al._2000_Instance_1","Paragraph":"Among various DM candidates, the most popular candidate is the weakly interacting massive particles (WIMPs; like the neutralino), which have mass in the GeV range (Jungman et al. 1996; Bertone et al. 2005; Hooper & Profumo 2007; Feng 2010). The WIMPs are non-relativistic at the epoch of decoupling from the interacting particles and have negligible free-streaming velocities. Therefore, they are \u201ccold,\u201d called cold dark matter (CDM). In  the CDM scenario, \u201chalos\u201d formed in small clumps, and then merged together into larger and massive objects. Galaxies formed in these halos because of the cooling of atomic hydrogen (H; Tegmark et al. 1997) or molecular hydrogen (H2, Ciardi et al. 2000; Haiman et al. 2000). On large cosmological scales (from the range \n\n\n\n\n\n down to \n\n\n\n\n\n), the CDM paradigm has had great success in explaining the observed universe and reproducing the luminous structures (Fixsen et al. 1996; Borgani & Guzzo 2001; Lange et al. 2001; Cole et al. 2005; Tegmark et al. 2006; Benson 2010; Hinshaw et al. 2013; Slosar et al. 2013; Wang 2013; Planck Collaboration et al. 2014; Wei et al. 2016). However, on small scales (\n\n\n\n\n\n), there are still some discrepancies between the CDM paradigm and observations. (a) The core\u2013cusp problem (Navarro et al. 1997; Subramanian et al. 2000). CDM simulations predict a cusp\u2013core DM halo, whereas the observations find them cored (Salucci et al. 2012). (b) The too big to fail problem (Boylan-Kolchin et al. 2012). CDM simulations predict a central DM density significantly higher than observation allows. (c) The \u201cmissing satellite problem.\u201d N-body simulations based on the CDM paradigm predict a number of subhalos larger than that of satellites found in our Galaxy (Klypin et al. 1999; Moore et al. 1999; Papastergis et al. 2011). Many methods have been proposed to solve these small-scale problems, such as modifying the nature of the DM from the CDM paradigm (Hu et al. 2000; Spergel & Steinhardt 2000; Su & Chen 2011; Menci et al. 2012), adding supernova feedback effect in simulation (Weinberg & Katz 2002; Mashchenko et al. 2006; Governato et al. 2010; Pontzen & Governato 2014), and considering the interplay between DM and baryons during the formation of the galaxy (El-Zant et al. 2001; Tonini et al. 2006; Pontzen & Governato 2014). However, these methods are insufficient to solve all of the above problems.","Citation Text":["Haiman et al. 2000"],"Functions Text":["Galaxies formed in these halos because of the cooling of","or molecular hydrogen"],"Functions Label":["Background","Background"],"Citation Start End":[[693,711]],"Functions Start End":[[548,604],[646,667]]}
{"Identifier":"2015ApJ...798...95B___2010_Instance_1","Paragraph":"The gravitational microlensing of lensed quasars has proven to be an effective tool for measuring the properties of quasar accretion disks, and is starting to become useful for studying the X-ray corona as well. The time-dependent microlensing magnification (or demagnification) of one or more images of a lensed quasar is moderated by the finite size of the source, which smooths the complicated caustic pattern of microlensing magnifications as the quasar passes over it. This allows us to use the microlensing magnifications to estimate the source size, and such work has shown that in general the accretion disks are larger than would be expected from either thin disk modeling or total flux arguments (Pooley et\u00c2 al. 2007; Anguita et\u00c2 al. 2008; Morgan et\u00c2 al. 2010; Hainline et\u00c2 al. 2012; Jim\u00c3\u00a9nez-Vicente et\u00c2 al. 2012). Since the effective temperature of the disk depends on radius, its apparent size depends on wavelength, leading to a chromatic dependence of the microlensing magnification, with the same quasar image experiencing larger variability at blue wavelengths than at red wavelengths. Several studies have used this to constrain the power-law slope of the size-wavelength relation, and the results have been consistent with each other and with the thin disk prediction that the size goes like the four-thirds power of the wavelength, mostly because of their large uncertainties (Poindexter et\u00c2 al. 2008; Bate et\u00c2 al. 2008; Eigenbrod et\u00c2 al. 2008; Floyd et\u00c2 al. 2009; Blackburne et\u00c2 al. 2011; Mosquera et\u00c2 al. 2011). Finally, efforts to put upper limits on the size of the X-ray regions have also been successful (Pooley et\u00c2 al. 2006, 2007; Chartas et\u00c2 al. 2009; Dai et\u00c2 al. 2010), and recently there have been attempts to constrain the direct and reflected components' sizes separately using color cuts or spectral decomposition (Chen et\u00c2 al. 2011; Blackburne et\u00c2 al. 2014; Morgan et\u00c2 al. 2012; Chen et\u00c2 al. 2012; Chartas et\u00c2 al. 2012; Mosquera et\u00c2 al. 2013).","Citation Text":["Morgan et\u00c2 al. 2010"],"Functions Text":["This allows us to use the microlensing magnifications to estimate the source size, and such work has shown that in general the accretion disks are larger than would be expected from either thin disk modeling or total flux arguments"],"Functions Label":["Background"],"Citation Start End":[[750,769]],"Functions Start End":[[474,705]]}
{"Identifier":"2018MNRAS.473.2020L__Preibisch_et_al._1998_Instance_1","Paragraph":"Upper Scorpius (hereafter UpSco) is part of the nearest OB association to the Sun, Scorpius Centaurus. The region is nearby, with a distance of \u223c145 pc from Hipparcos (de Bruijne et al. 1997) and a recent update from Gaia (144.2 \u00b1 17.6 pc: Fang, Herczeg & Rizzuto 2017). UpSco is young, with different age determinations and a possible spread among its members (Preibisch & Zinnecker 1999; Preibisch, Guenther & Zinnecker 2001; Pecaut, Mamajek & Bubar 2012; Song, Zuckerman & Bessell 2012; Pecaut 2016; Rizzuto et al. 2016). The members of UpSco exhibit a significant mean proper motion (\u03bc\u03b1cos\u2009\u03b4 = \u221210.5 and \u03bc\u03b4 = \u221223.2 mas yr\u22121, with a standard dispersion of about 6.4 mas yr\u22121: de Bruijne et al. 1997; de Zeeuw et al. 1999; Fang et al. 2017). The bright end of the UpSco population has been examined in X-rays (Walter et al. 1994; Kunkel 1999; Preibisch et al. 1998), astrometrically (de Bruijne et al. 1997; de Zeeuw et al. 1999) and spectroscopically (Preibisch & Zinnecker 2002). The low-mass and substellar population has been investigated over the past decade in great detail with the advent of modern detectors, permitting wide and\/or deep surveys (Ardila, Mart\u00edn & Basri 2000; Mart\u00edn, Delfosse & Guieu 2004; Slesnick, Carpenter & Hillenbrand 2006; Lodieu, Hambly & Jameson 2006; Lodieu et al. 2007, 2011b; Kraus et al. 2008; B\u00e9jar et al. 2008; Slesnick, Hillenbrand & Carpenter 2008; Lafreni\u00e8re, Jayawardhana & van Kerkwijk 2010; Dawson, Scholz & Ray 2011; Lafreni\u00e8re et al. 2011, 2014; Dawson et al. 2013, 2014; Lodieu 2013; Pe\u00f1a Ram\u00edrez, B\u00e9jar & Zapatero Osorio 2016; Best et al. 2017). The first transiting systems in the association have been announced in the last years, thanks to the Kepler K2 mission (Borucki et al. 2010; Lissauer, Dawson & Tremaine 2014; Batalha 2014), including a triple system composed of an F star and two solar-type stars (Alonso et al. 2015) and several M dwarf eclipsing binaries (Kraus et al. 2015; Lodieu et al. 2015; David et al. 2016). These systems are of prime importance, because they provide the first mass and radius measurements independent of models in this region. Finally, we should note the existence of a transiting Neptune-size planet candidate around an M3 member of UpSco (Mann et al. 2016; David et al. 2016).","Citation Text":["Preibisch et al. 1998"],"Functions Text":["The bright end of the UpSco population has been examined in X-rays"],"Functions Label":["Background"],"Citation Start End":[[845,866]],"Functions Start End":[[744,810]]}
{"Identifier":"2021AandA...651A..87O__Brunthaler_et_al._2021_Instance_1","Paragraph":"To complement our study, we also analyzed GLOSTAR continuum images toward sites with maser emission. A full description of the GLOSTAR continuum data calibration and imaging is given in Brunthaler et al. (2021), while the full analysis of continuum images of Cygnus X will be presented in a forthcoming paper. Here, we briefly discuss the imaging strategy. The calibration and imaging of the continuum data was performed with the Obit package (Cotton 2008). The 2 GHz bandwidth was first rearranged into nine frequency subbands, which were used to image each pointing individually. Then, for each frequency subband the pointings were combined into large individual mosaics to cover the entire observed area. Finally, we combined the different frequencies to obtain the image at the reference frequency, which has circular beams of 19\u2033 and 1.5\u2033 in the D and Bconfiguration, respectively. Continuum and methanol line maps from Effelsberg observations have also been obtained as part of the GLOSTAR survey (Brunthaler et al. 2021, Rugel et al., in prep.) We note that continuum images were constructed for Effelsberg data, the VLA D configuration, the VLA B configuration, a combination of the VLA D and B (D+B) configurations, and a combination of the VLA D configuration and Effelsberg observations. The central frequency of these images is 5.8 GHz. Here, we only use B-configuration continuum maps to study the region of the investigated methanol maser positions and D+B maps of the region around DR21 (see Sect. 4.5). Methanol line data from Effelsberg were also inspected to look for flux variations in the VLA-detected masers (Sect. 4.4). The noise in the continuum images is not uniform, but rather varies across the mapped region, and can be high around strong sources with complex or extended emission. We locally measured the noise in regions close to the maser locations, resulting in 1\u03c3 values in the range from 0.056 to 0.43 mJy beam\u22121 for B-configuration images. For the D configuration, the 1\u03c3 rms noise ranges from 0.10 to 2.6 mJy beam\u22121. The higher values measured in D-configuration data are due to bright extended emission, which is present across the Cygnus X region, and are resolved out by the array in the B configuration. The highest local rms noise occurs around the strong radio source, DR21, a compact HII region.","Citation Text":["Brunthaler et al. (2021)"],"Functions Text":["A full description of the GLOSTAR continuum data calibration and imaging is given in"],"Functions Label":["Uses"],"Citation Start End":[[186,210]],"Functions Start End":[[101,185]]}
{"Identifier":"2015ApJ...804..101Y__Hawley_et_al._1995_Instance_1","Paragraph":"It is interesting to note that for wind originating within \n\n\n\n\n\n, their value of Be becomes almost constant when \n\n\n\n\n\n. This corresponds to the slight decrease of the poloidal velocity beyond \n\n\n\n\n\n shown in the left panel of Figure 8. The reason why Be does not change beyond \n\n\n\n\n\n is because in this region turbulence has not well developed within our simulation time. Note that this radius is different from the inflow equilibrium radius, which is \n\n\n\n\n\n. Within \n\n\n\n\n\n, everything, especially the radial density profile, is fully reliable. Beyond this radius, the density profile is not reliable, but other properties, such as the level of turbulence and subsequently outflow properties, are still reliable up to a much larger \u201cturbulence radius,\u201d the limiting radius of turbulence steady state. This radius can be estimated as follows. Turbulence in accretion flow is because of MRI. The fastest growth rate of MRI at radius r is \n\n\n\n\n\n. More precisely, it takes 3\u20134 orbits for MRI to develop and \u223c10 orbits to saturate (Hawley et al. 1995). For our simulation time of \n\n\n\n\n\n, taking a timescale of 3 orbits, we can obtain that the \u201cturbulence radius\u201d is \n\n\n\n\n\n. This is close to the value of 800 rg mentioned above. Another way to understand the \u201cturbulence radius\u201d is as follows. For a geometrically thick disk, the largest turbulence eddies have size of order r. The corresponding eddy turnover time is \n\n\n\n\n\n, where \n\n\n\n\n\n is the rms turbulent velocity. Our simulation data show \n\n\n\n\n\n. If at a certain radius the eddy turnover time is substantially smaller than the duration of the simulation, then the local turbulence is likely to have reached quasi-steady state. Therefore, the \u201cturbulence radius\u201d should be some fraction of \n\n\n\n\n\n, which gives a similar result to the above estimation. Therefore, we think that the results beyond \n\n\n\n\n\n are not reliable. It is very likely that Be will keep changing and the poloidal velocity still remains constant beyond \n\n\n\n\n\n. This implies that wind can at least escape beyond the outer boundary of accretion flows. Simulations with longer run times can check this point.","Citation Text":["Hawley et al. 1995"],"Functions Text":["More precisely, it takes 3\u20134 orbits for MRI to develop and \u223c10 orbits to saturate"],"Functions Label":["Uses"],"Citation Start End":[[1029,1047]],"Functions Start End":[[946,1027]]}
{"Identifier":"2017AandA...601A.134M__Fung_&_Dong_(2015)_Instance_2","Paragraph":"Several predictions for planet(s) shaping the disk of SAO 206462 have been proposed (Muto et al. 2012; Garufi et al. 2013; Fung & Dong 2015; Bae et al. 2016; van der Marel et al. 2016a. Using linear equations from the spiral density wave theory, Muto et al. (2012) suggested two planets with separations beyond ~50 au by fitting independently the two spiral arms seen in Subaru\/HiCIAO data and with masses of ~0.5 MJ by using the amplitude of the spiral wave. Garufi et al. (2013) proposed that one planet of mass 5\u201313 MJ located inside the cavity at a separation of 17.5\u201320 au could be responsible for the different cavity sizes measured for the small and large dust grains. Fung & Dong (2015) presented scaling relations between the azimuthal separation of the primary and secondary arms and the planet-to-star mass ratio for a single companion on a circular orbit with a mass between Neptune mass and 16 MJ around a 1 M\u2299 star. They predicted with 30% accuracy that a single putative planet responsible for both spiral features of SAO 206462 would have a mass of ~6 MJ. Bae et al. (2016) presented dedicated hydrodynamical simulations of the SAO 206462 disk and proposed that both the bright scattered-light feature (Garufi et al. 2013) and the dust emission peak (P\u00e9rez et al. 2014) seen for the southwestern spiral arm result from the interaction of the spiral arm with a vortex, although a vortex alone can account for the S1 brightness peak. They suggested that a 10\u201315 MJ planet may orbit at 100\u2013120 au from the star. However, ALMA observations at two different frequencies seem to contradict a dust particle trapping scenario by a vortex (Pinilla et al. 2015). Stolker et al. (2016) performed new fitting of the spiral arms observed in SPHERE data and found a best-fit solution with two protoplanets located exterior to the spirals: r1\u2009~\u2009168 au, \u03b81\u2009~\u200952\u00b0 and r2\u2009~\u200999 au, \u03b82\u2009~\u2009355\u00b0. van der Marel et al. (2016a) proposed that the features seen in thermal emission in ALMA data and the scattered-light spiral arms are produced by a single massive giant planet located inside the cavity at a separation of ~30 au. Recently, Dong & Fung (2017) used the contrast of the spiral arms to predict a giant planet of ~5\u201310 MJ at ~100 au. ","Citation Text":["Fung & Dong (2015)"],"Functions Text":["presented scaling relations between the azimuthal separation of the primary and secondary arms and the planet-to-star mass ratio for a single companion on a circular orbit with a mass between Neptune mass and 16 MJ around a 1 M\u2299 star. They predicted with 30% accuracy that a single putative planet responsible for both spiral features of SAO 206462 would have a mass of ~6 MJ."],"Functions Label":["Background"],"Citation Start End":[[676,694]],"Functions Start End":[[695,1071]]}
{"Identifier":"2019ApJ...883...73C___2010_Instance_1","Paragraph":"Assuming that the force-field approach to the solution of the Parker (1965) cosmic-ray transport equation is valid, the connection between historic cosmic-ray intensities and the solar properties they encountered lies in the effective diffusion coefficient that is assumed in this approximation. Establishing such a connection, however, is no simple task. Many theories have been proposed to describe the scattering of cosmic rays in the heliosphere. The most likely candidates for this task, given their reasonable agreement with observations and numerical simulations of cosmic-ray diffusion coefficients, are the quasilinear theory (QLT) of Jokipii (1966), the weakly nonlinear theory (WNLT) of Shalchi et al. (2004b), and the nonlinear guiding center theory (NLGC; or one of its variants; see, e.g., Matthaeus et al. 2003; Shalchi 2006, 2009, 2010; Ruffolo et al. 2012). Shalchi (2009) provides in depth theoretical treatments of most the abovementioned theories. These scattering theories all require as a key input an expression for the power spectrum of the turbulent fluctuations of the HMF. These spectra depend upon basic turbulence quantities, such as the magnetic variance, and various correlation scales. Turbulence power spectra are discussed in detail by, e.g., Batchelor (1970) and Matthaeus et al. (2007), whereas more background on the abovementioned turbulence quantities can be found in, e.g., Matthaeus & Goldstein (1982), Petrosyan et al. (2010), Matthaeus & Velli (2011), and Bruno & Carbone (2013). These basic turbulence quantities have been observed to show a marked dependence on the solar cycle at Earth (see, e.g., Smith et al. 2006b; Burger et al. 2014; Zhao et al. 2018). It follows then that mean free paths derived from these scattering theories would be expected to depend on the solar cycle as well, and several studies have reported such a dependence. Chen & Bieber (1993) find from an analysis of cosmic-ray anisotropies and gradients as observed by means of NMs, that larger mean free paths are associated with solar minima, and smaller mean free paths with solar maxima. The authors also report a mean free path dependence on solar magnetic polarity. Nel (2016) and Zhao et al. (2018) both extensively analyze spacecraft observations, using the turbulence quantities so calculated as inputs for expressions for diffusion coefficients derived from the QLT and NLGC theories. Both authors report that the resulting mean free paths display solar cycle dependences.","Citation Text":["Matthaeus et al.","2010"],"Functions Text":["The most likely candidates for this task, given their reasonable agreement with observations and numerical simulations of cosmic-ray diffusion coefficients, are the quasilinear theory (QLT) of Jokipii (1966), the weakly nonlinear theory (WNLT) of Shalchi et al. (2004b), and the nonlinear guiding center theory (NLGC; or one of its variants; see, e.g.,"],"Functions Label":["Background"],"Citation Start End":[[804,820],[847,851]],"Functions Start End":[[451,803]]}
{"Identifier":"2018AandA...610A..44M__Kr\u00fcger_&_Dreizler_(1992)_Instance_2","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)"],"Functions Text":["Our analysis was rather easy, starting from a prediction based on","parameters."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1117,1141]],"Functions Start End":[[1051,1116],[1142,1153]]}
{"Identifier":"2021MNRAS.500..291B__Serafinelli_et_al._2019_Instance_1","Paragraph":"We have presented the analysis of the current X-ray observations of the disc wind in MCG-03-58-007. Here, multiple and variable wind components with velocities ranging from $\\sim \\! -0.08\\, c$ to $\\sim \\! - 0.2\\, c$ (and potentially up to $0.35\\, c$) are seen at different times. Multi-epoch observations of disc winds, like the one presented here, are crucial for revealing all the possible phases of the disc wind. For example, over a decade worth of observations of PDS\u2009456 revealed that the wind is most likely clumpy and\/or stratified with the ionization ranging from log\u2009($\\xi \/\\rm {erg\\, cm \\, s^{-1})}\\sim 2$\u2009erg cm s\u22121 up to log\u2009($\\xi \/\\rm {erg\\, cm \\, s^{-1})}=6$\u2009erg\u2009cm\u2009s\u22121 and velocities ranging from $\\sim \\! -0.2\\, c$ up to $\\sim \\! -0.46\\, c$ (Reeves et al. 2016, 2018a, 2020). It is possible, as suggested in other examples of ultra fast disc winds, that we are looking at a stratified wind, where multiple components are launched at different disc radii, but not all of them are always detected. This adds MCG-03-58-007 to the small but growing list of multiphase fast X-ray winds. Other examples of AGN with at least two variable phases of the X-ray winds are PG\u20091211+143 (Pounds et al. 2016; Reeves et al. 2018b), IRAS 13224-3809 (Parker et al. 2017; Chartas & Canas 2018; Pinto et al. 2018), 1H\u20090707-495 (Kosec et al. 2018), IRAS 17020+4544 (Longinotti et al. 2015), and PG 1114+445 (Serafinelli et al. 2019). In those cases, multiple phases with a common or different outflowing velocities are detected in the X-ray band. In contrast to most of the cases reported so far, neither of the two phases seen in MCG-03-58-007 requires a different ionization (aside from slice\u2009B) suggesting that we are seeing different streamlines of the same highly ionized flow. The only exception could be the eclipsing event seen in 2015, where a solution is found with a lower ionization for the Fe K intervening absorber. However, what we most likely see during this occultation event is a higher density and lower ionization clump of the wind, which could be faster because its higher opacity makes it easier to accelerate (Waters et al. 2017). Note that this does not imply that the soft X-ray wind components, like the ones seen for example in PDS\u2009456 or PG\u20091211+143, are not present; in contrast to the other examples, MCG-03-58-007 is seen through a relatively high column density (NH \u223c 2 \u00d7 1023\u2009cm\u22122, see Table 2) neutral absorber, therefore these phases may be hidden behind it. MCG-03-58-007 is not the only example where multiple Fe-K zones with the same ionization and outflowing with different velocities had been detected in a single observation. For instance, two simultaneous Fe-K phases were detected at least twice in PDS456 (Reeves et al. 2018a, 2020) and possibly in PG\u20091211+143 (Pounds et al. 2016) and IRAS\u200913349+2438 (Parker et al. 2020).","Citation Text":["Serafinelli et al. 2019"],"Functions Text":["Other examples of AGN with at least two variable phases of the X-ray winds are","and PG 1114+445"],"Functions Label":["Background","Background"],"Citation Start End":[[1404,1427]],"Functions Start End":[[1099,1177],[1387,1402]]}
{"Identifier":"2022AandA...662A..42M__N\u00f3brega-Siverio_et_al._2020b_Instance_1","Paragraph":"In the second part of the paper, we propose the set of self-similar solutions as tests for MHD numerical codes with ambipolar diffusion capabilities. To show their usefulness and validity, a battery of tests was carried out for the Bifrost code in two spatial dimensions starting from initial conditions with cylindrical symmetry (Sect. 4). We showed that the ambipolar diffusion module in Bifrost can cope with the passage of the solutions through the current sheets, with the level of accuracy increasing the higher the spatial resolution and in spite of the intrinsic singularity in them. Vice versa, the tests show that these functions can probe the capabilities of ambipolar diffusion modules to a larger extent than the simple ZKBP solution that has been used thus far (e.g. Masson et al. 2012; Vigan\u00f2 et al. 2019; N\u00f3brega-Siverio et al. 2020b). As test functions, the various harmonics proposed in our paper have the comparative advantage that they combine the B\u2004\u221d\u2004|\u2006\u03b4\u2006|1\/3 singularity at the internal nulls (with \u03b4 the distance to the null) with the finiteness of the nonlinear diffusive flux, and this combination must be sufficiently well reproduced by the code if it is to pass the test. The ZKBP solution, instead, has a null just at the outer front, and the singularity there is of a lower order (\u221d|\u03b4|1\/2), with zero diffusive flux across it. On the other hand, the scarcity of tests for the ambipolar diffusion term until now is in contrast with, for instance, the case of HD shocks, for which a whole category of exact solutions is available (the solutions of the Riemann problem) that have been used to develop sophisticated numerical schemes and tests (see Laney 1998; Toro 2009). The contrast to shocks, in fact, is interesting because of the differences in their mathematical and physical nature: in shocks, it is the (magneto)hydrodynamic evolution of the hyperbolic components of the PDE that leads to the formation of the singularities, which is then smoothed through simple diffusive phenomena (typically viscosity). In the ambipolar diffusion problem, wherever there is a null, it is the diffusive phenomenon itself that creates and maintains the singularity.","Citation Text":["N\u00f3brega-Siverio et al. 2020b"],"Functions Text":["We showed that the ambipolar diffusion module in Bifrost can cope with the passage of the solutions through the current sheets, with the level of accuracy increasing the higher the spatial resolution and in spite of the intrinsic singularity in them. Vice versa, the tests show that these functions can probe the capabilities of ambipolar diffusion modules to a larger extent than the simple ZKBP solution that has been used thus far (e.g."],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[821,849]],"Functions Start End":[[341,780]]}
{"Identifier":"2015ApJ...799..170C__Roxburgh_&_Vorontsov_2003_Instance_1","Paragraph":"Stellar properties were determined using three different techniques to model the oscillation frequencies extracted from the data. The first method relies on a dense grid of stellar models computed with the GARching STellar Evolution Code (GARSTEC; Weiss & Schlattl 2008) including the effects of microscopic diffusion, and on theoretical frequencies calculated using the Aarhus aDIabatic PuLSation code (Christensen-Dalsgaard 2008a). The results were obtained implementing a Bayesian scheme that uses the spectroscopic constraints and frequency ratios as the parameters in the fit (V. Silva Aguirre et al., submitted), the latter being almost insensitive to the surface effects in solar-like oscillators (Roxburgh & Vorontsov 2003; Silva Aguirre et al. 2011). Central values are given as the estimates of the stellar properties obtained in this manner. We also computed models using the ASTEC and YREC codes. In these cases, the fit was made to the individual frequencies after correcting for the surface effect with, respectively, an appropriately scaled version of the observed solar surface correction (Christensen-Dalsgaard 2012) and a solar-type correction as described in Carter et al. (2012). The stellar properties derived using the three techniques described above are consistent within the returned formal errors. Therefore, we added in quadrature the difference in central values of each property to the formal uncertainties determined from the GARSTEC Bayesian scheme as a measurement of the systematic spread arising from different codes and fitting techniques. In Table 2, we provide a precise estimate of the stellar age, t, from detailed frequency modeling. Values for the remaining fundamental stellar properties are consistent, within errors, with those obtained from grid-based modeling. In particular, no gain in precision was obtained for the stellar radius. On the other hand, the precision on the stellar mass is improved by nearly a factor of two, from 5.7% to 3.2%.","Citation Text":["Roxburgh & Vorontsov 2003"],"Functions Text":["The results were obtained implementing a Bayesian scheme that uses the spectroscopic constraints and frequency ratios as the parameters in the fit",", the latter being almost insensitive to the surface effects in solar-like oscillators"],"Functions Label":["Uses","Uses"],"Citation Start End":[[705,730]],"Functions Start End":[[434,580],[617,703]]}
{"Identifier":"2015AandA...577A..43S__Odstr\u010dil_&_Karlick\u00fd_(1997)_Instance_2","Paragraph":"The initialization of solar flares remains an unsolved problem. Early ideas on how the initialization might occur were described by Norman & Smith (1978). They argued that flare process cannot start in the entire flare volume at one instant, and proposed that the flare onset was localized in a small part of an active region, from which the energy release extends as dissipation spreading process throughout the flare volume. Two types of agents that may lead to this kind of a dissipation process were addressed: electron beams and shock waves. These agents can trigger flares at large distances from their initial locations, causing sympathetic (simultaneous) flares or leading to a sequential flare energy release in one active region (Liu et al. 2009; Zuccarello et al. 2009). These triggering processes were numerically studied by Karlick\u00fd & Jungwirth (1989) and Odstr\u010dil & Karlick\u00fd (1997). Karlick\u00fd & Jungwirth (1989) assumed that electron beams, penetrating into the current sheet in the magnetic reconnection region, generate Langmuir waves. Then, using the particle-in-cell model, the authors studied the effects of these electrostatic waves on the plasma system. Sufficiently strong Langmuir waves were found to be able to generate ion-sound waves through the three-wave decay process (B\u00e1rta & Karlick\u00fd 2000). These ion-sound waves increase electrical resistivity in the current sheet system, which results in the onset of the energy dissipation. Thus, the electron beams are able to cause magnetic reconnection. Odstr\u010dil & Karlick\u00fd (1997) studied the mechanism for the flare trigger by shock waves. They used a 2D magnetohydrodynamic model with the MHD shock wave propagating towards the current sheet. A portion of the shock wave passed through the sheet, and the rest was reflected. Nothing occurred at the very beginning of the wave-current sheet interaction. However, after some time, specific plasma flows around the current sheet were formed, which led to the start of magnetic reconnection. This shows that for reconnection to be triggered, the enhanced electrical resistivity as well as the plasma flows are important. ","Citation Text":["Odstr\u010dil & Karlick\u00fd (1997)"],"Functions Text":["studied the mechanism for the flare trigger by shock waves. They used a 2D magnetohydrodynamic model with the MHD shock wave propagating towards the current sheet. A portion of the shock wave passed through the sheet, and the rest was reflected. Nothing occurred at the very beginning of the wave-current sheet interaction. However, after some time, specific plasma flows around the current sheet were formed, which led to the start of magnetic reconnection.","This shows that for reconnection to be triggered, the enhanced electrical resistivity as well as the plasma flows are important."],"Functions Label":["Background","Motivation"],"Citation Start End":[[1524,1550]],"Functions Start End":[[1551,2009],[2010,2138]]}
{"Identifier":"2022AandA...659A.124H__Harrison_et_al._2018_Instance_1","Paragraph":"Active galactic nuclei (AGN) have drawn a lot of attention over the last decades because they have been beacons for the existence and demographics of super-massive black holes (BHs) throughout the history of the Universe (e.g., Soltan 1982; Kollmeier et al. 2006; Greene & Ho 2007; Schulze & Wisotzki 2010; Kelly & Shen 2013; Kormendy & Ho 2013; Schulze et al. 2015; Ba\u00f1ados et al. 2018). The release of gravitational binding energy via accretion of matter onto these BHs is expected to have a profound impact on the evolution of their host galaxies (e.g., Silk & Rees 1998; Granato et al. 2004; Di Matteo et al. 2005; Springel et al. 2005; Hopkins et al. 2008; Somerville et al. 2008; Fabian 2012; Harrison 2017; Harrison et al. 2018; Gaspari et al. 2019; Nelson et al. 2019). Large spectroscopic surveys such as the Sloan Digital Sky Survey (SDSS, York et al. 2000; Abazajian et al. 2009), the 2df QSO redshift survey (2QZ, Croom et al. 2004), the VIMOS VLT Deep Survey (VVDS, Le F\u00e8vre et al. 2013) or the VIMOS Public Extragalactic Redshift Survey (VIPERS, Scodeggio et al. 2018) in combination with several deep X-ray surveys taken with ROSAT (Voges et al. 1999), Chandra (Elvis et al. 2009; Xue et al. 2011), XMM-Newton (Pierre et al. 2016), and eROSITA (Predehl et al. 2021) as well as large radio surveys (e.g., Becker et al. 1995; Condon et al. 1998; Smol\u010di\u0107 et al. 2017; Shimwell et al. 2019; Lacy et al. 2020; Gordon et al. 2021) have provided an enormous data set to characterize the AGN population and its evolution with redshift in great detail. While the standard model for the AGN central engine has been successful in unifying the various classes of AGN appearance (Antonucci 1993; Urry & Padovani 1995; Padovani et al. 2017), some aspects such as changing-look AGN (CLAGN, e.g. MacLeod et al. 2016; Ruan et al. 2016; Graham et al. 2017; Noda & Done 2018) and tidal-disruption events (e.g., Gezari et al. 2009; Merloni et al. 2015; Auchettl et al. 2017) are just being explored more extensively in the time domain.","Citation Text":["Harrison et al. 2018"],"Functions Text":["The release of gravitational binding energy via accretion of matter onto these BHs is expected to have a profound impact on the evolution of their host galaxies (e.g.,"],"Functions Label":["Motivation"],"Citation Start End":[[714,734]],"Functions Start End":[[389,556]]}
{"Identifier":"2019AandA...623A.140G__Pohl_et_al._2017_Instance_1","Paragraph":"HD 169142 is a very young Herbig Ae-Be star with a mass of 1.65\u20132 M\u2299 and an age of 5\u201311 Myr (Blondel & Djie 2006; Manoj et al. 2007) that is surrounded by a gas-rich disk (i = 13\u00b0; Raman et al. 2006; PA = 5\u00b0; Fedele et al. 2017) that is seen almost face-on. The parallax is 8.77 \u00b1 0.06 mas (Gaia DR2 2018). Disk structures dominate the inner regions around HD 169142 (see, e.g., Ligi et al. 2018). Figure 1 shows the view obtained from polarimetric observations: the left panel shows the Q\u03a6 image in the J band obtained by Pohl et al. (2017) using SPHERE on a linear scale, and the two rings are clearly visible. The right panel shows a pseudo-ADI image of the inner regions obtained by differentiating the Q\u03a6 image (see Ligi et al. 2018, for more details). Biller et al. (2014) and Reggiani et al. (2014) discussed the possible presence of a point source candidate at small separation (0.2 arcsec from the star). However, the analysis by Ligi et al. (2018) based on SPHERE data does not support or refute these claims; in particular, they suggested that the candidate identified by Biller et al. (2014) might be a disk feature rather than a planet. Polarimetricimages with the adaptive optics system NACO at the Very Large Telescope (VLT; Quanz et al. 2013b), SPHERE (Pohl et al. 2017; Bertrang et al. 2018) and GPI (Monnier et al. 2017) show a gap at around 36 au, with an outer ring at a separation >40 au from the star. This agrees very well with the position of the rings obtained from ALMA data (Fedele et al. 2017); similar results were obtained from VLA data (Osorio et al. 2014; Mac\u00edas et al. 2017). We summarize this information about the disk structure in Table 1 and call the ring at 0.17\u20130.28 arcsec from the star Ring 1 and the ring at 0.48\u20130.64 arcsec Ring 2. We remark that in addition to these two rings, both the spectral energy distribution (Wagner et al. 2015) and interferometric observations (Lazareff et al. 2017; Chen et al. 2018) show an inner disk at a separation smaller than 3 au. This inner disk isunresolved from the star in high-contrast images and consistent with ongoing accretion from it onto the young central star.While the cavities between the rings seem devoid of small dust, some gas is present there (Osorio et al. 2014; Mac\u00edas et al. 2017; Fedele et al. 2017). Fedele et al. (2017) and Bertrang et al. (2018) have suggested the possibility that the gap between Rings 1 and 2 is caused by a planet with a mass slightly higher than that of Jupiter. However, this planet has not yet been observed, possibly because it is at the limit of or beyond current capabilities of high-contrast imagers. On the other hand, Bertrang et al. (2018) found a radial gap in Ring 1 at PA ~ 50\u00b0 that might correspond to a similar radial gap found by Quanz et al. (2013b) at PA ~ 80\u00b0. The authors noted that if this correspondence were real, then this gap might be caused by a planet at about 0.14 arcsec from the star. So far, this planet has not been unambiguously detected either.","Citation Text":["Pohl et al. (2017)"],"Functions Text":["Figure 1 shows the view obtained from polarimetric observations: the left panel shows the Q\u03a6 image in the J band obtained by","using SPHERE on a linear scale, and the two rings are clearly visible."],"Functions Label":["Uses","Uses"],"Citation Start End":[[523,541]],"Functions Start End":[[398,522],[542,612]]}
{"Identifier":"2022AandA...666A.190S___2014b_Instance_1","Paragraph":"For our dataset of absolute magnitudes, we used data collected at the Institute of Astronomy of V. N. Karazin Kharkiv National University within the long-term observational programme to study asteroid magnitude-phase curves (Shevchenko et al. 2010, 2012, 2014a, 2016; Slyusarev et al. 2012). We also used some observational data obtained within several other programmes (Belskaya et al. 2010; Chiorny et al. 2007, 2011; Dotto et al. 2009; Hahn et al. 1989; Kaasalainen et al. 2004; Lagerkvist et al. 1998; Michalowski et al. 1995; Mohamed et al. 1994, 1995; Oszkiewicz et al. 2021; Shevchenko et al. 1992, 2003, 2009, 2014b, 2021; Velichko et al. 1995; Wilawer et al. 2022). All magnitudes were measured in the Johnson V band and extrapolated to zero phase angle using the HG1G2 system proposed by Muinonen et al. (2010), with some modifications presented by Penttil\u00e4 et al. (2016). For computations, the online calculator1 for the HG1G2 photometric system was used. Since we derived absolute magnitudes in our data from the light curve maxima, and the definition of H is based on the rotationally averaged brightness, we added a half of the light curve amplitude corrected to zero phase angle to our results. We used the average correction coefficients from Zappala et al. (1990) for low- and moderate-albedo asteroids. This correction is typically very small because our light curve observations covered small phase angles. Absolute magnitudes obtained at different aspects were averaged. In such a manner, we obtained a homogeneous dataset of absolute magnitudes of about 400 asteroids up to H = 16.5 mag. Our database includes the absolute magnitude data, the G1 and G2 parameters, and the albedo and diameter values from different databases (such as Tedesco et al. 2002; Masiero et al. 2011, 2012; Nugent et al. 2015; Usui et al. 2011). The database is available at the CDS. Figure 1 shows the correlations of the absolute magnitudes from the largest datasets (MPC (HMPC), Pan-STARRS (HPS), and ATLAS (HATLAS)) with those of the Kharkiv dataset (HKH). For the ATLAS dataset, we used the absolute magnitudes in a cyan filter, since this filter overlaps the Johnson V band (Mahlke et al. 2021).","Citation Text":["Shevchenko et al.","2014b"],"Functions Text":["We also used some observational data obtained within several other programmes"],"Functions Label":["Uses"],"Citation Start End":[[582,599],[618,623]],"Functions Start End":[[292,369]]}
{"Identifier":"2017AandA...598A..21B__Formicola_et_al._(2004)_Instance_1","Paragraph":"In this section, we describe a few numerical experiments carried out to analyse the importance of various hypotheses used to compute structural kernels. All models were computed using the Cl\u00e9s stellar evolution code (Scuflaire et al. 2008b) with the following ingredients: the CEFF equation of state (Christensen-Dalsgaard & Daeppen 1992), the OPAL opacities from Iglesias & Rogers (1996), supplemented at low temperature by the opacities of Ferguson et al. (2005) and the effects of conductivity from Potekhin et al. (1999) and Cassisi et al. (2007). The nuclear reaction rates are those from the NACRE project (Angulo et al. 1999), supplemented by the updated reaction rate from Formicola et al. (2004) and convection was implemented using the classical, local mixing-length theory (B\u00f6hm-Vitense 1958). We also used the implementation of microscopic diffusion from Thoul et al. (1994), for which three groups of elements are considered and treated separately: hydrogen, helium and the metals (all considered to have diffusion speeds of 56Fe). The oscillation frequencies and eigenfunctions were computed using the Li\u00e8ge adiabatic oscillation code (Scuflaire et al. 2008a). We took much care to analyse the numerical quality of the eigenfunctions and the models before computing structural kernels. Irregularities and poor quality of the computed eigenfunctions can bias the results and lead to wrong structural kernels and thus wrong inferences from inverted results. From our experience in hare-and-hounds exercises and inversions, we have determined that adding seismic constraints to the model is very efficient at bringing the reference model into the linear regime thus validating the inversion process. In other words, fitting the average large and small frequency separations is already a big improvement in terms of linearity, although individual seismic constraints, such as individual frequency ratios and individual small frequency separations are the best way to maximise the chances of being in the linear regime. Individual large frequency separations can also be used, but due to their sensitivity to surface effects, they should not be used in observed cases. As such, since in this study we did not use very elaborate seismic fitting techniques, our tests serve the only purpose of isolating various contributions to the errors and to test various hypotheses usually done when carrying out structural inversions in the context of helio- and asteroseismology. ","Citation Text":["Formicola et al. (2004)"],"Functions Text":["The nuclear reaction rates are those from the NACRE project",", supplemented by the updated reaction rate from"],"Functions Label":["Uses","Uses"],"Citation Start End":[[681,704]],"Functions Start End":[[552,611],[632,680]]}
{"Identifier":"2017ApJ...835...25E__Rutten_1984_Instance_2","Paragraph":"We compare our results with a new reduction of observations from the Lowell Observatory SSS, which is running a long-term stellar activity survey complementary to the MWO HK Project. The SSS observes solar and stellar light with the same spectrograph, with the solar telescope consisting of an exposed optical fiber that observes the Sun as an unresolved source (Hall & Lockwood 1995; Hall et al. 2007). The basic measurement of SSS is the integrated flux in 1 \u212b bandpasses centered on the Ca ii H & K cores from continuum-normalized spectra, \u03d5HK, which can then be transformed to the S-index using a combination of empirical relationships derived from stellar observations:\n7\n\n\n\n\n\nwhere \n\n\n\n\n\n is the continuum flux scale for the Ca ii H & K wavelength region, which converts \u03d5HK to physical flux (erg cm\u22122 s\u22121). \n\n\n\n\n\n is a function of Str\u00f6mgren \n\n\n\n\n\n and is taken from Hall (1996). \n\n\n\n\n\n (simply K in other works) is the conversion factor from the MWO HKP-2 H & K flux (numerator of Equation (1)) to physical flux (Rutten 1984). Ccf is a factor that removes the color term from S and is a function of Johnson \n\n\n\n\n\n (Rutten 1984). Finally, Teff is the effective temperature. See Hall et al. (2007) and Hall & Lockwood (1995) for details on the extensive work leading to this formulation. What is important to realize about this method of obtaining S is that it requires three measurements of solar properties, \n\n\n\n\n\n, \n\n\n\n\n\n, and \n\n\n\n\n\n, along with the determination of one constant, \n\n\n\n\n\n. The solar properties are taken from best estimates in the literature, which vary widely depending on the source used, and can dramatically affect the resulting SSSS for the Sun. Hall et al. (2007) used \n\n\n\n\n\n, \n\n\n\n\n\n, and \n\n\n\n\n\n. The constant \n\n\n\n\n\n was empirically determined to be 0.97 \u00b1 0.11 erg cm\u22122 s\u22121 in Hall et al. (2007) as the value that provides the best agreement between SSSS and \n\n\n\n\n\n from Baliunas et al. (1995) for an ensemble of stars and the Sun. This combination of parameters resulted in a mean SSSS of 0.170 for the Sun using observations covering cycle 23. A slightly different calibration of SSS data in Hall & Lockwood (2004) used a flux scale \n\n\n\n\n\n based on Johnson \n\n\n\n\n\n, set to 0.65 for the Sun, and \n\n\n\n\n\n K. In Table 1 we estimated that this calibration resulted in a mean S = 0.168 for cycle 23. Hall et al. (2009), which included a revised reduction procedure and one year of data with the upgraded camera (see below), found \n\n\n\n\n\n.","Citation Text":["Rutten 1984"],"Functions Text":["Ccf is a factor that removes the color term from S and is a function of Johnson"],"Functions Label":["Uses"],"Citation Start End":[[1122,1133]],"Functions Start End":[[1034,1113]]}
{"Identifier":"2016AandA...587A.133G__Bruzual_&_Charlot_(2003)_Instance_1","Paragraph":"Notes. Intrinsic L\u03bd(1400)\/L\u03bd(900) ratios obtained from different stellar population models: the BC03, assuming different initial mass functions (Salpeter, Salp, and Chabrier, Chab), star-formation histories (constant, cSFR, exponentially declining with \u03c4 = 0.1Gyr, exponentially rising with \u03c4 = \u22120.1Gyr), metallicties (Z\u2299 and 0.02 \u00d7 Z\u2299), and stellar-population ages; the BPASS (Binary Population and Spectral Synthesis code Stanway et al. 2016); the Starburst99 models. We consider a range of wavelengths around 900 \u00c5 and 1400 \u00c5, comparable to that covered by the narrow-band and the HST\/ACS F606W filters in the rest frame, to estimate the average L\u03bd. As described in Bruzual & Charlot (2003), BC03 provides the spectral energy distribution of stars obtained from a comprehensive library of theoretical model atmospheres (KBFA in the table). The library consists of Kurucz (1996) spectra for O-K stars, Bessell et al. (1991) and Fluks et al. (1994) spectra for M giants, and Allard & Hauschildt (1995) spectra for M dwarfs. In BPASS the stellar evolutionary tracks contain a contribution from isolated stars and binary systems. Also, the stellar atmosphere models are selected from the BaSeL v3.1 library (Westera et al. 2002), supplemented by Wolf-Rayet stellar atmosphere models from the Potsdam PoWR group (Hamann & Gr\u00e4fener 2003). With Starburst99, it is possible to generate SEDs assuming a bunch of stellar evolutionary tracks, including stars with and without rotation (GenevaV40 and GenevaV00 respectively, Leitherer et al. 2014), and stellar atmospheres (the combination of model atmosphere from Pauldrach et al. 1998; Hillier & Miller 1998, is the recommended option). The 70%ROT+30%noROT entry indicates a model that is a combination of GenevaV40 for the 70% and GenevaV00 for the 30% (Levesque et al. 2012). The GenevaV40 and GenevaV00 tracks were released for metallicity equal to Z = 0.014 (~Z\u2299, Eldridge 2012). The change in the intrinsic ratio due to the different evolutionary tracks and stellar atmospheres is shown for a Salpeter IMF, Z\u2299, cSFR, and 100 Myr old stellar population. In the bottom part of the table, we report the ratios calculated from the best fit templates of AGN (Bongiorno et al. 2012), in which the SEDs are dominated by the radiation coming from an (un-)obscured AGN. (1) The intrinsic L\u03bd(1400)\/L\u03bd(900) ratio for an unobscured TypeII AGN can be as low as 2.35 (Richards et al. 2006). In the estimation of the ratios we take into account the HI absorption occurred in star atmospheres, but we neglect that within the interstellar and circum-galactic medium. ","Citation Text":["Bruzual & Charlot (2003)"],"Functions Text":["Intrinsic L\u03bd(1400)\/L\u03bd(900) ratios obtained from different stellar population models: the BC03, assuming different initial mass functions (Salpeter, Salp, and Chabrier, Chab), star-formation histories (constant, cSFR, exponentially declining with \u03c4 = 0.1Gyr, exponentially rising with \u03c4 = \u22120.1Gyr), metallicties (Z\u2299 and 0.02 \u00d7 Z\u2299), and stellar-population ages;","As described in","BC03 provides the spectral energy distribution of stars obtained from a comprehensive library of theoretical model atmospheres (KBFA in the table).","The library consists of Kurucz (1996) spectra for O-K stars, Bessell et al. (1991) and Fluks et al. (1994) spectra for M giants, and Allard & Hauschildt (1995) spectra for M dwarfs."],"Functions Label":["Uses","Uses","Uses","Background"],"Citation Start End":[[669,693]],"Functions Start End":[[7,366],[653,668],[695,842],[843,1024]]}
{"Identifier":"2021AandA...656A..64R__Ruiz-Lara_et_al._2020_Instance_1","Paragraph":"Finally, we confirm our hypothesis on the origin of the bi-modal A(Li) distribution, and of the isthmus, by comparing results from our GCE model with data. In Fig. 8 we show the normalised A(Li) distribution of all the stars in our sample (blue histogram) compared with the result from three models of GCE: (i) imposing a gently declining SFH (SFR_0b model, dashed red lines); (ii) a gently declining SFH plus a single star formation burst at a \u223c6 Gyr look-back time (SFR_1b model, dashed blue lines); and (iii) our best model (SFR_2b, black-solid lines), which contains a gently declining SFR plus two star formation bursts at \u223c6 Gyr and \u223c1.5 Gyr look-back times, the latter being the most impulsive. In our best model the first star formation peak is only 68% as intense as the latter and lasts 25% longer. The SFR_2b model is the one that best reproduces the observed A(Li) distribution while recovering the presently observed SFR and the thin-disk stellar and HI mass (Mor et al. 2019; Isern 2019; Ruiz-Lara et al. 2020). This result is robust as it only depends on the observed A(Li) distribution. Possible biases from the determination of the stellar ages do not have an impact on our conclusions. Figure 8 clearly shows that only a two-peaked SFH can explain the observed bi-modal A(Li) distribution when accounting for most known 7Li production and depletion mechanisms (solid black line). It is important to remark that we cannot quantify the exact intensity and length of the star formation bursts from our data and models, but we can only give a qualitative analysis. This is a consequence of: (i) the large age determination uncertainties still present in current data; (ii) the fact that the data we use is restricted to the solar neigbourhood and thus do not represent the entire Galaxy; and finally (iii) the fact that our model still misses some 7Li production and destruction processes that are currently under study (e.g. the effect of planet engulfment). In the future, new data from large surveys such as Galah (Buder et al. 2021) combined with better determinations of A(Li) in 7Li-poor low-luminosity cold dwarf stars will result in large, unbiased samples of MS stars. These new catalogues will allow us to better understand the many channels of 7Li production and destruction that occur inside dwarf stars. Also, future catalogues that include precise age determinations will allow us to better constrain the SFH of the MW, also by analysing its A(Li) distribution.","Citation Text":["Ruiz-Lara et al. 2020"],"Functions Text":["The SFR_2b model is the one that best reproduces the observed A(Li) distribution while recovering the presently observed SFR and the thin-disk stellar and HI mass"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1002,1023]],"Functions Start End":[[809,971]]}
{"Identifier":"2016MNRAS.457.2236K__Ouyed_&_Pudritz_1997_Instance_1","Paragraph":"Strong outflows are commonly associated with the early stages of stellar evolution. They are likely responsible for transporting excess angular momentum away from the star\u2013disc system and regulating the mass-accretion process and spin evolution of newly born stars (e.g. Hartmann & Stauffer 1989; Matt & Pudritz 2005; Bouvier et al. 2014). There are at least three possible types of outflows around young stellar objects: (1) a disc wind launched from the accretion disc (e.g. Blandford & Payne 1982; Ustyugova et al. 1995, 1999; Romanova et al. 1997; Ouyed & Pudritz 1997; K\u00f6nigl & Pudritz 2000; Pudritz et al. 2007), (2) an X-wind or a conical wind launched near the disc\u2013magnetosphere interaction region (e.g. Shu et al. 1994, 1995; Romanova et al. 2009) and (3) a stellar wind launched from open magnetic field lines anchored to the stellar surface (e.g. Decampli 1981; Hartmann & MacGregor 1982; Kwan & Tademaru 1988; Hirose et al. 1997; Strafella et al. 1998; Matt & Pudritz 2005; Romanova et al. 2005; Cranmer 2009). However, despite recent efforts, the exact launching mechanisms of the winds and outflows as well as the mechanism behind the collimation of the ejected gas into jets are still not well understood (e.g. Edwards et al. 2006; Ferreira, Dougados & Cabrit 2006; Kwan, Edwards & Fischer 2007; Tatulli et al. 2007b; Kraus et al. 2008; Eisner, Hillenbrand & Stone 2014). Hence, high-resolution interferometric observations which can resolve the wind-launching regions are crucial for addressing this issue. If a high spectral resolution is combined with a high spatial resolution, the emission-line regions near the base of the wind (e.g. Br\u03b3 emission regions in Herbig Ae\/Be stars) can be resolved in many spectral channels across the line (e.g. Weigelt et al. 2011; Kraus et al. 2012b,c; Garcia Lopez et al. 2015). This allows us to study the wavelength-dependent extent and the kinematics of the winds, which can be derived from the line visibilities and wavelength-dependent differential and closure phases. Such observations would help us to distinguish the different types of outflow scenarios.","Citation Text":["Ouyed & Pudritz 1997"],"Functions Text":["There are at least three possible types of outflows around young stellar objects: (1) a disc wind launched from the accretion disc (e.g."],"Functions Label":["Background"],"Citation Start End":[[552,572]],"Functions Start End":[[340,476]]}
{"Identifier":"2021AandA...647A.177D__Awad_&_Abu-Shady_2020_Instance_1","Paragraph":"An estimate of a sputtering cross-section can be inferred from our measurements with \u03c3s \u2248 V\u2215d, where V is the volume occupied by \n\n${Y}_{\\textrm{s}}^{\\infty}$Ys\u221e\n molecules and d the depth of sputtering. \n\n$\\sigma_{\\textrm{s}}\\approx {Y}_{\\textrm{s}}^{\\infty}\/l_{\\textrm{d}}\/\\textrm{ml}$\u03c3s\u2248Ys\u221e\/ld\/ml\n, where ml is the number of CO or CO2 molecules cm\u22122 in a monolayer (about 6.7 \u00d7 1014 cm\u22122 and 5.7 \u00d7 1014 cm\u22122, respectively,with the adopted ice densities). As is shown in Table 1, the sputtering radius rs would therefore be about 1.26 to 2.12 times larger than the radiolysis destruction radius rd in the case of the CO2 ice, and 2.03 to 2.36 for CO in the considered energy range (~0.5-1 MeV\/u). The net radiolysis is the combined effect of the direct primary excitations and ionisations, the core of the energy deposition by the ion, and the so-called delta rays (energetic secondary electrons) travelling at much larger distances from the core; that is, several hundreds of nanometres at the considered energies in this work (e.g. Mozunder et al. 1968; Magee & Chatterjee 1980; Katz et al. 1990; Moribayashi 2014; Awad & Abu-Shady 2020). The effective radiolysis track radius that we calculate is lower than the sputtering one, which points towards a large fraction of the energy deposited in the core of the track. The scatter on the ratio of these radii is due to the lack of more precise data. It nevertheless allows to put a rough constraint on the estimate of Nd in the absence of additional depth measurements, with \n\n${N}_{\\textrm{d}} \\lesssim {Y}_{\\textrm{s}}^{\\infty}\/\\sigma_{\\textrm{d}}$Nd\u2272Ys\u221e\/\u03c3d\n. If the rs\u2215rd ratio is high, a large amount of species come from the thermal sublimation of an ice spot less affected by radiolysis, and the fraction of ejected intact molecules is higher. The aspect ratio corresponding to these experiments evolves between about ten and twenty for CO2 and CO, whereas for water ice at a stopping power of Se \u2248 3.6 \u00d7 103eV\u22151015 H2 O molecules cm\u22122, we show that it is closer to one (Dartois et al. 2018). The depth of sputtering is much larger for CO and CO2 than for H2O at the same energy deposition, not only because their sublimation rate is higher, but also because they do not form OH bonds. For complex organic molecules embedded in ice mantles dominated by a CO or CO2 ice matrix, with the lack of OH bonding and the sputtering for trace species being driven by that of the matrix (in the astrophysical context), the co-desorption of complex organic molecules present in low proportions with respect to CO\/CO2 cannot only be more efficient, but will thus arise from deeper layers.","Citation Text":["Awad & Abu-Shady 2020"],"Functions Text":["The net radiolysis is the combined effect of the direct primary excitations and ionisations, the core of the energy deposition by the ion, and the so-called delta rays (energetic secondary electrons) travelling at much larger distances from the core; that is, several hundreds of nanometres at the considered energies in this work"],"Functions Label":["Uses"],"Citation Start End":[[1119,1140]],"Functions Start End":[[699,1029]]}
{"Identifier":"2021MNRAS.502.3179T__Sreenivasan,_Ramshankar_&_Meneveau_1989_Instance_1","Paragraph":"These properties can also be related to the fractal nature of radiative mixing layers. Recently, Fielding et al. (2020) showed that the area of the cooling surface in radiative mixing layer simulations obeys a fractal scaling, with\n(31)$$\\begin{eqnarray*}\r\n\\frac{A_{\\rm T}}{A_{\\rm L}} = \\left(\\frac{\\lambda }{L} \\right)^{2-D},\r\n\\end{eqnarray*}$$where \u03bb is the smoothing scale and D = 2.5 was the fractal dimension argued to hold by analogy with well-known fractals, and verified in their simulations. Turbulence combustion fronts are indeed well known to be fractals, due to the dynamical self-similarity of turbulence in the inertial range. Experimental measurements by e.g. instantaneous laser tomography have given values ranging from D = 2.1\u22122.4 in a variety of flow geometries, with a preferred value of D = 2.35 (Hentschel & Procaccia 1984; Sreenivasan, Ramshankar & Meneveau 1989); it has been argued that this fractal dimension is universal (Catrakis, Aguirre & Ruiz-Plancarte 2002; Aguirre & Catrakis 2005). From equation (19), the fractal dimension can be used to calculate the turbulent flame speed (Gouldin et al. 1986; Peters 1988). The fractal scaling and consequent increase in area AT should extend all the way down to the Gibson scale \u03bbG, which is defined to be the scale where the turbulent velocity equals the laminar flame speed, v(\u03bbG) = SL. This is often unresolved in simulations. If we use the Kolmogorov scaling v\u221d\u03bb1\/3, then we obtain:\n(32)$$\\begin{eqnarray*}\r\n\\frac{S_{\\rm T}}{S_{\\rm L}} = \\frac{A_{\\rm T}}{A_{\\rm L}} = \\left(\\frac{\\lambda _{\\rm G}}{L} \\right)^{2-D} = \\left(\\frac{u^{\\prime } }{S_{\\rm L}} \\right)^{3(D-2)} ,\r\n\\end{eqnarray*}$$where we have used equation (31) and v(\u03bbG) = SL. Thus, in equation (20), we have n = 3(D \u2212 2). The experimental value of D = 2.35 gives n = 1.05, in good agreement with Damk\u00f6hler\u2019s scaling, and fair agreement with the scaling in equation (23). The Fielding et al. (2020) value of D = 2.5 gives n = 1.5, or ST = u\u2032(u\u2032\/SL)1\/2. If one uses the laminar $S_L \\propto t_{\\rm cool}^{-1\/2}$ from our static simulations, this would imply $S_T \\propto t_{\\rm cool}^{1\/4}$. However, in the Fielding et al. (2020) model, the speed at which a cooling layer advances is $S_L \\propto t_{\\rm cool}^{1\/2}$, so they end up with $S_T \\propto t_{\\rm cool}^{-1\/4}$ as well. The scalings are sensitive to the fractal dimension D and the measurement error on D obtained from the simulations is unclear at this point. In addition, the cutoff scale of turbulence may not be the Gibson scale. We caution that fractal arguments have not proven to be fully robust in the turbulent combustion context. For instance, the measured fractal parameters fluctuate depending on the extraction algorithm, and have not been able to correctly predict the turbulent burning velocity (Cintosun, Smallwood & G\u00fclder 2007).","Citation Text":["Sreenivasan, Ramshankar & Meneveau 1989"],"Functions Text":["Experimental measurements by e.g. instantaneous laser tomography have given values ranging from D = 2.1\u22122.4 in a variety of flow geometries, with a preferred value of D = 2.35"],"Functions Label":["Uses"],"Citation Start End":[[847,886]],"Functions Start End":[[642,817]]}
{"Identifier":"2020AandA...644A..88K__Shapovalova_et_al._2010b_Instance_1","Paragraph":"Type 1 AGN (NLSy1, quasars) and aligned CB-SMBH models expected signatures. Two SMBH and their BLRs induce more rich and complex differential phase patterns. There are many configurations for which the aligned CB-SMBH differs between them and single SMBH. The features of synthetic spectra given in Figs. C.2b and C.2c show a similarity with those found in 3C 390.3 (Landt et al. 2014, see their Fig. A1). This object is well-known as a double-peaked line emitter in the optical band. The double-peaked profiles can be also associated with accretion disc emission (Eracleous & Halpern 1994, 2003; Gezari et al. 2007). However, if the binary model is appropriate for this object, then associated differential phases will have a complex double S-shaped structure. This model reflect the possibility of a non-coplanar CB-SMBH, with highly inclined orbits of clouds in both BLRs. If a single SMBH model with a BLR d of 95 light days is true (see Shapovalova et al. 2010b) a differential signal would be 0.9\u00b0. Even that spectral lines of objects can share some characteristics, our model predicts that corresponding differential phases are distinct because of different SMBH and the clouds\u2019 orbital parameters. The optically bright quasar PG 1211+143 (Landt et al. 2014) has a convex Pa\u03b1 shape slightly depressed of the centre, as in our synthetic case present in Fig. C.2a. The predicted differential phase would resemble an asymmetric double S shape, caused by non-coplanarity of the CB-SMBH system and high values for the clouds\u2019 orbital elements \u03a9c and \u03c9c. Moreover, the high rise spectral line in Fig. 5a is also observed in the spectrum of SDSSJ032213.90+005513.4 (Kim et al. 2010, see their Fig. 1). If the CB-SMBH model is appropriate for this object, then the expected differential phase would be similar in morphology with the previous case but distinct in their details because of different values of \u03a92. In addition, asymmetric two horn feature (blue line) found in Fig. C.3b, is also highly consistent with observed line in Mrk 79 (Landt et al. 2008, see their Fig. 13). Anticipated differential phases differ from those seen in previous cases because of large values, 230\u00b0\u2005 \u03a91,\u2006\u03a92\u2004 \u2004330\u00b0\u2005, and decreasing \u03c9k,\u2006k\u2004=\u20041,\u20062.","Citation Text":["Shapovalova et al. 2010b"],"Functions Text":["If a single SMBH model with a BLR d of 95 light days is true (see","a differential signal would be 0.9\u00b0."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[942,966]],"Functions Start End":[[876,941],[968,1004]]}
{"Identifier":"2020ApJ...903...46C__Banerjee_et_al._2009_Instance_1","Paragraph":"The large values of \u03b1 (significantly larger than unity) observed in clouds of low column densities or masses are often interpreted in terms of the presence of a large external confining pressure (P\/\n\n\n\n\n\n \n\n\n\n\n\n e.g., Keto & Myers 1986; Oka et al. 2001; Field et al. 2011; Leroy et al. 2015; Traficante et al. 2018a), although it is hard to imagine that such high pressures can be thermal in general, since the mean ambient thermal pressure in the interstellar medium (ISM) is rather low, \n\n\n\n\n\n \n\n\n\n\n\n, and large deviations from it occur very infrequently (e.g., Boulares & Cox 1990; Jenkins 2004; Jenkins & Tripp 2011). Instead, it is most likely that these values correspond to ram pressure, in which case they imply mass, momentum, and energy flux across Eulerian cloud boundaries, or a displacement of Lagrangian boundaries (Ballesteros-Paredes et al. 1999; Banerjee et al. 2009). Indeed, in a previous study (Camacho et al. 2016), we find, through measurement of the mean velocity divergence within the clouds in numerical simulations of cloud formation and evolution, that roughly half the clouds with an excess of kinetic energy are undergoing compression. This can be interpreted as the clouds being subject to a ram pressure (which amounts to an inertial compression) that is making them denser and smaller, so that they eventually will become gravitationally bound. The origin of this ram pressure can be large-scale turbulence, a large-scale potential well, or other instabilities. Furthermore, one important possibility is that clouds may be falling into the potential well of a stellar spiral arm, which is the main source of large-scale compression for the gas in the Galactic disk (Roberts 1969). Thus, this is not really a \u201cconfinement,\u201d since the clouds are not at rest. The same goes for the other half of the clouds, which are undergoing expansion. In this case, the excess of kinetic energy corresponds to the expansion motions, and again the cloud is not confined, so there is no need for a high confining pressure. In Camacho et al. (2016) and Ballesteros-Paredes et al. (2016, hereafter BP+18) it has been suggested that, for clouds formed by inertial compressions in the background medium (Ballesteros-Paredes et al. 1999), and which gradually become more strongly gravitationally bound, while the inertial compressive motions decay or dissipate, the kinetic energy transits from being dominated by the inertial motions to being dominated by the gravitationally driven motions (see also Collins et al. 2012). In that case, an initial decay of the Larson ratio and the virial parameter may be expected.","Citation Text":["Banerjee et al. 2009"],"Functions Text":["Instead, it is most likely that these values correspond to ram pressure, in which case they imply mass, momentum, and energy flux across Eulerian cloud boundaries, or a displacement of Lagrangian boundaries"],"Functions Label":["Uses"],"Citation Start End":[[863,883]],"Functions Start End":[[622,828]]}
{"Identifier":"2019MNRAS.488..803M__Long_2017_Instance_1","Paragraph":"The identification and confirmation of extragalactic SNRs are mainly done using data from the radio (Lacey, Duric & Goss 1997; Hyman et al. 2001; Lacey & Duric 2001), visible (Matonick & Fesen 1997; Matonick et al. 1997; Gordon et al. 1998; Blair & Long 2004; Sonbas, Akyuz & Balman 2009; Sonba\u015f et al. 2010), and X-Ray (Pence et al. 2001; Ghavamian et al. 2005) wavelength ranges. Only a few extragalactic SNRs have been identified in the infrared using the [Fe\u2009ii]\u20091.64\u2009$\\hbox{$\\mu $m}$ emission line (Greenhouse et al. 1997). The radiation of SNRs in different wavelength ranges is under the influence of biases as they cover different aspects of the ISM environment and SNR age and evolution (Pannuti et al. 2000; Leonidaki, Zezas & Boumis 2010; S\u00e1nchez-Ayaso et al. 2012; Long 2017). In the radio, only SNR candidates associated with H\u2009ii regions are identified, which means that radio SNR samples are biased towards star-forming regions. In the X-Ray, candidates are selected if they display a soft spectrum and if they are associated with an H\u2009ii region, which means that X-Ray samples have the same bias than the radio samples and are also biased against SNR candidates with hard spectra and no optical counterparts. In the optical, samples are under the influence of biases privileging the identification of SNRs located in low-density environments. A multiwavelength (X-Ray, optical, and radio) study of SNRs in NGC\u2009300 by Pannuti et al. (2000) revealed 16 new SNRs, 2 in the radio, and 14 in the X-Ray, in addition to the 28 SNRs previously identified in the optical. The lack of new optical detection is explained by the fact that optical SNRs can only be detected when they represent relatively low confusion with other H\u2009\u03b1 emission sources. The optical SNRs found here are generally located well away from star-forming regions. Consequently, SNR samples identified optically are often not complete. Another technique, based on the search for Large-Velocity-Width Sources (LVWS), is used to identify SNRs with optical spectroscopy (Chu & Kennicutt 1986; Castaneda, Vilchez & Copetti 1990). While H\u2009ii regions show low-velocity dispersion (\u03c3v \u2264 30\u2009km\u2009s\u22121; Melnick 1977; Gallagher & Hunter 1983), broad emission line widths observed in the LVWS can be caused by stellar winds or SN. Using this technique, Chu & Kennicutt (1986) discovered four LVWS inside the Giant H\u2009ii regions NGC\u20095471 A, B, C, and NGC\u20095461 in the nearby galaxy M101.","Citation Text":["Long 2017"],"Functions Text":["The radiation of SNRs in different wavelength ranges is under the influence of biases as they cover different aspects of the ISM environment and SNR age and evolution"],"Functions Label":["Background"],"Citation Start End":[[777,786]],"Functions Start End":[[529,695]]}
{"Identifier":"2018AandA...610A..10C__Montaigne_et_al._2005_Instance_1","Paragraph":"According to chemical models, the molecules HCS+ and CS are among the most closely related species present in the ISM, as they participate in a direct exchange in both the formation and destruction of CS (Drdla et al. 1989; Lucas & Liszt 2002). In the reaction network considered by Drdla et al. (1989) and Lucas & Liszt (2002), CS is believed to be formed primarily from the dissociative recombination reaction of (6)\\begin{equation} \\label{eq:HCSdisRecomb} \\text{HCS}^+ + {\\rm e}^- \\rightarrow \\text{CS} + \\text{H}, \\end{equation}HCS++e\u2212\u2192CS+H,where HCS+ is first formed by reactions beginning with S+. However, recent experiments indicate that this product channel occurs in only 19 percent of collisions, with 81 percent of collisions forming CH + S (Montaigne et al. 2005). Three mechanisms dominate the destruction of CS in these networks. In the first two, photoionization and ion-exchange reactions destroy CS to form CS+, by (7)\\begin{equation} \\label{eq:CSphotoIo} \\text{CS} + \\gamma \\rightarrow \\text{CS}^+ + {\\rm e}^- , \\end{equation}CS+\u03b3\u2192CS++e\u2212,and (8)\\begin{equation} \\label{eq:CSionExchange} \\text{CS} + X^+ \\rightarrow \\text{CS}^+ + X, \\end{equation}CS+X+\u2192CS++X,where X+ is a cationic species and X is the corresponding neutral species. CS+ then quickly reacts to form HCS+: (9)\\begin{equation} \\label{eq:CSio_recomb} \\text{CS}^+ + \\text{H}_2 \\rightarrow \\text{HCS}^+ + \\text{H} . \\end{equation}CS++H2\u2192HCS++H.In the third destruction route, CS reacts with H\\hbox{$_3^+$}+3 to form HCS+ directly: (10)\\begin{equation} \\label{eq:CSandH3plus} \\text{CS} + \\text{H}_3^+ \\rightarrow \\text{HCS}^+ + \\text{H}_2 . \\end{equation}CS+H3+\u2192HCS++H2.As CS is formed by a reaction of HCS+ and the dominant destruction pathways for CS form HCS+, we would expect the abundances of CS and HCS+ to be delicately balanced, and for the two abundances to track each other. While chemical models of sulfur-bearing chemistry in diffuse clouds underpredict the abundance of HCS+ by multiple orders of magnitude, Lucas & Liszt (2002) determined that if the observed abundance of HCS+ is injected into a diffuse cloud, it is possible to account for the CS abundances observed in their sample, further emphasizing the theoretical importance of NCS\/NHCS+. In clouds in the Galactic center, where H\\hbox{$_3^+$}+3 is \u227310 times more abundant compared to diffuse clouds in the disk (Oka et al. 2005), we might expect an offset in the ratio, tending toward a higher relative abundance of HCS+. ","Citation Text":["Montaigne et al. 2005"],"Functions Text":["However, recent experiments indicate that this product channel occurs in only 19 percent of collisions, with 81 percent of collisions forming CH + S"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[754,775]],"Functions Start End":[[604,752]]}
{"Identifier":"2020ApJ...904..117K___2014b_Instance_1","Paragraph":"Simultaneously solving the MHD equations and the global angle- and energy-dependent radiation transport equation, in general relativity, is both computationally expensive (typically prohibitively so) and technically challenging. Even so, significant progress has been made in the last decade, though the problem is usually made tractable by introducing at least one of the following simplifying assumptions: abandoning general relativity in favor of a pseudo-Newtonian description of the gravitational potential while performing realistic, multi-angle group radiation transport (Jiang et al. 2014a, 2014b, 2019a, 2019b); limiting the possible angular dependence of the radiation field by invoking either flux-limited diffusion (Zanotti et al. 2011; Roedig et al. 2012) or, more recently, the \u201cM1 closure\u201d relation in either axisymmetric (2D; Sadowski et al. 2014) or 3D simulations (Fragile et al. 2012, 2014; McKinney et al. 2014; Sadowski et al. 2016); or using Monte Carlo (Ryan et al. 2015) \/ hybrid Monte Carlo techniques (Ryan & Dolence 2020). Most attempts have eschewed energy-dependent transfer in favor of a \u201cgray\u201d atmosphere\u2014the radiation field is treated as monochromatic, coupled to the fluid only through the Rosseland mean opacity (Rybicki & Lightman 1986). The first of these approximations, the pseudo-Newtonian potential, is especially problematic in regions close to the black hole, where general relativistic effects play a critical role in determining both the structure of the accretion flow and the photon trajectories. The others are essentially variants of a diffusion approximation and are best suited to the cooler, denser, and optically thick body of the accretion disk, where the environment is similar to those found in stellar atmospheres\u2014the field from which these methods, and gray transfer, originate (Chandrasekhar 1960). With the exception of Monte Carlo methods (Ryan et al. 2015; Ryan & Dolence 2020), these are all especially poorly suited to the diffuse, hot, optically thin corona, especially at small radii near the black hole.","Citation Text":["Jiang et al.","2014b"],"Functions Text":["Even so, significant progress has been made in the last decade, though the problem is usually made tractable by introducing at least one of the following simplifying assumptions: abandoning general relativity in favor of a pseudo-Newtonian description of the gravitational potential while performing realistic, multi-angle group radiation transport"],"Functions Label":["Background"],"Citation Start End":[[579,591],[599,604]],"Functions Start End":[[229,577]]}
{"Identifier":"2022MNRAS.513.5377F__Blum_et_al._2017_Instance_2","Paragraph":"At each heliocentric distance rh, the activity model (Fulle et al. 2020b) is defined by five analytical equations fixing (i) the gas pressure P(s) depending on the depth s from the nucleus surface (Fig. 1 for the CO2 case), (ii) the gas flux Q from the nucleus surface, (iii) the temperature gradient \u2207T at depths of a few cm, (iv) the heat conductivity \u03bbs at depths of a few cm below the nucleus surface, and (v) the temperature Ts of the nucleus surface\n(3)$$\\begin{eqnarray*}\r\nP(s) = P_0 ~f(s) ~\\exp \\left[{- {T_0 \\over {T_s - s ~\\nabla T}}}\\right]\r\n\\end{eqnarray*}$$(4)$$\\begin{eqnarray*}\r\nQ = {{14 ~r ~P(R)} \\over {3 ~R}} \\sqrt{{2 ~m} \\over {\\pi k_B ~(T_s - R ~\\nabla T)}}\r\n\\end{eqnarray*}$$(5)$$\\begin{eqnarray*}\r\n\\nabla T = {\\sqrt{\\Lambda ~Q ~\/ ~\\sigma _B} \\over {8 ~(T_s - R ~\\nabla T) ~R}}\r\n\\end{eqnarray*}$$(6)$$\\begin{eqnarray*}\r\n\\lambda _s = {32 \\over 3} ~(T_s - R ~\\nabla T)^3 ~\\sigma _B ~R\r\n\\end{eqnarray*}$$(7)$$\\begin{eqnarray*}\r\n(1 - A) ~I_\\odot ~\\cos \\theta ~r_h^{-2} = \\epsilon \\sigma _B T_s^4 + \\lambda _s ~\\nabla T + \\Lambda ~Q\r\n,\r\n\\end{eqnarray*}$$where P0, T0, and \u039b values are listed in Table 1, s is the depth from the nucleus surface, $f(s) = 1 - (1 - {s \\over R})^4$ for s \u2264 R, f(s) = 1 elsewhere, r \u2248 50 nm and R \u2248 5 mm are the radii of the grains of which cometary dust consists (Levasseur-Regourd et al. 2018; G\u00fcttler et al. 2019; Mannel et al. 2019) and of the pebbles of which cometary nuclei consist (Blum et al. 2017; Fulle et al. 2020b), m is the mass of the gas molecule, kB is the Boltzmann constant, \u03c3B is the Stefan\u2013Boltzmann constant, A is the nucleus Bond albedo (e.g. A = 1.2 per cent measured at 67P; Fornasier et al. 2015), I\u2299 is the solar flux at the heliocentric distance of Earth, \u03b8 is the solar zenithal angle, and \u03f5 \u2248 0.9 is the nucleus emissivity. Since the gas originates from the superficial pebbles and is assumed to share the temperature Ts \u2212 s \u2207T of refractories and ices, the thermal diffusion due to gas convection is negligible with respect to the sublimation sink \u039b Q. A nucleus is active if the gas pressure overcomes the tensile strength S bonding dust particles to the nucleus surface (Skorov & Blum 2012), thus defining the activity onset for each ice (Table 2), occurring (i) at rh = 85 au for carbon monoxide (Fulle et al. 2020a); (ii) at rh = 60 au for molecular oxygen; (iii) at rh = 52 au for methane; (iv) at rh = 18 au for ethane; (v) at rh = 13 au for carbon dioxide (dotted line in Fig. 1); and (vi) at rh = 3.8 au for water (Fulle et al. 2020b; Ciarniello et al. 2021). The value R \u2248 5 mm has been constrained by several data collected at comet 67P, by laboratory experiments of dust accretion in conditions expected to occur in the solar protoplanetary disc and by observations of other protoplanetary discs (Blum et al. 2017). Other R-values would not provide the best fit of the 67P water-loss time-evolution (Ciarniello et al. 2021).","Citation Text":["Blum et al. 2017"],"Functions Text":["The value R \u2248 5 mm has been constrained by several data collected at comet 67P, by laboratory experiments of dust accretion in conditions expected to occur in the solar protoplanetary disc and by observations of other protoplanetary discs"],"Functions Label":["Uses"],"Citation Start End":[[2783,2799]],"Functions Start End":[[2543,2781]]}
{"Identifier":"2015ApJ...805..105C__Li_1995_Instance_1","Paragraph":"As we said in the previous section, the evolution of the magnetic fields in the two domains (the disk and the magnetosphere) are related. It is important to note here that there is an important complication that enters into the matching between the two domains and this has to do with the abrupt transition between the dense turbulent accreting flow and the almost empty ionized (probably also evaporating and outflowing) magnetosphere. This is very similar to the well known transition region in the Sun where, within a very thin layer of about 100 km, the temperature rises by almost two orders of magnitude to about one million degrees, and the matter density drops by about one order of magnitude. The solar transition region is the subject of ongoing intense theoretical and numerical investigations. The astrophysical disk transition region is at least equally complicated since it is not only the base of the disk corona, but also the origin of purported disk winds and outflows (e.g., Ferreira & Pelletier 1995; Li 1995). It is not the aim of the present study to solve in detail for the structure of the flow and magnetic field in this region. In practice, we define a transition region of two \u03b8-grid zones above \n\n\n\n\n\n. In that region we set\n14\n\n\n\n\n\nThis matching allows for magnetic field loops that reach the disk surface zones to escape to the magnetosphere. We also set the field lines in that region in rotation by introducing a magnetospheric poloidal electric field\n15\n\n\n\n\n\nFinally, we slightly modify the induction equation that we solve for the disk material in that region by introducing an extra phenomenological term that mimics various complex physical effects taking place in the surface layers of astrophysical accretion disks such as surface convection, buoyancy of magnetic field loops, and the fact that the surface layers are highly ionized (thus also fully conducting) due to cosmic ray irradiation, namely\n16\n\n\n\n\n\nwhere, \n\n\n\n\n\n is the disk magnetic field in the transition layer. The introduction of this transition layer effectively shields the disk interior: (a) it prevents any generated magnetospheric field from \u201centering\u201d the disk from above (especially in the case of turbulent high \u03b7 flow conditions), and (b) it \u201cpushes outward\u201d (in the \u03b8 direction) at a very small fraction of the speed of light any magnetic field loop that enters it from below. We implemented different fractions (\n\n\n\n\n\n, \n\n\n\n\n\n, etc.) and the results were qualitatively similar. As we said above, we are not claiming that we study in detail the disk-magnetosphere transition region. This is the reason we opted for a minimal two \u03b8-zone layer that only serves the purpose of shielding the disk interior from the disk magnetosphere. Note that Parfrey et al. (2015) ignored the important physical significance of a transition layer by directly coupling their magnetosphere to a prescribed distribution of the magnetic field on the surface of the accretion disk without solving for the magnetic field distribution in the disk interior.","Citation Text":["Li 1995"],"Functions Text":["The astrophysical disk transition region is at least equally complicated since it is not only the base of the disk corona, but also the origin of purported disk winds and outflows (e.g.,"],"Functions Label":["Motivation"],"Citation Start End":[[1020,1027]],"Functions Start End":[[806,992]]}
{"Identifier":"2019ApJ...871..257P__Gabuzda_et_al._2018_Instance_1","Paragraph":"We note that it is unlikely that poloidal magnetic fields are responsible for the observed RMs of M87 because in that case one expects \u03c1 \u221d r0 from B \u221d r\u22122, which is impossible to explain with the accretion models currently available (Yuan & Narayan 2014). However, there is indication of non-negligible poloidal fields as well as toroidal fields\u2014resulting in helical magnetic fields\u2014in the jet environment of other distant AGNs, which results in transverse RM gradients with no sign changes (e.g., Asada et al. 2002; Zamaninasab et al. 2013; G\u00f3mez et al. 2016; Gabuzda et al. 2018, see also Section 3.5.2). The existence of non-negligible poloidal fields was indicated even for the M87 jet at HST-1 from the observed moving knots with both fast and slow velocities, which could be explained by quad relativistic MHD shocks in a helical magnetic field permeating the jet (Nakamura et al. 2010; Nakamura & Meier 2014). In Sections 3.5.2 and 4.1, we explained that poloidal magnetic fields might be very weak at distances \u22735000 rs probed in this study and we concluded that hot accretion flows and winds are more probable to be the Faraday screen than the jet sheath. However, if the jet experiences recollimation, which may lead to formation of standing shocks (e.g., Daly & Marscher 1988; G\u00f3mez et al. 1995; Agudo et al. 2001; Mizuno et al. 2015; Mart\u00ed et al. 2016; Fuentes et al. 2018), then the strength of poloidal fields could be substantially enhanced. Indeed, the width of HST-1 is significantly smaller than expected from the parabolic (conical) width profile inside (outside) the Bondi radius (Asada & Nakamura 2012), which has been explained with a hydrodynamic recollimation shock (e.g., Stawarz et al. 2006; Bromberg & Levinson 2009; Asada & Nakamura 2012). Also, the core of blazars is often identified with a recollimation shock (e.g., Daly & Marscher 1988; Marscher 2008; Cawthorne et al. 2013). This may explain the presence of non-negligible poloidal fields in the sheath of blazar jets and in HST-1, but not in the M87 jet inside the Bondi radius.","Citation Text":["Gabuzda et al. 2018"],"Functions Text":["However, there is indication of non-negligible poloidal fields as well as toroidal fields\u2014resulting in helical magnetic fields\u2014in the jet environment of other distant AGNs, which results in transverse RM gradients with no sign changes (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[561,580]],"Functions Start End":[[256,497]]}
{"Identifier":"2016ApJ...822...15S__Brogaard_et_al._2012_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":["Brogaard et al. 2012"],"Functions Text":["Open clusters have been used to test the scaling relations for giants"],"Functions Label":["Background"],"Citation Start End":[[1507,1527]],"Functions Start End":[[1436,1505]]}
{"Identifier":"2020AandA...635A.154M__Hummels_et_al._(2017)_Instance_1","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"],"Functions Text":["RASCAS deploys a general two-step methodology (e.g."],"Functions Label":["Uses"],"Citation Start End":[[292,311]],"Functions Start End":[[240,291]]}
{"Identifier":"2022AandA...663A.172M___2012_Instance_2","Paragraph":"We note that Pavesi et al. (2019) derived a lower \u03bas parameter log(\u03bas)\u2004\u2248\u2004\u22121 for HZ10, by observing the CO(2-1) line, which implies a low star formation efficiency for this source. The conflict between the two results can be explained by the fact that Pavesi et al. (2019) estimated the gas mass by adopting a large CO-to-Mgas conversion factor \u03b1CO = 4.5 M\u2299 (K km s\u22121 pc2)\u22121, a value that is close to the Galactic conversion factor \u03b1CO = 4.36 M\u2299 (K km s\u22121 pc2)\u22121 (Bolatto et al. 2013). Although the Galactic conversion factor is a derived value for Milky Way and normal, star-forming galaxies in the local Universe, it may not be applicable for more extreme environments of starburst galaxies at high-z (see Carilli & Walter 2013 for a review). The conversion factor depends on the physical conditions of the gas in the ISM (temperature, surface density, dynamics, and metallicity), as well as the star formation and associated feedback (Narayanan et al. 2011, 2012; Genzel et al. 2012; Feldmann et al. 2012; Renaud et al. 2019; see, e.g., Bolatto et al. 2013 for a review). It is typically in the range between 0.8 and 4.36 M\u2299 (K km s\u22121 pc2)\u22121 (see, e.g., Bolatto et al. 2013; Carilli & Walter 2013, and Combes 2018 for reviews). Low metallicities (Z\u2004=\u20040.6\u2006Z\u2299 for HZ10) will drive \u03b1CO towards values higher than the Galactic value (Narayanan et al. 2012; Genzel et al. 2012; Popping et al. 2014), although \u03b1CO spans a broad range of values of \u03b1CO \u223c 0.4 \u2212 11 M\u2299 (K km s\u22121 pc2)\u22121, due to large uncertainties (see, e.g., Fig. 9 of Bolatto et al. 2013). On the other hand, high values of temperature, surface density, and velocity dispersion in a turbulent ISM of starbursts and merging systems will shift \u03b1CO towards lower values (Narayanan et al. 2011, 2012; Vallini et al. 2018). HZ10 has an extremely high value of the burstiness parameter log(\u03bas)\u2004\u223c\u20041.4 and high total density of the [C\u202fII] emitting gas log(n) \u223c 3.35 cm\u22123, and it is also a multi-component system (Jones et al. 2017, Carniani et al. 2018a). Thus, for this source we assumed \u03b1CO = 0.8 M\u2299 (K km s\u22121 pc2)\u22121, usually adopted for starburst galaxies (e.g., Downes & Solomon 1998; Bolatto et al. 2013). We obtained log(\u03bas)\u2004=\u20040.53\u2005\u00b1\u20050.34, which is within the 2\u03c3 uncertainties of the \n\n\n\nlog\n\n(\n\n\u03ba\ns\n\n)\n\n=\n1\n.\n\n43\n\n\u2212\n0.53\n\n\n+\n0.38\n\n\n\n\n$ \\log{(\\kappa_s)} = 1.43_{-0.53}^{+0.38} $\n\n\n, estimated exploiting the C\u202fIII] emission and the Vallini et al. (2020) model.","Citation Text":["Narayanan et al.","2012"],"Functions Text":["On the other hand, high values of temperature, surface density, and velocity dispersion in a turbulent ISM of starbursts and merging systems will shift \u03b1CO towards lower values"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1728,1744],[1751,1755]],"Functions Start End":[[1550,1726]]}
{"Identifier":"2022MNRAS.511.4946N__Myers_et_al._2015_Instance_1","Paragraph":"We validated our results against different selection effects due to redshift, luminosity, incompleteness of the FIRST survey, and incomplete radio counter-part identifications. However, we have not explicitly checked for any bias due to incomplete target selection in the DR16Q catalogue. As quasars are chosen for spectroscopic observations in SDSS primarily based on their photometric properties, our analysis can be biased due to the differential target selection completeness for BALQSOs and non-BALQSOs. As BALQSOs are significantly dust-reddened compared to non-BALQSOs, the completeness of BALQSOs is expected to be less. Allen et al. (2011) computed this difference in the selection probabilities using simulated BAL and non-BALQSO magnitudes and found that the completeness is very high for both BALQSOs and non-BALQSOs with z  2.1 and i  19.1, drops sharply for a narrow region close to z \u223c 2.6 and rises again for higher redshifts. This completeness difference can introduce some bias for analysis involving optically selected samples. To our rescue, eBOSS also targets all SDSS point sources that are within 1 arcsec of a radio detection in the FIRST point source catalogue (Myers et al. 2015). As our final sample is radio-selected BAL quasars and non-BAL quasars, we do not expect the behaviour of the BAL fraction to be severely influenced by any incompleteness effects. As with the case of dust-reddening, strong absorption troughs can make the broadband magnitude appear fainter for a BALQSO than an equivalent non-BALQSO. C\u2009iv BAL troughs fall in the SDSS i band for the redshift range 3.5  z  4.4. Our BAL quasar sample contains 82 quasars (4 per\u2009cent) in the redshift range where the i band luminosities are underestimated due to the presence of C\u2009iv BAL absorption. Additionally, a small fraction of quasars (1 per\u2009cent) should be LoBALs and FeLoBALs where the i-band luminosities are affected by Mg\u2009ii, Al\u2009iii, and Fe\u2009ii broad absorption. The orientation indicator, radio-to-optical core luminosity has the i-band optical luminosity in the denominator. Although the fraction of BALQSOs containing strong BAL features within the i-band wavelengths is less, any correction of the optical luminosities will further push these sources to lower log(RI) bins, thereby increasing the BAL fraction at higher orientation angles. Likewise, Hewett & Foltz (2003) suggest that optically bright BAL quasars are half as likely as non-BAL quasars to be detectable as S$_{1.4 \\,\\mathrm{ GHz}}\\, \\ge$ 1 mJy sources. To probe this, we divided the sample into three optical luminosity bins, (i) Log(Li band) \u2264 24 W\/Hz, (ii) 24 W\/Hz  Log(Li band) \u2264 25 W\/Hz, and (iii) Log(Liband) > 25 W\/Hz, and studied the distribution of radio luminosity of BAL and non-BAL quasars. We do not see any difference in the distribution of radio luminosity between BAL and non-BAL quasars for the lower optical luminosity bins. But, for the highest luminosity bin, we see that the BAL quasars have statistically lower radio luminosities as compared to the non-BAL quasars (p-value of KS test  0.05). Thus, a flux-limited survey like FIRST would have missed a fraction of optically bright BAL quasars that have radio fluxes below the FIRST sensitivity. As these missed BAL quasars have low radio luminosities and high optical luminosities, a correction for the missed quasars would only populate the lower log(RI) bins. This again will amplify the already seen trend of BAL fraction increase at high orientation angles. To ensure that the optically bright sources are not biasing the BAL fraction trend, we excluded the optically bright sources from our sample and then studied the variation of BAL fraction as a function of orientation. The pruned sample also follows the trend of high BAL fraction for high orientation angles.","Citation Text":["Myers et al. 2015"],"Functions Text":["To our rescue, eBOSS also targets all SDSS point sources that are within 1 arcsec of a radio detection in the FIRST point source catalogue"],"Functions Label":["Uses"],"Citation Start End":[[1187,1204]],"Functions Start End":[[1047,1185]]}
{"Identifier":"2021MNRAS.503.6170B__Santos-Santos,_Dom\u00ednguez-Tenreiro_&_Pawlowski_2020_Instance_1","Paragraph":"Flattened distributions of very likely co-orbiting satellite galaxies around the MW (Lynden-Bell 1976, 1982; Kroupa, Theis & Boily 2005; Pawlowski, Pflamm-Altenburg & Kroupa 2012) and M31 (Metz, Kroupa & Jerjen 2007; Ibata et al. 2013) have long posed a challenge to our understanding of galaxy formation in the \u039bCDM context. Recent proper motion data confirm that most of the classical MW satellites do indeed have a common orbital plane (Pawlowski, Kroupa & Jerjen 2013; Pawlowski & Kroupa 2020) aligned with the plane normal defined by the satellite positions alone (Santos-Santos, Dom\u00ednguez-Tenreiro & Pawlowski 2020). Their velocities show a very significant bias towards the tangential direction, as occurs for a rotating disc (Cautun & Frenk 2017). Proper motions of two M31 satellite plane members indicate that this structure is likely also coherently rotating (Sohn et al. 2020), as suggested by the RVs of satellites in this nearly edge-on structure (Ibata et al. 2013). After careful consideration of several proposed scenarios for how primordial CDM-rich satellites might end up in a thin plane, Pawlowski et al. (2014) concluded that none of them agree with observations for either the MW or M31. Structures as extreme as those observed are exceedingly rare in cosmological simulations (Ibata et al. 2014; Pawlowski & McGaugh 2014b), including hydrodynamical simulations (Ahmed, Brooks & Christensen 2017; Shao, Cautun & Frenk 2019; Pawlowski & Kroupa 2020) and simulations which model the effects of a central disc galaxy (Pawlowski et al. 2019). The arguments raised by Metz et al. (2009) and Pawlowski et al. (2014) against the group infall and filamentary accretion scenarios were later independently confirmed by Shao et al. (2018) using the EAGLE hydrodynamical cosmological simulation (Crain et al. 2015; Schaye et al. 2015). For a recent review of the satellite plane problem, we refer the reader to Pawlowski (2018), who considered both LG satellite planes and the recently discovered one around Centaurus A (M\u00fcller et al. 2018, 2021).","Citation Text":["Santos-Santos, Dom\u00ednguez-Tenreiro & Pawlowski 2020"],"Functions Text":["Recent proper motion data confirm that most of the classical MW satellites do indeed have a common orbital plane","aligned with the plane normal defined by the satellite positions alone"],"Functions Label":["Background","Background"],"Citation Start End":[[570,620]],"Functions Start End":[[326,438],[498,568]]}
{"Identifier":"2016MNRAS.463..512D__Sutter_et_al._2014_Instance_1","Paragraph":"The cosmic web, consisting of haloes, voids, filaments, and walls in large-scale structure is predicted by the cold dark matter model (Bond, Kofman & Pogosyan 1996; Pogosyan et al. 1998) and confirmed by large galaxy surveys (e.g. de Lapparent, Geller & Huchra 1986; Colless et al. 2003; Alam et al. 2015). Among these large-scale structures, the underdensities of the universe, i.e. cosmic voids, have been shown to have great potential for constraining dark energy and testing theories of gravity via several measurements. These measurements include: distance measurement via the Alcock\u2013Paczy\u0144ski test (Ryden 1995; Lavaux & Wandelt 2012; Sutter et al. 2014), weak gravitational lensing of voids (Krause et al. 2013; Clampitt & Jain 2015; Melchior et al. 2014; Gruen et al. 2016; S\u00e1nchez et al. 2016), the signal of the integrated Sachs\u2013Wolfe effect associated with voids (Sachs & Wolfe 1967; Granett, Neyrinck & Szapudi 2008; Nadathur, Hotchkiss & Sarkar 2012; Flender, Hotchkiss & Nadathur 2013; Ili\u0107, Langer & Douspis 2013; Cai et al. 2014; Planck Collaboration XIX 2014; Aiola, Kosowsky & Wang 2015; Kov\u00e1cs & Granett 2015; Planck Collaboration XXI 2015), void ellipticity as a probe for the dark energy equation of state (Lee & Park 2009; Lavaux & Wandelt 2010; Bos et al. 2012; Pisani et al. 2015; Sutter et al. 2015), void abundances and profiles for testing theories of gravity and cosmology (Li, Zhao & Koyama 2012; Clampitt, Cai & Li 2013; Barreira et al. 2015; Cai, Padilla & Li 2015; Lam et al. 2015; Massara et al. 2015; Zivick et al. 2015), coupled dark energy (Pollina et al. 2016), the nature of dark matter (Yang et al. 2015), baryon acoustic oscillations in void clustering (Kitaura et al. 2016; Liang et al. 2016), and redshift-space distortions in voids (Hamaus et al. 2015, 2016; Cai et al. 2016). Despite their popularity and great potential as a cosmological tool, a gap of knowledge between the evolution of individual voids through simulations and observations versus theory still persists. How voids evolve from the initial conditions and how dark energy or alternative theories of gravity shape this process still lacks a complete analytical understanding. As with the formation history of haloes, the initial conditions and evolution history of voids sets the base for their two fundamental properties: profile and abundance. As these are crucial for constraining cosmological parameters, it is therefore important to bridge the gap between theory and observations. This is the main goal of our study.","Citation Text":["Sutter et al. 2014"],"Functions Text":["Among these large-scale structures, the underdensities of the universe, i.e. cosmic voids, have been shown to have great potential for constraining dark energy and testing theories of gravity via several measurements. These measurements include: distance measurement via the Alcock\u2013Paczy\u0144ski test"],"Functions Label":["Motivation"],"Citation Start End":[[640,658]],"Functions Start End":[[307,603]]}
{"Identifier":"2020MNRAS.492.5152Z__Samarasinha_&_Mueller_2013_Instance_1","Paragraph":"Scheeres (2007) and Nesvorn\u1ef3 & Vokrouhlick\u1ef3 (2007) established averaging methods for YORP effects experienced by small bodies (mostly asteroids) taking real-time insolation and shape into account. However, unlike thermal radiation, as Fig. 2 shows, the CO (or H2O) production rate variation versus solar distance is too complex to be written into a linear or polynomial formula (over the entire range we want to investigate). Sidorenko, Scheeres & Byram (2008) applied the averaging method to obtain evolutionary equations in order to study the long-term variations in the nucleus spin state induced by outgassing torques, but the illumination variation is not considered in their simplified model. There are further analytical approaches described in the literature (Samarasinha & Mueller 2013; Steckloff & Jacobson 2016; Steckloff & Samarasinha 2018), although these models predicted a change in the spin state of the nucleus without taking into consideration the exact 3D shape model and its variation induced by mass loss. Therefore, to our knowledge, there exists no approximate framework that is able to combine spin and mass loss together with orbital evolution for the 3D shapes (accounting for shadowing and self-heating) we are exploring in this paper. Currently, only fully numerical schemes are appropriate to investigate sublimation torques in our context. At the same time, it is not feasible to include such numerically expensive considerations when exploring the parameter space of orbits, mass-loss function and axis orientations. Therefore, in this work, we consider the rotation of the nucleus only for the stable case, without taking into consideration the wobbling rotation excitation and the change in spin states as a result of sublimation torque. However, we plan to perform these in the future, relying on the knowledge gained from physically and scientifically interesting cases in this work. As a general remark, the spin (and\/or orientation) changes induced by sublimation torque are not expected to play a role only if the orientation variation time-scale is significantly larger than the time-scale of shape modification. However, if the time-scale is significantly shorter, new results are expected as the variation period of \u03b8 and \u03c6 might bring new periodical effects. A tumbling spin state with large procession angle or even a chaotic rotation may average the mass loss on the surface to yield an \u2018averaged\u2019 shape as a result. Therefore, if a nucleus exhibits an asymmetry shape believed to be driven by sublimation activity, it would indicate that the nucleus rotates in a stable stage for a long-term period in its dynamical history. In other words, this implies that the time-scale of wobbling excitation of such objects might be much longer than the one for shape modification.","Citation Text":["Samarasinha & Mueller 2013"],"Functions Text":["There are further analytical approaches described in the literature",", although these models predicted a change in the spin state of the nucleus without taking into consideration the exact 3D shape model and its variation induced by mass loss."],"Functions Label":["Background","Differences"],"Citation Start End":[[768,794]],"Functions Start End":[[699,766],[852,1026]]}
{"Identifier":"2022MNRAS.515.1568W__Weinberg_et_al._2003_Instance_1","Paragraph":"The temperature of the IGM can be measured using the effect of Doppler broadening on the forest of Ly\u2009\u03b1 absorption line seen in the spectra of high-redshift quasars and galaxies (Rauch 1998; Savaglio, Panagia & Padovani 2002). Traditionally, the absorbing clouds were treated as being distinct and fitting Voigt profiles (e.g. van de Hulst & Reesinck 1947; Carswell & Webb 2014; Webb, Carswell & Lee 2021) to features seen in spectra leads to two parameters, column densities and temperatures (e.g. Hu et al. 1995). However, since the identification of the Ly\u2009\u03b1 forest in hydrodynamic simulations (Cen et al. 1994; Zhang, Anninos & Norman 1995; Hernquist et al. 1996), it was realized that in the context of cosmological models for structure formation, the Ly\u2009\u03b1 forest is due to a fluctuating, continuous IGM (e.g. Bi 1993; Weinberg et al. 2003). Because of this, Hubble flow is the major source of broadening, and fitting the width of features with a thermally broadened profile of a discrete object will not yield the temperature directly. Many other techniques have therefore been developed to constrain the IGM temperature from the Ly\u2009\u03b1 forest. One is the relatively straightforward use of Voigt profile fit parameters to try to find the envelope that corresponds to a minimum IGM temperature (e.g. Garzilli, Theuns & Schaye 2020). Others include wavelet transforms (Meiksin 2000; Garzilli et al. 2012), or measuring statistical properties of the Ly\u2009\u03b1 forest flux, such as the power spectrum (Croft et al. 1998; Lai et al. 2006), which respond to the small-scale smoothing that results from increased IGM temperatures. More recent ideas include the reconstruction techniques of M\u00fcller, Behrens & Marsh (2021) that use 3D tomographic data. In the current work, we will use a neural network (NN, see below) trained on simulated Ly\u2009\u03b1 spectra to infer the temperature from the smoothness of absorption features. We will be looking in particular for discontinuous T jumps that could result from discrete sources of heating such as quasars reionizing He\u2009ii. We note that the temperature of gas will also influence the pressure smoothing scale (Gnedin & Hui 1998), which physically makes gaseous structures smoother than their underlying dark matter counterparts (Peeples et al. 2010). In our current preliminary work, we will not model this effect (see also the approach of McQuinn et al. 2011, who used a similar approximation), so that the smoothing of absorption features will be entirely attributed to Doppler broadening.","Citation Text":["Weinberg et al. 2003"],"Functions Text":["However, since the identification of the Ly\u2009\u03b1 forest in hydrodynamic simulations","it was realized that in the context of cosmological models for structure formation, the Ly\u2009\u03b1 forest is due to a fluctuating, continuous IGM (e.g."],"Functions Label":["Background","Background"],"Citation Start End":[[824,844]],"Functions Start End":[[516,596],[669,814]]}
{"Identifier":"2020ApJ...899..118C__Wang_et_al._2019_Instance_1","Paragraph":"The origin of giant pulses has been remaining a mystery since the discovery of giant pulses from Crab pulsar (Staelin & Reifenstein 1968). The generation of giant pulse activity was pointed to be an intrinsic phenomenon within the pulsar (Hankins 1971). The giant pulses are supposed to be the product of induced Compton scattering of the radio radiation off the plasma in the pulsar magnetosphere (Petrova 2006). The extremely high intensity is as well caused by an enhanced number of charges partaking in the nonthermal, coherent radiation processes (Hankins et al. 2003). Alternatively, the origination of giant pulses is proposed from the coherent instability of plasma near the magnetic equator of light cylinder (Wang et al. 2019). Singal & Vats (2012) suggested that the giant pulse emission and nulling may be opposite manifestations of the same physical process. The giant pulses are suggested to occur in pulsars with extremely high magnetic fields at the light cylinder of BLC > 105 G (Cognard et al. 1996). Therefore, the giant pulses are proposed to originate near the light cylinder (Istomin 2004). However, the giant pulses are also detected in the pulsars with ordinary magnetic fields at the light cylinder of BLC  100 G, such as PSRs B0031\u221207 (Kuzmin & Ershov 2004), B1112+50 (Ershov & Kuzmin 2003), J1752+2359 (Ershov & Kuzmin 2005), B0950+08 (Smirnova 2012), B0656+14 (Kuzmin & Ershov 2006), B1237+25 (Kazantsev & Potapov 2017), and B0301+19 (Kazantsev et al. 2019), and it does not seem to support the high BLC hypothesis. The Vela giant micropulse emission physics maybe independent on the high magnetic field at the light cylinder. Although Vela\u2019s BLC is about 20 times smaller than that of PSR B1937+21 and the Crab pulsar, it is still in the top 5% of pulsars with BLC estimate. The giant pulses from PSR J1824\u22122452A occur in narrow phase windows that correlate in phase with X-ray emission, and the two emission phenomena likely originate from the similar magnetospheric regions but not the same physical mechanism (Knight et al. 2006). In order to reveal the nature of the giant micropulses, simultaneous radio and X-ray observations on the Vela pulsar will be required.","Citation Text":["Wang et al. 2019"],"Functions Text":["Alternatively, the origination of giant pulses is proposed from the coherent instability of plasma near the magnetic equator of light cylinder"],"Functions Label":["Background"],"Citation Start End":[[719,735]],"Functions Start End":[[575,717]]}
{"Identifier":"2015MNRAS.452.2837M__Mendigut\u00eda_2013_Instance_1","Paragraph":"We constructed a sample of artificial stars representing the TT and HAeBe regime by using synthetic models of stellar atmospheres (Kurucz 1993). The properties of each object are provided in Table 1. Columns two and three show the stellar luminosity and effective temperature. From these, the stellar radii was derived, spanning between 0.7 and 4 R\u2299 (column 4). The stellar masses (column 5) were derived assuming log g = 4, and cover the 0.2\u20136 M\u2299 range. Magnetospheric accretion (MA) shock modelling was carried out for each star by adding (blackbody) accretion contributions to the photospheric (Kurucz) spectra (see e.g. the reviews in Calvet, Hartmann & Strom 2000; Mendigut\u00eda 2013). Two representative examples are presented in Fig. 2 (left-hand panel). The shock model was applied following the usual recipes for both the TTs and HAeBes, and we refer the reader to Calvet & Gullbring (1998); Mendigut\u00eda et al. (2011) and Fairlamb et al. (2015) for further details. Three different values for the UV excess in the Balmer region of the spectra (from \u223c3500 to 4000 \u212b, as defined in Mendigut\u00eda et al. 2013) were modelled for each object assuming typical values for the inward flux of energy carried by the accretion columns (1012 erg cm2 s\u22121) and the disc truncation radius (5R*): a \u2018maximum\u2019 excess (0.70 mag), whose corresponding accretion contribution is Lacc \u223c L* for L* \u2265 L\u2299; a \u2018minimum\u2019 excess (0.01 mag) representative of the observational limit, and whose corresponding accretion contribution is Lacc \u223c 0.01L* for L* \u2265 L\u2299; and finally, a \u2018typical\u2019 excess in-between the two previous (0.12 mag). The resulting accretion luminosities are shown in the last three columns of Table 1. These are plotted versus the corresponding L* values (blue diagonal dashed lines in Fig. 1), matching the overall distribution of data. We note that excesses larger than 0.70 mag could still be measured for the less luminous sources (L* \u2264 L\u2299) without reaching the upper bound (Lacc \u223c L*).","Citation Text":["Mendigut\u00eda 2013"],"Functions Text":["Magnetospheric accretion (MA) shock modelling was carried out for each star by adding (blackbody) accretion contributions to the photospheric (Kurucz) spectra (see e.g. the reviews in"],"Functions Label":["Uses"],"Citation Start End":[[670,685]],"Functions Start End":[[455,638]]}
{"Identifier":"2019MNRAS.490.5567Z__Aschwanden_et_al._1999_Instance_1","Paragraph":"Coronal loops are curvilinear structures in the outer layer of the solar atmosphere, and they are formed by thermal plasmas confined by magnetic fields and well reflect the coronal magnetic field configurations. Because a number of active phenomena in the solar atmosphere, such as flares, the solar wind, coronal mass ejection (CME) and coronal oscillation, are associated with the coronal magnetic field, instead of direct observation of the coronal magnetic field, solar physicists are able to indirectly obtain the structure and evolution of the coronal magnetic field by observing and studying the kinetic characteristics of coronal loops. On the other hand, when flares and CME occur, the coronal magnetic field structure will undergo drastic changes, which will result in the oscillation of coronal loops (Li, Liu & Tam 2017). Therefore, on the basis of the oscillation of coronal loops, especially their oscillation characteristics in frequency, phase and attenuation coefficient, solar physicists can further understand the relationships between the coronal magnetic field, flares and CME (Aschwanden et al. 1999; Jain, Maurya & Hindman 2015) and obtain corresponding physical parameters associated with oscillating coronal loops (Aschwanden & Schrijver 2011). Meanwhile, the study on coronal loops is also useful for us to well understand the properties of coronal plasmas (Parnell & De Moortel 2012), plasma motion inside coronal loops (Winebarger, DeLuca & Golub 2001; Uritsky et al. 2013; Zhou & Liang 2016), the propagation of MHD waves in coronal loops (Wang, Ofman & Davila 2013), the evolution of active regions (Yan et al. 2012) and coronal heating (Aschwanden et al. 2007a; Aschwanden & Peter 2017). In addition, in modelling the coronal magnetic field, the observed coronal loops can be taken as a constraint condition so as to obtain more realistic extrapolated results of the non-linear force-free (NLFF) magnetic field model (Su et al. 2009; Aschwanden 2013). In particular, along with the enrichment of observational data in multiple wavebands, the classification and study of different temperatures and densities of the coronal loops will be further expanded, and the meaningful physical mechanism of these phenomena will also be obtained (Aschwanden & Boerner 2011; Huang et al. 2012; Aschwanden et al. 2013a). However, it must be emphasized that the current observational data of coronal loops are just a two-dimensional (2D) projective result of the real shape of coronal loops on to the plane of sky. Therefore, to obtain the genuine shape of coronal loops, a three-dimensional (3D) reconstruction of coronal loops is needed (Aschwanden et al. 2008b). It is desirable to reconstruct the real shape of coronal loops, resulting from the fact that identical coronal loops from the 2D observation data of multiple viewing angles (the observation data of the Solar Terrestrial Relations Observatory (STEREO), Uritsky et al. 2013) can be detected and extracted.","Citation Text":["Aschwanden et al. 1999"],"Functions Text":["Therefore, on the basis of the oscillation of coronal loops, especially their oscillation characteristics in frequency, phase and attenuation coefficient, solar physicists can further understand the relationships between the coronal magnetic field, flares and CME","and obtain corresponding physical parameters associated with oscillating coronal loops"],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[1099,1121]],"Functions Start End":[[834,1097],[1152,1238]]}
{"Identifier":"2021MNRAS.501.2897G__Pribulla_&_Rucinski_2006_Instance_1","Paragraph":"EE Cet (ADS 2163 B) is the southern (slightly fainter) component of the visual binary WDS 02499+0856 (Mason et al. 2001). It was discovered by the HIPPARCOS mission (Perryman et al. 1997), by noticing the variability of the combined light of both visual components. Lampens et al. (2001) performed photometric measurements of the visual pair and gave (but only for one epoch) the following values V(A) = 9.47\u2009mag and V(B) = 9.83\u2009mag. Pribulla & Rucinski (2006) lists the orbital parameters (orientation and separation) of WDS 02499+0856 and gave \u03b8 = 194\u00b0, \u03c1 = 5.66\u2009arcsec and magnitude difference \u0394V = 0.07\u2009mag (the magnitude difference can be as large as \u0394V = 0.36\u2009mag, due to photometric variability of the eclipsing binary). WDS 02499+0856 turned out to be a quadruple system, when the northern component was found to be a double-lined (SB2) binary from the DDO spectroscopic observations (Pribulla & Rucinski 2006). D\u2019Angelo et al. (2006) re-confirmed the multiplicity of the system, and listed it among the contact binaries with additional components. Radial velocity observations from Rucinski et al. (2002) resulted in a well-defined circular orbit of the contact binary, with K1 = 84.05\u2009km\u2009s\u22121, K2 = 266.92\u2009km\u2009s\u22121 (q = 0.315), and an F8V spectral type. Karami & Mohebi (2007) using their own velocity curve analysis method, arrived at almost identical results for the mass ratio. Djura\u0161evi\u0107 et al. (2006) presented the first model, resulting in orbital inclination of i = 78.5\u00b0 and a fill-out factor of f = 32.69 per\u2009cent, T2 = 6314\u2009K, and T1 = 6095\u2009K, when spots were added. Their no-spot model resulted in very close value for the fill-out factor but slightly different geometrical and orbital parameters. The physical parameters derived in this study were: M1 = 1.37\u2009M\u2299, M2 = 0.43\u2009M\u2299, and mean radii R1 = 1.35\u2009R\u2299, R2 = 0.82\u2009R\u2299. It is worth noting here that the light curves analysed by these authors included the visual component in the photometric aperture with a contamination of about 54 per\u2009cent.","Citation Text":["Pribulla & Rucinski (2006)"],"Functions Text":["lists the orbital parameters (orientation and separation) of WDS 02499+0856 and gave \u03b8 = 194\u00b0, \u03c1 = 5.66\u2009arcsec and magnitude difference \u0394V = 0.07\u2009mag (the magnitude difference can be as large as \u0394V = 0.36\u2009mag, due to photometric variability of the eclipsing binary)"],"Functions Label":["Background"],"Citation Start End":[[434,460]],"Functions Start End":[[461,726]]}
{"Identifier":"2021MNRAS.503.3279S__Magrini_et_al._2017_Instance_2","Paragraph":"Among the several features, the distribution of chemical elements across the Galactic disc historically constitutes the most important constraint to chemo-dynamical models of our Milky Way. A number of studies (e.g. Tosi 1988; Hayden et al. 2014, 2015; Anders et al. 2017) have shown the spatial distributions of chemical abundances and their ratios across the Galactic disc. However, these studies are mainly based on field stars, which also include very old populations that had time to migrate significantly and redistribute the chemical elements across the Galaxy (e.g. Sellwood & Binney 2002; Ro\u0161kar et al. 2012; Martinez-Medina et al. 2016). Open clusters are a valuable alternative, being on average younger (Magrini et al. 2017), and therefore a better tracer of the gradients in the disc out of which the most recent stars formed. Since the work of Janes (1979), much observational evidence has established that the metallicity distribution (often abbreviated by the iron-to-hydrogen ratio [Fe\/H]) traced by clusters throughout the Milky Way disc shows a significant decrease with increasing distance from the Galactic Centre. This \u2018radial metallicity gradient\u2019 \u2013 in its apparent simplicity \u2013 reflects a complex interplay between several processes that are driving the evolution of our Galaxy, including star formation, stellar evolution, stellar migration, gas flows, and cluster disruption (Cunha & Lambert 1992, 1994; Friel 1995; Stahler & Palla 2004; Carraro et al. 2006; Boesgaard, Jensen & Deliyannis 2009; Magrini et al. 2009; Frinchaboy et al. 2013; Netopil et al. 2016; Anders et al. 2017; Spina et al. 2017; Bertelli Motta et al. 2018; Quillen et al. 2018). Complementary to the study of the overall metallicity distribution, the abundance ratios of several other elements, such as \u03b1-elements, iron peak, odd-z, and neutron capture, can provide deep insight into the variety of nucleosynthesis processes, with their production sites and time-scales (e.g. Carrera & Pancino 2011; Ting et al. 2012; Reddy, Lambert & Giridhar 2016; Duffau et al. 2017; Magrini et al. 2017, 2018; Donor et al. 2020; Casamiquela et al. 2020). Therefore, understanding the distribution of metals traced by clusters across the Galactic disc is fundamental for explaining the birth, life, and death of both stars and clusters, the recent evolution of our own Milky Way, and the evolution of other spiral galaxies (Boissier & Prantzos 2000; Bresolin 2019).","Citation Text":["Magrini et al. 2017"],"Functions Text":["Complementary to the study of the overall metallicity distribution, the abundance ratios of several other elements, such as \u03b1-elements, iron peak, odd-z, and neutron capture, can provide deep insight into the variety of nucleosynthesis processes, with their production sites and time-scales (e.g."],"Functions Label":["Background"],"Citation Start End":[[2068,2087]],"Functions Start End":[[1677,1973]]}
{"Identifier":"2018ApJ...854..167G__Meece_et_al._2017_Instance_1","Paragraph":"In the turbulent gaseous halos of clusters, groups, and galaxies (particularly massive ones), extended filaments and clouds condense out of the hot plasma in a top-down nonlinear23\n\n23\nThis nonlinear condensation process has properties significantly different from those of classic linear thermal instability (TI); the latter is mainly concerned with small overdensities overcoming buoyancy oscillations (e.g., Field 1965; Balbus & Soker 1989; Burkert & Lin 2000; Pizzolato & Soker 2005; McCourt et al. 2012\u2014more in Section 5).\n condensation cascade, forming a chaotic multiphase rain. The thermal state and kinematics of the progenitor hot plasma halo drive the formation and evolution of all the condensed structures, which inherit some of the parent properties. Part of the inner condensed gas eventually accretes onto the central supermassive black hole (SMBH), igniting the feedback response and efficiently self-regulating the entire atmosphere over several gigayears (e.g., Gaspari et al. 2011a, 2011b, 2012a, 2012b; Li & Bryan 2014; Barai et al. 2016; Soker 2016; Yang & Reynolds 2016; Meece et al. 2017; Voit et al. 2017). This feeding process is known as chaotic cold accretion (CCA; Gaspari et al. 2013b) and can intermittently boost the accretion rates up to 100\u00d7 the hot (Bondi) rate. If turbulence is subdominant, the halo tends instead to condense in a disk structure (due to the preservation of angular momentum), reducing feeding and feedback\u2014this regime is more important for low-mass, spiral galaxies.24\n\n24\nThe top-down rain differs from the bottom-up condensation in the disk of spiral galaxies, where the hot\/warm phase is created in situ by supernovae, which drive compressive, non-solenoidal turbulence (e.g., McKee & Ostriker 1977; Kim et al. 2013). Nevertheless, the two complement each other, producing multiphase gas in the more extended halo and in the disk, respectively. Massive galaxies, groups, and clusters, lacking an extended disk (e.g., Werner et al. 2014), typically reside in the top-down condensation regime.\n Finally, if the entropy of the halo (or cooling time) becomes too high, the whole atmosphere may simply prevent condensation and remain hot for an extended period of time, dramatically stifling the feedback response. Overall, assessing the dynamical state of the multiphase halos is crucial to understand the past and to predict the future evolution of cosmic structures.","Citation Text":["Meece et al. 2017"],"Functions Text":["The thermal state and kinematics of the progenitor hot plasma halo drive the formation and evolution of all the condensed structures, which inherit some of the parent properties. Part of the inner condensed gas eventually accretes onto the central supermassive black hole (SMBH), igniting the feedback response and efficiently self-regulating the entire atmosphere over several gigayears (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[1094,1111]],"Functions Start End":[[586,980]]}
{"Identifier":"2020AandA...642A.170V__Engelbrecht_&_Burger_2010_Instance_1","Paragraph":"Jovian electrons were used as test particles to model the charged-particle transport computationally (see e.g. Chenette et al. 1977; Conlon 1978; Fichtner et al. 2000; Zhang et al. 2007, and references therein) to ascertain the diffusion coefficients parallel and perpendicular to the Heliospheric Magnetic Field (HMF). This was usually done by comparing computed with measured electron intensities at Earth during periods of good or bad magnetic connection. Furthermore, given the demonstrated sensitivity of computed low-energy galactic electron intensities to various turbulence quantities (see Engelbrecht & Burger 2010, 2013; Engelbrecht 2019), it may be possible to draw conclusions from Jovian electrons to better understand the behaviour of those quantities in regions of the heliosphere where spacecraft observations of this character do not exist (see, e.g., Engelbrecht 2017). Since these transport parameters and the diffusion coefficients they depend on are spatially dependent, the time that particles reside in a certain part of the heliosphere may yield significant insights to the modulation of GCRs as well. Florinski & Pogorelov (2009) showed this dependency for GCR protons, investigating the time they spend in the heliotail, in the heliosheath, and in the solar wind within the termination shock, respectively. Utilising both galactic electrons and protons, Strauss et al. (2011a) focussed on the connection between thetotal propagation time and energy losses. They find a significant non-linear dependency on the total propagation times, which is strong enough to influence also the observations of Jovian electrons. These energy losses are entirely caused by adiabatic effects as other possible influences such as particle-particle interactions are negligible in the TPE due to a lack of significance in the interplanetary medium. As the adiabatic energy changes d E\u2215d t ~\u22122\u22153 \u22c5 EuSW\u2215r are connected to the radial position, the corresponding energy loss rate per step only depends on the temporal step size \u0394s and the radial position after the step. The radial direction of the step is thereby irrelevant. This leads to particles spending more simulation time at small radii losing more energy and implicitly to a statistical connection between the average energy losses and the particle\u2019s mean free paths.","Citation Text":["Engelbrecht & Burger 2010"],"Functions Text":["Furthermore, given the demonstrated sensitivity of computed low-energy galactic electron intensities to various turbulence quantities (see",", it may be possible to draw conclusions from Jovian electrons to better understand the behaviour of those quantities in regions of the heliosphere where spacecraft observations of this character do not exist"],"Functions Label":["Uses","Uses"],"Citation Start End":[[598,623]],"Functions Start End":[[459,597],[648,856]]}
{"Identifier":"2015ApJ...815....7V__Bale_et_al._2005_Instance_1","Paragraph":"Turbulence in plasmas is a complex phenomenon that is characterized by different regimes in different ranges of spatial and temporal scales. Turbulence in the solar wind has been extensively studied, both by detailed analyses of in situ measurements and from a theoretical point of view; see Bruno & Carbone (2005) for a review. Such studies often adopt complementary views that the turbulence may be described either as a collection of waves that interact nonlinearly, so-called wave turbulence, or else as a collection of broadband, essentially zero-frequency eddies or flux tubes that form a hierarchy of coherent structures. These approaches have been extensively reviewed (Barnes 1979; Matthaeus et al. 2015), and we do not attempt a critical comparison in the present work. Instead, we adopt mainly a wave taxonomy of the fluctuations based on linear theory in order to address a specific set of questions. As motivation, we note that a variety of observations in the solar wind (Bale et al. 2005; Sahraoui et al. 2012) have suggested that fluctuations near the end of the magnetohydrodynamic (MHD) inertial cascade range, and approaching the kinetic plasma range, may consist primarily of kinetic Alfv\u00e9n waves (KAWs). Here we address in particular the nature of fluctuations produced due to nonlinear interactions near the proton inertial length dp and investigate in some detail the basis for identifying them as KAWs. We show how phase mixing of large-scale parallel-propagating Alfv\u00e9n waves is an efficient mechanism for the production of KAWs at wavelengths close to dp and at a large propagation angle with respect to the magnetic field. To support the interpretation as KAWs, we perform and analyze MHD, Hall magnetohydrodynamic (HMHD), and hybrid Vlasov\u2013Maxwell (HVM) simulations that model the propagation of Alfv\u00e9n waves and their fully nonlinear interaction with a nonuniform plasma background. We will be able to characterize fluctuations produced by this phase-mixing-like interaction as highly oblique KAWs.","Citation Text":["Bale et al. 2005"],"Functions Text":["As motivation, we note that a variety of observations in the solar wind","have suggested that fluctuations near the end of the magnetohydrodynamic (MHD) inertial cascade range, and approaching the kinetic plasma range, may consist primarily of kinetic Alfv\u00e9n waves (KAWs)."],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[986,1002]],"Functions Start End":[[913,984],[1026,1224]]}
{"Identifier":"2018AandA...616A..11G__Forbes_et_al._2012_Instance_1","Paragraph":"In addition to secular evolutionary processes, a disc galaxy like ours is expected to have experienced several accretion events in its recent and early past (Bullock & Johnston 2005; De Lucia & Helmi 2008; Stewart et al. 2008; Cooper et al. 2010; Font et al. 2011; Brook et al. 2012; Martig et al. 2012; Pillepich et al. 2015; Deason et al. 2016; Rodriguez-Gomez et al. 2016). While some of these accretions are currently being caught in the act, like for the Sagittarius dwarf galaxy (Ibata et al. 1994) and the Magellanic Clouds (Mathewson et al. 1974; Nidever et al. 2010; D\u2019Onghia & Fox 2016), we need to find the vestiges of ancient accretion events to understand the evolution of our Galaxy and how its mass growth has proceeded over time. Events that took place in the far past are expected to have induced a thickening of the early Galactic disc, first by increasing the in-plane and vertical velocity dispersion of stars (Quinn et al. 1993; Walker et al. 1996; Villalobos & Helmi 2008, 2009; Zolotov et al. 2009; Purcell et al. 2010; Di Matteo et al. 2011; Qu et al. 2011; Font et al. 2011; McCarthy et al. 2012; Cooper et al. 2015; Welker et al. 2017), and second by agitating the gaseous disc from which new stars are born, generating early stellar populations with higher initial velocity dispersions than those currently being formed (Brook et al. 2004, 2007; Forbes et al. 2012; Bird et al. 2013; Stinson et al. 2013). These complementary modes of formation of the Galactic disc can be imprinted on kinematics-age and kinematics-abundance relations (Str\u00f6mberg 1946; Spitzer & Schwarzschild 1951; Nordstr\u00f6m et al. 2004; Seabroke & Gilmore 2007; Holmberg et al. 2007, 2009; Bovy et al. 2012a, 2016; Haywood et al. 2013; Sharma et al. 2014; Martig et al. 2016; Ness et al. 2016; Mackereth et al. 2017; Robin et al. 2017), and distinguishing between them requires full 3D kinematic information for several million stars, in order to be able to separate the contribution of accreted from in-situ populations, and to constrain impulsive signatures that are typical of accretions (Minchev et al. 2014) versus a more quiescent cooling of the Galactic disc over time. Accretion events that took place in the more recent past of our Galaxy can also generate ripples and rings in a galactic disc (G\u00f3mez et al. 2012b), as well as in the inner stellar halo (Jean-Baptiste et al. 2017). Such vertical perturbations of the disc are further complicated by the effect of spiral arms (D\u2019Onghia et al. 2016; Monari et al. 2016b), which together with the effect of accretion events might explain vertical wave modes as observed in SEGUE andRAVE (Widrow et al. 2012; Williams et al. 2013; Carrillo et al. 2018), as well as in-plane velocity anisotropy (Siebert et al. 2012; Monari et al. 2016b). Mapping the kinematics out to several kiloparsec from the Sun is crucial for understanding whether signs of these recent and ongoing accretion events are visible in the Galactic disc, to ultimately understand to what extent the Galaxy can be represented as a system in dynamical equilibrium (H\u00e4fner et al. 2000; Dehnen & Binney 1998), at least in its inner regions, or to recover the nature of the perturber and the time of its accretion instead from the characteristics and strength of these ringing modes (G\u00f3mez et al. 2012b).","Citation Text":["Forbes et al. 2012"],"Functions Text":["Events that took place in the far past are expected to have induced a thickening of the early Galactic disc,","and second by agitating the gaseous disc from which new stars are born, generating early stellar populations with higher initial velocity dispersions than those currently being formed"],"Functions Label":["Background","Background"],"Citation Start End":[[1373,1391]],"Functions Start End":[[746,854],[1163,1346]]}
{"Identifier":"2017AandA...608A..75C__Osten_et_al._2005_Instance_1","Paragraph":"The spectral energy distribution of the flare determines the altitude range in the planet\u2019s atmosphere that is affected by the flare. Intense flares on AU Mic display a strong continuum emission enhancement in the XUV (see Fig. 2c) which results in increased ionisation over a broad altitude range in the thermosphere. Large continuum enhancement during flares have also been seen in other wavebands. Strong increases in the FUV have been observed on AU Mic (Robinson et al. 2001), and also on another active M dwarf, AD Leo (Hawley & Pettersen 1991). More recently, Kowalski et al. (2010) showed that continuum emission may be the dominant luminosity source in the near ultraviolet (NUV) during flares. There have not been many simultaneous multi-wavelength studies of flaring stars (Hawley et al. 2003; Osten et al. 2005). There is a particular dearth of simultaneous data in the EUV which is critical for studies of exoplanetary upper atmospheres, but where measurements are difficult due to ISM absorption. More observations are needed, but if young, magnetically active stars display broadband continuum enhancements during flares, similar to AU Mic, this would mean that large altitude ranges in planetary atmospheres are affected. Programmes such as Measurements of the Ultraviolet Spectral Characteristics of Low-mass Exoplanetary Systems (MUSCLES) have begun to provide more information on the high-energy spectral shape of K and M dwarf stars, including in the EUV (France et al. 2016; Youngblood et al. 2016; Loyd et al. 2016). However, in these studies, the EUV flux is computed using reconstructed Lyman alpha flux assuming a spectral energy distributio similar to the Sun. This approach could lead to incorrect results for active stars. Further development of coronal models driven by multi-wavelength observations would provide the most accurate high energy spectra for active low-mass stars. Such observations should be provided by the Multiwavelength Observations of an eVaporating Exoplanet and its Star (MOVES) program which has begun to undertake a long-term study of the HD 189733 system (Fares et al. 2017). ","Citation Text":["Osten et al. 2005"],"Functions Text":["There have not been many simultaneous multi-wavelength studies of flaring stars","There is a particular dearth of simultaneous data in the EUV which is critical for studies of exoplanetary upper atmospheres, but where measurements are difficult due to ISM absorption. More observations are needed, but if young, magnetically active stars display broadband continuum enhancements during flares, similar to AU Mic, this would mean that large altitude ranges in planetary atmospheres are affected."],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[805,822]],"Functions Start End":[[704,783],[825,1237]]}
{"Identifier":"2017MNRAS.469.3108C__Schweizer_&_Seitzer_1992_Instance_1","Paragraph":"Having intermediate colours between the blue cloud and the red sequence, galaxies populating the so-called green-valley (e.g. Salim 2014; Schawinski et al. 2014) are generally considered as the transiting objects par excellence (Martin et al. 2007; Mendel et al. 2013; Salim 2014; Schawinski et al. 2014). Among these, the most interesting population certainly consists of those galaxies which have just entered the quenching phase (within a few Myr). Although hampered by the short duration of the quenching process, the search for galaxies in this critical phase of evolution has been carried on by several authors in the past decades. Galaxies characterized by both a tidally disturbed morphology and intermediate colours (e.g. Schweizer & Seitzer 1992; Tal et al. 2009) or low level of recent SF (Kaviraj 2010), young elliptical galaxies (Sanders et al. 1988; Genzel et al. 2001; Dasyra et al. 2006) and very recent post-merger remnants with strong morphological disturbances (Hibbard & van Gorkom 1996; Rothberg & Joseph 2004; Carpineti et al. 2012) have been considered as valid 'recent time' quenching candidates. Moreover, many attempts aimed at spectroscopically identifying quenching galaxies come from the investigations of the post-starburst (E+A or K+A) galaxies' UV and optical spectra, whose strong Balmer absorption lines and missing [O\u2009ii] \u03bb3727 (hereafter [O\u2009ii]) and H\u03b1 emission lines (Couch & Sharples 1987; Poggianti et al. 2004; Quintero et al. 2004; Balogh et al. 2011; Muzzin et al. 2012; Mok et al. 2013; Wu et al. 2014) have been interpreted as signs of a recent halt of the SF (Dressler & Gunn 1983; Zabludoff et al. 1996; Quintero et al. 2004; Poggianti et al. 2008; Wild et al. 2009). The scarcity of galaxies that are in the transition phase suggests that, whatever mechanism may be responsible for the SF shut-off, it has to happen on short time-scales (Tinker, Wechsler & Zheng 2010; Salim 2014). The relatively short duration of the quenching process is also suggested by the surprising identification of a significant number of galaxies that look already quiescent at z \u223c 4\u20135, when the Universe was only \u223c1\u20131.5 Gyr old (e.g. Mobasher et al. 2005; Wiklind et al. 2008; Juarez et al. 2009; Brammer et al. 2011; Marsan et al. 2015; Citro et al. 2016). However, a general and coherent picture concerning the SF quenching is still lacking. Very recently, evidence that the quenching of the SF could be a separated process with respect to the morphological transformation has come from the photometric and spectroscopic investigations of passive spiral galaxies (Fraser-McKelvie et al. 2016).","Citation Text":["Schweizer & Seitzer 1992"],"Functions Text":["Although hampered by the short duration of the quenching process, the search for galaxies in this critical phase of evolution has been carried on by several authors in the past decades. Galaxies characterized by both a tidally disturbed morphology and intermediate colours (e.g","have been considered as valid 'recent time' quenching candidates."],"Functions Label":["Background","Background"],"Citation Start End":[[731,755]],"Functions Start End":[[452,729],[1055,1120]]}
{"Identifier":"2022ApJ...929...32S__Singh_et_al._2017_Instance_1","Paragraph":"The 21 cm global spectrum experiments aim to measure the sky-averaged spectrum with high precision so as to probe the early epochs of the universe. There are a number of such ground-based experiments, including the Experiment to Detect the Global Epoch-of-Reionization Signature (EDGES; Bowman et al. 2008; Bowman & Rogers 2010; Monsalve et al. 2017), the Sonda Cosmol\u00f3gica de las Islas para la Detecci\u00f3nde Hidr\u00f3geno Neutro (SCI-HI; Voytek et al. 2014), the Probing Radio Intensity at high-z from Marion (PRIzM; Philip et al. 2019), the Shaped Antenna measurement of the background RAdio Spectrum (SARAS; Patra et al. 2013; Singh et al. 2017, 2018), the Cosmic Twilight Polarimeter (CTP; Nhan et al. 2019), the Broadband Instrument for Global Hydrogen Reionization Signal (BIGHORNS; Sokolowski et al. 2015), the Large-Aperture Experiment to Detect the Dark Age (LEDA; Bernardi et al. 2016; Bernardi 2018; Price et al. 2018), and the Radio Experiment for the Analysis of Cosmic Hydrogen (REACH; de Lera Acedo 2019). Compared to the 21 cm tomography experiments, measurement of the global 21 cm emission has a higher raw sensitivity and requires smaller collecting area, so that it could be conducted even with a single antenna. The EDGES (Bowman et al. 2018) reported the detection of a strong absorption feature at \u223c78 MHz, which has a cosmic dawn 21 cm spectrum interpretation, though it has an unexpectedly large amplitude (0.5 K) and an unusual flattened profile. If it originated from the cosmic 21 cm spectrum, this would suggest possibly new physics or astrophysics (Chen & Miralda-Escud\u00e9 2004, 2008; Creasey et al. 2011; Nelson et al. 2013; Barkana 2018; Fialkov et al. 2018; Fraser et al. 2018; Barkana et al. 2018; Slatyer & Wu 2018; Li et al. 2021; Houston et al. 2018; Hirano & Bromm 2018; Mu\u00f1oz et al. 2018; Yang 2021), though this result was not confirmed by a recent measurement of the SARAS experiment (Singh et al. 2022). It is imperative to check this result and improve upon it with further and more precise observations.","Citation Text":["Singh et al. 2017"],"Functions Text":["There are a number of such ground-based experiments, including","the Shaped Antenna measurement of the background RAdio Spectrum (SARAS;"],"Functions Label":["Background","Background"],"Citation Start End":[[624,641]],"Functions Start End":[[148,210],[533,604]]}
{"Identifier":"2022MNRAS.512.1629F__Fioroni,_Savage_&_DeYonker_2019_Instance_1","Paragraph":"The ORCA software (version 4.0.2) (Neese 2012) was used for all geometry minimizations, potential energy surface (PES), and vibrational frequency analyses using the global hybrid functional PW6B95 (Zhao & Truhlar 2005) coupled to the split valence triple-\u03b6 def2-TZVPP basis set with two sets of polarization functions (Weigend & Ahlrichs 2005) and the atom-pairwise dispersion correction energy with Becke\u2013Johnson damping (D3BJ) (Grimme et al. 2010; Grimme, Ehrlich & Goerigk 2011). The selected level of theory (PW6B95-D3BJ\/def2-TZVPP) is a reliable and accurate theoretical tool in the estimation of general main group thermochemistry, kinetics, and non-covalent interactions after the double hybrid functionals (Goerigk et al. 2017). The reliability of the used method is also underlined by the good qualitative agreement between the Density Functional Theory (DFT) and MP2-F12 calculations as found in previous works (Fioroni & DeYonker 2016; Fioroni et al. 2018; Fioroni, Savage & DeYonker 2019). To speed up calculations, the RI (Resolution of the Identity) (Neese 2003) and RIJCOSX (Neese et al. 2009b) algorithms were used coupling the Coulomb-fitting basis sets def2\/J (Weigend 2006). Because all the considered TSs involve a simple bond breaking\/formation or a dihedral rotation, after a PES search performed on the desired reaction coordinate, the eigenvector following method (Schlegel 1982; Horn et al. 1991; Eckert, Pulay & Werner 1997), as implemented in ORCA, was used. Furthermore, the computed structures were verified to represent a minimum or a transition state by analysis of harmonic vibrational frequency calculations. Finally, the visualization of the normal mode associated to the TS was analysed. The obtained enthalpies: (HTot = [EEl. + EZPE + EVib. + ERot. + ETrans.] + kBT) and S values (STot = SEl. + SVib. + SRot. + STrans.) were used to estimate the free energies (G) at T = 200 K. The selected T = 200 K is lower bound to the HCN polymerization to progress efficiently. When referring to astronomical bodies, such temperatures can be experienced, for example, by comets where temperature rises periodically by surface heating to release HCN and H molecules (Hoang et al. 2019) or by dust particles or greater bodies in the turbulent phase of a proto-planetary disc and planetary system evolution.","Citation Text":["Fioroni, Savage & DeYonker 2019"],"Functions Text":["The reliability of the used method is also underlined by the good qualitative agreement between the Density Functional Theory (DFT) and MP2-F12 calculations as found in previous works"],"Functions Label":["Similarities"],"Citation Start End":[[968,999]],"Functions Start End":[[737,920]]}
{"Identifier":"2016MNRAS.461.1719C__Clements,_Dunne_&_Eales_2010_Instance_1","Paragraph":"The early history of galaxy clusters is a poorly constrained aspect of galaxy and large-scale structure formation. Hierarchical clustering models predict that massive elliptical galaxies will form in the cores of what will become the most massive galaxy clusters today, but the epoch of the bulk of star formation for these galaxies is unclear. Observations of high redshift clusters (z = 1\u20131.5) by the ISCS project (IRAC Shallow Cluster Survey; Eisenhardt et al. 2008) suggest that this is at z > 3, and the presence of well-defined red sequences of galaxies in clusters out to z \u223c 2 (Andreon & Huertas-Company 2011; Gobat et al. 2011; Santos et al. 2011; Pierini et al. 2012) supports this conclusion. Theoretical models by Granato et al. (2004) suggest that forming clusters will go through a phase in which multiple members will undergo near-simultaneous massive bursts of star formation. The spectral energy distribution of these objects would be dominated by the far-IR, as is the case for local massive starbursts (e.g. Clements, Dunne & Eales 2010). A galaxy cluster or protocluster (we use the term protocluster to indicate a structure that has yet to become virialized) going through such a formative phase would appear as a clump of dusty protospheroidal galaxies, and might be detected by observations in the far-IR and submm bands. Hints of such objects may already have been found by Spitzer (Magliocchetti et al. 2007) and SCUBA (Ivison et al. 2000; Priddey, Ivison & Isaak 2008; Stevens et al. 2010 and references therein). A recent study by Ivison et al. (2013) has uncovered a group of HLIRGs and ULIRGs at z = 2.41 thought to be the progenitor of a 1014.6\u2009M\u2299 cluster. At still higher redshifts, the highest redshift protocluster currently known, at z \u223c 5.3 (Riechers et al. 2010; Capak et al. 2011), and a group of objects associated with a z = 5 quasar (Husband et al. 2013) both contain at least one submm-luminous object, while the best studied group of high z submm-luminous sources to date is probably the four objects associated with a source in the GOODS-North field designated GN20, all lying at z = 4 (Daddi et al. 2009; Carilli et al. 2011), though this group is extended over a broad range of redshifts \u0394z \u223c 0.1.","Citation Text":["Clements, Dunne & Eales 2010"],"Functions Text":["Theoretical models by Granato et al. (2004) suggest that forming clusters will go through a phase in which multiple members will undergo near-simultaneous massive bursts of star formation. The spectral energy distribution of these objects would be dominated by the far-IR, as is the case for local massive starbursts (e.g."],"Functions Label":["Background"],"Citation Start End":[[1027,1055]],"Functions Start End":[[704,1026]]}
{"Identifier":"2020ApJ...898L..33P__Delrez_et_al._2018_Instance_2","Paragraph":"For the TRAPPIST-1 system, data obtained by HST provide initial constraints on the extent and composition of the planet\u2019s atmospheres, suggesting that the four innermost planets do not have a cloud\/haze-free H2-dominated atmosphere (de Wit et al. 2016, 2018). However, follow-up work by Moran et al. (2018) have shown that HST data can also be fit to a cloudy\/hazy H2-dominated atmosphere. Complementary to HST, NASA\u2019s Spitzer Space Telescope\u2014which played a major role in the discovery and orbital determination of TRAPPIST-1d, e, f, and g (Gillon et al. 2017)\u2014has also allowed us to put additional constraints on the atmospheric composition of TRAPPIST-1b. Transit observations with Spitzer (Delrez et al. 2018) have found a +208 \u00b1 110 ppm difference between the 3.6 and 4.2 \u03bcm bands, suggesting CO2 absorption. Spitzer also showed that transit depth measurements do not show any hint of significant stellar contamination in the 4.5 \u03bcm spectral range. Morris et al. (2018) reached the same conclusion using a \u201cself-contamination\u201d approach based on the Spitzer data set. Spitzer's \u201cRed Worlds\u201d Program encompassed over 1000 hours of observations of the TRAPPIST-1 system, whose global results have been presented (Ducrot et al. 2020). HST and Spitzer measurements have also been combined with transit light curves obtained from space with K2 (Luger et al. 2017) and from the ground with the SPECULOOS-South Observatory (Burdanov et al. 2018; Gillon 2018) and Liverpool Telescope (Steele et al. 2004) where Ducrot et al. (2018) produced featureless transmission spectra for the planets in the 0.8\u20134.5 \u03bcm wavelength range, showing an absence of significant temporal variations of the transit depths in the visible. Additional ground-based observations with the United Kingdom Infra-Red Telescope, Anglo-Australian Telescope, and Very Large Telescope also show no substantial temporal variations of transit depths for TRAPPIST-1 b, c, e, and g (Burdanov et al. 2019). While the K2 optical data set detected a 3.3 day periodic 1% photometric modulation, it is not present in the Spitzer observations (Delrez et al. 2018). Further constraints on the molecular weight and presence\/absence of atmospheres on the TRAPPIST-1 planets will require additional observations with future facilities.","Citation Text":["Delrez et al. 2018"],"Functions Text":["While the K2 optical data set detected a 3.3 day periodic 1% photometric modulation, it is not present in the Spitzer observations"],"Functions Label":["Differences"],"Citation Start End":[[2097,2115]],"Functions Start End":[[1965,2095]]}
{"Identifier":"2019MNRAS.490..157M__Yu_&_Tremaine_2003_Instance_2","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"],"Functions Text":["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."],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[2317,2335]],"Functions Start End":[[2155,2316]]}
{"Identifier":"2022ApJ...927..138L__Lucas_et_al._2018_Instance_1","Paragraph":"A color\u2013color diagram in V \u2212 R versus B \u2212 V of the individual SSCs projected against the central spiral is shown in Figure 2 (left panel). The manner in which photometry of the individual stellar clusters was performed, including background subtraction, is described at length in Lim et al. (2020). In brief, we used the StarFinder algorithm (Diolaiti et al. 2000) to perform PSF-fitting photometry that includes subtraction of the local background. Only objects that exceed 4\u03c3\nlocal, where \u03c3\nlocal is the local root-mean-square noise level, are accepted as detections. Conversions to magnitude are based on the standard zero-points in the Vega system as described in the ACS Data Handbook (Lucas et al. 2018), and corrected for Galactic extinction. Two age loci, each of which is based on a model single stellar population (SSP; i.e., population of stars all having the same age and metallicity) as taken from Zackrisson et al. (2011), are plotted in Figure 2 for different metallicities of Z = 0.4 Z\n\u2299 (black) and Z = Z\n\u2299 (yellow). The model SSPs are based on those utilized in the Starburst99 software (Leitherer et al. 1999; V\u00e1zquez & Leitherer 2005), and their colors include both continuum and line emission from H ii regions corresponding to an adopted unity filling factor for the gas left over from star formation. The resulting line emission, in particular H\u03b1+[N ii], that is contained in the R band causes the steep rise in V \u2212 R starting below B \u2212 V \u2248 0.0 when the star clusters are younger than \u223c10 Myr. The two loci having different metallicities closely overlap, such that the inferred age at a given color differs little between the two loci. The  SSCs selected for study as indicated in Figure 1 (lower row) are concentrated around B \u2212 V \u2248 0.35, corresponding to an age around \u223c500 Myr (more details below). The spread in their colors is larger than can be accounted for by photometric uncertainties alone, although there may well be a degree of contamination by stellar clusters associated with the HVS, as well as the chance projection of SSCs lying beyond 5 kpc but positioned along the sightline toward the central spiral. By contrast, as shown in the color\u2013color diagram of Figure 2 (right panel), the SSCs located beyond the central spiral (omitting the region occupied by the HVS and the outer regions of NGC 1275 dominated by GCs, as described by Lim et al. 2020) span a continuous range of ages from a few Myr to at least \u223c1 Gyr, beyond which they cannot be distinguished from the even more numerous GCs around NGC 1275.","Citation Text":["Lucas et al. 2018"],"Functions Text":["Conversions to magnitude are based on the standard zero-points in the Vega system as described in the ACS Data Handbook"],"Functions Label":["Uses"],"Citation Start End":[[691,708]],"Functions Start End":[[570,689]]}
{"Identifier":"2018MNRAS.479..615M__Shakura_&_Sunyaev_1973_Instance_1","Paragraph":"We also detect a Compton hump at around 15\u201330 keV during the 2016 NuSTAR observation. The inclusion of the 10\u201350 keV spectral data does not affect the spectral model parameters (disc inclination angle, spin, emissivity indices, and break radius) obtained from fitting of the 0.3\u201310 keV spectrum only. The best-fitting value of the disc electron density as derived from the modelling of the broad-band (0.3\u201350 keV) spectral data is $n_{\\rm e}=5.2^{+5.2}_{-4.2}\\times 10^{16}$ cm-3. Mrk 1044 is known to be a highly accreting AGN with the dimensionless mass accretion rate of $\\dot{m}=\\frac{\\dot{M}c^{2}}{L_{\\rm E}}=16.6^{+25.1}_{-10.1}$ (Du et al. 2015). At high accretion rate, the inner region of a standard \u03b1-disc (Shakura & Sunyaev 1973) is radiation pressure dominated and the electron density of the disc can be written as (Svensson & Zdziarski 1994)\n(11)\r\n\\begin{eqnarray*}\r\nn_\\mathrm{ e}=\\frac{1}{\\sigma _{\\rm T}R_{\\rm S}} \\frac{256\\sqrt{2}}{27}\\alpha ^{-1}r^{3\/2}\\dot{m}^{-2} [1-(3\/r)]^{-1} (1-f)^{-3}.\r\n\\end{eqnarray*}\r\nwhere \u03c3T = 6.64 \u00d7 10\u221225 cm2 is the Thomson scattering cross section, RS = 2GMBH\/c2 is the Schwarzschild radius, MBH is the black hole mass, \u03b1 = 0.1 is the disc viscosity parameter, r = R\/RS, R is the characteristic disc radius, $\\dot{m}=\\frac{\\dot{M}c^{2}}{L_{\\rm E}}$ is the dimensionless mass accretion rate, and f is the fraction of the total power released by the disc into the corona. The variation of the disc electron density (ne) with the dimensionless mass accretion rate ($\\dot{m}$) for \u03b1 = 0.1, $M_{\\rm BH}=3\\times 10^{6}\\,\\mathrm{M}_{\\odot }$, f = 0.9, and r = 10 is shown as the solid curve in Fig. 17. As evident from Fig. 17, the assumption of constant disc density (ne = 1015 cm-3) is not physically realistic for low-mass AGNs even when the mass accretion rate is very high. The observed best-fitting value for the disc electron density of the source and its 90 per cent confidence limits are shown as the dotted and dashed lines in Fig. 17, respectively. The corresponding dimensionless mass accretion rate of Mrk 1044 estimated using equation (11) is $\\dot{m}\\approx 10-32$, which is in agreement with that found by Du et al. (2015). We further verified the SMBH mass with the use of X-ray variability techniques as pioneered by Ponti et al. (2012). The relation between the SMBH mass (MBH, 7) in units of $10^{7}\\,\\mathrm{M}_{\\odot }$ and normalized excess variance ($\\sigma _{\\rm NXS}^{2}$) in the 2\u201310 keV light curves of 10 ks segments and the bin size of 250 s, can be written as\n(12)\r\n\\begin{eqnarray*}\r\n\\log (\\sigma _{\\rm NXS}^{2})=(-1.83\\pm 0.1)+(-1.04\\pm 0.09)\\log (M_{\\rm BH,7}).\r\n\\end{eqnarray*}\r\nThe SMBH mass of Mrk 1044 measured using equation (12) is $M_{\\rm BH}=(4-5)\\times 10^{6}\\,\\mathrm{M}_{\\odot }$ which is close to that measured by Du et al. (2015).","Citation Text":["Shakura & Sunyaev 1973"],"Functions Text":["At high accretion rate, the inner region of a standard \u03b1-disc"],"Functions Label":["Uses"],"Citation Start End":[[717,739]],"Functions Start End":[[654,715]]}
{"Identifier":"2021AandA...649A.142A__Zhang_et_al._2019_Instance_1","Paragraph":"An et al. (1988) and Wu et al. (1990) investigated the effects of plasma injection on the formation of the Kippenhahn-Schl\u00fcter model of prominence in optimum conditions. These authors found that for high values of the plasma-\u03b2 parameter (the ratio of plasma pressure to magnetic pressure) the magnetic arcade develops a magnetic dip at the centre of the structure that supports the prominence plasma. However, comparing with our study, in the low plasma-\u03b2 regime (or under others injection conditions) they found that the dip is less deep and the system develops two additional plasma enhancements located at the lateral edges of the magnetic arcade. Recent works suggest that the deformation of the magnetic field lines is determined by the parameter \u03b4 (the ratio of the gravity to the magnetic pressure) (Zhou et al. 2018; Zhang et al. 2019). An et al. (1988) suggested that the steady lateral plasma accumulates because of both the injection process and because the field lines without dips do not geometrically contain the injected plasma, but Wu et al. (1990) proposed that the prominence mass is also supported by an increase in the pressure gradient. Since in this study we consider that magnetic field lines do not change owing to the presence of the dense thread we investigate in detail the results of An et al. (1988) and Wu et al. (1990) in part II of this paper. On the contrary, studies of the formation of 1D filament threads by chromospheric heating in the presence of non-adiabatic effects, such as radiative losses and thermal conduction, show that for magnetic loops without a dip, the plasma condenses but it streams along the magnetic field and disappears after falling to the footpoints (Antiochos et al. 2000; Karpen et al. 2006). Moreover, when the thread is initially in a thermal and force-balance equilibrium state but it is disturbed by a strong velocity perturbation, the prominence mass drains down to the chromosphere. Dense blobs of falling plasma have been habitually observed (Schrijver 2001; de Groof et al. 2005), therefore it seems that threads cannot be held static along vertical magnetic flux tubes in the corona. The coronal part of the tube can only slow down the falling blobs. M\u00fcller et al. (2004) proposed that the acceleration reduces because the pressure of the cooling plasma underneath the radiating blobs slows down the descent, and Oliver et al. (2014, 2016) and Mart\u00ednez-G\u00f3mez et al. (2020) argue that pressure gradient is the main force that opposes the action of gravity. Our study proposes that the pressure gradient can cause the equilibrium of threads in quasi-vertical flux tubes without dips even though it has not been corroborated by observations. Besides, our model is relatively simple to study the stability of filament threads in the magnetic field without dips. Other processes such as radiation and heat conduction must be considered, which might change the stability results.","Citation Text":["Zhang et al. 2019"],"Functions Text":["Recent works suggest that the deformation of the magnetic field lines is determined by the parameter \u03b4 (the ratio of the gravity to the magnetic pressure)"],"Functions Label":["Background"],"Citation Start End":[[825,842]],"Functions Start End":[[651,805]]}
{"Identifier":"2022MNRAS.517.4529B__Boruah,_Rozo_&_Fiedorowicz_2022_Instance_1","Paragraph":"The other criteria that we can use to categorize the reconstruction methods is whether the reconstruction is performed using forward-modelling or uses a direct inversion from the data. Inverting non-linear problems from partial, noisy, observations is an ill-posed inverse problem, which makes forward-modelled Bayesian methods particularly suitable for the task of reconstruction of high-dimensional fields. Bayesian reconstruction methods have become increasingly popular in cosmology and have been applied in a range of different applications such as initial conditions reconstruction (Jasche & Wandelt 2013; Modi, Feng & Seljak 2018; Jasche & Lavaux 2019), weak lensing (Fiedorowicz et al. 2022; Porqueres et al. 2021, 2022; Boruah, Rozo & Fiedorowicz 2022), and CMB lensing (Millea et al. 2021; Millea, Anderes & Wandelt 2020). Such methods have also been used for the local velocity field reconstruction. The simplest of such methods uses a Wiener filtering technique (Zaroubi, Hoffman & Dekel 1999). This approach assumes that the density\/velocity field is described as a Gaussian random field and the Wiener filtered reconstruction is the maximum-a-posteriori (MAP) solution for the problem. The Wiener filtering approach has been extended to account for uncertainties and biases in the reconstruction using a constrained realization approach (Hoffman & Ribak 1991; Hoffman, Courtois & Tully 2015; Hoffman et al. 2018; Lilow & Nusser 2021) An alternative way to account for the biases in the reconstruction in Wiener filtering is using the unbiased minimal variance approach (Zaroubi 2002). Another similar approach is the Bayesian hierarchical method, virbius (Lavaux 2016), which is based on the constrained realization approach but accounts for many different systematic effects in its analysis. This approach has been been applied to the Cosmicflows-3 (Tully, Courtois & Sorce 2016) data set by Graziani et al. (2019). A similar reconstruction code, hamlet, was introduced in Valade et al. (2022). However, these methods fail to account for the inhomogeneous Malmquist (IHM) bias which is an important source of systematic error in peculiar velocity analysis. The IHM bias arises from an incorrect assumption on the distribution of peculiar velocity tracers due to neglecting the line-of-sight inhomogeneities.","Citation Text":["Boruah, Rozo & Fiedorowicz 2022"],"Functions Text":["Bayesian reconstruction methods have become increasingly popular in cosmology and have been applied in a range of different applications such as","weak lensing"],"Functions Label":["Background","Background"],"Citation Start End":[[729,760]],"Functions Start End":[[409,553],[661,673]]}
{"Identifier":"2021ApJ...908..187W__Hayakawa_et_al._2020b_Instance_1","Paragraph":"Before 1500 A.D., when the magnetic latitude of China was higher than that at present, most of the observed aurorae were caused by CIRs and moderate CMEs. After 1500 A.D., the GNP moved to Canada. This caused the geomagnetic field latitude of China to decrease considerably. Meanwhile, the field intensity decreased to its minimum value in the last 2000 yr. Under these conditions, it would have been very rare to observe the aurora in most places in China after the 16th century, just like the present case (Wu et al. 2016a, 2016b; Hayakawa et al. 2020b, 2020c). However, this provides the best opportunity to investigate the existence of great-storm CMEs, which generate a wider auroral belt and equatorward shift of the EBAO and the EBAV (Boudouridis et al. 2003; Shue et al. 2009; Sigernes et al. 2011). If auroras were observed in China at lower magnetic latitudes than the EBAV related to the great storms, this indicates that the strong CMEs occurred. To quantify our results, we applied the classification of magnetic storms at the great level having Dst  \u2212350 nT (Loewe & Pr\u00f6lss 1997). According to the EBAV of the great level shown as the red dashed line in Figure 2(b), the auroras caused by great-storm CMEs could be identified, which were below the red dashed line. Surprisingly, there were many great-storm CMEs between 1500 and 1900, when Sp\u00f6rer minimum (1390\u20131550), Maunder minimum (1645\u20131715), and Dalton minimum (1797\u20131827) were inside (Lean et al. 1995; Solanki & Fligge 2000; Usoskin et al. 2015; Cliver & Herbst 2018; Hayakawa et al. 2020e). There were auroral sighting records in the same year, such as 1620 and 1646, from different places in China. As previously reported by Hayakawa et al. (2018b) there were simultaneous auroral records on the same day from different places in China in 1730. There are quite a few studies about the aurorae in September 1770 (Willis et al. 1996; Ebihara et al. 2017; Hayakawa et al. 2017a). Apart from the case analyses for extreme space weather events (Hattori et al. 2019; Isobe et al. 2019), our results indicate that there are more great-storm CMEs in the past 400 yr. Some of the great-storm CMEs happened in the solar maximum years, and some happened in the solar minimum years, as shown in Figure 7.","Citation Text":["Hayakawa et al. 2020b"],"Functions Text":["After 1500 A.D., the GNP moved to Canada. This caused the geomagnetic field latitude of China to decrease considerably. Meanwhile, the field intensity decreased to its minimum value in the last 2000 yr. Under these conditions, it would have been very rare to observe the aurora in most places in China after the 16th century, just like the present case"],"Functions Label":["Background"],"Citation Start End":[[533,554]],"Functions Start End":[[155,507]]}
{"Identifier":"2022MNRAS.515.5416Y__Wechsler_et_al._2022_Instance_1","Paragraph":"There are several different methods used in the literature for creating mock catalogues and light-cones. In purely empirical methods such as the JAdes extraGalactic Ultradeep Artificial Realizations (jaguar) models used to create mock catalogues in support of the JADES survey (Williams et al. 2018), observed galaxy properties are interpolated or extrapolated, and there is no underlying physics model nor setting within a \u039b cold dark matter (\u039bCDM) context. In what are sometimes called \u2018semi-empirical\u2019 methods [also called subhalo abundance matching (SHAM) or halo occupation distribution (HOD) models; see Wechsler & Tinker 2018], galaxy properties are mapped on to the properties of dark matter haloes such that a set of observational quantities is reproduced (Behroozi, Conroy & Wechsler 2010; Moster, Naab & White 2013, 2018; Behroozi et al. 2019; Wechsler et al. 2022). Both of these methods have the advantage that they are computationally efficient, are not dependent on a specific model for galaxy formation, and are guaranteed to match the observations that were used to calibrate them. However, they have the disadvantage that using them for forecasts for new observations is highly uncertain, and they are of limited use for interpretation. Semi-empirical models are typically calibrated using derived physical properties such as stellar masses and star formation rates (SFRs), which are highly uncertain at high redshifts, leading to models that are nominally calibrated on the same observations, but which have very different predictions for the link between galaxy and dark matter halo properties (see e.g. Yung et al. 2019b). We note that this is only a general overview for the semi-empirical modelling approach. These models are designed with different purposes in mind and adopt different calibration criteria. For example, UniverseMachine used high-redshift ultraviolet (UV) luminosity functions (LFs) to calibrate galaxy growth at the highest redshifts (Behroozi et al. 2019, 2020) and the semi-empirical model presented in Behroozi & Silk (2015) allows forecasts under the assumption that galaxy\u2013halo growth relationships are given by power laws.","Citation Text":["Wechsler et al. 2022"],"Functions Text":["In what are sometimes called \u2018semi-empirical\u2019 methods [also called subhalo abundance matching (SHAM) or halo occupation distribution (HOD) models; see Wechsler & Tinker 2018], galaxy properties are mapped on to the properties of dark matter haloes such that a set of observational quantities is reproduced"],"Functions Label":["Background"],"Citation Start End":[[855,875]],"Functions Start End":[[459,764]]}
{"Identifier":"2021ApJ...914...88P__Forgan_et_al._2018b_Instance_1","Paragraph":"Observing with instruments such as the Atacama Large Millimeter\/submillimeter Array (ALMA) is crucial to our understanding of planet-formation mechanisms, as we can observe at wavelengths that trace continuum emission from the cold midplane (e.g., Testi et al. 2014), where we expect planets to be forming or have already formed. In the case of midplane spiral structures, their origin may be linked to the presence of a companion\u2014stellar, fly-by, or planetary (Pohl et al. 2015; Bae & Zhu 2018a; Dong et al. 2018; Forgan et al. 2018b; Cuello et al. 2019; Keppler et al. 2020). Spirals may also be excited if the system is gravitationally unstable. Gravitationally instability is expected in cool and massive disks, where the disk-to-star mass ratio is larger than 0.1 (Bell et al. 1997; Gammie 2001; Lodato & Rice 2004; Cossins et al. 2009; Hall et al. 2016; Kratter & Lodato 2016; Rice 2016; Hall et al. 2019; Zhang & Zhu 2020). To date, not many spirals in dust continuum emission have a clear origin, except for those in multiple systems where the presence of spirals has been linked to stellar interactions (Kurtovic et al. 2018; Rosotti et al. 2020). On the other hand, there are disks where spirals have been reported at millimeter wavelengths and where no companion to which the spiral origin may be linked has been detected yet (to date these are Elias 27, IM Lup, WaOph 6, and MWC 758; P\u00e9rez et al. 2016; Dong et al. 2018; Huang et al. 2018c). If no companion is detected and the disk is massive compared to the host star mass, the gravitational instability (GI) scenario arises as a possible explanation for the origin of the observed spirals. Studying disks undergoing GI is important, as population synthesis models show that GI primarily ends up forming brown dwarf mass objects (Hall et al. 2017; Forgan et al. 2018a). It seems that giant-planet formation through GI is rare (Rice et al. 2015), but it may still be the formation mechanism for important systems like HR 8799 (Vigan et al. 2017).","Citation Text":["Forgan et al. 2018b"],"Functions Text":["In the case of midplane spiral structures, their origin may be linked to the presence of a companion\u2014stellar, fly-by, or planetary"],"Functions Label":["Background"],"Citation Start End":[[515,534]],"Functions Start End":[[330,460]]}
{"Identifier":"2017AandA...601A..87C__Falcke_(1996)_Instance_2","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} &amp;&amp;\\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}; \\\\ &amp;&amp;\\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)"],"Functions Text":["However, it differs from Eq. (2) in"],"Functions Label":["Differences"],"Citation Start End":[[1162,1175]],"Functions Start End":[[1126,1161]]}
{"Identifier":"2019MNRAS.482.3288G__Orazio,_Haiman_&_MacFadyen_2013_Instance_1","Paragraph":"The orbital decay of BSBHs may slow down or stall at \u223cpc scales (e.g. Begelman et al. 1980; Milosavljevi\u0107 & Merritt 2001; Zier & Biermann 2001; Yu 2002; Vasiliev, Antonini & Merritt 2014; Dvorkin & Barausse 2017; Tamburello et al. 2017), or the barrier may be overcome in gaseous environments (e.g. Gould & Rix 2000; Escala et al. 2004; Hayasaki, Mineshige & Sudou 2007; Hayasaki 2009; Cuadra et al. 2009; Lodato et al. 2009; Chapon, Mayer & Teyssier 2013; Rafikov 2013; del Valle et al. 2015), in triaxial or axisymmetric galaxies (e.g. Yu 2002; Berczik et al. 2006; Preto et al. 2011; Khan et al. 2013, 2016; Vasiliev, Antonini & Merritt 2015; Gualandris et al. 2017; Kelley, Blecha & Hernquist 2017a), and\/or by interacting with a third SMBH in hierarchical mergers (e.g. Valtonen 1996; Blaes, Lee & Socrates 2002; Hoffman & Loeb 2007; Kulkarni & Loeb 2012; Tanikawa & Umemura 2014; Bonetti et al. 2018). The accretion of gas and the dynamical evolution of BSBHs are likely to be coupled (Ivanov, Papaloizou & Polnarev 1999; Armitage & Natarajan 2002; Haiman, Kocsis & Menou 2009; Bode et al. 2010, 2012; Farris, Liu & Shapiro 2010, 2011; Kocsis, Haiman & Loeb 2012; Shi et al. 2012; D\u2019Orazio, Haiman & MacFadyen 2013; Shapiro 2013; Farris et al. 2014, 2015) such that the occurrence rate of BSBHs depends on the initial conditions and gaseous environments at earlier phases (e.g. thermodynamics of the host galaxy interstellar medium; Dotti et al. 2007, 2009; Dotti, Sesana & Decarli 2012; Fiacconi et al. 2013; Mayer 2013; Tremmel et al. 2018). Quantifying the occurrence rate of BSBHs at various merger phases is therefore important for understanding the associated gas and stellar dynamical processes. This is a challenging problem for three main reasons. First, BSBHs are expected to be rare (e.g. Foreman, Volonteri & Dotti 2009; Volonteri, Miller & Dotti 2009), and only a fraction of them accrete enough gas to be \u2018seen\u2019. Secondly, the physical separations of BSBHs that are gravitationally bound to each other (\u2272a few pc) are too small for direct imaging. Even VLBI cannot resolve BSBHs except for in the local universe (Burke-Spolaor 2011). CSO 0402+379 (discovered by VLBI as a double flat-spectrum radio source separated by 7 pc) remains the only secure case known (Rodriguez et al. 2006; Bansal et al. 2017, see Kharb, Lal & Merritt 2017; however, for a possible 0.35-pc BSBH candidate in NGC 7674). Thirdly, various astrophysical processes complicate their identification such as bright hot spots in radio jets (e.g. Wrobel, Walker & Fu 2014b). Until recently, only a handful cases of dual active galactic nuclei (AGNs) \u2013 galactic-scale progenitors of BSBHs \u2013 were known (Owen et al. 1985; Junkkarinen et al. 2001; Komossa et al. 2003; Ballo et al. 2004; Hudson et al. 2006; Max, Canalizo & de Vries 2007; Bianchi et al. 2008; Guidetti et al. 2008). While great strides have been made in identifying dual AGNs at kpc scales (e.g. Gerke et al. 2007; Comerford et al. 2009, 2012, 2015; Green et al. 2010; Liu et al. 2010, 2013, 2018; Fabbiano et al. 2011; Fu et al. 2011, 2012, 2015a,b; Koss et al. 2011, 2012, 2016; Rosario et al. 2011; Teng et al. 2012; Woo et al. 2014; Wrobel, Comerford & Middelberg 2014a; McGurk et al. 2015; M\u00fcller-S\u00e1nchez et al. 2015; Shangguan et al. 2016; Ellison et al. 2017; Satyapal et al. 2017), there is no confirmed BSBH at sub-pc scales (for recent reviews, see e.g. Popovi\u0107 2012; Burke-Spolaor 2013; Bogdanovi\u0107 2015; Komossa & Zensus 2016).","Citation Text":["D\u2019Orazio, Haiman & MacFadyen 2013"],"Functions Text":["The accretion of gas and the dynamical evolution of BSBHs are likely to be coupled","such that the occurrence rate of BSBHs depends on the initial conditions and gaseous environments at earlier phases"],"Functions Label":["Background","Background"],"Citation Start End":[[1187,1220]],"Functions Start End":[[908,990],[1262,1377]]}
{"Identifier":"2018ApJ...855...26A__Karim_et_al._2013_Instance_1","Paragraph":"The uncertainties in number counts were derived from Poisson statistics, which apply when event rates are calculated from small numbers of observed events (Gehrels 1986). Our results are given in Table 6 and plotted in Figure 7. For comparison, we also show the integral number counts from the lensing cluster surveys of Knudsen et al. (2008) and Johansson et al. (2011), from the SCUBA Half-Degree Extragalactic Survey (SHADES; Coppin et al. 2006), from LESS (Wei\u00df et al. 2009), from SCUBA-2 (Hsu et al. 2016), and from the high-resolution ALMA follow-up of LESS (Karim et al. 2013). We find that, for the intrinsic flux density range covered by our survey, results are consistent within uncertainties with previous single-dish surveys conducted in blank fields (e.g., Coppin et al. 2006), toward lensing clusters (Knudsen et al. 2008; Johansson et al. 2011), and combining both cluster and blank fields (Hsu et al. 2016). The exception is the data point at Sint = 6.3 mJy, which is fainter than the 4\u03c3 detection threshold for all clusters and therefore comprises sources that are necessarily magnified. The discrepancy is possibly due to the uncertainties in our analytical lens models, which may generally underestimate magnification factors relative to those derived from strong-lensing models, as suggested by the comparison presented in Figure 5. If magnification factors are minimally increased across the field, intrinsic flux densities and binned number counts do not vary significantly, but the detectable area Asource where \u03bc > 4\u03c3\/Sint is increased, thus affecting the number counts per unit area. To test this hypothesis, we repeated our calculations but slightly scaled our analytical magnification maps by a factor of \u223c1. We find that a satisfactory match between our resulting number counts at Sint = 6.3 mJy and previous surveys can be reached if all magnification estimates are varied by only \u223c5%, which is within the uncertainties obtained for \u03bc and listed in Table 5.","Citation Text":["Karim et al. 2013"],"Functions Text":["For comparison, we also show the integral number count","and from the high-resolution ALMA follow-up of LESS"],"Functions Label":["Uses","Uses"],"Citation Start End":[[565,582]],"Functions Start End":[[229,283],[512,563]]}
{"Identifier":"2022ApJ...936...16M__Dubrulle_et_al._1995_Instance_1","Paragraph":"Our assumption of a constant f\ndg does not account for vertical settling and depletion of dust in the disk atmosphere, which could in principle reduce the height of the different \u03c4 = 1 surfaces shown in Figure 6. Even though the Stokes number remains smaller than 1 for all considered particle sizes well above the \u03c4 = 1 surfaces for irradiation, the actual dust distribution in the atmosphere depends on the balance between turbulent stirring and vertical settling. The dust distribution away from the midplane is better captured by hydrodynamical simulations than by simple prescriptions relying on a parameterization of the turbulence strength (see, e.g., Dubrulle et al. 1995; Dullemond & Dominik 2004b), since it depends on the particular source of turbulence. An example of this is the anisotropic turbulence produced by the VSI, whose vertically elongated modes significantly increase the vertical stirring of dust with respect to models of homogeneous isotropic turbulence (Stoll & Kley 2016; Stoll et al. 2017). We note that these authors obtain a vertical Gaussian dust distribution of micron-sized particles with the same scale height as the gas in the entirety of a domain reaching z = 5H at 5 au, which supports the validity of our prescription. On the other hand, dust sedimentation of small grains in outer regions can even be affected by nonideal magnetohydrodynamical effects and magnetized winds (Riols & Lesur 2018; Booth & Clarke 2021; Hutchison & Clarke 2021). In principle, we do not expect dust settling effects to be relevant for our work as long as they occur in optically thin regions where radiative transport prevents any shadowing from occurring. Future disk models could test this hypothesis by studying the stability of disks in which the opacity is computed as a function of a dynamically evolving small dust distribution. We have also assumed instantaneous thermal equilibrium of gas and dust particles, which is not verified at high altitudes, as collisions between dust and gas particles become less frequent for smaller densities (see, e.g., Malygin et al. 2017; Pfeil & Klahr 2021). We do not expect this effect to alter our results, since at most it could increase the time it takes to form temperature perturbations in such regions.","Citation Text":["Dubrulle et al. 1995"],"Functions Text":["The dust distribution away from the midplane is better captured by hydrodynamical simulations than by simple prescriptions relying on a parameterization of the turbulence strength (see, e.g.,","since it depends on the particular source of turbulence."],"Functions Label":["Background","Background"],"Citation Start End":[[659,679]],"Functions Start End":[[467,658],[709,765]]}
{"Identifier":"2015ApJ...806...20B__Bartoli_et_al._2012a_Instance_1","Paragraph":"The ARGO-YBJ detector, hosted in a building at the YangBaJing Cosmic Ray Observatory (Tibet, China, 90\u00b031\u203250\u2033E, 30\u00b006\u203238\u2033N), 4300 m above sea level, has been designed for very high-energy (VHE) gamma-ray astronomy and cosmic-ray observations. It is made up of a single layer of resistive plate chambers (RPCs) operated in streamer mode, 2.850 m \u00d7 1.225 m each, organized in a modular configuration to cover a surface of about 5600 m2 with an active area of about 93%. The RPCs detect the charged particles in air showers with an efficiency \u226598%. To improve the shower reconstruction, other chambers are deployed around the central carpet for a total instrumented area of 100 m \u00d7 110 m. A highly segmented readout is performed by means of 55.6 cm \u00d7 61.8 cm external electrodes, called \u201cpads,\u201d whose fast signals are used for triggering and timing purposes. These pads provide the digital readout of the detector up to 22 particles m\u22122, allowing the count of the air shower charged particles without any significant saturation up to primary cosmic-ray energies of about 200 TeV (Bartoli et al. 2012a). In order to extend the dynamical range to PeV energies each RPC is also equipped with two large size pads (139 cm \u00d7 123 cm), allowing the collection of the total charge developed by the particles hitting the detector (Aielli et al. 2012). The digital output of each pad is splitted in two signals sent to the logic chain that builds the trigger and to the 18,360 multi-hit time-to-digital converters, which are routinely calibrated with 0.4 ns accuracy by means of an off-line method using cosmic-ray showers (He et al. 2007; Aielli et al. 2009a). More details about the detector and the RPC performance can be found in Aielli et al. (2006, 2009b). The detector is connected to two independent acquisition systems corresponding to two different operation modes, referred to as the shower mode and the scaler mode (Aielli et al. 2008). The data used in this paper were recorded by the digital readout in shower mode. This mode is implemented by means of an inclusive trigger based on the time correlation between the pad signals, depending on their relative distance. In this way the data acquisition is triggered when at least 20 pads in the central carpet are fired in a time window of 420 ns. By means of this trigger the energy threshold for gamma-induced showers can go down to 300 GeV with an effective area depending on the zenith angle (see Figure 1 in Bartoli et al. 2013).","Citation Text":["Bartoli et al. 2012a"],"Functions Text":["These pads provide the digital readout of the detector up to 22 particles m\u22122, allowing the count of the air shower charged particles without any significant saturation up to primary cosmic-ray energies of about 200 TeV"],"Functions Label":["Background"],"Citation Start End":[[1077,1097]],"Functions Start End":[[856,1075]]}
{"Identifier":"2021AandA...650A.164M__Davies_et_al._2012_Instance_3","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"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[2062,2080]],"Functions Start End":[[1887,2060]]}
{"Identifier":"2022MNRAS.513.5377F__Fulle_et_al._2020b_Instance_1","Paragraph":"Cometary activity is driven by the gas pressure P(s), which does not depend on the ice abundance, but only on the gas temperature: also minor species can drive cometary activity, provided that heat transfer inside a pebble is faster than ice depletion. This condition fixes an upper limit for the refractory-to-ice mass ratio \u03b4i (Fulle 2021)\n(8)$$\\begin{eqnarray*}\r\n\\delta _i \\lt {\\lambda _s \\over {3 ~Q ~R ~c_p}} - 1\r\n,\r\n\\end{eqnarray*}$$where cp \u2248 103 J kg\u22121 K\u22121 is the heat capacity of the pebbles (Blum et al. 2017). In order to fulfil equation (2) during the whole inbound orbit, cometary activity is driven by at least the five ices listed in Table 2, reporting the nucleus and gas coma parameters computed by means of the activity model (Fulle et al. 2020b). Thus, the less abundant the ice, the shorter the rh-range where it can drive cometary activity (e.g. ethane in Table 2). The upper limit of each water-to-ice mass ratio is \u03b4i\/\u03b4w, and in the protoplanetary disc \u03b4w \u2248 5 (Cambianica et al. 2020). The probable activity due to further ices not listed in Table 2 and the decrease of cp with the temperature (Takahashi & Westrum 1970; Shulman 2004; Bouziani & Jewitt 2022) significantly increase the upper limits of \u03b4i. However, e.g. formaldehyde has unknown thermodynamical parameters (Fray & Schmitt 2009), and e.g. ethylene and nitric oxide were not detected in 67P (Rubin et al. 2020), thus preventing a realistic computation of e.g. evolving cp values. According to the activity model, the distribution of water-ice is very inhomogeneous in the nucleus (Ciarniello et al. 2022), with water-rich pebbles, depleted of supervolatiles, embedded in a matrix of water-poor pebbles, which are rich of the ices listed in Table 2. Therefore, at rh > 3.8 au, dust is ejected from water-poor pebbles only, characterized by \u03b4w \u2248 50 (Fulle 2021), so that equation (2) is surely verified. In water-poor pebbles the water ice is less abundant than all other ices (excluded ethane, Table 2), so that these ices cannot be trapped inside a less abundant water ice. The observed activity of C\/2017K2 at rh = 23.7 au (Jewitt et al. 2017) excludes also that ices may be trapped inside CO2-ice, because the CO2-driven activity onsets at rh = 13 au (Table 2).","Citation Text":["Fulle et al. 2020b"],"Functions Text":["In order to fulfil equation (2) during the whole inbound orbit, cometary activity is driven by at least the five ices listed in Table 2, reporting the nucleus and gas coma parameters computed by means of the activity model"],"Functions Label":["Uses"],"Citation Start End":[[745,763]],"Functions Start End":[[521,743]]}
{"Identifier":"2016MNRAS.457.3191D__Martin_et_al._2005_Instance_1","Paragraph":"TYC 9486-927-1 was observed by Torres et al. (2006) as part of the Search for Associations Containing Young stars (SACY) programme (Torres et al. 2008). They assigned a spectral type of M1 and measured a radial velocity of vrad = 8.7 \u00b1 4.6 km s\u22121 from 10 observations. The large uncertainty is likely due to the star's high rotational velocity (vsin\u2009i = 43.5 \u00b1 1.2 km s\u22121); suggesting it is either a single rapid rotator or a spectroscopic binary with blended lines. TYC 9486-927-1 also shows signs of activity in X-ray (Thomas et al. 1998), H \u03b1 emission (Torres et al. 2006) and the UV (using GALEX data from Martin et al. 2005 we find log\u2009FFUV\/FJ = \u22122.49, log\u2009FNUV\/FJ = \u22122.11). 2MASS J2126\u22128140 is an L3 first identified by Reid et al. (2008, although referencing Cruz, Kirkpatrick & Burgasser 2009 as the discovery paper). Subsequently Faherty et al. (2013) classified it as a low gravity L3\u03b3 (using the gravity classification system of Cruz et al. 2009). Recent VLT\/ISAAC observations by Manjavacas et al. (2014) find it is a good match to the young, L3 companion CD-35 2722B (Wahhaj et al. 2011). These authors also used the spectral indices of Allers & Liu (2013) to confirm that the 2MASS J2126\u22128140 is an L3 and shows low gravity spectral features. Manjavacas et al. (2014) also used the BT-Settl-2013 atmospheric models (Allard, Homeier & Freytag 2012) to derive Teff = 1800 \u00b1 100 K, log\u2009g = 4.0 \u00b1 0.5 dex, albeit with better fits to supersolar metallicity models. Filippazzo et al. (2015) use photometry, a trigonometric parallax of 31.3 \u00b1 2.6 mas (referenced to Faherty et al., in preparation) and evolutionary models to derive an effective temperature of 1663 \u00b1 35 K. They also derived a mass of 23.80 \u00b11 5.19 MJ assuming a broad young age range of 10\u2013150 Myr. Gagn\u00e9 et al. (2014) listed 2MASS J2126\u22128140 as a high probability candidate member of Tucana\u2013Horologium association (TucHor) but noted that its photometric distance would be in better agreement with its TucHor kinematic distance if it were an equal mass binary.","Citation Text":["Martin et al. 2005"],"Functions Text":["using GALEX data from","we find log\u2009FFUV\/FJ = \u22122.49, log\u2009FNUV\/FJ = \u22122.11"],"Functions Label":["Uses","Uses"],"Citation Start End":[[610,628]],"Functions Start End":[[588,609],[629,677]]}
{"Identifier":"2015AandA...584A.103S__Potekhin_et_al._2013_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":["Potekhin et al. 2013"],"Functions Text":["Analytical fits of these neutron-star EoSs have been constructed in order to facilitate their inclusion in astrophysical simulations"],"Functions Label":["Background"],"Citation Start End":[[1209,1229]],"Functions Start End":[[1075,1207]]}
{"Identifier":"2017AandA...600A.123D__Coupeaud_et_al._(2011)_Instance_1","Paragraph":"Despite the differences of the dust analogues retained in previous studies (Mennella et al. 1998; Boudet et al. 2005; Coupeaud et al. 2011) and in the present one, the spectroscopic properties and temperature dependent behavior of all these analogues remain qualitatively similar: the MAC value is correlated with the dust temperature and has a complex spectral shape differing from a simple asymptotic behavior in \u03bb-2. At the same time, because of these differences between the samples in terms of internal structure, composition, porosity, homogeneity, the shape of the MAC in the FIR varies. Consequently, the MAC value at a given wavelength of all samples from this study and from previous studies spans a large range of values (see Table 1 for the MAC value of the samples studied here and Table 1 from Demyk et al. 2013, for previous studies at 1 mm). It must be noted that the MAC in Boudet et al. (2005) has been corrected for grain shape effect whereas in Coupeaud et al. (2011) and Mennella et al. (1998) and in the present study this correction was not applied (see Sect. 2.3). This partly explains why the MAC value are smaller in Boudet\u2019s study, together with the effect of differences in the samples in terms of composition and structure. At 100 \u03bcm, the MAC values of all the samples from this study and from previous work are similar at 300 and 10 K (50\u2013300 cm2\u2009g-1 at 300 K and 40\u2013260 cm2\u2009g-1 at 10 K). At 500 \u03bcm, the MAC is greater at 300 K than at 10 K (2\u201320 cm2\u2009g-1 at 300 K and 0.7\u201313 cm2\u2009g-1 at 10 K). This is also the case at 1 mm where the MAC is in the range 0.1\u201311 cm2\u2009g-1 at 300 K and 0.12\u20137 cm2\u2009g-1 at 10 K. We note that for each sample the MAC at 300 K is always greater than at 10 K. Furthermore, the fact that the range of values of the MAC at 300 K and 10 K overlaps results from the dispersion of the measured spectra. This diversity of the MAC value and spectral shape that vary as a function of the materials studied can be problematic for astrophysical applications since one has to decide what analog is the most relevant and should be used in the modeling. This is considered in more detail in Sect. 5. ","Citation Text":["Coupeaud et al. (2011)"],"Functions Text":["whereas in","and Mennella et al. (1998) and in the present study this correction was not applied"],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[965,987]],"Functions Start End":[[954,964],[988,1071]]}
{"Identifier":"2020ApJ...905..111Z__Jiri\u010dka_et_al._2001_Instance_3","Paragraph":"Surveys of radio bursts in decimetric wavelengths is presented in papers by Isliker & Benz (1994) and Jiri\u010dka et al. (2001), within 1\u20133 GHz and 0.8\u20132.0 GHz frequency ranges, respectively. Some of these bursts are still not well understood. This is a case of the slowly positively drifting bursts (SPDBs). They appear in groups or as single bursts, with a duration of an individual burst from 1 to several seconds and their frequency drift is lower than about 100 MHz s\u22121 (Jiri\u010dka et al. 2001). The SPDBs seem to be similar to the reverse type III bursts (Aschwanden 2002) but their frequency drift is much smaller. The majority of observed SPDBs are connected to solar flares (Jiri\u010dka et al. 2001), and they appear many times at the very beginning of the flares (Benz & Simnett 1986; Kotr\u010d et al. 1999; Kaltman et al. 2000; Karlick\u00fd et al. 2018). Kaltman et al. (2000) reported on several SPDBs observed during three solar flares in the 0.8\u20132 GHz frequency range. They found frequency drifts of the observed SPDBs to be within the 20\u2013180 MHz s\u22121 range. Kotr\u010d et al. (1999) studied one of those flares. By combining the radio and spectral plus imaging H\u03b1 observations, they explained the observed SPDBs as radio emission generated by downwards propagating shock waves. Based on numerical simulations of the formation of thermal fronts in solar flares, Karlick\u00fd (2015) proposed that SPDBs observed in the 1\u20132 GHz range could be a signature of a thermal front. Furthermore, Karlick\u00fd et al. (2018) reported the observation of an SPDB (1.3\u20132.0 GHz) observed during the impulsive phase of an eruptive flare. They found time coincidence between the SPDB occurrence, an appearance of an ultraviolet (UV)\/EUV multithermal plasma blob moving down along the dark H\u03b1 loop at approximately 280 km s\u22121, and the observed change of H\u03b1 profile at the footpoint of that dark loop. Combining these observations they concluded that observed SPDB was likely generated by the thermal front formed in front of the falling EUV blob.","Citation Text":["Jiri\u010dka et al. 2001"],"Functions Text":["The majority of observed SPDBs are connected to solar flares"],"Functions Label":["Background"],"Citation Start End":[[677,696]],"Functions Start End":[[615,675]]}
{"Identifier":"2020ApJ...896...12X__Rosenberg_&_Coleman_1969_Instance_1","Paragraph":"The QBOs are also found in the local wavelet power spectra of the Bx, By, and Bz that are displayed in the left panel of Figure 2. As this figure shows, the significant regions of above 95% confidence level in the range of 256\u2013512 days (peak at 1.01 yr) for Bx intermittently appear during 1975\u20132000, and those in the range of 256\u2013600 days (peak at 1.02 yr) for By are also intermittently found during 1975\u20132000. The QBOs for Bz only appear in solar cycle 22, which corresponds to the significant period of 0.90 yr, and also peak at 1.58 and 2.77 yr. The possible period of 3.45 yr for Bx does not show the corresponding significant region in the local wavelet power spectrum. The significant region that corresponds to the possible period of 2.37 yr for By appears near 1992. Early study proved that the solar QBOs can be transmitted into interplanetary space by the open magnetic flux (Lockwood 2001; Bazilevskaya et al. 2014); thus, the QBOs found in the Bx, By, and Bz components should be related to the solar QBOs. On the other hand, the annual variation in the IMF polarity in the band of 256\u2013512 days, which is referred to the Rosenberg\u2013Coleman effect (Rosenberg & Coleman 1969), only appears in the ascending phase of solar cycles by studying the IMF polarity date for the years 1927\u20132002 through the wavelet analysis method (Echer & Svalgaard 2004). The authors indicated that such a result should be attributed to a more stable and flat heliospheric current sheet that only appears in the ascending phase of solar cycles and strong disturbance of the heliospheric current sheet that is present in the declining phase and minimum time of solar cycles, and the Rosenberg\u2013Coleman effect had been confirmed by the IMF Bx component during 1964\u20132002. The wavelet power spectra of the IMF Bx and By components in Figure 2 show that the annual variation in two time series can also be found in some, but not all, ascending phases of solar cycles, which is partly similar to the wavelet map of the IMF polarity in Echer & Svalgaard (2004). Thus, the annual variation in the IMF Bx and By components is also modulated by the stability or strong disturbance of the heliospheric current sheet. However, the wavelet power spectra of the IMF Bz show a different result, which indicates that the annual variation in this time series should not be related to the heliospheric current sheet.","Citation Text":["Rosenberg & Coleman 1969"],"Functions Text":["On the other hand, the annual variation in the IMF polarity in the band of 256\u2013512 days, which is referred to the Rosenberg\u2013Coleman effect","The authors indicated that such a result should be attributed to a more stable and flat heliospheric current sheet that only appears in the ascending phase of solar cycles and strong disturbance of the heliospheric current sheet that is present in the declining phase and minimum time of solar cycles, and the Rosenberg\u2013Coleman effect had been confirmed by the IMF Bx component during 1964\u20132002."],"Functions Label":["Uses","Compare\/Contrast"],"Citation Start End":[[1161,1185]],"Functions Start End":[[1021,1159],[1360,1755]]}
{"Identifier":"2022MNRAS.509..314F__Spitkovsky,_Levin_&_Ushomirsky_2002_Instance_1","Paragraph":"Magnetohydrodynamic (MHD) shallow water equations are an alternative to complete system of MHD equations for plasma. In this approximation, a thin layer of plasma with a free boundary in a gravity field is studied (Gilman 2000). Shallow water flows with rotation are usually considered in case of large-scale flows in plasma astrophysics (Gilman 1967) and play an important role in understanding of various astrophysical objects. Such a model is used to study large-scale flows of the solar tachocline (a thin layer inside the Sun, located above the convective zone) (Gilman 2000; Dikpati & Gilman 2001; Miesch & Gilman 2004; Hughes, Rosner & Weiss 2007), flows of accreting matter in neutron stars (Inogamov & Sunyaev 1999, 2010; Spitkovsky, Levin & Ushomirsky 2002), dynamics of the atmospheres of neutron stars (Heng & Spitkovsky 2009), and magnetoactive tidally locked atmospheres of exoplanets (Cho 2008; Heng & Showman 2015; Batygin, Stanley & Stevenson 2017; Pierrehumbert & Hammond 2019). MHD shallow water equations are obtained by depth-averaging three-dimensional MHD equations of plasma layer with respect to pressure being hydrostatic and the height of the layer being much smaller than characteristic horizontal scale of motion. However, the depth-averaging procedure does not always lead to shallow water equations. Although the equations obtained by integrating over the height of the studied layer of the liquid will have a form similar to the equations of shallow water (Klimachkov & Petrosyan 2017b).1 Equations obtained in this paper are similar to MHD shallow water equations, but they have additional terms and equations referred to vertical magnetic field (Klimachkov & Petrosyan 2017a; Fedotova, Klimachkov & Petrosyan 2020). These new equations transform into common MHD shallow water equations (Gilman 2000) in case of zero vertical magnetic field. Due to this, hereinafter we name our approximation as quasi-two-dimensional MHD approximation and equations are named as quasi-two-dimensional equations for rotating astrophysical plasma with vertical magnetic field.2 Detailed review of wave processes in magnetohydrodynamics of astrophysical plasma with vertical magnetic field is presented in Petrosyan et al. (2020). Extensions of magnetohydrodynamics in plasma astrophysics for vertically stratified flows with vertical magnetic field can be found in Fedotova, Klimachkov & Petrosyan (2021).","Citation Text":["Spitkovsky, Levin & Ushomirsky 2002"],"Functions Text":["Such a model is used to study","flows of accreting matter in neutron stars"],"Functions Label":["Background","Background"],"Citation Start End":[[731,766]],"Functions Start End":[[430,459],[656,698]]}
{"Identifier":"2020ApJ...904..119F__Raaijmakers_et_al._2019_Instance_1","Paragraph":"In this work, the parameter ranges are for 68% credibility interval unless specifically mentioned. With Equation (1), the amount of the post-merger gravitational radiation can be reasonably\/qualitatively evaluated as long as \n\n\n\n\n\n (i.e., the EoS) is known. However, various EoS models have been proposed in the literature and it is not possible to be uniquely determined even in the foreseeable future. Fortunately, under the reasonable assumption that all NSs follow the same EoS, their properties can be jointly\/reliably constrained with the nuclear data, the GW data, the measured masses, and the estimated radii of some NSs (e.g., Lattimer & Prakash 2016; Abbott et al. 2017a; Tews et al. 2017; Most et al. 2018; Landry & Essick 2019). The masses of NSs in some binary systems have been accurately measured and there is a robust lower limit on MTOV \u2265 2 M\u2299 (Cromartie et al. 2020; Kandel & Romani 2020). The radii of NSs, however, usually are just evaluated indirectly and suffer from large systematical uncertainties. Thanks to the successful performance of the Neutron Star Interior Composition Explorer, the situation has changed and very recently the first-ever accurate measurement of mass and radius together for PSR J0030+0451, a nearby isolated quickly rotating NS, has been achieved (Miller et al. 2019; Riley et al. 2019), which favor a stiffer EoS than the data of GW170817. Hence, GW170817, PSR J0030+0451, some nuclear data as well as the lower limit on MTOV can be combined to reliably constrain the EoS as well as the bulk properties of NSs. This can be done either in the EoS parameterizing methods (Raaijmakers et al. 2019; Jiang et al. 2020) or the nonparametric approach (Essick et al. 2020; Landry et al. 2020), and the results are well consistent with each other. Here we directly adopt posterior samples of {MA, MB, \u039bA, \u039bB} obtained in Jiang et al. (2020) to calculate the \n\n\n\n\n\n for GW170817. Note that the region of \n\n\n\n\n\n represents the prompt black hole formation, which is irrelevant to GW170817 because of the delayed collapse of the remnant (Metzger 2019). So we neglect the posterior samples that give \n\n\n\n\n\n for GW170817. At 90% credible level, for the piecewise polytropic expansion method we have \n\n\n\n\n\n, while for the spectral decomposition method we have \n\n\n\n\n\n. We also adopt the method described in Kumar & Landry (2019) to evaluate the nonparametric posterior of PSRs+GWs+x-ray\/Riley case in Landry et al. (2020) and get \n\n\n\n\n\n for GW1708017. The incorporation of the strong phase transition possibility by Landry et al. (2020) favors lower k2T than ours (note that the \n\n\n\n\n\n region is excluded), but has an overall agreement with the piecewise result and the spectral result (as shown in upper panel of Figure 1). For this reason we combine an equal sample of k2T calculated from these three different parameterization methods to perform all the calculations in this work unless specially specified. Clearly, the inferred \n\n\n\n\n\n is well within the region that predicts the very prominent post-merger GW radiation, which is rather encouraging. The post-merger energy is then estimated to be \n\n\n\n\n\n (90% confidence level) when the fitting error is considered for the combined posterior (see the upper panel of Figure 2). For B1534+12, B2127+11C, J1757-1854, and J0453+1559, we expect that they would emit almost the same amount of energy as GW170817 in the post-merger phase. This is understandable considering the comparable total mass of these systems. For lighter BNS systems like J0514-4002A, the expected post-merger energy is relatively small because of the large \n\n\n\n\n\n. While for heavier BNS systems like GW190425, we predict a prompt collapse scenario and thus emit a small amount of energy in the post-merger phase (see the lower panels of Figures 1 and 2). To our knowledge, ours is the first study to combine the intriguing numerical finding of Zappa et al. (2018) with the EoS constrained with the multimessenger information of NSs and then to demonstrate that GW170817 is likely the most efficient post-merger GW emitter among the observed BNS GW events.","Citation Text":["Raaijmakers et al. 2019"],"Functions Text":["This can be done either in the EoS parameterizing methods","or the nonparametric approach","and the results are well consistent with each other."],"Functions Label":["Similarities","Similarities","Similarities"],"Citation Start End":[[1620,1643]],"Functions Start End":[[1561,1618],[1664,1693],[1736,1788]]}
{"Identifier":"2017MNRAS.470.4075L__Janev_&_Reiter_2004_Instance_1","Paragraph":"In order to determine the unknown rate constant that will be used in the network, we used a methodology developed in previous articles (Loison et al. 2014a,b, 2015) and summarized in Appendix A. This methodology includes an extensive literature review, various Density Functional Theory (DFT) and ab initio calculations for critical gas phase reactions, namely H + l-C3H2, H + t-C3H2, O + c-C3H2, O + l-C3H2, N + C3, N + c-C3H2, N + l-C3H2, H + C3O, O + C3O, O + c-C3H3+, OH + C3, OH + c-C3H2, H2 + l-C3H and H2 + c-C3H, to determine the presence, or not, of a barrier. When there is no barrier in the entrance valley and exothermic bimolecular exit channels, we chose to use the capture rate constant or sometimes a fraction of the capture rate constant by comparison with similar reactions (the capture rate is the upper limit of the rate constants for barrierless reactions). The dissociative recombination (DR) of c,l-C3H2+ and c,l-C3H3+ is an important source of c,l-C3H and is the main source c,l,t-C3H2 in our network. The first step of c,l-C3H2+ and c,l-C3H3+ DR is the formation of highly excited C3H2** and C3H3**, which leads to bond fragmentation. Angelova et al. (2004) have shown that the DR of c-C3H2+ leads to 87.5 per cent of C3Hx and 12.5 per cent of C2Hy + CHz, and the DR of c-C3H3+ leads to 90.7 per cent of C3Hx and 9.3 per cent of C2Hy + CHz. Moreover, in DR processes, the H ejection is in general favoured more than H2 ejection (Plessis et al. 2010, 2012; Janev & Reiter 2004). Considering the exothermicity for the ejection of two hydrogen atoms (endothermic for c-C3H3+ DR and only slightly exothermic for l-C3H2+, c-C3H2+ and l-C3H3+ DR, see Appendix B), this process will have a low branching ratio. Then, the dissociation of C3H2** and C3H3** will produce mainly C3H + H and C3H2 + H, both C3H and C3H2 species being also excited considering the exothermicity of the DR and the fact that the hydrogen atom will carry only a limited part of the available energy through kinetic energy. Part of the excited C3H* and C3H2* molecules will lead to dissociation when they are populated above the dissociation limit, but most of them will relax through radiative emission of an infrared photon. As noted by Herbst et al. (2000), the typical time-scale for isomeric conversion is much shorter than that for relaxation by one infrared photon. Thus, because radiative relaxation occurs slowly, isomeric conversion leads to equilibrated isomeric (c-C3H $\\leftrightarrows $ l-C3H, c-C3H2$\\leftrightarrows $ l-C3H2) abundances at each internal energy. The final balance is determined at or near the effective barrier to isomerization, which corresponds to the energy of the transition state. The ratio between the isomeric forms are then approximated by the ratio of the rovibrational densities of states of the isomers at the barrier to isomerization calculated using the MESMER program (Glowacki et al. 2012). Fig. 2 shows the isomerization pathway calculated at the DFT level. The calculated geometries of the stationary point can be found in Appendix A. The t-C3H2 has a triplet ground state and its production in the excited singlet ground state is neglected here. The production of t-C3H2 from the DR of c,l-C3H3+ is supposed to come from c,l-C3H3+ + e\u2212 \u2192 c,l-C3H3 \u2192 t-C3H2 + H, which is supposed to be a minor channel in comparison to c,l-C3H3+ + e\u2212 \u2192 c,l-C3H3 \u2192 c,l-C3H2 + H (t-C3H2 has a ground triplet state, in contrast to c,l-C3H2, which has a singlet ground state).","Citation Text":["Janev & Reiter 2004"],"Functions Text":["Moreover, in DR processes, the H ejection is in general favoured more than H2 ejection"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1481,1500]],"Functions Start End":[[1366,1452]]}
{"Identifier":"2016ApJ...831...63X__Bergh_2009_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":["van den Bergh 2009"],"Functions Text":["Currently, there are two possible scenarios on the origin of S0 galaxies.","The other is that S0 galaxies are intrinsically different from spiral galaxies since their formation"],"Functions Label":["Background","Background"],"Citation Start End":[[563,581]],"Functions Start End":[[209,282],[414,514]]}
{"Identifier":"2016MNRAS.456..512C__Kronberg_et_al._2004_Instance_1","Paragraph":"Extended radio emission in galaxies is associated with both radio jets and lobes and with outflows, seen often as aligned radio sources in the opposite directions with respect to the central compact radio core. Giant radio galaxies (GRG) are extreme cases of this phenomenology with jets and lobes extending on \u223c Mpc scales suggesting that they are either very powerful or very old site for electron acceleration. In this respect, GRGs have a crucial role in the acceleration of cosmic rays over large cosmic scales (e.g. Kronberg et al. 2004), in the feedback mechanism of AGNs into the intergalactic and intracluster medium (e.g. Subrahmanyan et al. 2008) and in the seeding of large-scale magnetic fields in the universe (e.g. Kronberg et al. 2004) and they are excellent sites to determine the total jet\/lobe energetics in AGN-dominated structures (see e.g. Colafrancesco 2008, Colafrancesco & Marchegiani 2011). To date our knowledge of GRGs (see e.g. Ishwara-Chandra & Saikia 1999, 2002; Lara et al. 2001; Machalski, Jamrozy & Zola 2001; Schoenmakers et al. 2001; Kronberg et al. 2004; Saripalli et al. 2005; Malarecki et al. 2013; Butenko et al. 2014) is limited by their sparse numbers and by the difficulty of detecting them over large areas of the sky. Low-frequency radio observations have an enhanced capacity to detect the extended old electron population in these objects (see e.g. the recent Low Frequency Array \u2013 LOFAR \u2013 observation of the GRG UGC095551), but high-frequency radio observations are less efficient in this task due to the steep-spectra of giant radio lobes. In this context these sources will be ideal targets for the next coming deep, wide-field surveys like, e.g. the ATLAS survey of the Australia Telescope Network Facility (ATNF; see Norris et al. 2009) or the Square Kilometre Array (SKA) deep surveys that will have the potential to study their population evolution up to high redshifts and thus clarifying their role on the feedback for the evolution of non-thermal processes in large-scale structures.","Citation Text":["Kronberg et al. 2004"],"Functions Text":["In this respect, GRGs have a crucial role in the acceleration of cosmic rays over large cosmic scales (e.g."],"Functions Label":["Motivation"],"Citation Start End":[[522,542]],"Functions Start End":[[414,521]]}
{"Identifier":"2018ApJ...852..112K__Neronov_&_Aharonian_2007_Instance_1","Paragraph":"The Virgo Cluster radio galaxy M87 (NGC 4486), located at a distance of \n\n\n\n\n\n Mpc (Mei et al. 2007) and believed to harbour a BH of mass \n\n\n\n\n\n, was the first extragalactic source detected at VHE energies (Aharonian et al. 2003). Given its proximity, M87 has been a prime target to probe scenarios for the formation of relativistic jets with high-resolution radio observations exploring scales down to some tens of rg, and much effort has recently been dedicated in this direction (e.g., Acciari et al. 2009; Doeleman et al. 2012; Hada et al. 2014, 2016; Akiyama et al. 2015, 2017; Kino et al. 2015). At VHE energies, M87 has revealed at least three active \u03b3-ray episodes, during which day-scale flux variability (i.e., \n\n\n\n\n\n) has been observed (Aharonian et al. 2006; Albert et al. 2008; Acciari et al. 2009; Abramowski et al. 2012; Aliu et al. 2012). The VHE spectrum is compatible with a relatively hard power law (photon index \u223c2.2) extending from 300 GeV to beyond 10 TeV, while the corresponding TeV output is relatively moderate, with an isotropic equivalent luminosity of \n\n\n\n\n\n erg s\u22121. The inner, parsec-scale jet in M87 is considered to be misaligned by \n\n\n\n\n\n\u201325\u00b0, resulting in modest Doppler boosting of its jet emission and creating challenges for conventional jet models to account for the observed VHE characteristics (see, e.g., Rieger & Aharonian 2012 for review and references). Gap-type emission models offer a promising alternative and different realizations have been proposed in the literature (e.g., Neronov & Aharonian 2007; Levinson & Rieger 2011; Broderick & Tchekhovskoy 2015; Vincent 2015; Ptitsyna & Neronov 2016). M87 is overall highly underluminous with characteristic estimates for its total nuclear (disk and jet) bolometric luminosity not exceeding \n\n\n\n\n\n erg s\u22121 by much (e.g., Owen et al. 2000; Whysong & Antonucci 2004; Prieto et al. 2016), suggesting that accretion onto its BH indeed occurs in a non-standard, advective-dominated (ADAF) mode characterized by an intrinsically low radiative efficiency (e.g., Di Matteo et al. 2003; Nemmen et al. 2014), with inferred accretion rates possibly ranging up to \n\n\n\n\n\n (e.g., Levinson & Rieger 2011) and a BH spin parameter close to its maximum one (e.g., Feng & Wu 2017). For these values of the accretion rate, the soft photon field (see Equations (2) and (3)) is sufficiently sparse that the maximum Lorentz factor \n\n\n\n\n\n of the magnetospheric particles is essentially determined by the curvature mechanism. The observed VHE variability is in principle compatible with \n\n\n\n\n\n, so that the different dependence of the gap power on \u03b2, Equation (20), does not necessarily (in the absence of other, intrinsic considerations of gap closure) imply a strong difference in the extractable gap powers. Figure 2 shows a representative point for M87 (taking \n\n\n\n\n\n). The observed VHE luminosity of M87 is some orders of magnitudes lower than the maximum possible gap power (given by the dotted line) and within the bound imposed by ADAF considerations (vertical line). The observed VHE flaring events thus appear consistent with a magnetospheric origin. VLBI observations of (delayed) radio core flux enhancements indeed provide support for the proposal that the variable VHE emission in M87 originates at the jet base very near to the BH (e.g., Acciari et al. 2009; Beilicke 2012; Hada et al. 2012, 2014).","Citation Text":["Neronov & Aharonian 2007"],"Functions Text":["Gap-type emission models offer a promising alternative and different realizations have been proposed in the literature (e.g.,"],"Functions Label":["Motivation"],"Citation Start End":[[1526,1550]],"Functions Start End":[[1400,1525]]}
{"Identifier":"2020MNRAS.494.2969T__Catelan_et_al._2001_Instance_1","Paragraph":"In the tidal torque model used in this work, the alignment of spiral galaxies is purely due to their orientation, which in turn is related to the angular momentum correlation of neighbouring galaxies relative to the line of sight (Croft & Metzler 2000; Crittenden et al. 2001). Angular momentum correlations are mainly build up at early times during structure formation and are thus due to initial correlations (Catelan & Theuns 1996, 1997; Theuns & Catelan 1997). The correlated angular momenta result into to correlated inclination angles of neighbouring galaxies and thus ultimately into correlated ellipticities (Catelan, Kamionkowski & Blandford 2001). Assuming that the symmetry axis of the galactic disc coincides with the direction of the angular momentum $\\hat{L} = \\boldsymbol{L}\/L$, the ellipticity can be written as the alignment of spiral galaxies is purely due to the orientation of their circular discs, which in turn is related to the angular momentum correlation of neighbouring galaxies relative to the line of sight (Croft & Metzler 2000; Crittenden et al. 2001). Angular momentum correlations are mainly build up at early times during structure formation and are thus due to initial correlations (Catelan & Theuns 1996, 1997; Theuns & Catelan 1997). The correlated angular momenta result into correlated inclination angles of neighbouring galaxies and thus ultimately into correlated ellipticities (Catelan et al. 2001). Assuming that the symmetry axis of the galactic disc coincides with the direction of the angular momentum $\\hat{L} = \\boldsymbol{L}\/L$, the ellipticity can be written as\n(17)$$\\begin{eqnarray*}\r\n\\epsilon = \\frac{\\hat{L}^2_x - \\hat{L}^2_y}{1+ \\hat{L}^2_z} + 2\\mathrm{i}\\frac{\\hat{L}_x\\hat{L}_y}{1+\\hat{L}^2_z}\\,\\, .\r\n\\end{eqnarray*}$$Angular momentum is generated by a torque exerted by the ambient large-scale structure on to the protogalactic halo, a mechanism called tidal torquing (White 1984; Barnes & Efstathiou 1987; Schaefer 2009; Stewart et al. 2013). For Gaussian random fields, the autocorrelation of angular momenta is given by (Lee & Pen 2001)\n(18)$$\\begin{eqnarray*}\r\n\\left\\langle \\hat{L}_\\alpha \\hat{L}_\\beta \\right\\rangle = \\frac{1}{3}\\left(\\frac{1+ A}{3}\\delta _{\\alpha \\beta } - A \\hat{\\Phi }_{\\alpha \\mu }\\hat{\\Phi }_{\\mu \\beta } \\right)\\,\\, .\r\n\\end{eqnarray*}$$The free parameter A determines the strength of the coupling between alignment and tidal torque. Since the correlation is determined by the traceless part of the shear tensor $\\hat{\\Phi }_{\\alpha \\beta }$ the resulting effect is clearly due to orientation effects only. For a Gaussian distribution $p(\\hat{L}|\\hat{\\Phi }_{\\alpha \\beta })\\mathrm{d}\\hat{L}$ and the use of equation (17), one can express the ellipticity in terms of the tidal field\n(19)$$\\begin{eqnarray*}\r\n\\epsilon (\\hat{\\Phi }) = \\frac{A}{2} \\left(\\hat{\\Phi }_{x\\alpha }\\hat{\\Phi }_{\\alpha x} - \\hat{\\Phi }_{y\\alpha }\\hat{\\Phi }_{\\alpha y} -2\\mathrm{i}\\hat{\\Phi }_{x\\alpha }\\hat{\\Phi }_{\\alpha y}\\right).\r\n\\end{eqnarray*}$$Correlations in the ellipticities can thus be traced back to the four-point function of the shear field, which is given in equation (11). For keeping a correct relative normalization of the shape correlations, we scale the resulting angular ellipticity spectra $C^{\\mathrm{s},II}_{ij}(\\ell)$ with the squared number of spiral galaxies $n_s^2$. It is remarkable that the shapes of spiral galaxies in the quadratic alignment model are in fact sensitive to tidal shear components parallel to the line of sight; in fact, those components determine the magnitude of the alignment effect, in contrast to the alignment of elliptical galaxies in the linear alignment model or to gravitational lensing, which reflect purely the tidal shear components perpendicular to the line of sight.","Citation Text":["Catelan et al. 2001"],"Functions Text":["The correlated angular momenta result into correlated inclination angles of neighbouring galaxies and thus ultimately into correlated ellipticities"],"Functions Label":["Uses"],"Citation Start End":[[1419,1438]],"Functions Start End":[[1270,1417]]}
{"Identifier":"2020MNRAS.498.6069P__Casertano_&_Hut_1985_Instance_1","Paragraph":"Galaxy colour is known to be sensitive to local density. Generally, sheets are denser than fields and filaments are denser than sheets. So the dependence of red and blue fractions on the geometry of large-scale environment shown in Fig. 5 may partly arise due to dependence of galaxy colour on local density. We need to decorrelate the effect of local density in order to test the role of large-scale structures on galaxy colours. We address this issue by calculating local number density of galaxies using kth nearest neighbour method (Casertano & Hut 1985). We compute the local number density using equation (2) with k = 5. We plot the local density against the local dimension of galaxies in three different volume-limited samples in Fig. 7. The top left, middle left, and bottom left panels of Fig. 7 show the relations between local dimension and local density for the three volume-limited samples when local dimensions are computed using $R_2=10 {\\, h^{-1}\\, {\\rm Mpc}}$. The results in these panels show that the environments with larger local dimension indeed tend to have a lower local density. However, these relationships show very large scatters. The environments with local dimension up to D = 2.5 can have a wide range of local densities and it is difficult to assign a specific density range to the environments with different local dimensions. The three panels in the middle column and three panels in the right column of Fig. 7 show the relations between local density and local dimension in these samples for $R_2=40 {\\, h^{-1}\\, {\\rm Mpc}}$ and $R_2=70 {\\, h^{-1}\\, {\\rm Mpc}}$. They show a similar trend as seen in the three panels in the left column of the same figure. It may be noted that galaxies with smaller local dimension are progressively absent when the geometry of environments are characterized on larger length-scales. This points out to the emergence of a homogeneous network of galaxies on larger length-scales as mentioned earlier.","Citation Text":["Casertano & Hut 1985"],"Functions Text":["We address this issue by calculating local number density of galaxies using kth nearest neighbour method"],"Functions Label":["Uses"],"Citation Start End":[[537,557]],"Functions Start End":[[431,535]]}
{"Identifier":"2020ApJ...898L..56L__Abdo_et_al._2010_Instance_1","Paragraph":"The 3 month \u03b3-ray light curve of 4FGL J1510.1+5702 reveals that it is at a high flux state in an epoch of several tens of days in 2018; meanwhile, optical flux densities of GB 1508+5714 in two bands rise in the same epoch. To further investigate the relationship between these two domains of emissions, a 3 day time bin \u03b3-ray light curve is presented, together with the zoomed-in ZTF light curves; see Figure 4. In spite of the limited statistics and no Fermi-LAT observation toward the target at the exact time of the optical flares, the time bin (i.e., centered at MJD 58288.5) with the largest TS value in the 3 day \u03b3-ray light curve is very close to the peaking time of the optical flares. Since optical variations are likely from the jet because of the large variability amplitude, based on the simultaneous \u03b3-ray and optical brightening, we conclude that 4FGL J1510.1+5702 is the \u03b3-ray counterpart of GB 1508+5714. There are three other time bins in the 3 month \u03b3-ray light curve with TS values \u226510, centered at MJD 54909, 57265, and 56540, respectively. The first two time intervals do not fall into the operation time range of iPTF\/ZTF. Moreover, no iPTF\/ZTF data of GB 1508+5714 around MJD 56,540 are available; see Figure 3. Theoretically, in the leptonic radiation scenario, the optical and GeV \u03b3-ray emissions of low synchrotron peaked blazars, including FSRQs, are proposed to be from the same population of emitting electrons. It is supported by correlated optical\/\u03b3-ray flares in FSRQs (e.g., Abdo et al. 2010; Bonning et al. 2012). Meanwhile, the redder-when-brighter spectral variability behavior has been detected in the optical wavelengths of \u03b3-ray FSRQs, which is explained by the influence of the blue and slowly varying accretion disk emission (e.g., Bonning et al. 2012; Fan et al. 2018). For GB 1508+5714, a similar trend is shown. The optical spectral color, rmag\u2013imag, is 0.66 mag in MJD 58286, while it is 0.43 mag in MJD 58272, despite the relatively large photometric uncertainties. The optical spectral color scaled by rmag corresponding to observations in different epochs are plotted, as shown in Figure 4. Its i-band variability amplitude is larger than that in the r band. All of these facts suggest that the contribution of the jet emission becomes significant at the optical wavelengths of GB 1508+5714 when the jet activity is intense with rising \u03b3-ray emission. Moreover, the significant \u03b3-ray emission together with the rapid optical variation in the r band provide information of the emitting jet blob. The radius of the emitting blob is constrained by the variability timescale, \n\n\n\n\n\n, where \n\n\n\n\n\n days for the current event. Meanwhile, assuming that the optical and \u03b3-ray photons of GB 1508+5714 are from the same region, to avoid serious absorption on \u03b3-rays from soft photons via the \n\n\n\n\n\n process, the corresponding optical depth should not be high,\n1\n\n\n\n\n\nwhere \n\n\n\n\n\n is the scattering Thomson cross section, \n\n\n\n\n\n is the differential comoving number density of the target photon per energy, \n\n\n\n\n\n is the energy of the target photon in dimensionless units, and \n\n\n\n\n\n is the absorption length (Dondi & Ghisellini 1995; Begelman et al. 2008). The soft photons from the jet itself could be responsible for the absorption. Since the highest energy of the detected \u03b3-ray photons of GB 1508+5714 is \u223c8 GeV, the corresponding soft photons are those detected at a few keV (3\n\n\n\n\n\n  erg s\u22121; Marcotulli et al. 2020), and the absorption length can be set the same as the radius of the emitting blob. A constraint of the Doppler factor of the jet blob, \n\n\n\n\n\n, is given.","Citation Text":["Abdo et al. 2010"],"Functions Text":["It is supported by correlated optical\/\u03b3-ray flares in FSRQs (e.g.,"],"Functions Label":["Similarities"],"Citation Start End":[[1508,1524]],"Functions Start End":[[1441,1507]]}
{"Identifier":"2022MNRAS.510.3039K__Gunell_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":["Gunell et al. 2018"],"Functions Text":["Thus, for planets in the Solar System, it was shown both in the observations","that the presence of a weak magnetic field can intensify atmospheric escape."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[960,978]],"Functions Start End":[[882,958],[1060,1136]]}
{"Identifier":"2018ApJ...857...98R__Ferraz-Mello_et_al._2008_Instance_1","Paragraph":"Numerous investigations of tidal activity on extrasolar planets have been conducted, with a range of topics from the behavior of gas giants (e.g., B\u011bhounkov\u00e1 et al. 2010, 2011; Remus et al. 2012a, 2012b; Storch & Lai 2014), to tidal alterations of system dynamics (e.g., Lecoanet et al. 2009; Matsumura et al. 2010; C\u00e9bron et al. 2011; Bolmont et al. 2015; Turbet et al. 2017), to tidal alterations of habitability (Barnes et al. 2008, 2013; Jackson et al. 2008a, 2008b; Heller & Armstrong 2014; Kopparapu et al. 2014), issues of spin dynamics (Correia et al. 2008; Ferraz-Mello et al. 2008; Efroimsky 2012b; Cunha et al. 2015), and the role of tides on exomoons (Namouni 2010; Heller & Barnes 2013). Many such studies naturally begin with frequency-independent internal models, but an increasing number consider viscoelastic models (Henning et al. 2009; B\u011bhounkov\u00e1 et al. 2010, 2011; Remus et al. 2012a, 2012b; Auclair-Desrotour et al. 2014; Correia et al. 2014; Henning & Hurford 2014; Makarov & Efroimsky 2014; Shoji & Kurita 2014; Driscoll & Barnes 2015; Makarov 2015). Countless more studies rely upon reasonable selections of tidal dissipation terms in order to inform simulations of system dynamics. For solid planetary objects, a detailed study is eventually needed to constrain which rheological models are best under the stress, pressure, and compositional conditions that are applicable to exoplanets and exomoons. Indeed, studies of the Earth tell us that multiple rheological models may be needed as one goes deeper into an exoplanet\u2019s interior. Higher pressures will surely change the microphysical mechanisms that govern the rheological response (Karato & Spetzler 1990). We currently must rely mainly on analytical and numerical modeling when exploring the interiors of extrasolar planets, particularly worlds in the super-Earth category not represented in our solar system (Valencia et al. 2007). It is not yet known how well laboratory results on the viscosity of peridotite can extend to high-pressure phases such as post-perovskite (Murakami 2004), which may play a large role in super-Earths.","Citation Text":["Ferraz-Mello et al. 2008"],"Functions Text":["Numerous investigations of tidal activity on extrasolar planets have been conducted, with a range of topics","issues of spin dynamics"],"Functions Label":["Background","Background"],"Citation Start End":[[566,590]],"Functions Start End":[[0,107],[520,543]]}
{"Identifier":"2019ApJ...874L..32C__Ogle_et_al._1997_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":["Ogle et al. 1997"],"Functions Text":["The inner few arcseconds is a complex mix of","polarized, broad optical emission lines due to scattering by dust"],"Functions Label":["Background","Background"],"Citation Start End":[[779,795]],"Functions Start End":[[261,305],[689,754]]}
{"Identifier":"2016ApJ...833...76B__Klimchuk_et_al._2008_Instance_1","Paragraph":"A significant limitation of the model is that it ignores the well-established hydrodynamic evolution of the loop during the cooling process, involving the substantial transfer of mass between the chromosphere and the corona. For large downward heat fluxes, the transition region is unable to radiate the supplied energy, resulting in the deposition of thermal energy in the dense chromosphere. The resulting two to three orders-of-magnitude temperature enhancements create a large pressure gradient that drives an upward enthalpy flux of \u201cevaporating\u201d plasma. However, as the loop cools, the decreased heat flux becomes insufficient to sustain the radiation emitted in the now-dense transition region and hence an inverse process of downward enthalpy flux starts to occur. It has been suggested (Klimchuk et al. 2008) that the enthalpy fluxes associated with both evaporating and condensing plasma are at all times in approximate balance with the excess or deficit of the heat flux relative to the transition region radiation loss rate. This basic idea has allowed the development of global \u201cEnthalpy-Based Thermal Evolution of Loops\u201d (EBTEL) models that describe the evolution of the average temperature and density in the coronal part of the loops; these models are generally in good agreement with one-dimensional hydrodynamic simulations (Klimchuk et al. 2008; Cargill et al. 2012a, 2012b). It is, in principle, possible to include the effects of a turbulence-controlled heat flux in EBTEL (or 1D hydrodynamic) models. If this heat flux is reduced sufficiently relative to its collisional value, then, for the reasons explained above, there will be a significant impact on the thermal evolution of the loop. Doing so, however, would still require a numerical treatment, which is beyond the scope of the present work (but which it is our intention to carry out in a future work). Instead, we adopt a simpler approach that allows a systematic and fairly transparent quantitative analysis of the impact of turbulence on the thermodynamics of post-flare loops.","Citation Text":["Klimchuk et al. 2008"],"Functions Text":["It has been suggested","that the enthalpy fluxes associated with both evaporating and condensing plasma are at all times in approximate balance with the excess or deficit of the heat flux relative to the transition region radiation loss rate."],"Functions Label":["Background","Background"],"Citation Start End":[[796,816]],"Functions Start End":[[773,794],[818,1036]]}
{"Identifier":"2018MNRAS.474..838D__and_2016_Instance_1","Paragraph":"The first step in searching for dynamically correlated minor bodies, particularly those resulting from break-ups, is to get a clear characterization of what the expectations may be. The outcome of cometary disruption is well documented through two well-studied examples, those of the comets 73P\/Schwassmann-Wachmann 3 and D\/1993 F2 (Shoemaker-Levy 9). For reasons that still remain unclear, comet 73P started to break apart in 1995 and dozens of fragments were observed in 2006 and 2007 (see e.g. Crovisier et al. 1996; Weaver et al. 2006; Reach et al. 2009; Hadamcik & Levasseur-Regourd 2016). Some of these fragments have been recovered in 2010\u20132011 (Harker et al. 2011, 2017; Sitko et al. 2011) and 2016\u20132017 (e.g. Kadota et al. 2017;1 Williams 2017); it may consist of hundreds of pieces now (68 of them have orbit determinations). This fragmentation process can be described as gentle and progressive. In striking contrast, comet Shoemaker-Levy 9 experienced a sudden, violent fragmentation event triggered by strong tidal forces during a close encounter with Jupiter in 1992 July (see e.g. Sekanina, Chodas & Yeomans 1994, 1998; Asphaug & Benz 1996; Sekanina 1997). Most fragments collided with Jupiter over a period of a week (1994 July 16\u201322); 21 of them have orbit determinations. Quite different may have been the collisional event that led to the formation of the Haumea family (Brown et al. 2007; Schlichting & Sari 2009; Leinhardt, Marcus & Stewart 2010; Lykawka et al. 2012; Ortiz et al. 2012) perhaps more than 1 Gyr ago (Ragozzine & Brown 2007; Volk & Malhotra 2012). Fragments of recently disrupted minor bodies must have very similar values of their semimajor axis, a, eccentricity, e, inclination, i, longitude of the ascending node, $\\mathit {\\Omega }$, argument of perihelion, \u03c9, and time of perihelion passage, \u03c4q, but $\\mathit {\\Omega }$, \u03c9 and \u03c4q tend to become increasingly randomized over time. In contrast, recently unbound pairs resulting from binary dissociation events might have relatively different values of a and e, but very similar values of i, $\\mathit {\\Omega }$ and \u03c9, the difference in \u03c4q may initially range from weeks to centuries, but grows rapidly over time (see e.g. de Le\u00f3n, de la Fuente Marcos & de la Fuente Marcos 2017; de la Fuente Marcos, de la Fuente Marcos & Aarseth 2017).","Citation Text":["Kadota et al. 2017"],"Functions Text":["Some of these fragments have been recovered","and 2016\u20132017"],"Functions Label":["Background","Background"],"Citation Start End":[[718,736]],"Functions Start End":[[595,638],[698,711]]}
{"Identifier":"2021ApJ...912..163B__Brennecka_et_al._2020_Instance_2","Paragraph":"Braukmuller et al. (2018) proposed that all elements fall into one of four categories based on their condensation temperature: refractory elements (50% condensation temperature, Tc,50 > 1400 K), which exhibit approximately uniform enrichments in their Si-normalized concentrations in CC chondrites compared to CI chondrites by a factor of \u223c1\u20131.4; main component elements (1300 K  Tc,50  1400 K), which have approximately the same Si-normalized elemental abundances in CC chondrites as CI chondrites (differ by a factor of \u223c0.8\u20131.1); slope-volatile elements (800 K  Tc,50  1300 K), which exhibit monotonically decreasing Si-normalized concentrations with decreasing Tc,50 compared to CI chondrites; and plateau volatile elements (Tc,50  800 K), which display uniform depletions in Si-normalized concentrations compared to CI chondrites by a factor of \u223c0.1\u20130.7 that are characteristic of each CC chondrite group. Given their uniform nature with Tc,50 and comparatively well-constrained isotopic and elemental compositions, we chose to focus on the concentrations of refractory, main component, and plateau volatile elements in this study. For the refractory and main component elements in CC chondrites, we examine the elemental and isotopic compositions of Ti and Cr, respectively, because these are lithophile elements whose isotopic compositions have been measured precisely for a large number of chondrites and their components (Trinquier et al. 2007, 2009; Qin et al. 2010; Olsen et al. 2016; Van Kooten et al. 2016; Gerber et al. 2017; Davis et al. 2018; Zhu et al. 2019; Schneider et al. 2020; Williams et al. 2020). For CC iron meteorites, we examine the isotopic compositions of Mo and Ni, respectively, because these are siderophile elements (so are therefore present in appreciable concentrations in iron meteorites, unlike Ti and Cr) whose compositions have also been relatively well studied in a number of iron meteorites as well as chondrites and their components (Burkhardt et al. 2011; Budde et al. 2016; Kruijer et al. 2017; Bermingham et al. 2018; Nanne et al. 2019; Budde et al. 2019; Worsham et al. 2019; Brennecka et al. 2020; Spitzer et al. 2020). For the plateau volatile elements, we examine the elemental compositions of six elements (Bi, Ag, Pb, Zn, Te, and Sn) that exhibit a number of desirable properties: their concentrations have been relatively well constrained in CC chondrites; they show a range of lithophile, siderophile, and chalcophile behaviors; their concentrations do not appear to be strongly dependent on redox state; they show minimal variability among NC chondrite groups. Our reasoning for not considering the isotopic compositions of these elements is discussed in Section 2.3. The adopted isotopic and chemical composition of each element used in this study in CC chondrites, CC iron meteorites, CAIs, CI chondrites, and NC chondrites are included in Table 1. Uncertainties on elemental concentrations have not been routinely reported throughout the literature, although these values are typically \u00b15 wt% (e.g., Lodders 2003; Palme et al. 2014). CAIs can be categorized into six groups based on their compositions (Stracke et al. 2012). For the purposes of this study, we adopt the composition of type I CAIs as the representative value of refractory objects because they are seemingly the most abundant type and lack the characteristic elemental depletions of other CAI groups (e.g., Stracke et al. 2012; Brennecka et al. 2020). We also focus largely on ordinary chondrites (OC) as representative NC meteorites rather than enstatite chondrites (EC) or Rumuruti chondrites (RC). This is because EC chondrites formed under more reducing conditions than OC and RC chondrites, which introduced a compositional signature for some elements to EC chondrites that is not present in OC, RC, or CC chondrites (presumably due to their formation in more oxidizing environments) so is not representative of large-scale mixing in the disk (Alexander 2019b). Additionally, the isotopic compositions of RC chondrites are sparsely measured compared to OC and EC chondrites. NC meteorites could have experienced a number of processes (e.g., mixing, chondrule formation, volatile loss, the addition of refractory materials, etc.) that gave these meteorites their specific chemical and isotopic signatures (Alexander 2019b). We do not explore these processes in this study and simply adopt the measured elemental and isotopic compositions of NC chondrites as potential end-members for the compositions of CC meteorites.","Citation Text":["Brennecka et al. 2020"],"Functions Text":["For the purposes of this study, we adopt the composition of type I CAIs as the representative value of refractory objects because they are seemingly the most abundant type and lack the characteristic elemental depletions of other CAI groups (e.g.,"],"Functions Label":["Uses"],"Citation Start End":[[3452,3473]],"Functions Start End":[[3183,3430]]}
{"Identifier":"2019AandA...627A..53H__Villar-Mart\u00edn_et_al._1999_Instance_1","Paragraph":"A spatial coincidence of the radio jet morphology and velocity dispersion of the ionised gas has already been reported for spatially-resolved spectroscopy of more luminous radio-quiet AGN (e.g. Husemann et al. 2013; Villar-Mart\u00edn et al. 2017) and powerful compact radio sources (e.g. Roche et al. 2016), but it has been correctly proposed that the fast moving plasma itself can lead to radio emission that mimics jet activity (Zakamska & Greene 2014; Hwang et al. 2018). In the case of HE 1353\u22121917 we can rule out that the ionised plasma is creating the radio emission because the high-velocity ionised gas traced by [O\u202fIII] is significantly displaced compared to the observed jet-like radio emission. Hence, we think that the radio jet is transferring its energy and momentum to the ambient medium through an extended shock front, which creates turbulence in a dense clumpy ISM. Such a great impact of the radio jet has been observationally shown in many cases (e.g. Villar-Mart\u00edn et al. 1999, 2014; O\u2019Dea et al. 2002; Nesvadba et al. 2006; Holt et al. 2008; Guillard et al. 2012; Harrison et al. 2015; Santoro et al. 2018; Tremblay et al. 2018; Jarvis et al. 2019) and theoretically supported through detailed hydrodynamic simulations (e.g. Krause & Alexander 2007; Sutherland & Bicknell 2007; Wagner & Bicknell 2011; Wagner et al. 2012; Cielo et al. 2018; Mukherjee et al. 2018). As we discussed in Sect. 3.6, the jet power alone is sufficient to energetically drive the outflow because only a small fraction of the AGN luminosity would impact the thin disc of the galaxy implying conversion efficiencies of more than 10% of Lbol. Hopkins & Elvis (2010) proposed a two-stage process for efficient radiation-driven outflows. They describe a scenario in which an initial weak wind in the hot gas phase, possibly initiated by an accretion disc wind or a radio jet, creates additional turbulence in the surrounding medium so that massive gas clouds will subsequently expand and disperse. This expansion of gas clouds would significantly increase their apparent cross-section with respect to incident radiation field of the AGN. Such a two-stage process may increase the coupling efficiency by an order of magnitude. While we cannot directly confirm this process with our observations, the close alignment of the jet axis and the ionisation cone greatly suggest that the outflow is driven jointly by both mechanical and radiative energy with an unknown ratio of the two. The open question is whether the same powerful outflow could have developed without the fast radio jet impacting the cold gas directly given its unique orientation.","Citation Text":["Villar-Mart\u00edn et al. 1999"],"Functions Text":["Hence, we think that the radio jet is transferring its energy and momentum to the ambient medium through an extended shock front, which creates turbulence in a dense clumpy ISM. Such a great impact of the radio jet has been observationally shown in many cases (e.g."],"Functions Label":["Similarities"],"Citation Start End":[[969,994]],"Functions Start End":[[703,968]]}
{"Identifier":"2018ApJ...854...73I__Bowler_et_al._2014_Instance_1","Paragraph":"The best-fit parameters are shown in Table 3. We find that the uncertainties in M* and \n\n\n\n\n\n are considerably large due to a degeneracy between the two parameters when all parameters are variable. Plotted in Figure 1 is the faint-end slope \u03b1 as a function of redshift. Our results indicate that the best-fit values of \u03b1 are about \u22122 at \n\n\n\n\n\n to 10, which are steeper than those at lower redshift (e.g., \n\n\n\n\n\n at \n\n\n\n\n\n in Bouwens et al. 2015). We show the fitting results at \n\n\n\n\n\n, 8, 9, and 10 in Figures 2, 3, 4, and 5, respectively. The top and bottom panels present the observed number densities and the best-fit luminosity functions in the image plane and the source plane, respectively. We also plot the results of previous blank-field surveys (Ouchi et al. 2009; Bradley et al. 2012; McLure et al. 2013; Oesch et al. 2013; Schenker et al. 2013; Bowler et al. 2014; Bouwens et al. 2015; Finkelstein et al. 2015; Calvi et al. 2016) and recent HFF results in other studies (Atek et al. 2015a; Laporte et al. 2016; McLeod et al. 2016). The best-fit parameters are consistent with those in previous studies. In the top panel of Figure 2, there may be an excess in the observed surface number density at \n\n\n\n\n\n. The reason for this excess is not clear, although using a size\u2013luminosity relation that gives smaller sizes at faint magnitudes may reduce this excess. At \n\n\n\n\n\n, the observed number densities at the bright end are slightly larger than the number densities in the simulation. This is probably due to the existence of an overdense region of \n\n\n\n\n\n dropouts in the Abell 2744 cluster field. We discussed the properties of the overdensity in Ishigaki et al. (2016; see also Zheng et al. 2014 and Atek et al. 2015b). At \n\n\n\n\n\n, although we detect no galaxies, we can place a constraint on the luminosity function from the non-detection. Based on the best-fit parameters where only \n\n\n\n\n\n is variable, \u223c1.4 galaxies are expected to be detected in the HFF fields. The middle panels of Figures 2\u20135 show histograms of the number of the dropouts. It is seen that our samples push the magnitude limits of the luminosity functions significantly by up to \u223c3 magnitude. The correlations between M* and \u03b1 at z \u223c6\u20137 and 8 are presented in Figure 6. \n \n","Citation Text":["Bowler et al. 2014"],"Functions Text":["We also plot the results of previous blank-field surveys"],"Functions Label":["Uses"],"Citation Start End":[[856,874]],"Functions Start End":[[697,753]]}
{"Identifier":"2019MNRAS.484..712D__Moriarty_et_al._2014_Instance_1","Paragraph":"The Gibbs free energy of the system, and thus the composition of the solids formed depends on the pressure and temperature at which condensation occurs. In order to consider reasonable pressures and temperatures for the inner regions of the PPD, and in order to convert these temperatures and pressures into radial locations within the disc and formation times for the solid condensates, we consider the simplest possible PPD model. We use the theoretical model derived in Chambers (2009), which models the viscous accretion of gas heated by the star. This model has been previously used for the modelling of planetesimal formation in PPDs (Moriarty et al. 2014; Harrison et al. 2018) and super-Earths (Alessi, Pudritz & Cridland 2017). This model ignores any vertical or radial mixing, and as will be discussed further later, any radial drift. All of these processes may be of critical importance in a realistic PPD. The Chambers model is a disc model with an alpha parameterization that divides the disc into three sections; an inner viscous evaporating region, an intermediate viscous region, and an outer irradiated region. For the calculations in this work, we have assumed disc parameters of $s_{0} = 33\\, \\mathrm{au}$, $\\kappa _{0} = 0.3\\, \\mathrm{m^{2}\\, kg}^{-1}$, \u03b1 = 0.01, \u03b3 = 1.7, \u03bc = 2.4, and $M_{*}=0.78\\, \\mathrm{M}_{\\odot }$ following Chambers (2009) and Motalebi et al. (2015). We also assume that the mass of the PPD is directly proportional to the mass of the host star according to M0 = 0.1M* (Chambers 2009; Andrews et al. 2013). The temperature and radius of the star in the PPD phase are assumed to be functions of the stellar mass in the form derived in Siess, Dufour & Forestini (2000). The relations used in this work to calculate the PPD mass, the initial stellar radius, and the initial stellar temperature as a function of stellar mass are consistent with the values given in Stepinski (1998) and Chambers (2009) for a solar mass star. The analytical expressions for the pressure and temperature of the mid-plane of the PPD as a function of radial location (a) and time (tdisc) are presented in Appendix C. The temperature\u2013radial location curves for the model disc around a star similar to HD 219134 are plotted as a function of time in Fig. 2. The pressure\u2013temperature space mapped out by the model disc for the case of a star similar to HD 219134 is displayed in the Appendix in Fig. C1. The pressure\u2013temperature space for the model disc of a solar mass star shows negligible differences compared to the HD 219134 case.","Citation Text":["Moriarty et al. 2014"],"Functions Text":["This model has been previously used for the modelling of planetesimal formation in PPDs"],"Functions Label":["Background"],"Citation Start End":[[641,661]],"Functions Start End":[[552,639]]}
{"Identifier":"2017MNRAS.469S..39F__Johansen_et_al._2015_Instance_1","Paragraph":"This conclusion can be quantified by the number of catastrophic collisions per comet (Rickman et al. 2015)\n(6)\r\n\\begin{equation}\r\nN_{{\\rm coll}} = N_{\\rm p} A_{\\rm p} u T \/ V,\r\n\\end{equation}\r\nwhere Np is the number of comets in the disc, Ap and u are the collision cross-section and speed, respectively, T \u2248 0.4 Gyr is the time spent from accretion to scattering beyond Neptune and V is the disc volume. A cumulative size distribution of comets with a power index of \u22122.0 and u \u2248 0.6 km s\u22121 provides Ncoll \u2248 3, assuming that \u22485 Earth masses are distributed in comets of diameters from 1 to 100 km (Rickman et al. 2015). The size distribution is probably shallower (Johansen et al. 2015), as confirmed by the shallowing size distribution of craters of diameter 10 km observed on Pluto and Charon (Singer et al. 2016; Robbins et al. 2017). We simulate the impacts on Pluto and Charon by means of Holsapple\u2019s model (http:\/\/keith.aa.washington.edu\/craterdata\/scaling\/index.htm) with the following target parameters: cold ice as geological type, lunar gravity, zero atmospheric pressure; and with the following projectile parameters: bulk density of 500 kg m\u22123, diameter of 0.1 km and impact speed of 1 km s\u22121. Tests performed assuming numerical values around those just listed provide craters always a factor 2 larger than the projectiles, so that the number of projectiles of 0.1 km radius is a factor from (0.1\/2.5)2.4 \u2212 3 = 7 to (0.1\/2.5)1.6 \u2212 3 = 90 below previous estimates (Morbidelli & Rickman 2015; Jutzi & Benz 2017; Jutzi et al. 2017), according to Charon\u2019s and Pluto\u2019s crater size distributions, respectively (Robbins et al. 2017). The same computations performed by Rickman et al. (2015), changing only the cumulative power index from \u22122.0 to \u22121.4, i.e. Charon\u2019s crater size distribution (Robbins et al. 2017), provide Ncoll \u2248 0.5, and Ncoll \u2248 0.05 with a cumulative power index of \u22120.6, i.e. Pluto\u2019s crater size distribution (Robbins et al. 2017), which make equation (6) consistent with the observed flux of 67P fractal particles. This conclusion is confirmed by the craters observed on Nix (Robbins et al. 2017), with a cumulative power index of \u22120.8, close to Pluto\u2019s one, a measured cumulative crater density of 0.01 km\u22122 at diameters >3 km (smaller projectiles would not destroy a parent comet, according to equation (5)) and an extrapolated cumulative crater density 0.1 km\u22122 at diameters >0.1 km, matching that extrapolated for Charon with a steeper cumulative index (Robbins et al. 2017). Nix and Charon are less affected than Pluto by a possible significant resurfacing and their crater density suggests a negligible number of catastrophic impacts during the Solar system lifetime on bodies of cross-section 10 km2.","Citation Text":["Johansen et al. 2015"],"Functions Text":["The size distribution is probably shallower"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[666,686]],"Functions Start End":[[621,664]]}
{"Identifier":"2021MNRAS.503.1734I__McKinney_et_al._2019_Instance_1","Paragraph":"We show in Fig. 6 the dependence of some Ly\u2009\u03b1 characteristics on absolute FUV magnitude. Our galaxies with detected LyC emission are shown by red filled circles and those with upper limits of LyC emission are represented by red open circles. All our low-mass galaxies are characterized by moderate Ly\u2009\u03b1 luminosities in a narrow range between 1042.19 and 1042.74\u2009erg\u2009s\u22121 (Fig. 6a), that are slightly below the values for confirmed LyC leakers (blue filled circles) by Izotov et al. (2016a,b, 2018a,b) and high-redshift galaxies (grey open circles) by Ouchi et al. (2008), Hashimoto et al. (2017), Jiang et al. (2018), Matthee et al. (2017), Matthee et al. (2018), Sobral et al. (2018), but are \u223c1 order of magnitude higher than those for GPs (black asterisks; Jaskot & Oey 2014; Henry et al. 2015; Jaskot et al. 2017; Yang et al. 2017a; McKinney et al. 2019). Our low-mass galaxies have high EW(Ly\u2009\u03b1) \u223c 65\u2013220\u2009\u00c5 (Table 7, red symbols in Fig. 6b), similar to those in other LyC leakers (blue symbols in Fig. 6b). They are at the high end of the EW(Ly\u2009\u03b1) values for high-z LAEs by Ouchi et al. (2008), Hashimoto et al. (2017), Jiang et al. (2018), Harikane et al. (2018), Caruana et al. (2018), Pentericci et al. (2018), Matthee et al. (2017, 2018), Sobral et al. (2018) (grey open circles) and GPs (black asterisks). However, contrary to expectations for galaxies with lower stellar masses and, likely, lower masses of the neutral gas, the separation between the Ly\u2009\u03b1 peaks is on average similar to that in higher mass LyC leakers (Fig. 6c), and it is higher in galaxies with upper limits of LyC emission (>400\u2009km\u2009s\u22121, red open circles). Ly\u2009\u03b1 escape fractions fesc(Ly\u2009\u03b1) in low-mass galaxies are also similar to those in higher mass LyC leakers, and they are lower in galaxies with upper limits of LyC emission (red open circles in Fig. 6d). We also note that the average ratio of blue and red peak fluxes of \u223c25 per cent (Table 7) is somewhat lower than the value of \u223c30 per cent quoted by Hayes et al. (2021) for $z\\, \\sim$ 0 galaxies.","Citation Text":["McKinney et al. 2019"],"Functions Text":["All our low-mass galaxies are characterized by moderate Ly\u2009\u03b1 luminosities in a narrow range between 1042.19 and 1042.74\u2009erg\u2009s\u22121 (Fig. 6a), that are slightly below the values for confirmed LyC leakers (blue filled circles) by Izotov et al. (2016a,b, 2018a,b) and high-redshift galaxies (grey open circles) by Ouchi et al. (2008), Hashimoto et al. (2017), Jiang et al. (2018), Matthee et al. (2017), Matthee et al. (2018), Sobral et al. (2018), but are \u223c1 order of magnitude higher than those for GPs (black asterisks;"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[836,856]],"Functions Start End":[[242,758]]}
{"Identifier":"2018MNRAS.476.2421G__Gallego_et_al._2018_Instance_1","Paragraph":"An alternative approach is to map the CGM through direct imaging of the Ly\u2009\u03b1 line. Theoretical models suggest that three main mechanisms should be able to generate circumgalactic Ly\u2009\u03b1 emission: cooling radiation of gravitationally heated gas (e.g. Haiman, Spaans & Quataert 2000; Yang et al. 2006; Dijkstra & Loeb 2009), ultraviolet (UV) photons produced through shock mechanisms (Taniguchi & Shioya 2000; Mori, Umemura & Ferrara 2004), and recombination radiation following photoionization (often referred as fluorescence) powered by UV sources (Cantalupo et al. 2005; Geach et al. 2009; Kollmeier et al. 2010). While the fluorescent signal powered by the diffuse metagalactic UV background (Hogan & Weymann 1987; Binette et al. 1993; Gould & Weinberg 1996; Haardt & Madau 1996), with an expected surface brightness (SB) of $\\rm SB_{Ly\\alpha }\\sim 10^{-20} \\,erg \\,s^{-1} \\,cm^{-2} \\,arcsec^{-2}$ (Cantalupo et al. 2005; Rauch et al. 2008), is still out of reach for current optical instrumentation (but see Gallego et al. 2018), Ly\u2009\u03b1 fluorescence is predicted to be boosted up into the detectable regime in the vicinity of bright ionizing sources, as luminous quasars (Rees 1988; Haiman & Rees 2001; Alam & Miralda-Escud\u00e9 2002; Cantalupo et al. 2005). This theoretical prediction has been confirmed by a number of surveys targeting the fluorescent Ly\u2009\u03b1 emission around luminous and radio-quiet quasars, using narrow-band (NB) filters on 8-metre class optical telescopes (e.g. Cantalupo, Lilly & Haehnelt 2012; Cantalupo et al. 2014; Martin et al. 2014; Arrigoni Battaia et al. 2016) and spectroscopic observations (e.g. Christensen et al. 2006; North et al. 2012; Herenz et al. 2015). However, these surveys have revealed giant Ly\u2009\u03b1 nebulae, spanning distances from the quasar larger than 100 physical kpc (pkpc), only in less than 10\u2009per\u2009cent of the targets (e.g. Cantalupo et al. 2014; Hennawi et al. 2015; maximum projected linear sizes >300 pkpc), and emission on smaller scales (R \u2272 50\u201360 pkpc) have been detected only in about 50\u2009per\u2009cent of the cases. This 50\u2009per\u2009cent detection rate is likely due to a combination of limits of the observational techniques, as for instance NB filter losses, spectroscopic slit losses, point spread function (PSF) losses and, most importantly, dilution of the signal of the Ly\u2009\u03b1 line into the continuum flux (both background and from the host galaxy) encompassed by the width of the filter (see also Borisova et al. 2016 for a discussion).","Citation Text":["Gallego et al. 2018"],"Functions Text":["While the fluorescent signal powered by the diffuse metagalactic UV background","with an expected surface brightness (SB) of $\\rm SB_{Ly\\alpha }\\sim 10^{-20} \\,erg \\,s^{-1} \\,cm^{-2} \\,arcsec^{-2}$","is still out of reach for current optical instrumentation (but see","Ly\u2009\u03b1 fluorescence is predicted to be boosted up into the detectable regime in the vicinity of bright ionizing sources, as luminous quasars"],"Functions Label":["Background","Background","Background","Background"],"Citation Start End":[[1009,1028]],"Functions Start End":[[613,691],[781,897],[942,1008],[1031,1169]]}
{"Identifier":"2022ApJ...927..149L__Cicone_et_al._2017_Instance_1","Paragraph":"These trends make physical sense and agree with the limited previous measurements. Physically, elevated SFR\/M\n\u22c6 may trace more intense radiation fields and stronger heating of the gas, suggesting higher temperatures. The anticorrelation with M\n\u22c6 may reflect the impact of dust shielding. Based on the existence of the mass\u2013metallicity relation (e.g., Tremonti et al. 2004; Kewley & Ellison 2008), we expect the low-mass members of our sample to also have lower dust-to-gas ratios (e.g., Leroy et al. 2011; R\u00e9my-Ruyer et al. 2014; Casasola et al. 2020) and more intense radiation fields. In literature studies, CO line ratios do appear enhanced in low-metallicity regions or galaxies (e.g., Lequeux et al. 1994; Bolatto et al. 2003; Druard et al. 2014; Kepley et al. 2016; Cicone et al. 2017, among many others). Higher SFR\/L\nCO may indicate poorly shielded, low-metallicity gas in which the CO persists only in the core of a molecular cloud (e.g., see discussion in Glover & Clark 2012; Schruba et al. 2012; Bolatto et al. 2013b; Rubio et al. 2015). Alternatively, higher SFR\/L\nCO can indicate more efficiently star-forming gas, which will often be denser gas with more nearby heating sources. These are both factors that can lead to higher line ratios, especially R\n32 and R\n31 (see Section 2). Given that our sample skews toward relatively massive and thus nearly solar metallicity targets, we expect that these density and heating effects likely represent the main drivers of the observed correlations. The correlations that we see agree with the results of Lamperti et al. (2020), who showed a correlation between SFR\/L\nCO,low and R\n31, and Yajima et al. (2021), who used a subset of the data we consider here and showed a correlation between R\n21 and SFR\/L\nCO,low. Qualitatively, Figure 5 echoes other results at low and high redshift that show a correlation between normalized star formation activity and excitation (e.g., Wei\u00df et al. 2005; Bolatto et al. 2013b; Liu et al. 2021).","Citation Text":["Cicone et al. 2017"],"Functions Text":["In literature studies, CO line ratios do appear enhanced in low-metallicity regions or galaxies (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[772,790]],"Functions Start End":[[587,689]]}
{"Identifier":"2022MNRAS.512.4280P__Umetsu_et_al._2016_Instance_1","Paragraph":"The fifth force, propagated by the scalar degree of freedom, affects the Poisson equations associated to the Newtonian potential \u03a6, as well as the relativistic one, \u03a8, according to (Kobayashi, Watanabe & Yamauchi 2015; Crisostomi & Koyama 2018; Dima & Vernizzi 2018),\n(1)$$\\begin{eqnarray*}\r\n\\frac{\\text{d} \\Phi (r)}{\\text{d}r} = \\frac{G M(r)}{r^2} \\left[1+\\frac{3}{4}Y_1\\left(\\frac{\\rho (r)}{\\bar{\\rho }(r)}\\right)\\left(2+\\frac{\\text{d}\\ln \\rho }{\\text{d}\\ln r}\\right)\\right],\r\n\\end{eqnarray*}\r\n$$(2)$$\\begin{eqnarray*}\r\n\\frac{\\text{d} \\Psi (r)}{\\text{d}r} =\\frac{G M(r)}{r^2}\\left[1-\\frac{15}{4}Y_2\\left(\\frac{\\rho (r)}{\\bar{\\rho }(r)}\\right)\\right].\r\n\\end{eqnarray*}\r\n$$In the above equations, we have assumed spherical symmetry. $\\bar{\\rho }(r)$ is the (spatially) average density at radius r from the centre of the galaxy cluster, and Y1, Y2 correspond to the dimensionless fifth-force couplings. Finally, G is the Newton\u2019s constant. Although the dynamics of member galaxies in the cluster is governed by the potential \u03a6, lensing is sourced by the combination\n(3)$$\\begin{eqnarray*}\r\n\\frac{\\mathrm{ d}}{\\mathrm{ d}r} \\Phi _{\\rm {lens}} = \\frac{1}{2}\\frac{\\mathrm{ d}}{\\mathrm{ d}r}(\\Phi + \\Psi).\r\n\\end{eqnarray*}\r\n$$Therefore, kinematical observations allow for contraints on Y1, while lensing constrains both Y1 and Y2. The right-hand side of above equation can be expressed in terms of the density profile \u03c1(r) according to the relevant equations for \u03a6 and \u03a8 above. The dominant source of pressureless matter density in the cluster comes from dark matter, which density we choose to model with a Navarro-Frenk-White (NFW) of Navarro, Frenk & White (1997) profile as\n(4)$$\\begin{eqnarray*}\r\n\\rho (r)=\\frac{\\rho _\\text{s}}{r\/r_\\text{s}(1+r\/r_\\text{s})^2},\r\n\\end{eqnarray*}\r\n$$with \u03c1s is a characteristic density and rs the radius at which the logarithmic derivative of the density profile takes the value \u22122. The NFW profile has been shown to provide an overall good agreement with observations and simulations over a broad range of scales in GR (e.g. Biviano et al. 2013; Umetsu et al. 2016; Peirani et al. 2017) and in MG (e.g. Lombriser et al. 2012a; Wilcox et al. 2016). Moreover, the GR analyses with lensing and internal kinematics of both clusters indicate that the total mass profile is well fitted by the NFW model (Biviano et al. 2013; Umetsu et al. 2016; Caminha et al. 2017; Sartoris et al. 2020). Under the assumption of a NFW profile, we can re-write the equation for the potential \u03a6 in an effective way as\n(5)$$\\begin{eqnarray*}\r\n\\frac{\\text{d}\\Phi }{\\text{d}r} \\equiv \\frac{G M_{\\text{dyn}}}{r^2}=\\frac{G}{r^2}\\left[ M_{\\rm {NFW}}(r)+M_1(r)\\right],\r\n\\end{eqnarray*}\r\n$$which serves as a definition of the dynamical mass Mdyn. Notice that, G here is still Newton\u2019s constant as measure in the Solar system. The fifth-force contribution M1 is defined in terms of the NFW parameters as\n(6)$$\\begin{eqnarray*}\r\nM_1(r)= M_{200}\\frac{Y_1}{4}\\frac{r^2(r_\\text{s}-r)}{(r_\\text{s}+r)^3}\\times [\\ln (1+c)- c\/(1+c)]^{-1}.\r\n\\end{eqnarray*}\r\n$$where c = r\/rs is the concentration and M200 is the mass of a sphere of radius r200 enclosing an average density 200 times the critical density of the universe at that redshift. In a similar fashion, the relevant expression for the lensing mass can be found by computing\n(7)$$\\begin{eqnarray*}\r\nM_{\\text{lens}}(r) =\\frac{r^2}{2G}\\left[\\frac{\\text{d}\\Psi }{\\text{d}r}+\\frac{\\text{d}\\Phi }{\\text{d}r}\\right].\r\n\\end{eqnarray*}\r\n$$$$\\begin{eqnarray*}\r\nM_{\\text{lens}}=M_{\\text{NFW}}+\\frac{r^2M_{200}\\left[Y_1(r_\\text{s}-r)-5Y_2(r_\\text{s}+r)\\right]}{4[\\log (1+c_{200})-c_{200}\/(1+c_{200})]}\\frac{1}{(r_\\text{s}+r)^{3}},\r\n\\end{eqnarray*}\r\n$$which can be effectively re-expressed in terms of the dynamical mass as\n(8)$$\\begin{eqnarray*}\r\nM_{\\text{lens}} \\equiv M_{\\text{dyn}}+M_2,\r\n\\end{eqnarray*}\r\n$$with M2 the contribution from the fifth force defined through\n(9)$$\\begin{eqnarray*}\r\nM_2=\\frac{r^2M_{200}}{8(r_\\text{s}+r)^{3}}\\frac{Y_1(r-r_\\text{s})-5Y_2(r_\\text{s}+r)}{[\\ln (1+c)-c\/(1+c)]}.\r\n\\end{eqnarray*}\r\n$$In view of the above equations, it is important to emphasize again that, although the fifth force effect enters the dynamical mass only through the coupling Y1, the lensing mass is affected by both Y1 and Y2. This is expected, since lensing is sourced by the combination of the two potentials \u03a6 and \u03a8, equation (3). Note also that, with gravitational lensing observations, one reconstructs the projected surface mass density profile \u03a3(R), where R is the projected radius from the cluster centre. We refer to e.g. Umetsu (2020) for an explicit discussion of the physics and mathematical framework.","Citation Text":["Umetsu et al. 2016"],"Functions Text":["The NFW profile has been shown to provide an overall good agreement with observations and simulations over a broad range of scales in GR"],"Functions Label":["Similarities"],"Citation Start End":[[2078,2096]],"Functions Start End":[[1914,2050]]}
{"Identifier":"2018AandA...615A..77L__XIII_2016_Instance_1","Paragraph":"The last decade and a half has seen a revolution in the study of overdensities in the early Universe. While the study and careful characterization of large associations of galaxies in the local Universe has been possible for nearly a century, and in the intermediate redshift Universe for a significant fraction of that time (e.g., Shapley & Ames 1926; Shapley 1930; Zwicky 1937; Abell 1958; Zwicky et al. 1961), the study of their progenitors presented several practical problems which have prevented their study until relatively recently. The primary problem, inherent to the study of nearly all galaxy populations in the early Universe, is the extreme apparent faintness of galaxy populations at these distances. While some phenomena exist in the early Universe, such as quasars or radio galaxies, which are so powerful and intrinsically bright that they have been able to serve as beacons to early searches near the epoch of H I reionization (z ~ 5.5\u201310, Becker et al. 2001, 2015; Planck Collaboration XIII 2016), the bulk of the galaxy population residing in the early Universe does not contain such phenomena (Miley & De Breuck 2008; Ouchi et al. 2008; Lemaux et al. 2014b; Ueda et al. 2014; Talia et al. 2017). As such, the first searches for these more typical primeval galaxies were largely doomed to failure (Davis & Wilkinson 1974; Partridge 1974; Pritchet & Hartwick 1987; Parkes et al. 1994). It was not until the advent of the 10m-class ground-based telescopes largely used in conjunction with the Hubble Space Telescope (HST) that the prospect of detecting and characterizing moderate samples of such galaxies became even remotely feasible (e.g., Steidel et al. 1999; Shapley et al. 2003; Giavalisco et al. 2004; Stanway et al. 2004; Malhotra & Rhoads 2004; Malhotra et al. 2005; Vanzella et al. 2005, Le F\u00e8vre et al. 2005). With this, the prospect of finding and characterizing analogs of the progenitors of the massive clusters and superclusters of galaxies scattered throughout the local Universe began to come within the realm of possibility.","Citation Text":["Planck Collaboration XIII 2016"],"Functions Text":["While some phenomena exist in the early Universe, such as quasars or radio galaxies, which are so powerful and intrinsically bright that they have been able to serve as beacons to early searches near the epoch of H I reionization (z ~ 5.5\u201310,","the bulk of the galaxy population residing in the early Universe does not contain such phenomena"],"Functions Label":["Background","Background"],"Citation Start End":[[985,1015]],"Functions Start End":[[716,958],[1018,1114]]}
{"Identifier":"2019MNRAS.489..855C__Husemann_et_al._2013_Instance_2","Paragraph":"The size of ENLRs have been defined in different ways in the literature. Bennert et al. (2002) and Schmitt et al. (2003b) used the Hubble Space Telescope (HST) to obtain narrow band images of $\\rm [O\\, III]$, and adopted the maximum 3\u03c3 detected radius as the radius of the ENLR. This method is subject to the instrumental sensitivity limit that could be very different in different observations. Studies with long-slit spectroscopic observations define the radius based on isophote (Greene et al. 2011; Hainline et al. 2013, 2014), or the distance at which the ionization state changes from AGN to star-forming activities (Bennert et al. 2006a,b). The long-slit based observations also have drawbacks: the morphology of ENLR is sometimes irregular so that the derived size depends on the orientation of slits (Greene et al. 2011; Husemann et al. 2013). We have compared the measured size based on the IFU and the mock long-slit observation in Fig. 4 following the method discussed below. In most cases, long-slit observations tend to underestimate the true size of ENLR. IFU spectroscopic data allow us to use 2D maps to define the sizes of ENLRs. Common definitions include the radius of a specified $\\rm [O\\, III]$ surface brightness isophote (Liu et al. 2013, 2014), or the $\\rm [O\\, III]$ flux weighted radius (Husemann et al. 2013, 2014; Bae et al. 2017). We followed the same method as Liu et al. (2013) but chose a different threshold. The isophote threshold of 10\u221215$\\rm erg\\, s^{-1}cm^{-2}arcsec^{-2}$ was used for quasars related studies. This is suitable for such bright objects but are not as useful for fainter Syferts in our sample as it will leave a large number of AGN undetected. The typical 3\u03c3 depth of the MaNGA observation in $\\rm [O\\, III]$ surface brightness can reach 10\u221217$\\rm erg\\, s^{-1}cm^{-2}arcsec^{-2}$. For our AGN sample, the majority of AGN spaxels have surface brightnesses above 10\u221216$\\rm erg\\, s^{-1}cm^{-2}arcsec^{-2}$ which is thus adopted in this work as the threshold to define the sizes of the ENLRs (hereafter R16). If all spaxels are above this threshold, we extrapolated the fitted $\\rm [O\\, III]$ surface brightness profile to determine R16 (see Section 3.4 for more detail). It should be noted that the surface brightness can be affected by cosmological dimming, which has a scale factor of (1 + z)4 (Liu et al. 2013; Hainline et al. 2014). That is important for works trying to compare sample with different redshift, especially for high redshift quasars.","Citation Text":["Husemann et al. 2013"],"Functions Text":["Common definitions include","or the $\\rm [O\\, III]$ flux weighted radius"],"Functions Label":["Uses","Uses"],"Citation Start End":[[1315,1335]],"Functions Start End":[[1148,1174],[1270,1313]]}
{"Identifier":"2019ApJ...871...58T__Charbonnel_&_Lagarde_2010_Instance_1","Paragraph":"We derived stellar parameters, [C\/M], and [N\/M] using SLAM. To avoid bad fits at the edges of the parameter space, we exclude stars with spectral S\/N in the g band less than 50, and metallicity less than \u22121.4. The derived C and N abundances are shown in Figure 8. Clearly, in the top panel, the CH-strong, CH-normal, and metal-poor field stars are separated, and their relative distribution in the N\u2013C parameter space is similar to the case of APOGEE abundances (left panel of Figure 7): (1) metal-poor field stars form a sequence in the lower left of the top panel. As evolved stars ascend the RGB, C and N abundances may be changed by first dredge-up (Iben 1964, 1967) and extra mixing (Gratton et al. 2000; Charbonnel & Lagarde 2010). Given that for a typical halo\/thick-disk star of 1 M\u2299, the first dredge-up occurs around Teff = 5200 K (Boothroyd & Sackmann 1999), and most of our sample stars have \n\n\n\n\n\n K and log g  2.5, we infer that most stars have already undergone first dredge-up. On the other hand, the C and N abundances of these stars could be altered by extra mixing. Stars with brighter K-band absolute magnitudes tend to have higher [N\/Fe] and lower [C\/Fe] (middle and bottom panels of Figure 8), which is consistent with extra-mixing theory and observation (Gratton et al. 2000; Charbonnel & Lagarde 2010); (2) CH-normal stars show an enhanced median N abundance and slightly depleted median C abundance. Clearly, the median N abundance of CH-normal stars is enhanced compared to that of normal metal-poor field stars with similar C abundances. In other words, the enhanced N abundances in CH-normal stars cannot be explained by the extra-mixing effect alone. We notice that a few CH-normal stars may have low N abundances, probably due to large uncertainties when spectra of a particular type are scarce in the training set, i.e., high-N metal-poor stars. The statistical similarity between APOGEE C and N abundances and LAMOST-derived C and N abundances further strengthens our statement above.","Citation Text":["Charbonnel & Lagarde 2010"],"Functions Text":["As evolved stars ascend the RGB, C and N abundances may be changed by first dredge-up","and extra mixing"],"Functions Label":["Background","Background"],"Citation Start End":[[710,735]],"Functions Start End":[[567,652],[671,687]]}
{"Identifier":"2021AandA...655A..99D__Hern\u00e1ndez_et_al._2010_Instance_1","Paragraph":"The giant planet metallicity correlation supports the core-accretion scenario for the formation of planets (Pollack et al. 1996; Ida & Lin 2004; Mordasini et al. 2009), in which it is assumed that planetesimals are formed by the condensation of heavy elements. The discovery of this correlation has led to an increased interest on the abundances of other elements in planet hosts(e.g. Sadakane et al. 2002; Bodaghee et al. 2003; Beir\u00e3o et al. 2005; Gilli et al. 2006; Ecuvillon et al. 2006; Bond et al. 2006, 2008; Robinson et al. 2006; Gonzalez & Laws 2007; Takeda et al. 2007; Delgado Mena et al. 2010; Gonz\u00e1lez Hern\u00e1ndez et al. 2010, 2013; Kang et al. 2011; Brugamyer et al. 2011; Adibekyan et al. 2015; da Silva et al. 2015; Mishenina et al. 2016; Su\u00e1rez-Andr\u00e9s et al. 2018). Using the same sample of this study, Adibekyan et al. (2012b) found that the [X\/Fe] ratios of Mg, Al, Si, Sc, and Ti both for giant and low-mass planet hosts are systematically higher than those of stars without detected planets at low metallicities ([Fe\/H] \u2272 from \u20130.2 to 0.1 dex depending on the element). Furthermore, this work confirmed the previous suggestion by Haywood (2009) that planets form preferentially in the thick disk rather than in the thin disk at lower metallicities. A plausible explanation for this behaviour is that if the amount of iron is low, it needs to be compensated with other elements that are important for planet formation, such as Mg and Si, and these elements are more abundant in the thick disk Adibekyan et al. (2012b,a). In a subsequent work, Delgado Mena et al. (2018) focused on heavier elements among the same sample of stars with and without planets, we found that planet hosts present higher abundances of Zn for [Fe\/H]  \u20130.1 dex, as a consequence of thick-disk stars having enhanced [Zn\/Fe] ratios (Delgado Mena et al. 2017). Moreover, Delgado Mena et al. (2018) also found a statistically significant underabundance of Ba for low-mass planet hosts that had been previously suggested by Mishenina et al. (2016).","Citation Text":["Gonz\u00e1lez Hern\u00e1ndez et al. 2010"],"Functions Text":["The discovery of this correlation has led to an increased interest on the abundances of other elements in planet hosts(e.g."],"Functions Label":["Motivation"],"Citation Start End":[[605,635]],"Functions Start End":[[261,384]]}
{"Identifier":"2022MNRAS.509.1959S__Nordlander_et_al._2019_Instance_1","Paragraph":"However, the transition between the two extremes of modern (metal-rich) and primordial (metal-poor) star formation, and in particular the role of dust coupling and stellar radiation feedback at low metallicity, has thus far received limited exploration. Krumholz (2011) present analytical models for radiation feedback and predict a weak scaling of IMF peak mass with metallicity, while Myers et al. (2011) and Bate (2014, 2019) carry out radiation-hydrodynamic simulations of star formation over a metallicity range from $0.01{-}3\\, Z_{\\rm {\\odot }}$ and find negligible effects on gas fragmentation. However, these studies do not explore lower metallicities, despite available evidence for the existence of a low-metallicity ISM in the past through the discovery of stars with metallicities as low as $10^{-4}\\, \\rm {Z_{\\odot }}$ (Caffau et al. 2011; Starkenburg et al. 2018), as well as several others with $\\rm {[Fe\/H]} \\lt -5$ (Christlieb et al. 2004; Keller et al. 2014; Frebel et al. 2015; Aguado et al. 2017, 2018; Ezzeddine et al. 2019; Nordlander et al. 2019). Coming from the opposite direction, Bromm et al. (2001), Omukai et al. (2005), and Omukai, Hosokawa & Yoshida (2010) consider the thermodynamics of gas of increasing metallicity, and find that dust and metal line cooling permits fragmentation to reach masses \u22721 M\u2299 only once the metallicity exceeds \u223c10\u22123.5 Z\u2299. Dust is a more efficient coolant than metal lines, and permits fragmentation to lower masses at lower metallicity (e.g. Meece, Smith & O\u2019Shea 2014; Chiaki & Yoshida 2020; Shima & Hosokawa 2021), but exactly by how much depends on the poorly known distribution of dust grain sizes in the early Universe (Schneider et al. 2006, 2012; Omukai et al. 2010; Schneider & Omukai 2010; Chiaki et al. 2015). However, the early Universe star formation models are fundamentally misanalogous to the modern ones that consider decreasing metallicity, in that the early Universe models consider dust solely as a coolant that enables fragmentation, whereas the modern ones assign it a more nuanced role, as both a source of cooling and later, once stellar feedback begins, a source of heating \u2013 a changeover that seems crucial to explaining why the IMF in the present-day Universe peaks at ${\\sim}0.2\\, \\rm {M_{\\odot }}$ rather than ${\\sim}10^{-2}\\, \\rm {M_{\\odot }}$ (Kroupa 2001; Chabrier 2003, 2005).","Citation Text":["Nordlander et al. 2019"],"Functions Text":["However, these studies do not explore lower metallicities, despite available evidence for the existence of a low-metallicity ISM in the past through the discovery of stars","as well as several others with $\\rm {[Fe\/H]} \\lt -5$"],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[1046,1068]],"Functions Start End":[[602,773],[879,931]]}
{"Identifier":"2022ApJ...939..117Z__Blandford_et_al._2019_Instance_1","Paragraph":"Blazars are a subclass of active galactic nuclei (AGNs) with relativistic jets of high-energy particles pointing near our line of sight (e.g., Urry & Padovani 1995). Their nonthermal emission is generally detected across the entire electromagnetic spectrum from radio to \u03b3-ray bands. Blazars are subclassified into flat-spectrum radio quasars (FSRQs) and BL Lac objects (BL Lacs), according to the equivalent width of the emission lines in their optical spectrum (Stickel et al. 1991; Stocke et al. 1991; Marcha et al. 1996). These two subclasses of blazars are thought to be intrinsically different, perhaps based on their accretion mode (Dermer & Giebels 2016). FSRQs have high luminosity and a thin and radiatively efficient black hole accretion disk (Malkan & Moore 1986), while BL Lacs are powered by an advection-dominated, low radiative efficiency accretion flow (Dermer & Giebels 2016; Blandford et al. 2019). The jet emission is relativistically beamed (Ghisellini 2019), with a Doppler boosting factor corresponding to a bulk Lorentz factor of several to greater than 10 (Pushkarev et al. 2009). In both cases, the broadband spectra consist of two broad humps, one peaking in the IR-to-X-ray regime and the other peaking in the \u03b3-ray regime. The low-energy peak is believed to be due to synchrotron emission, while the high-energy peak is likely due to inverse Compton scattering of low-energy photons of either the same synchrotron photons (for BL Lacs) or external photons from the disk\/BLR (for FSRQs) (e.g., Dutka et al. 2017). However, some blazars might not necessarily be detected in \u03b3-rays (e.g., Paliya et al. 2017). Indeed, a recent study showed that blazars undetected in \u03b3-rays are likely to have relatively smaller Doppler factors and more disk dominance (Paliya et al. 2017). In the case of strong Compton scattering, the beaming of \u03b3-rays could be larger than, e.g., that seen in the radio (Dermer 1995), leading to the possible nondetection (or reduced detection efficiency) of \u03b3-rays from sources not seen exactly pole-on.","Citation Text":["Blandford et al. 2019"],"Functions Text":["while BL Lacs are powered by an advection-dominated, low radiative efficiency accretion flow"],"Functions Label":["Background"],"Citation Start End":[[894,915]],"Functions Start End":[[777,869]]}
{"Identifier":"2022AandA...659L...1L__Lellouch_et_al._2013_Instance_2","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"],"Functions Text":["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"],"Functions Label":["Uses"],"Citation Start End":[[1003,1023]],"Functions Start End":[[815,1001]]}
{"Identifier":"2020MNRAS.492.4727C__Rees_1976_Instance_1","Paragraph":"A strong property of the isothermal gas that defines the interstellar medium (ISM) is the absence of a characteristic mass scale owing to the scale-free nature of gravity. From the collapse condition of an isothermal sphere, one can derive the Jeans mass as\n(2)$$\\begin{eqnarray*}\r\nM_{\\mathrm{J}} = \\frac{4 \\pi}{3} \\rho \\left(c_{\\mathrm{s}} t_{\\rm ff}\\right)^3\r\n\\end{eqnarray*}$$with cs being the sound speed, given by\n(3)$$\\begin{eqnarray*}\r\nc_{\\mathrm{s}} = \\sqrt{\\frac{k_{\\mathrm{B}} T}{\\mu \\, m_{\\mathrm{H}}}}\r\n\\end{eqnarray*}$$and the free-fall time tff is defined here as\n(4)$$\\begin{eqnarray*}\r\nt_{\\mathrm{ff}} = \\sqrt{\\frac{3\\pi }{32G \\rho }}.\r\n\\end{eqnarray*}$$One sees immediately that for an isothermally collapsing gas, we cannot choose a unique characteristic density that would result in a unique characteristic mass. The larger the density of the gas during collapse and fragmentation, the smaller the Jeans mass. At very large densities, though, owing to the increased dust absorption coefficient, the gas becomes opaque to its own radiation. This is sometimes referred to as the opacity limit (Low & Lynden-Bell 1976; Rees 1976). At this point, the fragmentation is halted, setting a scale below which no fragments form. It is possible to estimate the characteristic density at which this process occurs by requiring the optical depth of one Jeans radius to be unity\n(5)$$\\begin{eqnarray*}\r\n\\tau = \\kappa _{\\rm dust} \\rho R_{\\rm J} = 1\\,\\,\\,{\\rm with}\\,\\,\\,R_{\\rm J} = c_{\\mathrm{s}} t_{\\rm ff}.\r\n\\end{eqnarray*}$$For typical molecular clouds conditions in the Milky Way with T = 10 K, \u03bc = 2.2, and \u03badust = 0.1 cm2 g\u22121, we find the critical density that defines the opacity limit\n(6)$$\\begin{eqnarray*}\r\n\\rho _{\\rm crit} = \\frac{32 G}{3 \\pi \\kappa _{\\rm dust}^2 c_\\mathrm{ s}^2 } \\simeq 5 \\times 10^{-14} \\,{\\rm g}\\,{\\rm cm}^{-3} .\r\n\\end{eqnarray*}$$At this critical density, we obtain a unique Jeans mass of the order of\n(7)$$\\begin{eqnarray*}\r\nM_{\\mathrm{J}} = 0.670 \\frac{c_{\\mathrm{s}}^3}{\\sqrt{G^3 \\rho _{\\mathrm{crit}}}} \\approx 6 \\times 10^{-4} \\, \\mathrm{M_\\odot }.\r\n\\end{eqnarray*}$$Note that the critical density can also be derived using slightly more complicated arguments, leading to a very similar value for typical Milky Way conditions (Krumholz 2017). This mass corresponds to the smallest gas clumps that can overcome the pressure gradients and collapse. At this characteristic density, the gas evolution transitions from isothermal to adiabatic, which prevents the collapse of smaller fragments, even at higher densities. The Jeans mass at the opacity limit is obviously much smaller than the observed characteristic mass of the IMF. It does not even correspond to the minimal mass of a star, as the smallest of these fragments will not be able to collapse enough to reach stellar densities in their centres.","Citation Text":["Rees 1976"],"Functions Text":["At very large densities, though, owing to the increased dust absorption coefficient, the gas becomes opaque to its own radiation. This is sometimes referred to as the opacity limit","At this point, the fragmentation is halted, setting a scale below which no fragments form."],"Functions Label":["Background","Background"],"Citation Start End":[[1135,1144]],"Functions Start End":[[929,1109],[1147,1237]]}
{"Identifier":"2016AandA...591A..13V__Giovannini_et_al._2013_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":["Giovannini et al. 2013"],"Functions Text":["On scales up to a few Mpc from the nearest galaxy cluster, possibly along filaments, only a few diffuse synchrotron sources have been reported"],"Functions Label":["Background"],"Citation Start End":[[866,888]],"Functions Start End":[[660,802]]}
{"Identifier":"2021ApJ...915L...8D__Chen_et_al._2020_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":["Chen et al. 2020"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[726,742]],"Functions Start End":[[548,702]]}
{"Identifier":"2015ApJ...811L..32H__Liewer_et_al._2001_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":["Liewer et al. 2001"],"Functions Text":["The expanding box uses these comoving coordinates, approximating the spherical coordinates by the Cartesian ones"],"Functions Label":["Uses"],"Citation Start End":[[728,746]],"Functions Start End":[[614,726]]}
{"Identifier":"2021AandA...654A.132B__Davies_et_al._2014b_Instance_1","Paragraph":"We can use the equivalent width of H\u03b1 (EWH\u03b1, see Table 4) to put additional constraints on the recent star formation history of the AGNs. To do so, we first need to estimate the largest fraction of the H\u03b1 luminosity that could arise from star formation, and for this we use the group of AGNs with the highest EWH\u03b1. There are a number of reasons why AGNs could have lower EWH\u03b1, but the maximum values will only be found when both AGN and star formation contribution are maximal. The group of six AGNs with the highest EWH\u03b1 consists of ESO 137-G034, NGC 3081, NGC 2110, NGC 7582, NGC 2992, and NGC 5728, and their median EWH\u03b1 is 70 \u00c5. These also tend to have the highest log [OIII]\/H\u03b2, with a median ratio of 1.07. Similarly, their median log [NII]\/H\u03b1 is 0.01. To assess the star formation contribution to the H\u03b1 flux, we consider lines from the most extreme AGN photoionisation models of Groves et al. (2004), to the location of solar metallicity star-forming galaxies (see also Davies et al. 2014b). Since the AGNs are very close to the highest [OIII]\/H\u03b2 ratios that the models can produce, we find that at most 10% of their H\u03b1 flux could be due to star formation. Applying this correction reduces the EWH\u03b1 that is associated with star formation to  7 \u00c5. We also need to make a correction for the old stellar population, since EWH\u03b1 in models compares the line flux to the continuum due to stars associated with the line emission. We find that 5% of the continuum is associated with young stars. We therefore estimate that the maximum EWH\u03b1 associated with the most recent star-forming episode is  7 \u00c5\/5%, which is  140 \u00c5. Comparing this to Starburst99 models (specifically Figs. 83 and 84 in Leitherer et al. 1999) rules out continuous star formation since it would imply a timescale exceeding several Gyr, inconsistent with the stellar population synthesis results. Instead, it means that the recent star formation must have ceased. Based on the rate at which the EWH\u03b1 falls after star formation stops, we estimate that the end of star formation must have typically happened at least 6 Myr prior.","Citation Text":["Davies et al. 2014b"],"Functions Text":["To assess the star formation contribution to the H\u03b1 flux, we consider lines from the most extreme AGN photoionisation models of Groves et al. (2004), to the location of solar metallicity star-forming galaxies (see also"],"Functions Label":["Uses"],"Citation Start End":[[978,997]],"Functions Start End":[[759,977]]}
{"Identifier":"2022AandA...662A...8M__Shu_et_al._1987_Instance_1","Paragraph":"Even before studying the relationship between the IMF and the CMF, it is important to realize that how the IMF originates from the observed CMF depends directly on the definition of the cores, assumed to be the gas mass reservoir used for the formation of each star or binary system. As shown by Louvet et al. (2021), defining this mass reservoir may seem obvious in the observed map of a cloud, but core characteristics (size, mass) depend heavily on the spatial scales probed by the observations. In addition, the theoretical definition of cores also depends on whether the star-formation scenario is quasi-static or dynamic. In the former scenario, cores are gas condensations sufficiently dense to be on the verge of gravitational collapse, and they convert the core gas into stars (Shu et al. 1987; Chabrier 2003; McKee & Ostriker 2007; Andr\u00e9 et al. 2014). After a quasi-static phase of concentration of the cloud gas into cores, cores become distinct from their surrounding cloud and start to collapse, and their future stellar content becomes independent of the properties of the parental cloud. In the latter scenario, dynamics play a major role during all phases of the star-formation process (e.g., Ballesteros-Paredes et al. 2007; Hennebelle & Falgarone 2012; Padoan et al. 2014). In particular, global infall of filament networks and gas inflow toward cores are expected to be important drivers of star formation (e.g., Smith et al. 2009; V\u00e1zquez-Semadeni et al. 2019; Padoan et al. 2020). In this framework, filaments, cores, and stellar embryos simultaneously accrete gas, and the gas reservoir associated with star formation largely exceeds the extent of the observed cores. This so-called clump-fed scenario was proposed in various recent papers and described in detail in the review by Motte et al. (2018a, see references therein). One of the main objectives of the ALMA-IMF Large Program is to discriminate between the quasi-static and dynamic scenarios by quantifying the role of cloud kinematics in defining core mass and in possibly changing it over time.","Citation Text":["Shu et al. 1987"],"Functions Text":["In addition, the theoretical definition of cores also depends on whether the star-formation scenario is quasi-static or dynamic. In the former scenario, cores are gas condensations sufficiently dense to be on the verge of gravitational collapse, and they convert the core gas into stars"],"Functions Label":["Background"],"Citation Start End":[[787,802]],"Functions Start End":[[499,785]]}
{"Identifier":"2021AandA...654A.126B__Hotta_(2017)_Instance_1","Paragraph":"Analytical and semianalytical approaches have been developed to describe this process and estimate the width of the overshooting layer (e.g., Schmitt et al. 1984; Zahn 1991; Rempel 2004). With the improvement of computational methods and resources, an increasing number of studies have been devoted to numerical simulations of convection and overshooting using realistic stellar conditions (geometry, luminosity, thermal diffusivity, equation of state, opacities, etc.). A commonly used tactic to increase the efficiency and improve the stability of these simulations is to artificially increase the luminosity (or nuclear energy for convective cores or burning shells) and to modify the thermal diffusivity of the reference stellar model. This tactic is common and has been used, for example, in Rogers et al. (2006, 2013), Meakin & Arnett (2007), Tian et al. (2009), Brun et al. (2011, 2017), Hotta (2017), Cristini et al. (2017), Edelmann et al. (2019), Horst et al. (2020). This approach is used to increase the Mach number of the convective flow, reducing the disparity between advective and acoustic timescales, improving the efficiency of time-explicit codes limited by the Courant\u2013Friedrich\u2013Levy constraint. It is also used to provide numerical stability or to accelerate the thermal relaxation. But no examination of its potentially far-reaching impact has been conducted. Rempel (2004) pointed out that enhanced energy flux in numerical simulations could lead to unrealistically vigourous convection, which could impact the properties of the overshooting layer and could explain some of the discrepancies between analytical models and numerical simulations. Numerical simulations also suggest an increase in the overshooting depth with increasing flux (Hotta 2017; K\u00e4pyl\u00e4 2019). Determining scaling laws of the overshooting depth as a function of the energy input could thus allow for an extrapolation of the results to more realistic stellar conditions and help to estimate the overshooting depth in real stars, as suggested by Hotta (2017), for example, for the Sun. But K\u00e4pyl\u00e4 (2019) also shows that an artificial modification of the heat conductivity in the radiative and overshooting regions could impact the overshooting process. In some computational studies, both the luminosity and the thermal diffusivity are enhanced by the same factor to ensure that the thermal structure is unchanged compared to the reference stellar structure and with the expectation that the larger thermal diffusivity counterbalances the larger energy flux. This procedure has been proposed as a way to provide a good representation of the true dynamics of the system (e.g., Rogers et al. 2006, 2013; Tian et al. 2009). But this expectation has never been demonstrated. Another expectation concerns internal waves, excited by convective motions and by flows penetrating the convective boundary. Simulations with artificially modified luminosity and thermal diffusivity are also used to perform the analysis of internal waves, either for convective envelopes (e.g., Rogers et al. 2006; Brun et al. 2011; Alvan et al. 2014) or for convective cores (e.g., Rogers & McElwaine 2017; Edelmann et al. 2019; Horst et al. 2020). None of these works have examined whether the wave spectrum of a realistic star is accurately predicted by such simulations.","Citation Text":["Hotta (2017)","Hotta 2017","Hotta (2017)"],"Functions Text":["A commonly used tactic to increase the efficiency and improve the stability of these simulations is to artificially increase the luminosity (or nuclear energy for convective cores or burning shells) and to modify the thermal diffusivity of the reference stellar model. This tactic is common and has been used, for example, in","This approach is used to increase the Mach number of the convective flow, reducing the disparity between advective and acoustic timescales, improving the efficiency of time-explicit codes limited by the Courant\u2013Friedrich\u2013Levy constraint. It is also used to provide numerical stability or to accelerate the thermal relaxation. But no examination of its potentially far-reaching impact has been conducted."],"Functions Label":["Background","Background"],"Citation Start End":[[895,907],[1763,1773],[2039,2051]],"Functions Start End":[[471,796],[978,1381]]}
{"Identifier":"2019ApJ...882...97P__Warmuth_&_Mann_2016_Instance_1","Paragraph":"The RADYN code (Carlsson & Stein 1992, 1995, 1997; Allred et al. 2015) is a one-dimensional radiative hydrodynamic code that can be used to study the interaction of particle beams with the solar atmosphere. It uses the Fokker\u2013Planck formalism (McTiernan & Petrosian 1990), which takes into account the beam energy losses due to Coulomb collisions and pitch-angle diffusion when incorporating relativistic effects. RADYN includes a six-level hydrogen atom, a nine-level helium atom, and a six-level calcium atom. A return current has also been included in the simulations. We generated a set of models that simulates the conditions in the solar atmosphere during weak to intense WLFs. The beam fluxes used had values of 3 \u00d7 109, 1 \u00d7 1010, and 3 \u00d7 1010 erg cm\u22122 s\u22121, while the low-energy cutoff EC covered the parameter space where flare values are usually found (20\u2013120 keV; Warmuth & Mann 2016). The spectral index \u03b4 was equal to 3 for all models. The beams were applied continuously and the outputs were analyzed at t = 20 s. This value is consistent with the X-ray analysis of the best observed type II WLF to date (Proch\u00e1zka et al. 2018). The initial atmosphere used in this work has the transition region placed at a height of 1200 km above the photospheric floor and has a coronal temperature of 3 MK at 10 Mm (QS.SL.HT loop described in Allred et al. 2015). The beams were injected at the top of a half-loop with a Gaussian distribution with an HWHM of 235. Line synthesis was carried out using the RH code (Uitenbroek 2001) incorporating partial redistribution that is particularly important for resonance line profiles. We used the 20-level hydrogen atom and the 6-level calcium atom to produce synthetic spectra for the Lyman and Balmer line and continuum diagnostics. A VOIGT profile was used for Ca ii K line, Lyman \u03b3 the higher-order Lyman lines. The Ca ii H and Ly\u03b1 and \u03b2 profiles were modeled in PRD. The Balmer lines had profiles of type VOIGT_VCS_STARK, which incorporates the unified Stark effect theory (Kowalski et al. 2017). The spectra were synthesized by setting a minimal spectral resolution of 0.05 nm.","Citation Text":["Warmuth & Mann 2016"],"Functions Text":["The beam fluxes used had values of 3 \u00d7 109, 1 \u00d7 1010, and 3 \u00d7 1010 erg cm\u22122 s\u22121, while the low-energy cutoff EC covered the parameter space where flare values are usually found (20\u2013120 keV;"],"Functions Label":["Uses"],"Citation Start End":[[874,893]],"Functions Start End":[[684,873]]}
{"Identifier":"2017ApJ...835..101H__Vanderbeke_et_al._2014_Instance_1","Paragraph":"In the present survey of GCSs in BCGs, we use the color index \n\n\n\n\n\n (from here on we drop the accents on the SDSS indices). In the following discussion it will be useful to have a calibration of this index versus cluster metallicity [Fe\/H]. To do this, we would ideally need to have GC photometry of the same clusters in both the Kron-Cousins and SDSS systems, in addition to spectroscopically based metallicity measurements. At present, there are no ideal solutions to that problem. Galaxies satisfying all three of these criteria are rare; in principle the Milky Way GC databases could be used, but cluster-to-cluster foreground reddenings differ strongly, the published SDSS indices (Vanderbeke et al. 2014) show considerable scatter versus metallicity, and the variety of studies from which the UBVRI indices were derived are completely different from the SDSS survey, so that aperture-size mismatches are significant. Similar problems affect the M31 GC sample. The best option at the present time for developing a \n\n\n\n\n\n transformation is likely to be from the nearby early-type giant galaxy NGC 5128: here, UBVRI photometry is available from Peng et al. (2004), griz photometry from Sinnott et al. (2010), and [Fe\/H] values derived through \n\n\n\n\n\n from Woodley et al. (2010); these [Fe\/H] values are in turn well correlated with the Sloan-system spectroscopic index [MgFe]\u2019 (see Woodley et al.). We have extracted the GCs in common from these three catalogs, with the results shown in Figure 1. The great majority of these GCs lie well outside the central few kiloparsecs of NGC 5128 and thus are unaffected by the well-known dust lane. We have therefore applied only the foreground reddening of the Galaxy, for which we adopt \n\n\n\n\n\n (Cardelli et al. 1989) to obtain the intrinsic colors. We note, however, that the UBVRI measurements were done on \n\n\n\n\n\n aperture diameters corrected to \n\n\n\n\n\n through median curves of growth (Peng et al. 2004), while the griz measures were done through \n\n\n\n\n\n apertures (Sinnott et al. 2010), which means that a small aperture mismatch may exist here as well that affects the zero-point of \n\n\n\n\n\n.","Citation Text":["Vanderbeke et al. 2014"],"Functions Text":["Galaxies satisfying all three of these criteria are rare; in principle the Milky Way GC databases could be used, but cluster-to-cluster foreground reddenings differ strongly, the published SDSS indices","show considerable scatter versus metallicity, and the variety of studies from which the UBVRI indices were derived are completely different from the SDSS survey, so that aperture-size mismatches are significant."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[688,710]],"Functions Start End":[[485,686],[712,923]]}
{"Identifier":"2020ApJ...904...81G__Nariyuki_&_Hada_2007_Instance_1","Paragraph":"Alfv\u00e9n waves of arbitrary amplitude with constant total pressure are known to provide an exact solution to the compressible magnetohydrodynamic (MHD) system in a homogeneous plasma, in that nonlinearities are turned off and there no couplings with compressible modes. However, such a dynamical system is linearly unstable to parametric instabilities and large-amplitude Alfv\u00e9n waves are known to decay into compressible and secondary Alfv\u00e9nic modes through three- or four-wave resonances that lead to a variety of parametric instabilities, depending on the plasma beta and dispersive effects, such is the case of parametric decay (Galeev & Oraevskii 1973; Derby 1978), modulational, and beat instabilities (Sakai & Sonnerup 1983; Wong & Goldstein 1986; Jayanti & Hollweg 1993; Nariyuki & Hada 2007). Parametric instabilities of Alfv\u00e9n waves (or of a spectrum of Alfv\u00e9n waves) have been widely studied over the years through theoretical approaches (Goldstein 1978; Jayanti & Hollweg 1993; Malara & Velli 1996), and numerical simulations adopting both MHD (Ghosh et al. 1994; Malara et al. 2000; Del Zanna et al. 2001; Tenerani et al. 2017) and kinetic models (Terasawa et al. 1986; Matteini et al. 2010; Nariyuki et al. 2012; Verscharen et al. 2012), although most often in one-dimensional setups. In particular, the traditional parametric decay instability has attracted much attention over the years in the context of both turbulence and plasma heating. This type of decay is most efficient at low values of the plasma beta and it essentially involves the decay of a pump Alfv\u00e9n wave into a lower-frequency reflected Alfv\u00e9n wave and a forward sound wave. For this reason, parametric decay remains an appealing process because it provides a natural mechanism for the production of reflected modes, which is essential for the triggering of a turbulent cascade. Indeed, recently it has been proposed as a viable mechanism to initiate the turbulent cascade in the solar wind acceleration region (Chandran 2018; R\u00e9ville et al. 2018), while global MHD simulations of the solar wind have also shown that the parametric decay instability can contribute substantially to solar wind heating and acceleration, thanks to the generation of compressible modes that, in the absence of kinetic effects, naturally steepen into shocks (see, e.g., Shoda et al. 2019). The traditional parametric decay has been also invoked as a possible source for the generation of inward modes and solar wind turbulence in the inner heliosphere, where an increasing content of reflected waves (cross-helicity) and an evolving turbulent spectrum are observed with increasing heliocentric distance (Bavassano et al. 2000). However, expansion effects are known to inhibit its development, essentially because the parametric decay process is strongly suppressed as the plasma beta increases at larger heliocentric distances (Tenerani & Velli 2013, 2020; Del Zanna et al. 2015). Temperature anisotropies can destabilize the parametric decay at values of the plasma beta approaching unity and above, but the anisotropy in the solar wind is not large enough to affect the instability significantly (Tenerani et al. 2017).","Citation Text":["Nariyuki & Hada 2007"],"Functions Text":["However, such a dynamical system is linearly unstable to parametric instabilities and large-amplitude Alfv\u00e9n waves are known to decay into compressible and secondary Alfv\u00e9nic modes through three- or four-wave resonances that lead to a variety of parametric instabilities, depending on the plasma beta and dispersive effects, such is the case of parametric decay"],"Functions Label":["Background"],"Citation Start End":[[777,797]],"Functions Start End":[[268,629]]}
{"Identifier":"2020MNRAS.498.1801K__Gundlach_et_al._2011_Instance_1","Paragraph":"Water ice is ubiquitous in the cold regions of the Universe, owing to the fact that hydrogen and oxygen are the two most abundant elements to form a solid such as icy dust particles and comets. It is, therefore, commonly accepted that the essential component of dust particles and planetesimals in protoplanetary discs is water ice beyond the so-called snow line, at which the temperature of gas is low enough for water vapour to condense into ices (e.g. Cyr, Sears & Lunine 1998). Reactive accretion of water ice from hydrogen and oxygen atoms on the surface of dust particles takes place in the dense core of molecular clouds where the growth of dust particles has been observed by scattering of stellar radiation (Steinacker et al. 2010). It is worthwhile noting that laboratory experiments on the coagulation growth of water-ice particles have a long history outside astronomy and planetary science, since coagulation is observed in daily life and is a plausible route to the formation of snowflakes (e.g. Faraday 1860; Hosler, Jensen & Goldshlak 1957). Recent works on laboratory measurements of cohesion between crystalline water-ice particles at vacuum conditions provided encouraging results that dust particles composed of water ice might be much more cohesive than previously believed (Gundlach et al. 2011; Gundlach & Blum 2015; Jongmanns et al. 2017). Form a theoretical point of view, Chokshi, Tielens & Hollenbach (1993) demonstrated that the JKR theory of elastic contact formulated by Johnson, Kendall & Roberts (1971) is a powerful tool for better understanding of dust coagulation. Numerical simulations incorporating the JKR theory have shown that dust aggregates consisting of submicrometre-sized water-ice particles proceed with coagulation growth even at a collision velocity of 50 m s\u22121 (Wada et al. 2009, 2013). As a result, the majority of recent studies on dust coagulation and planetesimal formation assume that silicate aggregates are disrupted by mutual collision at a velocity of vdisrupt \u223c 1 m s\u22121, but icy aggregates at vdisrupt \u223c 10 m s\u22121 (e.g. Birnstiel, Dullemond & Brauer 2010; Vericel & Gonzalez 2019). Such a trendy assumption led Dr\u0105\u017ckowska & Alibert (2017) to propose planetesimal formation by the \u2018traffic jam\u2019 effect at the snow line, provided that sticky water-ice particles grow faster and thus drift toward the central star faster than less-sticky bare silicate particles, implying that aggregates of the former catch up the latter at the snow line, which results in a traffic jam. However, we argue that the importance of water ice to dust coagulation is still open to debate, since water ice is not necessarily stickier than other materials such as silicates and complex organic matter (Kimura et al. 2015, 2020a; Musiolik & Wurm 2019).","Citation Text":["Gundlach et al. 2011"],"Functions Text":["Recent works on laboratory measurements of cohesion between crystalline water-ice particles at vacuum conditions provided encouraging results that dust particles composed of water ice might be much more cohesive than previously believed"],"Functions Label":["Background"],"Citation Start End":[[1296,1316]],"Functions Start End":[[1058,1294]]}
{"Identifier":"2016MNRAS.463.3204R__Trautman_1967_Instance_1","Paragraph":"Since the perturbation in equation (22) is not conservative, we shall focus on the extension of Noether's theorem to non-conservative systems by Djukic & Vujanovic (1975). It must be $F_i-q^\\prime _if\\ne 0$ for the conservation law to hold (Vujanovic, Strauss & Jones 1986). For the case of non-conservative systems the generators Fi, f, and the gauge \u03a8 need to satisfy the following relation:\n\n(29)\n\n\\begin{eqnarray}\n&amp;&amp;{\\sum _i\\left\\lbrace \\left( \\frac{\\mathrm{\\partial} {\\scr {L}}}{\\mathrm{\\partial} {q_i}} \\right) F_i + \\left(\\frac{\\mathrm{\\partial} {\\scr {L}}}{\\mathrm{\\partial} {q^\\prime _i}} \\right)( {F}^\\prime _i - q^\\prime _i{f}^\\prime ) + Q_i(F_i-q^\\prime _if)\\right\\rbrace} \\nonumber\\\\\n&amp;&amp;{\\quad+ f^\\prime \\scr {L} + f\\frac{\\mathrm{\\partial} {\\scr {L}}}{\\mathrm{\\partial} {\\tau }} = {\\Psi }^\\prime .}\n\\end{eqnarray}\n\nThis equation and the condition $F_i-q^\\prime _if\\ne 0$ furnish the generalized NBH equations (Trautman 1967; Djukic & Vujanovic 1975; Vujanovic et al. 1986). The NBH equations involve the full derivative of the gauge function and the generators with respect to \u03c4, meaning that equation (29) depends on the partial derivatives of \u03a8, Fi, and f with respect to time, the coordinates, and the velocities. By expanding the convective terms the NBH equations decompose in the system of Killing equations:\n\n(30)\n\n\\begin{eqnarray}\n&amp;&amp;{ \\scr {L} \\frac{\\mathrm{\\partial} {f}}{\\mathrm{\\partial} {q_j^\\prime }} +\\sum _i \\frac{\\mathrm{\\partial} {\\scr {L}}}{\\mathrm{\\partial} {q^\\prime _i}}\\left( \\frac{\\mathrm{\\partial} {F_i}}{\\mathrm{\\partial} {q_j^\\prime }} - q^\\prime _i\\frac{\\mathrm{\\partial} {f}}{\\mathrm{\\partial} {q_j^\\prime }} \\right) = \\frac{\\mathrm{\\partial} {\\Psi }}{\\mathrm{\\partial} {q_j^\\prime }},} \\nonumber \\\\\n&amp;&amp;{ \\frac{\\mathrm{\\partial} {}}{\\mathrm{\\partial} {\\tau }}(f\\scr {L} - \\Psi ) + \\sum _i \\Bigg\\lbrace \\frac{\\mathrm{\\partial} {\\scr {L}}}{\\mathrm{\\partial} {q_i}}F_i + \\scr {L}\\frac{\\mathrm{\\partial} {f}}{\\mathrm{\\partial} {q_i}}q_i^\\prime + Q_i(F_i-q_i^\\prime f)+ \\frac{\\mathrm{\\partial} {\\scr {L}}}{\\mathrm{\\partial} {q_i^\\prime }}} \\nonumber \\\\\n&amp;&amp;{\\times\\, \\left[ \\frac{\\mathrm{\\partial} {F_i}}{\\mathrm{\\partial} {\\tau }} - q_i^\\prime \\frac{\\mathrm{\\partial} {f}}{\\mathrm{\\partial} {\\tau }} + \\sum _j\\left( \\frac{\\mathrm{\\partial} {F_i}}{\\mathrm{\\partial} {q_j}}q_j^\\prime - q_i^\\prime q_j^\\prime \\frac{\\mathrm{\\partial} {f}}{\\mathrm{\\partial} {q_j}} \\right) \\right] - \\frac{\\mathrm{\\partial} {\\Psi }}{\\mathrm{\\partial} {q_i}}q_i^\\prime \\Bigg \\rbrace = 0.}\\nonumber\\\\\n\\end{eqnarray}\nThe system (30) decomposes in three equations that can be solved for the generators F\u03c1, F\u03b8, and f given a certain gauge. If the transformation defined in equation (26) satisfies the NBH equations, then the system admits the integral of motion (28).","Citation Text":["Trautman 1967"],"Functions Text":["This equation and the condition $F_i-q^\\prime _if\\ne 0$ furnish the generalized NBH equations"],"Functions Label":["Uses"],"Citation Start End":[[938,951]],"Functions Start End":[[843,936]]}
{"Identifier":"2017MNRAS.469.2662D__Morales_et_al._2012_Instance_1","Paragraph":"A well-studied way to look at antenna spectral requirements is from the perspective of foreground avoidance in power spectrum space. In the avoidance scheme, smooth-spectrum foregrounds should, in the ideal case, occupy a wedge\u2013shaped region of the two\u2013dimensional power spectrum space (where the wavenumber k can be decomposed into k\u22a5 and $k_\\Vert$ that are the transverse and line-of-sight wave numbers, respectively) whereas the remaining area \u2013 the so-called EoR window \u2013 is dominated by the 21 cm emission (Datta, Bowman & Carilli 2010; Morales et al. 2012; Trott, Wayth & Tingay 2012; Vedantham, Udaya Shankar & Subrahmanyan 2012; Thyagarajan et al. 2013; Liu, Parsons & Trott 2014a,b). Sources have most of their emission at low k\u2225 values although, due to the inherent chromatic interferometric response, this area increases with baseline length to resemble a characteristic wedge-like shape (Parsons et al. 2012; Trott et al. 2012; Vedantham et al. 2012; Liu et al. 2014a,b; Thyagarajan et al. 2015). In the most optimistic scenario, the EoR power spectrum can be directly measured in the EoR window, whose boundaries are set by the so-called horizon limit, i.e. the maximum delay that an astrophysical signal can experience, given by the separation of the two receiving elements. In practice, the boundaries of the EoR window can be narrowed by a number of mechanisms that spread power from the foreground dominated region into the EoR window, in particular calibration errors (Barry et al. 2016; Patil et al. 2016), leakage of foreground polarization (Bernardi et al. 2010; Jeli\u0107 et al. 2010; Moore et al. 2013; Asad et al. 2015, 2016) and intrinsic chromaticity of the instrumental response. Recent attention has indeed been given to simulate and characterize the element response, particularly for the HERA (DeBoer et al. 2017), realizing that it may be one of the critical items responsible for spilling power from the wedge into the EoR window.","Citation Text":["Morales et al. 2012"],"Functions Text":["In the avoidance scheme, smooth-spectrum foregrounds should, in the ideal case, occupy a wedge\u2013shaped region of the two\u2013dimensional power spectrum space (where the wavenumber k can be decomposed into k\u22a5 and $k_\\Vert$ that are the transverse and line-of-sight wave numbers, respectively) whereas the remaining area \u2013 the so-called EoR window \u2013 is dominated by the 21 cm emission"],"Functions Label":["Background"],"Citation Start End":[[542,561]],"Functions Start End":[[133,510]]}
{"Identifier":"2019ApJ...875...90L__Velli_et_al._2015_Instance_3","Paragraph":"When energy flows from the interior of the Sun outward into the solar atmosphere, why is the Sun\u2019s outer atmosphere, the corona, much hotter than the inner atmosphere, the underlying chromosphere and photosphere? This is the long-standing problem of the coronal heating, which is one of the eight key mysteries in modern astronomy (Kerr 2012). For about 80 yr since the discovery of the extremely hot corona around the late 1930s (Grotian 1939; Edlen 1945), people have worked hard on addressing this issue, and great advances have been made in observation and theoretical studies (Parnell & De Moortel 2012; Amari et al. 2015; Arregui 2015; Cargill et al. 2015; De Moortel & Browning 2015; Jess et al. 2015; Klimchuk 2015; Longcope & Tarr 2015; Peter 2015; Schmelz & Winebarger 2015; Velli et al. 2015; Wilmot-Smith 2015). Especially during recent decades, high-resolution observations of solar super-fine structures indicate that small spicules, minor hot jets along small-scale magnetic channels from the low atmosphere upwards to the corona, petty tornados and cyclones, and small explosive phenomena such as mini-filament eruptions and micro- and nano-flares\u2014all of these small-scale magnetic activities contribute greatly to coronal heating (De Pontieu et al. 2011; 2018; Zhang & Liu 2011; Parnell & De Moortel 2012; Klimchuk 2015; Peter 2015; Schmelz & Winebarger 2015; Henriques et al. 2016; Li et al. 2018a). Additionally, contributions of MHD waves to heating the corona have been observationally illustrated (van Ballegooijen et al. 2011; Jess et al. 2015; Kubo et al. 2016; Morton et al. 2016; Soler et al. 2017; Morgan & Hutton 2018). Meanwhile, with the progress of observational studies, two groups of theoretical models, magnetic reconnection models and magnetohydrodynamic wave models, have traditionally attempted to explain coronal heating, but so far no models can address the key mystery perfectly (van Ballegooijen et al. 2011; Arregui 2015; Cargill et al. 2015; Peter 2015; Velli et al. 2015; Wilmot-Smith 2015). Maybe we do not need to intentionally take to heart such the classical dichotomy, because waves and reconnections may interact with each other (De Moortel & Browning 2015; Velli et al. 2015). Additionally, statistical studies may look at coronal heating from a comprehensive perspective. Li et al. (2018b) found that the long-term variation of the heated corona, which is represented by coronal spectral irradiances, and that of small-scale magnetic activity are in lockstep, indicating that the corona should statistically be effectively heated by small-scale magnetic activity. Observational and theoretical model studies through heating channels and modes, and statistical studies by means of heating effect (performance of the heated corona), both suggest that coronal heating originates from small-scale magnetic activity.","Citation Text":["Velli et al. 2015"],"Functions Text":["Maybe we do not need to intentionally take to heart such the classical dichotomy, because waves and reconnections may interact with each other"],"Functions Label":["Motivation"],"Citation Start End":[[2208,2225]],"Functions Start End":[[2036,2178]]}
{"Identifier":"2015MNRAS.452.2731S__Stroe_et_al._2013_Instance_1","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"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[1251,1268]],"Functions Start End":[[1057,1225]]}
{"Identifier":"2020AandA...644A..97C__Leroy_et_al._2013_Instance_2","Paragraph":"Major nearby galaxy cold gas mapping surveys (Regan et al. 2001; Wilson et al. 2009; Rahman et al. 2011; Leroy et al. 2009; Donovan Meyer et al. 2013; Bolatto et al. 2017; Sorai et al. 2019; Sun et al. 2018) have focused on observations of the molecular gas (through CO lines). Despite a few notable exceptions (e.g. Alatalo et al. 2013; Saintonge et al. 2017), these surveys observed mainly spiral or infrared-bright galaxies (i.e. galaxies with significant star formation) and have furthered our understanding of how star formation happens, rather than how it stops. This boils down to quantifying the relation between molecular gas and star formation rate (SFR), which appears nearly linear in nearby discs (Kennicutt 1998; Bigiel et al. 2008; Leroy et al. 2013; Lin et al. 2019). This relationship is often parametrised via the ratio between the SFR and the molecular gas mass (Mmol), which is called the molecular star formation efficiency (SFE = SFR\/Mmol = 1\u2215\u03c4dep), where the inverse of the SFE is the depletion time, \u03c4dep. The depletion time indicates how much time is necessary to convert all the available molecular gas into stars at the current star formation rate. On kpc scales and in the discs of nearby star-forming galaxies, \u03c4dep is approximately constant around 1\u20132 Gyr (Bigiel et al. 2011; Rahman et al. 2012; Leroy et al. 2013; Utomo et al. 2017), and it appears to weakly correlate with many galactic properties such as stellar mass surface density or environmental hydrostatic pressure (Leroy et al. 2008; Rahman et al. 2012). Nevertheless, small but important deviations for a constant SFE have been noticed, which can be the first hints of star formation quenching. In some galaxies, the depletion time in the centres appear shorter (Leroy et al. 2013; Utomo et al. 2017) or longer (Utomo et al. 2017) with respectto their discs. These differences may correlate with the presence of a bar or with galaxy mergers (Utomo et al. 2017; see also Muraoka et al. 2019) and do not seem to be related to unaccounted variation in the CO-to-H2 conversion factor (Leroy et al. 2013; Utomo et al. 2017). Spiral arm streaming motions have also been observed to lengthen depletion times (Meidt et al. 2013; Leroy et al. 2015).","Citation Text":["Leroy et al. 2013"],"Functions Text":["On kpc scales and in the discs of nearby star-forming galaxies, \u03c4dep is approximately constant around 1\u20132 Gyr"],"Functions Label":["Background"],"Citation Start End":[[1327,1344]],"Functions Start End":[[1176,1285]]}
{"Identifier":"2022ApJ...930...70H__Vogl_et_al._2020_Instance_1","Paragraph":"There are some existing studies applying machine learning to transient studies. For example, the spectral types of the SNe can be classified based on their light-curve data (M\u00f6ller et al. 2016; Muthukrishna et al. 2019a; Takahashi et al. 2020; Villar et al. 2020), and transients can be identified from the astronomical survey images (Goldstein et al. 2015; Mahabal et al. 2019; G\u00f3mez et al. 2020). The light curves of SNe Ia can be well modeled by functional principal component analysis (FPCA; He et al. 2018), where it was shown remarkably that a set of FPCA eigenvectors that are independent of the photometric filters can be derived from the observed light curves of SNe Ia. There are a few studies of the application of deep learning neural networks to the spectral data of SNe. For example, Muthukrishna et al. (2019b) used a convolutional neural network (CNN) for automated SN type classification based on SN spectra. Several other works (Chen et al. 2020; Vogl et al. 2020; Kerzendorf et al. 2021) applied a Gaussian process, principal component analysis (PCA), and deep learning neural networks to radiative transfer models of SNe. Sasdelli et al. (2016) used unsupervised learning algorithms to investigate the subtypes of SNe Ia. Stahl et al. (2020) developed neural networks to predict the photometric properties of SNe Ia (phase and \u0394m\n15) based on spectroscopic data. Saunders et al. (2018) used PCA to find low dimensional representations of the spectral sequences of 140 well-observed SNe Ia. Chen et al. (2020), in particular, built an artificial intelligence assisted inversion (AIAI) of radiative transfer models and used that to link the observed SN spectra with theoretical models. The AIAI is able to retrieve the elemental abundances and density and temperature profiles from observed SN spectra. The AIAI approach has the potential for quantitatively coupling complex theoretical models with the ever-increasing amount of high-quality observational data.","Citation Text":["Vogl et al. 2020"],"Functions Text":["Several other works","applied a Gaussian process, principal component analysis (PCA), and deep learning neural networks to radiative transfer models of SNe."],"Functions Label":["Background","Background"],"Citation Start End":[[965,981]],"Functions Start End":[[926,945],[1007,1141]]}
{"Identifier":"2021AandA...650A.203G__Chiosi_1980_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":["Chiosi 1980"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[1045,1056]],"Functions Start End":[[862,1043]]}
{"Identifier":"2022AandA...659A..54E__Kendall_et_al._1992_Instance_1","Paragraph":"Exploration of the stationary points of the [H, P, S, O] molecular system was initially carried out using M\u00f8ller\u2013Plesset second order Perturbation Theory (MP2) (M\u00f8ller & Plesset 1934; Frisch et al. 1990a,b; Head-Gordon & Head-Gordon 1994; Head-Gordon et al. 1988; S\u00e6b\u00f8 & Alml\u00f6f 1989) with the 6-311++G(d, p) basis set (McLean & Chandler 1980; Krishnan et al. 1980). Following the initial survey, transition state structures connecting various isomers of [H, P, S, O] were located using MP2\/6-311++G(d, p) and confirmed via an intrinsic reaction coordinate (IRC) calculation. To achieve a quantitative picture of the ground state potential energy surface, coupled cluster theory with single, double, and perturbative triple excitations [CCSD(T)] (Knowles et al. 1993, 2000) single point calculations employing the aug-cc-pV(Q+d)Z basis set were performed at the MP2 equilibrium geometries. Entrance channel pathways were explored via relaxed scans and individual searches for transition states. Following this, higher level calculations on the electronic structure, rotational constants, and harmonic vibrational frequencies of each isomer were performed using CCSD(T) together with the Dunning basis sets aug-cc-pV(X+d)Z (X = T, Q, 5) (Dunning 1989; Woon & Dunning 1993; Kendall et al. 1992), including additional tight d functions on the sulfur and phosphorous atoms (Dunning et al. 2001). CCSD(T) including all electrons was then performed using the second order Douglas-Kroll\u2013Hess Hamiltonian in conjunction with the contracted relativistic Douglas\u2013Kroll aug-cc-pwCVTZ-DK basis set to correct for scalar relativistic and core-correlation effects (Jorge et al. 2009). The aforementioned combinationof level of theory and basis set will be referred to as CCSD(T)-AE\/TZ-DK. Finally, explicit correlation of the electrons was included using the CCSD(T)-F12b method (Werner et al. 2011) with the Dunning basis sets aug-cc-pVXZ (X = T, Q, 5). Single-reference character of the wavefunction was confirmed via single-state complete active spaceself-consistent field theory in addition to the coupled cluster T1 diagnostic. Energies and geometry parameters were then extrapolated to the complete basis set limit using the two-point (Q, 5) extrapolation scheme E(x)= ECBS + AX\u22123, where E(x) is the value calculated using the basis set of cardinal number X, ECBS is the extrapolated energy, and A is a parameter fit in the least-squares fitting procedure. Mulliken population analysis was performed on cis-HOPS using MOLPRO. Corrections to the rotational constants, spectroscopic data, and anharmonic vibrational data were calculated for the four lowest energy isomers at the CCSD(T)\/aug-cc-pV(T+d)Z level of theory using CFOUR (Matthews et al. 2020; Stanton et al. 2019). Permanent dipole moments for each isomer were calculated using the finite field method (field strengths of 0, 0.005, and \u22120.005 au) as implemented in MOLPRO.","Citation Text":["Kendall et al. 1992"],"Functions Text":["Following this, higher level calculations on the electronic structure, rotational constants, and harmonic vibrational frequencies of each isomer were performed using CCSD(T) together with the Dunning basis sets aug-cc-pV(X+d)Z (X = T, Q, 5)"],"Functions Label":["Uses"],"Citation Start End":[[1271,1290]],"Functions Start End":[[994,1234]]}
{"Identifier":"2019MNRAS.484.1359Y__Melrose_&_Yuen_2014_Instance_1","Paragraph":"The discrete change in spin-down rate between the on and off states suggests a switching behaviour in the pulsar magnetosphere. While explanation for the phenomenon remains inconclusive, conventional treatments typically involve referring each emission state to different condition in the magnetosphere in such a way that a plasma-filled magnetosphere is related to the emission-on state and a sudden depletion of charge particles is responsible for the decrease in the spin-down rate during the emission-off state (Contopoulos 2005; Kramer et al. 2006; Beskin & Nokhrina 2007; Li et al. 2012). Kramer et al. (2006) and Beskin & Nokhrina (2007) suggest that energy loss of PSR B1931 + 24 is due only to the rotating dipole during the emission-off state where there is no current flowing, whereas current loss is responsible for the energy loss during the emission-on state. In this scenario, pulses are undetectable due to decrease in the plasma density in the emitting region to a magnetospheric state so that the pulse intensity of the pulsar decreases when it jumps to that state and resumes to its normal intensity when it returns to its initial state. While the approach has revealed unprecedented information about the pulsars, it implicitly treats the two emission states as independent entities with presumed vacuum condition for the emission-off state. Our investigation is based on a magnetospheric model that synthesizes the corotating magnetosphere model and the vacuum model (Melrose & Yuen 2014, 2016), in which multiple quasi-stable magnetospheric states are described by the parameter y such that the charge filling fraction of a state corresponds to a y value and a change in the charge filling fraction corresponds to a change of state which is signified by a change in y. The presence of plasma in the corotating magnetosphere model implies a global current flow that closes at the stellar surface within the polar cap region. By considering a simple framework for the spin-down torque as induced by the component of the global current that crosses the magnetic fields at the stellar surface, we show that discrete changes in the spin-down rates may be attributed to changes in the pulsar torque caused by variations in the global current flow when jumping takes place between different magnetospheric states in the pulsar magnetosphere. Using this idealized model, characteristics of intermittent pulsars are studied by jointly considering the two emission states based on the parameters y and \u03b1, the obliquity angle between the magnetic and rotation axes, through the ratio of the respective spin-down rate whose changes are described by the ratio of the corresponding spin-down torque in the two states. The yet unidentified links between intermittent and ordinary pulsars imply that, in addition to the parameters already under detailed examination in the literature, there may exist unknown factors that are distinctly related to intermittency and likely to be pulsar-specific. An obvious example is the large differences in the emission-on and -off cycles which range from days (Kramer et al. 2006) to years (Camilo et al. 2012). Since the correlation between these parameters and intermittency is still unclear, before including such factors, it is useful to explore the implications of the idealized model.","Citation Text":["Melrose & Yuen 2014"],"Functions Text":["Our investigation is based on a magnetospheric model that synthesizes the corotating magnetosphere model and the vacuum model",", in which multiple quasi-stable magnetospheric states are described by the parameter y such that the charge filling fraction of a state corresponds to a y value and a change in the charge filling fraction corresponds to a change of state which is signified by a change in y."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1489,1508]],"Functions Start End":[[1362,1487],[1515,1790]]}
{"Identifier":"2018MNRAS.474.2277D__Oliveira,_Dottori_&_Bica_1998_Instance_2","Paragraph":"There are three possible explanations for the origin of these systems: (1) they formed from the fragmentation of the same molecular cloud (Elmegreen & Elmegreen 1983), (2) they were generated in distinct molecular clouds and then became bound systems after a close encounter leading to a tidal capture (Vallenari, Bettoni & Chiosi 1998; Leon, Bergond & Vallenari 1999), or (3) they are the result of division of a single star-forming region (Goodwin & Whitworth 2004; Arnold et al. 2017). Their subsequent evolution may also have different outcomes. Dynamical models and N-body simulations (see, e.g., Barnes & Hut 1986; de Oliveira, Dottori & Bica 1998, and references therein) have shown that, depending on the initial conditions, a bound pair of clusters may either become unbound, because of significant mass-loss in the early phases of stellar evolution, or merge into a single and more massive cluster on a short time-scale (\u224860\u2009Myr) due to loss of angular momentum from escaping stars (see Portegies Zwart & Rusli 2007). The final product of a merger may be characterized by a variable degree of kinematic and morphologic complexity, mostly depending on the values of the impact parameter of the pre-merger binary system (de Oliveira, Bica & Dottori 2000; Priyatikanto et al. 2016). In some cases, the stellar system resulting from the merger event may show significant internal rotation (in fact, for many years this has been the preferred dynamical route to form rotating star clusters; see Sugimoto & Makino 1989; Makino, Akiyama & Sugimoto 1991; Okumura, Ebisuzaki & Makino 1991; de Oliveira, Dottori & Bica 1998). Merger of cluster pairs has been sometimes invoked to interpret the properties of particularly massive and dynamically complex clusters (e.g. see the study of \u03c9 Centauri by Lee et al. 1999, G1 by Baumgardt et al. 2003 and NGC\u2009 2419 by Br\u00fcns & Kroupa 2011), and, more in general, as an avenue to form clusters with multiple populations with different chemical abundances both in terms of iron and light elements (e.g. van den Bergh 1996; Catelan 1997; Amaro-Seoane et al. 2013; Gavagnin, Mapelli & Lake 2016; Hong et al. 2017).","Citation Text":["de Oliveira, Dottori & Bica 1998"],"Functions Text":["In some cases, the stellar system resulting from the merger event may show significant internal rotation (in fact, for many years this has been the preferred dynamical route to form rotating star clusters; see"],"Functions Label":["Background"],"Citation Start End":[[1591,1623]],"Functions Start End":[[1290,1499]]}
{"Identifier":"2016ApJ...817...12P__Chamandy_et_al._2014_Instance_1","Paragraph":"Large-scale magnetic fields with strength of the order of 1\u201310 \u03bcG have been observed in disk galaxies (e.g., Beck et al. 1996; Fletcher 2010; Beck 2012; Beck & Wielebinski 2013; Van Eck et al. 2015). The origin of these fields can be explained through mean-field dynamo theory (Ruzmaikin et al. 1988; Beck et al. 1996; Brandenburg & Subramanian 2005a; Kulsrud & Zweibel 2008). The conservation of magnetic helicity is one of the key constraints in these models, and also leads to the suppression of the \u03b1-effect. The operation of the mean-field dynamo automatically leads to the growth of magnetic helicity of opposite signs between the large-scale and small-scale magnetic fields (Pouquet et al. 1976; Gruzinov & Diamond 1994; Blackman & Field 2002). To avoid catastrophic suppression of the dynamo action (\u03b1-quenching), the magnetic helicity due to the small-scale magnetic field should be removed from the system (Blackman & Field 2000, 2001; Kleeorin et al. 2000). Mechanisms suggested to produce these small-scale magnetic helicity fluxes are: advection of magnetic fields by an outflow from the disk through the galactic fountain or wind (Shukurov et al. 2006; Sur et al. 2007; Chamandy et al. 2014), magnetic helicity flux from anisotropy of the turbulence produced by differential rotation (Vishniac & Cho 2001; Subramanian & Brandenburg 2004, 2006; Sur et al. 2007; Vishniac & Shapovalov 2014), and through diffusive flux (Kleeorin et al. 2000, 2002; Brandenburg et al. 2009; Mitra et al. 2010; Chamandy et al. 2014). The outflow of magnetic helicity from the disk through dynamo operation leads to the formation of a corona (Blackman & Field 2000). According to Taylor's hypothesis, an infinitely conducting corona would resistively relax to force-free field configurations under the constraint of global magnetic helicity conservation (Woltjer 1960; Taylor 1974; Finn & Antonsen 1983; Berger & Field 1984; Mangalam & Krishan 2000). In this paper, we include advective and diffusive fluxes in a simple semi-analytic model of a galactic dynamo that transfers magnetic helicity outside the disk and consequently builds up a force-free corona in course of time. We first solve the time-dependent dynamo equations by expressing them as separable in variables r and z. The radial part of the dynamo equation is solved using an eigenvector expansion constructed using the steady-state solutions of the dynamo equation. The eigenvalues of the z part of the solution are obtained by solving a fourth-order algebraic equation, which primarily depends upon the turbulence parameters and the magnetic helicity fluxes. Once the dynamo solutions are written out as parametric functions of these parameters, the evolution of the mean magnetic field is computed numerically by simultaneously solving the dynamical equations for \u03b1-quenching and the growth of large-scale coronal magnetic helicity. Since the large-scale magnetic field lines cross the boundary between the galactic disk and the corona, the magnetic helicity of the large-scale magnetic field in the disk volume is not well defined. Hence we use the concept of gauge-invariant relative helicity (Finn & Antonsen 1983; Berger & Field 1984; Berger 1985) to estimate the large-scale magnetic helicity in the disk and the corona. Here the gauge-invariant relative helicity for the cylindrical geometry is calculated using the prescription given in Low (2006, 2011). We then investigate the dependence of the saturated mean magnetic field strength and its geometry on the magnetic helicity fluxes within the disk and the corresponding evolution of the force-free field in the corona.","Citation Text":["Chamandy et al. 2014"],"Functions Text":["Mechanisms suggested to produce these small-scale magnetic helicity fluxes are: advection of magnetic fields by an outflow from the disk through the galactic fountain or wind"],"Functions Label":["Background"],"Citation Start End":[[1184,1204]],"Functions Start End":[[969,1143]]}
{"Identifier":"2021ApJ...921...18K__Kushwaha_et_al._2018a_Instance_3","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"],"Functions Text":["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"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1883,1904]],"Functions Start End":[[1611,1881]]}
{"Identifier":"2018MNRAS.478.2541F__Smith_&_Tombleson_2015_Instance_2","Paragraph":"Utilizing the full range of peak absolute magnitudes observed in LBV stars [M \u2243 \u2009(\u221213)\u2013(\u22129)\u2009mag; e.g. Smith et al. 2011b] provides a range of peak apparent magnitudes of m \u2243 \u200916.5\u201320.5\u2009mag, for a distance of D \u2243 8\u2009Mpc (Table 1). Hence, if the transient source (m \u2248 21.0\u2009mag; Section 2.1.2) is an LBV star, it must have been observed a short time after its peak (Fig. 2; left-hand column; rows 4\u20135); quiescent LBV stars can have absolute magnitudes as low as M \u2243 \u22126\u2009mag (e.g. Smith et al. 2011b), or m \u2243 23.5\u2009mag (for a distance of D \u2243 8\u2009Mpc; Table 1). An LBV star provides an adequate explanation for the transient time-scale, the isolation (e.g. Smith & Tombleson 2015; Smith 2016), the lack of a host H\u2009ii region (Fig. 2 and 3; see below, however), and the transient\/main source offset (d \u2248 0.3\u2009kpc, for D \u2243 8\u2009Mpc; Table 1; see below, however). In addition, (net) fading of the LBV peak event by \u0394m \u2272 3\u2009mag (over a period of approximately 40 yr) may possibly produce an apparent brightness centroid shift\/morphology variation in the main source (Sections 2.1.2 and 2.1.3). However, this scenario has its caveats. First, it is difficult to interpret the flux variability of the main source (Sections 2.1.4 and 2.3.1) within such a context. Secondly, although a fraction of LBV stars are isolated (e.g. Smith & Tombleson 2015; Smith 2016), in the only XMP with a documented LBV star, the LBV star appears embedded within an H\u2009ii region (DDO 68; Pustilnik et al. 2017); no clear H\u2009ii region at the location of the transient source is discernible in the HST images (Fig. 3; top), although it should be detectable given the typical lifetime of an H\u2009ii region (few Myr; e.g. Alvarez, Bromm & Shapiro 2006). Thirdly, a (net) magnitude variation of \u0394m \u22723\u2009mag of the transient over a period of approximately 60 yr would result in a quiescent LBV source (m \u2272 23.5\u2009mag) that should have been detectable in the HST images (Fig. 3; top). Lastly, as LBV stars brighten, they become redder (e.g. Sterken 2003), which appears to contradict the POSS data (Table 2; Fig. 2; rows 2\u20135). Consequently, as the LBV scenario explains some of the observables challenged by other scenarios (e.g. transient timeline), it remains a contender for the observed phenomenon.","Citation Text":["Smith & Tombleson 2015"],"Functions Text":["Secondly, although a fraction of LBV stars are isolated (e.g."],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1303,1325]],"Functions Start End":[[1241,1302]]}
{"Identifier":"2018AandA...619A..13V__Saviane_et_al._2012_Instance_4","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)"],"Functions Text":["For internal consistency, all EWs in this work have been adjusted to the measurement scale of","by using this relation."],"Functions Label":["Uses","Uses"],"Citation Start End":[[2136,2157]],"Functions Start End":[[2042,2135],[2158,2181]]}
{"Identifier":"2020ApJ...895...82V__Fryer_et_al._2018_Instance_1","Paragraph":"The shock is then revived by adding an energy injection following the parameterized method of Fryer et al. (2018). In this model, roughly \n\n\n\n\n\n was deposited into the inner \n\n\n\n\n\n in the first \n\n\n\n\n\n. Some of this energy is lost through neutrino emission and the total explosion energy at late times for this model is \n\n\n\n\n\n. This explosion is then mapped into our three-dimensional calculations, using one million SPH particles. The mapping took place when the supernova shock had moved out of the iron core and propagated into the Si\u2013S rich shell at \n\n\n\n\n\n. We note that our 1D methods employed for modeling the collapse, core bounce, and initial explosion do not capture the full physics of the central engine (for a discussion, see Fryer et al. 2018), and this is a source of uncertainty in our yield calculations. The details of the engine change the shock trajectories, and neutrino chemistry can change Ye values (Saez et al. 2018; Fujimoto & Nagakura 2019). The nature of the shock affects mostly the yields after the shock falls below NSE (before it falls out of NSE, the yields are set by the equilibrium values, not the time-dependent evolution). Our model captures one instance of the range of asymmetric trajectories, and it should be noted that no model at this time is sufficiently accurate to dictate exactly the properties of the asymmetries (Janka et al. 2016). In addition, any model that does not include convection-driven asymmetries from the progenitor star cannot properly capture the asymmetries (Arnett et al. 2015). The 3D explosion model used here also displays stochastic asymmetries, implying that any manner of convective asymmetry could generate similar results. If this behavior is universal, it could have important implications. These points taken together indicate that nucleosynthetic patterns arising from convection-like behavior are robust, regardless of the driver. As discussed below, this increases the utility of NSE nucleosynthesis, particularly of \n\n\n\n\n\n and \n\n\n\n\n\n, as diagnostics of the conditions in the progenitor star.","Citation Text":["Fryer et al. (2018)"],"Functions Text":["The shock is then revived by adding an energy injection following the parameterized method of"],"Functions Label":["Uses"],"Citation Start End":[[94,113]],"Functions Start End":[[0,93]]}
{"Identifier":"2017ApJ...844...14L__Martell_et_al._2008_Instance_1","Paragraph":"Figure 2 shows the measured spectral indices of stars as functions of K magnitude, obtained from the 2MASS catalog. The CN, HK\u2032, and CH indices increase with decreasing magnitude because the brighter RGB stars have lower temperatures and the strengths of these molecular bands generally increase with decreasing temperature. Therefore, the chemical abundances of stars are compared on the \u03b4-index versus magnitude diagrams. It is important to note that the observed stars show a large spread in \u03b4-index that is at least several times larger than the measurement error. The standard deviations for all sample stars are 0.23 for CN, 0.07 for HK\u2032, and 0.07 for CH. In particular, the CN index distribution shows the largest spread. Note that a bimodality or a large spread in CN distribution is generally observed in most GCs (Norris et al. 1981; Norris 1987; Briley et al. 1992; Harbeck et al. 2003; Kayser et al. 2008; Martell et al. 2008).2\n\n2\nAlthough the evolutionary mixing effect can also contribute to the large spread in CN index distribution among bright RGB stars (Sweigart & Mengel 1979), this effect alone cannot explain a discrete distribution and a wide spread in the unevolved stars (see, e.g., Kayser et al. 2008).\n Therefore, we have divided subpopulations of RGB stars in NGC 5286 on the histogram of the \u03b4CN index (see Figure 3). It is clear from this histogram that RGB stars are divided into three subpopulations: CN-weak (\u03b4CN  \u22120.2; blue circles), CN-intermediate (\u22120.2 \u2264 \u03b4CN  0.1; green circles), and CN-strong (0.1 \u2264 \u03b4CN; red circles). The distribution of CN index into three or more subpopulations is similar to that reported in NGC 1851 (Campbell et al. 2012; Lim et al. 2015; Simpson et al. 2017). This is also consistent with the recent results from population models and spectroscopic observations that show that most GCs host three or more subpopulations (see, e.g., Jang et al. 2014; Carretta 2015). The presence of multiple populations is also observed from recent photometry using UV filters, which are mainly sensitive to N abundance (Milone et al. 2015; Piotto et al. 2015). In this regard, further observations are required to check that the trimodal CN distribution, observed in NGC 1851 and NGC 5286, is a ubiquitous feature in other GCs as well.","Citation Text":["Martell et al. 2008"],"Functions Text":["In particular, the CN index distribution shows the largest spread. Note that a bimodality or a large spread in CN distribution is generally observed in most GCs"],"Functions Label":["Similarities"],"Citation Start End":[[918,937]],"Functions Start End":[[662,822]]}
{"Identifier":"2019MNRAS.487..475C__Ward-Thompson_et_al._2009_Instance_1","Paragraph":"We attempted to find a correlation between the mean magnetic field and the outflow and minor axis of the cloud CB 17. Relative orientations between various quantities of CB 17 are presented in the first row of Table 6, along with a comparative study of the same for some dark clouds. The first column gives the cloud ID and columns 2\u20136 give the position angles of the mean magnetic field at the envelope ($\\lt \\theta ^{\\rm env}_B\\gt $), mean magnetic field at the core ($\\lt \\theta ^{\\rm core}_B\\gt $), outflow axis (\u03b8out), minor axis (\u03b8min) of the core of the cloud and Galactic plane (\u03b8GP), respectively. $\\lt \\theta ^{\\rm env}_B\\gt $ of CB 17 is found to be almost aligned along the Galactic plane over that region of the sky (column 7), which indicates the dominance of the Galactic magnetic field over the envelope magnetic field of the cloud and thus we cannot infer much about the magnetic field structure from the optical study. A similar feature has also been observed in the cases of CB 34 (Das et al. 2016), L328, L673-7 (Soam et al. 2015), CB 26 (Halder et al. 2019), CB 3 and CB 246 (Ward-Thompson et al. 2009). However, $\\lt \\theta ^{\\rm core}_B\\gt $ of CB 17 (obtained by submm polarimetry) turned out to be perpendicular to $\\lt \\theta ^{\\rm env}_B\\gt $ (column 8); a similar phenomenon has been observed in the case of CB 34-C1 (Das et al. 2016), IRAM 04191 (Soam et al. 2015) and CB 54 (Wolf et al. 2003). Since, in the case of CB 17, $\\lt \\theta ^{\\rm env}_B\\gt $ is along the Galactic plane orientation, this implies that only $\\lt \\theta ^{\\rm core}_B\\gt $ (denser region) is linked with the ongoing physical phenomena in the cloud. $\\lt \\theta ^{\\rm core}_B\\gt $ is oriented perpendicular to \u03b8GP as well (column 9) and a similar orientation has been observed in the case of CB 34-C1 (Das et al. 2016), IRAM 04191 (Soam et al. 2015), CB 230 and CB 244 (Wolf et al. 2003) as well. Moreover, $\\lt \\theta ^{\\rm core}_B\\gt $ is found to be almost aligned along the minor axis of the core of the cloud; the angular offset is nearly 5.9\u00b0 (column 10). The alignment of $\\lt \\theta ^{\\rm core}_B\\gt $ with the minor axis of the cloud fits the magnetically regulated star formation model, in which the magnetic field should lie along the minor axis of the cloud (Mouschovias & Morton 1991; Li 1998), and the same feature has also been observed for the clouds CB 34-C1 (Das et al. 2016) and IRAM 04191 (Soam et al. 2015). The angular offset between $\\lt \\theta ^{\\rm core}_B\\gt $ and the outflow axis is found to be 80.9\u00b0 (column 11), that is, the core-scale magnetic field is oriented almost perpendicular to the outflow direction and a similar phenomenon has also been observed in the case of CB 34 (Das et al. 2016), CB 68 (Bertrang et al. 2014), B335, CB 230, CB 244 (Wolf et al. 2003) and CB 3 (Ward-Thompson et al. 2009). The angular offset between \u03b8out and \u03b8min is found to be 75\u00b0 and the same feature has been observed for CB 34-C1 (Das et al. 2016).","Citation Text":["Ward-Thompson et al. 2009"],"Functions Text":["A similar feature has also been observed","CB 3 and CB 246"],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[1097,1122]],"Functions Start End":[[937,977],[1080,1095]]}
{"Identifier":"2018MNRAS.476.5417S__McCray_&_Kafatos_1987_Instance_1","Paragraph":"According to our interpretation, the observed high velocity dispersion region in zone \u2018F\u2019 associated with the zone totally empty of gas between the two opposite velocity clouds (that the 3Dbarolo model would instead predict to be filled) resembles an expanding superbubble. Superbubbles are known to be associated with very massive stars (OB-associations): strong stellar winds and subsequent SN explosions from those stars inject energy and mass into the ambient ISM, creating shock fronts that sweep-up the ISM (e.g. Castor, McCray & Weaver 1975). As the SN explosions start occurring within the cavity formed by the stellar wind bubbles, super-bubbles are created (e.g. McCray & Kafatos 1987), which may eventually reach kiloparsec sizes. These explosions never form a visible SNR, but instead expend their energy in the hot interior as sound waves. Both stellar winds and stellar explosions thus power the expansion of the superbubble in the ISM. The interstellar gas swept up by superbubbles generally cools, forming a dense shell around the cavity (observed in Hi and H\u03b1). The bubble interior contains hot (>106\u2009K), rarefied material, usually associated with extended diffuse X-ray emission (thus appearing as an empty cavity in cold and molecular gas). In this interpretation, the two blobs of gas above and below the observed empty cavity in region \u2018F\u2019 (at +1.5\u2009arcsec offset, with velocities of +70 and \u221270\u2009km\u2009s\u22121, corresponding to the two peaks in Fig. 6), can be easily explained as gas surrounding the \u2018superbubble\u2019 being thrown away in different directions by the expanding spherical shell (see Kamphuis, Sancisi & van der Hulst 1991; Boomsma et al. 2008, for detected velocities up to 100\u2009km\u2009s\u22121 associated with \u2018holes\u2019 in Hi). In Fig. 13, we show different CO(2\u20131) velocity channels, where the \u2018superbubble\u2019, surrounded by material moving in opposite directions (e.g. peaks at \u221270 and 70\u2009km\u2009s\u22121, not present at the systemic velocities), shows up as an empty area (encircled in the plot), with a clear shock front at its left-hand side, peaking around the systemic velocities (\u22430\u2009km\u2009s\u22121).","Citation Text":["McCray & Kafatos 1987"],"Functions Text":["As the SN explosions start occurring within the cavity formed by the stellar wind bubbles, super-bubbles are created (e.g.","), which may eventually reach kiloparsec sizes."],"Functions Label":["Background","Background"],"Citation Start End":[[673,694]],"Functions Start End":[[550,672],[694,741]]}
{"Identifier":"2018AandA...609A..13K__Mucciarelli_et_al._(2017)_Instance_2","Paragraph":"Gaia 1 is a star cluster that was recently discovered by Koposov et al. (2017) in the first Gaia data release (Gaia Collaboration 2016), alongside with another system of lower mass. Its observation and previous detections were seriously hampered by the nearby bright star Sirius, which emphasized the impressive discovery power of the Gaia mission. This object was first characterized as an intermediate-age (6.3 Gyr) and moderately metal-rich (\u22120.7 dex) system, based on isochrone fits to a comprehensive combination of Gaia, 2MASS (Cutri et al. 2003), WISE (Wright et al. 2010), and Pan-STARRS1 (Chambers et al. 2016) photometry. Hence, this object was characterized by Koposov et al. (2017) as a star cluster, most likely of the globular confession. Further investigation of Gaia 1 found a metallicity higher by more than 0.5 dex, which challenged the previous age measurement and rather characterized it as a young (3 Gyr), metal-rich (\u22120.1 dex) object, possibly of extragalactic origin given its orbit that leads it up to ~1.7 kpc above the disk (Simpson et al. 2017). Subsequently, Mucciarelli et al. (2017) measured chemical abundances of six stars in Gaia 1, suggesting an equally high metallicity, but based on their abundance study, the suggestion of an extragalactic origin was revoked. While a more metal-rich nature found by the latter authors conformed with the results by Simpson et al. (2017), the evolutionary diagrams of both studies are very dissimilar and could not be explained by one simple isochrone fit. In particular, it was noted that \u201cthe Simpson et al. (2017) stars do not define a red giant branch in the theoretical plane, suggesting that their parameters are not correct\u201d (Fig. 1 of Mucciarelli et al. 2017). Such an inconsistency clearly emphasizes that a clear-cut chemical abundance scale is inevitable for fully characterising Gaia 1, and to further allow for tailored age determinations, even more so in the light of the seemingly well-determined orbital characteristics, Thus, this work focuses on a detailed chemical abundance analysis of four red giant members of Gaia 1, based on high-resolution spectroscopy, which we complement with an investigation of the orbital properties of this transition object. Combined with the red clump sample of Mucciarelli et al. (2017) and reaching down to the subgiant level (Simpson et al. 2017), stars in different evolutionary states in Gaia 1 are progressively being sampled. ","Citation Text":["Mucciarelli et al. 2017"],"Functions Text":["In particular, it was noted that \u201cthe Simpson et al. (2017) stars do not define a red giant branch in the theoretical plane, suggesting that their parameters are not correct\u201d (Fig. 1 of"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1714,1737]],"Functions Start End":[[1528,1713]]}
{"Identifier":"2021MNRAS.503.2108P__Andresen_et_al._2017_Instance_1","Paragraph":"CCSNe are also of interest for GW astronomy as targets in their own right. As the sensitivity of GW detectors increases, they will begin to detect not only binary mergers but also other lower amplitude sources of GWs such as CCSNe. Accurate knowledge of the GW emission from CCSNe will be essential for detection and parameter estimation. The GW signal from rotational core bounce has already been well covered in the literature (e.g. Dimmelmeier et al. 2008; Abdikamalov et al. 2014; Fuller et al. 2015; Richers et al. 2017). In the non-rotating case, the GW emission from the post-bounce phase has been studied using self-consistent 3D simulations by many groups (Kuroda, Kotake & Takiwaki 2016; Andresen et al. 2017, 2019; Kuroda et al. 2017, 2018; Powell & M\u00fcller 2019, 2020; Radice et al. 2019; Andresen, Glas & Janka 2020; Mezzacappa et al. 2020; Pan et al. 2020). The structure of the GW emission has shown common features in different simulations from recent years. The dominant emission feature in the GW emission is due to the quadrupolar surface f\/g mode 1 of the proto-neutron star (PNS), which produces GW frequencies rising in time from a few hundred Hz up to a few kHz (M\u00fcller, Janka & Wongwathanarat 2012; Sotani et al. 2017; Kuroda et al. 2018; Morozova et al. 2018; Torres-Forn\u00e9 et al. 2018, 2019). In addition, some models (Kuroda et al. 2016, 2017; Andresen et al. 2017; Mezzacappa et al. 2020; Powell & M\u00fcller 2020) exhibit low-frequency GW emission due to the standing accretion shock instability (SASI; Blondin, Mezzacappa & DeMarino 2003; Blondin & Mezzacappa 2006; Foglizzo et al. 2007). In rapidly rotating models, very strong GW emission can also occur during the post-bounce phase due to a corotation instability (Takiwaki & Kotake 2018). The emerging understanding of the GW emission features has led to the formulation of universal relations for the GW emission (Torres-Forn\u00e9 et al. 2019) and paved the way for phenomenological modelling for CCSN signals (Astone et al. 2018). Further work is still needed, however, to extend these models to fully explore CCSN GW signals from across the progenitor parameter space. The majority of 3D simulations that include GW emission are for progenitor stars below $30\\, \\mathrm{M}_{\\odot }$. In this paper, we perform simulations of high-mass Population III (Pop-III) stars in the pulsational pair instability regime to expand the parameter space coverage of 3D simulations and to provide further insights into the massive and very massive star remnant BH population.","Citation Text":["Andresen et al. 2017"],"Functions Text":["In the non-rotating case, the GW emission from the post-bounce phase has been studied using self-consistent 3D simulations by many groups"],"Functions Label":["Background"],"Citation Start End":[[698,718]],"Functions Start End":[[527,664]]}
{"Identifier":"2017MNRAS.469.3270V__Chakrabarti,_Jin_&_Arnett_1987_Instance_1","Paragraph":"Observations show that the core temperatures of powerful AGN jets are estimated to be quite high (Moellenbrock et al. 1996). So the jets are hot to start with in this paper too. The advective disc model, as in most disc models, comes with a variety of inner disc temperatures. Simulations of advective discs for high viscosity parameter produced T \u2273 1012K in the PSD (Lee et al. 2016). Moreover, in the presence of viscous dissipation in curved space\u2013time, the Bernoulli parameter (\u2212hut) may increase by more than 20\u2009per\u2009cent of its value at large distance and produce very high temperatures in the PSD (Chattopadhyay & Kumar 2016). For highly rotating BHs too, the temperatures of the inner disc easily approach 1012 K. It must also be remembered that inner regions of the accretion disc can be heated by the Ohmic dissipation, reconnection, turbulence heating or MHD wave dissipation may heat up the inner disc or the base of the jet (Pudritz 2003). High temperatures in the accretion disc can induce exothermic nucleosynthesis too (Chakrabarti, Jin & Arnett 1987; Hu & Peng 2008). All these processes taken together in an advective disc will produce very hot jet base. We do not specify the exact processes that will produce very hot jet base, but would like to emphasize that it is quite possible to achieve so. One may also wonder that if the jets are indeed launched from the disc, and how justified is it to consider non-rotating jets. Phenomenologically speaking, if jets have a lot of rotation then it would not flow around the axis of symmetry and, therefore, either it has to be launched with less angular momentum or it has to lose most of the angular momentum with which it is launched. It has been shown that viscous transport removes significant angular momentum of the collimated outflow close to the axis (Lee et al. 2016). Since the jet is launched with low angular momentum and it is further removed by viscosity or by the presence of magnetic field; therefore, the assumption of non-rotating hot jet is quite feasible. Incidentally, similar to this study, there are many theoretical studies of jets that have been undertaken under similar assumptions of non-rotating, hot jets at the base (Fukue 1987b; Falcke 1996; Memola et al. 2002).","Citation Text":["Chakrabarti, Jin & Arnett 1987"],"Functions Text":["High temperatures in the accretion disc can induce exothermic nucleosynthesis too","All these processes taken together in an advective disc will produce very hot jet base."],"Functions Label":["Background","Background"],"Citation Start End":[[1035,1065]],"Functions Start End":[[952,1033],[1084,1171]]}
{"Identifier":"2022MNRAS.509..314F__Dikpati_et_al._2020_Instance_1","Paragraph":"Magneto-Rossby waves arise due to the inhomogeneity of the Coriolis force depending on latitude on a sphere in rotating astrophysical plasma by analogy with a neutral fluid (Petviashvili & Pokhotelov 1992; Onishchenko et al. 2004; Vallis 2006; Onishchenko, Pokhotelov & Astafieva 2008; Zeitlin 2018). It should be noted that owing to the presence of Lorentz force full vorticity is no longer conserved and in the case of a strong magnetic field magneto-Rossby waves tend to Alfv\u00e9n wave solutions almost on all wavenumber range (Zaqarashvili et al. 2021). Magneto-Rossby waves are basic mechanism in variability of various objects in plasma astrophysics. Magneto-Rossby waves determine the large-scale dynamics of the Sun and stars (Hughes et al. 2007; Zaqarashvili et al. 2007, 2011; Dikpati et al. 2020; Mandal & Hanasoge 2020; Raphaldini et al. 2020), dynamics of magnetoactive atmospheres of exoplanets captured by tides from the host star (Cho 2008), flows in accretion discs of neutron stars (Inogamov & Sunyaev 2010). Despite the difficulty of observing Rossby waves in astrophysical plasma, they have recently been detected on the Sun (McIntosh et al. 2017; Zaqarashvili & Gurgenashvili 2018; Loeptien et al. 2018; Liang et al. 2019). We also note a number of important studies on the effect of Magneto-Rossby waves on solar seasons (Lou 2000; Dikpati et al. 2017, 2018) and space weather (Dikpati & McIntosh 2020; Dikpati, McIntosh & Wing 2021). In addition, magneto-Rossby waves play a key role in the occurrence of zonal flows in two-dimensional magnetohydrodynamic turbulence (Tobias, Diamond & Hughes 2007; Zinyakov & Petrosyan 2018, 2020). The weakly non-linear theory of magneto-Rossby waves in the quasi-two-dimensional MHD single-layer model has been developed in Klimachkov & Petrosyan (2017a) and in the presence of large-scale compressibility in Klimachkov & Petrosyan (2018). In Fedotova et al. (2020), MHD shallow water approximation has been extended to the case of external vertical magnetic field and stratified flows. Linear and weakly non-linear theory of magneto-Rossby waves in a rotating stratified plasma in the quasi-two-dimensional MHD approximation in two-layer model has been developed. In Klimachkov & Petrosyan (2017a, 2018) and Fedotova et al. (2020), the phase matching conditions have been investigated and non-linear interactions of three magneto-Rossby waves have been found. In Raphaldini & Raupp (2015), Raphaldini et al. (2019, 2020) influence of non-linear dynamics of magneto-Rossby waves on cyclic nature of solar magnetic activity is studied. Non-linear interactions of magneto-Rossby waves are associated with the duration of the solar cycle. In Raphaldini et al. (2020) is shown that irregular transitions in wave amplitudes resemble the observed time series of solar activity. In Klimachkov & Petrosyan (2017a, 2018), Fedotova et al. (2020) magneto-Rossby waves have been investigated in the \u03b2-plane approximation for the Coriolis force. The \u03b2-plane approximation describes rotating spherical plasma flows in a local Cartesian coordinate system. In this case, the Coriolis parameter changes little with small changes in latitude and expands in a series up to the first order in latitude.","Citation Text":["Dikpati et al. 2020"],"Functions Text":["Magneto-Rossby waves are basic mechanism in variability of various objects in plasma astrophysics. Magneto-Rossby waves determine the large-scale dynamics of the Sun and stars"],"Functions Label":["Background"],"Citation Start End":[[784,803]],"Functions Start End":[[555,730]]}
{"Identifier":"2021AandA...648A...5M__Windhorst_et_al._(1990)_Instance_1","Paragraph":"Another important consistency check regards the angular size distribution of the sources. Figure 6 shows the cumulative size distributions of the final catalogs combined together, in four flux density bins (yellow solid lines). Such distributions can be considered reliable only down to a flux-dependent minimum intrinsic size (see vertical gray lines), below which most of the sources cannot be reliably deconvolved and they are conventionally assigned \u0398 = 0. The observed distributions are compared with various realizations of the cumulative distribution function described by Eq. (6), obtained by varying either the function exponent q (left and right columns respectively) or the assumed median size \u2013 flux relations (see various black lines).The original function proposed by Windhorst et al. (1990) (Eq. (6) with q = 0.62, see left column) does provide a good approximation of the observed distributions, when assuming the original \u0398med \u2212 S relation described by Eq. (7), only at flux densities S150 MHz\u227310 mJy (see long-dashed lines). This is perhaps not surprising considering that this relation was calibrated at 1.4 GHz down to a few mJy fluxes. At the lowest flux densities (S150 MHz\u22721 mJy) we need to assume a steepening of the parameter m (see Eq. (8)), to get a good match with observations (dotted line in the top left panel). This is consistent with what proposed for higher frequency deep surveys (as discussed earlier in this section). At intermediate fluxes (S150 MHz ~ 1\u221210) mJy, on the other hand, none of the discussed median size \u2013 flux relations can reproduce the observed size distribution (see second-row panel on the left). It is interesting to note, however, that if we assume a steeper exponent for the distribution function described by Eq. (7) (i.e., q = 0.80), we get a very good match with observations at all fluxes, when assuming a flux-dependent scaling factor (k = k(S); see Eq. (9)) for the Windhorst et al. (1990) median size \u2013 flux relation (black solid lines on the right). The median sizes derived from the T-RECS simulated catalogs (Bonaldi et al. 2019) also provide good results for q = 0.80 (dot-dashed lines on the right), except again at intermediate fluxes (S150 MHz ~ 1\u221210), where they show strong discrepancies with observations also in Fig. 5. This seems to indicate that the number density of extended radio galaxies in this flux density range is over-estimated in the T-RECS simulated catalogs.","Citation Text":["Windhorst et al. (1990)"],"Functions Text":["The original function proposed by","(Eq. (6) with q = 0.62, see left column) does provide a good approximation of the observed distributions, when assuming the original \u0398med \u2212 S relation described by Eq. (7), only at flux densities S150 MHz\u227310 mJy (see long-dashed lines). This is perhaps not surprising considering that this relation was calibrated at 1.4 GHz down to a few mJy fluxes. At the lowest flux densities (S150 MHz\u22721 mJy) we need to assume a steepening of the parameter m (see Eq. (8)), to get a good match with observations (dotted line in the top left panel)."],"Functions Label":["Differences","Differences"],"Citation Start End":[[782,805]],"Functions Start End":[[748,781],[806,1342]]}
{"Identifier":"2021MNRAS.507.1229P__Lyman_et_al._2016_Instance_1","Paragraph":"In this work, we present well-calibrated optical photometric (\u22120.2 to +413 d), polarimetric (\u22122 to +31 d) and optical (\u22125 to +391 d), NIR (\u22125 to +22 d) spectroscopic studies of SN 2012au, based on data obtained using many observational facilities around the globe. Analysis based on our photometric observations suggests that SN 2012au appears to be one of the most luminous SNe Ib (MB, peak = \u221218.06 \u00b1 0.12 mag), though fainter than the threshold limit of SLSNe I (M$_\\mathrm{ g}\\, \\lt -$19.8 mag; Quimby et al. 2018). The MR, peak (\u223c \u201318.67\u2009\u00b1\u20090.11 mag) of SN 2012au is brighter than the average values of SNe Ib and Ic, but closer to those reported for SNe Ic-BL (Drout et al. 2011). Similarly, the peak bolometric luminosity of SN 2012au (\u223c\u2009[6.56 \u00b1 0.70] \u00d7 1042 erg s\u22121) is higher than the mean peak luminosities of SNe Ib and Ic, but still lower than those of SNe Ic-BL (Lyman et al. 2016). Using the early bolometric light curve of SN 2012au, the estimated values of Mej, Ek, MNi, and T0 are \u223c5.1 \u00b1 0.7 M\u2299, \u223c\u2009(4.8 \u00b1 0.6) \u00d7 1051 erg, \u223c0.27\u22120.30 M\u2299, and \u223c66.0 \u00b1 9.4 d, respectively. These physical parameters of SN 2012au are close to those inferred for SN 2009jf (a bright SN Ib: Sahu et al. 2011) and \u2013 on average \u2013 larger than for classical SNe Ib\/c but smaller than for some SNe Ic-BL. SN 2012au manifests larger Mej and MNi in comparison with most of the SNe IIb, Ib, and Ic, which may be the prime reason behind the luminous peak of SN 2012au, as seen in the case of SLSNe I (Nicholl et al. 2015). On the other hand, light-curve decline rates of SN 2012au (at phases \u2265+40 d) in all the optical bands are shallower than typically observed in the case of SNe Ib and slow-decaying SLSNe I, and theoretically predicated for 56Co $\\rightarrow \\, ^{56}$fe decay. As SN 2012au exhibits comparatively larger Mej, a larger optical depth resulting in a larger diffusion time-scale (for the trapped energy to cross the outer envelope) could broaden the light curve. Therefore, high trapping of gamma-rays at late phases or higher opacity of massive ejecta are among the plausible interpretations for the modest luminosity decline rate of SN 2012au in comparison with other SNe Ib (Clocchiatti & Wheeler 1997). However, smoothly distributed circumstellar media up to a larger radius could be another possibility behind the late-time shallower decay rate for SN 2012au, but an absence of the CSMI in the late-time spectra ruled out this scenario (Milisavljevic et al. 2018). The late-time bolometric light curve of SN 2012au is better constrained by $L~\\varpropto\\, t^{-2}$, a conventional magnetic dipole equation. Hence for SN 2012au the shallower decay of the late-time light curve might be a potential indicator of a central engine powering source that is accelerating the inner ejecta.","Citation Text":["Lyman et al. 2016"],"Functions Text":["Similarly, the peak bolometric luminosity of SN 2012au (\u223c\u2009[6.56 \u00b1 0.70] \u00d7 1042 erg s\u22121) is higher than the mean peak luminosities of SNe Ib and Ic, but still lower than those of SNe Ic-BL"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[875,892]],"Functions Start End":[[686,873]]}
{"Identifier":"2022AandA...665A.118F__Carlsson_et_al._2016_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":["Carlsson et al. 2016"],"Functions Text":["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.","active or enhanced regions"],"Functions Label":["Background","Background"],"Citation Start End":[[887,907]],"Functions Start End":[[579,796],[859,885]]}
{"Identifier":"2020ApJ...903L..22T__Vuitton_et_al._2007_Instance_3","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"],"Functions Text":["CH3C3N itself may form the protonated species, C4H3NH+, through reactions with the HCNH+ and C2H5+ ions producing HCN and C2H4, respectively"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1673,1692]],"Functions Start End":[[1531,1671]]}
{"Identifier":"2020AandA...636A.103H__Xu_&_Borovsky_2015_Instance_1","Paragraph":"Solar wind categorization schemes rely on different solar wind properties to identify the solar wind type and adopt mainly one of the following two approaches: (1) Composition-based schemes exploit the oxygen and carbon charge-state composition of the solar wind (Zhao & Fisk 2010; von Steiger et al. 2000). Based on the assumption that the charge states observed in the solar wind are determined in the solar corona and are not significantly changed thereafter, lower (higher) charge states are associated with source regions of the respective solar wind stream with comparatively low (high) electron temperatures. Because the charge state is not expected to change during the travel time of the solar wind, composition-based criteria are well suited to identify different solar source regions. Transport effects due to, for instance, compression regions in stream interaction regions and wave-particle interaction are not directly reflected in the charge-state composition. Stream interaction regions tend to be characterized by a (gradual or abrupt) transition in the oxygen charge-state compositions. From the charge-state information alone, stream interaction regions cannot be uniquely identified, and their solar source regions cannot be unambiguously determined without additional or context information. (2) Proton plasma properties provide an alternative to determine the solar wind type (Xu & Borovsky 2015; Camporeale et al. 2017). A clear advantage of this approach is that the required observables are available from more spacecraft. Unlike the charge-state composition, the proton speed, proton density, proton temperature, and magnetic field strength (and derived quantities such as the specific entropy and the Alfv\u00e9n speed) are all susceptible to transport effects. In particular, these quantities show radial gradients throughout the heliosphere (Marsch et al. 1982; Bale et al. 2019; Kasper et al. 2019). Thus, a solar wind categorization based on threshold values for these quantities can be expected to depend on position. In particular, the solar wind proton temperature is not a tracer of the coronal (electron) temperature. The solar wind proton temperatures show the opposite effect (von Steiger et al. 2000): high solar wind proton temperatures are observed for coronal hole wind (which originates from comparatively cool coronal regions), while low solar wind proton temperatures appear in the slow solar wind (which likely originates in hot coronal regions). The solar wind proton temperature is probably strongly influenced by transport effects, in particular, by wave-particle interactions. In addition, proton temperature, proton density, and magnetic field strength all show characteristic variations in stream interaction regions. Solar wind categorization schemes based on proton plasma properties are therefore well suited to assess the effect of solar wind evolution during its travel time. However, these transport effects can blur the tracers of the solar origin of the solar wind. For approaches based on charge-state composition and on proton plasma, the respective threshold values are usually determined heuristically and vary in the literature. Mainly as a result of their availability, solar wind electron data (e.g., Lin et al. 1995; Wilson et al. 2018) are typically not considered for solar wind classifications, although their properties, for example, the electron temperature and the electron-proton collisional age, can be expected to be informative in this context. Future improvements on solar wind classification would most likely benefit considerably from including electron data. The collisional age (or Coulomb number) has been proposed as an ordering parameter for the solar wind in Kasper et al. (2008), Tracy et al. (2016), and Maruca et al. (2013). The collisional age can be interpreted as counting the number of 90\u00b0-equivalent collisions during the travel time from the Sun to the observing spacecraft. This notion of the collisional age relies on the simplifying assumption that the solar wind parameters are constant during the solar wind travel time. Maruca et al. (2013) introduced an improvement in the computation of the collisional age that takes into account that the underlying quantities are not constant during the travel time of the solar wind.","Citation Text":["Xu & Borovsky 2015"],"Functions Text":["Proton plasma properties provide an alternative to determine the solar wind type","A clear advantage of this approach is that the required observables are available from more spacecraft."],"Functions Label":["Background","Background"],"Citation Start End":[[1399,1417]],"Functions Start End":[[1317,1397],[1444,1547]]}
{"Identifier":"2016ApJ...822....7K__Kumar_et_al._2012_Instance_1","Paragraph":"To investigate the particle precipitation or transport sites during the flare, we used HXR 25\u201350 keV and NoRH 17\/34 GHz images. We chose the Pixon algorithm (Metcalf et al. 1996) for the RHESSI image reconstruction. The Pixon method is considered to be the most accurate algorithm (Hurford et al. 2002). The integration time for each image was 20 s. We utilized NoRH 5 s cadence intensity images (R+L) at 17 and 34 GHz. Figure 4 displays the HXR 25\u201350 keV (blue) and NoRH 17 GHz (red) contours overlaid on the AIA 1600 \u212b images at \u223c23:51 UT. These images are used at the peak time of each of the bursts observed in the HXR 25\u201350 keV and microwave channels. Figure 4(a) shows a small filament rising at the flare site. We can easily identify the two legs of the filament (marked by N and S in the figure). Generally, the filaments are observed in the chromospheric images (e.g., H\u03b1 and AIA 304 \u212b). However, if the kink-unstable filament is heated during magnetic reconnection, it is often observed in the AIA 1600 \u212b channels (e.g., Kumar et al. 2012 and Kumar & Cho 2014). The locations of both the 25\u201350 keV sources and the 17 GHz source are almost at the quasi-circular ribbon. However, their centroids constructed at the 90% level of the peak intensity are not cospatial. We note that the 25\u201350 keV sources (centroid) are located close to the legs of the filament, whereas the 17 GHz source is located at the northern leg of the filament. It seems that these sources are at the footpoints of an underlying flare loop. To identify the location of the flare loop, we selected the AIA 94 \u212b hot channel image at 23:53 UT. In this image (panel b), we show the NoRH 34 GHz contours (yellow) overlaid on AIA 94 \u212b images at \u223c23:53 UT. This is done to show the coronal loops associated with the eruption of the small filament. The overlying 34 GHz source is cospatial with the small loop located above the quasi-circular ribbon. This 34 GHz emission may be the evidence of trapped nonthermal electrons in the loop. The rising filament is heated during reconnection with the ambient fields and then decays into an untwisting jet, as seen in Figures 4(b)\u2013(e). Furthermore, during \u223c23:53\u201323:59 UT, we see the footpoint HXR sources. The NoRH 17 GHz sources cover the quasi-circular ribbon (Figures 4(c)\u2013(d)) and could be the emission from the trapped electrons in the loop. This suggests that the bursts are most likely caused by the same population of nonthermal electrons with different energies.","Citation Text":["Kumar et al. 2012"],"Functions Text":["However, if the kink-unstable filament is heated during magnetic reconnection, it is often observed in the AIA 1600 \u212b channels (e.g.,"],"Functions Label":["Uses"],"Citation Start End":[[1031,1048]],"Functions Start End":[[897,1030]]}
{"Identifier":"2015MNRAS.446.1293C__Hauser_&_Dwek_2001_Instance_1","Paragraph":"The contribution to the submm background from LBGs is still poorly constrained; however, our data can finally address this question, since we now have a robust detection of the average 850\u2009\u03bcm flux density of LBGs at three epochs, at least those with ultraviolet luminosities of L1700 \u2248 1029 erg s\u22121 Hz\u22121, characteristic of galaxies in our sample (Table 2). Using our average stacked 850\u2009\u03bcm flux densities, we estimate surface brightness densities of 1700, 600, and 100 mJy deg\u22122 of LBGs at z \u223c 3, 4, and 5, respectively. By comparison, the total background at 850\u2009\u03bcm inferred from COBE-FIRAS is 3.1\u20134.4 \u00d7 104\u2009mJy\u2009deg\u22122 (Puget et al. 1996; Fixsen et al. 1998; Hauser & Dwek 2001; Lagache, Puget & Dole 2005). Summing these separate surface brightness densities together, we find that the LBGs with L1700 > 1029 erg s\u22121 Hz\u22121 in our z \u223c 3, 4, and 5 samples comprise around 6\u20138 per cent of the submm background at 850\u2009\u03bcm (where the range of values simply reflects the uncertainty in the COBE-FIRAS result). However, the true contribution from LBG-like galaxies to the submm background will come from a wider range in redshift, not just from the rather narrow redshift slices we have sampled (see Fig. 1), and from sources falling out of our samples due to incompleteness. To determine the total (corrected) contribution from LBGs over 3  z  5, we assume that LBGs have a constant comoving number density (which is a reasonable assumption since the bright end of the luminosity function for LBGs shows little evolution over 3  z  5; Bouwens et al. 2007; McLure et al. 2009; Reddy & Steidel 2009). Starting with the z \u223c 3 sample (which is our most complete subsample of LBGs), we can integrate over the comoving volume element for this redshift range, scaling the LBG submm background contribution accordingly. We find that the total contribution to the submm background from LBGs over the redshift range 3  z  5 is likely to be closer to 14\u201320 per cent. This result is consistent with Webb et al. (2003), who estimated an upper limit to the contribution to the submm background from 1  z  5 of less than 20 per cent.","Citation Text":["Hauser & Dwek 2001"],"Functions Text":["By comparison, the total background at 850\u2009\u03bcm inferred from COBE-FIRAS is 3.1\u20134.4 \u00d7 104\u2009mJy\u2009deg\u22122","Summing these separate surface brightness densities together, we find that the LBGs with L1700 > 1029 erg s\u22121 Hz\u22121 in our z \u223c 3, 4, and 5 samples comprise around 6\u20138 per cent of the submm background at 850\u2009\u03bcm (where the range of values simply reflects the uncertainty in the COBE-FIRAS result)."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[659,677]],"Functions Start End":[[521,618],[708,1002]]}
{"Identifier":"2022MNRAS.517.2383B__Strat,_Morgenstern_et_al._2009_Instance_1","Paragraph":"The UK Chemistry and Aerosol (UKCA) model (Morgenstern et al. 2009; O\u2019Connor et al. 2014; Archibald et al. 2020) is a framework that we use to describe the global atmospheric chemical composition of our simulated exoplanet. UKCA includes aerosol and gas-phase chemistry and is coupled to the UM dynamics. It uses the UM components for large-scale advection, convective transport, and boundary layer mixing of its aerosol and chemical tracers (O\u2019Connor et al. 2014; Archibald et al. 2020). UKCA contains a large number of gas-phase and heterogeneous chemical reactions, some of which we have included in our chemical network. Furthermore, the chemistry schemes in UKCA describe wet and dry deposition (Giannakopoulos et al. 1999). In this study, we use the Stratospheric (Strat, Morgenstern et al. 2009) and Stratospheric-Tropospheric (StratTrop, Archibald et al. 2020) chemistry schemes. Originally, StratTrop includes 75 chemical species that are connected by 283 reactions (Archibald et al. 2020). We used a reduced version of the UKCA chemistry schemes (Table 2) to quantify the impact of the different chemical mechanisms on the atmospheric chemistry of a tidally-locked exoplanet. First, we use a simple network that describes the Chapman mechanism of ozone formation (Chapman 1930), following Yates et al. (2020). Second, we add the reactive hydrogen (HOx) catalytic cycle, where HOx denotes the ensemble of atomic hydrogen (H), the hydroxyl radical (OH) and the hydroperoxy radical (HO2). We include this cycle to account for ozone chemistry following the oxidation and photolysis of water vapour. Lastly, we add the nitrogen oxide (NOx) catalytic cycle, including NO and nitrogen dioxide (NO2) to the network. We also include other oxidized nitrogen species, such as nitrate (NO3), nitrous oxide (N2O), and the reservoirs nitric acid (HNO3), and dinitrogen pentoxide (N2O5). Collectively, these nitrogen species belong to the NOy family and can also influence ozone chemistry. In our simulations, lightning is the main source of NO that initiates further NOy chemistry, as described in Section 2.4. In the upper atmosphere, the slow termolecular reaction between N2 and O(1D) provides another source of NOy, but this does not impact the lightning-induced chemistry that occurs at altitudes below 20\u2009km.","Citation Text":["Strat, Morgenstern et al. 2009"],"Functions Text":["In this study, we use the Stratospheric","chemistry schemes."],"Functions Label":["Uses","Uses"],"Citation Start End":[[771,801]],"Functions Start End":[[730,769],[869,887]]}
{"Identifier":"2019AandA...623A...1I__Rudick_et_al._(2009)_Instance_1","Paragraph":"From the observational side, the deep FDS data further confirm that the bulk of the gravitational interactions between galaxies happened on the W-NW sub-clump of the cluster. In fact, this is the only region of the cluster, inside the virial radius, where the intra-cluster baryons (diffuse light and GCs) are found, that is, the bridge between NGC 1399 and FCC 184 (Iodice et al. 2016), the intra-cluster light and GCs between FCC 184, FCC 170, and FCC 161 (Iodice et al. 2017b), and the new faint filaments between FCC 143 and FCC 147 (see Sect. 5.1 and Fig. 13). The gravitational interactions could have also modified the structure of the galaxy outskirts and produced the intra-cluster baryons. Compared with simulations by Rudick et al. (2009), the diffuse form observed for the ICL is consistent with the scenario where this component formed by stripped material from the outskirts of a galaxy in a close passage with the cD (Iodice et al. 2017b). In this area of the cluster, the stellar envelope of some ETGs is asymmetric, appearing more elongated and twisted in one direction, while the outskirts of galaxies at larger distances from the cluster centre have a more regular shape. Mastropietro et al. (2005) show how harassment can induce twists in the outer isophotes of dwarf galaxies. More massive and luminous galaxies, like the elliptical galaxies in this region of the Fornax cluster, have a deeper potential well, therefore the stripping of stars by harassment implies even stronger tidal forces. A disturbed morphology in galaxy outskirts could also result from the ongoing accretion of smaller satellites. Simulations on the mass accretion and stellar halo formation for different stellar masses (1010\u2005\u2212\u20051013\u2006M\u2299) show that the outskirts of galaxies (for \u03bcr\u2004\u223c\u200427\u2005\u2212\u200531 mag arcsec\u22122) appear with a quite disturbed morphology and with an overall elongated shape (Michel-Dansac et al. 2010; Cooper et al. 2015; Monachesi et al. 2018). The structure of the stellar envelope, as well as the shape of the SB profile, depends on the mass and number of the accreted progenitors. The ETGs showing asymmetric and diffuse envelopes (FCC 161, FCC 167, FCC 184, see Sect. 5.1 and Fig. 13) are in range of stellar mass (\u223c3\u2005\u2212\u200510\u2005\u00d7\u20051010\u2006M\u2299) comparable with simulations and they could still build up their envelope. As noticed by Iodice et al. (2017b), a fraction of the ICL population in this region of the cluster could also come from lower mass dwarf galaxies that are tidally disrupted in the potential well of the massive galaxies, which are therefore contributing to the mass assembly in their halo. This is further supported by a recent study from Venhola et al. (2017), based on FDS data, that found a drop in the number density of LSB galaxies at cluster-centric distances smaller than \u223c180 kpc.","Citation Text":["Rudick et al. (2009)"],"Functions Text":["Compared with simulations by","the diffuse form observed for the ICL is consistent with the scenario where this component formed by stripped material from the outskirts of a galaxy in a close passage with the cD"],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[729,749]],"Functions Start End":[[700,728],[751,931]]}
{"Identifier":"2022MNRAS.516.1539O__Yang_et_al._2022_Instance_1","Paragraph":"At intermediate energies, Fig. 1 shows that non-thermal bremsstrahlung processes become more important. We find this dominates the keV X-ray emission during the first \u223c2 Myr. While OY22 demonstrated this non-thermal emission would likely not be detectable in external galaxies, X-ray observations towards the Galactic Fermi bubbles reveal substantially higher emission than that computed with our model at a similar bubble age (see the comparison in OY22; also Snowden et al. 1997; Kataoka et al. 2013; Predehl et al. 2020). Although this may suggest better X-ray detection prospects for external galaxy bubbles than our results would imply, we note that these Galactic observations may include a very significant thermal bremsstrahlung contribution from all the gas in the Milky Way halo, which likely extends to a radius of \u223c250\u2009kpc (i.e. far larger than the size of our simulation box, see e.g. Blitz & Robishaw 2000; Grcevich & Putman 2009), the Galactic bulge, and gas heated by the shocks associated with the bubble (Yang et al. 2022; see also Zhang & Guo 2021), none of which is included in our model (and could evolve as the bubble ages). Moreover, features external to the Galactic bubbles (e.g. the North Polar Spur; Kataoka et al. 2013) are not included in our model, but may make a contribution to the observed X-ray emission. We thus consider a direct comparison between the X-ray emission from our model and that observed from the Galactic Fermi bubbles\/halo to be complicated by substantial thermal emission and emission from structures not associated with the bubbles. Addressing these additional contributions is non-trivial, and is not necessarily informative to predict the observational prospects of bubbles around external galaxies where thermal emission contributions could differ greatly. After the first few Myr, non-thermal bremsstrahlung X-rays are surpassed by inverse Compton emission, which dominates the X-ray emission from the bubble by 7 Myr. At these later times, a significant non-thermal bremsstrahlung component can also be seen to emerge in TeV \u03b3-rays. This is attributed to the concentration of gas and CR energy density near the top of the bubble (cf. Section 3.3), although we note that the physical strength of this emission may not be properly resolved.20 This will be explored in more detail in future work (see also Section 4.3).","Citation Text":["Yang et al. 2022"],"Functions Text":["Although this may suggest better X-ray detection prospects for external galaxy bubbles than our results would imply, we note that these Galactic observations may include a very significant thermal bremsstrahlung contribution from all the gas in the Milky Way halo, which likely extends to a radius of \u223c250\u2009kpc","the Galactic bulge, and gas heated by the shocks associated with the bubble","none of which is included in our model (and could evolve as the bubble ages)."],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1023,1039]],"Functions Start End":[[525,834],[946,1021],[1069,1146]]}
{"Identifier":"2020ApJ...901...45Z__Gopalswamy_et_al._2004_Instance_1","Paragraph":"In this paper, we are focusing on the association of the acceleration and release of SEPs with different types of CME\u2013CME interaction. In fact, the particle acceleration process is a complex one involving many factors, e.g., (1) the CME (or shock) speeds (Richardson et al. 2015; Papaioannou et al. 2016; Kouloumvakos et al. 2019), (2) the shock parameters, e.g., Mach numbers, compression ratios, and shock geometry (Lee 1983; Kozarev et al. 2015; Kouloumvakos et al. 2019), (3) the magnetic connectivity between the spacecraft and the source (Richardson et al. 2014; Xie et al. 2019), (4) the level of seed particles, and (5) the level of turbulence (e.g., Wanner & Wibberenz 1993; Laitinen et al. 2018; Strauss & le Roux 2019). Figure 2 shows a relationship between the CME 3D speed and the SEP Ip in NP3D group, which is consistent with the results in Kouloumvakos et al. (2019). The penultimate point has been widely discussed (e.g., Gopalswamy et al. 2004; Kahler & Vourlidas 2014; Ding et al. 2015), involving the role of the flare material and the interacting CMEs on seed particles. In Ding et al. (2015), the authors used a Fe\/O ratio of 2.0 as the threshold for the presence of flare material and found that all events except one with Fe\/O > 2.0 are in the \u201ctwin-CME\u201d scenario, indicating the presence of flare seed material that are possibly from pre-flares. In Kahler & Vourlidas (2014), they proposed that higher SEP Ip could be explained by increases in both CME rates and seed particles during times of high solar activity instead of being due to CME\u2013CME interaction. In Figure 7, the higher Ip in the cases of (nearly) interacting CMEs gives a hint that more seed particles will be accelerated if the priCME eruption is close to the preCME eruption, but it is still necessary to figure out which factor, i.e., the solar activity, the flare, or the preCME, mainly controls the level of the seed populations. The launch of the Parker Solar Probe (PSP) can contribute to this investigation in the future, because the spacecraft will reach as close as 8.86 Rs from the solar surface by 2024. During its first two perihelion passages, a series of SEP events have been studied, focusing on the acceleration mechanisms and seed population preconditioning (see McComas et al. 2019; Schwadron et al. 2020).","Citation Text":["Gopalswamy et al. 2004"],"Functions Text":["The penultimate point has been widely discussed (e.g.,","involving the role of the flare material and the interacting CMEs on seed particles."],"Functions Label":["Background","Background"],"Citation Start End":[[939,961]],"Functions Start End":[[884,938],[1007,1091]]}
{"Identifier":"2016MNRAS.460.3554A__Reville_&_Bell_2014_Instance_1","Paragraph":"We consider that the amplified hotspot magnetic field B is turbulent, and that the large-scale background field downstream of the reverse shock is Bjd, nearly perpendicular to the shock normal because the perpendicular component is compressed and enhanced by a factor of 4 to 7 (i.e. Bjd \u223c rBj). In such a case, to accelerate particles up to an energy Ec via a diffusive mechanism, the mean-free path \u03bbc \u223c rg(\u03b3c, B)2\/s in the shock downstream region, where B is a small-scale field, has to be smaller than Larmor radius in Bjd (Lemoine & Pelletier 2010; Reville & Bell 2014).4 The condition \u03bbc \u2272 rg(\u03b3c, Bjd) is satisfied when the magnetic-turbulence scalelength is\n\n(23)\n\r\n\\begin{equation}\r\ns \\ge \\frac{E_{\\rm c}}{eB}\\left(\\frac{B_{\\rm jd}}{B}\\right) = r_{\\rm g}(\\gamma _{\\rm s},B),\r\n\\end{equation}\r\n\nwhere rg(\u03b3s, B) is the Larmor radius of protons with energy\n\n(24)\n\r\n\\begin{eqnarray}\r\nE_{\\rm s} &amp;=&amp; E_{\\rm c}\\left(\\frac{B_{\\rm jd}}{B}\\right) = 0.07 E_{\\rm c}\\left(\\frac{r}{7}\\right) \\left(\\frac{B_{\\rm j}}{\\rm \\mu G}\\right) \\left(\\frac{B}{100 \\,\\rm \\mu G}\\right)^{-1}\\nonumber\\\\\r\n&amp;&amp;\\sim 10 \\left(\\frac{r}{7}\\right) \\left(\\frac{\\nu _{\\rm c}}{10^{14}\\,{\\rm Hz}}\\right)^{\\frac{1}{2}} \\left(\\frac{B_{\\rm jd}}{\\rm \\mu G}\\right) \\left(\\frac{B}{\\rm 100\\,\\mu G}\\right)^{{-}\\frac{5}{2}}\\,\\,{\\rm GeV},\r\n\\end{eqnarray}\r\n\nwhere we take B \u223c 100 \u03bcG and Bj \u223c \u03bcG as characteristic values. Note that\n\n(25)\n\r\n\\begin{equation}\r\n\\frac{s}{\\rm cm} > 5\\times 10^{11}\\left(\\frac{r}{7}\\right) \\left(\\frac{\\nu _{\\rm c}}{10^{14}\\,{\\rm Hz}}\\right)^{\\frac{1}{2}} \\left(\\frac{B_{\\rm jd}}{\\rm \\mu G}\\right) \\left(\\frac{B}{\\rm 100\\,\\mu G}\\right)^{{-}\\frac{5}{2}}\r\n\\end{equation}\r\n\nis greater than c\/\u03c9pi in equation (14), as required. Note however that this limit, s \u2273 500\u2009c\/\u03c9pi for typical values considered in this paper, cannot be fulfilled by Weibel-generated turbulence with scale \u223cc\/\u03c9pi. Therefore, the maximum energy achieved by electrons in the jet reverse shock, Ec, cannot be constrained by Weibel instabilities.","Citation Text":["Reville & Bell 2014"],"Functions Text":["In such a case, to accelerate particles up to an energy Ec via a diffusive mechanism, the mean-free path \u03bbc \u223c rg(\u03b3c, B)2\/s in the shock downstream region, where B is a small-scale field, has to be smaller than Larmor radius in Bjd"],"Functions Label":["Uses"],"Citation Start End":[[554,573]],"Functions Start End":[[296,526]]}
{"Identifier":"2015MNRAS.446.3631A__Brown,_Bildsten_&_Rutledge_1998_Instance_1","Paragraph":"Note that the evolution scenario in Fig. 2 is qualitatively different from previous expectations that assumed standard viscous damping and a large saturation amplitude, see e.g. Levin (1999). There it was proposed that sources slowly spin-up at low temperatures outside of the instability region (where cooling becomes slow) followed by quick r-mode heating once the source enters the instability region and similarly fast spin-down and cooling segments which complete a cycle. There would then be no gravitational wave emission from known radio pulsars since sources would leave the instability region very quickly (Levin 1999). However, this scenario is not compatible with the well-established hypothesis that LMXBs are the progenitors of millisecond pulsars, because their observed large temperatures and frequencies (Haskell et al. 2012) place them firmly inside the instability region, as shown e.g. in fig. 1 of Alford & Schwenzer (2013). These astrophysical observations are furthermore theoretically explained by deep crustal heating due to pycnonuclear reactions (Brown, Bildsten & Rutledge 1998). For large r-mode saturation amplitudes LMXBs either could not spin-up to the high frequencies of observed radio pulsars or would be spun down very quickly, which should rule out this scenario. A modified scenario that might account for the observed spin limit of pulsars would be the presence of enhanced damping that does not increase with temperature, for instance due to a viscous boundary layer (Andersson et al. 2000) or mutual friction in a superfluid and superconducting core (Haskell, Andersson & Passamonti 2009; Haskell et al. 2012). In this case r-modes could be completely damped up to the frequencies of observed radio pulsars and it would be very unlikely that any of the known pulsars currently emits gravitational waves (Andersson et al. 2000). However, estimates on mutual friction are still rather uncertain (Haskell et al. 2009, 2012), and we have recently shown that even in the benevolent scenario of a thin viscous boundary layer the damping is very likely not sufficient to explain the fastest spinning sources (Alford & Schwenzer 2013). In Fig. 2 we show the case that crustal heating dominates r-mode heating whereby a source spins up along a vertical trajectory. If r-mode heating dominates, a source would spin-up along roughly the same trajectory along which it subsequently spins down. Another option is that the spin-up stalls since a steady state is reached where the accretion spin-up is balanced by r-mode spin-down (Wagoner 2002). However these different scenarios during the LMXB phase lead to the same qualitative evolution once the accretion ends.","Citation Text":["Brown, Bildsten & Rutledge 1998"],"Functions Text":["These astrophysical observations are furthermore theoretically explained by deep crustal heating due to pycnonuclear reactions"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1074,1105]],"Functions Start End":[[946,1072]]}
{"Identifier":"2020MNRAS.492.2528L__Prieto_et_al._2017_Instance_1","Paragraph":"Massive black holes (MBHs) are ubiquitous in the Universe, and inhabit all massive galaxies (e.g. Ferrarese & Merritt 2000). MBHs are typically observed via their accretion-powered radiation as active galactic nuclei (AGNs), whose impact on to the host galaxy is invoked to explain massive galaxy quenching (e.g. Silk & Rees 1998; Di Matteo, Springel & Hernquist 2005) and the emergence of the MBH\u2013galaxy correlations (G\u00fcltekin et al. 2009; Kormendy & Ho 2013). However, the demographics of MBHs are not yet well constrained (see e.g. Reines & Comastri 2016; Greene, Strader & Ho 2019, for reviews) and the role of AGN feedback in dwarf galaxies is just starting to be explored (Penny et al. 2018; Mackay Dickey et al. 2019; Manzano-King, Canalizo & Sales 2019). In particular, theoretical studies have shown that the MBH\u2013halo occupation fraction can be low in dwarf galaxies (e.g. Greene 2012; Miller et al. 2015; Habouzit, Volonteri & Dubois 2017), and even in systems hosting MBHs, their growth is strongly suppressed because of the typically low gas densities in the host (compared to more massive disc galaxies) and the strong impact of supernova explosions, both at low and high redshift (Dubois et al. 2014; Angl\u00e9s-Alc\u00e1zar et al. 2017; Prieto et al. 2017; Trebitsch et al. 2018). Observationally, this could be partially reflected in MBHs in dwarfs being inactive for most of their life, hence in the small number of low-luminosity AGNs found (Reines, Greene & Geha 2013; Baldassare et al. 2016, 2018, 2019; Mezcua et al. 2016, 2018; Reines et al. 2019), although there are also strong observational biases that could make these AGNs difficult to find. On the other hand, if MBHs were present in many dwarfs, and efficiently grew, their feedback could significantly affect their host (see e.g. Penny et al. 2018; Mackay Dickey et al. 2019; Manzano-King et al. 2019), could play a role in the reionization of the Universe (Volonteri & Gnedin 2009), could help removing gas from massive disc progenitors (e.g. Peirani et al. 2012), and also mitigate the \u2018too-big-to-fail\u2019 problem (Garrison-Kimmel et al. 2013). Recently, Kaviraj, Martin & Silk (2019; K19 hereafter) tried to better assess the role of AGN feedback in dwarfs by jointly analysing the Hyper-Supreme Cam Subaru Strategic Program (Aihara et al. 2018b) and the WISE (Wright et al. 2010) surveys, finding that AGNs in dwarfs could exhibit very large bolometric luminosities, hence they could play a significant role in their host evolution. However, a clear consensus on the identification of AGNs in dwarfs is still missing (see e.g. Satyapal et al. 2014; Sartori et al. 2015; Marleau et al. 2017), in particular because of the low resolution of WISE (\u223c6\u2009\u2009arcsec) relative to current optical surveys, resulting in a strong source overlap, and the possible contamination of infrared emission by star formation activity that could mimic AGN activity (Hainline et al. 2016; Satyapal, Abel & Secrest 2018), especially in dwarf galaxies (Hainline et al. 2016). In this paper, we build-up on the K19 work by re-analysing their dwarf sample in more detail in the aim at better disentangling plausible AGNs in dwarf from star-forming galaxies by taking into account possible source overlapping in the sample, and also assess the MBH properties.","Citation Text":["Prieto et al. 2017"],"Functions Text":["In particular, theoretical studies have shown that the MBH\u2013halo occupation fraction can be low in dwarf galaxies","and even in systems hosting MBHs, their growth is strongly suppressed because of the typically low gas densities in the host (compared to more massive disc galaxies) and the strong impact of supernova explosions, both at low and high redshift"],"Functions Label":["Background","Background"],"Citation Start End":[[1243,1261]],"Functions Start End":[[763,875],[951,1193]]}
{"Identifier":"2022MNRAS.511.1439F__Gaensler_et_al._2002_Instance_1","Paragraph":"As discussed before, a fraction of PSRs is bound to emerge from the progenitor SNR before our fiducial final time tend = 105yr, due to the high average kick velocity that characterizes the pulsar population. The typical escape time can be estimated by matching the PSR displacement due to its kick velocity (Vpsr) with the size of the SNR in the Sedov\u2013Taylor phase:\n(14)$$\\begin{eqnarray}\r\nt_{\\rm esc} \\simeq 725\\,\\, \\mbox{kyr} \\left[\\left(\\frac{E_{\\rm sn}}{10^{51} \\, \\mbox{erg}}\\right)\\! \\left(\\frac{\\rho _0}{1\\, \\mbox{part\/cm^{3}}}\\right)^{-1}\\! \\left(\\frac{V_{\\rm psr}}{100\\, \\mbox{km\/s}}\\right)^{-5} \\right]^{\\, 1\/3}\\!\\!\\! . \\\\\r\n\\end{eqnarray}$$Considering the mean (median) value of the PSR velocity distribution of 380 (330) km s\u22121 and of the ISM number density of 0.7 (0.25) particles cm\u22123, we obtain a mean (median) escape time of tesc \u2243 88 (160) kyr. Even taking into account that transition of SNRs to the radiative phase is expected at 35 (60) kyr, the escape time only slightly reduces to tesc \u2243 77 (120) kyr. Since tend = 100 kyr, only a fraction of the sources will then escape the SNR by the end of the simulation. For those systems with tesc tend, the runaway PSR will give rise to the formation of a bow shock nebula. These nebulae, whose first examples were detected in H\u03b1 (Chevalier, Kirshner & Raymond 1980; Kulkarni & Hester 1988; Cordes, Romani & Lundgren 1993; Bell et al. 1995; van Kerkwijk & Kulkarni 2001; Jones, Stappers & Gaensler 2002; Brownsberger & Romani 2014; Dolch et al. 2016; Romani, Slane & Green 2017), more recently have been discovered and observed in X-rays and sometimes in radio (Gaensler et al. 2002, 2004; Arzoumanian et al. 2004; Chatterjee et al. 2005; Yusef-Zadeh & Gaensler 2005; Hui & Becker 2007; Hui & Becker 2008; Kargaltsev & Pavlov 2008; Misanovic, Pavlov & Garmire 2008; de Rosa et al. 2009; Ng et al. 2010; De Luca et al. 2011; Ng et al. 2012; Marelli et al. 2013; Jakobsen et al. 2014; Auchettl et al. 2015; Klingler et al. 2016; Kargaltsev et al. 2017; Posselt et al. 2017; Kim et al. 2020). They are characterized by a cometary shape, with a tiny head typically of the order of 1016 cm, whose size is set by ram pressure balance between the PSR wind and the incoming (in the PSR frame) ISM, followed by a long tail opposite to the PSR motion, which can extend for very long distances up to a few pc. Given their limited spatial extension and low residual luminosity (Kargaltsev et al. 2017), bow shock nebulae will not probably be statistically relevant in \u03b3-rays and so far have not been detected (Abdalla et al. 2018c). Runaway PSRs, however, have recently been associated to extended TeV haloes (Abeysekara 2017; Sudoh et al. 2019), most likely due to escaping pairs (Bykov et al. 2017; Evoli, Linden & Morlino 2018; Olmi & Bucciantini 2019c; Di Mauro, Manconi & Donato 2020; Evoli et al. 2021). However, the formation and properties of \u03b3-ray haloes are still poorly understood, and different interpretations lead to very different expectations in terms of the possible detection of these sources in the next future (Sudoh et al. 2019; Giacinti et al. 2020). The modelling of these complex sources is outside the scopes of the present work, so we simply keep trace of the position of escaped PSRs for possible future implementations. The fraction of PWNe escaped from their parent SNR at the end of the simulation is represented in the right-hand panel of Fig. 1, where evolved PWNe are shown on top of the initial distribution of PWNe + SNRs.","Citation Text":["Gaensler et al. 2002"],"Functions Text":["These nebulae, whose first examples were detected in H\u03b1","more recently have been discovered and observed in X-rays and sometimes in radio"],"Functions Label":["Background","Background"],"Citation Start End":[[1624,1644]],"Functions Start End":[[1236,1291],[1542,1622]]}
{"Identifier":"2017MNRAS.464.1065G__Opher_et_al._2015_Instance_1","Paragraph":"Finally in this paper, we restricted ourselves to a very limiting case when the ISM is at rest with respect to the star. This two-jet solution can, in principle, be generalized by adding the interstellar flow. Let us consider an arbitrary plane perpendicular to z-axis. This plane cuts a circle from the astropause. In the case of subsonic ISM flow, we can consider planar potential solutions around circles for each plane. According to the d'Alembert paradox, the force acting on each circle is zero. Therefore, the tube of the astropause should not be deflected into the tail, although the circle could be deformed to the ellipsoidal shape in the self-consistent solution. This scenario works, if we consider the interstellar flow to be ideal and incompressible. However, numerical results (e.g. Opher et al. 2015) show some bending of the jets towards the tail. This bending in numerical models (for slow incompressible ISM flows) is connected with the numerical dissipation inherent in the numerical schemes. Numerical viscosity may cause the boundary layer breakage on the surface of the astropause, which consequently causes the pressure asymmetry that deflects the astropause. This may be the explanation for the fact that the tube of the astropause is always deflected to the tail in the numerical models. The above-described numerical effects have nothing to do with physical dissipation phenomena responsible for the bending of real astrospheres. The physical dissipation mechanisms (e.g. magnetic reconnection, finite resistivity, Hall effects) should be explored as a possible cause of the astropause bending in the models with slow subsonic ISM flow. For the fast supersonic ISM flow, the BS is formed around the astropause. The post-shock ISM flow is vortical, and the d'Alembert paradox does not work in this case. Therefore, the bending of the astrospheric jets into the tail direction is easier to justify for the supersonic relative ISM\/SW motion.","Citation Text":["Opher et al. 2015"],"Functions Text":["However, numerical results (e.g.","show some bending of the jets towards the tail. This bending in numerical models (for slow incompressible ISM flows) is connected with the numerical dissipation inherent in the numerical schemes. Numerical viscosity may cause the boundary layer breakage on the surface of the astropause, which consequently causes the pressure asymmetry that deflects the astropause. This may be the explanation for the fact that the tube of the astropause is always deflected to the tail in the numerical models. The above-described numerical effects have nothing to do with physical dissipation phenomena responsible for the bending of real astrospheres.","The physical dissipation mechanisms (e.g. magnetic reconnection, finite resistivity, Hall effects) should be explored as a possible cause of the astropause bending in the models with slow subsonic ISM flow."],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Future Work"],"Citation Start End":[[798,815]],"Functions Start End":[[765,797],[817,1456],[1457,1663]]}
{"Identifier":"2017AandA...604A..80M__Propris_et_al._(2013)_Instance_3","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)"],"Functions Text":["In addition,","claim that previous estimates of the evolution in the RS GLF","were also due to SB effects."],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1651,1675]],"Functions Start End":[[1638,1650],[1676,1736],[1788,1816]]}
{"Identifier":"2020ApJ...904..185O__Price_et_al._2018_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":["Price et al. 2018"],"Functions Text":["Moreover, the disk formation process has been revealed to be much more complicated for binary and multiple cases,","and in numerical simulations (e.g.,"],"Functions Label":["Background","Background"],"Citation Start End":[[905,922]],"Functions Start End":[[441,554],[748,783]]}
{"Identifier":"2017ApJ...838..132O___2015_Instance_1","Paragraph":"As one of the previous studies on the atomic gas in the Perseus region, Stanimirovi\u0107 et al. (2014) calculated the \n\n\n\n\n\n optical depth toward 26 extra-Galactic radio continuum sources (such as quasars). They calculated \n\n\n\n\n\n as a function of velocity toward each radio source by using the method described in Heiles & Troland (2003). The calculated \n\n\n\n\n\n spectra were fitted with a sum of Gaussian functions, and the peak \n\n\n\n\n\n values and the Gaussian FWHMs (full width at half maximum) of each velocity component were derived. However, they found the optically thick \n\n\n\n\n\n gas only toward \n\n\n\n\n\n of lines of sight. We test the results in Stanimirovi\u0107 et al. (2014) and Fukui et al. (2014, 2015), which differ by several factors. To do this, we compare \n\n\n\n\n\n derived in Stanimirovi\u0107 et al. (2014) and that obtained in Section 3.5. Since our \n\n\n\n\n\n corresponds to the average value within given velocity width (\n\n\n\n\n\n; see Section 3.4), we compare the following two values toward each radio source located in the region we analyzed: (1) the sum of the areas of each Gaussian component of the \n\n\n\n\n\n profiles calculated in Stanimirovi\u0107 et al. (2014), and (2) the products of our \n\n\n\n\n\n and \n\n\n\n\n\n. The results are listed in columns 8 and 9 of Table 2. Note that 4C +30.04, 3C 092, and 3C 093.1 are located at the masked area; therefore \n\n\n\n\n\n cannot be calculated by our method toward them. Although there is a rough positive correlation between them, they differ \n\n\n\n\n\n times as a concrete numerical value. The correlation between them is plotted in Figure 16. The solid line indicates the one-to-one relationship (\n\n\n\n\n\n), and the dashed line is the best-fit regression line through the origin (\n\n\n\n\n\n). The results in Stanimirovi\u0107 et al. (2014) are systematically smaller than our results, and it is obvious that there is a discrepancy between these two results. This discrepancy can, however, be explained by characteristics of the data used to derive \n\n\n\n\n\n in the present study and Stanimirovi\u0107 et al. (2014).","Citation Text":["Fukui et al.","2015"],"Functions Text":["We test the results in Stanimirovi\u0107 et al. (2014) and","which differ by several factors."],"Functions Label":["Differences","Differences"],"Citation Start End":[[674,686],[694,698]],"Functions Start End":[[620,673],[701,733]]}
{"Identifier":"2020MNRAS.491.1498C__Zhang_&_Wang_2019_Instance_1","Paragraph":"We use the largest sample of FRB 121102, which is observed by GBT at 4\u20138\u2009GHz (Zhang et al. 2018). This sample contains 21 pulses reported in Gajjar et al. (2018) and 72 pulses identified by machine learning. These pulses were observed within a 6\u2009h observation. They share the same observation conditions and were observed by the same telescope. Therefore, we can put them together to analysis and ignore complex selection effects. Power-law distributions of energy \u03b1E = 1.63 \u00b1 0.06 and distributions \u03b1T = 1.57 \u00b1 0.13 for these 93 FRB 121102 bursts are shown in Fig. 5. Gourdji et al. (2019) discovered a low-energy sample with 41 bursts for FRB 121102 and found \u03b1E \u223c 1.7 if all bursts are included (see their fig. 5). However, if the low-energy bursts are discarded, a steeper \u03b1E \u223c 2.8 is found. Wang & Zhang (2019) also found that six samples of FRB 121102 bursts observed by different telescopes at different frequencies show a universal energy distribution with \u03b1E around 1.7. Meanwhile, a similar power-law index of energy distribution for non-repeating FRBs is also found (Lu & Piro 2019; Zhang & Wang 2019). The waiting time distribution of FRB 121102 can also be described by a non-stationary Poisson process with mean burst rates $\\lambda _{0}=1.23^{+0.80}_{-0.38} \\times 10^{-5} \\,\\rm ms^{-1}$. Zhang et al. (2018) found that the rate of detection is not stationary, and the distribution of waiting time cannot be well fitted using Poissonian distribution for the same sample. For a small sample of waiting times of FRB 121102, Oppermann, Yu & Pen (2018) modelled the distribution of waiting times as Weibull distribution, which can describe non-Poissonian distributions with clustering. It must be noted that the non-stationary Poissonian distribution used in this paper is similar to the Weibull distribution. Because the rate of bursts in a non-stationary Poisson process also varies with time (Wheatland et al. 1998), the mean burst rate is 1.23$^{+0.80}_{-0.38}\\times 10^{-2}$\u2009s\u22121. Using the same data, Zhang et al. (2018) found that the rate is 0.05\u2009s\u22121 for Poissonian distribution. Using a sparse waiting time sample, Oppermann et al. (2018) derived a mean repetition rate of $5.7^{+3.0}_{-2.0}$\u2009d\u22121. The large discrepancy between the two rates is that the waiting times used in this paper are derived from 93 bursts in 5\u2009h observation, comparing to 17 bursts in about 74\u2009h observation in Oppermann et al. (2018).","Citation Text":["Zhang & Wang 2019"],"Functions Text":["Meanwhile, a similar power-law index of energy distribution for non-repeating FRBs is also found"],"Functions Label":["Similarities"],"Citation Start End":[[1094,1111]],"Functions Start End":[[980,1076]]}
{"Identifier":"2016ApJ...832..183K__Kuiper_1941_Instance_1","Paragraph":"The results presented here are based on the assumptions that the CBPs do not interact with the material ejected from their binary star. However, numerical studies have indicated that this material is neither lost isotropically from the binary during the CE phase, nor does it all become unbound. Instead, 1\u201310% of the ejecta may fall back into a CB disk according to Kashi & Soker (2011), and Passy et al. (2012) suggest that \u223c80% of the ejected material may remain gravitationally bound to the binary (also see Kuiper 1941; Shu et al. 1979; and Pejcha et al. 2016 for mass loss outflows through the L2 Lagrange point24\n\n24\nMass loss through L2 results in several possible outcomes, e.g., isotropic or equatorial wind, CB disk; for details, see Table 1 and Figures 12 and 13 of Pejcha et al. (2016).\n). Either of these scenarios would significantly complicate the dynamical evolution of the system as the CBP could accrete material and gain mass and also experience migration similar to that during planetary formation.25\n\n25\nThere are, however, two potential benefits of the former in terms of detection: a more luminous planet would be more amenable to direct imaging efforts, and a more massive planet would cause stronger ETVs.\n Such accretion of material of a different specific angular momentum will change the orbital evolution of the planet, as will gravitational interaction with the bulk gas in an accretion-favorable environment such as a CB disk. Additionally, interactions between CBPs and the CE ejecta of close binary stars (including ejection or destruction of the planets) may also play an important role in the elusive mechanisms responsible for the shaping, morphology, and chemistry of planetary nebulae (see Bear & Soker 2016 for triple star origin of asymmetric planetary nebulae; for recent reviews see Zijlstra 2014 and Jones 2015, and references therein). We note that the CE-triggered dynamical disturbances discussed here occur on timescales of \u223c1 year and are thus \u201cinstantaneous\u201d compared to any subsequent planet migration, which occurs on much longer timescales.","Citation Text":["Kuiper 1941"],"Functions Text":["The results presented here are based on the assumptions that the CBPs do not interact with the material ejected from their binary star. However, numerical studies have indicated that this material is neither lost isotropically from the binary during the CE phase, nor does it all become unbound. Instead, 1\u201310% of the ejecta may fall back into a CB disk according to Kashi & Soker (2011), and Passy et al. (2012) suggest that \u223c80% of the ejected material may remain gravitationally bound to the binary (also see"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[512,523]],"Functions Start End":[[0,511]]}
{"Identifier":"2020MNRAS.493.4950S__Haines_et_al._2015_Instance_2","Paragraph":"In the framework of the hierarchical formation of structures, clusters of galaxies are continuously accreting galaxies. It has been suggested that in this process of falling, galaxies could undergo different physical processes that could affect the star formation even before they reach the cluster. Consequently, to fully understand what the cluster environment produces in galaxies, it is of key importance to have a throughout characterisation of the population of galaxies in the outskirts of clusters. Several observations have shown that properties of galaxies such as star formation, gas content, and colour are affected by the cluster environment at large clustercentric distances (e.g. Lewis et al. 2002; Solanes et al. 2002; G\u00f3mez et al. 2003; Braglia et al. 2009; Hansen et al. 2009; Park & Hwang 2009; von der Linden et al. 2010; Haines et al. 2015; Rhee et al. 2017). In particular, spiral galaxies with low star formation rates were found in the outskirts of clusters in early studies such as Couch et al. (1998) or Dressler et al. (1999). In recent years, a deficit of star-forming galaxies in the infalling region of clusters has been reported (e.g. Wetzel et al. 2013; Haines et al. 2015; Bianconi et al. 2018). This has been reproduced in simulations by Bah\u00e9 et al. (2013). These results can be explained by the presence of environmental effects accelerating the consumption of the gas reservoir before galaxies enter in a cluster, a process known as pre-processing (e.g. Fujita 2004; Mihos 2004). An important fraction of the cluster galaxies has spent time in groups or filaments before they fall into the cluster (e.g. McGee et al. 2009; De Lucia et al. 2012; Wetzel et al. 2013; Hou, Parker & Harris 2014). The population of galaxies in the outskirts of clusters includes not only galaxies that have not yet entered the cluster but also backsplash galaxies, i.e. galaxies that have passed close to the centre of the cluster since their infall and are now beyond the virial radius (e.g. Mamon et al. 2004; Gill, Knebe & Gibson 2005; Mahajan, Mamon & Raychaudhury 2011). For an adequate characterisation of the properties of galaxies that are falling into clusters, it is important to take into account the contamination by backsplash galaxies, which, having orbited through the inner regions of a cluster, could have been affected by the physical processes present in that extreme environment. The backsplash scenario in the evolution of galaxies has also been explored in Rines & Diaferio (2005), Pimbblet et al. (2006), Aguerri & S\u00e1nchez-Janssen (2010), and Muriel & Coenda (2014).","Citation Text":["Haines et al. 2015"],"Functions Text":["In recent years, a deficit of star-forming galaxies in the infalling region of clusters has been reported (e.g."],"Functions Label":["Background"],"Citation Start End":[[1186,1204]],"Functions Start End":[[1054,1165]]}
{"Identifier":"2015ApJ...799..138S__Zaritsky_et_al._1994_Instance_1","Paragraph":"We present these results with one very important caveat. Accurately determining metallicities at different redshifts is of key importance to studying the evolution of the MZR. In the local universe, relationships between strong emission line ratios and metallicity can be calibrated to \u00e2\u0080\u009cdirect\u00e2\u0080\u009d electron temperature-determined metallicities from measuring auroral lines such as [O\u00e2\u0080\u0089iii]\u00c2 \u00ce\u00bb4363 (Pettini & Pagel 2004; Pilyugin & Thuan 2005) or photoionization models of star-forming regions (Zaritsky et al. 1994; Kewley & Dopita 2002; Kobulnicky & Kewley 2004; Tremonti et al. 2004). At redshifts above z \u00e2\u0088\u00bc 1, it is nearly impossible to detect weak auroral lines for directly determining metallicity (but see Yuan & Kewley 2009; Rigby et al. 2011; Brammer et al. 2012a; Christensen et al. 2012; Maseda et al. 2014). Creating photoionization models that suitably represent high-redshift star-forming regions requires knowledge of physical parameters which have been poorly constrained up to this point. Thus, it is unknown if local metallicity calibrations hold at high redshifts. Figure 6 shows a comparison between metallicities determined using the O3N2 indicator and the N2 indicator for both local SDSS galaxies (grey points) and MOSDEF z \u00e2\u0088\u00bc 2.3 galaxies (black points). The black dashed line indicates a one-to-one relationship. If local calibrations do indeed hold at high redshifts, then the relationship between metallicities determined from different indicators should not evolve with redshift. It is clear that the z \u00e2\u0088\u00bc 2.3 galaxies are offset below the local galaxies. The dotted line is the best-fit line of slope unity to the individual z \u00e2\u0088\u00bc 2.3 galaxies, yielding an offset of \u00e2\u0088\u00920.1 dex from a one-to-one correspondence, over twice that displayed by the SDSS sample. Steidel et al. (2014) found an offset slightly larger than this at z \u00e2\u0088\u00bc 2.3. This offset demonstrates that the two metallicity indicators are not evolving in the same way with redshift, and shows the need of metallicity calibrations appropriate for high-redshift galaxies.","Citation Text":["Zaritsky et al. 1994"],"Functions Text":["In the local universe, relationships between strong emission line ratios and metallicity can be calibrated to \u00e2\u0080\u009cdirect\u00e2\u0080\u009d electron temperature-determined metallicities from","or photoionization models of star-forming regions"],"Functions Label":["Background","Background"],"Citation Start End":[[497,517]],"Functions Start End":[[176,349],[446,495]]}
{"Identifier":"2022AandA...658A..77N__Dudzevi\u010di\u016bt\u0117_et_al._(2020)_Instance_1","Paragraph":"Finally, we attempted to estimate the total obscured star formation within the Mpc scale environments of ELAN, which can be achieved by assuming that the excess number of submillimeter sources are all associated with the respective ELAN. To compute the SFR densities we first seek a proper conversion from S850 to SFRs. We utilized a sample of ALMA-identified SMGs in the UKIDSS-UDS field, which has been studied in detail with proper SED fittings (Dudzevi\u010di\u016bt\u0117 et al. 2020). Similarly to our observations, the sample of Dudzevi\u010di\u016bt\u0117 et al. (2020) was drawn from a flux-limited sample produced by the SCUBA-2 Cosmology Legacy Survey (Geach et al. 2017). Their results should therefore be representative, on average, of SCUBA-2 sources uncovered in other fields. These authors found a linear correlation of log10[SFR(M\u2299 yr\u22121)] = (0.42 \u00b1 0.06) \u00d7 log10[S870(mJy)] + (2.19 \u00b1 0.03) for their SMGs, which span a flux range of \u223c1\u2005\u2212\u200510 mJy, appropriate for the SCUBA-2 sources discovered in our target ELAN fields. We then computed the total SFR densities by integrating over a given S850 range in which the SFR contribution at a given S850 is the product of the excess number density of the submillimeter sources and their corresponding SFR based on the conversion. For each field, by considering a flux range of S850\u2004=\u20041\u2005\u2212\u200520 mJy and assuming a sphere with an equivalent circular radius of the corresponding effective area, we obtain SFR densities of 1100 \u00b1 500, 1100 \u00b1 500, 2300 \u00b1 1100, and 1400 \u00b1 1100 M\u2299 yr\u22121 Mpc\u22123 for Fabulous, Slug, Jackpot, and MAMMOTH-1 ELAN, and a weighted average SFR density of \u03a3SFR = 1200 \u00b1 300 M\u2299 yr\u22121 Mpc\u22123. We plot the results in Fig. 5, showing that they are consistent with those found in the Mpc-scale environments of other quasar samples or proto-clusters at z\u2004\u223c\u20042\u2005\u2212\u20053 (Clements et al. 2014; Dannerbauer et al. 2014; Kato et al. 2016). This result suggests that the star formation activities around ELANe are at a similar level of other dense systems in this redshift range, or, in other words, at a factor of about 300 greater than the cosmic mean.","Citation Text":["Dudzevi\u010di\u016bt\u0117 et al. 2020","Dudzevi\u010di\u016bt\u0117 et al. (2020)"],"Functions Text":["We utilized a sample of ALMA-identified SMGs in the UKIDSS-UDS field, which has been studied in detail with proper SED fittings","Similarly to our observations, the sample of","was drawn from a flux-limited sample produced by the SCUBA-2 Cosmology Legacy Survey"],"Functions Label":["Uses","Similarities","Similarities"],"Citation Start End":[[449,473],[521,547]],"Functions Start End":[[320,447],[476,520],[548,632]]}
{"Identifier":"2015ApJ...810...96S___2010_Instance_1","Paragraph":"Most of our knowledge about reconnection comes from effectively 2D reconnection experiments, and only recently have efforts to understand 3D magnetic reconnection been pursued (Priest 2011; Pontin 2011; Shepherd & Cassak 2012; Janvier et al. 2014), revealing a much wider range of dynamics. Very little is known about the actual properties of the solar plasma during the reconnection process on the Sun. Many authors have studied the global and local reconnection rates and timescales of solar reconnection during solar flares based on analysis of the ribbon motions in the chromosphere (Fletcher & Hudson 2001; Qiu et al. 2002, 2010; Isobe et al. 2005; Jing et al. 2005; Saba et al. 2006; Miklenic et al. 2007; Xie et al. 2009). Generally, they use measurements of the magnetic flux at the photosphere, the velocity of the ribbons parallel to the PIL, and the 2D approximation of the standard flare model to determine the reconnected flux and the electric field and Poynting flux at the reconnection site. In a future paper, we plan to test the validity of this approximation. However, these studies are purely based on the ribbon dynamics and are not able to capture the dynamics of the 3D coronal magnetic field that is actually involved in the reconnection process. For example, to calculate the energy release rate from reconnection, one needs to combine observations with knowledge of the location and size (i.e., length\u2014the width is generally below the resolution) of the reconnection current sheet, which can only be provided by data-constrained magnetic field models and extrapolations or by data-driven MHD simulations with the appropriate NLFFF initial conditions (and is still only an approximation). Aulanier et al. (2000, 2012) and Schrijver et al. (2011) utilize idealized MHD simulations to interpret the appearance of the flare ribbons. Although such studies show a qualitative similarity of the QSLs and flare ribbons and are valuable for determining the basic topology of the region, no actual estimates of the reconnection parameters can be obtained, which could be improved by the use of data-constrained and data-driven magnetic field modeling.","Citation Text":["Qiu et al.","2010"],"Functions Text":["Many authors have studied the global and local reconnection rates and timescales of solar reconnection during solar flares based on analysis of the ribbon motions in the chromosphere"],"Functions Label":["Background"],"Citation Start End":[[612,622],[629,633]],"Functions Start End":[[404,586]]}
{"Identifier":"2018AandA...618A..62P__Cantalupo_et_al._2010_Instance_1","Paragraph":"The recovery of the sky signal from these huge, noisy time streams, a process called map-making, represents one of the most important steps in CMB data analysis and, if the detector noise properties and scanning strategy are known, map-making becomes a linear inverse problem. The generalized least-squares (GLS) equation provides an unbiased solution to map-making for an arbitrary choice of weights given by a symmetric and positive definite matrix (Tegmark 1997a). Moreover, if we consider the inverse covariance of the time domain noise as the weights, the GLS estimate is also a minimum variance and a maximum likelihood solution to the problem. However, computation of the solution in such a case may require either an explicit factorisation of a huge, dense matrix (Tegmark 1997a; Borrill 1999; Stompor et al. 2002) or an application of some iterative procedure (Wright 1996; Oh et al. 1999; Dor\u00e9 et al. 2001; de Gasperis et al. 2005; Cantalupo et al. 2010). These latte typically involve several matrix-vector multiplications at each iteration step. What makes the map-making problem particularly challenging are the sizes of the current and forthcoming CMB data sets which are directly related to the number of floating point operations (flops) needed to achieve the solution and to the memory requirements due to the sizes of the arrays required for it. Both these factors set the requirements on computational resources and indeed many current CMB data analysis pipelines opt for massively parallel computing platforms. However, even in such circumstances, efficient algorithms are necessary to ensure that the analysis can indeed be performed. The computational complexity of the algorithms involving an explicit matrix inversion is \n\n\n\nO\n(\n\n\nN\n\n\np\n\n\n3\n\n\n)\n\n\n$ \\textstyle\\mathcal O(\\mathbf N_\\mathrm p^3) $\n\n\n flops, where Np is the number of pixels in the map, and therefore they are only suitable for cases where the estimated sky maps do not involve many sky pixels. Nonetheless, whenever feasible, the direct approaches can yield high-precision, unbiased estimates of the sky signal (e.g. Poletti et al. 2017, for a recent example). The next generations of the ground experiments, CMB-Stage III (Arnold et al. 2014; Henderson et al. 2016; Benson et al. 2014) and IV (Abazajian et al. 2016), however, are expected to observe significant fractions of the entire sky with high resolution, thus resulting in maps with Np \u2243 \ud835\udcaa(106), rendering the direct approaches prohibitive even for the largest forthcoming supercomputers.","Citation Text":["Cantalupo et al. 2010"],"Functions Text":["Moreover, if we consider the inverse covariance of the time domain noise as the weights, the GLS estimate is also a minimum variance and a maximum likelihood solution to the problem. However, computation of the solution in such a case may require","or an application of some iterative procedure","These latte typically involve several matrix-vector multiplications at each iteration step."],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[942,963]],"Functions Start End":[[468,714],[823,868],[966,1057]]}
{"Identifier":"2016MNRAS.461..839K__Galazutdinov,_LoCurto_&_Krelowski_2008b_Instance_1","Paragraph":"Diffuse interstellar bands (DIBs) are absorption features seen in the spectra of many astronomical objects in the visible and infrared wavelength regions. The total number of these features exceeds 400. These bands are observed in absorption in star light crossing translucent clouds and the carriers of only three such features have been proposed; none of them certain. The origin of the DIBs is as puzzling as since their first detection, more than 90 years ago. The central wavelengths of DIBs are not readily identified with any known atomic or molecular spectral lines. The profiles of DIBs are not just Gaussians as shown for the first time by Westerlund & Kre\u0142owski (1988) which may suggest their molecular carriers. Moreover, precisely determined DIB profiles (Galazutdinov, LoCurto & Krelowski 2008b), showing a specific substructure pattern each, should facilitate their carrier's identification. The better-substantiated hypothesis is that DIBs arise from absorption produced by polycyclic aromatic hydrocarbons (PAHs), although Cox (2011) and Salama & Ehrenfreund (2014) doubt this statement and no specific PAH has been identified until recently (Salama et al. 2011). The other possibility is that carbon chain molecules, similar to those observed in star-forming regions, may carry DIBs. but this idea also has no clear proof (Motylewski et al. 2000). Two DIBs, observed in the near-infrared, were tentatively ascribed to the bucky ball (C$_{60}^{+}$) molecule (e.g. Foing & Ehrenfreund 1994; Campbell et al. 2015). Nevertheless there is an indication that this identification was proposed but not finally proven. Recently it has been demonstrated that DIB profiles (at least some of them) change in unison with the rotational temperature of simple homonuclear molecules C2 and C3 (Ka\u017amierczak et al. 2009). The relation between temperatures of simple carbon chains and those of DIB carriers may suggest that there is a common chemical history and that the DIB carrier may be polyatomic centrosymmetric (non-polar) molecules as well. This is likely why the earlier attempts to match the electronic spectra of known polar molecules (identified in microwave region of wavelengths) and DIBs failed (Motylewski et al. 2000). In a certain sense, this phenomenon still remains an enigma for astrophysicists. For recent reviews see Cami & Cox (2014).","Citation Text":["Galazutdinov, LoCurto & Krelowski 2008b"],"Functions Text":["Moreover, precisely determined DIB profiles","showing a specific substructure pattern each, should facilitate their carrier's identification."],"Functions Label":["Uses","Uses"],"Citation Start End":[[769,808]],"Functions Start End":[[724,767],[811,906]]}
{"Identifier":"2022ApJ...940...72R__Camilo_et_al._2006_Instance_1","Paragraph":"Several studies have discussed the radio luminosity of GLEAM-X J1627 during its radio outburst in comparison with the limits of its rotational energy (Erkut 2022; Hurley-Walker et al. 2022). In particular, assuming isotropic emission, the radio luminosity of the brightest single peaks (L\nradio \u223c 1030\u20131031 erg s\u22121; Hurley-Walker et al. 2022) exceeds the limits on the rotational power of the source by a few orders of magnitude. Figure 6 shows those peak radio luminosities and the rotational power of GLEAM-X J1627 in comparison with other pulsars, rotating radio transients (RRATs) and radio-loud magnetars. For the radio-loud magnetars, given their large variability, we have chosen the brightest radio pulses reported in the literature (data collected from Camilo et al. 2006, 2007; Weltevrede et al. 2011; Deller et al. 2012; Lynch et al. 2015; Majid et al. 2017; Pearlman et al. 2018; Lower et al. 2020, and Esposito et al. 2021). It is well known that assuming isotropic radio emission is not realistic, and a beaming factor necessarily has to be present (see, e.g., Erkut 2022). However, the relation between the duty cycle and the spin period of canonical pulsars has a large spread (Manchester et al. 2005). Moreover, it is observed that this relationship does not apply to radio-loud magnetars, which in general show larger duty cycles than what one would expect from the extrapolation of this tentative relation for radio pulsars to magnetars (see, e.g., Camilo et al. 2006, 2007). To avoid the uncertainty of beaming models, which for magnetars are mostly unknown even theoretically, we plotted the isotropic radio luminosity for all the different pulsar classes in Figure 6. From this plot, at variance with canonical radio pulsars, we see how the brightest single peaks for radio-loud magnetars might exceed their rotational powers, in line with what is possibly observed for GLEAM-X J1627. While not resolving uncertainties related to the exact mechanism of radio emission or the beaming factor, Figure 6 shows that, under the assumption of isotropic emission, even for magnetars the brightest single peaks exceed their rotational energy budget. Considering all the uncertainties in the assumptions used to derive the radio luminosities plotted in Figure 6, GLEAM-X J1627\u2019s radio luminosity excess over its rotational power cannot be used as an argument for or against its neutron star nature.","Citation Text":["Camilo et al. 2006"],"Functions Text":["For the radio-loud magnetars, given their large variability, we have chosen the brightest radio pulses reported in the literature (data collected from"],"Functions Label":["Uses"],"Citation Start End":[[762,780]],"Functions Start End":[[611,761]]}
{"Identifier":"2021AandA...654A.132B__Esquej_et_al._(2014)_Instance_1","Paragraph":"A third approach to assessing the central star formation is based on the presence of poly-aromatic hydrocarbons (PAHs), which are known to trace young stars, but over a wider range of ages (\u223c100 Myr) than the ionised line emission. However, there is a debate about whether the PAHs are excited by stars or by the AGN itself, or whether small PAHs are destroyed by the hard radiation from an AGN (Siebenmorgen et al. 2004; Smith et al. 2007; Sales et al. 2010). This may depend on which PAH feature one considers: Diamond-Stanic & Rieke (2010) found that in AGNs on kpc scales the 6.2, 7.7, and 8.6 \u03bcm features were suppressed, while the 11.3 \u03bcm feature was not. Partly for this reason, recent studies of PAHs close to AGNs have focussed on the feature at 11.3 \u03bcm. Esquej et al. (2014) detected PAHs on subarcsec scales in about half of their sample of 29 nearby Seyfert galaxies. They argued that the high column densities in the torus around the AGN would provide shielding that enables PAHs to survive, and the implied central star formation rate density was much higher than in the circum-nuclear region. A similar conclusion about PAH survival was reached by Esparza-Arredondo et al. (2018) for their sample of 19 AGNs, based on the lack of relation between the X-ray luminosity and a central PAH deficit \u2013 while PAHs were detected in most AGNs, the equivalent width was lower in the centre. They attributed this to low star formation rates in that region. On the other hand, while Jensen et al. (2017) also detected PAHs in a sample of 13 AGNs, they argued that the similarity and slope of the radial profiles from tens to hundreds of parsecs point towards an origin in a single compact source of excitation. In order to shed more light on the excitation of PAHs in this context, Alonso-Herrero et al. (2020) compared PAH line ratios in 22 AGNs to models of PAH excitation and analysed this in the context of the measured molecular gas content. They concluded that PAHs can be shielded from the hard AGN radiation.","Citation Text":["Esquej et al. (2014)"],"Functions Text":["detected PAHs on subarcsec scales in about half of their sample of 29 nearby Seyfert galaxies.","They argued that the high column densities in the torus around the AGN would provide shielding that enables PAHs to survive, and the implied central star formation rate density was much higher than in the circum-nuclear region.","A similar conclusion about PAH survival was reached by Esparza-Arredondo et al. (2018) for their sample of 19 AGNs, based on the lack of relation between the X-ray luminosity and a central PAH deficit \u2013 while PAHs were detected in most AGNs, the equivalent width was lower in the centre."],"Functions Label":["Background","Compare\/Contrast","Similarities"],"Citation Start End":[[764,784]],"Functions Start End":[[785,879],[880,1107],[1108,1395]]}
{"Identifier":"2021MNRAS.506.1045M__Kotani_et_al._1994_Instance_1","Paragraph":"Discovered in 1977 from its bright H \u03b1 emission (Stephenson & Sanduleak 1977), SS433\u2019s defining characteristics are undoubtedly the helical motion of highly collimated jets of plasma launched from its innermost regions, and mass-loaded, non-polar outflows (Fabian & Rees 1979; Margon et al. 1979) which together inflate the surrounding W50 supernova remnant. Knots in SS433\u2019s jet can be resolved at radio frequencies using very long baseline interferometry (VLBI) and indicate the presence of highly relativistic electrons (Vermeulen et al. 1987), while the baryon content is revealed by emission lines ranging from H and He lines in the optical through to highly ionized Fe lines in the X-rays (Kotani et al. 1994; Marshall et al. 2013). The Doppler shifts of the lines indicate precession of the accreting system with a period of \u2248 162 d, also seen in optical (He ii) emission lines originating from the non-polar wind (Fabrika 1997). Both the jets and winds carry a large kinetic luminosity (>1038 erg s\u22121, e.g. Marshall et al. 2002), which requires extraction of energy via accretion on to a compact object. While the nature of the compact object in SS433 remains somewhat unknown (although dynamical arguments suggest the presence of a black hole \u2013 Blundell, Bowler & Schmidtobreick 2008), the rate of mass transfer from the companion star, as inferred from the IR excess (Shkovskii 1981; Fuchs et al. 2006), is thought to be \u223c1 \u00d7 10\u22124 M\u2299 yr\u22121, orders of magnitude in excess of the Eddington limit for any plausible stellar remnant (>300 times the Eddington mass accretion rate for a typical stellar mass black hole of around 10 M\u2299). Classical theory and radiation magnetohydrodynamic (RMHD) simulations agree that such \u2018super-critical\u2019 rates of accretion will lead to a radiatively supported, large scale height (H\/R \u2248 1, where H is the height of the disc at distance R from the compact object) accretion disc with powerful winds launched from the surface at mildly relativistic speeds (Shakura & Sunyaev 1973; Poutanen et al. 2007; Ohsuga & Mineshige. 2011; Takeuchi et al. 2013; Jiang et al. 2014; Sadowski et al. 2014).","Citation Text":["Kotani et al. 1994"],"Functions Text":["Knots in SS433\u2019s jet can be resolved at radio frequencies using very long baseline interferometry (VLBI) and indicate the presence of highly relativistic electrons",", while the baryon content is revealed by emission lines ranging from H and He lines in the optical through to highly ionized Fe lines in the X-rays"],"Functions Label":["Background","Background"],"Citation Start End":[[696,714]],"Functions Start End":[[359,522],[546,694]]}
{"Identifier":"2020MNRAS.494.5110B__Troja_et_al._2018_Instance_1","Paragraph":"Following the short gamma-ray burst (sGRB) associated with this event, GRB 170817A (Abbott et al. 2017a,b; Goldstein et al. 2017), radio emission was anticipated as the associated merger outflow interacted with the circum-merger medium. Monitoring the radio emission could therefore provide crucial information on the energetics and geometry of the outflow, as well as the ambient environment. At radio frequencies, telescopes were observing the Advanced LIGO\u2013Virgo probability region for GW170817 within 29\u2009min post-merger (Callister et al. 2017a), and subsequent monitoring of AT 2017gfo resulted in an initial radio detection 16\u2009d after the event (Abbott et al. 2017a; Hallinan et al. 2017). Further monitoring, predominantly at frequencies between 0.6 and 15\u2009GHz, has since taken place (e.g. Alexander et al. 2017, 2018; Corsi et al. 2018; Dobie et al. 2018; Margutti et al. 2018; Mooley et al. 2018a,b,c;Resmi et al. 2018; Troja et al. 2018, 2019). At these frequencies, a general picture emerged in which the radio light curve was first observed to steadily rise, before it turned over and began a more rapid decay. Using a compilation of 0.6\u201310\u2009GHz radio data from 17 to 298\u2009d post-merger, Mooley et al. (2018c) derived both a fitted time for the radio peak of 174$^{+9}_{-6}$ d and a fitted 3-GHz peak flux density of 98$^{+8}_{-9}\\, \\mu$Jy (also see similar analyses in Dobie et al. 2018 and Alexander et al. 2018). The fitted radio spectral index \u03b11 from this study is \u22120.53 \u00b1 0.04, consistent with broad-band spectral indices determined using radio, optical, and X-ray data at various epochs, where the typical value is approximately \u22120.58 (e.g. Alexander et al. 2018; Margutti et al. 2018; Troja et al. 2018, 2019; Hajela et al. 2019). Mooley et al. (2018c) also found power-law dependencies for the rise and decay phases of approximately t0.8 and t\u22122.4, respectively, where t is the time since the merger. Within the associated uncertainties, these results are consistent with the broad-band evolution of AT 2017gfo (e.g. Alexander et al. 2018; Hajela et al. 2019; Lamb et al. 2019; Troja et al. 2019).","Citation Text":["Troja et al. 2018"],"Functions Text":["Further monitoring, predominantly at frequencies between 0.6 and 15\u2009GHz, has since taken place (e.g."],"Functions Label":["Background"],"Citation Start End":[[928,945]],"Functions Start End":[[695,795]]}
{"Identifier":"2020MNRAS.499.1788W__Malhotra_et_al._2001_Instance_1","Paragraph":"Owing to their brightness at rest-frame FIR wavelengths, the ionized and neutral species of Carbon, Nitrogen, and Oxygen are powerful diagnostic lines for tracing the ISM of nearby and distant galaxies. When combined with photodissociation region (PDR) models (Tielens & Hollenbach 1985), measurement of the emission from different lines provide a means to constrain quantities such as the ionization rate and metallicity of the ISM of galaxies. Multilevel FIR transition lines have now been widely surveyed in local galaxies, originally by ISO (e.g. Luhman et al. 1998; Malhotra et al. 2001) and more recently using the PACS spectrometer (Poglitsch et al. 2010) on the ESA Herschel Space Observatory (Pilbratt et al. 2010). The [C\u2009ii]158\u2009\u03bcm line is typically the brightest in star-forming galaxies, arising from ionized, and even neutral gas where it is the main coolant (Wolfire et al. 2003). Another commonly observed FIR line tracer of the ionized ISM is [N\u2009ii]122\u2009\u03bcm. It has the advantage that it can be found associated with lower excitation gas, close to that observed in our own Galaxy (e.g. Goldsmith et al. 2015; Herrera-Camus et al. 2016). Another major coolant of the ISM is [O\u2009i]63\u2009\u03bcm (Wolfire et al. 2003). Owing to its high excitation temperature and critical density, it can dominate the cooling in regions of starburst activity (Kaufman et al. 1999; Kaufman, Wolfire & Hollenbach 2006; Brauher, Dale & Helou 2008; Croxall et al. 2012; Narayanan & Krumholz 2017). When combined with measurements of the [C\u2009ii]158\u2009\u03bcm line intensity and FIR luminosity, the [O\u2009i]63\u2009\u03bcm line intensity can constrain the FUV field, G, and the gas density using PDR models. The luminosity in these FIR lines generally exhibits a deficit in the most FIR luminous galaxies compared to the trend expected from lower luminosity galaxies (e.g. Malhotra et al. 2001; Graci\u00e1-Carpio et al. 2011; D\u00edaz-Santos et al. 2017). This has made the emission from lines like [O\u2009i]63\u2009\u03bcm more challenging to detect at high-redshifts.","Citation Text":["Malhotra et al. 2001"],"Functions Text":["Multilevel FIR transition lines have now been widely surveyed in local galaxies, originally by ISO (e.g."],"Functions Label":["Background"],"Citation Start End":[[571,591]],"Functions Start End":[[446,550]]}
{"Identifier":"2020MNRAS.493.4960T__Young,_Ross_&_Fabian_1998_Instance_1","Paragraph":"The albedo profile is then computed by cloudy after having specified the values of temperature T(r) and density \u03c1(r) for each radial patch (with r the radial distance). Results for the energy-dependent albedo profile A(Eloc) obtained for different radial patches in the cases of a = 0 (left), 0.9 (centre), and 0.998 (right) are shown in Fig. 8. As it can be clearly seen, in all the cases explored the simplifying assumption of 100 per\u2009cent albedo turns out to be a good approximation only at very high energies, i.e. Eloc \u2248 10\u2013100 keV. Elsewhere, the albedo significantly deviates from unity (except for some values of r at lower energies), especially in the 0.1\u201310 keV band, which is indeed the working energy range of the forthcoming X-ray polarimeters like IXPE. Here different line features appear, more or less visible depending on the BH spin and radial distance, such as the clear iron absorption edge that occurs at around \u223c6\u20137 keV (see e.g. Young, Ross & Fabian 1998). Plots in Fig. 9 give a more exhaustive view on how the albedo profile depends on the density of the slab in which calculations are performed. In this case the outputs are obtained for a = 0.998, three different values of the temperature, i.e. those corresponding to r = 2 (left), 5 (middle), and $10\\, r_{\\rm g}$ (right) according to the Novikov & Thorne (1973) temperature profile, and different values of the total hydrogen density n(H) between 1015 and 1023 cm\u22123. The plots show that the dependence of A(Eloc) on the density is stronger at low energies, where it exhibits an increasing behaviour by decreasing n(H). In particular, it attains values close to 0 at around 0.1 keV for particle densities in excess of 1022\u20131023 cm\u22123, especially for lower temperatures (i.e. for larger radial distances, see the right-hand panel). On the other hand, at higher energies ($\\gtrsim 10$\u201320 keV), the albedo tends to reach the same value (\u22481) in the entire range of densities explored.","Citation Text":["Young, Ross & Fabian 1998"],"Functions Text":["Here different line features appear, more or less visible depending on the BH spin and radial distance, such as the clear iron absorption edge that occurs at around \u223c6\u20137 keV (see e.g."],"Functions Label":["Uses"],"Citation Start End":[[952,977]],"Functions Start End":[[768,951]]}
{"Identifier":"2015MNRAS.454.1468K__Winckel_2003_Instance_1","Paragraph":"Owing to their dusty circumstellar environments, a large mid-infrared (mid-IR) excess is a characteristic feature of post-AGB stars and a detection of cold circumstellar material using mid-IR photometry can be used to identify these objects. The first extensive search for these objects was initiated in the mid-80's using results from the Infrared Astronomical Satellite (Neugebauer et al. 1984) which enabled the identification of post-AGB stars in our Galaxy (Kwok 1993). The Toru$\\acute{\\rm n}$ catalogue (Szczerba et al. 2007) for Galactic post-AGB stars lists around 391 very likely post-AGB objects. The Galactic sample of post-AGB stars have been found to be a very diverse group of objects (Van Winckel 2003). Studies showed that the majority of the optically visible Galactic post-AGB stars could be classified based on their spectral energy distributions (SEDs) into two groups: shell-sources and disc-sources (Van Winckel 2003). The shell-sources show a double-peaked SED with the hot central star peaking at shorter wavelengths while the cold, detached, expanding dust shell peaks at longer wavelengths. This type of SED is considered to be characteristic of objects that follow the single-star evolution scenario mentioned above. The disc-sources do not show two distinct flux peaks in the mid-IR but they do display a clear near-infrared (near-IR) excess indicating that circumstellar dust must be close to the central star, near sublimation temperature. It is now well established that this feature in the SED indicates the presence of a stable compact circumbinary disc, and therefore these sources are referred to as disc-sources (de Ruyter et al. 2006; Deroo et al. 2007; Gielen et al. 2011a; Hillen et al. 2013). The rotation of the disc was resolved with the ALMA array (Bujarrabal et al. 2013a) in one object and using single dish observations Bujarrabal et al. (2013b) confirmed that disc rotation is indeed widespread. Moreover, these disc-sources are confirmed to be binaries and show orbital periods between 100 and 2000 d (Van Winckel et al. 2009; Gorlova et al. 2014). In contrast, for the Galactic shell-sources long-term radial velocity monitoring efforts have not yet resulted in any clear detected binary orbit (Hrivnak et al. 2011), which either confirms the single-star nature of these objects or introduces a possibility that these systems can have companions on very wide orbits.","Citation Text":["Van Winckel 2003"],"Functions Text":["The Galactic sample of post-AGB stars have been found to be a very diverse group of objects"],"Functions Label":["Background"],"Citation Start End":[[700,716]],"Functions Start End":[[607,698]]}
{"Identifier":"2020MNRAS.499.3792B__Pimbblet_2011_Instance_2","Paragraph":"For consistency, we adopt the translation of this to the absolute velocities of cluster galaxies normalized by their respective galaxy cluster velocity dispersions into the range $0.3 \\lt |\\Delta \\mathrm{V}|\/\\sigma _{r_{200}} \\lt 0.5$ as deduced by Pimbblet (2011). Thus, if the mode of the standardised velocities for a sub-population has its foci at around $0.3 \\lt |\\Delta \\mathrm{V}|\/\\sigma _{r_{200}} \\lt 0.5$ for values around the virial radius, which we assume to be Rvirial \u223c r200, said sub-population would be classified as infalling. In contrast, a sub-population of backsplash cluster galaxies would be expected to peak significantly at $|\\Delta \\mathrm{V}|\/\\sigma _{r_{200}}\\sim 0$ for values at or beyond our definition of the virial radius, with their fraction reaching zero at some upper limit (e.g. Mamon et al. 2004; Pimbblet 2011; Bah\u00e9 et al. 2013; Haggar et al. 2020). Therefore, with respect to Fig. 5, we see that the column of our non-merging sub-populations across both bins of radius do not show any significant difference in the distributions of velocities with the exception of those that lie \u2264r200, which show the non-AGN sub-population to occupy a mode within the range that nominally represents infallers, most likely for cluster galaxies 0.5 \u2264 r200  1.0 (Gill et al. 2005). Additionally, the AGN sub-population slightly deviates from the non-AGN velocity distribution with a mode centred at $|\\Delta \\mathrm{\\it V}|\/\\sigma _{r_{200}}\\sim 0.8$, which could indicate stronger infalling. In contrast, the column of our merging AGN sub-populations shows the strongest deviations from the distribution of non-AGN, especially with the >r200 bin showing a significant centrally dominated AGN sub-population, where such a central dominance in relative velocity corresponds to a sub-population that were predominantly backsplash cluster galaxies. However, the dependence of this being the true nature of the sub-population relies upon more precise definitions of the radii since there is a natural upper limit a bound cluster galaxy can extend outward to with respect to its galaxy cluster\u2019s potential, known as the splashback radius (More et al. 2015, 2016). In addition, Haggar et al. (2020) show that the fraction of backsplash galaxies diminishes by 2r200 and 2.5r200 for massive (\u223c\u00d71015M\u2299) merging and non-merging cluster systems, respectively, thus demonstrating that merging cluster environments experience a greater decrease in the fraction of harbouring backsplash galaxies as one continues to extend beyond r200. Indeed, the sub-populations of the merging cluster galaxies present in the \u2264r200 bin show more variations in their general distributions with the modes of both the AGN and non-AGN sub-populations lying around $0.3 \\lt |\\Delta \\mathrm{V}|\/\\sigma _{r_{200}} \\lt 0.5$, which eludes to mostly infalling sub-populations rather than those associated with backsplash. Finally, if one considers the equivalent peak of the AGN density histogram at $\\Delta \\mathrm{V}|\/\\sigma _{r_{200}}\\sim 1.7$ it could be possible there is a mix of recently accreted cluster galaxies and those that are relaxing on to a common potential. Although it should be noted that not much information can be confidently derived from the AGN sub-populations within the bins that possess small samples size (N \u2272 100), especially with the merging AGN-hosting cluster galaxies at \u2264r200 that only has N = 15.","Citation Text":["Pimbblet 2011"],"Functions Text":["In contrast, a sub-population of backsplash cluster galaxies would be expected to peak significantly at $|\\Delta \\mathrm{V}|\/\\sigma _{r_{200}}\\sim 0$ for values at or beyond our definition of the virial radius, with their fraction reaching zero at some upper limit (e.g."],"Functions Label":["Uses"],"Citation Start End":[[834,847]],"Functions Start End":[[544,814]]}
{"Identifier":"2022ApJ...935..137K__Whittet_et_al._2001_Instance_1","Paragraph":"In Figure 2, the reduced AKARI IRC spectra of all protostars are presented; the absorption features of the H2O, CO2, and CO ices are clearly detected. All of our targets show deep and broad absorption features of H2O ice in the wavelength range 2.7\u20133.4 \u03bcm. In the case of AFGL 7009S, strong extinction toward the source saturates the absorption features throughout the wavelengths from 2.7 to 3.6 \u03bcm, including H2O ice. Other ice features, such as CH4 (Lacy et al. 1991; Boogert et al. 2004) and CH3OH (Grim et al. 1991; Brooke et al. 1996), were observed at 3.3\u22123.5 \u03bcm, but it is difficult to extract their absorption profiles from the blended features due to the low spectral resolution of AKARI IRC. An absorption feature of the CO2 ice around 4.27 \u03bcm was clearly detected toward all targets. At the wavelength around 4.6 \u03bcm for Perseus 1 and 3, RNO 91, and AFGL 7009S, there is a hint for another ice component overlapping with the CO absorption feature at 4.67 \u03bcm. Many near-infrared observations have revealed the same feature on the blue wing part of the CO absorption feature (Tegler et al. 1995; Chiar et al. 1998; Whittet et al. 2001; van Broekhuizen et al. 2005; Aikawa et al. 2012), which was suggested as an absorption feature of XCN ice. Lacy et al. (1984) and Pendleton et al. (1999) reported that the 4.62 \u03bcm absorption feature of XCN ice consists of a nitrile group and an unknown component \u201cX\u201d. Many laboratory studies have suggested that UV photolysis or cosmic ray irradiation of ice mantle could make the solid state OCN\u2212 on grain surfaces (Lacy et al. 1984; Grim & Greenberg 1987; Bernstein et al. 2000; Palumbo et al. 2000; Hudson et al. 2001; van Broekhuizen et al. 2004). In addition to these ice components, there are some absorption features around 4.8 and 4.9 \u03bcm. For Perseus 3 and the background star, the absorption features with a peak position around 4.78 \u03bcm are likely associated with 13CO ice (Boogert et al. 2002; Pontoppidan et al. 2003). We also detected another absorption feature at 4.83 \u03bcm toward the low-luminosity targets. However, we could not find any corresponding ice features from previous studies. The 4.9 \u03bcm absorption feature detected toward all targets was identified as solid carbonyl sulfide (OCS) ice. OCS ice can be produced when the interstellar ices containing CO and CO2 are exposed to UV photons or cosmic rays (Palumbo et al. 1997).","Citation Text":["Whittet et al. 2001"],"Functions Text":["At the wavelength around 4.6 \u03bcm for Perseus 1 and 3, RNO 91, and AFGL 7009S, there is a hint for another ice component overlapping with the CO absorption feature at 4.67 \u03bcm. Many near-infrared observations have revealed the same feature on the blue wing part of the CO absorption feature","which was suggested as an absorption feature of XCN ice."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1124,1143]],"Functions Start End":[[796,1083],[1195,1251]]}
{"Identifier":"2015AandA...582A..26G__model,_Young_(2014)_Instance_1","Paragraph":"In addition, if 26Al in the solar system was inherited from a molecular cloud, one would expect the solar system to have a (60Fe\/26Al)0 ratio identical to that of the ISM. If the collapse timescales \u2013 depending on the amount of turbulence, the intensity of magnetic fields, and other complex parameters (McKee & Ostriker 2007) \u2013 exceed a few 26Al half lives, the solar system\u2019s initial (60Fe\/26Al)0 ratio is even expected to be higher than that of ISM because 26Al decays much faster than 60Fe. Both radionuclides have now been identified in the ISM via \u03b3-ray astronomy (Diehl 2014). The measured ISM flux ratio 60Fe\/26Al ratio is 0.15 (Wang et al. 2007), which translates into a mass ratio of 0.35. The ISM 60Fe\/26Al mass ratio is therefore two orders of magnitude higher than the solar value of 3.9 \u00d7 10-3 (see Sect. 1), ruling out that both radionuclides were inherited from the ISM. Despite the observational evidence, and to keep alive the inherited model, Young (2014) proposed to decrease the theoretical 60Fe abundance in molecular clouds by assuming that stars more massive than 30 M\u2299 do not explode as SNe. These stars would not contribute to the 60Fe inventory, bringing the theoretical ISM 60Fe\/26Al ratio closer to that of the solar system. Because stars more massive than 30 M\u2299 contribute to less than 50% of the 60Fe production2, this proposition cannot resolve the two orders of magnitude discrepancy between the model and observations. Furthermore, the recent discovery of a supernova whose progenitor mass was far above the threshold of 30 M\u2299 is at odds with that solution (Gal-Yam et al. 2014). Finally, if 26Al was inherited from the natal molecular cloud, it would be homogeneously distributed in the SPD, which is contrary to observations (Krot et al. 2012). As discussed by Gounelle & Meynet (2012), some heterogeneity in the 26Al distribution is expected in the dense shell model because the shell collapse timescale (a few 105 yr) is comparable to the 26Al half life. ","Citation Text":["Young (2014)"],"Functions Text":["Despite the observational evidence, and to keep alive the inherited model,","proposed to decrease the theoretical 60Fe abundance in molecular clouds by assuming that stars more massive than 30 M\u2299 do not explode as SNe. These stars would not contribute to the 60Fe inventory, bringing the theoretical ISM 60Fe\/26Al ratio closer to that of the solar system. Because stars more massive than 30 M\u2299 contribute to less than 50% of the 60Fe production2, this proposition cannot resolve the two orders of magnitude discrepancy between the model and observations."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[962,974]],"Functions Start End":[[887,961],[975,1452]]}
{"Identifier":"2018ApJ...867...90N__Hobbs_et_al._2006_Instance_1","Paragraph":"We constructed the gamma-ray folded light curve of 3FGL J2039.6\u20135618 using the orbital period from J. Strader et al. (2018, in preparation) and the two best periods in our searches using the full data and the data before MJD 57,040. The phase zero is defined to be the epoch of the optical maximum (T0 = 57,603.95787 in MJD). In order to maximize the signal-to-noise ratio, we use the method of photon weighting (Kerr 2011) instead of an aperture selection to produce the folded light curve. The probability that the photon is emitted by 3FGL J2039.6\u20135618 is assigned to each of the photons using gtsrcprob in the Science Tools. The orbital phase is calculated using the TEMPO2 package (Hobbs et al. 2006) with the Fermi pulg-in (Ray et al. 2011). In the timing model, we adopted the main-sequence\/pulsar binary model (MSS; Wex 1998). Using the photon probabilities as the weights, the resulting weighted light curves are shown in Figure 4. From top to bottom, the orbital periods used in folding the light curves are PStrader = 0.22798177 days, Pfull = 0.2279757 days, and Pbefore = 0.2279808 days, respectively. Two orbital periods are shown for clarity. The light curve indicates that the modulation is a single peak. The FWHM of the peak spans from about \u03d5 = 0.25 to \u03d5 = 0.7, which occupies almost half of the orbital period. Although the best period found from the Rayleigh test has different values with different time spans, the folded light curves in Figure 4 generally show similar signal structures. As we speculate that the gamma-ray modulation disappears after MJD 57,040, it is not reliable to use the best period Pfull obtained from the full data set, which includes the no-signal duration. On the other hand, the best period Pbefore is obtained from the data only containing the portion before MJD 57,040; therefore, it may be biased to be applied to the study of the full data set. Therefore, with a negligible difference in the light curves and modulation significances (\n\n\n\n\n\n vs. \n\n\n\n\n\n for data before MJD 57,040), we adhere to the orbital period PB = PStrader = 0.2279817(7) days, which is reported from the independent optical observation by J. Strader et al. (2018, in preparation), for the rest of this study. Figure 5 shows the energy-dependent orbital light curves of 3FGL J2039.6\u20135618 in the energy ranges of 0.1\u2013500 GeV (top), 0.1\u20133 GeV (middle), and 3\u2013500 GeV (bottom). It is clear that the orbital modulation is mostly contributed by the lower part of the energy.","Citation Text":["Hobbs et al. 2006"],"Functions Text":["The orbital phase is calculated using the TEMPO2 package"],"Functions Label":["Uses"],"Citation Start End":[[687,704]],"Functions Start End":[[629,685]]}
{"Identifier":"2018AandA...611A..74R__Grady_et_al._2013_Instance_2","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"],"Functions Text":["Although a cavity of 55 astronomical units (au) in radius has been inferred from dust millimeter emission","infrared polarized intensity observations have found no clear evidence for a cavity in scattered light"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1373,1390]],"Functions Start End":[[1140,1245],[1269,1371]]}
{"Identifier":"2021AandA...653A..36M__Goulding_&_Alexander_(2009)_Instance_5","Paragraph":"The SFG sample was constructed using the Great Observatories All-Sky LIRG Survey (GOALS sample, Armus et al. 2009), from which we extracted 158 galaxies, with data from Inami et al. (2013), who report the fine-structure lines at high resolution in the 10\u2005\u2212\u200536\u2006\u03bcm interval, and Stierwalt et al. (2014), who include the detections of the H2 molecular lines and the PAH features at low spectral resolution. For those galaxies in the GOALS sample that have a single IRAS counterpart, but more than one source detected in the emission lines, we have added together the line or feature fluxes of all components, to consistently associate the correct line or feature emission to the total IR luminosity computed from the IRAS fluxes. To also cover lower luminosity galaxies, as the GOALS sample only includes luminous IR galaxies (LIRGs) and ultra-luminous IR galaxies (ULIRGs), we included 38 galaxies from Bernard-Salas et al. (2009) and Goulding & Alexander (2009), to reach the total sample of 196 galaxies with IR line fluxes in the 5.5\u2005\u2212\u200535\u2006\u03bcm interval in which an AGN component is not detected. For the Bernard-Salas et al. (2009), Goulding & Alexander (2009), and the GOALS samples, we excluded all the composite starburst-AGN objects identified as those with a detection of [NeV] either at 14.3 or 24.3 \u03bcm. It is worth noting that the original samples from Goulding & Alexander (2009) and Bernard-Salas et al. (2009) have spectra solely covering the central region of the galaxies. To estimate the global SFR, we corrected the published line fluxes of the Spitzer spectra by multiplying them by the ratio of the continuum reported in the IRAS point source catalogue to the continuum measured on the Spitzer spectra extracted from the CASSIS database (Lebouteiller et al. 2015). We assumed here that the line emission scales (at first order) with the IR brightness distribution. In particular, we considered the continuum at 12\u2006\u03bcm for the [NeII]12.8\u2006\u03bcm and [NeIII]15.6\u2006\u03bcm lines, and the continuum at 25 \u03bcm for the [OIV]25.9\u2006\u03bcm, [FeII]26\u2006\u03bcm, [SIII]33.5\u2006\u03bcm, and [SiII]34.8\u2006\u03bcm lines. This correction was not needed for the AGN sample and the GOALS sample because of the greater average redshift of the galaxies in the 12MGS and GOALS samples. In particular, the 12MGS active galaxy sample has a mean redshift of 0.028 (Rush et al. 1993), while the GOALS sample has a mean redshift of 0.026. The galaxies presented by Bernard-Salas et al. (2009) have instead an average redshift of 0.0074, while the sample by Goulding & Alexander (2009) has an average redshift of 0.0044. For the other lines in the 10\u2005\u2212\u200536\u2006\u03bcm interval, Goulding & Alexander (2009) did not report a detection, and we used the data presented in Bernard-Salas et al. (2009) for a total of 15 objects. Both Bernard-Salas et al. (2009) and Goulding & Alexander (2009) reported data from the high-resolution Spitzer-IRS spectra. Data in the 50\u2005\u2212\u2005205\u2006\u03bcm interval were taken from D\u00edaz-Santos et al. (2017). For the GOALS sample, 20 starburst galaxies were taken from Fern\u00e1ndez-Ontiveros et al. (2016), and 23 objects were taken from the ISO-LWS observations of Negishi et al. (2001). As a result, we obtained a total sample of 193 objects. Lastly, the PAH features\u2019 fluxes were measured from the low-resolution Spitzer-IRS spectra by Brandl et al. (2006), including 12 objects from the sample of Bernard-Salas et al. (2009) and 179 objects from Stierwalt et al. (2014).","Citation Text":["Goulding & Alexander (2009)"],"Functions Text":["For the other lines in the 10\u2005\u2212\u200536\u2006\u03bcm interval,","did not report a detection, and we used the data presented in Bernard-Salas et al. (2009) for a total of 15 objects."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[2618,2645]],"Functions Start End":[[2570,2617],[2646,2762]]}
{"Identifier":"2022ApJ...924...56S__Madau_&_Fragos_2017_Instance_1","Paragraph":"The formation and evolution of black holes (BHs) in the universe is one of the major issues to be addressed by the modern research in astrophysics and cosmology. In the mass range m\n\u2022 \u223c 5\u2013150 M\n\u2299, BHs are originated from the final, often dramatic stages in the evolution of massive stars (possibly hosted in binary systems). These compact remnants can produce luminous X-ray binaries (e.g., Mapelli et al. 2010; Farr et al. 2011; Inoue et al. 2016), can constitute powerful sources of gravitational waves for ground-based detectors like the current LIGO\/Virgo facility (e.g., Belczynski et al. 2010; Dominik et al. 2015; Spera & Mapelli 2017; Boco et al. 2019; Spera et al. 2019; Abbott et al. 2021a, 2021b), can possibly energize short gamma-ray bursts and associated kilonovas (e.g., Abbott et al. 2020, 2021c; Ackley et al. 2020; Gompertz et al. 2020), can inject strong energy inputs in the primeval universe (e.g., Mirabel et al. 2011; Justham & Schawinski 2012; Artale et al. 2015; Madau & Fragos 2017; Lehmer et al. 2021), and can provide light seeds for the subsequent growth of more massive BHs (e.g., Madau et al. 2014; Volonteri et al. 2015; Lupi et al. 2016; Pacucci et al. 2017; Boco et al. 2020; Das et al. 2021). At the other end, in the range M\n\u2022 \u223c 106\u20131010\nM\n\u2299, supermassive BHs grow mainly by gaseous accretion that energizes the spectacular broadband emission of active galactic nuclei (AGNs). Such an activity can have a profound impact on galaxy evolution (e.g., Alexander & Hickox 2012; Lapi et al. 2014, 2018), as testified by the strict relationships between the relic BH masses and the physical properties of the hosts (e.g., Kormendy & Ho 2013; Shankar et al. 2016, 2020; Zhu et al. 2021). The intermediate-mass range m\n\u2022 \u223c 103\u2013106\nM\n\u2299 is the most uncertain. So far, only tentative evidence of these systems has been identified (see Paynter et al. 2021). However, the chase is open in view of their astrophysical relevance. Most noticeably, they can provide heavy seeds for quick (super)massive BHs growth (e.g., Mayer & Bonoli 2019; Boco et al. 2020), as it seems required by the puzzling observations of an increasing numbers of giant monsters M\n\u2022 \u2273 109\nM\n\u2299 when the age of the universe was shorter than \u22720.8 Gyr (e.g., Mortlock et al. 2011; Venemans et al. 2017; Banados et al. 2018). Moreover, such intermediate-mass BHs will constitute important targets for space-based gravitational wave detectors like LISA and DECIGO (see eLisa Consortium 2013; Kawamura et al. 2021; also Barausse & Lapi 2021; Boco et al. 2021b).","Citation Text":["Madau & Fragos 2017"],"Functions Text":["These compact remnants","can inject strong energy inputs in the primeval universe (e.g.,"],"Functions Label":["Background","Background"],"Citation Start End":[[988,1007]],"Functions Start End":[[325,347],[856,919]]}
{"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)"],"Functions Text":["Five clusters from the sample of","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."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1210,1231]],"Functions Start End":[[1177,1209],[1232,1435]]}
{"Identifier":"2015ApJ...815....7V__Davila_1987_Instance_1","Paragraph":"There are also simplified problems, less complex than fully developed turbulence, in which one finds the formation of small scales in the direction perpendicular to an applied magnetic field \n\n\n\n\n\n It is well known that this effect appears in the context of MHD when imposed parallel-propagating waves interact with an inhomogeneous background consisting either of pressure-balanced structures or velocity shears (Ghosh et al. 1998). Phase mixing of torsional Alfv\u00e9n waves in axisymmetric-equilibrium magnetic configurations have been studied in Ruderman et al. (1999), where possible applications to the solar corona and solar wind have been proposed. In two-dimensional (2D) equilibria, where the Alfv\u00e9n velocity varies in directions perpendicular to the magnetic field, two mechanisms have been investigated in detail: (1) phase mixing (Heyvaerts & Priest 1983), in which differences in group velocity at different locations progressively bend wave fronts, and (2) resonant absorption, which concentrates the wave energy in a narrow layer where the wave frequency locally matches a characteristic frequency (Alfv\u00e9n or cusp). These processes have been studied both by investigating normal modes of the inhomogeneous structure (Kappraff & Tataronis 1977; Mok & Einaudi 1985; Steinolfson 1985; Davila 1987; Hollweg 1987; Califano et al. 1990, 1992) and by considering the evolution of an initial disturbance (Lee & Roberts 1986; Malara et al. 1992, 1996). Effects of density stratification and magnetic line divergence (Ruderman et al. 1998), nonlinear coupling with compressive modes (Nakariakov et al. 1997, 1998), and evolution of localized pulses (Tsiklauri & Nakariakov 2002; Tsiklauri et al. 2003) have been considered. The propagation of MHD waves in inhomogeneous magnetic fields containing null points has also been studied in detail (Landi et al. 2005; see also McLaughlin et al. 2010 for a review). Phase mixing in 3D inhomogeneous equilibria has also been considered in the small-wavelength limit (Similon & Sudan 1989) using a WKB approximation (Petkaki et al. 1998; Malara et al. 2000), also within the problem of coronal heating (Malara et al. 2003, 2005, 2007). Particle acceleration in phase mixing of Alfv\u00e9n waves in a dispersive regime has been studied by Tsiklauri et al. (2005) using particle-in-cell simulations, both in 2D (Tsiklauri 2011) and in 3D (Tsiklauri 2012) configurations. Finally, instabilities generating KAWs in a plasma with transverse density modulations have been considered by Wu & Chen (2013). Similar ideas involving dissipative mechanisms related to the interaction of Alfv\u00e9n waves or KAWs and phase mixing have been examined in the context of the magnetospheric plasma sheet (Lysak & Song 2011) and in coronal loops (Ofman & Aschwanden 2002).","Citation Text":["Davila 1987"],"Functions Text":["These processes have been studied both by investigating normal modes of the inhomogeneous structure"],"Functions Label":["Background"],"Citation Start End":[[1294,1305]],"Functions Start End":[[1128,1227]]}
{"Identifier":"2018MNRAS.480.1174H__Jehin_et_al._2009_Instance_1","Paragraph":"The origin of nitrogen in the Solar system is still an open question. More specifically, the main repository of nitrogen in the protosolar nebula (PSN) is still unclear, although there is some consensus that it may be atomic, N, or molecular, ${\\rm N_2}$ (Schwarz & Bergin 2014). Furthermore, the large variations of the isotopic ratio of nitrogen (${\\rm ^{14}N}\/{\\rm ^{15}N}$), as measured in various carriers within different types of Solar system objects, remain unexplained (Al\u00e9on 2010; Hily-Blant et al. 2013a, 2017; F\u00fcri & Marty 2015). One striking problem is the ${\\rm ^{14}N}\/{\\rm ^{15}N}$ isotopic ratio of nitrogen in comets. Its average value, 144 \u00b1 3 (Jehin et al. 2009; Bockel\u00e9e-Morvan et al. 2015; Shinnaka et al. 2016; Hily-Blant et al. 2017), is three times lower than the bulk ratio of 441 \u00b1 6 in the protosun as inferred from solar wind measurements (Marty et al. 2011). The reasons for these different ratios remain elusive, casting doubts on our understanding of the origin of the composition of comets and more generally of the origin of nitrogen in the Solar system. Several possibilities (not mutually exclusive) could explain the discrepancy: (i) the tracers of nitrogen observed so far in comets \u2013 HCN, CN, and ${\\rm NH_2}$ \u2013 are minor reservoirs of cometary nitrogen and thus naturally do not reflect the bulk ratio in the PSN, (ii) efficient fractionation processes in the PSN at the time of comet formation, (iii) efficient fractionation processes in the parent interstellar cloud, and (iv) exchange processes within cometary ices since their formation. Recently, it was shown that protoplanetary discs \u2013 or equivalently PSN analogues \u2013 carry at least two isotopic reservoirs of nitrogen, traced, respectively, by CN and HCN, with HCN probing a secondary, fractionated, reservoir (Hily-Blant et al. 2017). Furthermore, the isotopic reservoirs traced by HCN and CN are found to be in a 1:3 ratio, respectively, reminiscent of the factor of 3 between the cometary and bulk isotopic ratios (144:441) in the PSN. It follows that exchange processes in parent bodies [possibily (iv) above] are not necessary. The PSN hypothesis is supported by models of selective photodissociation of ${\\rm N_2}$ in protoplanetary discs (Heays et al. 2014) that predict a strong enrichment of HCN in ${\\rm ^{15}N}$, but also of CN, in contrast with observations (Hily-Blant et al. 2017). At present, clear-cut observational evidences supporting the PSN or interstellar hypothesis are still lacking.","Citation Text":["Jehin et al. 2009"],"Functions Text":["One striking problem is the ${\\rm ^{14}N}\/{\\rm ^{15}N}$ isotopic ratio of nitrogen in comets. Its average value, 144 \u00b1 3",", is three times lower than the bulk ratio of 441 \u00b1 6 in the protosun as inferred from solar wind measurements","The reasons for these different ratios remain elusive, casting doubts on our understanding of the origin of the composition of comets and more generally of the origin of nitrogen in the Solar system."],"Functions Label":["Motivation","Motivation","Motivation"],"Citation Start End":[[664,681]],"Functions Start End":[[542,662],[757,867],[889,1088]]}
{"Identifier":"2019AandA...630A..37S__Behar_et_al._2017_Instance_3","Paragraph":"Solar wind velocity distribution moments are described in Behar et al. (2017). The ion density nsw is the moment of order 0, and the ion bulk velocity usw (a vector) appears in the moment of order 1, the flux density \n\n$n_{\\mathrm{sw}} \\ \\underline{\\mathbf{u}_{\\mathrm{sw}}}$\n\n\n\nn\n\nsw\n\n\u2009\n\n\nu\n\nsw\n\n\n_\n\n\n\n. The bulk speed can be defined as the norm of the bulk velocity, that is, \n\n$u_{\\mathrm{sw}} = |\\underline{\\vec{u}_{\\mathrm{sw}}}$\n\n\n\nu\n\nsw\n\n=|\n\n\nu\n\nsw\n\n\n_\n\n\n\n|. However, this bulk speed is representative of single-particle speeds as long as the velocity distribution function is compact (e.g., a Maxwellian distribution). Complex velocity distribution functions were observed by RPC-ICA within the atmosphere of 67P. For instance, partial ring distributions were frequently observed for solar wind protons at intermediate heliocentric distances, when the spacecraft approached the SWIC (Behar et al. 2017). To illustrate the effect of such distorted distributions, a perfect ring (or shell) distribution centered on the origin of the plasma reference frame can be imagined, in which all particles have the same speed of 400 km s\u22121. The norm of the bulk velocity in this case would be 0 km s\u22121, whereas the mean speed of the particles is 400 km s\u22121, which is the relevant speed for SWCX processes. This mean speed, noted Usw, of the particles is calculated by first summing the differential number flux over all angles, and then taking the statistical average (Behar 2018). Over the entire mission, the deceleration of the solar wind using the mean speed of the particles is much more limited than the deceleration shown by the norm of the bulk velocity (Behar et al. 2017): there is more kinetic energy in the solar wind than the bulk velocity vector would let us think. This is the main difference with the paradigm used at previously studied (and more active) comets (Behar et al. 2018b). These complex, nonthermal velocity distribution functions also prevent us from reducing the second-order moment (the stress tensor) to a single scalar value, which, for a Maxwellian distribution, could be identified with a plasma temperature. In the context of 67P and for an important part of the cometary orbit around the Sun, the temperature of the solar wind proton has no formal definition.","Citation Text":["Behar et al. 2017"],"Functions Text":["Over the entire mission, the deceleration of the solar wind using the mean speed of the particles is much more limited than the deceleration shown by the norm of the bulk velocity",": there is more kinetic energy in the solar wind than the bulk velocity vector would let us think."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1659,1676]],"Functions Start End":[[1478,1657],[1677,1775]]}
{"Identifier":"2015ApJ...814...84D__Ellison_et_al._2008_Instance_1","Paragraph":"We first need to discuss that the environmental dependence of the metallicity of star-forming galaxies may depend on the scale at which the environment is defined. At small scales (a few tens of kpc), there is substantial evidence from observations and simulations for a decrement of metallicity and an enhancement of star-formation activity in galaxy close pairs, merging, and interacting systems compared to isolated, field galaxies, mostly attributed to the interaction-induced inflow of metal-poor gas from the periphery of interacting galaxies to the center, diluting their metal content, and increasing their gas fuel for star-formation (Mihos & Hernquist 1996; Kewley et al. 2006; Ellison et al. 2008, 2013; Michel-Dansac et al. 2008; Rupke et al. 2010; Sol Alonso et al. 2010; Perez et al. 2011; Scudder et al. 2012b; Ly et al. 2014). At intermediate group scales where galaxy interactions are more common compared to cluster and field environments (Perez et al. 2009; Tonnesen & Cen 2012)\u2014owing to a combination of (1) a lower velocity dispersion of group galaxies relative to their cluster counterparts and (2) a higher number density of group galaxies compared to the field systems, which provide an ideal condition for interactions)\u2014there is also evidence for a deficit of metals in group galaxies compared to control samples in the field (Lara-L\u00f3pez et al. 2013b), possibly owing to a higher fraction of interacting galaxies. At larger filamentary and cluster scales, the slight metal enhancement of galaxies relative to the field might be due to (1) the inflow of already-enriched interafilamentary or interacluster gas into galaxies, as observations and simulations have shown a more metal-enriched intergalactic medium in cluster and filament environments compared to the field (Arnaud et al. 1994; Aracil et al. 2006; Stocke et al. 2006, 2007; Sato et al. 2007; Cen & Chisari 2011; Dav\u00e9 et al. 2011; Oppenheimer et al. 2012); (2) the environmental strangulation (Larson et al. 1980; Peng et al. 2015) of low-metallicity diluting gas falling from the surrounding LSS cosmic web into galaxies; (3) the environmental ram pressure stripping (Gunn & Gott 1972; Abadi et al. 1999) of the metal-poor diluting gas in the periphery of galaxies; and (4) trapping and recycling of metal-enriched outflows due to the hotter environment of filaments and clusters (Aracil et al. 2006; Cen & Ostriker 2006; Werner et al. 2008) compared to the field.","Citation Text":["Ellison et al. 2008"],"Functions Text":["At small scales (a few tens of kpc), there is substantial evidence from observations and simulations for a decrement of metallicity and an enhancement of star-formation activity in galaxy close pairs, merging, and interacting systems compared to isolated, field galaxies, mostly attributed to the interaction-induced inflow of metal-poor gas from the periphery of interacting galaxies to the center, diluting their metal content, and increasing their gas fuel for star-formation"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[688,707]],"Functions Start End":[[164,642]]}
{"Identifier":"2021AandA...654A..88W__Cai_et_al._2017_Instance_1","Paragraph":"This paper focuses on the population of high-redshift radio galaxies (HzRGs; L500\u2006MHz\u2004>\u20041026 W Hz\u22121Miley & De Breuck 2008), which are some of the most massive galaxies known at any redshift (with a narrow range in stellar masses of (1\u2005\u2212\u20056)\u00d71011\u2006M\u2299 for 1\u2004\u2004z\u2004\u20045.2; De Breuck et al. 2010). Their energetic radio jets are unique markers of concomitant powerful AGN activity, which place them amongst the most active sources at and near cosmic noon. High-redshift radio galaxies have furthermore been shown to be powerful beacons of dense (proto-)cluster environments in the early Universe (e.g., Le Fevre et al. 1996; Stern et al. 2003; Venemans et al. 2002, 2003, 2004, 2005, 2007; Wylezalek et al. 2013). The quasar-level AGN activity (Miley & De Breuck 2008) at the center is blocked by the thick dusty torus acting as the \u201ccoronograph\u201d (Vernet et al. 2001); this makes HzRGs true obscured type-2 quasars, allowing us to probe their host galaxies and CGM without strong AGN contamination (e.g., for unobscured quasars, see Arrigoni Battaia et al. 2019, and for radio-quite type-2 sources, see Cai et al. 2017). Comprehensive studies using near-infrared integral field unit (IFU) instruments show that the ionized gas in HzRGs is highly perturbed (FWHM\u2004\u223c\u20041000 km s\u22121) at kiloparsec scales and is aligned with the radio jets (Nesvadba et al. 2006, 2007, 2008, 2017a,b; Collet et al. 2015, 2016). This implies that the energy and momentum transfer between the central quasar and their ISM is likely due to the jets. Radio-mode feedback may therefore play a fundamental role during the evolution of HzRGs. Recently, Falkendal et al. (2019) combined infrared and millimeter data and deduced a more robust result of a relatively low star formation rate (SFR) for a sample of HzRGs, suggesting evidence of rapid quenching compared to previous studies (e.g., Drouart et al. 2014). Using a small sample of HzRGs, Nesvadba et al. (2011) shows that they are going through a transition phase from active to passive. These observations indicate that HzRGs are on a different track of evolution compared to radio-quiet objects, assembling most of their stellar mass early (z\u2004\u223c\u20043; Seymour et al. 2007; De Breuck et al. 2010), and that radio jets may actively affect their quenching. However, there is also circumstantial evidence showing that the jet can induce star formation. Humphrey et al. (2006) found that HzRGs (z\u2004>\u20042 in the sample) with smaller radio sources and more perturbed gas (emission line) kinematics show lower UV continuum polarization, which could be due to the presence of more luminous young stellar populations and can possibly be explained by the interaction between radio jets and the ISM that enhances star formation. Besides, there is also an anticorrelation between the rest frame submillimeter flux density and radio size in HzRGs (Humphrey et al. 2011), although it is not clear if the physics behind this is feedback-induced star formation, a simultaneous triggering of star formation and the radio-loud AGN activity, or simply environmental effects. Some well-studied HzRGs show evidence of having high SFRs (e.g., 4C41.47 and PKS 0529\u2212549; Nesvadba et al. 2020; Falkendal et al. 2019). In these sources, we may interestingly be witnessing both the jets compressing the gas, leading to enhanced SFRs (e.g., Fragile et al. 2017), and the feedback from the AGN and star formation quenching it (Man et al. 2019).","Citation Text":["Cai et al. 2017"],"Functions Text":["The quasar-level AGN activity","at the center is blocked by the thick dusty torus acting as the \u201ccoronograph\u201d","this makes HzRGs true obscured type-2 quasars, allowing us to probe their host galaxies and CGM without strong AGN contamination (e.g.","for radio-quite type-2 sources, see"],"Functions Label":["Motivation","Motivation","Motivation","Motivation"],"Citation Start End":[[1092,1107]],"Functions Start End":[[703,732],[758,835],[858,992],[1056,1091]]}
{"Identifier":"2018ApJ...854..120M__Baxter_et_al._2016_Instance_1","Paragraph":"The main purpose of this paper is to calibrate the mass\u2013observable relation from a joint measurement of the abundance (number counts) and the stacked cluster weak lensing profiles. We develop and apply our method to the Sloan Digital Sky Survey (SDSS) red-sequence Matched-filter Probabilistic Percolation (redMaPPer) cluster catalog that is constructed by identifying overdensities of red sequence galaxies with similar colors to galaxy clusters from the SDSS ugriz photometries (Rozo & Rykoff 2014; Rykoff et al. 2014; Rozo et al. 2015a, 2015b; see most recently Rykoff et al. 2016, for the details of the method). Since the cluster finder gives an estimation of the optical richness, \u03bb, for each cluster, we will constrain the scaling relation between the optical richness and mass for the clusters. In this paper we develop a forward modeling approach, where we constrain the probability distribution of richness for a given halo mass, \n\n\n\n\n\n. This is in contrast with previous studies (Baxter et al. 2016; Jimeno et al. 2017; Melchior et al. 2017; Simet et al. 2017), where the backward modeling approach is employed to constrain the probability of mass for a given richness, \n\n\n\n\n\n. The forward modeling approach has several advantages. First, we can use the abundance measurements more easily to constrain the mass\u2013observable relation, as Saro et al. (2015) constrained \n\n\n\n\n\n from the abundance measurements of SZ-selected clusters after matching to redMaPPer clusters. Second, \n\n\n\n\n\n can be inferred from \n\n\n\n\n\n once the halo mass function \n\n\n\n\n\n is given, based on the Bayes theorem, while the opposite transformation, i.e., inferences of \n\n\n\n\n\n from \n\n\n\n\n\n, is not straightforward, because this requires knowledge of the richness function \n\n\n\n\n\n over the whole range of \u03bb, which is not generally available, or is at least very noisy (and possibly affected by contamination), for richness below a threshold richness in cluster catalogs. Third, the forward modeling is convenient to generate mock catalogs of clusters by populating halos in N-body simulations with galaxies, e.g., to test systematics in a cluster-finding algorithm.","Citation Text":["Baxter et al. 2016"],"Functions Text":["This is in contrast with previous studies","where the backward modeling approach is employed to constrain the probability of mass for a given richness,"],"Functions Label":["Differences","Differences"],"Citation Start End":[[991,1009]],"Functions Start End":[[948,989],[1073,1180]]}
{"Identifier":"2022AandA...659A.180G__Kutsenko_et_al._2018_Instance_1","Paragraph":"In the last few decades, the dynamic properties of the quiet Sun have been thoroughly investigated using a range of substantially different techniques, allowing us to elaborate a consistent picture of the photospheric dynamics by approaching the problem from different points of view. Particularly interesting and promising are the studies involving the tracking of small-scale magnetic fields in the quiet photosphere. Such investigations reveal features that still cannot be captured by theoretical models and\/or simulations because of the complexity of the system and the simultaneous coupling of a wide range of spatial and temporal scales. These studies are based on the hypothesis that magnetic fields are passively transported by the plasma flow and provide a characterisation of advection and diffusion processes in the quiet photosphere from granular to supergranular scales (see, e.g. Wang 1988; Berger et al. 1998; Cadavid et al. 1998, 1999; Hagenaar et al. 1999; Lawrence et al. 2001; S\u00e1nchez Almeida et al. 2010; Abramenko et al. 2011; Manso Sainz et al. 2011; Lepreti et al. 2012; Giannattasio et al. 2013, 2014a,b; Giannattasio et al. 2019; Keys et al. 2014; Caroli et al. 2015; Del Moro et al. 2015; Yang et al. 2015a,b; Roudier et al. 2016; Abramenko 2017; Kutsenko et al. 2018; Agrawal et al. 2018; Giannattasio & Consolini 2021). These studies reveal an anomalous scaling of magnetic field transport with a superdiffusive character consistent with a Levy walk inside supergranules and a more Brownian-like motion in their boundary. This has provided constraints on the magnetic flux emergence and evolution that models have to consider in order to fully explain the dynamics governing these environments. At the same time, the scaling laws affecting magnetic fields in the quiet Sun are crucial to understanding how such fields vary with scale size, despite the small scales at which dissipation occurs still being inaccessible with the currently available observations (see, e.g. Lawrence et al. 1994; Stenflo 2012, and references therein). For example, in the milestone work by Stenflo (2012), the spectrum of magnetic flux density in the quiet Sun was found to be consistent with a Kolmogorov power-law scaling. The scale at which scale invariance is broken lies below the current resolution limit. This latter author argued that the collapse of magnetic fields in Kilogauss flux tubes injects energy that is expected to cascade down because of the flux decay occurring via interchange instability, and to fragment into weaker \u2018hidden\u2019 fields at smaller scales (down to \u223c10 km). As far as we know, no other studies have focused on the scaling properties characterising the magnetic fields in a quiet Sun region within a range of spatial and temporal scales from (sub)granular to supergranular in the time domain. In this work, for the first time we apply the structure function analysis typical of complex systems (Frisch 1995) to fill this gap. This approach complements the studies based on feature tracking mentioned above. The main difference is that, while in those works statistical properties of the photospheric plasma flows are investigated via the transport of small-scale magnetic fields in a frozen-in condition, here we directly study the magnetic field variations emerging from magnetogram time-series. The paper is organised as follows. Section 2 describes the data set used and the analysis techniques applied. Section 3 describes the obtained results, while Sect. 4 is devoted to their discussion in the light of current literature. Finally, in Sect. 5, we present our conclusions and present future perspectives.","Citation Text":["Kutsenko et al. 2018"],"Functions Text":["These studies are based on the hypothesis that magnetic fields are passively transported by the plasma flow and provide a characterisation of advection and diffusion processes in the quiet photosphere from granular to supergranular scales (see, e.g","These studies reveal an anomalous scaling of magnetic field transport with a superdiffusive character consistent with a Levy walk inside supergranules and a more Brownian-like motion in their boundary."],"Functions Label":["Background","Background"],"Citation Start End":[[1274,1294]],"Functions Start End":[[645,893],[1349,1550]]}
{"Identifier":"2022ApJ...926...21B__Viviani_et_al._2018_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":["Viviani et al. 2018"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[1489,1508]],"Functions Start End":[[1001,1247]]}
{"Identifier":"2018ApJ...869..143P__Wills_&_Browne_1986_Instance_1","Paragraph":"In order to understand this discrepancy, we consider the BL Lac object nature of NGC 1275 (Veron 1978). The BL Lac object properties of 3C 84 have been noted in connection with the synchrotron optical emission that hides the accretion generated continuum. The BL Lac object aspect can be very pronounced with optical polarization that changes dramatically in amplitude and position angle, with the largest optical polarization reaching 6% (Angel & Stockman 1980). Such extreme blazar-like properties suggest a small LOS to the jet axis (Lind & Blandford 1985). We qualify this statement by noting that the evidence does not support extreme BL Lac object behavior. The range of polarization is 1%\u20136% in Angel & Stockman (1980) and more extensive polarization data9\n\n9\nFrom http:\/\/www.bu.edu\/blazars\/VLBAproject.html.\n covering the time period 2011\u20132018 indicates an optical polarization that is usually below 2% and extremely rare (2) instances of polarization >3% were reported. This suggests a slightly off angle BL Lac object. The H\u03b2 BEL is considered to be rotating gas in a flat \u201cpancake-like\u201d region in which the normal to the BEL disk is parallel to the jet axis (Wills & Browne 1986). Thus, the FWHM that appears in Equation (4) depends on the LOS. For polar lines of sight, Equation (4) will under estimate Mbh. In the virial mass estimation,\n5\n\n\n\n\n\nwhere G is the gravitational constant, RBLR is the orbital radius of the BLR (broad line region) and vBLR is the velocity of the BEL gas. In order to relate vBLR to an observed quantity, one defines the de-projection factor, f,\n6\n\n\n\n\n\nEquations (5) and (6) indicate that \n\n\n\n\n\n. The de-projection factor has been estimated for various classes of objects that are believed to be differentiated by an LOS (Antonucci 1993; DeCarli et al. 2011). The method of DeCarli et al. (2011) was to estimate Mbh from the bulge luminosity of the host galaxy. Using this estimate to set the value of Mbh in the virial formula, they were able to estimate f for various classes of objects,\n7\n\n\n\n\n\nEquation (4) is derived based on assuming an isotropic distribution of BEL gas velocity. Thus, we adopt a correction factor for the estimate in Equation (4) of \n\n\n\n\n\n for BL Lac object orientations such as the one that exists in 3C 84. Taking the nominal value of DeCarli et al. (2011) in Equation (7), we expect a correction factor of 64. The orientation corrected central black hole mass estimate based on the data in Table 1 yields \n\n\n\n\n\n. Alternatively, if we use the blazar correction of 42 associated with the nominal value of fblazars in Equation (7) instead, we get \n\n\n\n\n\n.","Citation Text":["Wills & Browne 1986"],"Functions Text":["The H\u03b2 BEL is considered to be rotating gas in a flat \u201cpancake-like\u201d region in which the normal to the BEL disk is parallel to the jet axis","Thus, the FWHM that appears in Equation (4) depends on the LOS."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1170,1189]],"Functions Start End":[[1029,1168],[1192,1255]]}
{"Identifier":"2021ApJ...908...95H__Harrington_et_al._2016_Instance_1","Paragraph":"Strong gravitational lensing of high-z star-forming galaxies offers a unique way to examine highly magnified molecular gas. The method for selecting strongly lensed dusty galaxy candidates, at z > 1, is primarily based on unusually bright (sub)millimeter fluxes compared to the expected steep drop-off in (sub)millimeter number counts (e.g., Negrello et al. 2007, 2010). This method has since identified a large number across the extragalactic sky, i.e., more than 100 lensed candidates at z > 1 (Ivison et al. 2010; Vieira et al. 2010, 2013; Bussmann et al. 2013, 2015; Wardlow et al. 2013; Wei\u00df et al 2013; Ca\u00f1ameras et al. 2015; Harrington et al. 2016; Strandet et al. 2016; D\u00edaz-S\u00e1nchez et al. 2017; Negrello et al. 2017; Bakx et al. 2018). The lensed population of dusty star-forming galaxies selected by the South Pole Telescope (SPT), Herschel Space Observatory, and Planck has now been detected in more than two CO transitions (e.g., Spilker et al. 2016; Strandet et al. 2017; Yang et al. 2017; Bakx et al. 2020; this work). The Herschel-selected, strongly lensed galaxy sample (Bussmann et al. 2013) offered the first systematic approach to producing a statistically significant sample of CO\/[C i] lines (Yang et al. 2017), followed by a compilation in 11 Planck- and Herschel-selected lensed galaxies (Ca\u00f1ameras et al. 2018b), including four galaxies with both [C i] lines detected (Nesvadba et al. 2019). The IR-to-CO luminosity relations of local starbursts and high-z star-forming galaxies explored by Greve et al. (2014) indicate that the ISM radiation field is an important component to consider when understanding CO line excitation, yet this investigation was limited to 23 unlensed and 21 lensed dusty star-forming systems\u2014all with more than three frequency measurements of the dust continuum and usually a single CO line detection. Most previous studies used only single- and\/or double-component gas-emitting regions to reproduce the observed CO emission, excluding the simultaneous modeling of the available [C i] emission, but also ignoring the role of the dust continuum emission as a heating source of the gas.","Citation Text":["Harrington et al. 2016"],"Functions Text":["This method has since identified a large number across the extragalactic sky, i.e., more than 100 lensed candidates at z > 1"],"Functions Label":["Background"],"Citation Start End":[[632,654]],"Functions Start End":[[371,495]]}
{"Identifier":"2019ApJ...880...92J__Yang_et_al._2013_Instance_1","Paragraph":"Studies of molecules play a prominent role in explaining the physical, chemical, and kinematic properties of the interstellar medium (ISM) in galaxies (Omont 2007; Tielens 2013). One such molecule is H2O, the third most abundant molecule in the warm dense ISM after H2 and CO (Neufeld et al. 1995). As an asymmetric rotor with a large electric dipole moment, H2O has a rich and complex spectrum giving rise to emission and absorption lines mainly in the submillimeter and far-infrared (FIR) regimes of the electromagnetic spectrum. Observations from local galaxies (van der Werf et al. 2010; Wei\u00df et al. 2010; Rangwala et al. 2011; Yang et al. 2013), high-redshift ultra-luminous infrared galaxies (ULIRGs; Omont et al. 2013; Yang et al. 2016), and active galactic nuclei (AGNs; van der Werf et al. 2011) have shown H2O emission to be ubiquitous with intensities as bright as CO lines. Modeling has shown that, in addition to infrared pumping where H2O is excited by FIR photons, collisions also contribute to the intensities of low-excitation transitions (e.g., Gonz\u00e1lez-Alfonso et al. 2010, 2012). This is best represented in Figure 3 from Liu et al. (2017), which shows prominent H2O lines in different ISM components. The low-excitation lines become weaker or completely disappear in the warm and hot regions (>40 K) where infrared pumping dominates over collisions. The higher excitation transitions that require strong far-infrared radiation density are mainly found in the hotter regions (100\u2013200 K) of the galaxy. The cascading emission lines, \n\n\n\n\n\n (Eup = 100.8 K,\u03bdrest = 987.927 GHz), \n\n\n\n\n\n (Eup = 137 K, \u03bdrest = 752.033 GHz), and p-H2O (22,0 \u2212 21,1) (Eup = 196 K, \u03bdrest = 1228.789 GHz) are pumped by 101 \u03bcm photons from the base 11,1 level and are primarily excited in the warm regions of the galaxy. The collisional excitation of the low-lying levels (11,1 and 20,2) in optically thin or high-density hot regions might also contribute to the emission of the \n\n\n\n\n\n line. Hence, H2O transitions probe the infrared radiation field density and physical properties of the ISM such as gas density and kinetic temperature (e.g., Wei\u00df et al. 2010; Gonz\u00e1lez-Alfonso et al. 2014; Liu et al. 2017).","Citation Text":["Yang et al. 2013"],"Functions Text":["Observations from local galaxies","have shown H2O emission to be ubiquitous with intensities as bright as CO lines."],"Functions Label":["Background","Background"],"Citation Start End":[[632,648]],"Functions Start End":[[532,564],[805,885]]}
{"Identifier":"2016ApJ...821...74J__Tokovinin_et_al._2013_Instance_1","Paragraph":"Recent theoretical work has suggested that the presence, or lack thereof, of long-period giant planets could affect the formation of such systems. Batygin & Laughlin (2015) argued that the migration of Jupiter within our own solar system might have disrupted a massive primordial inner protoplanetary disk that could have formed multiple short-period super-Earths; they predicted that systems like the Kepler short-period multiple systems should typically lack long-period giant planets. A related question is, how common are planetary systems broadly similar in architecture to our solar system, with small close-in planets and more distant giant planets? We can begin to answer these questions in the near future through the combination of searches for short-period super-Earths and data from the long-term RV programs that have been monitoring many bright FGK stars for well over a decade. Super-Earths can be found with either high-precision RV observations or space-based transit searches. Such high-precision RV surveys include those being undertaken currently with HARPS (e.g., D\u00edaz et al. 2016), HARPS-N (M15), APF (Vogt et al. 2014), and CHIRON (Tokovinin et al. 2013), and in the near future with MINERVA (Swift et al. 2015), CARMENES (Quirrenbach et al. 2014), ESPRESSO (M\u00e9gevand et al. 2014), and SPIRou (Artigau et al. 2014). The major upcoming space-based transit survey is that of TESS (Ricker et al. 2015). Long-term RV programs include the McDonald Observatory Planet Search (e.g., Endl et al. 2016), the Anglo-Australian Planet Search (e.g., Jones et al. 2010), the Lick-Carnegie Exoplanet Survey (e.g., Rowan et al. 2016), the CORALIE planet search (Marmier et al. 2013), and the planet search at ESO (e.g., Zechmeister et al. 2013). Long-period giant planets will also be found by Gaia, which will produce a huge sample of astrometrically detected planets (Perryman et al. 2014). While most of the Kepler sample is too faint to have been observed previously by long-term RV surveys (e.g., Coughlin et al. 2015), Gaia will be able to astrometrically detect long-period planets around many of these stars. Our own McDonald Observatory Planet Search program now has a baseline of 12\u201315 years for \u223c200 FGKM stars, and a handful of stars also have lower precision observations dating back more than 25 years. HD 219134 is one of these stars, and here we present an analysis of our RV observations of this star, as well as our data on the stellar activity.","Citation Text":["Tokovinin et al. 2013"],"Functions Text":["Such high-precision RV surveys include those being undertaken currently with","and CHIRON"],"Functions Label":["Background","Background"],"Citation Start End":[[1155,1176]],"Functions Start End":[[995,1071],[1143,1153]]}
{"Identifier":"2018AandA...612A..77M__Gromadzki_&_Miko\u0142ajewska_(2009)_Instance_1","Paragraph":"\u201cWiggling\u201d outflows are often observed among young stellar jets and protostellar molecular outflows (Eisloffel et al. 1996; Terquem et al. 1999). Terquem et al. (1999) investigated such binary systems where the accretion disk, from which the jet originates, is inclined to the binary orbital plane. They concluded that the observed jet \u201cwiggling\u201d is a consequence of the jet precession caused by tidal interactions in such non-coplanar binary systems. Nichols & Slavin (2009) as well as Hollis & Michalitsianos (1993) suggested that the precession of the accretion disk around the WD may be responsible for the bending of the wide-angle outflow found in the previous studies. In analogy to these observations of young stellar jets, we suggest that the \u201cwiggling\u201d that we also find here for the R Aqr jet may result from disk precession as well. We estimated the precession period using the theory from Terquem et al. (1999) and the binary parameters taken from Gromadzki & Miko\u0142ajewska (2009) \u2013 Mh = 0.8M\u2299 (the mass of the hot WD companion), Mp\u2215Mh = 1.65 (where Mp is the mass of the primary companion), D = 15.5 AU (the mean semi-major axis of the system), and e = 0.25 (eccentricity). The value of Rd is unknown in our case and we used Rd = 5 AU giving D\u2215Rd \u2248 3 which corresponds to the average value of 2 \u2264 D\u2215Rd \u2264 4 (the range taken from Terquem et al. 1999). For this calculation, we assumed that the angle \u03b4 between the disk plane and that of the binary orbit is small enough (10\u00b0) and we adopted cos\u03b4 = 1. Using Eq. (1) from Gromadzki & Miko\u0142ajewska (2009), we derived the precession time of T \u2248 530 yr. This value is quite large for the wiggling waves that we see. We estimated the projected spatial wavelength \u03bbproj of the \u201cwiggling\u201d wave according to \u03bb = \u03bbproj\u2215sin i, where i is the angle between the jet symmetry axis and the line of sight, and T = \u03bb\u2215\u03c5, where \u03c5 is the jet velocity, from Gromadzki & Miko\u0142ajewska (2009). Using i = 72\u00b0 and \u03c5 ~ 100 km\/s, we derive \u03bbproj \u2248 10 500 AU which is more than 20 times larger than the projected length of the observed wiggling outflow (2\u2032\u2032 \u2248 440 AU). However, we should note that the precessing time strongly depends on the D\u2215Rd; the T decreases significantly with increasing R. It may also be the case that the \u201cwiggling\u201d model developed for YSO jets is not fitting for R Aqr which consists of evolved objects, and both the WD and the disk where the jet probably forms are much hotter than YSO systems. Furthermore, we cannot exclude that the steady wiggling might be a sequence of dynamical interactions of the two collimated flows tilted to each other.","Citation Text":["Gromadzki & Miko\u0142ajewska (2009)"],"Functions Text":["We estimated the precession period using the theory from Terquem et al. (1999) and the binary parameters taken from","\u2013 Mh = 0.8M\u2299 (the mass of the hot WD companion), Mp\u2215Mh = 1.65 (where Mp is the mass of the primary companion), D = 15.5 AU (the mean semi-major axis of the system), and e = 0.25 (eccentricity)."],"Functions Label":["Uses","Uses"],"Citation Start End":[[961,992]],"Functions Start End":[[845,960],[993,1186]]}
{"Identifier":"2018MNRAS.479.3438G___2005_Instance_1","Paragraph":"For instance, due to its dissipative nature, gas can be very efficient in absorbing and transporting outwards the angular momentum of the pair, likely leading to a rapid evolution and eventual coalescence. From observations, as well as numerical simulations, it has been established that in gas-rich galaxy mergers there is a large inflow of gaseous material to the central kiloparsec of the galactic remnant, often resulting in a massive circumnuclear disc (Barnes & Hernquist 1992; Sanders & Mirabel 1996; Mayer et al. 2007). Driven by dynamical friction and global torques from this disc, the pair of MBHs decays very efficiently down to separations of the order of \u223c1\u20130.1 pc, where it forms a gravitationally bound binary (Escala et al. 2004, 2005; Mayer et al. 2007; Fiacconi et al. 2013; del Valle et al. 2015; Ro\u0161kar et al. 2015). At these sub-parsec scales, most theoretical and numerical studies have focused on the evolution of binaries surrounded by a gaseous circumbinary disc, often either co-rotating (see e.g. Ivanov, Papaloizou & Polnarev 1999; Armitage & Natarajan 2005; Cuadra et al. 2009; Haiman, Kocsis & Menou 2009; Lodato et al. 2009; Nixon et al. 2011; Roedig et al. 2011, 2012; Kocsis, Haiman & Loeb 2012; Amaro-Seoane, Brem & Cuadra 2013; D\u2019Orazio, Haiman & MacFadyen 2013; Mu\u00f1oz & Lai 2016; Miranda, Mu\u00f1oz & Lai 2017; Tang, MacFadyen & Haiman 2017) or counter-rotating (see e.g. Roedig & Sesana 2014; Nixon & Lubow 2015) with respect to the binary\u2019s orbital motion. These discs are generally assumed to be well-defined, smooth, and relaxed, with no attempt to link their presence to the gaseous environment around the binary, nor to the fuelling mechanisms that bring gas to the nucleus. Furthermore, all these idealized scenarios are subject to the disc consumption problem, namely, if the disc dissolves through some process (e.g. star formation, active galactic nucleus (AGN)\/supernovae feedback, Lupi et al. 2015), the evolution of the binary orbit stops. In fact, the evolution of MBHBs in gas-rich environments is intimately related to the unsolved problem of gas supply to the centre of galactic nuclei.","Citation Text":["Escala et al.","2005"],"Functions Text":["Driven by dynamical friction and global torques from this disc, the pair of MBHs decays very efficiently down to separations of the order of \u223c1\u20130.1 pc, where it forms a gravitationally bound binary"],"Functions Label":["Background"],"Citation Start End":[[727,740],[747,751]],"Functions Start End":[[528,725]]}
{"Identifier":"2021AandA...654A.124W__Tanvir_et_al._2017_Instance_1","Paragraph":"The first multi-messenger GW event was discovered on 17 August, 2017. About 1.7 s after the GW170817 signal detected by LIGO and Virgo (Abbott et al. 2017a), the Fermi Gamma-ray Burst Monitor was successfully triggered by GRB 170817A (Abbott et al. 2017b; Goldstein et al. 2017; Zhang et al. 2018) and subsequently a large number of follow-up observations monitored the afterglow emission in different electromagnetic bands from the radio to X-rays (Alexander et al. 2017; Hallinan et al. 2017; Margutti et al. 2017; Troja et al. 2017; D\u2019Avanzo et al. 2018; Ghirlanda et al. 2019; Lazzati et al. 2018; Lyman et al. 2018) and the kilonova AT 2017gfo in the ultraviolet\u2013optical\u2013infrared band (Abbott et al. 2017c; Andreoni et al. 2017; Arcavi et al. 2017; Chornock et al. 2017; Coulter et al. 2017; Covino et al. 2017; Cowperthwaite et al. 2017; Evans et al. 2017; Hu et al. 2017; Kilpatrick et al. 2017; Lipunov et al. 2017; Nicholl et al. 2017; Smartt et al. 2017; Soares-Santos et al. 2017; Tanvir et al. 2017). The observations of GRB 170817A and its afterglows robustly confirmed the long-standing hypothesis that SGRBs can originate from compact binary mergers. Moreover, it became possible to explore the angular structure of the SGRB jet from an off-axis view (Lamb & Kobayashi 2017; Granot et al. 2018; Lazzati et al. 2018; Mooley et al. 2018a,b; Li et al. 2019). Meanwhile, the observations of AT 2017gfo indicated the existence of the merger ejecta, which suggests that the progenitor binary should at least contain one NS. In more detail, the existence of a \u201cblue\u201d and possibly also a \u201cpurple\u201d component in the AT 2017gfo emission further indicated that the merger product of the GW170817 event is very likely to be a hypermassive NS, which lasted for at least a few hundred milliseconds, as an immediately formed black hole can only be associated with a \u201cred\u201d kilonova1 (Cowperthwaite et al. 2017; Perego et al. 2017; Tanaka et al. 2017; Tanvir et al. 2017; Villar et al. 2017; Kawaguchi et al. 2018). Therefore, in summary, the progenitor of the GW170817 event can be identified as a DNS system, which is consistent with the result of the GW analysis.","Citation Text":["Tanvir et al. 2017"],"Functions Text":["and the kilonova AT 2017gfo in the ultraviolet\u2013optical\u2013infrared band"],"Functions Label":["Background"],"Citation Start End":[[992,1010]],"Functions Start End":[[621,689]]}
{"Identifier":"2022AandA...658A.194P__Khata_et_al._2020_Instance_2","Paragraph":"The stellar photospheric parameters we collected from literature for the benchmark stars are summarized in Table A.1. Although most benchmark stars have v sini  2 km s\u22121 (Reiners et al. 2018), there are two stars with larger values: J07558+833 (12.1 km s\u22121) and J13005+056 (16.4 km s\u22121). These stars are useful to investigate the performance of the algorithms when dealing with higher rotational velocities. The literature values were derived with different methods. These methods include: interferometry to estimate the stellar radius and Teff (Boyajian et al. 2012; S\u00e9gransan et al. 2003; von Braun et al. 2014; Berger et al. 2006; Newton et al. 2015), synthetic model fitting using BT-Settl models to determine Teff (Gaidos et al. 2014; L\u00e9pine et al. 2013; Gaidos & Mann 2014; Mann et al. 2015) and log g (L\u00e9pine et al. 2013), empirical relations to derive stellar mass in the form of mass-luminosity relations (Mann et al. 2015; Khata et al. 2020; Boyajian et al. 2012; Berger et al. 2006; S\u00e9gransan et al. 2003), along with the mass-magnitude relations (Maldonado et al. 2015), mass-radius relations (von Braun et al. 2014), mass\u2013Teff relations (Gaidos & Mann 2014; Gaidos et al. 2014), empirical relations to derive the stellar radius in the form of mass-radius relations (Maldonado et al. 2015) and Teff\u2013radius relations (Gaidos & Mann 2014; Gaidos et al. 2014; Houdebine et al. 2019), pEW measurements to determine Teff (Maldonado et al. 2015; Neves et al. 2014; Newton et al. 2015) and [Fe\/H] (Maldonado et al. 2015; Neves et al. 2014; Gaidos et al. 2014; Mann et al. 2015), the definition of spectral indices such as the H2O-K2 index to estimate Teff (Rojas-Ayala et al. 2012), as well as the combination of the H2O-K2 index with pEWs to derive [Fe\/H] (Rojas-Ayala et al. 2012; Khata et al. 2020), the stellar radius and Teff (Khata et al. 2020), and spectral curvature indices for the determination of Teff (Gaidos & Mann 2014). Additionally, [Fe\/H] was derived by using color-magnitude metallicity relations (Dittmann et al. 2016), atomic line strength relations (Gaidos & Mann 2014), and spectral feature relations (Terrien et al. 2015). Terrien et al. (2015) used K-band magnitudes and the Dartmouth Stellar Evolution Program (Dotter et al. 2008) to derive the stellar radius, whereas Mann et al. (2015) employed the Boltzmann equation with Teff determined from synthetic model fits. Last, but not least, Houdebine et al. (2019) derived Teff from photometric colors. For more details on the individual methods, we refer to the descriptions in the corresponding works.","Citation Text":["Khata et al. 2020"],"Functions Text":["as well as the combination of the H2O-K2 index with pEWs to derive [Fe\/H]"],"Functions Label":["Background"],"Citation Start End":[[1788,1805]],"Functions Start End":[[1688,1761]]}
{"Identifier":"2022AandA...659A..44S__Gandolfi_et_al._2010_Instance_1","Paragraph":"Discovery of exoplanetary systems has presented a rather more complex picture of planetary architectures. Transiting exoplanets, those that cross the visible disk of their host stars from our vantage point, permit the measurement of the spin-orbit misalignment between the planetary orbital plane and the stellar equatorial plane that is projected onto the sky, as well as other crucial characteristics (Triaud 2018). This is what is referred to as the sky-projected obliquity angle (\u03bb hereafter). Its measurement is performed through the observations of the Rossiter-McLaughlin effect (Rossiter 1924; McLaughlin 1924). For exoplanets, it entails observations of the stellar radial velocity (RV) during the planetary transit. The anomaly in the measured RV values arises from the deformation of absorption lines from which they are determined, which is caused by the transiting planet occulting either the blue- or red-shifted portion of the spinning stellar disk. The measurement of this effect is possible thanks to precision, high dispersion spectrographs at large telescopes that allow for one to obtain high resolution and large signal-to-noise ratio (S\/N) spectra at relatively high temporal sampling. The measurement of \u03bb has been performed for a large number of transiting exoplanets, full details of which can be found in TEPCAT1 (Southworth 2011). These have revealed a surprising diversity in the orbital alignments (for example Queloz et al. 2000; Winn et al. 2005, 2006, 2009; Triaud et al. 2009; Gandolfi et al. 2010; Mancini et al. 2018; Yu et al. 2018; Lendl et al. 2020; Sedaghati et al. 2021), which is in contrast to the Laplacian ideals of planets forming inside a flat disk, coplanar with the stellar equator and staying there (de Laplace 1796). A surprising picture that has emerged is that a significant fraction of those close-in hot-Jupiter regime planets are on misaligned orbits, as is evident in panel b of Fig. 1 (Albrecht et al. 2012; Dawson 2014). Furthermore, the spectral type of the host also appears to play a role, whereby giant planets around hot stars seem to exist on more oblique orbits (panel a of Fig. 1), perhaps pointing to a different, more chaotic formation history, as compared to their cooler counterparts. The relation between the obliquity and the host star temperature was observed by Winn et al. (2010), who placed the boundary between the two regimes at T\u22c6 = 6250 K (namely theKraft break; Kraft 1967). H\u00e9brard et al. (2011) also point out a lack of planets with masses > 3 MJup on retrograde orbits, the distribution for which is shown in panel c of Fig. 1. Tidal interactions over time with the host star are also expected to realign orbits of close-in, massive planets (Zahn 1977). Attempts have been made to study the impact of stellar age on the obliquity of planetary orbits (for example Safsten et al. 2020; Anderson et al. 2021), with Triaud (2011) finding that hot-Jupiters around younger A stars are more misaligned, setting the age barrier at 2.5 Gyr. However, a lack of precision in the measured stellar ages and the absence of uniform and homogeneous studies estimating those ages have hindered any concrete conclusions being drawn with regard to the impact of stellar ages on planetary orbital alignments. This fact is evident in panel d of Fig. 1.","Citation Text":["Gandolfi et al. 2010"],"Functions Text":["These have revealed a surprising diversity in the orbital alignments (for example","which is in contrast to the Laplacian ideals of planets forming inside a flat disk, coplanar with the stellar equator and staying there"],"Functions Label":["Background","Compare\/Contrast"],"Citation Start End":[[1510,1530]],"Functions Start End":[[1358,1439],[1612,1747]]}
{"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"],"Functions Text":["Following our previous work","we have adopted here a G+L profile fit."],"Functions Label":["Uses","Uses"],"Citation Start End":[[792,811]],"Functions Start End":[[738,765],[813,852]]}
{"Identifier":"2015AandA...578A.124N__Jansen_et_al._(1995)_Instance_1","Paragraph":"The Orion Bar is an ideal source for probing the excitation and chemistry of molecules in PDRs, thanks to its close distance of 414 pc (Menten et al. 2007) and its well-known structure and geometry. The Orion Bar is located between the Orion molecular cloud and an Hii region illuminated by the Trapezium cluster. The FUV radiation field of the Trapezium cluster at the location of the Orion Bar is equivalent to (1\u22124) \u00d7 104\u03c70 in Draine (1978) units. Its orientation changes from face-on to nearly edge-on where the molecular emissions peak. The observations presented in this paper also correspond to the nearly edge-on orientation part of the Orion Bar. The geometrical enhancement of the column densities toward the nearly edge-on part of the Orion Bar was derived by multiple studies and is in the range between 4 and 20. The tilt angle compared to a completely edge-on orientation was suggested to be 3\u00b0 in the model of Hogerheijde et al. (1995) and Jansen et al. (1995). A tilt angle of 3\u00b0 is equivalent to an enhancement factor of 20 for the measured column densities. Based on Oi 1.317 \u03bcm emission, Walmsley et al. (2000) find a model that requires a geometrical enhancement factor of 5 to convert the observed column densities into face-on values. Neufeld et al. (2006) find the geometrical enhancement factor to be 4 based on measured C+ column densities. Using a clumpy 3D PDR model, Andree-Labsch et al. (2014) successfully reproduced the Orion Bar stratification using a clumpy edge-on cavity wall, and they claim that a model of a convex filament fails to describe the structure of the Orion Bar. The average kinetic temperature was estimated to be 85 K (Hogerheijde et al. 1995). Closer to the ionization front, higher temperatures are also measured; for example, OH transitions observed with Herschel\/PACS are consistent with 160\u2212220 K gas (Goicoechea et al. 2011) and CH+ observations with temperatures around 500 K (Nagy et al. 2013). Part of the molecular line emission measured toward the Orion Bar corresponds to an \u201cinterclump medium\u201d with densities between a few 104 and 2 \u00d7 105 cm-3 (Simon et al. 1997). It has been suggested that other molecular lines originate in clumps with densities in the range between 1.5 \u00d7 106 and 6 \u00d7 106 cm-3 (Lis & Schilke 2003). ","Citation Text":["Jansen et al. (1995)"],"Functions Text":["The tilt angle compared to a completely edge-on orientation was suggested to be 3\u00b0 in the model of Hogerheijde et al. (1995) and","A tilt angle of 3\u00b0 is equivalent to an enhancement factor of 20 for the measured column densities."],"Functions Label":["Uses","Uses"],"Citation Start End":[[955,975]],"Functions Start End":[[826,954],[977,1075]]}
{"Identifier":"2022AandA...663A.105P__Bonafede_et_al._2012_Instance_1","Paragraph":"Cluster radio relics are usually found in the outskirts of merging galaxy clusters. They exhibit elongated morphologies and high degrees of polarisation above 1 GHz (up to 70%, Ensslin et al. 1998; Bonafede et al. 2014; Loi et al. 2019; de Gasperin et al. 2022). The resolved spectral index in radio relics shows a gradient: it steepens towards the cluster centre and flattens towards the outskirts. Their size can reach up to \u223c2 Mpc, and high-resolution observations have revealed filamentary structures within relics themselves (Di Gennaro et al. 2018; Rajpurohit et al. 2020, 2022a,b; de Gasperin et al. 2022). The Largest Linear Sizes (LLS) and radio powers of relics are correlated, as well as the integrated spectral index and the radio power (van Weeren et al. 2009b; Bonafede et al. 2012; de Gasperin et al. 2014). Relics trace ICM shock waves with relatively low (M\u2004 \u20043) Mach numbers (Finoguenov et al. 2010; Akamatsu et al. 2013; Shimwell et al. 2015; Botteon et al. 2016). The acceleration of electrons is believed to proceed via diffusive shock acceleration (DSA) in the ICM (Ensslin et al. 1998; Roettiger et al. 1999), in which particles scatter back and forth across the shock front gaining energy at every crossing. Nevertheless, this mechanism has been shown to be rather inefficient in accelerating electrons from the thermal pool (Vazza & Br\u00fcggen 2014; Vazza et al. 2016; Botteon et al. 2020a; Br\u00fcggen & Vazza 2020; see Brunetti & Jones 2014 for a review). Recently, it has been suggested that seed electrons could originate from the tails and lobes (driven by AGN outflows) of cluster radio galaxies (Bonafede et al. 2014; van Weeren et al. 2017; Stuardi et al. 2019), which alleviates the requirements of high acceleration efficiencies at cluster shocks (e.g., Markevitch et al. 2005; Kang et al. 2012, 2017; Botteon et al. 2016; Eckert et al. 2016). In some cases, double relics have been detected on opposite sides of the cluster centre (e.g., Rottgering et al. 1997; van Weeren et al. 2010, 2012b; Bonafede et al. 2012; de Gasperin et al. 2015a). In these clusters it is possible to constrain the merger history, providing important information about the formation processes of relics.","Citation Text":["Bonafede et al. 2012"],"Functions Text":["The Largest Linear Sizes (LLS) and radio powers of relics are correlated, as well as the integrated spectral index and the radio power"],"Functions Label":["Background"],"Citation Start End":[[775,795]],"Functions Start End":[[614,748]]}
{"Identifier":"2019MNRAS.484.2605M__Umeda_et_al._2016_Instance_1","Paragraph":"Remnant BHs of Population III stars with ${\\sim }10^{2}\\, \\mathrm{M}_{\\odot }$ are one of the candidates for such seeds (e.g. Madau & Rees 2001). They can grow SMBHs in the available time if either with continuous accretion at the Eddington limit or with short episode of super-Eddington growth. In reality, however, the BH\u2019s growth can easily be hindered by its own radiative feedback (Milosavljevi\u0107, Couch & Bromm 2009; Park & Ricotti 2011; Orofino, Ferrara & Gallerani 2018; Sugimura et al. 2018). An alternative and attractive pathway for seed BH formation is via the so-called direct collapse (e.g. Bromm & Loeb 2003), where a supermassive star (SMS) with ${\\sim } 10^{5}\\, \\mathrm{M}_{\\odot }$ collapses into a BH with a similar mass by general relativistic instability (Umeda et al. 2016; Woods et al. 2017; Haemmerl\u00e9 et al. 2018). Here, SMSs are supposed to form from a primordial gas in some peculiar sites. Unlike in ordinary first star formation, which is driven by H2 cooling (see e.g. Bromm & Larson 2004; see also Glover 2013 for a review), H2 is dissociated by strong external far-ultraviolet (UV) radiation from nearby galaxies and the contraction of clouds is solely caused by H atomic cooling in the most intensively studied channel for SMS formation (Omukai 2001; Omukai, Schneider & Haiman 2008; Shang, Bryan & Haiman 2010; Regan, Johansson & Wise 2014; Sugimura, Omukai & Inoue 2014). Such clouds contract almost isothermally at \u223c104 K without experiencing vigorous fragmentation. The protostar formed at the centre subsequently accretes the gas at a high rate of $10^{-1}\\, \\mathrm{M}_{\\odot }\\, \\mathrm{yr}^{-1}$ due to high temperature (Latif et al. 2013; Inayoshi, Omukai & Tasker 2014; Becerra et al. 2015; Chon, Hosokawa & Yoshida 2018). Such rapidly accreting protostar inflates greatly in radius with effective temperature of several 1000 K and grows supermassive with ${\\gtrsim } 10^{5}\\, \\mathrm{M}_{\\odot }$ avoiding ionization radiation feedback on the accretion flow (Hosokawa, Omukai & Yorke 2012; Hosokawa et al. 2013), before collapsing by general relativistic instability.","Citation Text":["Umeda et al. 2016"],"Functions Text":["An alternative and attractive pathway for seed BH formation is via the so-called direct collapse","where a supermassive star (SMS) with ${\\sim } 10^{5}\\, \\mathrm{M}_{\\odot }$ collapses into a BH with a similar mass by general relativistic instability"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[777,794]],"Functions Start End":[[501,597],[624,775]]}
{"Identifier":"2020AandA...640A..56R__Hosseinzadeh_et_al._2017_Instance_1","Paragraph":"SN2006jc, PS15dpn and other narrow-line SNe. Two out of the three SNe we considered above are super-luminous, however the final collapse of a PPI+CC progenitor or PISNe does not need to be superluminous (Woosley 2017). The PPI is just one possible mechanism to create CSM, which can produce extreme luminosities by generating radiation from the kinetic energy of the ejecta and\/or narrow emission lines (even if the luminosity does not reach extreme values). The detection of narrow H lines determines the classification of a SN as a type IIn, while the detection of narrow He emission lines determines the classification as type Ibn. Both kinds of event are too common to be entirely explained with PPI+CC progenitors, and it is likely that both classes contain events with a diversity of physical mechanisms (e.g., Pastorello et al. 2008 but see also Hosseinzadeh et al. 2017). Nevertheless, it is possible that at least some of these events might correspond to the observational counterpart of the death of PPI+CC progenitors. In particular, our simulations can produce several solar masses of H-free CSM moving at a few thousand km\u2006s\u22121, which correspond to the width of the He lines detected in some SN Ibn without any fine-tuning required. Even if the detection of narrow lines alone is not sufficient to associate a specific SN to a PPI event, combining evidences from previous coincident transients, large ejecta masses or long lightcurve durations, large 56Ni yields, an extremely young surrounding stellar population, and\/or nucleosynthetic signatures might strengthen the case for associating specific event with this scenario. Possible examples of SN Ibn that might correspond to PPI+CC are SN2006jc and PS15dpn. The former showed relatively narrow He lines possibly hinting to asphericity of the CSM (Foley et al. 2007) and was spatially coincident with an unexplained outburst two years earlier (e.g., Pastorello et al. 2007; Foley et al. 2007). For the latter, Wang & Li (2019) proposed to fit the light curve by combining CSM interaction and radioactive decay, and inferred CSM and 56Ni masses of \u223c0.8\u2006M\u2299 and \u223c0.1\u2006M\u2299, respectively, in good agreement with our models.","Citation Text":["Hosseinzadeh et al. 2017"],"Functions Text":["Both kinds of event are too common to be entirely explained with PPI+CC progenitors, and it is likely that both classes contain events with a diversity of physical mechanisms (e.g.,"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[853,877]],"Functions Start End":[[635,816]]}
{"Identifier":"2018MNRAS.476..184A__Rocha-Pinto_et_al._2000_Instance_1","Paragraph":"The task of estimating stellar ages has been addressed by several authors and lots of different methods are found in the literature. For instance, there are (i) empirical methods, which use a deterministic relation between a given parameter and the age of a star (first proposed by Skumanich 1972). This is the case of gyrochronology (e.g. Barnes 2003, 2007; Mamajek & Hillenbrand 2008; Collier Cameron et al. 2009), decay of chromospheric activity (e.g. Soderblom, Duncan & Johnson 1991; Rocha-Pinto & Maciel 1998; Rocha-Pinto et al. 2000; Pace & Pasquini 2004; Lyra & Porto de Mello 2005; Pace et al. 2009; Zhao et al. 2011; Pace 2013), lithium depletion (Sestito & Randich 2005; Jackson & Jeffries 2014; Carlos, Nissen & Melendez 2016), and \u2018magnetochronology\u2019 (proposed by Vidotto et al. 2014). (ii) Model-dependent methods, which are based on the comparison between measurable physical quantities and the ones expected from stellar structure models that use age as one of its parameters. Isochrone fitting (Edvardsson et al. 1993; Nordstr\u00f6m et al. 2004; Pont & Eyer 2004; J\u00f8rgensen & Lindegren 2005; Silaj & Landstreet 2014; Maxted, Serenelli & Southworth 2015) and asteroseismology (Cunha et al. 2007; Vauclair 2009; Metcalfe et al. 2010; Silva Aguirre et al. 2017) are classified in this category. (iii) Semifundamental methods, those that are based on well-known fundamental physics and employ only few assumptions. These are the cases for the method of cluster expansion (e.g. Makarov 2007) and nucleocosmocronology, known to predict unreliable ages (Ludwig et al. 2010). (iv) Statistical methods, which use statistical relations, like the age\u2013metallicity relation (AMR) and the age\u2013velocity dispersion relation (AVR), between a given property and the age. These relations have not been much explored in the literature as a direct tool to estimate stellar ages. Some few examples of its usage are found in Lachaume et al. (1999) and Maciel, Rodrigues & Costa (2011) (for the AVR) and Spina et al. (2016) (for the AMR, especially [Y\/Mg]\u2009\u00d7\u2009age and [Y\/Al]\u2009\u00d7\u2009age).","Citation Text":["Rocha-Pinto et al. 2000"],"Functions Text":["The task of estimating stellar ages has been addressed by several authors and lots of different methods are found in the literature. For instance, there are (i) empirical methods, which use a deterministic relation between a given parameter and the age of a star (first proposed by Skumanich 1972). This is the case of","decay of chromospheric activity"],"Functions Label":["Background","Background"],"Citation Start End":[[516,539]],"Functions Start End":[[0,318],[417,448]]}
{"Identifier":"2015MNRAS.450.3458C__Cichowolski_et_al._2001_Instance_1","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)"],"Functions Text":["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","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."],"Functions Label":["Uses","Uses"],"Citation Start End":[[444,469]],"Functions Start End":[[234,443],[471,588]]}
{"Identifier":"2020AandA...641A.139D__Dvorak_et_al._(2015)_Instance_1","Paragraph":"The use of N-body simulations that include fragmentation allows us to perform a more detailed study of the final composition of the planets formed. In particular, we can study the water loss and\/or accretion of the final planets more realistically than in the classic models of accretion. Marcus et al. (2010) presented two empirical models for the mantle stripping in differentiated planetary embryos after a collision. The authors set a simple planet structure of two layers, assuming differentiation in core and mantle, where the mantle could be composed by silicate or ice. In this work, the authors concluded that the more energetic the collision, the more mass from the mantle is lost. Therefore, for violent collisions, water could be more easily removed. Dvorak et al. (2015) performed SPH (smoothed particle hydrodynamics) simulations and studied water loss in planetary embryos and water retained in significant fragments after a collision. They concluded that the impact velocity and the impact angle play a key role in the water loss of a planetary embryo after a collision. The investigations developed by Marcus et al. (2010) and Dvorak et al. (2015) suggest that incorporating a realistic model of volatile transport and removal in an N-body code, may lead to reduced water contents on the resulting terrestrial-like planets, in comparison with those derived from classical models that assume perfect mergers. Burger et al. (2018) studied the volatile loss and transfer. The authors focused on hit-and-run encounters using SPH simulations. They concluded that the cumulative effect of several hit-and-run collisions could efficiently strip off volatile layers of protoplanets. Driven by this, Dugaro et al. (2019) studied the water delivery in planets formed in the habitable zone (HZ), using the mantle stripping models derived by Marcus et al. (2010) in their N-body simulationswith fragmentation. The authors showed that fragmentation is not a barrier for the surviving of water worlds in the HZ, and fragments may be important in the final water content of the potentially habitable terrestrial planets formed in situ.","Citation Text":["Dvorak et al. (2015)"],"Functions Text":["performed SPH (smoothed particle hydrodynamics) simulations and studied water loss in planetary embryos and water retained in significant fragments after a collision. They concluded that the impact velocity and the impact angle play a key role in the water loss of a planetary embryo after a collision."],"Functions Label":["Background"],"Citation Start End":[[763,783]],"Functions Start End":[[784,1086]]}
{"Identifier":"2022MNRAS.515...22J__Newman_et_al._2013_Instance_2","Paragraph":"In Fig. 5, we consider how the velocity dispersion profile scales with radius. Specifically, we plot the power-law index (\u03b7) versus the central velocity dispersion (\u03c30). The vast majority of the galaxies with \u03c30 \u2272 2.45 are BGGs and these have negative \u03b7 values. This includes most of the Romulus galaxies (red filled and open circles), the L18 BGGs (blue crosses) and the early-type galaxies that comprise the SAURON sample (Cappellari et al. 2006; grey line and shaded area). In contrast, nearly all of simulated BGGs with \u03c30 \u2272 2.45) from the DIANOGA Hydro-10x simulations Marini et al. (2021) have positive \u03b7 values. For \u03c30 \u2273 2.45, the spread of \u03b7 for the observed galaxies (e.g. L18 and Newman et al. 2013 BCGs) broadens and spans both positive and negative \u03b7 values. In fact, majority of the galaxies tend to have positive \u03b7s. This change in behaviour is well known. A number of studies have noted that on the group-scale and lower, the stellar velocity dispersion profile of the central galaxies tend to decrease with increasing radius. On the cluster-scale, the BCGs typically have rising velocity dispersion profiles with increasing radius (Von Der Linden et al. 2007; Bender et al. 2015; Veale et al. 2017). The origin of this flip is still not well understood. We leave a more detailed investigation of this change to future work. Here, we simply mention two possible explanations: The change in slope may be a reflection of the differences in the dynamical state (e.g. mass-to-light ratio; M\/L) at the outskirts of BCGs (Dressler 1979; Fisher, Illingworth & Franx 1995; Sembach & Tonry 1996; Carter et al. 1999; Kelson et al. 2002; Loubser et al. 2008; Newman et al. 2013; Schaller et al. 2015; Marini et al. 2021), or it could be due to increased contribution from the intragroup\/intracluster light along the line-of-sight and the increased leverage of tangential orbits (Loubser et al. 2020). All of these effects are linked to the increased frequency of galaxy\u2013galaxy interactions and more specifically, central-satellite interactions, implicated in the build-up of extended diffuse stellar component. And, as discussed by Schaye et al. (2015), Oppenheimer et al. (2021), and the EAGLE simulations clearly show that the extended stellar halo becomes increasingly more important, and hosts a non-trivial fraction of the total stellar mass towards the cluster scale.","Citation Text":["Newman et al. 2013"],"Functions Text":["We leave a more detailed investigation of this change to future work. Here, we simply mention two possible explanations: The change in slope may be a reflection of the differences in the dynamical state (e.g. mass-to-light ratio; M\/L) at the outskirts of BCGs"],"Functions Label":["Future Work"],"Citation Start End":[[1663,1681]],"Functions Start End":[[1270,1529]]}
{"Identifier":"2018AandA...610A..38F__Bisterzo_et_al._2017_Instance_1","Paragraph":"Similarly to the [\u03b1\/Fe] ratio, the ratio of the slow (s-) neutron capture process elements to iron can be regarded as a cosmic clock. Ba, Sr, La, and Y are mainly s-process elements produced on long timescales by low mass AGB stars (Matteucci 2012). Since a low mass star must evolve to the AGB phase before the s-process can occur, the s-process elements are characterized by a delay in the production, much like the delay of iron production by SNe Ia relative to the \u03b1 elements production by core collapse SNe. Among the four s-process elements mentioned above, GES provides the abundances of Y II (the first s-process peak element) and Ba II (the second s-process peak element) for all our sample stars. Their abundances behave differently in the Galactic thick and thin discs (Bensby et al. 2005, 2014; Israelian et al. 2014; Bisterzo et al. 2017; Delgado Mena et al. 2017). Unlike the Galactic thick disc stars, which show an almost constant [Ba\/Fe] abundance close to the solar value, the Galactic thin disc stars have their [Ba\/Fe] abundances increasing with [Fe\/H] and reaching their maximum values around solar metallicity, after which a clear decline is seen (see also Cristallo et al. 2015a,b, for the most recent s-process calculation in AGB yields). The same trend is observed in our sample. In Fig. 13 we display the Li-[Ba\/Fe], [Ba\/Fe] as a function of [Fe\/H], and the evolution of absolute Ba abundance A(Ba), as derived from Ba II lines. Similar figures are also plotted for yttrium (Y II). [Ba\/Fe] and [Y\/Fe] values here are derived from MCMC simulations, taking into account the measurement uncertainties of A(Ba II)\/A(Y II) and [Fe\/H]. By applying the same MCMC setups used for [\u03b1\/Fe] (see Sect. 3.1), we calculate the mean values of [Ba\/Fe] and [Y\/Fe] for each star. These values, together with their corresponding 1\u03c3 uncertainties, are listed in Table 1. In the literature there are several theoretical works on the evolution of [Ba\/Fe] and [Y\/Fe] in the Galactic thin disc (e.g. Pagel & Tautvaisiene 1997; Travaglio et al. 1999, 2004; Cescutti et al. 2006; Maiorca et al. 2012; Bisterzo et al. 2017). For comparison, we show in Fig. 13 the predictions of the most recent one (Bisterzo et al. 2017) where the updated nuclear reaction network was used.","Citation Text":["Bisterzo et al. 2017"],"Functions Text":["Among the four s-process elements mentioned above, GES provides the abundances of Y II (the first s-process peak element) and Ba II (the second s-process peak element) for all our sample stars. Their abundances behave differently in the Galactic thick and thin discs"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[830,850]],"Functions Start End":[[513,779]]}
{"Identifier":"2021ApJ...923..106Z__Gabici_et_al._2009_Instance_1","Paragraph":"Diffusive shock acceleration (DSA) operating at expanding shock waves of supernova remnants (SNRs) is widely believed to be the mechanism converting the kinetic energy released by supernova explosions into the energy of cosmic rays (CRs) (e.g., Malkov & Drury 2001). In the DSA theory, CRs being accelerated at shocks must be scattered by self-generated magnetic turbulence. Since the highest-energy CRs in the shock precursor are prone to lack self-generated turbulence, they are expected to escape the shock. The DSA theory generally predicts that a substantial fraction of the shock energy is carried away by escaping CRs. In the presence of molecular clouds surrounding the SNR, escaping CRs can illuminate the clouds through pp interactions, producing gamma-ray emission with a flux depending on the amount of nuclear CRs released by an SNR and the diffusion coefficient in the interstellar medium (ISM; Aharonian & Atoyan 1996; Aharonian et al. 2004; Rodriguez Marrero et al. 2008; Gabici et al. 2009). High Energy Stereoscopic System (H.E.S.S.) observations reveal a complex of sources (HESS J1800-240A, B, and C) \u223c0.5\u00b0 south of the SNR W28, coincident with molecular clouds in the field, and the Large Area Telescope (LAT) on board the Fermi satellite reveals a similar structure in gigaelectronvolt energies (Abdo et al. 2010; Hanabata et al. 2014). The gigaelectronvolt\u2013teraelectronvolt gamma-ray emission around W28 can be regarded as a realization of this scenario (Aharonian et al. 2008). Another example is the detection of two extended gamma-ray structures located at two opposite edges of SNR W44 by Fermi-LAT (Uchiyama et al. 2012; Peron et al. 2020). The gamma-ray emission coincides with the molecular cloud complex that surrounds SNR W44. The gamma-ray emission that appears to come from the surrounding molecular cloud complex can be ascribed to the CRs that have escaped from W44. The total kinetic energy channeled into the escaping CRs is estimated to be larger than a few 1049 erg in both W28 and W44, although the exact number depends on the value of the diffusion coefficient of escaping CRs.","Citation Text":["Gabici et al. 2009"],"Functions Text":["In the presence of molecular clouds surrounding the SNR, escaping CRs can illuminate the clouds through pp interactions, producing gamma-ray emission with a flux depending on the amount of nuclear CRs released by an SNR and the diffusion coefficient in the interstellar medium (ISM;"],"Functions Label":["Background"],"Citation Start End":[[988,1006]],"Functions Start End":[[626,908]]}
{"Identifier":"2019AandA...622A.106M__Maddox_et_al._2018_Instance_1","Paragraph":"The standard single-frequency detection methods for point sources in the CMB and far IR are based on wavelet techniques (Vielva et al. 2003; Barnard et al. 2004; Gonz\u00e1lez-Nuevo et al. 2006) or on the matched filter (or MF hereafter, Tegmark & de Oliveira-Costa 1998; Herranz et al. 2002; Barreiro et al. 2003; L\u00f3pez-Caniego et al. 2006, see also Herranz & Vielva 2010 for a review.). Wavelets are well suited for the detection of compact sources due to their good position-scale determination properties, whereas the MF is the optimal linear detector-estimator because it provides the maximum signal-to-noise (S\/N) amplification for a source with a known shape (usually the point-spread function, or PSF hereafter, of the telescope) embedded in statistically homogeneous and spatially correlated noise. By default, these techniques are applicable only to single-frequency sky images: even for multiwavelength observatories such as the Herschel Space Observatory (Pilbratt et al. 2010) or Planck (Tauber et al. 2010), the standard detection pipelines have produced individual source catalogs for each frequency band (see e.g., Planck Collaboration VII 2011; Planck Collaboration XXVIII 2014; Planck Collaboration XXVI 2016; Maddox et al. 2018). The next logical step is to boost the signal of faint sources by combining the different bands into a single detection, that is, \u201cmultifrequency detection\u201d. Most of the blind component separation algorithms that are used for diffuse components in microwave and far IR astronomy can not deal with the high diversity of spectral behaviors associated to the different populations of extragalactic compact sources (see for example Leach et al. 2008). However, over the last few years a number of multifrequency compact source detection techniques have been proposed in the literature (Herranz & Sanz 2008; Herranz et al. 2009; Lanz et al. 2010, 2013; Planck Collaboration Int. LIV 2018). A review on the topic can be found in Herranz et al. (2012). In particular, if the spatial profile and the spectral energy distribution (SED) of the sources are known, and if the cross-power spectrum is known, or can be estimated from the data, the optimal linear detection method is the matched multifilter (or MMF hereafter, Herranz et al. 2002). Lanz et al. (2010) also showed that the MMF can be generalized for the case where the SED of the sources is not known. This generalization outperforms the single-frequency MF in terms of S\/N and can be used to infer the spectral index of synchrotron-dominated radio sources, as shown in Lanz et al. (2013). However, in this paper we will incorporate a specific SED to the MMF in order to derive a photometric redshift estimation of dusty galaxies and high-redshift star forming galaxies detected in the IR part of the spectrum1. We will do so by applying the multifrequency MMF filter to the first and second data releases of the Herschel Astrophysical Terahertz Large Area Survey (the Herschel-ATLAS or H-ATLAS, Eales et al. 2010), the largest single key project carried out in open time with the Herschel Space Observatory. We restrict our multifrequency analysis to the three wavelength bands covered by the SPIRE instrument aboard Herschel (Griffin et al. 2010), centered around 250, 350 and 500 \u03bcm. As discussed in Hopwood et al. (2010), Lapi et al. (2011), Gonz\u00e1lez-Nuevo et al. (2012), Pearson et al. (2013) and Donevski et al. (2018), the SPIRE bands are ideal for capturing the peak in the SED corresponding to dust emission of star-forming galaxies at z\u2004\u223c\u20042, that is redshifted from its rest-frame wavelength around 70\u2013100 \u03bcm to the SPIRE wavelengths: This is the redshift range where galaxies have formed most of their stars. At higher redshifts, dusty star-forming galaxies (DSFGs) occupy the most massive halos and are among the most luminous objects found at z\u2004\u2273\u20044 (Micha\u0142owski et al. 2014; Oteo et al. 2016; Ikarashi et al. 2017). These high-redshift DSFGs have markedly red colors as seen by SPIRE, with rising flux densities from 250 to 500 \u03bcm (the so-called \u201c500 \u03bcm-risers\u201d), and have received a great deal of attention in the recent years (see for example Ivison et al. 2016; Negrello et al. 2017; Strandet et al. 2017). The DSFGs, and particularly the 500 \u03bcm risers uncovered by Herschel, are providing much insight into the early star forming history of the universe. However, sensitivity and limited angular resolution severely constrain the power of this type of objects as astrophysical probes. The sensitivity of SPIRE allows for the direct detection of only the brightest, and thus rarest objects, at the bright end of the luminosity function. By means of our multifrequency MMF technique, we intend to enhance the detectability and statistical significance of very faint red objects in the H-ATLAS source catalog and so expand the list of reliable 500 \u03bcm-riser candidates.","Citation Text":["Maddox et al. 2018"],"Functions Text":["By default, these techniques are applicable only to single-frequency sky images: even for multiwavelength observatories such as the Herschel Space Observatory","or Planck","the standard detection pipelines have produced individual source catalogs for each frequency band (see e.g.,","The next logical step is to boost the signal of faint sources by combining the different bands into a single detection, that is, \u201cmultifrequency detection\u201d."],"Functions Label":["Background","Background","Background","Motivation"],"Citation Start End":[[1223,1241]],"Functions Start End":[[803,961],[985,994],[1017,1125],[1244,1400]]}
{"Identifier":"2018ApJ...858...91Y__Leroy_et_al._2008_Instance_1","Paragraph":"Our spectroscopic coverage includes the damped Ly\u03b1 absorption (DLA) feature from neutral hydrogen. Analysis of this feature is complicated by blending between the host and the companion, as well as uncertainties in the continuum modeling, and it is difficult to constrain the individual contributions of the two component systems to the overall line. Modeling the line as two separate, blended Voigt profiles, we can robustly constrain the total column density of the two systems together to 5 \u00d7 1019 cm\u22122 \u2264 N(H i) \u2264 9 \u00d7 1019 cm\u22122; our best-fit estimate is (6.7 \u00b1 1.2) \u00d7 1019 cm\u22122, making the system technically a sub-DLA (P\u00e9roux et al. 2003). This spread in N(H i) is due to uncertainties in the relative contributions of the two systems and continuum placement. The derived N(H i) is an order of magnitude less than a typical value through a disk of a spiral galaxy with stellar mass >1010 M\u2299, measured by The H i Nearby Galaxy Survey (THINGS; Walter et al. 2008; Leroy et al. 2008). This result has two implications. One is that SN2017egm may have exploded on the near side of the galaxy, where there is less neutral H i material along the line of sight. One may argue for the second possibility that UV fluxes from the SN explosion could photoionize a large fraction of neutral H i in the disk of NGC 3191, and SN2017egm could be anywhere in the disk. To validate this hypothesis, we calculate the time required, tphot, to photoionize N(H i) \u223c 1021 cm\u22122 over a scale height R of an H i disk. H i-ionizing flux, JUV, at the time of SN explosion is poorly constrained by observation. Let us take JUV as a fraction of fUV of the peak bolometric luminosity (Lbol); thus we have \n\n\n\n\n\n. The Milky Way thin disk (stellar) has a scale height of 100 pc.7\n\n7\nThe H i gas can be more extended (Marasco & Fraternali 2011).\n As the most conservative assumption, let us take R = 50 pc if the SN is at the mid-plane; thus we have \n\n\n\n\n\n cm\u22122 s\u22121. In this equation, without any knowledge of the early-time UV fluxes from SN 2017egm, we assume it is only 10% of the estimated bolometric luminosity (fUV = 0.01). Here, the maximum Lbol is 2 \u00d7 1044 erg s\u22121. To photoionize a column of 1021 cm\u22122 H i atoms, we need the same number of H i-ionizing photons, thus, over a timescale of tphot, we have JUVtphot \u2243 1021 cm\u22122, leading to tphot = 3 \u00d7 109 s. This is two orders of magnitude longer than the time lag between the explosion and the HST spectroscopy date for SN2017egm. Since \n\n\n\n\n\n, a larger fUV value will shorten the photoionization timescale tphot by a factor of a few, but not enough to change our conclusion. This result suggests that photoionization due to the SN explosion is probably localized within a 5 pc region.","Citation Text":["Leroy et al. 2008"],"Functions Text":["The derived N(H i) is an order of magnitude less than a typical value through a disk of a spiral galaxy with stellar mass >1010 M\u2299, measured by The H i Nearby Galaxy Survey (THINGS"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[966,983]],"Functions Start End":[[764,944]]}
{"Identifier":"2021MNRAS.503.3279S__Magrini_et_al._2017_Instance_1","Paragraph":"Among the several features, the distribution of chemical elements across the Galactic disc historically constitutes the most important constraint to chemo-dynamical models of our Milky Way. A number of studies (e.g. Tosi 1988; Hayden et al. 2014, 2015; Anders et al. 2017) have shown the spatial distributions of chemical abundances and their ratios across the Galactic disc. However, these studies are mainly based on field stars, which also include very old populations that had time to migrate significantly and redistribute the chemical elements across the Galaxy (e.g. Sellwood & Binney 2002; Ro\u0161kar et al. 2012; Martinez-Medina et al. 2016). Open clusters are a valuable alternative, being on average younger (Magrini et al. 2017), and therefore a better tracer of the gradients in the disc out of which the most recent stars formed. Since the work of Janes (1979), much observational evidence has established that the metallicity distribution (often abbreviated by the iron-to-hydrogen ratio [Fe\/H]) traced by clusters throughout the Milky Way disc shows a significant decrease with increasing distance from the Galactic Centre. This \u2018radial metallicity gradient\u2019 \u2013 in its apparent simplicity \u2013 reflects a complex interplay between several processes that are driving the evolution of our Galaxy, including star formation, stellar evolution, stellar migration, gas flows, and cluster disruption (Cunha & Lambert 1992, 1994; Friel 1995; Stahler & Palla 2004; Carraro et al. 2006; Boesgaard, Jensen & Deliyannis 2009; Magrini et al. 2009; Frinchaboy et al. 2013; Netopil et al. 2016; Anders et al. 2017; Spina et al. 2017; Bertelli Motta et al. 2018; Quillen et al. 2018). Complementary to the study of the overall metallicity distribution, the abundance ratios of several other elements, such as \u03b1-elements, iron peak, odd-z, and neutron capture, can provide deep insight into the variety of nucleosynthesis processes, with their production sites and time-scales (e.g. Carrera & Pancino 2011; Ting et al. 2012; Reddy, Lambert & Giridhar 2016; Duffau et al. 2017; Magrini et al. 2017, 2018; Donor et al. 2020; Casamiquela et al. 2020). Therefore, understanding the distribution of metals traced by clusters across the Galactic disc is fundamental for explaining the birth, life, and death of both stars and clusters, the recent evolution of our own Milky Way, and the evolution of other spiral galaxies (Boissier & Prantzos 2000; Bresolin 2019).","Citation Text":["Magrini et al. 2017"],"Functions Text":["Open clusters are a valuable alternative, being on average younger","and therefore a better tracer of the gradients in the disc out of which the most recent stars formed."],"Functions Label":["Background","Background"],"Citation Start End":[[716,735]],"Functions Start End":[[648,714],[738,839]]}
{"Identifier":"2020ApJ...895..128M__Zaldarriaga_et_al._2018_Instance_2","Paragraph":"We analyze the 10 BBH mergers reported by LIGO and Virgo in their O1 and O2 observing runs (Abbott et al. 2019a; LIGO Scientific Collaboration & Virgo Collaboration 2019). Before discussing results, it is useful to review expectations from the literature for the spin distributions resulting from different formation scenarios. Isolated binary evolution is predicted to yield black holes with spins preferentially aligned with their orbit. Although spin misalignments may be introduced by natal supernova kicks, episodes of mass transfer and tidal torques serve to realign component spins before the formation of the final black hole binary (Rodriguez et al. 2016; Zevin et al. 2017; Gerosa et al. 2018; Qin et al. 2018; Zaldarriaga et al. 2018; Bavera et al. 2020). The black holes\u2019 spin magnitudes in this scenario are much more uncertain. Recent work indicates that angular momentum is efficiently transported away from stellar cores, leaving black holes with natal spins as low as a \u223c 10\u22122 (Qin et al. 2018; Fuller & Ma 2019). While tides on the progenitor of the second-born black hole can spin up the progenitor star (Zaldarriaga et al. 2018), this effect can be counteracted by mass loss in stellar winds, and more detailed simulations find only low or moderate spin increases due to tides (Qin et al. 2018; Bavera et al. 2020). Meanwhile, dynamically formed systems in dense stellar clusters have no a priori preferred axis, and so are likely to have random spin configurations (Rodriguez et al. 2016, 2018, 2019; Doctor et al. 2020). Once again, however, the expected spin magnitudes are largely unknown, subject to the same uncertainties mentioned above regarding natal black hole spins. One firm prediction of the dynamical scenario concerns the spins of second-generation binaries, whose components were themselves formed from previous mergers. Regardless of their component spins, black hole mergers generally yield remnants with a \u223c 0.7; thus the effective spin of two such second-generation binaries may be large (Fishbach et al. 2017; Gerosa & Berti 2017; Rodriguez et al. 2018, 2019; Doctor et al. 2020).","Citation Text":["Zaldarriaga et al. 2018"],"Functions Text":["While tides on the progenitor of the second-born black hole can spin up the progenitor star","this effect can be counteracted by mass loss in stellar winds, and more detailed simulations find only low or moderate spin increases due to tides"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1124,1147]],"Functions Start End":[[1031,1122],[1150,1296]]}
{"Identifier":"2015MNRAS.450..630S__Chung_et_al._2010_Instance_1","Paragraph":"Many studies often have environmental classes simply divided into (relaxed) \u2018clusters\u2019 or \u2018fields\u2019. However, in a \u039b cold dark matter (\u039bCDM) Universe, most clusters are expected to be the result of group\/smaller cluster mergers \u2013 some of which can be extremely violent. Little is known about the role of cluster and group mergers in galaxy formation and evolution, and whether they could be important in setting the environmental trends which have now been robustly measured and described. It is particularly important to understand if cluster mergers trigger star formation (e.g. Miller & Owen 2003; Ferrari et al. 2005; Owen et al. 2005; Hwang & Lee 2009; Wegner, Chu & Hwang 2015), if they quench it (e.g. Poggianti et al. 2004), or, alternatively, if they have no direct effect (e.g. Chung et al. 2010). Results from Umeda et al. (2004), studying a merging cluster at z \u223c 0.2 (Abell 521) found tentative evidence that merging clusters could perhaps trigger star formation. More recently, Stroe et al. (2014a) conducted a wide field H\u03b1 narrow-band survey over two merging clusters with a simple geometry, with the merger happening in the plane of the sky. Stroe et al. (2014a) find a strong boost in the normalization of the H\u03b1 luminosity function of the CIZA J2242.8+5301 (\u2018Sausage\u2019) cluster, several times above the field and other clusters. The authors suggest that they may be witnessing star formation enhancement or triggered due to the passage of the shock wave seen in the radio and X-rays. Interestingly, Stroe et al. (2014a) do not find this effect on the other similar merging cluster studied (\u2018Toothbrush\u2019), likely because it is a significantly older merger (about 1 Gyr older; cf. Stroe et al. 2014a, 2015), and thus displays only the final result (an excess of post-starburst galaxies instead of H\u03b1 emitters). The results are in very good agreement with simulations by Roediger et al. (2014) and recent observational results by Pranger et al. (2014).","Citation Text":["Chung et al. 2010"],"Functions Text":["It is particularly important to understand if cluster mergers","or, alternatively, if they have no direct effect (e.g."],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[787,804]],"Functions Start End":[[489,550],[732,786]]}
{"Identifier":"2015ApJ...799...42D__Roberge_et_al._2012_Instance_1","Paragraph":"The possible presence of dust in the habitable zones of nearby main-sequence stars is considered a major threat for the direct imaging and characterization of Earth-like extrasolar planets (exo-Earths) with future dedicated space-based telescopes. Several independent studies have addressed this issue and concluded that visible to mid-infrared direct detection of exo-Earths would be seriously hampered in the presence of dust disks 10 to 20\u00c2 times brighter than the solar zodiacal cloud assuming a smooth brightness distribution (e.g., Beichman et al. 2006; Defr\u00c3\u00a8re et al. 2010; Roberge et al. 2012). The prevalence of exozodiacal dust at such a level in the terrestrial planet region of nearby planetary systems is currently poorly constrained and must be determined to design these future space-based instruments. So far, only the bright end of the exozodi luminosity function has been measured on a statistically meaningful sample of stars (Lawler et al. 2009; Kennedy & Wyatt 2013). Based on WISE observations and extrapolating over many orders of magnitude, Kennedy & Wyatt (2013) suggest that at least 10% of gigayear-old main-sequence stars may have sufficient exozodiacal dust to cause problems for future exo-Earth imaging missions. To determine the prevalence of exozodiacal dust at the faint end of the luminosity function, NASA has funded the Keck Interferometer Nuller (KIN) and the Large Binocular Telescope Interferometer (LBTI) to carry out surveys of nearby main-sequence stars. Science observations with the KIN started in 2008 and the results were reported recently (Millan-Gabet et al. 2011; Mennesson et al. 2014). One of their analyses focused on a sample of 20 solar-type stars with no far infrared excess previously detected (i.e., no outer dust reservoir). Assuming a log-normal luminosity distribution, they derived the median level of exozodiacal dust around such stars to be below 60\u00c2 times the solar value with high confidence (95%). Yet, the state-of-the-art exozodi sensitivity achieved per object by the KIN is approximately one order of magnitude larger than that required to prepare future exo-Earth imaging instruments.","Citation Text":["Roberge et al. 2012"],"Functions Text":["The possible presence of dust in the habitable zones of nearby main-sequence stars is considered a major threat for the direct imaging and characterization of Earth-like extrasolar planets (exo-Earths) with future dedicated space-based telescopes. Several independent studies have addressed this issue and concluded that visible to mid-infrared direct detection of exo-Earths would be seriously hampered in the presence of dust disks 10 to 20\u00c2 times brighter than the solar zodiacal cloud assuming a smooth brightness distribution (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[582,601]],"Functions Start End":[[0,537]]}
{"Identifier":"2016MNRAS.455.2959D__Church_et_al._2014_Instance_1","Paragraph":"Here, we invoke an extend accretion disc corona (ADC) model to explain the long-term time lags detected in NS-LMXBs, including the long-term time lags of GX 349+2 that we derive in this work. Analysing the dip and non-dip spectra of NS-LMXBs, Church & Ba\u0142uci\u0144ska-Church (1993, 1995) proposed a Birmingham model which consists of a blackbody component interpreted as the emission from a point source, i.e. the NS, and a power-law component that might be resulted from the Comptonization of thermal emission in an ADC above the accretion disc. Through dip ingress time technique, Church & Ba\u0142uci\u0144ska-Church (2004) measured the radius of the ADC and developed the Birmingham model into an extended ADC model. The measured radius of a thin, hot corona above the accretion disc varies in the range of \u223c(2\u201370)\u00d7 104 km. Therefore, the corona is very extended and the disc is substantially covered by the corona. In the extended ADC model, almost all soft X-ray photons from the accretion disc are inversely Comptonized by the energetic electrons from the extended ADC, which produces the observed hard X-rays, while the observed soft X-rays are interpreted as the emission from the NS; the accretion disc is illuminated by the emission of the NS, leading to the production of the extended ADC above the disc. This model was successfully applied to Cyg-like Z sources (Church, Halai & Ba\u0142uci\u0144ska-Church 2006; Jackson, Church & Ba\u0142uci\u0144ska-Church 2009; Ba\u0142uci\u0144ska-Church et al. 2010) and Sco-like Z sources (including GX 349+2) (Church et al. 2012), as well as atoll sources (Church et al. 2014), so it could be a universal model for NS-LMXBs. Since the extended ADC model is a unified model for NS-LMXBs, we try to interpret the long-term time lags in NS-LMXBs with the help of this model. The extended ADC and the NS are two independent emitting regions, which satisfies the request that the hard and soft X-rays for long-term time lags are emitted from two distant regions, as discussed above. In order to explain the long-term time lags detected in NS-LMXBs in terms of the extended ADC model, we introduce two time-scales. One is the Comptonization time-scale during which the disc seed photons are inversely Comptonized by the high-energy electrons in the extended ADC, and another is the viscous time-scale in the order of hundreds of seconds (Lei et al. 2008), during which the accreting matter flows from the disc on to the NS. The hard X-ray time lags will be produced if the Comptonization time-scale is less than the viscous time-scale, and, contrarily, the soft X-ray time lags will be observed if the Comptonization time-scale is larger than the viscous time-scale. It is noted that a minority of positively correlated short-term time lags (1 s) are listed in Tables 1 and 3, which are derived with the CCF method in our work. These short-term time lags cannot be explained by the models reviewed in section 4.1, because those models are used to interpret the short-term time lags produced in two adjacent energy intervals, while these short-term time lags obtained in this work are derived from two distant energy intervals, i.e. 2\u20135 kev and 16\u201330 keV energy intervals. In the frame of the extended ADC model, we propose that these short-term time lags will be observed under the circumstance that the Comptonization time-scale is comparable with the viscous time-scale.","Citation Text":["Church et al. 2014"],"Functions Text":["This model was successfully applied to","as well as atoll sources","so it could be a universal model for NS-LMXBs."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[1566,1584]],"Functions Start End":[[1302,1340],[1540,1564],[1587,1633]]}
{"Identifier":"2021MNRAS.500.2577K__Mathys_2017_Instance_1","Paragraph":"The binary characteristics of early-type magnetic stars may provide crucial clues, allowing one to test alternative fossil field hypotheses. The non-magnetic chemically peculiar stars of Am (A-type stars with enhanced lines of Fe-peak elements) and HgMn (late-B stars identified by strong lines of Hg and\/or Mn) types are frequently found in close binaries (Gerbaldi, Floquet & Hauck 1985; Ryabchikova 1998; Carquillat & Prieur 2007), including eclipsing systems (Nordstrom & Johansen 1994; Strassmeier et al. 2017; Takeda et al. 2019). In contrast, only about ten close (Porb  20\u2009d) spectroscopic binaries containing at least one magnetic ApBp star are known (Landstreet et al. 2017). The overall incidence rate of magnetic upper main sequence stars in close binaries is less than 2 per cent (Alecian et al. 2015), although this fraction is significantly higher if one includes wide long-period systems (Mathys 2017). This low incidence of magnetic ApBp stars in close binaries is frequently considered as an argument in favour of the stellar merger origin of fossil fields (de Mink et al. 2014; Schneider et al. 2016). In this context, confirmation of magnetic ApBp stars in short-period binary systems gives support to alternative theories or, at least, demonstrates that early-type stars may acquire magnetic fields through different channels. In addition, detached close binary stars, particularly those showing eclipses, are valuable astrophysical laboratories that provide model-independent stellar parameters and allow one to study pairs of co-evolving stars formed in the same environment. Until recently, no early-type magnetic stars in eclipsing binaries were known. The first such system, HD\u200966051, was identified by Kochukhov et al. (2018). The second system, HD\u200962658 containing twin components of which only one is magnetic, was found by Shultz et al. (2019). Several other eclipsing binaries containing candidate ApBp stars were proposed (Hensberge et al. 2007; Gonz\u00e1lez, Hubrig & Castelli 2010; Skarka et al. 2019), but the magnetic nature of these stars has not been verified by direct detections of their fields using the Zeeman effect. In this paper, we put a spotlight on another candidate eclipsing magnetic Bp star, which received little attention prior to our work despite being significantly brighter than the confirmed magnetic eclipsing systems HD\u200962658 and HD\u200966051.","Citation Text":["Mathys 2017"],"Functions Text":["The overall incidence rate of magnetic upper main sequence stars in close binaries is less than 2 per cent","although this fraction is significantly higher if one includes wide long-period systems"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[905,916]],"Functions Start End":[[686,792],[816,903]]}
{"Identifier":"2021MNRAS.506.5836R__Lyne_&_Manchester_1988_Instance_1","Paragraph":"Below, we describe a simple emission model that can explain most of the key emission features of the pulsars presented in this paper. It draws significantly from the model proposed in Timokhin (2010) that explains nulling and mode changing in pulsars by shrinking and expanding the magnetosphere of the neutron star. The model assumes a dipole magnetic field line structure and a schematic diagram is shown in Fig. 10. The emission beam is assumed to be a \u2018fan beam\u2019 with multiple flux tubes along the open field lines (Karastergiou & Johnston 2007; Oswald, Karastergiou & Johnston 2019). We assume the fan beam to be \u201cpatchy\u201d and partially filled meaning that only a portion of the tube is emitting radiation. The patchy beam scenario is one of the common beam models that has been used to explain the emission properties of pulsars (see Lyne & Manchester 1988; Manchester et al. 2010). As shown in Fig. 10, it is assumed that the emission patch is fixed within the polar cap region, i.e. the region consists of open magnetic field lines with respect to the boundary of the comoving magnetosphere (known as the light cylinder). When the magnetosphere shrinks the polar cap region expands accordingly with the dipole field lines and changes the orientation of the emission flux tube produced from the patchy region (see Fig. 10b). Due to these changes, our line of sight (LoS) encounters a different part of the flux tube, resulting in variation in the observed pulse profile shape, switching the pulsar from the normal to swoosh emission. It has been postulated that low-frequency radio emission is produced at a higher altitude compared to the high-frequency emission in the magnetosphere (i.e. \u2013 RFM, Cordes 1978; Oswald et al. 2019). Therefore, due to the curvature of the dipole magnetic field lines, the flux tube of the patchy beam at a lower frequency can be oriented in a direction that is out of our LoS (see Fig. 10b). This results in apparent nulls at low frequencies, which is consistent with the observations of PSR B0919+06 (see Fig. 1, and also Shaifullah et al. 2018).","Citation Text":["Lyne & Manchester 1988"],"Functions Text":["The patchy beam scenario is one of the common beam models that has been used to explain the emission properties of pulsars (see"],"Functions Label":["Background"],"Citation Start End":[[839,861]],"Functions Start End":[[711,838]]}
{"Identifier":"2020MNRAS.497...52H__Sereno_&_Umetsu_2011_Instance_1","Paragraph":"In Fig. 11, we compared the WL-derived masses (MWL) with those derived from the dynamical analysis (Mdyn) and listed in Table 1. The dynamical masses were computed under the assumption of the singular isothermal model and velocity dispersions (Haines et al. 2018). We obtained the WL masses for the four massive clusters (A\u20093556, A\u20093558, A\u20093560, and A\u20093562) with the MCMC method. On the other hand, we only gave the upper limit of the masses for six out of seven low-mass clusters due to the large shape noise in the WL analysis. In the figure, we labelled the low-mass clusters as black crosses, AS\u20090726 as black box and the four massive clusters as cyan circles. For AS\u20090726, we only indicatively adopted the average mass for the low-mass clusters derived using the stacked shear profile. We notice that (i) the masses obtained from the dynamical and WL analysis are consistent within 1\u03c3 for all clusters except A\u20093554, A\u20093558, and AS\u20090726; (ii) the dynamical masses turn out to be systematically higher than the WL-derived ones. Previous studies showed that WL masses obtained from the tangential shear fitting were biased low up to 10 per\u2009cent with a scatter of \u223c25 per\u2009cent (Oguri et al. 2005; Sereno & Umetsu 2011; Sereno & Ettori 2015). The main source of the bias are due to substructures and triaxiality. When a cluster whose major axis is perpendicular to the line of sight, i.e elongated in the sky-plane, a mass obtained with the spherical NFW profile is typically underestimated. The presence of substructures around clusters and uncorrelated large-scale structures along the line of sight also generate the biases for the estimation of the WL masses (Meneghetti et al. 2010; Becker & Kravtsov 2011; Giocoli et al. 2012, 2014). In addition, we point out that the dynamical mass of AS\u20090726 is actually an upper limit since the velocity distribution of member galaxies is strongly bimodal, and certainly not Gaussian (see fig. 14 in Haines et al. 2018). This system probably consists of two groups with velocity dispersions \u223c300\u2009km\u2009s\u22121 rather than one single system with \u03c3 \u223c600\u2009km\u2009s\u22121. This would reduce the mass estimate by a factor of 4. The complex structure of A\u20093558 and, in general, the dynamical activity in the SSC core (see Bardelli et al. 1998; Ettori et al. 2000; Finoguenov et al. 2004; Rossetti et al. 2007a) may explain the systematic differences between the two mass determinations quoted above, since the virial mass tends to overestimate the mass of unrelaxed clusters. Moreover, Rossetti et al. (2007b) studied A\u20093558 with an X-ray observation and showed the possibility of the presence of substructures along the line of sight, which interact with the cluster. Since such structures along the line of sight broaden velocity distribution, masses obtained from the dynamical analysis can be overestimated up to 100 per\u2009cent (Takizawa, Nagino & Matsushita 2010; Pratt et al. 2019). Fig. 14 in Haines et al. (2018) showed the broad velocity distribution of galaxies for the cluster. This indicates the presence of the substructures along the line of sight and the possibility for overestimating the dynamical mass. A\u20093554 also shows strong substructures in the velocity distribution diagram.","Citation Text":["Sereno & Umetsu 2011"],"Functions Text":["Previous studies showed that WL masses obtained from the tangential shear fitting were biased low up to 10 per\u2009cent with a scatter of \u223c25 per\u2009cent","The main source of the bias are due to substructures and triaxiality. When a cluster whose major axis is perpendicular to the line of sight, i.e elongated in the sky-plane, a mass obtained with the spherical NFW profile is typically underestimated."],"Functions Label":["Background","Background"],"Citation Start End":[[1199,1219]],"Functions Start End":[[1032,1178],[1244,1492]]}
{"Identifier":"2020AandA...637A..59A__Massalkhi_et_al._2019_Instance_1","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"],"Functions Text":["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"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[743,764]],"Functions Start End":[[460,668]]}
{"Identifier":"2021AandA...656A.148R__Wiel_et_al._2019_Instance_1","Paragraph":"After the gravitational collapse and if the total mass of the individual cloud is approximately the mass of the Sun (2 \u00d7 1030 Kg; (van Dishoeck 2014), a new astrophysical system forms that is dominated gravitationally by a low-mass protostar known as a young stellar object (YSO). The protostar is in the center of the system and is surrounded by a Keplerian-rotating envelope of dust and gas, that is gravitationally connected and in which the angular momentum is conserved (Cassen & Moosman 1981). The distance to the star will define the surrounding energy, which is dominantly thermal and capable of heating the more distant dustgrain mantles. The processes progressively inject chemical constituents into the gas-phase (Ceccarelli et al. 2001)and determine what types of physical-chemical processes govern in every region of the disk. Circumstellar envelopes oflow-mass protostars (CELMP) are environments that are extraordinarily rich in organic molecules as H2CO or HNCO and in iCOMs such as CH3CN, CH3CHO, and C2 H5OH (Sch\u00f6ier et al. 2002), in addition to other O- and N-bearing complexes (J\u00f8rgensen et al. 2012). IRAS 16293-2422 (hereafter IRAS 16293) was chosen for this work as the quintessential example of such an object. IRAS 16293 is adequate for this role because formaldehyde has previously been detected with notable variations in abundances in three differentiated regions of its protodisk (Ceccarelli et al. 2001; J\u00f8rgensen et al. 2012; Jaber et al. 2014; van der Wiel et al. 2019). Additionally, glycolaldehyde (C2 H4O2), the simplest form of sugar and one of the first intermediates in the formose reaction (Larralde et al. 1995), has also been detected in this object for the first time in space in 2012 (J\u00f8rgensen et al. 2012). The more distant regions of the disk (>4000 AU) contain species such as CO, H2CO, CH3OH, or H2O (van Dishoeck et al. 1995) at low temperatures (10\u201320 K) where the gas is slightly warmer than the dust grains as they are tightly coupled to them. Energy from collisions among dust and gas is therefore considered to be the main heating mechanism in this region. Nevertheless, the thermal energy generated by those collisions is not enough to activate the formation of molecules with barriers at or below 20 K (0.0397 kcal mol\u22121). The gas column density of molecular hydrogen has been estimated to be N(H2) = 1.3 \u00d7 1023 cm\u22122 (Ward-Thompson et al. 1999), with a fractional abundance of formaldehyde \u2013 N(H2CO)\/N(H2) \u2013 ~ 4 \u00d7 10\u221210 cm\u22123 in the gas-phase (Ceccarelli et al. 2001). When the temperature rises above ~ 20 K, CO starts to desorb from the ice grains and enters the gas-phase with an increase of ~ 103 cm\u22123 in detected densities with respectto H2 (Cassen & Moosman 1981; Aikawa et al. 2015). Formaldehyde starts to deplete from frozen grains at around ~ 40 K and it is fully desorbed at ~ 60 K (Ceccarelli et al. 2001). The additional H2CO mixes with the existing circumstellar mass of gas, which may justify why at ~ 700 AU from the core and at gas temperatures between 80 and 100 K ~ 50 kcal mol\u22121), the detected fractional abundances of formaldehyde reach N(H2CO)\/N(H2) = ~ 4.0 \u00d7 10\u22129 cm\u22123 (Ceccarelli et al. 2001). This implies an H2 column density that is estimated to be N(H2) = ~ 5.0 \u00d7 1021 cm\u22122 (Bottinelli et al. 2014). The inner part of the envelope at ~ 150 AU in a region with temperatures of 100\u2013150 K has a higher density of formaldehyde with an N(H2CO)\/N(H2) of ~ 10 \u00d7 10\u22127 cm\u22123, as well as an increase in fractional abundances for H2O (which desorbs from iced mantles at ~80 K) (Ceccarelli et al. 2001). This region also produces new molecules principally due to the thermal energy emitted from the YSO (99.99 kcal mol\u22121). One example is trans-HONO. This chemical compound has recently been detected for first time in space and in this part of the disk (Coutens et al. 2019). Its proposed formation has inspired some reactions proposed in this work that may lead to H2CO. In this region, a column density for molecular hydrogen is considered like that in region II (N(H2) = ~ 5.0 \u00d7 1021 cm\u22122). The three regions of IRAS 16293 dictate the physical parameters for the presentation of our computations which are defined as follows:\n\nRegion I\/d ~ 4000 AU, Tgas = 20 K, Pgas = 2.29 \u00d7 107 K cm\u22123;\nRegion II\/d~ 700 AU, Tgas = 80 K, Pgas = 7.06 \u00d7 107 K cm\u22123;\nRegion III\/d~150 AU, Tgas = 150 K, Pgas = 6.08 \u00d7 108 K cm\u22123.","Citation Text":["van der Wiel et al. 2019"],"Functions Text":["IRAS 16293-2422 (hereafter IRAS 16293) was chosen for this work as the quintessential example of such an object. IRAS 16293 is adequate for this role because formaldehyde has previously been detected with notable variations in abundances in three differentiated regions of its protodisk"],"Functions Label":["Motivation"],"Citation Start End":[[1476,1500]],"Functions Start End":[[1122,1408]]}
{"Identifier":"2021MNRAS.500.5009M__Criscienzo_et_al._2006_Instance_1","Paragraph":"RR Lyrae are old low-mass stars that, during the central helium-burning phase, show mainly radial pulsation while crossing the classical instability strip in the colour\u2013magnitude diagram. From the observational point of view, they represent the most numerous class of pulsating stars in the Milky Way and, being associated with old stellar populations, are typically found in globular cluster and abundant in the Galactic halo and bulge. The investigation of RR Lyrae properties is motivated by their important role both as distance indicators and tracers of old stellar populations. In particular, evolving through the central helium-burning phase, they represent the low-mass, Population II counterparts of Classical Cepheids, as powerful standard candles and calibrators of secondary distance indicators. In particular, they can be safely adopted to infer distances to Galactic globular clusters (see e.g. Coppola et al. 2011; Braga et al. 2016, 2018, and references therein), the Galactic centre (see e.g. Contreras Ramos et al. 2018; Marconi & Minniti 2018; Griv, Gedalin & Jiang 2019), and Milky Way satellite galaxies (see e.g. Coppola et al. 2015; Mart\u00ednez-V\u00e1zquez et al. 2019; Vivas et al. 2019, and references therein). Being associated with old stellar populations, they represent the basis of an alternative Population II distance scale (see e.g. Beaton et al. 2016, to the traditionally adopted Classical Cepheids), more suitable to calibrate secondary distance indicators that are not specifically associated with spiral galaxies (e.g. the globular cluster luminosity function, see Di Criscienzo et al. 2006, and references therein). The properties that make RR Lyrae standard candles are (i) the well-known relation connecting the absolute visual magnitude MV to the metal abundance [Fe\/H] (see e.g. Sandage 1993; Caputo et al. 2000; Cacciari & Clementini 2003; Catelan, Pritzl & Smith 2004; Di Criscienzo, Marconi & Caputo 2004; Federici et al. 2012; Marconi 2012; Marconi et al. 2015, 2018; Muraveva et al. 2018, and references therein); (ii) the period\u2013luminosity relation in the near-infrared (NIR) filters and in particular in the K 2.2 \u03bcm band (see e.g. Longmore et al. 1990; Bono et al. 2003; Dall\u2019Ora et al. 2006; Coppola et al. 2011; Ripepi et al. 2012; Coppola et al. 2015; Marconi et al. 2015; Muraveva et al. 2015; Braga et al. 2018; Marconi et al. 2018, and references therein). In spite of the well-known advantage of using NIR filters (see e.g. Marconi 2012; Coppola et al. 2015, and references therein), in the last decades there has been a debate on the coefficient of the metallicity term of the KB and PL relation (see e.g. Bono et al. 2003; Sollima, Cacciari & Valenti 2006; Marconi et al. 2015, and references therein). On the other hand, it is interesting to note that many recent determinations (see e.g. Sesar et al. 2017; Muraveva et al. 2018) seem to converge towards the predicted coefficient by Marconi et al. (2015), with values in the range 0.16-0.18 mag\u2009dex\u22121. As for the optical bands, our recently developed theoretical scenario (Marconi et al. 2015) showed that, apart from the MV\u2212[Fe\/H] relation that is affected by a number of uncertainties (e.g. a possible non-linearity, the metallicity scale with the associated \u03b1 elements enhancement and helium abundance variations, as well as evolutionary effects, see Caputo et al. 2000; Marconi et al. 2018, for a discussion), the metal-dependent Period\u2013Wesenheit (PW) relations are predicted to be sound tools to infer individual distances. In particular, for the B-, V- band combination, Marconi et al. (2015) demonstrated that the inferred PW relation is independent of metallicity. In order to test this theoretical tool, we need to compare the predicted individual distances with independent reliable distance estimates, for example, the astrometric ones recently obtained by the Gaia satellite (Gaia Collaboration 2016). To this purpose, in this paper we transform the predicted light curves derived for RR Lyrae models with a wide range of chemical compositions (Marconi et al. 2015, 2018) into the Gaia bands, derive the first theoretical PW relations in these filters and apply them to Gaia Data Release 2 Data base (hereinafter Gaia DR2; Gaia Collaboration 2018; Clementini et al. 2019; Ripepi et al. 2019). The organization of the paper is detailed in the following. In Section 2, we summarize the adopted theoretical scenario, while in Section 3 we present the first theoretical light curves in the Gaia filters. From the inferred mean magnitudes and colours, the new theoretical PW relations are derived in Section 4 that also includes a discussion of the effects of variations in the input chemical abundances. In Section 5, the obtained relations are applied to Gaia Galactic RR Lyrae with available periods, parallaxes, and mean magnitudes to infer independent predictions on their individual parallaxes, to be compared with Gaia DR2 results. The conclusions close the paper.","Citation Text":["Di Criscienzo et al. 2006"],"Functions Text":["Being associated with old stellar populations, they represent the basis of an alternative Population II distance scale","more suitable to calibrate secondary distance indicators that are not specifically associated with spiral galaxies (e.g. the globular cluster luminosity function, see","and references therein)."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[1596,1621]],"Functions Start End":[[1230,1348],[1429,1595],[1623,1647]]}
{"Identifier":"2016MNRAS.461..248S__Munari_et_al._2013_Instance_3","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)","Munari et al. (2013)"],"Functions Text":["The difference between the Saro et al. (2013) and","galaxy scaling relations depends on the details of the semi-analytic and hydrodynamical implementations used in Saro et al. (2013) and","respectively."],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[2266,2286],[2422,2442]],"Functions Start End":[[2216,2265],[2287,2421],[2444,2457]]}
{"Identifier":"2017MNRAS.469.2859S__Ramella_et_al._2002_Instance_1","Paragraph":"The difficulty resides in the definition of \u2018group\u2019 itself. If on the simulation side, groups are well defined thanks to an access to the entire three-dimensional information, on the observational side, calling an ensemble of galaxies a group constitutes a great challenge because of a restricted access to the information. In observations, knowing precisely the fraction of collapsed material becomes quite problematic. Still several schemes have been developed to define groups within galaxy catalogues. They mainly invoke Friends of Friends (FoF) like algorithms based on projected separation, radial velocities and even luminosities to identify what are called \u2018groups\u2019 of galaxies (e.g. Huchra & Geller 1982; Geller & Huchra 1983; Ramella et al. 2002; Eke et al. 2004; Yang et al. 2005; Crook et al. 2007; Lavaux & Hudson 2011; Makarov & Karachentsev 2011; Old et al. 2014; Tempel et al. 2014; Old et al. 2015). This paper does not aim at scrutinizing in detail the methods used to group catalogues. It aims at testing two recently released versions of groups for galaxies in the local Universe to understand the differences in the reconstructions generated by two various grouping schemes as described below:\nTempel et al. (2016) introduced a new grouping method (hereafter Tempel Grouping scheme). This method is based on a widely used FoF percolation method, where different linking lengths in radial (along the line of sight) and in transversal (in the plane of the sky) directions are used but the conventional FoF groups are refined using multimodality analysis. More precisely, Tempel et al. (2016) use a model-based clustering analysis to check the multimodality of groups found by the FoF algorithm and they separate nearby\/merging systems. In the current paper, we use published catalogues of groups detected using this new method.Tully Grouping scheme is based on literature groups and in that respect is not a systematic scheme. Within 30 Mpc, groups are those identified by Tully (1987), further away groups are those given in the literature like Abell\u2019s catalogue (Abell, Corwin & Olowin 1989). Recently, Tully (2015a,b) published a more systematic way of deriving groups based on radii of second turn around and iterations. After comparisons, we find that the catalogue grouped with this last scheme is an intermediate between the catalogues obtained with Tully and Tempel Grouping schemes and as such will result in more mitigated conclusions would we compare it to Tempel Grouping scheme. In addition, Tully Grouping scheme has been used so far with the second catalogue to build constrained initial conditions. We thus stick to Tully Grouping scheme in the rest of the paper.4","Citation Text":["Ramella et al. 2002"],"Functions Text":["Still several schemes have been developed to define groups within galaxy catalogues. They mainly invoke Friends of Friends (FoF) like algorithms based on projected separation, radial velocities and even luminosities to identify what are called \u2018groups\u2019 of galaxies (e.g."],"Functions Label":["Background"],"Citation Start End":[[736,755]],"Functions Start End":[[421,691]]}
{"Identifier":"2022MNRAS.511.1714T__Quataert,_Jiang_&_Thompson_2022_Instance_1","Paragraph":"Outflows present similarly confounding puzzles, e.g. the existence and survival of cool and low-ionization gas moving at high speeds in observed outflows (e.g. Veilleux, Cecil & Bland-Hawthorn 2005; Tremonti, Moustakas & Diamond-Stanic 2007; Zhang et al. 2017; Cashman et al. 2021) and high-velocity clouds (HVCs; e.g. Putman et al. 2012; Richter et al. 2017). These clouds should be shredded and destroyed before they can attain the observed velocities if they are accelerated by the ram pressure of a hot wind (e.g. Zhang et al. 2017). Possibly the material is accelerated by radiation pressure (e.g. Murray, Quataert & Thompson 2005; Murray, M\u00e9nard & Thompson 2011; Hopkins, Quataert & Murray 2012) or cosmic rays (CRs; e.g. Everett, Zweibel & Benjamin 2008; Br\u00fcggen & Scannapieco 2020; Quataert, Jiang & Thompson 2022), but alternatively the high-speed cool gas may be due rapid radiative cooling in some situations. If a cloud is sufficiently large, then mixing of the cold cloud with the hotter wind can lead to a region of mixed material with a cooling time less than the cloud-destruction time, and this cooling region can cause a cloud\u2019s mass to grow (Armillotta, Fraternali & Marinacci 2016; Gronke & Oh 2018, 2020; Schneider, Robertson & Thompson 2018; Kanjilal, Dutta & Sharma 2021, but see caveats in S5.4 of Schneider et al. 2020). This multiphase cloud growth could explain the presence of cool gas at high speeds and the pervasive low- and high-ionization metals found in the CGM (e.g. Tumlinson et al. 2011; Werk et al. 2013; Burchett et al. 2019). Similar conclusions have been reached about inflowing gas: galactic accretion is now recognized to sometimes occur in cold, filamentary streams (Kere\u0161 et al. 2005; Dekel et al. 2009), and Mandelker et al. (2020) have shown that if a cold stream has a large enough radius, it can grow in mass due to cooling in a mixing layer. Models of HVCs moving through a galactic halo also find that the HVCs may be disrupted (or grow very little) if they are small (Heitsch & Putman 2009; Marinacci et al. 2010), but they can grow substantially if they are large enough (Fraternali et al. 2015; Gritton, Shelton & Galyardt 2017). Br\u00fcggen & Scannapieco (2016) have shown that thermal conduction, while causing evaporation, can actually extend the lifetime of cold clouds (see also Vieser & Hensler 2007 in a somewhat different context).","Citation Text":["Quataert, Jiang & Thompson 2022"],"Functions Text":["Possibly the material is accelerated by radiation pressure","or cosmic rays (CRs; e.g.","but alternatively the high-speed cool gas may be due rapid radiative cooling in some situations."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[790,821]],"Functions Start End":[[538,596],[702,727],[824,920]]}
{"Identifier":"2015MNRAS.447.3832D__Gandhi_et_al._2008_Instance_1","Paragraph":"GX 339-4 is a recurrent X-ray transient and the system is a confirmed black hole X-ray binary with a low-mass companion star. Although the black hole mass, the system inclination angle and distance are still unknown, they range between 5.8 and 10 M\u2299 (Hynes et al. 2003; Mu\u00f1oz-Darias, Casares & Mart\u00ednez-Pais 2008; Shidatsu et al. 2011), 20\u00b0 and 50\u00b0 (Miller et al. 2006; Done & Diaz Trigo 2010; Shidatsu et al. 2011) and 6 and 15 kpc (Hynes et al. 2004; Zdziarski et al. 2004; Shidatsu et al. 2011), respectively. The source exhibits multiwavelength variability on a broad range of time-scales (Motch, Ilovaisky & Chevalier 1982; Fender, Hanson & Pooley 1999; Corbel et al. 2003, 2013; Dunn et al. 2008; Gandhi 2009; Casella et al. 2010). In addition, it also shows evidence of relativistic jets (Fender et al. 1997; Corbel et al. 2000; Markoff et al. 2003; Gandhi et al. 2008). The observations we used in this work are part of a multiwavelength study of GX 339-4 (Cadolle Bel et al. 2011; Corbel et al. 2013), and in particular of the first mid-infrared study of the source published in Gandhi et al. (2011) and obtained on 2010 March 11. GX 339-4 was observed with theWide-field Infrared Survey Explorer (WISE; Wright et al. 2010) satellite in four bands (1.36 \u00d7 1013, 2.50 \u00d7 1013, 6.52 \u00d7 1013 and 8.82 \u00d7 1013 Hz, respectively, W4, W3, W2 and W1), at 13 epochs, sampled at multiples of the satellite orbital period of 95 min and with a shortest sampling interval of 11 s, when WISE caught the source on two consecutive scans. Radio data were obtained with the Australian Telescope Compact Array during two days \u2013 close to but not simultaneous with WISE data \u2013 on 2010 March 7 and 14. The mean fluxes are 9.1 \u00b1 0.1 and 9.7 \u00b1 0.1 mJy at 5.5 and 9 GHz, respectively. X-ray data were quasi-simultaneous with WISE, taken between epochs 12 and 13 with the Rossi X-ray Timing Explore (RXTE) satellite. Gandhi et al. (2011) confirm the detection in the mid-infrared of a synchrotron break associated with the compact jet in GX 339-4 (Corbel & Fender 2002), and report the first clear detection of its strong variability. This detection of the jet's intrinsic variability and the overall properties of GX 339-4 make it the ideal source to test our model.","Citation Text":["Gandhi et al. 2008"],"Functions Text":["In addition, it also shows evidence of relativistic jets"],"Functions Label":["Background"],"Citation Start End":[[857,875]],"Functions Start End":[[738,794]]}
{"Identifier":"2017ApJ...839...72A__Emsellem_et_al._1994_Instance_2","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"],"Functions Text":["Although a supermassive BH can be modeled by adding a Keplerian potential, it is much simpler to model the BH as this small Gaussian"],"Functions Label":["Uses"],"Citation Start End":[[1822,1842]],"Functions Start End":[[1688,1820]]}
{"Identifier":"2022ApJ...925..203J__Combes_et_al._2019_Instance_1","Paragraph":"The ALMA results mentioned earlier (Alonso-Herrero et al. 2020; Garc\u00eda-Burillo et al. 2021) add yet another scale given that the KDC contains both an r \u223c 150\u2013200 pc molecular gas ring and an inner r  50 pc torus-like structure. This smaller nuclear disk\/ring may correspond to the structure predicted by a radiation-driven fountain model, where AGN radiation feedback induces vertical gas flows that result in a geometrically thick torus (Wada 2012). This mechanism would lead to tori that are a few tens of parsec wide and are dynamic, evolving structures as proposed to interpret several recent observational studies and compilations (e.g., Ramos Almeida & Ricci 2017; Combes et al. 2019; H\u00f6nig 2019). Furthermore, Alonso-Herrero et al. (2021) analyzed high-resolution mid-IR imaging of NGC 7582 from VLT\/VISIR and found both an unresolved component and an extended polar dust component. These authors argue that these observations can be interpreted with a torus+wind model, according to which IR radiation pressure creates a dusty wind component that contributes to AGN obscuration (e.g., Venanzi et al. 2020). Recent models with a realistic 3D distribution of clumpy dusty material can also reproduce polar mid-IR emission starting with a standard clumpy model, depending on the torus opening angle and scale height (Nikutta et al. 2021). Subparsec-resolution observations will be needed to better constrain torus parameters. For example, NGC 1068 observations with the GRAVITY instrument on the European Southern Observatory Very Large Telescope Interferometer revealed a thin ring with an inner radius of r = 0.24 pc, close to the sublimation radius and inconsistent with a geometrically and optically thick torus on these small scales (Gravity Collaboration et al. 2020). Future high-resolution observations of the nuclear regions of AGN hosts will augment our understanding but will likely still need to be combined with probes on a range of physical scales to establish the full picture of AGN obscuration and AGN outflows.","Citation Text":["Combes et al. 2019"],"Functions Text":["This mechanism would lead to tori that are a few tens of parsec wide and are dynamic, evolving structures as proposed to interpret several recent observational studies and compilations (e.g.,"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[671,689]],"Functions Start End":[[451,642]]}
{"Identifier":"2020ApJ...890...89G__Bieler_et_al._2015_Instance_1","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"],"Functions Text":["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","of molecular oxygen","all indicate that comet 67P formed at low temperature and did not experience any substantial global-scale heating after its formation."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[613,631]],"Functions Start End":[[235,447],[592,611],[1025,1159]]}
{"Identifier":"2022MNRAS.512.3137Z__Katz_et_al._1999_Instance_1","Paragraph":"However, it is not straightforward to explain H2 formation in astronomical sources even when the catalytic roles of dust grains are introduced into models. Interstellar species are believed to be formed on cold grain surfaces via the so called Langmuir\u2013Hinshelwood mechanism (Watson & Salpeter 1972; Pickles & Williams 1977; Hasegawa, Herbst & Leung 1992). To form H2, H atoms accrete on dust grains and then bind weakly with surfaces, which is known as physisorption. They can overcome the diffusion barrier and move on the grain surfaces via quantum tunnelling or thermal hopping. However, laboratory and theorectical studies showed that quantum tunnelling does not contribute much to the mobility of H atoms on silicate, carbonaceous or solid amorphous water (ASW) surfaces (Pirronello et al. 1997, 1999; Katz et al. 1999; Nyman 2021). If H atoms encounter other H atoms, then H2 molecules are formed. But H atoms can also desorb and leave grain surfaces. A hydrogen atom must reside on a grain long enough to find a partner H atom to form H2. As the dust temperature increases, the H atom desorption and diffusion rates also increase. So the temperature of grain surfaces must be sufficiently low so that an H atom can encounter another one before it desorbs. On the other hand, the temperature of grain surfaces must be high enough so that H atoms can diffuse on the grain surface. The parameter that measures how strongly species are to bound to grain surfaces is called desorption energy. It was found that if we assume a single H desorption energy in models, the dust temperature range over which efficient H2 formation occurs is narrow (6\u201310 K for olivine grains) (Katz et al. 1999). Moreover, it was found that the highest dust temperature at which H2 can be formed efficiently is less than 17\u2009K (Katz et al. 1999). However, the grain surface temperature in the unshielded diffuse clouds, where hydrogen molecules are believed to be efficiently formed, is around 20 K (Li & Draine 2001).","Citation Text":["Katz et al. 1999"],"Functions Text":["However, laboratory and theorectical studies showed that quantum tunnelling does not contribute much to the mobility of H atoms on silicate, carbonaceous or solid amorphous water (ASW) surfaces"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[808,824]],"Functions Start End":[[583,776]]}
{"Identifier":"2016ApJ...822L..10I__Imanishi_&_Dudley_2000_Instance_1","Paragraph":"We measured the total continuum flux in the SW star-forming region (rectangular with a width of 07 in the east\u2013west \u00d7 a width of 21 in the north\u2013south) (02h42m4059\u20134064, \u221200\u00b000\u20324789\u20135000)J2000 to be \u223c10 mJy. If this is solely due to the thermal free\u2013free emission inside star-forming H ii-regions, then the corresponding far-infrared (40\u2013500 \u03bcm) luminosity becomes \u223c6 \u00d7 1043 (erg s\u22121), or there is a star formation rate of \u223c3 M\u2299 yr\u22121 (Kennicutt 1998), when Equation (1) of Nakanishi et al. (2005) is used. If dust thermal radiation contributes in some way to the continuum flux there, the estimated star formation luminosity will be smaller. Previous infrared spectroscopy failed to clearly detect the signatures of ongoing active star formation at the central few arcsec region of the NGC 1068 nucleus due to the lack of polycyclic aromatic hydrocarbon (PAH) emission features (Imanishi et al. 1997; Le Floc\u2019h et al. 2001), which are a good probe of star formation activity (Moorwood 1986; Imanishi & Dudley 2000). It was also argued that the molecular gas properties at the NGC 1068 nucleus are dominated by AGN radiation rather than star formation (Usero et al. 2004; Garcia-Burillo et al. 2010). Thus, our continuum emission map provides a new signature for the detectable amount of star formation activity in the nuclear few arcsec region of NGC 1068, thanks to the high sensitivity of ALMA. The PAH emission flux within the nuclear 38 \u00d7 38 region in infrared spectroscopy is estimated to be 2.7 \u00d7 1040 (erg s\u22121) (Imanishi 2002), which corresponds to star formation-originated far-infrared luminosity with 2.7 \u00d7 1043 (erg s\u22121) (Mouri et al. 1990). This upper limit is apparently smaller than the above estimate but can be reconciled because (1) not all of the SW star-forming region was covered by previous infrared slit spectroscopy (Imanishi et al. 1997), (2) the estimated star formation luminosity is reduced if dust thermal radiation contributes to the observed continuum emission at \u223c266 GHz, and (3) some fraction of the PAHs can potentially be destroyed by the AGN's strong X-ray radiation in the close vicinity of the AGN in NGC 1068 (Voit 1992).","Citation Text":["Imanishi & Dudley 2000"],"Functions Text":["Previous infrared spectroscopy failed to clearly detect the signatures of ongoing active star formation at the central few arcsec region of the NGC 1068 nucleus due to the lack of polycyclic aromatic hydrocarbon (PAH) emission features","which are a good probe of star formation activity"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[991,1013]],"Functions Start End":[[642,877],[925,974]]}
{"Identifier":"2020ApJ...898...25T__Spera_et_al._2019_Instance_1","Paragraph":"Recent detections of gravitational waves (GWs) have shown evidence for a high rate of black hole (BH)\u2013BH and neutron star (NS)\u2013NS mergers in the universe (Abbott et al. 2016a, 2016b, 2017a, 2017b, 2017c, 2019a; Zackay et al. 2019, 2020; Venumadhav et al. 2020). However, the proposed astrophysical pathways to mergers remain highly debated. Possible compact-object merger pathways include isolated binary evolution (Dominik et al. 2012; Kinugawa et al. 2014; Belczynski et al. 2016, 2017; Breivik et al. 2016; Giacobbo et al. 2018; Bavera et al. 2019; Spera et al. 2019) accompanied by mass transfer (Inayoshi et al. 2017a; Pavlovskii et al. 2017; van den Heuvel et al. 2017), common-envelope ejection (e.g., Paczynski 1976; Ivanova et al. 2013), envelope expansion (Tagawa et al. 2018), or chemically homogeneous evolution in a tidally distorted binary (de Mink & Mandel 2016; Mandel & de Mink 2016; Marchant et al. 2016), evolution of triple or quadruple systems (Antonini et al. 2017; Liu & Lai 2017, 2018, 2019; Silsbee & Tremaine 2017; Arca-Sedda et al. 2018; Hoang et al. 2018b; Randall & Xianyu 2018; Fragione & Kocsis 2019; Fragione et al. 2019; Michaely & Perets 2019), gravitational capture (O\u2019Leary et al. 2009; Kocsis & Levin 2012; Gond\u00e1n et al. 2018b; Rodriguez et al. 2018; Rasskazov & Kocsis 2019; Zevin et al. 2019; Samsing et al. 2020), dynamical evolution in open clusters (Banerjee 2017, 2018a, 2018b; Bouffanais et al. 2019; Kumamoto et al. 2019; Rastello et al. 2019) and dense star clusters (e.g., Portegies Zwart & McMillan 2000; O\u2019Leary et al. 2006, 2016; Samsing et al. 2014; Ziosi et al. 2014; Mapelli 2016; Rodriguez et al. 2016a, 2016b; Askar et al. 2017; Fujii et al. 2017; di Carlo et al. 2019; Zevin et al. 2019; Zhang et al. 2019), and dynamical interaction in gas-rich nuclear regions (McKernan et al. 2012, 2014, 2018; Bellovary et al. 2016; Bartos et al. 2017; Stone et al. 2017; Leigh et al. 2018; Tagawa & Umemura 2018; Yi et al. 2018; Secunda et al. 2019; Yang et al. 2019a, 2019b; Gayathri et al. 2020; McKernan et al. 2020; Tagawa et al. 2020).","Citation Text":["Spera et al. 2019"],"Functions Text":["However, the proposed astrophysical pathways to mergers remain highly debated. Possible compact-object merger pathways include isolated binary evolution"],"Functions Label":["Background"],"Citation Start End":[[552,569]],"Functions Start End":[[262,414]]}
{"Identifier":"2021ApJ...923L..22A__Rosado_et_al._2015_Instance_2","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"],"Functions Text":["individual periodic signals or continuous waves (CWs"],"Functions Label":["Background"],"Citation Start End":[[1032,1050]],"Functions Start End":[[865,917]]}
{"Identifier":"2022AandA...659A.124H__Liu_et_al._(2013b)_Instance_1","Paragraph":"Combining different samples from various instruments at different redshifts therefore inevitably introduces ENLR size\u2013luminosity relations with different slopes \u03b1 depending on the details of target selection and analysis approaches. Slopes ranging from \u03b1\u2004=\u20040.22\u2005\u00b1\u20050.04 (Greene et al. 2012), \u03b1\u2004=\u20040.25\u2005\u00b1\u20050.02 (Liu et al. 2013b), \u03b1\u2004\u223c\u20040.3\u20130.4 (Hainline et al. 2013; Chen et al. 2019a), to \u03b1\u2004\u223c\u20040.5 (Bennert et al. 2002; Husemann et al. 2014) are reported in the literature. The slopes solely inferred from the CARS data are consistent with those reported by Greene et al. (2012) and Liu et al. (2013b) and are therefore on the shallower side of previous estimates. Nevertheless, the scatter in the observed relation is significant and measured slope variations might be entirely attributed to the observationally induced biases as discussed above. A slope of \u03b1\u2004=\u20040.5 is reminiscent of the BLR size-luminosity relation, but would require a constant ionization parameter U that demands a constant density with radius. This is not observed for the ENLR on kiloparsec scales (e.g., Bennert et al. 2006; Kakkad et al. 2018) and more detailed photoionization calculations are required to predict the shallower slopes inferred for most studies (Dempsey & Zakamska 2018). We cannot study the radial variations of U as our snapshot MUSE observations are not deep enough to map the electron density given the too low S\/N of the [S\u202fII] doublet on kpc scales. However, the photoionization calculations do not take into account variable ionizing flux from AGN on 105 yr time scales (Schawinski et al. 2015) and the various geometrical intersections of the ionizing radiation field with the gas distribution of the galaxies. The CARS survey is least biased with regard to RENLR given the narrow redshift range and large dynamic range offered by MUSE (see Fig. 13). Therefore, the CARS survey is one of the best data set to explore the origin of the significant scatter in ENLR size\u2013luminosity relation and search for additional factors or more fundamental parameters controlling the ENLR size.","Citation Text":["Liu et al. 2013b"],"Functions Text":["Combining different samples from various instruments at different redshifts therefore inevitably introduces ENLR size\u2013luminosity relations with different slopes \u03b1 depending on the details of target selection and analysis approaches. Slopes ranging from","\u03b1\u2004=\u20040.25\u2005\u00b1\u20050.02","are reported in the literature."],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[308,324]],"Functions Start End":[[0,252],[291,306],[437,468]]}
{"Identifier":"2016ApJ...831...41W__Dai_&_Liu_2012_Instance_1","Paragraph":"Our results have immediate stimulations for further research. First, although here we have studied the SNe Ic-BL not associated with GRBs, the main conclusion can be equally applied to GRB-SNe. It is usually believed that the central engines of GRBs are black holes (Popham et al. 1999; Narayan et al. 2001; Kohri & Mineshige 2002; Liu et al. 2007; Song et al. 2016) or magnetars (Usov 1992; Dai & Lu 1998a, 1998b; Zhang & Dai 2008, 2009, 2010; Giacomazzo & Perna 2013; Giacomazzo et al. 2015). However, since it is currently infeasible to identify the GRB central engine directly because of the cosmological distance scales of GRBs (Kumar & Zhang 2015), the researchers instead pursue indirect signatures of black holes (Geng et al. 2013; Wu et al. 2013; Yu et al. 2015) and magnetars (Dai et al. 2006; Gao et al. 2013; Wang & Dai 2013; Yu et al. 2013; Zhang 2013; Metzger & Piro 2014; Wang et al. 2015a, 2016a; Li & Yu 2016; Liu et al. 2016) that power the energetic GRBs. Growing indirect observational evidence suggests that magnetars could act as the central engines of both LGRBs and short-duration GRBs (Dai et al. 2006; Rowlinson et al. 2010, 2013; Dai & Liu 2012; Wang & Dai 2013; Wu et al. 2014; Gao et al. 2015; Greiner et al. 2015). However, because of the high mass of 56Ni needed to heat the GRB-SNe, magnetars are doubted as the candidate central engine of GRBs. With our demonstration that this high mass of 56Ni is actually not the case, such a concern is removed.8\n\n8\nWe note that Cano et al. (2016a) drew the conclusion that GRB-SNe are powered by 56Ni decay under the assumption that the central engine of GRB-SNe is a magnetar. In drawing this conclusion, Cano et al. (2016a) assume that the mangetar\u2019s rotational energy is equally divided between GRB afterglow and SNe. This assumption is somewhat ad hoc and unjustified. At the least, although the jet launching could be the result of the magnetar spin-down, it seems more likely to be the result of accretion onto the magnetar (Zhang & Dai 2008, 2009, 2010). Furthermore, Cano et al. (2016a) do not consider the origin of the kinetic energy of the GRB-SNe in their model. If we accept the assumption that a huge amount of kinetic energy of the GRB-SNe comes from the rotational energy of the magnetar, the initial rotational period of the magnetar cannot be as long as given by Cano et al. (2016a). Finally, as we mentioned above, our conclusion that the 56Ni mass for the tail modeling is usually much smaller than the peak modeling is consistent with the finding by Maeda et al. (2003).\n\n","Citation Text":["Dai & Liu 2012"],"Functions Text":["Growing indirect observational evidence suggests that magnetars could act as the central engines of both LGRBs and short-duration GRBs"],"Functions Label":["Future Work"],"Citation Start End":[[1157,1171]],"Functions Start End":[[975,1109]]}
{"Identifier":"2021AandA...655A..22M__Hopkins_et_al._2014_Instance_1","Paragraph":"Despite the consensus on the importance of galactic winds in dwarf galaxy formation, large uncertainties persist in how the winds are driven. As the ISM physics is typically not resolved in full cosmological simulations (although some simulations are getting close, see e.g., Wheeler et al. 2019; Agertz et al. 2020), subgrid feedback models are often used to alleviate the numerical \u2018overcooling\u2019 problem, whereby energy injected into the ISM is rapidly cooled away because of unresolved ISM structures. Many studies have investigated supernova (SN) feedback, perhaps the most important process to drive galactic wind in dwarf galaxies. Different approaches of subgrid models were explored, ranging from kinetic feedback models where wind mass loading and velocities are predetermined (Springel & Hernquist 2003; Oppenheimer & Dav\u00e9 2006; Dalla Vecchia & Schaye 2008), to explicit models where radiative cooling is temporarily shut off (Stinson et al. 2006; Agertz et al. 2011; Teyssier et al. 2013), to mechanical feedback where the momentum boost during the unresolved adiabatic phase is calibrated through small-scale ISM simulations and injected into the ISM (Hopkins et al. 2014; Kimm & Cen 2014; Smith et al. 2018). Simulations of dwarf galaxies in recent years have often adopted one of these approaches for SN feedback. Although many are considered successful in reproducing observations, few attempts have been made to compare these models in full cosmological simulations. It is therefore important to gauge how the evolution of simulated dwarf galaxies (especially at the low mass end) is dependent on feedback models, in order to provide a more robust model with realistic uncertainties for future observations. Moreover, most existing SN feedback models take into account individual SN explosions. However, recent studies such as Nath & Shchekinov (2013) and Sharma et al. (2014), have pointed out that stars form in clusters and, through multiple repeated SN explosions, generate the so-called superbubbles, which contain hot gas, and the sharp temperature gradients between hot and cold gas make thermal conduction an extremely important process. The evolution of superbubbles and the importance of thermal conduction have been investigated using ISM simulations with direct treatments of thermal conduction (e.g., El-Badry et al. 2019; Steinwandel et al. 2020). However, the resolution requirement for directly modelling conduction is still unattainable for cosmological galaxy formation simulations.","Citation Text":["Hopkins et al. 2014"],"Functions Text":["Different approaches of subgrid models were explored","to mechanical feedback where the momentum boost during the unresolved adiabatic phase is calibrated through small-scale ISM simulations and injected into the ISM","Simulations of dwarf galaxies in recent years have often adopted one of these approaches for SN feedback. Although many are considered successful in reproducing observations, few attempts have been made to compare these models in full cosmological simulations. It is therefore important to gauge how the evolution of simulated dwarf galaxies (especially at the low mass end) is dependent on feedback models, in order to provide a more robust model with realistic uncertainties for future observations."],"Functions Label":["Background","Background","Motivation"],"Citation Start End":[[1164,1183]],"Functions Start End":[[638,690],[1001,1162],[1222,1723]]}
{"Identifier":"2019AandA...621A..27F__Delvecchio_et_al._2015_Instance_1","Paragraph":"SMBHs and host galaxies share several properties. Both SMBHs and galaxies have exponential cut-offs at the high mass end of their co-moving space densities (e.g., Shankar et al. 2009; Ilbert et al. 2013; Kelly & Shen 2013). The population of both SMBHs and galaxies also exhibit mass downsizing whereby the oldest, in the case of galaxies, and the most massive of SMBHs grew early and rapidly (e.g., Thomas et al. 2005, 2010; Merloni & Heinz 2008). However, there is a mismatch in both the shape and co-moving number density between galaxies and dark matter halos, especially at the low and high mass ends of these functions (Benson et al. 2003). Because powerful AGN can have a mechanical and radiative energy output similar to or exceeding that of the binding energy of a massive galaxy and dark matter halo, AGN are thought to play a key role in regulating the baryonic growth of galaxies. Both observations and simulations have suggested that there may be a positive trend between the mean black hole accretion rate and star-formation rate (SFR; e.g., Delvecchio et al. 2015; McAlpine et al. 2017), while the mean SFR as a function of black hole accretion rate shows no correlation for low luminosity sources (e.g., Stanley et al. 2015; McAlpine et al. 2017). One should be cautious when interpreting both theoretical and observational results in the definition of what exactly AGN feedback is and how AGN affect their host galaxies to explain the properties of an ensemble of galaxies (Scholtz et al. 2018). The strength and nature of AGN feedback \u2013 the cycle whereby the SMBH regulates both its own growth and that of its host \u2013 depends on galaxy mass and morphology. For example, the most massive elliptical galaxies are generally metal-rich and old, while less massive lenticular galaxies, which make up the bulk of the early-type galaxy population, have star formation histories that lasts significantly longer (Thomas et al. 2005, 2010; Emsellem et al. 2011; Krajnovi\u0107 et al. 2011). Clearly, if AGN feedback plays a crucial role in shaping the ensemble of galaxies, its impact on massive dispersion dominated galaxies must result in somewhat different characteristics in these galaxies compared to rotationally-dominated and predominately less massive lenticular galaxies.","Citation Text":["Delvecchio et al. 2015"],"Functions Text":["Both observations and simulations have suggested that there may be a positive trend between the mean black hole accretion rate and star-formation rate (SFR; e.g.,"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1056,1078]],"Functions Start End":[[893,1055]]}
{"Identifier":"2016ApJ...824..142Q__Galvin_et_al._2008_Instance_1","Paragraph":"Current observations of the interstellar neutral helium trajectories by the Interstellar Boundary EXplorer (IBEX) (McComas et al. 2009) determined the inflow longitude of the interstellar wind to be \n\n\n\n\n\n (Leonard et al. 2015), \u223c757 (McComas et al. 2015), \n\n\n\n\n\n (Bzowski et al. 2015), and \n\n\n\n\n\n (Schwadron et al. 2015). Recent studies using the Ulysses\/GAS instrument made observations of the interstellar helium distribution by detecting the sputtered ions from a lithium fluoride coated surface, finding the inflow direction to be \n\n\n\n\n\n (Bzowski et al. 2014) and \n\n\n\n\n\n (Wood et al. 2015). However, recent observations of the pickup helium focusing cone using the Plasma and Suprathermal Ion Composition (PLASTIC; Galvin et al. 2008) on board the Solar and Terrestrial Relations Observatory Ahead (STEREO A) spacecraft measured the peak density and determined the inflow longitude to be \n\n\n\n\n\n (Drews et al. 2012). In addition, Gershman et al. (2013) measured the peak density of the pickup helium focusing cone using the Fast Imaging Plasma Spectrometer (FIPS; Andrews et al. 2007) instrument on board the MErcury, Surface, Space ENvironment, GEochemistry, and Ranging (MESSENGER) spacecraft and the Solar Wind Ion Composition Spectrometer (Gloeckler et al. 1998) instrument on board the Advanced Composition Explorer spacecraft and determined inflow longitudes of \n\n\n\n\n\n and \n\n\n\n\n\n, respectively. Although these pickup ion measurements fall within the 1\u03c3 uncertainties of the IBEX observations, they appear to show consistently small deviations in the same direction. Refer to Figure 1 for a diagram of the \n\n\n\n\n\n longitudinal difference between the observations focused on in this study of Drews et al. (2012) and Schwadron et al. (2015). The difference between the pickup helium and neutral helium observations suggests that the transport of pickup ions inside 1 au may play a role\u2014a concept that was addressed but not factored into the derived peak longitudes in the previously mentioned pickup ion studies and has recently been up for debate (Chalov 2014; Lallement & Bertaux 2014; Frisch et al. 2015).  \n","Citation Text":["Galvin et al. 2008"],"Functions Text":["However, recent observations of the pickup helium focusing cone using the Plasma and Suprathermal Ion Composition (PLASTIC","on board the Solar and Terrestrial Relations Observatory Ahead (STEREO A) spacecraft measured the peak density and determined the inflow longitude to be"],"Functions Label":["Uses","Uses"],"Citation Start End":[[720,738]],"Functions Start End":[[596,718],[740,892]]}
{"Identifier":"2019ApJ...871L..22W__Alexandrova_2008_Instance_2","Paragraph":"In analogy to the hydrodynamic case, the nonlinear coherent vortex structure also plays an important role in plasma dynamics and transport processes (Hasegawa & Mima 1978; Shukla et al. 1985; Petviashvili & Pokhotelov 1992; Horton & Hasegawa 1994). These vortices tend to have a long lifetime and are widely observed in space, laboratory, and numerical simulation of plasma (Chmyrev et al. 1988; Burlaga 1990; Volwerk et al. 1996; Stasiewicz et al. 2000; Sundkvist et al. 2005; Alexandrova et al. 2006; Alexandrova 2008; Alexandrova & Saur 2008; Vianello et al. 2010; Servidio et al. 2015). An essential subset of these plasma vortices is known as Alfv\u00e9n vortices, which can be viewed as the cylindrical analog of the nonlinear Alfv\u00e9n wave (Petviashvili & Pokhotelov 1992). The Alfv\u00e9n vortices have an axis that is nearly parallel to the unperturbed magnetic field, along which the shape is generally invariant. Thus, these vortices are quasi-2D structures. The associated perpendicular magnetic fluctuations are linearly related with the perpendicular velocity fluctuations, but their relative amplitudes are not obligatorily equal (as is the case in an Alfv\u00e9n wave): \n\n\n\n\n\n, where \u03be is not necessarily equal to 1. In addition, Alfv\u00e9n vortices do not propagate along \n\n\n\n\n\n in the plasma frame, and they hardly propagate in the perpendicular plane when the axis of the vortex is inclined with respect to \n\n\n\n\n\n that are in contrast with Alfv\u00e9n wave (Wang et al. 2012). After first being reported in the Earth\u2019s magnetosheath (Alexandrova et al. 2006; Alexandrova 2008), multiscale quasi-bidimensional Alfv\u00e9n vortices (with \n\n\n\n\n\n) have been identified in numerous space environments: in slow solar wind (Perrone et al. 2016; Roberts et al. 2016), in fast solar wind (Lion et al. 2016; Perrone et al. 2017), and in Saturn\u2019s magnetosheath (Alexandrova & Saur 2008). It seems that the intermittent structures in fast solar wind are dominated by Alfv\u00e9n vortices (Perrone et al. 2017), which agrees with the 2D MHD turbulence model (Zank et al. 2017).","Citation Text":["Alexandrova 2008"],"Functions Text":["After first being reported in the Earth\u2019s magnetosheath"],"Functions Label":["Background"],"Citation Start End":[[1552,1568]],"Functions Start End":[[1470,1525]]}
{"Identifier":"2016ApJ...817..152X__Vreeswijk_et_al._2004_Instance_1","Paragraph":"The connection between long-duration GRBs (LGRBs) and SNe was predicted theoretically (Colgate 1974; Woosley 1993) and has been verified observationally (e.g., Galama et al. 1998; Hjorth et al. 2003; see a review in Woosley & Bloom 2006). They usually happen in the star formation regions of the galaxies (e.g., Paczy\u0144ski 1998; see the reviews in Woosley & Bloom 2006 and Kumar & Zhang 2015). The immensely bright afterglows illuminate the gas and dust within the star-forming regions of the host galaxy and intervening intergalactic medium along the GRB line of sight. Their spectra are usually featureless power-laws or broken power-laws, which can be well described by the synchrotron radiations of relativistic electrons. Therefore, GRB afterglows are good probes of burst environment and the interstellar dust and gas in distant, star-forming galaxies (Metzger et al. 1997; Jensen et al. 2001; Savaglio et al. 2003; Vreeswijk et al. 2004; Chen et al. 2005; Prochaska et al. 2007a; Schady et al. 2007, 2010; Watson et al. 2007; Fox et al. 2008; Starling et al. 2008; Jang et al. 2011; Xin et al. 2011). GRB afterglow spectra with Ly\u03b1 absorption features indicate the presence of large column densities of cold neutral gas within GRB host galaxies, and their hydrogen column densities (\n\n\n\n\n\n) are usually larger than \n\n\n\n\n\n cm\u22122 (e.g., Prochaska et al. 2007b; Schady 2012). The damped Ly\u03b1 systems (DLA) may represent the ISM near the GRBs in a few kiloparsecs (Kpc), but not gas directly local to the GRB (Prochaska et al. 2007b). Thus GRB optical afterglows may be used as probes of the ISM in their host galaxies, as the ISM observed is less affected by the GRB or its progenitor (Watson et al. 2007). The visual dust extinctions (\n\n\n\n\n\n) along the GRB lines of sight of many GRBs are low. As shown in Greiner et al. (2011), about 50% of GRBs observed with GROND after the launch of Swift mission have \n\n\n\n\n\n. In addition, the early optical light curves of about one-third of GRBs show a smooth onset bump (Li et al. 2012). It may be due to the deceleration of the GRB fireball by the ambient medium (Sari & Piran 1999; Kobayashi & Zhang 2007). In this scenario, the rising slope of the bump is determined by the medium density profile (\n\n\n\n\n\n) and the spectrum index of the accelerated electrons \n\n\n\n\n\n], says, \n\n\n\n\n\n (Liang et al. 2013). Hence, The afterglow onset bumps would be also an ideal probe to study the properties of the fireball and the profile of the circumburst medium. Liang et al. (2013) found that \n\n\n\n\n\n (see also Watson et al. 2007; Jin et al. 2012).","Citation Text":["Vreeswijk et al. 2004"],"Functions Text":["Therefore, GRB afterglows are good probes of burst environment and the interstellar dust and gas in distant, star-forming galaxies"],"Functions Label":["Motivation"],"Citation Start End":[[921,942]],"Functions Start End":[[726,856]]}
{"Identifier":"2017ApJ...835...94O__Martin_et_al._1994_Instance_1","Paragraph":"In this paper, we performed statistical analyses on the chirality and the magnetic configurations (inverse-polarity versus normal polarity) of the solar filaments that erupted on the solar disk from 2010 May 13 to 2015 December 31 covering both the rising phase and the beginning of the declining phases of solar cycle 24. The chirality is determined by an indirect method proposed by Chen et al. (2014), i.e., left-\/right-skewed drainage corresponds to dextral\/sinistral chirality. The determination of the magnetic configuration is also based on a method proposed by Chen et al. (2014), i.e., those filaments that follow Martin's Rule (Martin et al. 1994) are of the inverse-polarity type, and those that disobey Martin's Rule are of the normal-polarity type. By studying a sample of 571 filaments, we obtained the following results.\n\n(1)\nAbout 94.8% of the filaments in the northern hemisphere have negative helicity, and 87.4% of the filaments in the southern hemisphere have positive helicity, indicating a significant hemispheric preference of helicity. As a whole, 91.6% of our sample of erupting filaments follows the hemispheric rule of helicity sign. With the improved method for determining the filament chirality, the strength of the hemispheric rule is higher than that in previous studies. It should be noted that the statistical result is based on the erupting filaments. Those filaments that do not erupt during the disk passage are not included in our sample.\n\n\n(2)\nFollowing convention, we divided the filaments into three types, the quiescent type, the intermediate type, and the active-region type. It is shown that the strength of the hemispheric rule is 93% for the quiescent filaments, 95% for the intermediate filaments, and 83% for the active-region filaments.\n\n\n(3)\nRegarding the cyclic behavior of the hemispheric preference, it is found that the strength of the quiescent filaments decreases slightly from \u223c97% in the rising phase to \u223c85% in the declining phase, whereas the strength of the intermediate filaments keeps a high value around 96 \u00b1 4%. Only the active-region filaments show significant variations. Their strength of the hemispheric rule rises from \u223c63% to \u223c95% in the rising phase, and keeps a high value of 82% \u00b1 5% during the declining phase. However, during a half-year period around the solar maximum, the hemispheric preference totally vanishes.\n\n\n(4)\nIt is found that in our sample of erupting filaments, 89% are inverse-polarity filaments that are magnetically supported by a flux rope, whereas 11% are normal-polarity filaments that are magnetically supported by a sheared arcade.\n\n\n","Citation Text":["Martin et al. 1994"],"Functions Text":["The determination of the magnetic configuration is also based on a method proposed by Chen et al. (2014), i.e., those filaments that follow Martin's Rule","are of the inverse-polarity type, and those that disobey Martin's Rule are of the normal-polarity type."],"Functions Label":["Uses","Uses"],"Citation Start End":[[638,656]],"Functions Start End":[[483,636],[658,761]]}
{"Identifier":"2021ApJ...910...86R__Zitrin_et_al._2015_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":["Zitrin et al. 2015"],"Functions Text":["Complementing these observations, the spectroscopic confirmation","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."],"Functions Label":["Background","Background"],"Citation Start End":[[1064,1082]],"Functions Start End":[[948,1012],[1264,1484]]}
{"Identifier":"2021MNRAS.506...52O__Carroll,_Press_&_Turner_1992_Instance_1","Paragraph":"The SFG population bias factor, bSFG may be calculated from the ratio of galaxy to dark matter correlation functions, i.e:\n(27)$$\\begin{eqnarray*}\r\nb_{\\rm SFG}^2(z) &amp;=&amp; \\frac{\\xi _{\\rm g}(r, z)}{\\xi _{\\rm DM}(r, z)} \\nonumber \\\\\r\n&amp;=&amp; \\left(\\frac{r_0(z)}{8} \\right)^{\\iota } \\, \\frac{J_2}{\\sigma _8^2 \\,\\, \\mathcal {G}^2(z)}\r\n\\end{eqnarray*}$$(e.g. Kaiser 1984; Bardeen et al. 1986; Lindsay et al. 2014), where the matter fluctuation amplitude \u03c38 = 0.811 \u00b1 0.006 (Planck Collaboration 2020), and $\\mathcal {G}(z) = g(z)\/g_0$, with g(z) as the growth factor at redshift z and g0 = g(z = 0) (e.g. Carroll, Press & Turner 1992).7 Additionally, J2 = 72\/([3 \u2212 \u03b9][4 \u2212 \u03b9][6 \u2212 \u03b9]2\u03b9), and $r_0(z) = r^{\\rm c}_0 (1+z)^{p}$ with p = 1 \u2212 (3 + w)\/\u03b9 (Lindsay et al. 2014). Here, the choice of the parameter w reflects the clustering model adopted. In this demonstrative model, we consider only linear clustering (Overzier et al. 2003) where clustering growth is set by linear perturbation theory and w = \u03b9 \u2212 1. We leave the investigation of alternative clustering growth models to future work \u2013 for example, stable clustering (where clusters have a fixed physical size and w = 0), co-moving clustering (where clusters have fixed co-moving size and w = \u03b9 \u2212 3) and decaying clustering (which implies a rapid clustering decay) are also considered in the literature (Overzier et al. 2003; Kim et al. 2011; Elyiv et al. 2012). The remaining parameters in equation (27) are the power-law slope of the two-point correlation function of galaxies, \u03b9, and the galaxy clustering length $r_0^{\\rm c}$. Both of these may be estimated empirically for SFGs, and we adopt the best-fitting values of Hale et al. (2018): \u03b9 = 1.8 and $r_0^{\\rm c} = 6.1\\, \\text{Mpc}\\, h^{-1}$. These were computed from radio-selected SFGs in the COSMOS field using deep Karl G. Jansky Very Large Array (VLA) data at 3\u2009GHz, reaching redshifts as high as z \u223c 5, thus covering our range of interest (z \u2264 3).8 The resulting bias factor from these parameter choices is higher than those computed for SFGs at other wavelengths (e.g. Gilli et al. 2007; Starikova et al. 2012; Magliocchetti et al. 2013), but this is attributed to the greater extent of the redshift distribution of the sources.","Citation Text":["Carroll, Press & Turner 1992"],"Functions Text":["and $\\mathcal {G}(z) = g(z)\/g_0$, with g(z) as the growth factor at redshift z and g0 = g(z = 0) (e.g."],"Functions Label":["Uses"],"Citation Start End":[[610,638]],"Functions Start End":[[507,609]]}
{"Identifier":"2016AandA...592A..74S__Sobolewska_&_Papadakis_(2009)_Instance_1","Paragraph":"In Fig. B.1 we plot the soft X-ray light curves for our candidate highly variable AGN using available X-ray data taken by the satellite missions Einstein, ROSAT, XMM, Suzaku and Swift. The count rates were obtained from different archives including HEASARC, the XMM Science Archive, the Swift UKSSDC and from our own Swift XRT data analysis, and for upper limits the 1SXPS catalogue (Evans et al. 2014) and the XMM upper limit server2 were queried. The count rates of the different satellites were converted to fluxes between 0.2\u20132.0\u2009keV using PIMMS3 assuming a power law with a photon index of 1.7 as a spectral shape taking into account Galactic extinction as given by Willingale et al. (2013). Sobolewska & Papadakis (2009) found a positive correlation between flux and spectral slope for a sample of bright RXTE AGN in the 2\u201310\u2009keV band. This could affect the relative fluxes seen in our sample which are plotted in Fig. B.1. We have attempted to quantify this for the different detectors used in the creation of our light curves. The sample of Sobolewska & Papadakis (2009) showed spectral changes with observed power-law slope varying between 1.0 and 2.0 (see their Fig. 7). For a typical Galactic absorption of 3 \u00d7 1020 cm-2 the change from slope of 1.0 to 2.0 would alter our estimated fluxes by \u221214% (ROSAT), \u221213% (XMM-Newton), +7% (Swift-XRT), +76% (Suzaku), +25% (Einstein-IPC). The change is large for Suzaku observations since in this case we use the count rate between 2.0\u201310.0\u2009keV and extrapolate it to the soft band. All other satellites are sensitive in the soft band and hence the fluxes are less dependent upon the assumed spectral index. Six sources within our sample (XMMSL1\u2009J024916.6-041244, J034555.1-355959, J045740.0-503053, J051935.5-323928, J070841.3-493305, and J193439.3+490922) display variation in flux of a factor of ten or greater between at least one pair of XMM and Swift observations, on timescales of months to years. The ratio between the soft X-ray flux observed with Swift and that observed with XMM for the remaining sources is typically a factor of a few. We observed the two TDE candidates with XRT, and found that both had faded significantly, following expectations from previous and later fluxes and upper limits (Figs. 1p and h). ","Citation Text":["Sobolewska & Papadakis (2009)"],"Functions Text":["found a positive correlation between flux and spectral slope for a sample of bright RXTE AGN in the 2\u201310\u2009keV band. This could affect the relative fluxes seen in our sample which are plotted in Fig. B.1. We have attempted to quantify this for the different detectors used in the creation of our light curves."],"Functions Label":["Uses"],"Citation Start End":[[697,726]],"Functions Start End":[[727,1034]]}
{"Identifier":"2016ApJ...826..117Y__Roux_&_Webb_2009_Instance_1","Paragraph":"Traditionally, the Parker transport equation (Parker 1965) was used to model pickup ion acceleration at the SWTS when using a transport theory approach. However, interesting Voyager results, such as strongly fluctuating pitch-angle anisotropies upstream, the detection of highly anisotropic intensity spikes at the SWTS, the average upstream anisotropy peaking at a surprisingly high energy far above the upstream flow energy, and energetic particle spectra with multiple power-law slopes with breaking points in between that are highly variable upstream (Decker et al. 2005, 2008b; Cummings et al. 2006), suggest that transport modeling should be modified in at least two ways. First, the turbulent nature of the magnetic field conditions at the SWTS should be taken into account, and second, a transport theory that applies when particle distributions are highly anisotropic is needed, given that the Parker transport equation only holds for nearly isotropic particle distributions. In response, shock acceleration transport models were developed in recent years based on the numerical solutions of the focused transport equation (K\u00f3ta & Jokipii 2004; le Roux et al. 2007; Florinski et al. 2008a, 2008b; le Roux & Webb 2009) to take advantage of the fact that focused transport is not restricted to small pitch-angle anisotropies. This is especially advantageous at lower suprathermal particle energies upstream, where particle distributions can be sporadically strongly anisotropic (Decker et al. 2006), allowing one to model particle injection into diffusive shock acceleration (DSA) naturally at those energies. Furthermore, statistical variations in the observed magnetic field direction near the SWTS were included as a time series to model time variations in the injection and DSA of pickup ions to simulate the highly volatile nature of actual DSA at the SWTS (Florinski et al. 2008a, 2008b; le Roux & Webb 2009; Arthur & le Roux 2013). This focused transport approach should be seen as complementary to more sophisticated self-consistent shock acceleration models based on hybrid codes (Kucharek & Scholer 1995; Giacalone 2005) and particle-in-the-cell models (Scholer et al. 2003; Lembege et al. 2004), but it has the virtue of relative simplicity because different statistical plasma parameters can easily be studied separately and in combination at the SWTS to gain a more clear conceptional understanding of the role of such statistics on pickup acceleration at the SWTS.","Citation Text":["le Roux & Webb 2009"],"Functions Text":["In response, shock acceleration transport models were developed in recent years based on the numerical solutions of the focused transport equation","to take advantage of the fact that focused transport is not restricted to small pitch-angle anisotropies."],"Functions Label":["Background","Background"],"Citation Start End":[[1206,1225]],"Functions Start End":[[985,1131],[1227,1332]]}
{"Identifier":"2017MNRAS.470..713M__Farinelli_et_al._2012_Instance_1","Paragraph":"Thereafter, we proceeded to fit the broad-band spectrum with the analytical model \u2018COMPMAG\u2019 and its updated version. The \u2018COMPMAG\u2019 model has many spectral parameters that were impossible to constrain simultaneously. Especially, \u03b2 if left free preferred a value >0.9 in the fit, which is greater than the maximum possible terminal velocity at the neutron star surface. Therefore, we tested the two-component \u2018COMPMAG\u2019 model, assuming reasonable parameter values applicable to low-luminosity accretion-powered pulsars, i.e. free-fall velocity profile (\u03b7 = 0.5), and \u03b2 = 0.5 that corresponds to the maximum terminal velocity at the neutron star surface, and a pencil beam emission pattern. The radius of the accretion column (r0) was set to 0.25 corresponding to a radius of \u223c1 km for a neutron star of mass 1.4 M\u2299. The Albedo at the neutron star surface (A) was set to 1. This provided an acceptable fit to the data with acceptable and physically viable parameter values. Addition of a CRSF feature was not required for this model. The high \u03b2 preferred by the fit indicated that bulk motion Comptonization (BMC) dominates the Comptonization process in X Persei. The fit also provided a low value of electron temperature (kT) consistent with the high value of \u03b2 obtained (Farinelli et al. 2012). The dominance of BMC in low-luminosity pulsars, especially X-Persei has been predicted before and is consistent with our results (Becker & Wolff 2005, 2007). In Becker & Wolff (2007), the authors qualitatively described the observed spectrum of X-Persei with a pure BMC model. The obtained reduced \u03c72 value with \u2018COMPMAG\u2019 was higher than that obtained with the \u2018newhcut\u2019 model (1.84 for 569 dof), although there were no systematic pattern in the residuals, and the higher \u03c72 was mainly contributed by a higher variance in a few energy bins of the XIS spectra. Fit with the updated \u2018COMPMAG2\u2019 model was difficult with the given statistical quality of the data, as it required fitting many more additional parameters. Moreover as discussed in Farinelli et al. (2016), the \u2018COMPMAG\u2019 model provides an adequate description of the spectrum for low-luminosity pulsars, where the blackbody emission provides the major source of photons for Comptonization, appropriate for the case of X Persei. Considering the above, and as \u2018COMPMAG\u2019 is a more physical model to understand the continuum spectra of the source, we used this as our best-fitting model for the rest of the paper. As a CRSF feature is not required for this model, we cannot claim the presence of a CRSF in the average spectrum of X Persei. Table 1 summarizes the best-fitting broad-band spectral parameters obtained using the \u2018newhcut\u2019 (without addition of the CRSF) and \u2018COMPMAG\u2019 models. Fig. 3 shows the best-fitting unfolded spectrum for the best-fitting model, showing the model components.","Citation Text":["Farinelli et al. 2012"],"Functions Text":["The fit also provided a low value of electron temperature (kT) consistent with the high value of \u03b2 obtained"],"Functions Label":["Similarities"],"Citation Start End":[[1269,1290]],"Functions Start End":[[1160,1267]]}
{"Identifier":"2021AandA...650A..56R__Johnson_et_al._2017_Instance_1","Paragraph":"The GRMHD model library is limited in that it has only a few discrete parameter values for the magnetization, black hole spin, and electron temperature distributions. Apart from increasing the number of discrete values, one could think about ways to interpolate between these using, for example, machine learning techniques (van der Gucht et al. 2020; Yao-Yu Lin et al. 2020) or fit to semi-analytic models (e.g., Broderick et al. 2016). The variability within the GRMHD models was found to be an important limitation for constraining black hole parameters, as attested by, for example, the small difference in recovered parameters between the EHT2021 and EHT2021+ arrays. The analysis pipeline may be extended to include a characterization of the source variability as part of the model selection process (e.g., Kim et al. 2016; Roelofs et al. 2017; Medeiros et al. 2017, 2018; Johnson et al. 2017; Bouman et al. 2018; Wielgus et al. 2020), which could improve the constraining power beyond the averaging method introduced here. EHT expansions are expected to make the large-scale jet visible in reconstructed images of the black hole shadow due to an increased dynamic range (Doeleman et al. 2019; Roelofs et al. 2020; Raymond et al. 2021). This connection between event-horizon scales and the extended jet has not been taken into account in the parameter estimation framework used here, as the GRMHD library images have a small field of view (160 \u03bcas). With the development of GRMHD simulations that have the ability to connect large (e.g., Fromm et al. 2017, 2018, 2019; Liska et al. 2018; Chatterjee et al. 2019) and small scales at different wavelengths and of an extended fitting framework, the constraining power is expected to improve especially for EHT extensions and space arrays. For a mass measurement, feature extraction techniques such as a ring fit (Event Horizon Telescope Collaboration 2019d,f) may be used, potentially in combination with fitting the more extended (variable) structure (Broderick et al. 2020b). Models and analysis techniques for Sgr A* and polarization could be considered as well. These possible avenues for further simulation and fitting framework development mean that the parameter constraints presented in this paper should not be interpreted as set limits on the constraining power of the considered arrays. Rather, they show what is achievable with the current state of the art.","Citation Text":["Johnson et al. 2017"],"Functions Text":["The variability within the GRMHD models was found to be an important limitation for constraining black hole parameters, as attested by, for example, the small difference in recovered parameters between the EHT2021 and EHT2021+ arrays. The analysis pipeline may be extended to include a characterization of the source variability as part of the model selection process (e.g.,","which could improve the constraining power beyond the averaging method introduced here."],"Functions Label":["Future Work","Future Work"],"Citation Start End":[[879,898]],"Functions Start End":[[438,812],[942,1029]]}
{"Identifier":"2015MNRAS.451..353R__Naray_et_al._2006_Instance_1","Paragraph":"The \u039b cold dark matter(\u039bCDM) paradigm (Blumenthal et al. 1984) predicted cuspy DM density profiles in the centre of galaxies (Navarro, Frenk & White 1996) and recent results with higher resolution seem to suggest a shallower DM density profile (Navarro et al. 2004, 2010). None the less, the past and newer developments are in contradiction with observational results of cored DM density profiles in dwarf galaxies (de Blok, McGaugh & Rubin 2001; de Blok & Bosma 2002; Swaters et al. 2003; Chemin, de Blok & Mamon 2011). This discrepancy is known as the cuspy-core problem. Many other authors, in an attempt to contribute to the solution of the cuspy-core problem, performed the RC decomposition considering a constant M\/L (Carignan & Freeman 1985, 1988; Jobin & Carignan 1990; Persic & Salucci 1990; Martimbeau, Carignan & Roy 1994; Persic, Salucci & Stel 1996; Blais-Ouellette, Amram & Carignan 2001; de Blok, McGaugh & Rubin 2001; de Blok & Bosma 2002; Swaters et al. 2003; Kuzio de Naray et al. 2006; Fuentes-Carrera et al. 2007; Spano et al. 2008; Kuzio de Naray, McGaugh & de Blok 2008; Repetto et al. 2010; Chemin et al. 2011) and reinforce the evidence of a cored DM distribution in the central part of galaxies. Conversely, other authors pursue spectrophotometric and SPS studies to derive the M\/L of the stellar component (Salucci, Yegorova & Drory 2008; de Denus-Baillargeon et al. 2013; Repetto et al. 2013) or at least to constrain the stellar component (Dutton et al. 2005) to avoid the disc\u2013halo degeneracy. Salucci et al. (2008) accomplished RCs mass modelling and spectral energy distribution (SED) fitting with spectrophotometric models to obtain the disc masses of 18 Sa spiral galaxies (principally bulgeless systems) finding that by decomposing the RCs with the spectrophotometric disc masses the results are consistent with the corresponding maximum disc solution. Repetto et al. (2013) employed the Zibetti et al. (2009) method to optical (SDSS) and NIR (2MASS) images of NGC 5278 (KPG 390A) to obtain the stellar disc mass profile of that galaxy from SPS models to reproduce the RC of KPG 390A. The new strategy relied on fitting only the DM RC, obtained by subtracting the SPS baryonic disc from the observed RC of NGC 5278. The most important finding of Repetto et al. (2013) is that the favoured DM distribution is cored when the disc mass approximate to the maximum disc solution in agreement with the general belief. de Denus-Baillargeon et al. (2013) used a chemo-spectrophotometric galactic evolution model to determine the stellar M\/L and perform the RC decomposition of 10 spiral and dwarf irregular galaxies from the SINGS survey (Kennicutt et al. 2003). The authors employed the settled baryonic disc as a weighting function to fit the model DM halo to the RCs of the studied subsample of galaxies, and concluded that the stellar discs obtained from their chemo-spectrophotometric models were compatible with the MDH. The few examples presented indicate that a growing effort to break the disc\u2013halo degeneracy exists; however, the most significant contributions trying to address the cuspy-core problem still rely on the general assumption of considering the stellar M\/L constant along the galactic disc. In general, it is still missing a significant endeavour to earn the DM distribution, through RC fitting, determining the disc stellar mass from SPS studies. For this reason, it is worth considering the cuspy-core problem with a different observational approach, focusing on the formulation of a general procedure to better constrain the baryonic disc mass in galaxies.","Citation Text":["Kuzio de Naray et al. 2006"],"Functions Text":["Many other authors, in an attempt to contribute to the solution of the cuspy-core problem, performed the RC decomposition considering a constant M\/L"],"Functions Label":["Background"],"Citation Start End":[[977,1003]],"Functions Start End":[[574,722]]}
{"Identifier":"2016ApJ...826..168X__Ilgner_&_Nelson_2008_Instance_1","Paragraph":"MRI is considered to be the most promising mechanism driving angular-momentum transport in protoplanetary disks (Balbus & Hawley 1991; Brandenburg et al. 1995; Hawley et al. 1995; Balbus et al. 1996; Balbus & Hawley 1998). However, protoplanetary disks are cold, dense, and, therefore, poorly ionized. The low level of ionization tends to decouple the disk gas from magnetic fields, which generates non-ideal MHD effects: Ohmic dissipation, ambipolar diffusion (AD), and the Hall effect (e.g., Armitage 2011; Turner et al. 2014). These effects quench MRI in different ways: Ohmic dissipation originates from collisions between electrons and neutrals, AD from collisions between ions and neutrals, and the Hall effect from drift between electrons and ions (Fleming et al. 2000; Sano & Stone 2002; Bai & Stone 2011). Ohmic dissipation operates in high-density regions with weak field, AD dominates in highly ionized and low-density regions, and the Hall effect lies in between (Fleming & Stone 2003; Bai & Stone 2011; Bai 2014). So far, the effect of Ohmic dissipation has been best studied. Investigations show the layered accretion in the inner disk, where the midplane region is \u201cdead\u201d due to low ionization while the surface layer is \u201cactive\u201d due to sufficient ionization (Gammie 1996; Jin 1996; Fleming et al. 2000; Fleming & Stone 2003; Turner et al. 2007; Ilgner & Nelson 2008; Oishi & Mac Low 2009; Okuzumi & Hirose 2011). Recent works that take into account both Ohmic dissipation and AD show that AD may render the surface layer and portions of the outer disk inactive (Bai & Stone 2011; Landry et al. 2013; Kalyaan et al. 2015). Bai & Stone (2013) find that MRI is completely suppressed in the inner disk and a strong magnetocentrifugal wind is launched. Three-dimensional simulations that include all three non-ideal MHD effects are also performed (Bai 2014, 2015; Lesur et al. 2014; Simon et al. 2015). In the inner disk, the influence of the Hall effect on midplane angular-momentum transport depends on the orientation of the vertical magnetic field with the disk rotation axis. When the field is aligned with the axis, the enhanced Maxwell stress promotes angular-momentum transport. When the field is anti-aligned with the axis, the midplane remains quiescent. In the outer disk, the Hall effect has little influence on the disk turbulence. Although the inclusion of AD and the Hall effect substantially changes the level of turbulence in the protoplanetary disks, the feature that the viscosity is low in the inner disk and high in the outer disk is still valid. In this study, we assume that gas giant planets form in situ via the core accretion scenario, which implies that their formation locations are always in the low-viscosity region. Since in this study we focus on the relation between photoevaporation and planet formation and gap opening by planets in the disk, we adopt Ohmic dissipation to represent the non-ideal MHD effects on the MRI. We consider that this simplification has little influence on our main calculation results.","Citation Text":["Ilgner & Nelson 2008"],"Functions Text":["Investigations show the layered accretion in the inner disk, where the midplane region is \u201cdead\u201d due to low ionization while the surface layer is \u201cactive\u201d due to sufficient ionization"],"Functions Label":["Background"],"Citation Start End":[[1361,1381]],"Functions Start End":[[1090,1273]]}
{"Identifier":"2017ApJ...846...28R__Kopparapu_et_al._2013_Instance_1","Paragraph":"A question of relevance to life on other planets is whether the Snowball catastrophe occurs at lower radiative forcing at low or high obliquity. There is some conflict in the literature on the role of the Snowball bifurcation in planetary habitability. On the one hand, defining habitability in terms of surface liquid water suggests that a Snowball climate is uninhabitable, which has led various authors to propose metrics of fractional or seasonal habitability for planets with partial ice cover (e.g., Williams & Pollard 2003; Spiegel et al. 2008). On the other hand, not only did photosynthesis persist through Snowball events in Earth\u2019 history, but the events may have crucially shaped the subsequent evolution of complex life (e.g., Hoffman & Schrag 2002; Hoffman et al. 2017; Laasko & Schrag 2017). The traditional habitable zone concept assumes a planet with a functioning silicate-weathering feedback and a positive CO2 greenhouse effect (Kasting et al. 1993; Kopparapu et al. 2013). Global glaciation may be triggered on such a planet through a rapid drawdown of atmospheric CO2 that reduces q below the thresholds at which the non-Snowball states disappear. In Earth history, this seems to have occurred through accidents of tectonics (Hoffman et al. 2017). These events are self-terminating through the suppression of silicate weathering and accumulation of a strong CO2 greenhouse. Such transitions can in principle occur anywhere within the habitable zone. Our simple model is unsuited to the tasks of diagnosing the inner edge of habitability (where the relevant physics are the runaway water vapor greenhouse and hydrogen loss to space) or the outer edge (where the relevant physics are CO2 condensation and Rayleigh scattering). However, we can compute the q value at which the bifurcation occurs as a function of obliquity and other model parameters, and whether the transition into the Snowball state occurs from a partially ice-covered state (cap or belt) or directly from the ice-free state.","Citation Text":["Kopparapu et al. 2013"],"Functions Text":["The traditional habitable zone concept assumes a planet with a functioning silicate-weathering feedback and a positive CO2 greenhouse effect","Global glaciation may be triggered on such a planet through a rapid drawdown of atmospheric CO2 that reduces q below the thresholds at which the non-Snowball states disappear."],"Functions Label":["Background","Background"],"Citation Start End":[[970,991]],"Functions Start End":[[807,947],[994,1169]]}
{"Identifier":"2022MNRAS.515...22J__Newman_et_al._2013_Instance_1","Paragraph":"In Fig. 5, we consider how the velocity dispersion profile scales with radius. Specifically, we plot the power-law index (\u03b7) versus the central velocity dispersion (\u03c30). The vast majority of the galaxies with \u03c30 \u2272 2.45 are BGGs and these have negative \u03b7 values. This includes most of the Romulus galaxies (red filled and open circles), the L18 BGGs (blue crosses) and the early-type galaxies that comprise the SAURON sample (Cappellari et al. 2006; grey line and shaded area). In contrast, nearly all of simulated BGGs with \u03c30 \u2272 2.45) from the DIANOGA Hydro-10x simulations Marini et al. (2021) have positive \u03b7 values. For \u03c30 \u2273 2.45, the spread of \u03b7 for the observed galaxies (e.g. L18 and Newman et al. 2013 BCGs) broadens and spans both positive and negative \u03b7 values. In fact, majority of the galaxies tend to have positive \u03b7s. This change in behaviour is well known. A number of studies have noted that on the group-scale and lower, the stellar velocity dispersion profile of the central galaxies tend to decrease with increasing radius. On the cluster-scale, the BCGs typically have rising velocity dispersion profiles with increasing radius (Von Der Linden et al. 2007; Bender et al. 2015; Veale et al. 2017). The origin of this flip is still not well understood. We leave a more detailed investigation of this change to future work. Here, we simply mention two possible explanations: The change in slope may be a reflection of the differences in the dynamical state (e.g. mass-to-light ratio; M\/L) at the outskirts of BCGs (Dressler 1979; Fisher, Illingworth & Franx 1995; Sembach & Tonry 1996; Carter et al. 1999; Kelson et al. 2002; Loubser et al. 2008; Newman et al. 2013; Schaller et al. 2015; Marini et al. 2021), or it could be due to increased contribution from the intragroup\/intracluster light along the line-of-sight and the increased leverage of tangential orbits (Loubser et al. 2020). All of these effects are linked to the increased frequency of galaxy\u2013galaxy interactions and more specifically, central-satellite interactions, implicated in the build-up of extended diffuse stellar component. And, as discussed by Schaye et al. (2015), Oppenheimer et al. (2021), and the EAGLE simulations clearly show that the extended stellar halo becomes increasingly more important, and hosts a non-trivial fraction of the total stellar mass towards the cluster scale.","Citation Text":["Newman et al. 2013"],"Functions Text":["For \u03c30 \u2273 2.45, the spread of \u03b7 for the observed galaxies (e.g.","broadens and spans both positive and negative \u03b7 values. In fact, majority of the galaxies tend to have positive \u03b7s. This change in behaviour is well known."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[690,708]],"Functions Start End":[[619,681],[715,870]]}
{"Identifier":"2018AandA...616A..72W__K\u00e4pyl\u00e4_et_al._2012_Instance_1","Paragraph":"The Sun, our nearest late-type star exhibits a magnetic activity cycle with a period of around 11 yr. The cyclic magnetic field is generated by a dynamo operating below the surface, where it converts the energy of rotating convective turbulence into magnetic energy. We are still far from fully understanding the solar dynamo mechanism (e.g., Ossendrijver 2003; Charbonneau 2014). One reason is the limited information about the dynamics in the solar convection zone. Helioseismology has provided us with the profile of temperature and density stratification and the differential rotation (e.g., Schou et al. 1998) in the interior. Further information, such as the meridional circulation profiles, convective velocity strength, or even magnetic field distributions, are currently inconclusive or not even possible (e.g., Basu 2016; Hanasoge et al. 2016). One way to investigate how important differential rotation, meridional circulation, and turbulent convective velocities are for the solar dynamo is to use numerical simulations. Since the early simulations by Gilman (1983), the increase of computing resources has prompted several advances using numerical simulations. Nowadays, global simulations of convective dynamos are able to reproduce cyclic magnetic fields and dynamo solutions resembling many features of the solar magnetic field evolution (Ghizaru et al. 2010; K\u00e4pyl\u00e4 et al. 2012; Warnecke et al. 2014; Augustson et al. 2015), even long-time evolution (Augustson et al. 2015; K\u00e4pyl\u00e4 et al. 2016; Beaudoin et al. 2016). The cyclic magnetic field in these simulations can be well understood in terms of Parker\u2013Yoshimura rule (Parker 1955; Yoshimura 1975; Warnecke et al. 2014), in which a propagating \u03b1\u03a9 dynamo wave is excited; see also Gastine et al. (2012). The \u03b1 effect (Steenbeck et al. 1966) describes the magnetic field enhancement from helical turbulence and the \u03a9 effect the shearing of magnetic field caused by differential rotation. The propagation direction of the dynamo wave depends on the sign of \u03b1 and shear: to generate an equatorward propagating wave, the product of \u03b1 and the radial gradient of \u03a9 must be negative (positive) in the northern (southern) hemisphere. Explaining the solar equatorward propagation of the sunspot appearance by the Parker\u2013Yoshimura rule therefore requires either invoking the near-surface-shear layer (Brandenburg 2005), because the radial gradient is only negative in that layer (Barekat et al. 2014) and \u03b1 is positive, or changing the sign of \u03b1 in the bulk of the convection zone (Duarte et al. 2016) where the radial shear is positive. Furthermore, to understand the magnetic field evolution fully in the global numerical simulation one needs suitable analysis tools to extract the important contribution of turbulent dynamo effects. One of these tools is the test-field method (Schrinner et al. 2005, 2007; Warnecke et al. 2018). This method allows the determination of the turbulent transport coefficients directly from the simulations. This includes the measurement of tensorial coefficients such as \u03b1, turbulent pumping, and turbulent diffusion. The first application to global convection simulations of solar-like dynamo has already revealed that the turbulent effects can have a significant impact on large-scale magnetic field dynamics (Warnecke et al. 2018; Gent et al. 2017).","Citation Text":["K\u00e4pyl\u00e4 et al. 2012"],"Functions Text":["Nowadays, global simulations of convective dynamos are able to reproduce cyclic magnetic fields and dynamo solutions resembling many features of the solar magnetic field evolution"],"Functions Label":["Similarities"],"Citation Start End":[[1376,1394]],"Functions Start End":[[1174,1353]]}
{"Identifier":"2021ApJ...909...18F__Hon_et_al._2020_Instance_1","Paragraph":"The mechanisms of CL-AGNs are still debated. Early explanations mainly focus on the change in obscuration (Bianchi et al. 2005), while some recent studies favor the change of accretion rate (Stern et al. 2018; Sniegowska et al. 2020). Photoionization research shows that the CL behavior in CL quasars can be fully explained by the photoionization responses of the BELs to the extreme variability of the ionizing continuum (Guo et al. 2020). Besides, several objects can be interpreted as tidal disruption events (Merloni et al. 2015; Ricci et al. 2020). All of the models can generate significant variance in the observed intensity of continuum and emission lines. These models might be restricted by statistical studies after extending the sample of CL-AGNs. Several groups have been devoted to searching for new CL-AGNs, and tens of candidates have been discovered based on large optical and X-ray surveys, such as the Sloan Digital Sky Survey (e.g., LaMassa et al. 2015; MacLeod et al. 2016; Rumbaugh et al. 2018; Hon et al. 2020), intermediate Palomar Transient Factory (Gezari et al. 2017), Pan-STARRS1 (MacLeod et al. 2019), Catalina Real-time Transient Survey (Yang et al. 2018), and XMM-Newton slew survey (Zetzl et al. 2018; Kollatschny et al. 2020). In the future, the Large Synoptic Survey Telescope will regularly monitor millions of AGNs, and the number of CL-AGNs will further increase. The evolution of the AGN type can be well studied if we obtain continuous multiwavelength data of the complete changing process, and the development of time-domain surveys makes it possible to achieve this. The brightening of NGC 2617 (Shappee et al. 2014) and 1ES 1927+654 (Trakhtenbrot et al. 2019) triggered the alert of the All-Sky Automated Survey for Supernovae,7\n\n7\n\nhttp:\/\/www.astronomy.ohio-state.edu\/asassn\n\n and follow-up observations confirmed these two CL-AGNs. However, so far, there are only small available data sets for most detected CL-AGNs, especially around the critical regime of transition between types 1 and 2. Moreover, previous studies rarely gave the change of geometry and kinematics in broad-line regions (BLRs).","Citation Text":["Hon et al. 2020"],"Functions Text":["Several groups have been devoted to searching for new CL-AGNs, and tens of candidates have been discovered based on large optical and X-ray surveys, such as the Sloan Digital Sky Survey (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[1017,1032]],"Functions Start End":[[760,952]]}
{"Identifier":"2021ApJ...909...86X__Mou\u00ebl_et_al._2017_Instance_1","Paragraph":"The level of solar activity is typically represented by the number or area of sunspots. On the other hand, the \u201crush-to-the-pole\u201d behavior of PCFs is believed to be one of the precursors for the next solar maximum (Altrock 2014). Figure 3 shows the maximum sunspot numbers (SILSO World Data Center 1914\u20132020) in comparison with the migration speed of PCFs (orange dots) in the northern and southern hemisphere for cycles 16\u201324. Within these nine solar cycles, the maximum sunspot number varies in two large periods, cycles 16\u201320 and 20\u201324. Note that even in longer time spans, such as from cycle 1 to 24, the sunspot number seems to vary periodically in different periods, like the often-discussed Gleissberg cycle (Le Mou\u00ebl et al. 2017). The migration speeds of PCFs in the northern hemisphere can be separated into two groups, before and after cycle 20, which is similar to the trend found in the variation of sunspot numbers. However, the migration speed in the southern hemisphere is not correlated with the sunspot number and the migration speed in the northern hemisphere. As shown in Table 2, the Pearson correlation coefficients are below 0.5 in both hemispheres, in which the value for the northern hemisphere is relatively higher. To investigate the trend of temporal variation, the derivatives of migration speeds and sunspot number are calculated. The correlation between these derivatives represents the relationship between the temporal trends. It shows that the periodical variation of PCF migration in the northern hemisphere correlates better with the variation of sunspot numbers, but the absolute values of coefficients are not high (\u223c0.5). Based on this result, a causal connection between the temporal variations of PCF migration and the maximum number of sunspots cannot be established. In the SILSO World Data Center (1914\u20132020) archive of sunspot numbers, clear NS asymmetry after cycle 20 is noticed. But the overall levels of the sunspot numbers are the same in the northern and southern hemispheres. Therefore, the NS asymmetry, which is relatively small compared to the absolute sunspot number, will not affect the results shown in Figure 3.","Citation Text":["Le Mou\u00ebl et al. 2017"],"Functions Text":["Note that even in longer time spans, such as from cycle 1 to 24, the sunspot number seems to vary periodically in different periods, like the often-discussed Gleissberg cycle"],"Functions Label":["Background"],"Citation Start End":[[716,736]],"Functions Start End":[[540,714]]}
{"Identifier":"2021ApJ...912..106Y__Minchev_et_al._2013_Instance_2","Paragraph":"Our analysis on the LAMOST-RC stars by dissecting the MAPs shows that the chemical bimodality is observed throughout the Galactic disk, and the high- and low-[\u03b1\/Fe] sequences are corresponding to the thick and thin disks of the Milky Way, respectively. How to explain the formation mechanism of the stellar thin and thick disks is beyond the scope of this paper, but our results provide some observational constraints to the model of the chemodynamical evolution of the Milky Way disk. Our flared vertical profiles for the thin and thick disks are in good agreement with the prediction of the thin+thick flaring disk model (L\u00f3pez-Corredoira & Molg\u00f3 2014), and are consistent with the number simulations of the chemodynamical evolution in Galactic disks formed in the cosmological context (Minchev et al. 2013, 2014, 2015, 2017), as well as the cosmological zoom simulation of VINTERGATAN (Agertz et al. 2021). These model simulations suggest that the vertical flaring trends are a natural consequence of inside-out, upside down growth coupled with disk flaring (see also Bird et al. 2013; Garc\u00eda de la Cruz et al. 2021), which allows for the low-[\u03b1\/Fe] stars to exist several kpc above the disk\u2019s midplane. As analyzed by B16, the exponential flaring profiles for the low-[\u03b1\/Fe] MAPs suggests that radial migration played an important role in the formation and evolution of the thin disk. Radial migration of stars via cold torquing, also known as \u201cchurning,\u201d by a bar and spiral waves (Minchev et al. 2013) then allows for the populations to spatially overlap in the solar neighborhood. Similar to the flared thin disk, the flaring profile for the high-[\u03b1\/Fe] MAP indicates the radial migration has occurred in the formation of the thick disk as suggested by model simulations (e.g., Sch\u00f6nrich & Binney 2009; Minchev et al. 2015; Li et al. 2018). Of course, we cannot rule out the other formation scenarios of the thick disk, such as the accreted gas from satellites (Brook et al. 2004), accreted stars from galaxy mergers (Abadi et al. 2003), or from disk-crossing satellites heating up the thin disk (Read et al. 2008). On the other hand, the broken exponential radial profiles for the thin and thick disks cannot be explained by any model of the galactic disks. In fact, nearly all the models we mentioned above present a single-exponential profile decreasing with the increasing of R (e.g., Minchev et al. 2015; Li et al. 2018; Agertz et al. 2021). And the smooth downtrend of radial profile in the outer disk R > Rpeak, as shown in Figure 10, means that there is no cut-off of the stellar component at R = 14\u201315 kpc as stated by Ruphy et al. (1996), which is also discovered by B16.","Citation Text":["Minchev et al. 2013"],"Functions Text":["Radial migration of stars via cold torquing, also known as \u201cchurning,\u201d by a bar and spiral waves","then allows for the populations to spatially overlap in the solar neighborhood."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1487,1506]],"Functions Start End":[[1389,1485],[1508,1587]]}
{"Identifier":"2021AandA...653A..32D__Kriek_et_al._2006_Instance_1","Paragraph":"The choice of SFH to infer evolutionary parameters intrinsically carries a degeneracy with the adopted functional form. A more direct approach is to quantify the light-weighted contribution of recent star formation by measuring the relative contribution to the integrated stellar spectrum of short-lived massive stars with respect to long-lived lower-mass stars. Balmer absorption lines reach their maximum strength in A-type stars with a spectral break at 3646 \u00c5. Stars of lower mass and lower effective temperature produce metal absorption lines (CaII H & K, Fe, and Mg) which result in a sharp spectral break at 4000 \u00c5. Moreover, the underlying continuum changes shape with time, progressively losing emission in the NUV-blue spectral range while flattening in the NIR. The different evolutionary rates of the stars producing the lines and their fractional contribution to the optical light at fixed mass enable tracing the evolutionary stage of a galaxy. In Fig. 11 (upper panels) we show the evolution of the spectral break measured through the DB definition (Kriek et al. 2006) and the Dn4000 definition (Balogh et al. 1999), as well as the relative strength (the ratio) between the two (lower left panel). Lighter shaded curves show the variation with increasing duration of star formation. The ratio is shown as a function of age of composite templates built with a short truncated SFH. The ratio is only mildly dependent on dust reddening because the two indices share a similar wavelength range. Moreover, the two indices are fairly robust against low resolution. The ratio varies strongly during the first 1 Gyr or so, reaching its maximum around 0.3\u20130.5 Gyr. Eventually, it drops below 1 when the light-weighted contribution from A-type stars fades away. Constant star formation instead results in a ratio of \u223c1.1 that is rather constant with time. Varying the metallicity of the input templates has the effect of anticipating the transition to DB\/Dn4000\u2004>\u20041. This effect is strongest when supersolar metallicity (2.5Z\u2299) is adopted in which case the transition is reached at 0.9 Gyr. We suggest that this ratio could be used to spot PSB galaxies with high dust attenuation along and across the UVJ diagram when high-resolution spectra are unavailable.","Citation Text":["Kriek et al. 2006"],"Functions Text":["In Fig. 11 (upper panels) we show the evolution of the spectral break measured through the DB definition","as well as the relative strength (the ratio) between the two (lower left panel)."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1065,1082]],"Functions Start End":[[959,1063],[1132,1212]]}
{"Identifier":"2016AandA...588A...2L__M\u00e4tzler_(1998)_Instance_2","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"],"Functions Text":["We use the Mishima et al. (1983) formulation (see Appendix of"],"Functions Label":["Uses"],"Citation Start End":[[762,774]],"Functions Start End":[[700,761]]}
{"Identifier":"2018MNRAS.477..957D__Slee_et_al._2001_Instance_1","Paragraph":"The role of cluster mergers in producing shocks underlying radio relics at cluster peripheries is well supported by observational evidence (Giovannini, Tordi & Feretti 1999; Bagchi et al. 2006; van Weeren et al. 2009; Venturi et al. 2013). A different class of relics, proposed to be fading radio galaxy lobes or such lobes revived due to adiabatic compression (En\u00dflin & Gopal-Krishna 2001), are also well known to occur in clusters. Examples of relics in this category are those in Abell 4038 (Slee & Roy 1998; Kale & Dwarakanath 2012) and Abell 85 (Slee et al. 2001). In A168 we find an arc-like relic at the periphery (B) and a smaller steep-spectrum relic (A) in its wake, as seen projected in the plane of the sky. The intracluster medium (ICM) in the cluster is elongated along the north\u2013south direction and the orientation of B is perpendicular to this direction. Merger-shock-related relics have been found to be oriented preferentially perpendicular to the elongation axis of the ICM (van Weeren et al. 2011). The flat spectral index of B, arc-like morphology and orientation thus support the scenario that a cluster merger along the north\u2013south direction led to an outgoing merger shock that accelerated electrons, which are detected as relic B. Relic A has a steep and curved spectrum (Fig. 3) indicative of an ageing population of relativistic electrons. We propose that relic A is a candidate adiabatically compressed lobe of a radio galaxy, the compression having been caused by the outgoing shock at B. Simulations of adiabatically compressed cocoons of radio galaxies have shown that, as compression proceeds, the cocoon is torn into filamentary structures, which can appear like a single torus or multiple tori at late stages (En\u00dflin & Br\u00fcggen 2002). The morphology of relic A is complex and filamentary, which compares well with the predicted structures of compressed cocoons. The morphology of B is not smooth at the outer boundary, but shows kink-like features connected to A, implying a possible distortion due to the presence of a radio cocoon in the path of the outgoing shock that led to the formation of A.","Citation Text":["Slee et al. 2001"],"Functions Text":["Examples of relics in this category are those in","and Abell 85"],"Functions Label":["Background","Background"],"Citation Start End":[[551,567]],"Functions Start End":[[434,482],[537,549]]}
{"Identifier":"2016ApJ...830...15J__Mathieu_et_al._1997_Instance_1","Paragraph":"If PTFO 8-8695 is accreting material onto the star from a gas disk devoid of small grains, the excess H\u03b1 emission may result entirely from the accreting material whether or not there is a planetary companion present. This accretion-related emission would presumably be similar to H\u03b1 emission seen in other CTTSs, many of which also have close companions. If there is a low-mass companion to PTFO 8-8695, accretion from a disk may be through accretion streams such as those proposed by Artymowicz & Lubow (1996) (see also G\u00fcnther & Kley 2002). At this time, it is not known if a planetary mass companion can excite accretion streams such as those modeled by Artymowicz & Lubow (1996) and G\u00fcnther & Kley (2002). A few CTTSs binaries are thought to potentially be accreting through accretion streams. These include DQ Tau (Basri et al. 1997; Mathieu et al. 1997), UZ Tau E (Jensen et al. 2007), AK Sco (Alencar et al. 2003), KH 15D (Hamilton et al. 2012), and the eclipsing binary system CoRoT 223992193 in NGC 2264 (Gillen et al. 2014). None of these stars shows the type of H\u03b1 variations seen in PTFO 8-8695 where the accretion-related emission appears to move from one side of the line profile to the other as it spirals onto one or both of the stars. This type of line profile behavior is also not seen in the H\u03b1 profile variations of single CTTSs in extensive studies of their line profile variability (e.g., Giampapa et al. 1993; Johns & Basri 1995a, 1995b; Johns-Krull & Basri 1997; Oliveira et al. 1998; Alencar et al. 2001), nor is it predicted from theoretical models of magnetospheric accretion such as those shown in Kurosawa & Romanova (2013). While we cannot completely rule out accretion from a tenuous disk as the source of the excess H\u03b1 emission observed in PTFO 8-8695, we argue that this is not the most likely explanation of the observed emission. Deep mid-IR or millimeter continuum observations, or a deep search for close circumstellar disk gas emission (e.g., H2 emission, see France et al. 2012), could shed light on whether there is a tenuous disk around this star feeding accretion onto it.","Citation Text":["Mathieu et al. 1997"],"Functions Text":["A few CTTSs binaries are thought to potentially be accreting through accretion streams. These include DQ Tau","None of these stars shows the type of H\u03b1 variations seen in PTFO 8-8695 where the accretion-related emission appears to move from one side of the line profile to the other as it spirals onto one or both of the stars."],"Functions Label":["Background","Differences"],"Citation Start End":[[839,858]],"Functions Start End":[[710,818],[1035,1251]]}
{"Identifier":"2018AandA...617A..86L__Tian_et_al._2016_Instance_1","Paragraph":"Using AIA intensity images in six EUV bandpasses, a differential emission measure (DEM) analysis (Cheng et al. 2012; Shen et al. 2015) is performed for the solar flare at 17:03:37 UT, when the oscillations are pronounced. Figure 8 panels b and c plot the DEM profiles in the flaring loop and the background corona, respectively. They contain the same region with an FOV of 3\u2033 \u00d7 3\u2033, as enclosed by the red boxes in Fig. 1e. The black profile in each panel is the best-fit DEM solution to the observed fluxes. The colored rectangles represent the errors of the DEM curve, which are calculated from 100 Monte Carlo (MC) realizations of the observational data (Cheng et al. 2012; Tian et al. 2016; Li et al. 2017b). The average temperature (T) and emission measure (EM) inside and outside (background corona) of the flaring loop are also estimated according to their errors, respectively. For example, the confident temperature (log T) range inside the flaring loop is 6.0\u22127.5, while that outside of the flaring loop is 5.8\u22127.1, since the temperatures in solar flare are much higher than that in the background corona. Therefore, the number density inside the flaring loop can be estimated with \n\n\n\n\nn\ne\n\n\n=\n\n\n\nE\nM\n\/\nw\n\n\n\n\n$ n_e\\,{=}\\,\\sqrt{EM\/w} $\n\n\n by assuming a filling factor of 1.0 (Tian et al. 2016; Li et al. 2017b), and we can obtain a lower limited density inside the flaring loop of \u223c4.7 \u00d7 1010 cm\u22123. On the other hand, the effective LOS depth (\n\n\n\nl\n\n\u2248\n\n\nH\n\u03c0\nr\n\n\n\n\u223c\n\n4\n\n\u00d7\n\n\n10\n10\n\n\n\n$ l\\,{\\approx}\\sqrt{H\\pi r}\\,{\\sim}\\,4\\,{\\times}\\, 10^{10} $\n\n\n cm), instead of the loop width, is applied to calculate the number density outside of the flaring loop (Zhang & Ji 2014; Zucca et al. 2014; Su et al. 2016; Li et al. 2017b), and we get 9.1 \u00d7 108 cm\u22123. Finally, a number density ratio (rd = n0\/ne) of \u223c0.02 between outside and inside of the flaring loop is determined, which is very close to the density contrast from recent observations (Tian et al. 2016; Li et al. 2017b).","Citation Text":["Tian et al. 2016","Tian et al. 2016"],"Functions Text":["The colored rectangles represent the errors of the DEM curve, which are calculated from 100 Monte Carlo (MC) realizations of the observational data","Finally, a number density ratio (rd = n0\/ne) of \u223c0.02 between outside and inside of the flaring loop is determined, which is very close to the density contrast from recent observations"],"Functions Label":["Uses","Similarities"],"Citation Start End":[[676,692],[1942,1958]],"Functions Start End":[[508,655],[1756,1940]]}
{"Identifier":"2022AandA...666A.141M__Hopkins_et_al._2006_Instance_1","Paragraph":"Many mechanisms that probably lead to quenching have been proposed, and surely we need some parameters to discriminate between these mechanisms. One of such critical parameters we are looking for is the quenching timescale, which varies from one mechanism to another. Most of these timescales are obtained from simulations (Wright et al. 2019; Wetzel et al. 2013; Walters et al. 2022). Among mass-dependent quenching mechanisms, the supernova feedback is a strong physical process that quickly quenches star formation in \u223c0.1 Gyr (Ceverino & Klypin 2009). On the other hand, the AGN feedback timescale is still under debate. Quasar mode AGN feedback is a strong process. Quasars\u2019 lifetimes are not long ( 0.1 Gyr, Hopkins et al. 2006), and the gas outflow driven by quasars can quench star formation at million-year-level timescales (Smethurst et al. 2021). On the other hand, the radio mode of AGN feedback is a slow process (Best et al. 2005); it takes up to and beyond 1 Gyr to establish a balance between cooling and heating to reach a low star formation rate phase (Fabian 2012). In Schawinski et al. (2014) the authors maintain that the timescale of AGN feedback leading to a quenching can be related to the host galaxy type. Furthermore, Hirschmann et al. (2017) simulated the star formation history (SFH) of galaxies with AGN and found that the quenching timescale can have a wide range, from a few hundred Myr to a few Gyr. For environmental quenching mechanisms, strangulation is a long-term process that lasts a few billion years (Peng et al. 2015). Merger-driven quenching does not happen directly after galaxy merger events; a median delay time of 1.5 Gyr is expected, and the timescale varies over a wide range (Rodr\u00edguez Montero et al. 2019). Ram pressure stripping is a rapid type of quenching, with a timescale of \u223c0.2 Gyr (Steinhauser et al. 2016). Thus, we can use the quenching timescale to determine some of the physical mechanisms involved.","Citation Text":["Hopkins et al. 2006"],"Functions Text":["Quasars\u2019 lifetimes are not long ( 0.1 Gyr,"],"Functions Label":["Background"],"Citation Start End":[[714,733]],"Functions Start End":[[671,713]]}
{"Identifier":"2015ApJ...808...56M__Beaulieu_et_al._2011_Instance_2","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"],"Functions Text":["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.,"],"Functions Label":["Background"],"Citation Start End":[[1928,1948]],"Functions Start End":[[1694,1904]]}
{"Identifier":"2022MNRAS.513.2194I___2017_Instance_1","Paragraph":"There are several assumptions related to the FB origin. One of the first intriguing assumptions of the FB origin was the scenario of the evolution of the relativistic jet remnant. Presumably, relativistic jets existed during the recent activities of Sgr A* about 107 yr ago. These assumptions lead to many interesting discussions about FB\u2019s inheritance from relativistic jets (Guo et al. 2012; Yang et al. 2012; Zhang & Guo 2020). The FBs are presumably filled with matter from jets originating from the supermassive black hole (SMBH) vicinity. The expansion process could occur at the front of a shock wave that has collided with the external environment and formed the boundaries of the currently observed FBs. On the other hand, Su & Finkbeiner (2012) provided the evidence of the possibility of a gamma jet existence, which could also lead to the filling of such a structure as FBs by hot particles. There is another possibility of the FB formation due to the tidal disruption of stars (Cheng et al. 2011; Chernyshov et al. 2014, 2017) and the acceleration of individual particles by the arising shock waves. The FB primary structure originating from supernova explosions in the central molecular zone (CMZ) is also possible (Lacki 2014). The young stars\u2019 formation and the stellar winds may also have a role in the FB origin, including the CMZ life cycle (Armillotta et al. 2019). Cheng et al. (2011) also described a scenario of FBs\u2019 periodic feeding by supernova explosions. Various interesting scenarios for the FB formation were also discussed by Mertsch & Petrosian (2019). Comparative analysis of the possible scenarios of FB origin, in general, has been of great interest for explaining its nature since discovery in 2010. It is also interesting to note that all scenarios that are mentioned above may complement each other. FBs may have really formed from relativistic jets and energized by intensive periodic (on average, once every 104 yr) processes within. Stellar plasma injected by supernova shock waves may be one of the important additional channels for the supply of energy to a stationary bubble formed from a jet.","Citation Text":["Chernyshov et al.","2017"],"Functions Text":["There is another possibility of the FB formation due to the tidal disruption of stars","and the acceleration of individual particles by the arising shock waves."],"Functions Label":["Background","Background"],"Citation Start End":[[1010,1027],[1034,1038]],"Functions Start End":[[904,989],[1040,1112]]}
{"Identifier":"2019ApJ...885..168O__Thomas_et_al._2004_Instance_1","Paragraph":"Tidal heating of Io has been shown to be responsible for its widespread volcanism. The tidal heating rate of Jupiter\u2019s tidally locked moon, \n\n\n\n\n\n, driven by forced eccentricities, e, locked by Europa and Ganymede\u2019s Laplace resonance with Io, is the dominant interior heating source. Similarly, the tidal heating of an exomoon will likely dominate the interior energy budget due to the additional stellar tide. Consequently, the tidal heating rate is orders of magnitude higher than at Io, which for an exo-Io of similar rheological properties (\n\n\n\n\n\n, Rs = RIo, \u03c1s = \u03c1Io) can be written as (Cassidy et al. 2009; Equations (19) and (20))\n3\n\n\n\n\n\nwhere \u03c5 = 3 \u00d7 10\u22127 cm3 erg\u22121, and \u03c4s = \u03c4p\/5 based on the tidal stability criterion discussed in Section 2. For utility, we describe the exo-Io\u2019s tidal efficiency as \n\n\n\n\n\n, which can readily be computed for any three-body system as tabulated in Table 4. The enhanced tidal heating described in Equation (3) will also contribute to the surface temperature T0 = Teq + \u0394T0, which is very roughly approximated as\n4\n\n\n\n\n\nwhere \u03c3sb is the Stefan\u2013Boltzmann constant and Teq. At Io, the total neutral volcanic content (SO2, SO, NaCl, KCl, Cl, and dissociation products) ejected to space (Section 4.2.1) by the incident plasma is estimated to be, on average, \u223c1000 kg s\u22121 (e.g., Thomas et al. 2004), varying within an order of magnitude over decades of observations (Burger et al. 2001; Wilson et al. 2002; Thomas et al. 2004). While the source of the dominant gas SO2 is ultimately tidally driven volcanism, the near-surface atmosphere is mostly dominated by the sublimation of SO2 frost (Tsang et al. 2016). By observing the atmospheric evolution of the SO2 column density with heliocentric distance, Tsang et al. (2013) estimated the direct volcanic component to be Nvolc \u223c 6.5 \u00d7 1016 cm\u22122, typically \n\n\n\n\n\n of the total observed SO2 column density. Ingersoll (1989) demonstrated the relative contributions due to both sublimation and volcanic sources in maintaining Io\u2019s atmosphere and established a relationship relating the volcanic source rate to the volcanically supplied atmospheric pressure:\n5\n\n\n\n\n\nThis expression also gives the volcanic column density \n\n\n\n\n\n, where g is the acceleration due to gravity. Adopting an observed atmospheric temperature of Tatm = 170 K by Lellouch et al. (2015) corresponding to an atmospheric scale height of H = 12 km, a thermal velocity \n\n\n\n\n\n equal to 150 m s\u22121, and a sticking coefficient \u03b1 = 0.5 for the SO2 mass of 64 amu yields a volcanic source rate of \n\n\n\n\n\n \u223c 6.9 \u00d7 106 kg s\u22121 of SO2 integrated over Io\u2019s mass MIo. The average volumetric mixing ratio for NaCl to SO2 at Io is observed to be XNaCl \u223c 3 \u00d7 10\u22123 (Lellouch et al. 2003). This leads to a source rate of \n\n\n\n\n\n \u223c 7.4 \u00d7 103 kg s\u22121 of NaCl, somewhat larger than but reasonably consistent with the direct measurement of the NaCl volcanic source rate of (0.8\u20133.1) \u00d7 103 kg s\u22121 (Lellouch et al. 2003). From these estimates, we will adopt \u223c3 \u00d7 103 kg s\u22121 of Na i as the volcanic source rate for Io.","Citation Text":["Thomas et al. 2004"],"Functions Text":["At Io, the total neutral volcanic content (SO2, SO, NaCl, KCl, Cl, and dissociation products) ejected to space (Section 4.2.1) by the incident plasma is estimated to be, on average, \u223c1000 kg s\u22121 (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[1315,1333]],"Functions Start End":[[1113,1314]]}
{"Identifier":"2020ApJ...901...41S__Laursen_et_al._2013_Instance_1","Paragraph":"Observations have shown that the shape of the Ly\u03b1 line is diverse. It includes broad damped absorption profiles, P-Cygni profiles, double-peak profiles, pure symmetric emission profiles, and combinations thereof (Kunth et al. 1998; Mas-Hesse et al. 2003; Shapley et al. 2003; M\u00f8ller et al. 2004; Noll et al. 2004; Tapken et al. 2004; Venemans et al. 2005; Wilman et al. 2005). This variety can be understood through a detailed radiative transfer calculation, which is analytically solvable only for simple cases (e.g., a static, plane-parallel slab by Harrington 1973 and Neufeld 1990, and a static uniform sphere by Dijkstra et al. 2006). Later, numerical algorithms based on Monte Carlo techniques were developed to solve radiative transfer for more general cases. Now theoretical studies mostly rely on them (e.g., Spaans 1996; Loeb & Rybicki 1999; Ahn et al. 2000, 2002; Zheng & Miralda-Escud\u00e9 2002; Richling 2003; Cantalupo et al. 2005; Dijkstra et al. 2006; Hansen & Oh 2006; Tasitsiomi 2006; Verhamme et al. 2006, 2015; Laursen et al. 2013; Behrens et al. 2014; Duval et al. 2014; Gronke et al. 2015; Smith et al. 2019; Lao & Smith 2020; Michel-Dansac et al. 2020). Meanwhile, a galaxy model needs to be constructed to perform such a radiative transfer calculation. One can adopt a realistic galaxy model from hydrodynamical simulations. Galaxies from such simulations can be useful for performing a statistical study of Ly\u03b1 properties, but they cannot be directly used to model individual galaxies in observations. Therefore it would be better to adopt a simple but manageable toy model for the purpose of reproducing observations. One example for such models is a shell model, in which a central Ly\u03b1 source is surrounded by a constantly expanding, homogeneous, spherical shell of H i medium with dust. Although this shell model has surprisingly well reproduced many observed Ly\u03b1 line profiles (e.g., Ahn 2004; Schaerer & Verhamme 2008; Verhamme et al. 2008; Schaerer et al. 2011; Gronke et al. 2015; Yang et al. 2016; Gronke 2017; Karman et al. 2017), because of its extreme simplicity and contrivance, there is still room for improvement (e.g., see Section 7.2 in Yang et al. 2016; Orlitov\u00e1 et al. 2018).","Citation Text":["Laursen et al. 2013"],"Functions Text":["Later, numerical algorithms based on Monte Carlo techniques were developed to solve radiative transfer for more general cases. Now theoretical studies mostly rely on them (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[1027,1046]],"Functions Start End":[[640,817]]}
{"Identifier":"2022MNRAS.514.1961R__Prochaska_&_Zheng_2019_Instance_2","Paragraph":"Along with the time-domain detections, we identified J173438.35-504550.4 as a potential host galaxy for FRB 20201123A using robust statistical treatment given the relatively small localization error region. At face value, the low redshift of J173438.35-504550.4 appears at odds with the large dispersion measure for FRB 20201123A (${\\rm DM}_{\\rm FRB}\\approx 434 \\, {\\rm pc \\, cm^{-3}}$). Our Galaxy, however, contributes ${\\rm DM}_{\\rm ISM}\\approx 200 \\, {\\rm pc \\, cm^{-3}}$ (NE2001 gives\u2009229\u2009${\\rm pc \\, cm^{-3}}$ and YMW16 gives 162 ${\\rm pc \\, cm^{-3}}$) from its interstellar medium and a presumed ${\\rm DM}_{\\rm Halo}\\sim 50 \\, {\\rm pc \\, cm^{-3}}$ from its halo (Prochaska & Zheng 2019). This leaves ${\\approx}180 \\, {\\rm pc \\, cm^{-3}}$ for the cosmos (DMcosmic) and the host (DMhost). At z = 0.05, the average cosmic contribution is $\\langle {\\rm DM}_{\\rm cosmic}\\rangle \\sim 42\\, {\\rm pc \\, cm^{-3}}$ (Macquart et al. 2020) but the intrinsic scatter in this quantity is predicted to be large. Adopting the best-fitting model to the Macquart relation by Macquart et al. (2020), the 95 per cent confidence interval is ${\\rm DM}_{\\rm cosmic}= [15,125] \\, {\\rm pc \\, cm^{-3}}$. Allowing for the maximum value of this interval (which would imply a significant foreground galaxy halo), we recover a minimum host contribution of ${\\rm DM}_{\\rm host, min} \\approx 60~\\rm pc~cm^{-3}$. This is consistent with estimates for host galaxy contributions from theoretical and empirical treatments (Prochaska & Zheng 2019; James et al. 2022). For a true DMcosmic value of this sightline closer to (or below) the mean, the host contribution would exceed $100 \\, {\\rm pc \\, cm^{-3}}$. Such values are inferred for other FRB hosts (e.g. FRB\u200920121102A; Tendulkar et al. 2017). In conclusion, we find no strong evidence to rule out the association with J173438.35-504550.4 based on its redshift and DMFRB. The significant host contribution to the DM, combined with the scattering in FRB 20201123A possibly originating in the host, shows that it shares similarities with other highly active, repeating FRBs like FRB 20121102A and FRB 20190520A and potentially resides in a turbulent and dense environment within the host.","Citation Text":["Prochaska & Zheng 2019"],"Functions Text":["Allowing for the maximum value of this interval (which would imply a significant foreground galaxy halo), we recover a minimum host contribution of ${\\rm DM}_{\\rm host, min} \\approx 60~\\rm pc~cm^{-3}$. This is consistent with estimates for host galaxy contributions from theoretical and empirical treatments"],"Functions Label":["Similarities"],"Citation Start End":[[1493,1515]],"Functions Start End":[[1184,1491]]}
{"Identifier":"2022ApJ...927..145S__Cui_et_al._2012_Instance_1","Paragraph":"It is generally thought that the stellar halo of the Milky Way (MW) contains the fossil record of the MW\u2019s formation history imprinted in its kinematical and chemical properties. Since Eggen et al. (1962) showed a model of the MW\u2019s formation process from the kinematic analysis of halo stars in the solar neighborhood, many studies of this Galactic fossil component have been carried out using a large number of newly available halo samples, with their assembly requiring considerable observational efforts over the past decades. These include, for example, the Hipparcos Catalog, obtained from the first astrometry satellite (Perryman et al. 1997), and spectroscopic catalogs such as RAVE (Steinmetz et al. 2006), SEGUE (Yanny et al. 2009), LAMOST (Cui et al. 2012; Zhao et al. 2012), GALAH (De Silva et al. 2015), APOGEE (Majewski et al. 2017), H3 (Conroy et al. 2019), and more. Perhaps the most significant impacts on this field of research have been brought about by the second astrometry satellite, Gaia. The Gaia catalog (Gaia Collaboration et al. 2016, 2018, 2021) provides trigonometric parallaxes and proper motions for billions of Galactic stars with unprecedented high accuracy. Based on the astrometry data of the stars in the Gaia catalog combined with that of the spectroscopic catalogs, the MW\u2019s new dynamical maps have been drawn (e.g., Belokurov et al. 2018; Helmi et al. 2018; Myeong et al. 2018a, 2018b; Beane et al. 2019; Hagen et al. 2019; Iorio & Belokurov 2019; Anguiano et al. 2020; Cordoni et al. 2021; Koppelman et al. 2021), and new aspects of the MW\u2019s formation and evolution history have been revealed (e.g., Antoja et al. 2018; Sestito et al. 2019; Wyse 2019). In particular, the Gaia catalog has enabled the discovery of new substructures in the MW\u2019s stellar halo (e.g., Koppelman et al. 2019; Myeong et al. 2019; Li et al. 2020) that are remnants of past merging\/accretion events associated with the MW\u2019s formation history (e.g., Fern\u00e1ndez-Trincado et al. 2020; Naidu et al. 2020, 2021).","Citation Text":["Cui et al. 2012"],"Functions Text":["Since Eggen et al. (1962) showed a model of the MW\u2019s formation process from the kinematic analysis of halo stars in the solar neighborhood, many studies of this Galactic fossil component have been carried out using a large number of newly available halo samples, with their assembly requiring considerable observational efforts over the past decades. These include, for example,","LAMOST","and more."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[750,765]],"Functions Start End":[[179,557],[742,748],[872,881]]}
{"Identifier":"2020MNRAS.498..385J__Roman-Duval_et_al._2010_Instance_1","Paragraph":"In the upper panel of Fig. 17, we demonstrate that the mass distribution of GMCs in each galaxy reproduces the upper limit of \u223c3 to 8 \u00d7 106\u2009M\u2299 observed by Rosolowsky et al. (2003) in M33, by Freeman et al. (2017) in M83 and by Miville-Desch\u00eanes et al. (2017) and Colombo et al. (2019) in the Milky Way. This upper limit has been predicted to arise due to a combination of centrifugal forces and stellar feedback (Reina-Campos & Kruijssen 2017). We also find a turnover in the mass spectrum between 104.8 and 105\u2009M\u2299, consistent with the behaviour of the GMC mass distribution in the Milky Way (Miville-Desch\u00eanes et al. 2017), although we cannot rule out the possibility that the turnover we see in the simulations is influenced by their limited mass resolution. Above the turnover, the GMC mass function has a power-law form with \u03b2 \u223c 1.9, close to the observed range of \u03b2 \u2208 [1.6, 1.8] for clouds in the Milky Way (Solomon et al. 1987; Williams & McKee 1997; Heyer et al. 2009; Roman-Duval et al. 2010; Miville-Desch\u00eanes et al. 2017; Colombo et al. 2019) over the same mass range (log\u2009M \u2208 [4.8, 6.5]). In the lower panel of Fig. 17, we display the spectrum of GMC sizes for each simulated galactic disc, given by the effective cloud radius \u2113eff, such that\n(42)$$\\begin{eqnarray*}\r\n\\ell _{\\rm eff} = 1.91 \\sqrt{\\Delta \\ell _{\\rm maj}^2 + \\Delta \\ell _{\\rm min}^2},\r\n\\end{eqnarray*}$$where \u0394\u2113maj and \u0394\u2113min are the second moments of an ellipse fitted to the footprint of each cloud in the galactic mid-plane, using astrodendro. We adopt this definition of the cloud size in order to make a direct comparison to works in the existing observational literature (e.g. Solomon et al. 1987; Bertoldi & McKee 1992; Rosolowsky & Leroy 2006; Colombo et al. 2019). The factor of 1.91 is the correction first defined by Solomon et al. (1987) for converting the RMS cloud extent to an estimate of the spherical cloud size. The smallest resolved cloud has a diameter of 18\u2009pc, so we do not capture the observed turnover of the distribution at \u223c30\u2009pc (Miville-Desch\u00eanes et al. 2017). Likewise, our largest clouds slightly exceed the truncation size of 70\u2009pc observed by Colombo et al. (2019), with a maximum diameter of \u223c200\u2009pc. Importantly, we do approximately reproduce the observed power-law slope of $\\mathrm{d}N\/\\mathrm{d}R \\sim R^{-\\beta _\\ell }$ with \u03b2\u2113 \u223c 2.8 (Colombo et al. 2019). This is given by the black line in Fig. 17, while our best fit to the simulation data over the observed range of cloud sizes \u2113eff \u2208 [18, 70] pc is given by the purple line, with a slightly shallower slope of \u03b2\u2113 = 2.43 \u00b1 0.06.","Citation Text":["Roman-Duval et al. 2010"],"Functions Text":["Above the turnover, the GMC mass function has a power-law form with \u03b2 \u223c 1.9, close to the observed range of \u03b2 \u2208 [1.6, 1.8] for clouds in the Milky Way"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[976,999]],"Functions Start End":[[761,911]]}
{"Identifier":"2016MNRAS.459.3585G__Tiengo_&_Mereghetti_2007_Instance_1","Paragraph":"In the following, we consider an NS with mass MNS = 1.5\u2009M\u2299 and radius RNS = 12 km, which is compatible with expectations from modern equations of state such as APR or BSk21 models (Akmal, Pandharipande & Ravenhall 1998; Goriely, Chamel & Pearson 2010). The value of the radius is also in agreement with the estimates derived by Sartore et al. (2012) and Ho et al. (2007), assuming a source distance of 120 pc (Walter et al. 2010). This choice translates into a gravitational redshift factor at the star surface 1 + z = 1.26. The rotational period of RX J1856 is P = 7 s and the X-ray pulsed fraction is the lowest among the XDINSs, \u223c1.3\u2009per\u2009cent (Tiengo & Mereghetti 2007). The polar strength of the dipole field is taken to be Bp = 1013 G, a value which is somehow intermediate between the spin-down measure and the estimates from spectral fitting (Ho et al. 2007; van Kerkwijk & Kaplan 2008). We assume that the magnetic field is dipolar (see Section 2) and that the surface temperature distribution is that induced by the core-centred dipole. Since for fields higher than \u223c1011\u2009G, electron conduction across the field lines is strongly suppressed, the meridional temperature variation is Tdip \u2243 Tp|cos\u2009\u03b8B|1\/2, where Tp is the polar value of the temperature (e.g. Greenstein & Hartke 1983). We checked that this simple expression for Tdip differs only slightly ( \u2272 6\u2009\u2009per\u2009\u2009cent) from the more accurate formula by Potekhin, Pons & Page (2015) for \u03b8 \u2272 80\u00b0. However, taken face value, the previous expression for Tdip yields vanishingly small values near the magnetic equator. The analysis of Sartore et al. (2012) has shown that the X-ray spectrum of RX J1856 is best modelled in terms of two blackbody components with $kT^\\infty _\\mathrm{h}\\sim 60$ eV and $kT^\\infty _\\mathrm{c}\\sim 40$ eV. To account for this in a simple way, we actually adopt a temperature profile given by Ts = max\u2009(Tdip, Tc) with Tp = Th, where $T_\\mathrm{h,c}=T^\\infty _\\mathrm{h,c}\/(1+z)$.","Citation Text":["Tiengo & Mereghetti 2007"],"Functions Text":["The rotational period of RX J1856 is P = 7 s and the X-ray pulsed fraction is the lowest among the XDINSs, \u223c1.3\u2009per\u2009cent"],"Functions Label":["Background"],"Citation Start End":[[647,671]],"Functions Start End":[[525,645]]}
{"Identifier":"2020MNRAS.491.5406T__Chisholm_et_al._2017_Instance_1","Paragraph":"We show the mass-loading factors required to simultaneously reproduce the stellar metallicity and SFR of passive galaxies in the bottom panel of Fig. 9. We find that the mass-loading factor strongly decreases with increasing stellar mass. This anticorrelation between stellar mass and \u03bbeff is qualitatively consistent with the mass dependence in theoretical models, which predict that $\\lambda \\propto M_*^{-1\/3}$ (Murray, Quataert & Thompson 2005) for momentum-driven winds and $\\lambda \\propto M_*^{-2\/3}$ for energy-driven winds (e.g. Dekel & Silk 1986), as well as other observational evidence (e.g. Heckman et al. 2015; Chisholm et al. 2017; Fluetsch et al. 2019). Our results indicate that \u2018effective\u2019 outflows (which are capable of permanently removing gas from the galaxy) are, together with starvation, of increasing importance in low-mass galaxies. In particular, since the rate at which gas is locked up into long-lived stars is given by (1 \u2212 R)\u03a8 = 0.575\u03a8, and the rate at which gas is ejected from the galaxy is given by \u03bbeff\u03a8, then, for galaxies with log\u2009(M*\/M\u2299)  10.2, the rate at which gas is lost through galactic winds is roughly 1\u20133 times larger than the rate at which gas is locked up into long-lived stars. Clearly outflows play an essential role in depleting the gas reservoir of low-mass galaxies during the quenching phase. Outflows are relatively weaker (\u03bbeff \u2264 0.6) in more massive galaxies (log\u2009(M*\/M\u2299) > 10.2), with starvation becoming the dominant quenching mechanism, as illustrated in Fig. 9. However, outflows still play an important role in quenching star formation, as comparable amounts of gas are removed through galactic winds and through the formation of long-lived stars. Although these massive galaxies may be ejecting large amounts of gas in the form of outflows (i.e. a large \u03bb), our results suggest that these ejection events are short lived and\/or most of the outflowing gas does not escape the galaxy and is instead recycled (i.e. a relatively small \u03bbeff).","Citation Text":["Chisholm et al. 2017"],"Functions Text":["This anticorrelation between stellar mass and \u03bbeff is qualitatively consistent","as well as other observational evidence (e.g."],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[625,645]],"Functions Start End":[[239,317],[558,603]]}
{"Identifier":"2022MNRAS.512.4280P__Umetsu_et_al._2016_Instance_2","Paragraph":"The fifth force, propagated by the scalar degree of freedom, affects the Poisson equations associated to the Newtonian potential \u03a6, as well as the relativistic one, \u03a8, according to (Kobayashi, Watanabe & Yamauchi 2015; Crisostomi & Koyama 2018; Dima & Vernizzi 2018),\n(1)$$\\begin{eqnarray*}\r\n\\frac{\\text{d} \\Phi (r)}{\\text{d}r} = \\frac{G M(r)}{r^2} \\left[1+\\frac{3}{4}Y_1\\left(\\frac{\\rho (r)}{\\bar{\\rho }(r)}\\right)\\left(2+\\frac{\\text{d}\\ln \\rho }{\\text{d}\\ln r}\\right)\\right],\r\n\\end{eqnarray*}\r\n$$(2)$$\\begin{eqnarray*}\r\n\\frac{\\text{d} \\Psi (r)}{\\text{d}r} =\\frac{G M(r)}{r^2}\\left[1-\\frac{15}{4}Y_2\\left(\\frac{\\rho (r)}{\\bar{\\rho }(r)}\\right)\\right].\r\n\\end{eqnarray*}\r\n$$In the above equations, we have assumed spherical symmetry. $\\bar{\\rho }(r)$ is the (spatially) average density at radius r from the centre of the galaxy cluster, and Y1, Y2 correspond to the dimensionless fifth-force couplings. Finally, G is the Newton\u2019s constant. Although the dynamics of member galaxies in the cluster is governed by the potential \u03a6, lensing is sourced by the combination\n(3)$$\\begin{eqnarray*}\r\n\\frac{\\mathrm{ d}}{\\mathrm{ d}r} \\Phi _{\\rm {lens}} = \\frac{1}{2}\\frac{\\mathrm{ d}}{\\mathrm{ d}r}(\\Phi + \\Psi).\r\n\\end{eqnarray*}\r\n$$Therefore, kinematical observations allow for contraints on Y1, while lensing constrains both Y1 and Y2. The right-hand side of above equation can be expressed in terms of the density profile \u03c1(r) according to the relevant equations for \u03a6 and \u03a8 above. The dominant source of pressureless matter density in the cluster comes from dark matter, which density we choose to model with a Navarro-Frenk-White (NFW) of Navarro, Frenk & White (1997) profile as\n(4)$$\\begin{eqnarray*}\r\n\\rho (r)=\\frac{\\rho _\\text{s}}{r\/r_\\text{s}(1+r\/r_\\text{s})^2},\r\n\\end{eqnarray*}\r\n$$with \u03c1s is a characteristic density and rs the radius at which the logarithmic derivative of the density profile takes the value \u22122. The NFW profile has been shown to provide an overall good agreement with observations and simulations over a broad range of scales in GR (e.g. Biviano et al. 2013; Umetsu et al. 2016; Peirani et al. 2017) and in MG (e.g. Lombriser et al. 2012a; Wilcox et al. 2016). Moreover, the GR analyses with lensing and internal kinematics of both clusters indicate that the total mass profile is well fitted by the NFW model (Biviano et al. 2013; Umetsu et al. 2016; Caminha et al. 2017; Sartoris et al. 2020). Under the assumption of a NFW profile, we can re-write the equation for the potential \u03a6 in an effective way as\n(5)$$\\begin{eqnarray*}\r\n\\frac{\\text{d}\\Phi }{\\text{d}r} \\equiv \\frac{G M_{\\text{dyn}}}{r^2}=\\frac{G}{r^2}\\left[ M_{\\rm {NFW}}(r)+M_1(r)\\right],\r\n\\end{eqnarray*}\r\n$$which serves as a definition of the dynamical mass Mdyn. Notice that, G here is still Newton\u2019s constant as measure in the Solar system. The fifth-force contribution M1 is defined in terms of the NFW parameters as\n(6)$$\\begin{eqnarray*}\r\nM_1(r)= M_{200}\\frac{Y_1}{4}\\frac{r^2(r_\\text{s}-r)}{(r_\\text{s}+r)^3}\\times [\\ln (1+c)- c\/(1+c)]^{-1}.\r\n\\end{eqnarray*}\r\n$$where c = r\/rs is the concentration and M200 is the mass of a sphere of radius r200 enclosing an average density 200 times the critical density of the universe at that redshift. In a similar fashion, the relevant expression for the lensing mass can be found by computing\n(7)$$\\begin{eqnarray*}\r\nM_{\\text{lens}}(r) =\\frac{r^2}{2G}\\left[\\frac{\\text{d}\\Psi }{\\text{d}r}+\\frac{\\text{d}\\Phi }{\\text{d}r}\\right].\r\n\\end{eqnarray*}\r\n$$$$\\begin{eqnarray*}\r\nM_{\\text{lens}}=M_{\\text{NFW}}+\\frac{r^2M_{200}\\left[Y_1(r_\\text{s}-r)-5Y_2(r_\\text{s}+r)\\right]}{4[\\log (1+c_{200})-c_{200}\/(1+c_{200})]}\\frac{1}{(r_\\text{s}+r)^{3}},\r\n\\end{eqnarray*}\r\n$$which can be effectively re-expressed in terms of the dynamical mass as\n(8)$$\\begin{eqnarray*}\r\nM_{\\text{lens}} \\equiv M_{\\text{dyn}}+M_2,\r\n\\end{eqnarray*}\r\n$$with M2 the contribution from the fifth force defined through\n(9)$$\\begin{eqnarray*}\r\nM_2=\\frac{r^2M_{200}}{8(r_\\text{s}+r)^{3}}\\frac{Y_1(r-r_\\text{s})-5Y_2(r_\\text{s}+r)}{[\\ln (1+c)-c\/(1+c)]}.\r\n\\end{eqnarray*}\r\n$$In view of the above equations, it is important to emphasize again that, although the fifth force effect enters the dynamical mass only through the coupling Y1, the lensing mass is affected by both Y1 and Y2. This is expected, since lensing is sourced by the combination of the two potentials \u03a6 and \u03a8, equation (3). Note also that, with gravitational lensing observations, one reconstructs the projected surface mass density profile \u03a3(R), where R is the projected radius from the cluster centre. We refer to e.g. Umetsu (2020) for an explicit discussion of the physics and mathematical framework.","Citation Text":["Umetsu et al. 2016"],"Functions Text":["Moreover, the GR analyses with lensing and internal kinematics of both clusters indicate that the total mass profile is well fitted by the NFW model"],"Functions Label":["Similarities"],"Citation Start End":[[2351,2369]],"Functions Start End":[[2180,2328]]}
{"Identifier":"2020MNRAS.493.5413K__Dubey_et_al._2009_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":["Dubey et al. 2009"],"Functions Text":["We test in Section 2 two available one-dimensional (1D) codes:","and a modified version of the Eulerian, 1D hydrodynamic flash4.0 code with thermonuclear burning","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)."],"Functions Label":["Uses","Uses","Uses","Compare\/Contrast"],"Citation Start End":[[711,728]],"Functions Start End":[[416,478],[592,688],[731,757],[758,994]]}
{"Identifier":"2016MNRAS.462.1415C__Lu_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":["Lu et al. 2014"],"Functions Text":["This progress has led to a general consensus about the main physical ingredients required to describe the evolution of the galaxy population (e.g."],"Functions Label":["Background"],"Citation Start End":[[790,804]],"Functions Start End":[[615,761]]}
{"Identifier":"2017ApJ...834L..13W__Chang_et_al._2012_Instance_1","Paragraph":"Thanks to their short spectral lags, cosmological distances, and very high-energy photons, GRBs have been viewed as the most promising sources for studying  LIV effects (Amelino-Camelia et al. 1998; Ellis et al. 2006; Jacob & Piran 2008). In the past, various limits on LIV have been obtained from the spectral time lags of individual GRBs or a large sample of GRBs (e.g., Amelino-Camelia et al. 1998; Coleman & Glashow 1999; Schaefer 1999; Ellis et al. 2003, 2006; Boggs et al. 2004; Kahniashvili et al. 2006; Jacob & Piran 2008; Abdo et al. 2009a, 2009b; Biesiada & Pi\u00f3rkowska 2009; Xiao & Ma 2009; Shao et al. 2010; Chang et al. 2012, 2016; Nemiroff et al. 2012; Ellis & Mavromatos 2013; Kosteleck\u00fd & Mewes 2013; Vasileiou et al. 2013, 2015; Pan et al. 2015; Zhang & Ma 2015; Wei et al. 2016). In particular, Abdo et al. (2009a) used the time lag of the highest energy (13.2 GeV) photon from GRB 080916C to constrain the linear LIV energy scale (\n\n\n\n\n\n\nE\n\n\nQG\n,\n1\n\n\n\n\n) and presented a stringent limit of \n\n\n\n\n1.3\n\u00d7\n\n\n10\n\n\n18\n\n\n\n\n GeV, improving the previous limits by at least one order of magnitude. Abdo et al. (2009b) set the current strictest limits on both the linear and quadratic LIV energy scales by analyzing the arrival time delay between a 31 GeV photon and the low-energy (trigger) photons from GRB 090510. The limits set are \n\n\n\n\n\n\nE\n\n\nQG\n,\n1\n\n\n>\n(\n1\n\u2212\n10\n)\n\n\nE\n\n\nPl\n\n\n\n\n and \n\n\n\n\n\n\nE\n\n\nQG\n,\n2\n\n\n>\n1.3\n\u00d7\n\n\n10\n\n\n11\n\n\n\n\n GeV. However, these limits were based on the rough time lag of a single GeV-scale photon. It is necessary to consider using the true spectral time lags of bunches of high-energy photons (i.e., the lags of high-quality, high-energy light curves) to constrain the LIV. Furthermore, since the emission mechanism of GRBs is still poorly understood, it is difficult to distinguish an intrinsic time delay at the source from a delay induced by propagation in a vacuum to the observer. That is, the method of the flight-time difference used for testing LIV is hindered by our ignorance concerning the intrinsic time delay in different energy bands (see, e.g., Ellis et al. 2006; Biesiada & Pi\u00f3rkowska 2009).","Citation Text":["Chang et al. 2012"],"Functions Text":["In the past, various limits on LIV have been obtained from the spectral time lags of individual GRBs or a large sample of GRBs (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[619,636]],"Functions Start End":[[239,372]]}
{"Identifier":"2022MNRAS.512.1499R__Holmbeck_et_al._2019_Instance_1","Paragraph":"The differences in composition are reflected in the final abundances after r-process nucleosynthesis, shown in Fig. 15. The abundances are obtained using a grid of pre-computed trajectories with SkyNet (Lippuner & Roberts 2017), as discussed in detail in Radice et al. (2018b). We normalize the relative abundances by fixing the height of the third r-process peak (A \u2243 190). We also report Solar r-process abundances from Arlandini et al. (1999) in the same figure. However, we emphasize that even if NS mergers were the sole contributor of r-process elements, there is no reason to expect that every merger should produce ejecta with relative abundances close to Solar. Indeed, variability between the yields of different mergers is required to explain observed abundances in metal-poor stars (Holmbeck et al. 2019). Overall, the simulations span a factor \u223c2 in the ratio of A \u2243 100 to third r-process peak. However, the difference between the M0 and M1 at the SR resolution, which is the resolution we use for production simulation, are modest compared to the systematic uncertainties from the unknown NS EOS and to the variability due to the binary mass ratio (Radice et al. 2018b; Nedora et al. 2021b). Clearly, strong conclusions cannot be drawn from this limited study alone, but our simulations suggest that the uncertainties in the yields from mergers arising from neutrino radiation treatment are modest. This is also supported by the results of Foucart et al. (2020). They compared M1 and Monte Carlo neutrino transport in the context of NS mergers and reported only a modest ${\\sim }10{{\\ \\rm per\\ cent}}$ difference in the Ye of the ejecta between the two schemes. Interestingly, they reported that M1 systematically overestimates the Ye of the ejecta, so we cannot exclude that the M0+Leakage results are actually more accurate than the results obtained with THC_M1. That said, it is important to emphasize that this comparisons has only been made for the dynamical ejecta and not for the secular ejecta, which we discuss in Section 6.5.","Citation Text":["Holmbeck et al. 2019"],"Functions Text":["Indeed, variability between the yields of different mergers is required to explain observed abundances in metal-poor stars"],"Functions Label":["Uses"],"Citation Start End":[[795,815]],"Functions Start End":[[671,793]]}
{"Identifier":"2021AandA...656A..79K__Matsumoto_et_al._(1990)_Instance_1","Paragraph":"Theoretical single-ionization cross-sections for the Ar2+ ion are compared to the experimental data in Fig. 2. The cross-sections for other two levels, 3P1 and 3P0, are in close agreement with the values from 3P2 and are therefore not presented. The cross-sections of the direct process are evaluated in the potential of the ionized ion. Experimental data obtained by Diserens et al. (1988) and Man et al. (1993) are only presented in Fig. 2. These experimental data showed the lowest contribution from the metastable fraction in the ion beam. Measurements from Muller et al. (1980), Danjo et al. (1984), Mueller et al. (1985), and Matsumoto et al. (1990) demonstrated onsets below the ionization threshold. In addition, the CADW calculations (Loch et al. 2007) are presented for comparison in Fig. 2. The DW cross-sections for the ground level are in good agreement with the experimental results at low energies (Diserens et al. 1988; Man et al. 1993). However, the calculations are below the measurements at the high energies. The theoretical cross-sections are above the experimental values for the two highest levels of the ground configuration from the ionization threshold up to 200\u2013300 eV, which is well beyond the peak region. On the other hand, good agreement with the measurements for these two levels is obtained at the higher energies. It should be noted that the CADW cross-sections (Loch et al. 2007) underestimate the measurements at the low electron energies. Excitations from the 2p subshell are investigated in the CADW calculations of these latter authors, but it is not clear whether the EA channels corresponding to the excitations from the 3s and 3p subshells are included in their study. On the other hand, decays of the excited Ar2+ 2p 53s 23p4 nl (n \u2265 4, l n) configurations lead to states of the Ar4+ ion. The total excitation cross-sections from the 2p subshell amount to ~1 Mb at the peakenergy of 280 eV in our study. Therefore, the excitations from the 2p subshell are not investigated in this work.","Citation Text":["Matsumoto et al. (1990)"],"Functions Text":["Measurements from Muller et al. (1980), Danjo et al. (1984), Mueller et al. (1985), and","demonstrated onsets below the ionization threshold."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[632,655]],"Functions Start End":[[544,631],[656,707]]}
{"Identifier":"2015MNRAS.447.3936M__Grigorieva_et_al._2007_Instance_1","Paragraph":"We now analyse whether release of this entrapped CO is possible, and quantify it using simple assumptions. In the Fomalhaut ring, the rate at which large planetesimals in the belt are being ground down, which is the same as the rate at which small dust grains are being replenished, has been estimated as 6.0 \u00d7 1017 kg yr\u22121 (Wyatt & Dent 2002). Assuming the planetesimal composition to be similar to that observed in Solar system comets (ice\/rock ratio of about 1, CO\/H2O ice ratios between 0.4 and 30\u2009per\u2009cent), we can use it to estimate the rate of CO release, which corresponds to between 1.9 \u00d7 1015 and 8.0 \u00d7 1016 kg yr\u22121. The mechanisms that can be responsible for this release are collisions and photodesorption. Collisions can contribute through dust vaporization (Czechowski & Mann 2007), planetesimal breakup (Zuckerman & Song 2012) and giant impacts (Jackson et al. 2014). Photodesorption, on the other hand, will affect the H2O ice on the surface of solids, in turn exposing CO and allowing it to escape on very short time-scales. The rate at which water vapour will be released is da\/dt \u223c 1.5 \u00d7 10\u22123 \u03bcm yr\u22121 (scaled to the distance of the ring, from the result of Grigorieva et al. 2007), where a is the vertical thickness of the layer. The total water mass released in this manner will then be\n\n(19)\n\n\\begin{equation}\n\\frac{\\mathrm{d}M_{\\rm H_2O}}{\\mathrm{d}t} = 4\\sigma _{\\rm tot}\\rho _{\\rm gr} \\frac{\\mathrm{d}a}{\\mathrm{d}t},\n\\end{equation}\n\nwhere \u03c1gr is the grain density, which we take as that of water ice (\u223c1 \u00d7 10\u221215 kg \u03bcm\u22123) and \u03c3tot is the total cross-section of icy grains in the Fomalhaut ring (in \u03bcm2). Under the assumption that all the grains are fully icy, we can use the total cross-sectional area of the Fomalhaut ring (33.7 au2; Wyatt & Dent 2002) to obtain an H2O production rate of 4.6 \u00d7 1018 kg yr\u22121. This is higher than the rate at which mass is being passed down the collisional cascade (again, 6.0 \u00d7 1017 kg yr\u22121), pointing towards a more realistic scenario where grains are not fully icy, but made up of a mixture of ice and rock. We therefore conclude that photodesorption of H2O might contribute significantly to the release of trapped CO gas from planetesimals in the ring, though the extent of this contribution depends on how much of the cross-sectional area of the Fomalhaut disc is icy. In any case, the exact mechanism for CO gas production is unimportant as long as there is one that can feasibly explain its release.","Citation Text":["Grigorieva et al. 2007"],"Functions Text":["The rate at which water vapour will be released is da\/dt \u223c 1.5 \u00d7 10\u22123 \u03bcm yr\u22121 (scaled to the distance of the ring, from the result of","), where a is the vertical thickness of the layer."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1176,1198]],"Functions Start End":[[1042,1175],[1198,1248]]}
{"Identifier":"2021MNRAS.503...13Z__Asplund_et_al._2000_Instance_1","Paragraph":"Although the method developed here is not identical to the means by which SONG determines radial velocity from the observed spectra, it simulates the SONG observations sufficiently well. First, the set of fictitious Fe\u2009i lines carefully chosen in this work is able to represent the properties of most Fe\u2009i lines seen between 4400 and 6900 \u00c5, which constitute a large part of all lines in this wavelength interval. Secondly, the procedure to extract radial velocity from theoretical spectral lines is similar to how radial velocities are typically obtained from observed spectra. Thirdly, the evaluation of our final radial velocity amplitude includes the information of many spectral lines that span the whole range in observation. The major uncertainty in our method is associated with the linear relationship between radial velocity amplitude and equivalent width. Due to the complicated physical processes involved in spectral line formation in a 3D atmosphere (e.g. Asplund et al. 2000), it is difficult to quantify higher order effects beyond the linear relation between $\\mathfrak {v}$ and W\u03bb; that is, the systematic uncertainty of the linear fitting. Nevertheless, it is still illuminating to provide the statistical uncertainty. The statistical uncertainty is quantified using the bootstrap method. The data set considered here is the radial velocity amplitude and equivalent width of 49 fictitious Fe\u2009i lines. We conduct 10\u2009000 bootstrap samplings, that is, generating 10\u2009000 data sets each containing 49 randomly sampled $\\mathfrak {v}$ and W\u03bb pairs. A linear regression between equivalent width and radial velocity is then performed for each re-sampled data set. For each fitting, we compute the equivalent width weighted mean radial velocity amplitude for all selected Fe\u2009i lines. The bootstrap method therefore results in 10\u2009000 weighted mean radial velocity amplitudes, their mean and variance is the desired final radial velocity amplitude and its statistical uncertainty, which is 72.2 \u00b1 0.5\u2009m\u2009s\u22121 for the 3D solar model and 93.2 \u00b1 0.3\u2009m\u2009s\u22121 for the \u03f5 Tau model.","Citation Text":["Asplund et al. 2000"],"Functions Text":["Due to the complicated physical processes involved in spectral line formation in a 3D atmosphere (e.g.","it is difficult to quantify higher order effects beyond the linear relation between $\\mathfrak {v}$ and W\u03bb; that is, the systematic uncertainty of the linear fitting."],"Functions Label":["Uses","Compare\/Contrast"],"Citation Start End":[[970,989]],"Functions Start End":[[867,969],[992,1158]]}
{"Identifier":"2022MNRAS.510.6085L__Tafalla_&_Hacar_2015_Instance_1","Paragraph":"Filamentary structures have been found at almost all size scales in the Galaxy. Massive, long filamentary dark clouds are commonly found inside giant molecular clouds (GMCs; e.g. Bergin & Tafalla 2007; Andr\u00e9 et al. 2014, and references therein), such as the dark clouds L1495 in the Taurus cloud complex (e.g. Chapman et al. 2011) and the Serpens South cloud in the Serpens region (e.g. Dhabal et al. 2018). Filamentary clouds of 4\u20136 pc length are common, and possibly longer than 10 pc. Some of these clouds are dark at infrared wavelengths. The line\u2013width size relation observed for molecular gas indicates that the thermal Mach number would exceed 10 at such size scales. The long-term survival of these filamentary structures requires a reinforcing mechanism. As shown in the ideal magnetohydrodynamical (MHD) simulations of Li & Klein (2019), a moderately strong, large-scale magnetic field (Alfv$\\acute{\\rm e}$n Mach number, ${{\\cal M}_{\\rm A}}\\sim 1$) can provide such a mechanism. In the weak-field model with ${{\\cal M}_{\\rm A}}=10$, the appearance of molecular clouds is clumpy, rather than the long and slender filamentary clouds seen in moderately strong field models. High-resolution images of massive molecular clouds from the Herschel space telescope reveal complex filamentary substructures (e.g. Andr\u00e9 et al. 2014). The characteristic inner width of molecular filaments found with Herschel is about \u223c0.1 pc (Arzoumanian 2011; Arzoumanian et al. 2019). Dense cores, where stars form, are located along or at the intersections of some of these fine substructures (e.g. K\u00f6nyves et al. 2015; Tafalla & Hacar 2015). From these observations of molecular cloud structures at different size scales, one can visualize an evolutionary sequence of star formation starting from highly supersonic, magnetized GMCs, continuing on to filamentary dark clouds that form within them, and then on to finer filamentary substructures. Fragmentation of these filamentary structures and substructures leads to the clumps and dense cores that form protostellar clusters and protostars. Knowing the physical conditions inside filamentary clouds would provide crucial information on the formation of filamentary substructures and dense cores, and on the origin of the initial mass function (IMF) and the star formation rate. Particularly important is the characterization of the physical properties of transcritical filamentary structures whose mass per unit length is within a factor of \u223c2 of the critical line mass ${M_{\\rm crit,\\, th,\\, \\ell }}=2\\, c_{\\rm s}^2\/G$ of nearly isothermal cylindrical filaments (e.g. Ostriker 1964; Inutsuka & Miyama 1997), where cs is the isothermal sound speed. Indeed, Herschel observations suggest that transcritical filamentary structures dominate the mass function of star-forming filaments and that their fragmentation may set the peak of the prestellar core mass function and perhaps ultimately the peak of the IMF (Andr\u00e9 et al. 2019). In this paper, we report the results of polarimetric observations of the pristine section B211 of one such transcritical filament, the Taurus B211\/B213 filament, using the High-resolution Airborne Wideband Campera plus (HAWC+) onboard Stratospheric Observatory For Infrared Astronomy (SOFIA). We determine the magnetic field structure inside a filamentary cloud with filamentary substructures.","Citation Text":["Tafalla & Hacar 2015"],"Functions Text":["Dense cores, where stars form, are located along or at the intersections of some of these fine substructures (e.g."],"Functions Label":["Background"],"Citation Start End":[[1605,1625]],"Functions Start End":[[1469,1583]]}
{"Identifier":"2022MNRAS.517.1313M__Krumholz_&_McKee_2005_Instance_1","Paragraph":"Star formation is an inefficient process, as evidenced by observed gas depletion times,1 which are two orders of magnitude above the dynamical time, both in galaxies (e.g. Leroy et al. 2017; Utomo et al. 2018), and in individual giant molecular clouds (GMCs) (e.g. Krumholz & Tan 2007; Evans, Heiderman & Vutisalchavakul 2014; Heyer et al. 2016; Pokhrel et al. 2020; Hu et al. 2022). Theoretical models explain this inefficiency through a combination of mechanisms that provide support against gravitational collapse, including turbulence, magnetic fields, stellar feedback, and dynamical stabilization (Krumholz & McKee 2005; Ostriker, McKee & Leroy 2010; Federrath & Klessen 2012; Krumholz, Klein & McKee 2012b; Federrath 2013b; Padoan et al. 2014; Federrath 2015; Burkhart 2018; Meidt et al. 2018; Krumholz & Federrath 2019; Evans, Kim & Ostriker 2022). Recent progress in both theory and observations have highlighted the pivotal role that feedback, especially due to massive (main-sequence) stars, plays in star\/star-cluster formation (Krumholz et al. 2014; Krumholz, McKee & Bland-Hawthorn 2019), and the lifecycle of GMCs (see Chevance et al. 2020, 2022a for reviews). This massive-star feedback has been suggested to be largely responsible for limiting the integrated star formation efficiency (\u03f5*) to low values in typical environments, where \u03f5* is given by\n(1)$$\\begin{eqnarray}\r\n\\epsilon _* = \\frac{M_{*}}{M_{\\mathrm{gas}}},\r\n\\end{eqnarray}$$which quantifies the net efficiency of star formation over the lifetime of a GMC, i.e. the ratio of the final stellar mass M* and the available gas mass in the parent molecular cloud Mgas. Feedback achieves this by (i) disrupting GMCs in order \u223c unity dynamical time-scales, through the momentum and energy carried by feedback processes (e.g. Grudi\u0107 et al. 2018), and (ii) driving turbulent motions that could further provide support against collapse (e.g. Mac Low & Klessen 2004; Krumholz, Matzner & McKee 2006; Elmegreen 2009; Gritschneder et al. 2009; Federrath et al. 2010; Wibking, Thompson & Krumholz 2018; Gallegos-Garcia et al. 2020; Menon, Federrath & Kuiper 2020; Menon et al. 2021).","Citation Text":["Krumholz & McKee 2005"],"Functions Text":["Theoretical models explain this inefficiency through a combination of mechanisms that provide support against gravitational collapse, including turbulence, magnetic fields, stellar feedback, and dynamical stabilization"],"Functions Label":["Background"],"Citation Start End":[[604,625]],"Functions Start End":[[384,602]]}
{"Identifier":"2022MNRAS.515.5495M__Gallazzi_et_al._2008_Instance_1","Paragraph":"The stellar metallicity in the Universe evolves with redshift (Mannucci et al. 2010; Sommariva et al. 2012; Krumholz & Dekel 2012; Dayal, Ferrara & Dunlop 2013; Madau & Dickinson 2014). The metallicity at a high redshift (z > 2) is much smaller in comparison to the low redshift Universe z  2. The first-generation stars contaminate the interstellar medium and cause a chemical evolution of the Universe. We can treat the metallicity evolution with redshift by a relation \n(2)$$\\begin{eqnarray*}\r\n\\log _{10}(Z(z))= \\gamma z +\\zeta ,\r\n\\end{eqnarray*}$$where \u03b3 captures the redshift dependence and \u03b6 captures the metallicity value at z = 0 (Mannucci et al. 2010; Madau & Dickinson 2014). This relation captures the metallicity of the parent star or the gas cloud from which a star has formed. It is written to express only a mean evolution of the metallicity. Along with the mean metallicity evolution of the Universe, there is going to be a scatter in the metallicity depending on the galaxy properties. Such a source of uncertainty brings additional stochasticity to the metallicity relation. Currently, a limited number of observations (Gallazzi et al. 2008; Mannucci et al. 2010; Krumholz & Dekel 2012) are available to explore the environment dependence of the metallicity, and most of our current understandings are based on simulations(Genel 2016; Torrey et al. 2019). These studies show that the overall median metallicity dependence of the galaxies at different redshifts can be explained by power form (Pei, Fall & Hauser 1999; Young & Fryer 2007; Torrey et al. 2019). Several studies of GW merger rates and mass distribution are performed (Belczynski et al. 2002; Dominik et al. 2012, 2015; Mapelli et al. 2017; Giacobbo, Mapelli & Spera 2018; Toffano et al. 2019; van Son et al. 2022) which are motivated by these studies and show that the black hole mass distribution can exhibit a redshift dependence. The existence of any stochasticity in the galaxy metallicity distribution will also influence the mass distribution but is currently not well known. However, as the relation given in equation (1) is in terms of the logarithm of metallicity, so the impact of fluctuation around the median value depending on the individual galaxy properties is going to be a small (logarithmic) change. As we are unable to measure the host of the BBH due to a large sky localization error of the BBH, we cannot directly associate the properties of galaxies with BBH source properties. So, we can only infer an ensemble average mass distribution from the GW data and the additional stochasticity (which will depend on the host properties) will appear as an additional uncertainty in the measurement of MPISN. As a result, we consider a median distribution of galaxy metallicity and the dependence of MPISN on it.","Citation Text":["Gallazzi et al. 2008"],"Functions Text":["Currently, a limited number of observations","are available to explore the environment dependence of the metallicity"],"Functions Label":["Background","Background"],"Citation Start End":[[1138,1158]],"Functions Start End":[[1093,1136],[1205,1275]]}
{"Identifier":"2019AandA...625A.114J__Tacconi_et_al._2018_Instance_1","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"],"Functions Text":["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.,"],"Functions Label":["Background"],"Citation Start End":[[261,275],[282,286]],"Functions Start End":[[0,216]]}
{"Identifier":"2015AandA...576A...5C__Christen_&_M\u00fcller_(2003)_Instance_1","Paragraph":"Using the CASSIS2 software, we detected 8 lines of glycolaldehyde, 31 lines of the aGg\u2032 conformer of ethylene glycol, and 26 lines of methyl formate (see Table 1). The glycolaldehyde and methyl formate transitions are taken from the JPL spectroscopic database (Pickett et al. 1998), while the ethylene glycol transitions are from the CDMS catalog (M\u00fcller et al. 2001, 2005). The predictions are based on experimental data from Butler et al. (2001), Widicus Weaver et al. (2005) and Carroll et al. (2010) for glycolaldehyde, Christen et al. (1995) and Christen & M\u00fcller (2003) for ethylene glycol, and Ilyushin et al. (2009) for methyl formate. The frequencies of five of the detected glycolaldehyde lines were directly measured in the laboratory (Butler et al. 2001). Some of the lines result from a blending of several transitions of the same species. The lines that are strongly blended with other species are not listed in Table\u20091. All three species are emitted very compactly at the position of the continuum peak (\\hbox{$\\alpha_{2000}=03^{\\rm h}28^{\\rm m}55\\fs57$}\u03b12000=03h28m55.s57, \\hbox{$\\delta_{2000}=31\\degr14\\arcmin37\\farcs1$}\u03b42000 = 31\u00b014\u203237 .\u030b 1). The angular sizes obtained with a circular Gaussian fit in the (u, \u03bd) plane vary from a point source to a maximum of 1\u2033 depending on the transition. The line fluxes listed in Table 1 were measured at the continuum peak position with the CASSIS software using a Gaussian fitting method (Levenberg-Marquardt algorithm). The lines that are contaminated in the wings by other transitions are consequently fitted with a sum of Gaussians. We carefully checked that the derived full widths at half maximum (FWHM) are consistent with the other line measurements. The average FWHM is about 4.5 km\u2009s-1 at 317 GHz, and 5.0 km\u2009s-1 at 225 and 242 GHz. The widths of the methyl formate lines at 87 GHz are quite broad (~12 km\u2009s-1). It is consequently difficult to completely exclude an additional flux contribution from other species. The variation of FWHM with the frequency can be explained by the spectral resolution of the observations that decreases toward the lower frequencies. ","Citation Text":["Christen & M\u00fcller (2003)"],"Functions Text":["The predictions are based on experimental data from","for ethylene glycol"],"Functions Label":["Uses","Uses"],"Citation Start End":[[551,575]],"Functions Start End":[[375,426],[576,595]]}
{"Identifier":"2017ApJ...837..109L__parsec,_Cisternas_et_al._2013_Instance_1","Paragraph":"It has been generally believed (see Kormendy & Ho 2013; Kormendy 2016, for recent reviews) that, unlike massive galaxies at high redshifts (e.g., \n\n\n\n\n\n) whose evolution is driven by a major merger, low-redshift galaxies largely evolve through secular evolution (i.e., in a slow and gentle manner) driven by internal processes within galactic disks and\/or by environmental effects such as harassment or nurturing (see Kormendy & Kennicutt 2004 for details; see also Cowie et al. 1996; Conselice et al. 2014 for the cosmic evolution). It has also been generally believed accordingly that most activity (namely active galactic nuclei\u2014AGNs) of supermassive black holes (BHs) at the centers of galaxies at low redshifts (e.g., \n\n\n\n\n\n, Cisternas et al. 2011; and even up to \n\n\n\n\n\n, Kocevski et al. 2012) appear to be fueled by the random accretion of gas via internal, secular processes working close to the BH (say, within a few hundred parsec, Cisternas et al. 2013; also see Hopkins & Hernquist 2006). In these cases, there is no connection between AGN activity and major mergers of their host galaxies. This non-connection appears to be particularly true for low-z AGNs hosting intermediate-mass black holes (IMBHs):12\n\n12\n Following Greene & Ho (2007) and Dong et al. (2012), hereinafter we refer to BHs with \n\n\n\n\n\n \n\n\n\n\n\n at the centers of galaxies as \u201clow-mass\u201d or \u201cintermediate-mass\u201d BHs; accordingly, for the ease of narration wherever it is not ambiguous, hereinafter we refer to AGNs hosting low-mass BHs as low-mass AGNs or IMBH AGNs. Normally, we prefer \u201cintermediate-mass BHs\u201d to \u201clow-mass BHs\u201d because of the possible confusion of the latter with the stellar-mass BHs in low-mass X-ray binaries (LMXBs).\n according to the analysis of their images by the Hubble Space Telescope (HST), the majority of low-mass AGNs live in late-type disk galaxies without a classical bulge (Greene et al. 2008; Jiang et al. 2011). Indeed, the accretion rate for an IMBH is so tiny (0.05 \n\n\n\n\n\n yr\u22121 even if at the maximum Eddington accretion) that a steady supply of fuel, in the form of stellar mass loss from evolved stars or Bondi accretion of hot gas in the innermost regions, is available readily much more than required (Ho 2008). \u201cIn fact, the paradox for local BHs is not whether there is enough fuel to light them up. Rather, the puzzle is how to keep them so dim despite the ready abundance of in situ gas.\u201d (Kormendy & Ho 2013; see also Ho 2008). That is, the real problem seems to be this (see Ho 2009): there must be some mechanism (yet to know) to hinder the innermost fueling process.13\n\n13\n According to W.-M. Gu (2016, private communication), AGN outflows, which can be launched even at very low accretion rates (e.g., Wang et al. 2013; Gu 2015), provide such a mechanism.\n\n","Citation Text":["Cisternas et al. 2013"],"Functions Text":["It has also been generally believed accordingly that most activity (namely active galactic nuclei\u2014AGNs) of supermassive black holes (BHs) at the centers of galaxies at low redshifts","appear to be fueled by the random accretion of gas via internal, secular processes working close to the BH (say, within a few hundred parsec,"],"Functions Label":["Background","Background"],"Citation Start End":[[941,962]],"Functions Start End":[[534,715],[799,940]]}
{"Identifier":"2021AandA...656A.148R__in_2012_Instance_1","Paragraph":"After the gravitational collapse and if the total mass of the individual cloud is approximately the mass of the Sun (2 \u00d7 1030 Kg; (van Dishoeck 2014), a new astrophysical system forms that is dominated gravitationally by a low-mass protostar known as a young stellar object (YSO). The protostar is in the center of the system and is surrounded by a Keplerian-rotating envelope of dust and gas, that is gravitationally connected and in which the angular momentum is conserved (Cassen & Moosman 1981). The distance to the star will define the surrounding energy, which is dominantly thermal and capable of heating the more distant dustgrain mantles. The processes progressively inject chemical constituents into the gas-phase (Ceccarelli et al. 2001)and determine what types of physical-chemical processes govern in every region of the disk. Circumstellar envelopes oflow-mass protostars (CELMP) are environments that are extraordinarily rich in organic molecules as H2CO or HNCO and in iCOMs such as CH3CN, CH3CHO, and C2 H5OH (Sch\u00f6ier et al. 2002), in addition to other O- and N-bearing complexes (J\u00f8rgensen et al. 2012). IRAS 16293-2422 (hereafter IRAS 16293) was chosen for this work as the quintessential example of such an object. IRAS 16293 is adequate for this role because formaldehyde has previously been detected with notable variations in abundances in three differentiated regions of its protodisk (Ceccarelli et al. 2001; J\u00f8rgensen et al. 2012; Jaber et al. 2014; van der Wiel et al. 2019). Additionally, glycolaldehyde (C2 H4O2), the simplest form of sugar and one of the first intermediates in the formose reaction (Larralde et al. 1995), has also been detected in this object for the first time in space in 2012 (J\u00f8rgensen et al. 2012). The more distant regions of the disk (>4000 AU) contain species such as CO, H2CO, CH3OH, or H2O (van Dishoeck et al. 1995) at low temperatures (10\u201320 K) where the gas is slightly warmer than the dust grains as they are tightly coupled to them. Energy from collisions among dust and gas is therefore considered to be the main heating mechanism in this region. Nevertheless, the thermal energy generated by those collisions is not enough to activate the formation of molecules with barriers at or below 20 K (0.0397 kcal mol\u22121). The gas column density of molecular hydrogen has been estimated to be N(H2) = 1.3 \u00d7 1023 cm\u22122 (Ward-Thompson et al. 1999), with a fractional abundance of formaldehyde \u2013 N(H2CO)\/N(H2) \u2013 ~ 4 \u00d7 10\u221210 cm\u22123 in the gas-phase (Ceccarelli et al. 2001). When the temperature rises above ~ 20 K, CO starts to desorb from the ice grains and enters the gas-phase with an increase of ~ 103 cm\u22123 in detected densities with respectto H2 (Cassen & Moosman 1981; Aikawa et al. 2015). Formaldehyde starts to deplete from frozen grains at around ~ 40 K and it is fully desorbed at ~ 60 K (Ceccarelli et al. 2001). The additional H2CO mixes with the existing circumstellar mass of gas, which may justify why at ~ 700 AU from the core and at gas temperatures between 80 and 100 K ~ 50 kcal mol\u22121), the detected fractional abundances of formaldehyde reach N(H2CO)\/N(H2) = ~ 4.0 \u00d7 10\u22129 cm\u22123 (Ceccarelli et al. 2001). This implies an H2 column density that is estimated to be N(H2) = ~ 5.0 \u00d7 1021 cm\u22122 (Bottinelli et al. 2014). The inner part of the envelope at ~ 150 AU in a region with temperatures of 100\u2013150 K has a higher density of formaldehyde with an N(H2CO)\/N(H2) of ~ 10 \u00d7 10\u22127 cm\u22123, as well as an increase in fractional abundances for H2O (which desorbs from iced mantles at ~80 K) (Ceccarelli et al. 2001). This region also produces new molecules principally due to the thermal energy emitted from the YSO (99.99 kcal mol\u22121). One example is trans-HONO. This chemical compound has recently been detected for first time in space and in this part of the disk (Coutens et al. 2019). Its proposed formation has inspired some reactions proposed in this work that may lead to H2CO. In this region, a column density for molecular hydrogen is considered like that in region II (N(H2) = ~ 5.0 \u00d7 1021 cm\u22122). The three regions of IRAS 16293 dictate the physical parameters for the presentation of our computations which are defined as follows:\n\nRegion I\/d ~ 4000 AU, Tgas = 20 K, Pgas = 2.29 \u00d7 107 K cm\u22123;\nRegion II\/d~ 700 AU, Tgas = 80 K, Pgas = 7.06 \u00d7 107 K cm\u22123;\nRegion III\/d~150 AU, Tgas = 150 K, Pgas = 6.08 \u00d7 108 K cm\u22123.","Citation Text":["J\u00f8rgensen et al. 2012"],"Functions Text":["Additionally, glycolaldehyde (C2 H4O2), the simplest form of sugar and one of the first intermediates in the formose reaction",", has also been detected in this object for the first time in space in 2012"],"Functions Label":["Background","Background"],"Citation Start End":[[1728,1749]],"Functions Start End":[[1503,1628],[1651,1726]]}
{"Identifier":"2021ApJ...909...65K__Liu_et_al._2014_Instance_1","Paragraph":"Various groups around the world have proposed different models to explain the formation of these two peculiar classes of SNe Ia. Sub-Chandrasekhar limiting-mass WDs were believed to be formed by merging two sub-Chandrasekhar mass WDs (double-degenerate scenario), leading to another sub-Chandrasekhar mass WD, exploding due to accretion of a helium layer (Hillebrandt & Niemeyer 2000; Pakmor et al. 2010). On the other hand, the super-Chandrasekhar WDs were often explained by incorporating different physics, such as a double-degenerate scenario (Hicken et al. 2007), presence of magnetic fields (Das & Mukhopadhyay 2013, 2014), presence of a differential rotation (Hachisu et al. 2012), presence of charge in the WDs (Liu et al. 2014), ungravity effect (Bertolami & Mariji 2016), lepton number violation in magnetized WD (Belyaev et al. 2015), generalized Heisenberg uncertainty principle (Ong 2018), and many more. However, none of these theories can self-consistently explain both of the peculiar classes of WDs. Moreover, each of these has some caveats or incompleteness, mostly based on the stability (Komatsu et al. 1989; Braithwaite 2009). Furthermore, numerical simulations showed that a merger of two massive WDs could never lead to a mass as high as 2.8M\u2299 owing to the off-center ignition and formation of a neutron star rather than an (overluminous) SN Ia (Saio & Nomoto 2004; Martin et al. 2006). Hence, all the conventional pictures failed to explain the inferred masses of both the sub- and super-Chandrasekhar progenitor WDs and also both classes of progenitor WDs simultaneously by invoking the same physics. Moreover, each of the theories can explain only one regime of SN Ia, but it seems more likely that the nature would prefer only one scenario\/physics to exhibit the same class of SNe. Whether it be an under- or overluminous SN Ia, other physics such as the presence of Si, etc., remains the same. Therefore, we seem to require just one theory to explain all the SNe Ia.","Citation Text":["Liu et al. 2014"],"Functions Text":["On the other hand, the super-Chandrasekhar WDs were often explained by incorporating different physics, such as a","presence of charge in the WDs"],"Functions Label":["Background","Background"],"Citation Start End":[[720,735]],"Functions Start End":[[406,519],[689,718]]}
{"Identifier":"2015MNRAS.451.1528P__Brun_et_al._2014_Instance_1","Paragraph":"Currently, there is an extensive data base of observations of magnetic activity on the main-sequence stars (see e.g. reviews by Donati & Landstreet 2009; Reiners 2012). Cool stars with outer convective envelope are of particular interests because they are Sun-like. It is believed that the magnetic activity on the solar-like stars results from large-scale dynamo processes driven by turbulent convection and rotation (Brandenburg & Subramanian 2005). Observations (e.g. B\u00f6hm-Vitense 2007; Donati & Landstreet 2009; Katsova et al. 2010; Saar 2011; Katsova, Livshits & Mishenina 2013; Marsden et al. 2014; Vidotto et al. 2014), as well as the 2D mean-field models of the angular momentum balance (Ruediger 1989; Kitchatinov & R\u00fcdiger 1999; Kitchatinov 2013) and the 3D numerical simulations (Miesch & Toomre 2009; Hotta & Yokoyama 2011; Guerrero et al. 2013b; Brun et al. 2014; K\u00e4pyl\u00e4, K\u00e4pyl\u00e4 & Brandenburg 2014) show that parameters of the differential rotation and convection, e.g. the typical size and turnover time of convective flows, depend on the general stellar parameters, such as mass, age, the spectral class and the rotation rate. The mass of a star and its Rossby number, which is the ratio of the period of rotation and a typical turnover time of convection, are likely the most important parameters governing the stellar dynamo (Donati & Landstreet 2009; Morin et al. 2013). The diagram 3 in the paper by Donati & Landstreet (2009) shows an increase of the magnetic activity with decrease of the Rossby number and the mass of a star. These parameters determine the topology of the large-scale magnetic field, as well. It is found that the axisymmetric solar-type dynamo can operate in stars with mass about 1\u2009M\u2299, and with periods of rotation longer than 10 d. Observations, also show pieces of evidence that solar analogues with period of rotation smaller than 10 d, may have the substantial non-axisymmetric components of the large-scale magnetic field (Donati & Landstreet 2009; Folsom et al. 2014).","Citation Text":["Brun et al. 2014"],"Functions Text":["and the 3D numerical simulations","show that parameters of the differential rotation and convection, e.g. the typical size and turnover time of convective flows, depend on the general stellar parameters, such as mass, age, the spectral class and the rotation rate."],"Functions Label":["Background","Background"],"Citation Start End":[[859,875]],"Functions Start End":[[757,789],[912,1141]]}
{"Identifier":"2015AandA...579A..51B__Egan_et_al._1998_Instance_1","Paragraph":"Even though the understanding of high-mass star formation has made tremendous progress over the past decade (Beuther et al. 2007; Zinnecker & Yorke 2007; Klessen 2011; Tan et al. 2014), the initial conditions are still poorly constrained. It is known, that stars predominantly form in clusters (Lada & Lada 2003), which is especially true for high-mass stars (e.g., de Wit et al. 2005; Gvaramadze et al. 2012). On observational grounds, the first evolutionary stage of such clusters spawning future high-mass stars might in general be termed pre-protocluster cores (Evans et al. 2002). Their appearance and morphology can be different depending on the environment. Objects of this sort have been found in several (sub-)millimeter surveys (e.g., Klein et al. 2005; Beuther & Sridharan 2007). Also the ISO satellite mission has resulted in a list of such objects revealed by far-infrared observations at 170 \u03bcm within the ISO Serendipity Survey (ISOSS, Bogun et al. 1996). For these clumps, cold dust and gas temperatures have been established in follow-up investigations (Krause et al. 2003, 2004). The most prominent variety of young massive clumps are the infrared dark clouds (IRDCs). They were discovered as dark silhouettes against the galactic background at 8 and 15 \u03bcm with the Midcourse Space Experiment (MSX, Egan et al. 1998) and ISO (Perault et al. 1996). ISO, MSX, the Spitzer Space telescope, and further (sub-)millimeter observations have helped for studying these objects in detail (e.g., Simon et al. 2006; Hennemann et al. 2008; Vasyunina et al. 2009; Ragan et al. 2009; Peretto & Fuller 2009). Because IRDCs can only be seen in absorption against a strong infrared background, their location is mainly within the disk toward the inner quadrants of our Milky Way, whereas ISOSS sources are more widely distributed. In fact, the Milky Way midplane had to be avoided for ISOSS because of saturation. IRDCs can have masses up to several thousand solar masses, and the more massive ones are explicitly thought to be progenitors of star clusters (e.g., Wyrowski 2008; Rathborne et al. 2010; Henning et al. 2010; Ragan et al. 2012a). On average the ISOSS sources have masses less than 500 M\u2299 (Hennemann et al. 2008). While their masses often are somewhat lower than for large high-contrast IRDCs, the general characteristics such as low temperature and low turbulent linewidth are similar to IRDCs. Therefore, ISOSS sources are also considered to be early evolutionary stages of intermediate- to high-mass star formation in more isolated regions. Consequently, the better characterization of such clumps (ISOSS and IRDCs) is an important step toward understanding the initial conditions of high-mass star formation. ","Citation Text":["Egan et al. 1998"],"Functions Text":["They were discovered as dark silhouettes against the galactic background at 8 and 15 \u03bcm with the Midcourse Space Experiment (MSX,"],"Functions Label":["Uses"],"Citation Start End":[[1317,1333]],"Functions Start End":[[1187,1316]]}
{"Identifier":"2019MNRAS.488.5029H__Malhotra_et_al._2001_Instance_2","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":["Malhotra et al. 2001"],"Functions Text":["The comparison sample is compiled from the literature up to z \u223c 3"],"Functions Label":["Uses"],"Citation Start End":[[1410,1430]],"Functions Start End":[[1343,1408]]}
{"Identifier":"2019AandA...630A.131M__Molina_et_al._2009_Instance_1","Paragraph":"X-rays emerging from active galactic nuclei (AGNs) are the result of an inverse-Compton process occurring in the proximity of the central black hole (BH), where optical-UV photons arising from the accretion disc are inverse-Compton scattered by hot electrons in an optically thin, compact corona (e.g. Haardt & Maraschi 1991, 1993; Haardt et al. 1994, for details on the two-phase model). Such a Comptonisation mechanism accounts for the power-law-like shape of the X-ray primary continuum emission and the high-energy roll-over observed in various nearby AGNs (e.g. Nicastro et al. 2000; Perola et al. 2002; De Rosa et al. 2002; Molina et al. 2009, 2013; Malizia et al. 2014; Ricci et al. 2018). Broadband X-ray spectral investigations are of primary importance in studying AGN Comptonisation properties. Indeed, as extensively discussed in the literature (e.g. Ghisellini 2013) both the photon index (\u0393) and high-energy cut-off (Ec) of the X-ray primary emission depend on the intrinsic properties of the Comptonising medium, namely its temperature, optical depth, and geometry. Therefore, the interplay between coronal parameters and the AGN X-ray spectral shape has been the object of several investigations, especially with observatories capable of detecting hard X-rays. Dadina (2007), using BeppoSAX data, collected and studied the photon index and Ec of a sample of AGNs (see also Perola et al. 2002, for previous results), while similar works were performed on INTEGRAL data (e.g. Bassani et al. 2006; Molina et al. 2009; Malizia et al. 2014), and, in the context of the BAT AGN Spectroscopic Survey, by Ricci et al. (2018). Subsequently, NuSTAR (Harrison et al. 2013), thanks to its unprecedented effective area above 10 keV, greatly helped in studying the exponential cut-offs of the nuclear continuum in several AGNs (see e.g. Fabian et al. 2015, 2017; Tortosa et al. 2018a). These space missions gave rise to a substantial corpus of high-energy cut-off and photon index measurements.","Citation Text":["Molina et al. 2009","Molina et al. 2009"],"Functions Text":["Such a Comptonisation mechanism accounts for the power-law-like shape of the X-ray primary continuum emission and the high-energy roll-over observed in various nearby AGNs (e.g.","while similar works were performed on INTEGRAL data (e.g."],"Functions Label":["Background","Background"],"Citation Start End":[[630,648],[1511,1529]],"Functions Start End":[[389,566],[1432,1489]]}
{"Identifier":"2021ApJ...906..105C__Huang_et_al._2019_Instance_1","Paragraph":"Even though \u03a9 is only a function of \u03a8, the physical connection between the two at the foot-point is not very obvious.9\n\n9\nThey may be constrained by proper boundary conditions (e.g., Contopoulos et al. 2013).\n Likely most researchers, here we assume that they generally follow the ansatz\n13\n\n\n\n\n\nIn the region approaching the polar axis, one expects that the magnetic flux vanishes and thus \n\n\n\n\n\n. Therefore, in the region near the polar axis, mathematically \u03bb \u2265 0 is required to guarantee a finite value of \u03a9 (notice that \u03bb > 0 is likely unphysical). Magnetic field lines just around the polar axis would connect to the central CO, which is expected to rotate with a roughly constant angular velocity10\n\n10\nIn the case of threading a BH, magnetic field lines at different polar angles may not rotate with exactly the same frequency; see discussion in Section 8.\n (see details below), implying \u03bb = 0 for the case of magnetic field lines threading the CO. In the region where magnetic field lines are threading the AD, one usually expects \u03bb  0. The choice of \u03a6, which determines the toroidal magnetic field, is important (e.g., Camenzind 1987; Sulkanen & Lovelace 1990). Physically, it is the rotation that develops a toroidal magnetic field and a poloidal electric field in the lab frame (or the inertial frame), which seems to imply that \u03a6 would not be chosen arbitrarily and would be self-consistently determined by the MHD equations, given the \u03a8 and \u03a9 specified (e.g., Beskin & Tchekhovskoy 2005; Contopoulos et al. 2013; Huang et al. 2019, 2020). To solve the rotation term, Equation (11), one can see that in the case of\n14\n\n\n\n\n\nthe function \u03a8 can be variable-separated,\n15\n\n\n\n\n\nwhere the subscript \u201cr\u201d denotes \u201crotation.\u201d The negative sign in Equation (14) exists to guarantee that there is always a swept-back magnetic field line with respect to the direction of rotation. This relation implies that the strength of the toroidal magnetic field would be proportional to the angular velocity and the enclosed magnetic flux, which can be reasonably understood because it is the rotation of the poloidal magnetic field (related to the magnetic flux) that produces the toroidal magnetic field. In terms of Equation (15), Equation (11) can be expressed as\n16\n\n\n\n\n\nwhere \n\n\n\n\n\n, \n\n\n\n\n\n, \n\n\n\n\n\n, and \n\n\n\n\n\n. The left-hand side of the above equation is only a function of r (the r component), and the right-hand side is only a function of \u03b8 (the \u03b8 component). Both equal the same constant. Let us set this constant as \n\n\n\n\n\n, a choice making the solution of the r component equation concise:\n17\n\n\n\n\n\nWe introduce a new variable,\n18\n\n\n\n\n\nwhich makes the \u03b8-component equation become\n19\n\n\n\n\n\nwhere \n\n\n\n\n\n and \n\n\n\n\n\n. Let us first consider its asymptotic properties. In the case of \u03b8 \u226a 1 (keeping in mind \n\n\n\n\n\n when \n\n\n\n\n\n), the leading-order terms give\n20\n\n\n\n\n\nIt is clear that this equation has a solution of the form\n21\n\n\n\n\n\nTherefore, a magnetic field line forms a \u201cgeneral parabolic\u201d configuration11\n\n11\nIn the case of \u03bd = 0, \u03a8 is only a function of \u03b8, which presents a monopole solution. The case of \u03bd = \u03b2 leads to \u03a8 \u221d R, which gives a cylindrical solution.\n at \u03b8 \u226a 1, that is, \u03a8 \u221d r\u03bd\u03b8\u03b2. Now, let us consider what value \u03b2 might take. In order to do this, we have to consider the higher-order terms of Equation (19), which may be comparable to the nonrotation term. Therefore, we have to consider the original Equation (10). Let us substitute a general form of \n\n\n\n\n\n (with coefficients a1, a2, a3... to be determined) into the original Equation (10). One gets the first two leading-order terms (note that a1 = 0):\n22\n\n\n\n\n\nIt can be seen that, for any values of \u03b2, \u03bb, and \u03bd, one always has an a2 to make the second term vanish, whereas the first term yields \u03b2 \u2248 2 (\u03b2 = 0 corresponds to the nonrotation case). We note that this choice cannot guarantee that the higher-order terms vanish, so this solution is only an approximation. In the case of \u03b2  2\u03bd (i.e., \u03bd \u2273 1), the first term dominates over the second term in Equation (22), whereas in the case of \u03b2 > 2\u03bd (i.e., \u03bd \u2272 1), the second term dominates over the first one. We therefore expect that the approximation \u03b2 \u2248 2 would be more efficient in the former case (i.e., \u03bd \u2273 1). In another aspect, from Equation (2), one has \n\n\n\n\n\n, which implies that the choice of \u03b2 = 2 can avoid singularity or vanishing magnetic fields on the magnetic polar axis (either B\u03b8 or B\u03d5 vanishes). This choice is also supported by some numerical simulations, which showed that this choice corresponds to a minimum torque (i.e., the least amount of toroidal magnetic fields) and is the one picked by a \u201creal\u201d system (see, e.g., Michel 1969; Contopoulos 1995; Narayan et al. 2007; Tchekhovskoy et al. 2008). Although the coefficient \u03b2 \u2248 2 is derived asymptotically (\u03b8 \u226a 1), which must hold throughout the jet region because \u03a6, \u03a9, and \u03a8 are each conserved along magnetic field lines. In Appendix A, we come back to this question and present another proof of the relation \u03a6 \u2248 \u22122\u03a9\u03a8 based on a more physical consideration, showing that this relation is only valid in the limit of a highly magnetized jet flow (e.g., Lyubarsky 2009).","Citation Text":["Huang et al. 2019"],"Functions Text":["Physically, it is the rotation that develops a toroidal magnetic field and a poloidal electric field in the lab frame (or the inertial frame), which seems to imply that \u03a6 would not be chosen arbitrarily and would be self-consistently determined by the MHD equations, given the \u03a8 and \u03a9 specified (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[1526,1543]],"Functions Start End":[[1171,1472]]}
{"Identifier":"2018AandA...611A..74R__Grady_et_al._2013_Instance_1","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"],"Functions Text":["Based on its SED, MWC 758 has been classified as a pre-transition disk"],"Functions Label":["Background"],"Citation Start End":[[1120,1137]],"Functions Start End":[[1048,1118]]}
{"Identifier":"2022AandA...661A..10B__Ghirardini_et_al._2021a_Instance_4","Paragraph":"It is also possible that these clusters have a smaller extent and can just be missed by our extent selection as our detection algorithm sets the extent to zero if it is smaller than 6 (Brunner et al. 2022). Following the method presented in Ghirardini et al. (2021a), we estimated several dynamical properties of the clusters in the point source sample and compared them with the extent-selected sample presented in Ghirardini et al. (2021a). In Fig. 5 we compare the distributions of the core radii (Rcore) constrained by the V06 model and the concentration parameter (cSB) between these two samples. The concentration parameter is defined as the ratio of the surface brightness within 0.1 R500 to the surface brightness within R500 (Ghirardini et al. 2021a; Santos et al. 2008; Maughan et al. 2012). Intuitively, the expectation is that the smaller the core radius, the more compact the cluster. The left panel of Fig. 5 clearly shows that the clusters in the point source sample have relatively smaller core radii, hence the emission is more concentrated in a smaller area. Consistently, the concentration of the point source sample shows a clear excess in higher values than the extent-selected sample, indicating that a significantly larger fraction of cool-core clusters and clusters host a central AGN. We performed the same experiment by applying cuts in flux 1.5 \u00d7 10\u221214 ergs s\u22121 cm\u22122 and in detection likelihood to test whether the clusters are missed by the extent selection because they are fainter and\/or more compact than the extent-selected clusters. The distribution of number density, core radius, and concentration parameter remains the same, indicating that the population of clusters in the point source catalog is more compact than the extent-selected sample. The extent-selected sample does not show a clear bias toward cool-core clusters or clusters with a central AGN, but contains the fraction of cool-cores is similar to that of SZ surveys (Ghirardini et al. 2021a). In this sample, we observe the opposite trend.","Citation Text":["Ghirardini et al. 2021a"],"Functions Text":["The distribution of number density, core radius, and concentration parameter remains the same, indicating that the population of clusters in the point source catalog is more compact than the extent-selected sample. The extent-selected sample does not show a clear bias toward cool-core clusters or clusters with a central AGN, but contains the fraction of cool-cores is similar to that of SZ surveys","In this sample, we observe the opposite trend."],"Functions Label":["Compare\/Contrast","Differences"],"Citation Start End":[[1967,1990]],"Functions Start End":[[1566,1965],[1993,2039]]}
{"Identifier":"2017ApJ...849..109P__Lee_et_al._2014_Instance_1","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"],"Functions Text":["and have coupled the code to SN ejecta models"],"Functions Label":["Uses"],"Citation Start End":[[575,590]],"Functions Start End":[[528,573]]}
{"Identifier":"2022MNRAS.509.5155R__Richard_et_al._2009_Instance_1","Paragraph":"In addition to being useful laboratories for studying the evolution of galaxies in dense environments, galaxy clusters can be effective gravitational lenses (for a review, see Kneib & Natarajan 2011). Gravitational lensing is the deflection of light by intervening mass, that can produce highly magnified and distorted images of background galaxies. The large mass and solid angle covered by highly concentrated galaxy clusters make them ideal gravitational lenses, which can be used as \u2018cosmic telescopes\u2019 to observe very distant galaxies (e.g. Richard et al. 2009). The amplification of sources through gravitational lensing has been extremely effective in observing faint, distant sources across the electromagnetic spectrum, including continuum and spectral line emission in the radio domain (e.g. Carilli & Walter 2013, and references therein). While molecular gas has been studied across the Universe up to z \u223c 1 and beyond, H\u2009i emission remains undetected through lensing, with two searches for galaxy\u2013galaxy lensed H\u2009i sources at z \u223c 0.4 (Hunt, Pisano & Edel 2016; Blecher et al. 2019). Gravitational lensing conserves surface brightness while increasing the solid angle of the source, boosting the observed flux. This amplification \u03bc can facilitate the detection of unresolved lensed sources, which maximises their detection probability, and reduces the integration time needed for a given source by \u03bc2. Next-generation cm-wavelength interferometers are now sensitive enough to observe the higher redshift H\u2009i Universe, and the detection of gravitationally lensed H\u2009i is probable in new surveys (Deane, Obreschkow & Heywood 2015). The detection of lensed H\u2009i behind intermediate-redshift galaxy clusters will provide a deep cosmic view of H\u2009i emission in galaxies, pre-SKA era, within a fraction of the observation time of unlensed detections. Successful lensed H\u2009i detections, along with readily detected CO emission lines, will constrain the H\u2009i\/H2 ratio at these redshifts, an important parameter in understanding galaxy evolution over cosmic time (Obreschkow & Rawlings 2009).","Citation Text":["Richard et al. 2009"],"Functions Text":["The large mass and solid angle covered by highly concentrated galaxy clusters make them ideal gravitational lenses, which can be used as \u2018cosmic telescopes\u2019 to observe very distant galaxies (e.g."],"Functions Label":["Uses"],"Citation Start End":[[546,565]],"Functions Start End":[[350,545]]}
{"Identifier":"2020MNRAS.498.6069P__Dressler_1980_Instance_1","Paragraph":"The present Universe is full of galaxies that come in various shapes and sizes with different mass, luminosity, colour, star formation rate (SFR), metallicity, and H\u2009i content. Understanding the galaxy properties and their evolution is an important goal of cosmology. The modern galaxy surveys, 2dF Galaxy Redshift Survey (2dFGRS; Colless et al. 2001) and Sloan Digital Sky Survey (SDSS; Strauss et al. 2002) reveal that the galaxies are distributed in the cosmic web (Bond, Kofman & Pogosyan 1996), which is an interconnected web-like network comprising of different types of environments such as filaments, sheets, knots, and voids. The galaxy properties vary across the different environments in the cosmic web. For example, the well-known density\u2013morphology relation reveals that the ellipticals are preferably found inside the dense groups and clusters, whereas the spirals are intermittently distributed in the fields (Hubble 1936; Zwicky 1968; Oemler 1974; Dressler 1980; Goto et al. 2003). These findings are further supported by other studies with two-point correlation function (Willmer, da Costa & Pellegrini 1998; Brown, Webstar & Boyle 2000; Zehavi et al. 2005), genus statistics (Hoyle et al. 2002; Park et al. 2005), and filamentarity (Pandey & Bharadwaj 2005, 2006) of the galaxy distribution. It is now well known that many other galaxy properties are strongly sensitive to their environment (Davis & Geller 1976; Guzzo et al. 1997; Zehavi et al. 2002; Blanton et al. 2003; Einasto et al. 2003a; Hogg et al. 2003; Kauffmann et al. 2004; Mouhcine, Baldry & Bamford 2007; Koyama et al. 2013). The formation and evolution of galaxies are known to be driven by accretion, tidal interaction, merger, and various other secular processes. These physical processes are largely determined by the environment of the galaxies. The environment thus play a central role in the formation and evolution of galaxies and the study of the environmental dependence of the galaxy properties provides crucial inputs to the theories of galaxy formation and evolution.","Citation Text":["Dressler 1980"],"Functions Text":["For example, the well-known density\u2013morphology relation reveals that the ellipticals are preferably found inside the dense groups and clusters, whereas the spirals are intermittently distributed in the fields"],"Functions Label":["Background"],"Citation Start End":[[964,977]],"Functions Start End":[[715,923]]}
{"Identifier":"2015MNRAS.446.3002B__Diemand,_Moore_&_Stadel_2004_Instance_1","Paragraph":"We recall that the velocity anisotropy profile is given by a combination of the radial and tangential velocity dispersion:\n\n(15)\n\n\\begin{equation}\n\\beta _{\\rm ani}(r)\\equiv 1-\\frac{\\bar{v_{\\theta }^2}(r)}{\\bar{v_r^2}(r)}\\,.\n\\end{equation}\n\nDue to the lack of observational constraints on this quantity, the first anisotropy profiles discussed in the literature were based on analytical studies aiming at building dynamical models (in spherical symmetry) with self-consistent stellar phase-space distribution functions. Many such models have simple anisotropy profiles that are either constant or change from isotropic near the centre to radial at large radius (e.g. Osipkov 1979; Merritt 1985, see below). More recently, indications of radial anisotropy in the outer regions of DM haloes have been obtained from numerical simulations (e.g. Diemand, Moore & Stadel 2004). In the inner region, a strong anisotropy can be generated by dynamical formation and evolution processes. To better describe these profiles, Baes & van Hese (2007) introduced a technique to construct dynamical models with arbitrary mass density and anisotropy profiles. These three different families of anisotropy profiles are described below and will be explored in Section 5.\n\nThe constant anisotropy modelling (e.g. used by Charbonnier et al. 2011) simply reads\n\n(16)\n\n\\begin{equation}\n\\beta _{\\rm ani}^{\\rm Cst}(r)=\\beta _0.\n\\end{equation}\n\n\nThe Osipkov\u2013Merritt profile (Osipkov 1979; Merritt 1985) is parametrized as\n\n(17)\n\n\\begin{equation}\n\\beta _{\\rm ani}^{\\rm Osipkov}(r)=\\frac{r^2}{r^2+r_a^2},\n\\end{equation}\n\nwith a single free scale parameter ra which locates the transition from \u03b2ani = 0 in the inner parts (isotropic) to 1 at large radii (full radial anisotropy).\nThe Baes and van Hese profile (Baes & van Hese 2007) is more general and is written as\n\n(18)\n\n\\begin{equation}\n\\beta _{\\rm ani}^{\\rm Baes}(r) =\\frac{\\beta _0 + \\beta _\\infty (r\/r_a)^\\eta }{1+(r\/r_a)^\\eta }\\,,\n\\end{equation}\n\nwhere the four parameters are the central anisotropy \u03b20, the anisotropy at large radii \u03b2\u221e, and the sharpness of the transition \u03b7 at the scale radius ra. The Osipkov\u2013Merritt profile is recovered when using \u03b20 = 0, \u03b2\u221e = 1 and \u03b7 = 2.\n","Citation Text":["Diemand, Moore & Stadel 2004"],"Functions Text":["More recently, indications of radial anisotropy in the outer regions of DM haloes have been obtained from numerical simulations (e.g."],"Functions Label":["Background"],"Citation Start End":[[840,868]],"Functions Start End":[[706,839]]}
{"Identifier":"2017MNRAS.471.3057M__Bovy_et_al._2016b_Instance_1","Paragraph":"We have performed the first detailed dissection of the stellar populations of the Milky Way disc in age, [Fe\/H] and $\\mathrm{[ \\alpha \\mathrm{\/Fe]}}$ space, bridging the gap between the detailed observational understanding of MAPs (e.g. Bovy et al. 2012b, 2016b) and the plethora of studies of co-eval stellar populations in simulated galaxies (e.g Bird et al. 2013; Stinson et al. 2013; Martig et al. 2014a). We have placed novel constraints on models for the formation of the Milky Way disc by combining detailed density models fit to the mono-age, mono-[Fe\/H] populations of the low and high $\\mathrm{[ \\alpha \\mathrm{\/Fe]}}$ disc, with surface mass density contributions calculated on the basis of these density fits and stellar evolution models. We summarize our key results as follows:\nRadial and vertical profiles: The mono-age, mono-[Fe\/H] populations of the $\\mathrm{[ \\alpha \\mathrm{\/Fe]}}$ poor disc are well fitted by a radially broken exponential, with a peak radius, Rpeak, that varies as a function of age and [Fe\/H]. We find that the distance between Rpeak's of the low and high [Fe\/H] populations increases with age, which we interpret as evidence for a decreasing [Fe\/H] gradient with time (e.g. Anders et al. 2017). The radial variation of the stellar surface density of the high $\\mathrm{[ \\alpha \\mathrm{\/Fe]}}$ mono-age populations is found to have insignificant breaks, and they are better fit by a single exponential in this disc region. As these populations are the oldest, this may be a sign of the disc evolution washing out the density peak over time, or may point to a different formation scenario for high $\\mathrm{[ \\alpha \\mathrm{\/Fe]}}$ stars, where no density peak ever existed. These findings are in good agreement with earlier studies of MAPs (Bovy et al. 2016b). We measure an average high $\\mathrm{[ \\alpha \\mathrm{\/Fe]}}$ population scalelength of hR, in = 1.9 \u00b1 0.1\u2009kpc, and find scaleheights between 600 and 1000\u2009pc, in good agreement with current measures of the $\\mathrm{[ \\alpha \\mathrm{\/Fe]}}$ rich disc scalelength and scaleheight (e.g. those outlined in Bland-Hawthorn & Gerhard 2016).Profile broadening: We show that the radial surface density profile of the low $\\mathrm{[ \\alpha \\mathrm{\/Fe]}}$ populations broadens with age in a given [Fe\/H] bin, which we interpret as evidence of the gradual dispersal of mono-[Fe\/H] populations, presumably due to radial migration and radial heating. The variation in shape of the broken exponential profile changes differently depending on the population [Fe\/H], with low [Fe\/H] populations inner profiles flattening faster, whereas the high [Fe\/H] outer profiles flatten faster. We interpret this effect as tentative evidence for [Fe\/H] dependent radial migration arising from pre-existing [Fe\/H] gradients in the star-forming disc. We showed that our results qualitatively reproduce those of Hayden et al. (2015), finding a skewed MDF that varies as a function of R.Flaring: We find that flaring seems to be present in almost all mono-age populations, at differing levels. We have shown that the inverse flaring scalelength Rflare\u2212 1 increases with age, meaning that the youngest populations flare most strongly. This finding appears inconsistent with that above, under the assumption that flaring is the result of radial migration. However, these results may be reconciled by invoking a more active accretion history in the early life of the disc, which could have suppressed flaring (e.g. Minchev et al. 2014b).The surface-mass density at R0: We have measured the surface mass density at the solar radius for each mono-age, mono-[Fe\/H] population, finding a total surface mass density of $\\Sigma _{R_0, {\\rm tot}} = 20.0_{-2.9}^{+2.4}\\mathrm{(stat.)}_{-2.4}^{+5.0}\\mathrm{(syst.)}\\ \\mathrm{M_{{\\odot }} \\ pc^{-2}}$. Before allowing for systematics, this value is less than current estimates (e.g. Flynn et al. 2006; Bovy et al. 2012a; McKee et al. 2015), however, the systematic uncertainties are large, mainly due to a mismatch between the log\u2009g scales in APOGEE and the PARSEC models, and as such, we find our value to be consistent within the uncertainties. The relative contribution of high to low $\\mathrm{[ \\alpha \\mathrm{\/Fe]}}$ populations, $f_\\Sigma$, is 18\u2009per\u2009cent \u00b1 5\u2009per\u2009cent, which is consistent with existing measurements (e.g. Bland-Hawthorn & Gerhard 2016).The hZ distribution at R0: The shape of the mass-weighted hZ distribution found by this study is in good agreement with that of Bovy et al. (2012a), calling into question the existence of a vertical structural discontinuity in the Milky Way disc. The reconciliation of this finding with the discontinuity in chemical space (e.g. the bimodality in $\\mathrm{[ \\alpha \\mathrm{\/Fe]}}$ at fixed [Fe\/H]: Nidever et al. 2014; Hayden et al. 2015) may shed new light on our understanding of the formation of the Galactic disc.The surface-mass density profile of the Milky Way: We have found the combined (from mono-age, mono-[Fe\/H] populations at low and high $\\mathrm{[ \\alpha \\mathrm{\/Fe]}}$) surface-mass density-weighted profiles of the Milky Way disc as a function of $\\mathrm{[ \\alpha \\mathrm{\/Fe]}}$, age and [Fe\/H], and found that the total surface density is also described by a broken exponential. We find that our results fail to determine the sign of the inner exponential to high significance out to \u223c10\u2009kpc, but detect a turnover to a declining exponential, at high significance, thereafter. We find evidence of a radial mean age and [Fe\/H] gradient driven by the changing dominant population as a function of radius. A detailed comparison of these findings with numerical simulations is necessary for a proper interpretation. Our finding of a decline in stellar density may be consistent with that found in other studies (e.g. Reyl\u00e9 et al. 2009; Sale et al. 2010), albeit at shorter radii.","Citation Text":["Bovy et al.","2016b"],"Functions Text":["We have performed the first detailed dissection of the stellar populations of the Milky Way disc in age, [Fe\/H] and $\\mathrm{[ \\alpha \\mathrm{\/Fe]}}$ space, bridging the gap between the detailed observational understanding of MAPs (e.g."],"Functions Label":["Extends"],"Citation Start End":[[237,248],[256,261]],"Functions Start End":[[0,236]]}
{"Identifier":"2021AandA...652A.117B__Rieder_&_Kenworthy_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":["Rieder & Kenworthy 2016"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[1357,1380]],"Functions Start End":[[1137,1355]]}
{"Identifier":"2020AandA...644A.159C__Titov_&_Lambert_2013_Instance_1","Paragraph":"In order to try to get insights into such systematics, we produced several variants of the ICRF3 S\/X band frame by changing the reference epoch of the catalog or alternately by not considering Galactic acceleration in the modeling. Interestingly, the D2 and D3 glide terms for these variants were found to vary by several tens of microarcseconds in the comparison to ICRF2, in line with the level of the systematics observed for those terms. Such findings are not unexpected since Galactic acceleration manifests itself as a dipolar deformation in the source coordinates (e.g., Titov & Lambert 2013). Moving further, and noting that Mignard et al. (2016) mentioned this phenomenon as a possibility for explaining the observed glide between ICRF2 and the Gaia Data Release 1 (Gaia DR1) auxiliary quasar solution, we decided to reproduce an equivalent of ICRF2 by considering only the stretch of data used for ICRF2 (i.e., including only the VLBI sessions up to March 2009 in the solution) and to make a variant that adds Galactic acceleration in the modeling, as implemented for ICRF3. To guarantee the maximum consistency, those two analyses were conducted by employing the same software package as that used for ICRF2, namely CALC-SOLVE (see Fey et al. 2015). Looking at the results, we first observed that our \u201creproduced\u201d ICRF2 shows similar deformations as the original ICRF2 when compared to ICRF3, hence ruling out the possibility that ICRF2 was in error. Most importantly, the ICRF2 variant that incorporates Galactic acceleration modeling was found to have much reduced glide terms compared to the original or reproduced ICRF2. This is illustrated by the bar chart in Fig. 16 which shows that the D2 term has now vanished while the D3 term has been cut by more than half (down to a value of \u221239\u2005\u00b1\u20054 \u03bcas) in this variant, hence indicating that the deformations between the two frames, in large part, stem from Galactic acceleration not accounted for in ICRF2. This is somehow not surprising since the data set for ICRF2 already covered 30 years, enough for Galactic acceleration effects to emerge, even though the accuracy of the frame was lower than that of ICRF3.","Citation Text":["Titov & Lambert 2013"],"Functions Text":["Interestingly, the D2 and D3 glide terms for these variants were found to vary by several tens of microarcseconds in the comparison to ICRF2, in line with the level of the systematics observed for those terms. Such findings are not unexpected since Galactic acceleration manifests itself as a dipolar deformation in the source coordinates (e.g.,"],"Functions Label":["Similarities"],"Citation Start End":[[578,598]],"Functions Start End":[[232,577]]}
{"Identifier":"2019ApJ...870...39V__Veron-Cetty_&_Veron_1986_Instance_1","Paragraph":"NGC 3557 is a bona fide southern-sky elliptical galaxy (E3) at a distance of 40 Mpc (1\u2033 is 198 pc at this distance7\n\n7\nFor distance derivations, we have assumed a cosmology model with H0 =73 km s\u22121 Mpc\u22121, \u03a9matter = 0.27, and \u03a9vacuum = 0.73.\n), and a member of a small group of galaxies (Brough et al. 2006). It has been classified as a LINER (e.g., Annibali et al. 2010) and as a flat-spectrum radio galaxy (Healey et al. 2007), with a jet that bends at distances of a few arcmin from the center (Schmitt et al. 2002). Detections of several of the ISM components of this object have been reported in the literature. Dust has been observed both as FIR emission (Pasquale et al. 2009) and as absorption against the central stellar continuum (Lauer et al. 2005). Nuclear optical line emission has been reported from spectroscopy (Veron-Cetty & Veron 1986; Rampazzo et al. 2005) as well as narrowband photometry (Goudfrooij et al. 1994). The MIR spectrum of NGC 3557 is of the most common Class-2 type, which is currently associated with a post-star-formation phase (Vega et al. 2009). Regarding the cold gas component, atomic gas (H i) has been reported as nondetection in several works (Serra & Osterloo 2010), while molecular gas emission has been detected in single-dish observations (Prandoni et al. 2010). NGC 3557 was included in our sample because of its relative proximity and CO(2\u20131) brightness (as reported by Prandoni et al. 2010), which suggested the possibility to study the molecular structures in detail. Since the object is considered to be in a stage where little current star formation may be happening, it was considered important as a representative of the molecular gas structures to be expected in post-star-formation scenarios. Furthermore, the disk-like nuclear dust distribution seen in the photometry by Lauer et al. (2005) indicated the possible presence of a molecular gas disk, which is a very interesting dynamical structure linked to the presence of organized angular momentum at those spatial scales.","Citation Text":["Veron-Cetty & Veron 1986"],"Functions Text":["Nuclear optical line emission has been reported from spectroscopy"],"Functions Label":["Background"],"Citation Start End":[[827,851]],"Functions Start End":[[760,825]]}
{"Identifier":"2020ApJ...889...15Y__Yang_et_al._2016b_Instance_1","Paragraph":"Although each one of the four aforementioned mechanisms has some observational support in certain systems, there is not a single mechanism that can explain all observed polarization in protoplanetary disks. Alignment with respect to the local radiation anisotropy (\u201ck-RAT alignment\u201d thereafter) is best supported by the azimuthal polarization pattern observed at ALMA Band 3 in the HL Tau system (Kataoka et al. 2017). However, it predicts a strong azimuthal variation of polarization and circular pattern (rather than elliptical pattern) (Yang et al. 2019). There is some tentative evidences for alignment with respect to the magnetic field, through either Radiative Alignment Torques (\u201cB-RAT alignment\u201d; Lazarian & Hoang 2007), or recently proposed Mechanical Alignment Torques (Hoang et al. 2018), in, e.g., the IRAS 4A system at cm wavelengths (Cox et al. 2015; Yang et al. 2016b) and BHB07-11 (Alves et al. 2018) at (sub)millimeter wavelengths. But there is no well-resolved system that matches the theoretical expectations (see, e.g., Cho & Lazarian 2007; Yang et al. 2016b; Bertrang et al. 2017) assuming the widely expected disk toroidal magnetic field yet (Flock et al. 2015). Mechanical alignment has recently received some attention. Hoang et al. (2018) claims that under MATs, grains can be aligned with respect to local dust-gas streaming direction, in the case of a weak or zero magnetic field, even if the velocity difference is sub-sonic. Within this picture, Kataoka et al. (2019) investigated the direction of streaming velocities for dust grains with different Stokes numbers, and the resulting polarization orientations. They found that their polarization pattern in the order-of-unity Stokes number case resembles that observed by Alves et al. (2018) in BHB07-11. The BHB07-11, however, is a binary system, and we expect more complicated velocity fields than the simple one assumed in Kataoka et al. (2019). Yang et al. (2019) investigated the observational features of another mechanical alignment mechanism, the Gold mechanism (Gold 1952), to address the circular versus elliptical pattern problem in the ALMA Band 3 polarization observations of HL Tau disk. However, they failed to explain the nonexistence of strong azimuthal variation, and suggested the scattering by dust grains aligned under the Gold mechanism may be the origin of the polarization at ALMA Band 3 in the HL Tau system.","Citation Text":["Yang et al. 2016b"],"Functions Text":["There is some tentative evidences for alignment with respect to the magnetic field","in, e.g., the IRAS 4A system at cm wavelengths"],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[866,883]],"Functions Start End":[[559,641],[801,847]]}
{"Identifier":"2018ApJ...860...24P__Warmuth_2015_Instance_1","Paragraph":"Figure 13 shows the temporal evolution of the density, \u03c1, plasma flow velocity, vx, position of the wave crest, PosA, phase speed, vw, and magnetic field component in the z-direction, Bz, for the primary waves in every different case of initial amplitude, \u03c1IA. In Figure 13(a), we observe that the amplitude of the density remains approximately constant at their initial values until the time when the shock is formed and the density amplitude of the primary wave starts decreasing (see Vr\u0161nak & Luli\u0107 2000), i.e., at t \u2248 0.03 (blue), t \u2248 0.04 (red), and t \u2248 0.055 (green). For the case of \u03c1IA = 1.3 (magenta), a decrease of the amplitude of the primary can hardly be observed, as expected for low-amplitude wave (Warmuth 2015). One can see that the larger the initial amplitude, \u03c1IA, the stronger the decrease of the primary wave\u2019s amplitude, which is consistent with observations (Warmuth & Mann 2011; Muhr et al. 2014; Warmuth 2015). The amplitudes decrease to values of \u03c1 \u2248 1.6 (blue), \u03c1 \u2248 1.5 (red), and \u03c1 \u2248 1.4 (green) until the primary wave starts entering the CH. Due to the fact that the waves with larger initial amplitude enter the CH earlier than those with small initial amplitude, we can see in Figure 13(a) that the tracking of the parameters of the faster waves stops at an earlier time than the one for the slower waves. A similar behavior to the one of the density, \u03c1, can be observed for the plasma flow velocity, vx, in Figure 13(b) and the magnetic field component, Bz, in Figure 13(e). Here, the amplitudes decrease from vx = 0.75, Bz = 1.9 (for \u03c1IA = 1.9, blue), vx = 0.6, Bz = 1.7 (for \u03c1IA = 1.7, red), vx = 0.45, Bz = 1.5 (for \u03c1IA = 1.5, green), and vx = 0.27, Bz = 1.3 (for \u03c1IA = 1.3, magenta) to vx = 0.55, Bz = 1.6 (for \u03c1IA = 1.9, blue), vx = 0.46, Bz = 1.5 (for \u03c1IA = 1.7, red), vx = 0.36, Bz = 1.4 (for \u03c1IA = 1.5, green), and vx = 0.25, Bz = 1.25 (for \u03c1IA = 1.3, magenta). Figure 13(c) shows how the primary waves propagate in the positive x-direction. In all four cases of different initial amplitude, \u03c1IA, the phase speed decreases slighty (consistent with observations; see Warmuth et al. 2004 and Warmuth 2015) until the waves enter the CH at different times, i.e., the values for the phase speed start at vw \u2248 2.2 (for \u03c1IA = 1.9, blue), vw \u2248 1.9 (for \u03c1IA = 1.7, red), vw \u2248 1.7 (for \u03c1IA = 1.5, green), and vw \u2248 1.4 (for \u03c1IA = 1.3, magenta) and decrease to vw \u2248 1.5 (for \u03c1IA = 1.9, blue), vw \u2248 1.39 (for \u03c1IA = 1.7, red), vw \u2248 1.2 (for \u03c1IA = 1.5, green), and vw \u2248 1.13 (for \u03c1IA = 1.3, magenta).","Citation Text":["Warmuth 2015"],"Functions Text":["For the case of \u03c1IA = 1.3 (magenta), a decrease of the amplitude of the primary can hardly be observed, as expected for low-amplitude wave"],"Functions Label":["Similarities"],"Citation Start End":[[714,726]],"Functions Start End":[[574,712]]}
{"Identifier":"2019AandA...623A..11P__Simm_et_al._2016_Instance_1","Paragraph":"Another clear difference for Ark 120 with the NT thin disc predictions is the variability. From the best simultaneous fit of the 2013 and 2014 observations, we infer a significant increase in mass accretion rate through the disc from 0.03 to 0.07 Lbol\/LEdd in only one year, but a standard thin disc around a SMBH cannot vary on such timescale. Indeed, the radial mass accretion rate change via viscous processes has a time-scale of torb(R)\/[\u03b1(H\/R)2]. For Ark,120, assuming an orbital time-scale at R\u2004=\u2004100 Rg, a viscosity parameter of 0.1 and a H\/R (H is the height scale of the disc) value of 0.1, this corresponds to \u223c150 years. Moreover, the optical-UV flux significantly changes in less than a year, varying for example by a factor of 50% in the UVW2 band between 2013 and 2014 (e.g. see Fig. 1, Lobban et al. 2018). Similar rapid changes in the optical-UV flux are typically seen in other BLS1s, especially those at low Lbol\/LEdd (e.g. MacLeod et al. 2010; Koz\u0142owski 2016; Simm et al. 2016; Rakshit & Stalin 2017). These are generally assumed to be from reprocessing, where the X-ray flux illuminates the outer disc (e.g. Buisson et al. 2017), and adds to the intrinsic emission. An additional reprocessed component in the optical-UV would lead us to overestimate the value of the intrinsic \u1e40, so to an underestimate of BH spin via the efficiency argument (Kubota & Done 2018). Nonetheless, changes as large as about 50% in the UV flux are unlikely to be driven by X-ray reprocessing, as the UV flux in Ark 120 is much higher than in the X-ray band. Besides, detailed models of the expected optical-UV variability from X-ray reprocessing fail to fit the excellent long term simultaneous optical-UV-X-ray datasets, and would imply a larger disc size than expected by standard thin disc (e.g. NGC 5548: McHardy et al. 2014; Edelson et al. 2015; Gardner & Done 2017; NGC 4151: Edelson et al. 2017; Ark 120: Gliozzi et al. 2017; Fairall 9: Pal et al. 2017; NGC 4593: Cackett et al. 2018; Pal & Naik 2018; Microlensing studies: Morgan et al. 2010; Dai et al. 2010). Such large discs should be significantly brighter than observed and this discrepancy may be explained for example by a flatter temperature profile than in NT discs, from scattering of a significant part of the optical flux on larger scales, by electron scattering in the disc atmosphere (Dai et al. 2010; Morgan et al. 2010; Hall et al. 2018).","Citation Text":["Simm et al. 2016"],"Functions Text":["Similar rapid changes in the optical-UV flux are typically seen in other BLS1s, especially those at low Lbol\/LEdd (e.g."],"Functions Label":["Similarities"],"Citation Start End":[[979,995]],"Functions Start End":[[822,941]]}
{"Identifier":"2021ApJ...914L..19Z__Metzger_et_al._2008_Instance_1","Paragraph":"Thermonuclear explosions and AICs of WDs in AGN disks are potential sources for future joint gravitational wave (GW), EM, and neutrino multi-messenger observations. In addition to AIC of WDs, another possible formation channel of millisecond magnetars in AGN disks is mergers of BNSs. Such a merger may produce a millisecond magnetar if the NS equation of state is stiff enough (e.g., Ai et al. 2020). This would also drive a magnetar-powered explosion as discussed above. The mechanism by which a nascent millisecond magnetar might power a GRB jet has been studied in detail over the past decade (Bucciantini et al. 2008, 2012; Metzger et al. 2008, 2011; Siegel et al. 2014).8\n\n8\nSome indirect evidence for the magnetar mechanism in GRBs are the presence of a plateau in the early X-ray afterglow of long-duration GRBs, which may be explained as the continuous energy injection from the spindown of a magnetar (Dai & Lu 1998; Zhang & M\u00e9sz\u00e1ros 2001; Zhang et al. 2006).\n A putative jet may thus be launched from the AIC of WDs and BNS mergers occurring in AGN disks. However, Zhu et al. (2021b, 2021a) and Perna et al. (2021a) recently showed that GRB jets embedded in AGN disks would be usually choked by the dense material of the accretion disks. Although it is hard to observe gamma-ray signals, Zhu et al. (2021a) suggested that these choked GRBs may effectively produce TeVPeV neutrinos that could be detected by IceCube and IceCube-Gen2. Furthermore, the dissipation of magnetic energy during the merger of BWDs is expected to accelerate cosmic rays and produce high-energy neutrinos (Xiao et al. 2016). In the GW channel, BWD mergers are promising astrophysical GW sources for space-borne GW observatories, e.g., LISA (Amaro-Seoane et al. 2017), TaiJi (Ruan et al. 2020), and TianQin (Luo et al. 2016), while the GW signals from BNS mergers are readily detected with the ground-based GW detectors such as LIGO and Virgo (Abbott et al. 2017). Future joint observations of GW, EM, and neutrino signals can reveal the existence of WD and NS populations in AGN disks.","Citation Text":["Metzger et al. 2008"],"Functions Text":["The mechanism by which a nascent millisecond magnetar might power a GRB jet has been studied in detail over the past decade"],"Functions Label":["Background"],"Citation Start End":[[629,648]],"Functions Start End":[[473,596]]}
{"Identifier":"2022AandA...668A..10F__Walker_et_al._2008_Instance_1","Paragraph":"The most common approaches to observe SPI are spectral observations of typical emission lines for stellar chromospheres and coronae. The earliest approach to observe SPI dates back to the early 2000\u2019s and focused on non-thermal chromospheric Ca II H and K emissions (Shkolnik et al. 2003). Subsequent studies by Shkolnik et al. (2005, 2008) of the system HD 179949 saw enhanced emissions in four out of six observational epochs that appeared to be periodic with the planetary orbital period. Other authors observed different systems and likewise saw chro-mospheric excess emissions (Walker et al. 2008; Staab et al. 2017). In addition, Cauley et al. (2019) estimated possible magnetic field strengths of hot Jupiters based on chromospheric excess emissions. The derived energy fluxes from chromospheric Ca II H and K emissions are larger than typical fluxes derived from Alfv\u00e9n wing models and magnetohydrostatic models (Saur et al. 2013; Lanza 2015) and the observed periodicity only appears in some epochs, which was suggested to be alternatively explained by star spots (Miller et al. 2012). Recently, Strugarek et al. (2019) investigated how modeling and observations could reconcile by including the magnetic topology of stellar coronae based on the example of Kepler-78. In the UV range, France et al. (2016, 2018) conducted large surveys with the Hubble Space Telescope, however, the authors could not identify any signals related to SPI. At coronal X-ray wavelengths, several studies have investigated the influence of planets on stellar X-ray activity, such as Kashyap et al. (2008); Scharf (2010); Poppenhaeger et al. (2010). Some studies found correlations between planets and X-ray activity in stars and some did not. Poppenhaeger & Schmitt (2011a,b) explain previously observed correlations, attributed to SPI, as a result of selection effects due to planet detection methods and the limitations in X-ray observations. The radio wavelength range is particularly promising to search for SPI. Recent observations with LOFAR showed strong hints for the existence of SPI in several systems (Vedantham et al. 2020; Callingham et al. 2021; P\u00e9rez-Torres et al. 2021). In addition, a recent campaign in the radio range provided the first tentative direct observational hints of an intrinsic magnetic field on an exoplanet (Turner et al. 2021).","Citation Text":["Walker et al. 2008"],"Functions Text":["Other authors observed different systems and likewise saw chro-mospheric excess emissions"],"Functions Label":["Background"],"Citation Start End":[[583,601]],"Functions Start End":[[492,581]]}
{"Identifier":"2022MNRAS.514.2974M__Dadina_2007_Instance_1","Paragraph":"Active galactic nuclei (AGNs) are extragalactic sources that emit across the whole electromagnetic spectrum. Such systems are composite and each sub-structure has its own role in shaping the emerging spectrum (see Padovani et al. 2017, for a comprehensive review). It is ubiquitously accepted that the X-ray emission originates in the very inner regions of AGNs, near the central supermassive black hole (SMBH). Accretion of matter infalling on to the SMBH is responsible for the enormous amount of optical-UV photons, a fraction of which can be further energized via inverse-Compton (Sunyaev & Titarchuk 1980) off thermal electrons (the so-called hot corona: Haardt & Maraschi 1991, 1993; Madejski et al. 1995; Zdziarski et al. 1995) up to the X-rays. The maximum energy gain for these seed photons is mainly set by the hot plasma\u2019s temperature, and, to a lower extent, by its opacity (e.g. Rybicki & Lightman 1979; Beloborodov 1999; Middei et al. 2019). In fact, the X-ray continuum in AGNs is well modelled by a power law with a high energy roll-over (e.g. Perola et al. 2002; Dadina 2007; Molina et al. 2009, 2013; Malizia et al. 2014; Fabian et al. 2015, 2017; Ricci et al. 2018; Tortosa et al. 2018). AGN X-ray spectra may show additional features due to reprocessing of the primary X-ray emission by the circumnuclear material. A fluorescence emission line from the Fe K-shell is commonly observed as the most prominent feature (e.g. Bianchi et al. 2009) and its analysis carries a wealth of information on the physics of the reflecting material. This emission line has an intrinsically narrow profile that can undergo distortions, such as broadening, due to special and general relativistic effects. In particular, the closer to the SMBH the reflectors, the more distorted (i.e. the broader) the neutral or ionized Fe line profile (e.g. Fabian et al. 1995). On the contrary, at larger distance, these effects are negligible, thus the Fe K\u03b1 shape is consistent with a narrow profile. Additionally, in the case of Compton-thick reflectors (i.e. NH \u2273 1.5 \u00d7 1024 cm\u22122) the X-ray spectra show a typical emission excess around 30 keV, the so-called Compton-hump (e.g. Matt, Fabian & Ross 1993). The effect of any absorbing matter crossing our line of sight can significantly attenuate the observed number of photons, especially in the soft X-rays (e.g. Cappi et al. 1999; Awaki et al. 2000; Matt 2002; Bianchi et al. 2009; Middei et al. 2021).","Citation Text":["Dadina 2007"],"Functions Text":["In fact, the X-ray continuum in AGNs is well modelled by a power law with a high energy roll-over (e.g."],"Functions Label":["Background"],"Citation Start End":[[1080,1091]],"Functions Start End":[[956,1059]]}
{"Identifier":"2017AandA...608A..16R__Ravindranath_et_al._2006_Instance_1","Paragraph":"The measurement of clumps at redshifts beyond the peak of cosmic star formation is notoriously difficult. The identification of clumpy galaxies in the dominant population of irregular galaxies at high redshift dates back to the first deep Hubble Space Telescope (HST) images (e.g., Williams et al. 1996). Subsequent studies at z ~ 1\u22123 revealed that clumpy galaxies are more numerous than in the local Universe (e.g., Abraham et al. 1996; van den Bergh et al. 1996; Giavalisco et al. 1996; Elmegreen et al. 2004, 2007, 2008, 2009; Elmegreen & Elmegreen 2006; Elmegreen et al. 2013; Kubo et al. 2013, 2016; Glazebrook 2013; Tadaki et al. 2014; Murata et al. 2014; Guo et al. 2015; Garland et al. 2015; Bournaud 2016). The fraction of clumpy galaxies was studied from z = 0 up to the most distant galaxies identified today at z ~ 10 (Ravindranath et al. 2006; Guo et al. 2012, 2015; Shibuya et al. 2016). In the sample of galaxies with 0.5 z  3 in the CANDELS survey, Guo et al. (2015) find that galaxies with low mass log\u2009(M\u22c6\/M\u2299)  9.8 have a high fraction of off-center clumps fclumpy ~ 60% constant over the observed redshift range, while for intermediate and massive galaxies fclumpy decreases from ~40% at z ~ 3 to ~15% at z ~ 0.5. Combining deep HST imaging from the 3D-HST, CANDELS, HUDF and HFF surveys Shibuya et al. (2016) claim that \\hbox{$f_{\\rm clumpy}^{\\rm UV}$}fclumpyUV follows an evolution similar to that of the star formation rate density (e.g., Madau & Dickinson 2014), increasing from z \u2243 8 to z \u2243 1\u22123 and decreasing from z \u2243 1 to z \u2243 0. In comparing these observed trends with the predictions of simulations, Guo et al. (2015) conclude that VDI are likely responsible for fclumpy at high mass and that minor mergers are a viable explanation for fclumpy at intermediate mass for z  1.5, while Shibuya et al. (2016) conclude that VDI is the main origin of all clumps. Both these studies exclude major mergers as a possible contribution to the evolution of the fraction of clumps. This is somewhat surprising as the major merger fraction for star-forming galaxies is observed to increase up to ~20% at z \u2243 1\u22122 (e.g., Lotz et al. 2011; L\u00f3pez-Sanjuan et al. 2011, 2013), staying high at least to z ~ 3\u22124 (e.g., Conselice et al. 2008; Tasca et al. 2014a), and one would therefore expect that a fraction of the clumps is related to ex situ major merging. ","Citation Text":["Ravindranath et al. 2006"],"Functions Text":["The fraction of clumpy galaxies was studied from z = 0 up to the most distant galaxies identified today at z ~ 10"],"Functions Label":["Background"],"Citation Start End":[[831,855]],"Functions Start End":[[716,829]]}
{"Identifier":"2020MNRAS.499.1788W__Wolfire_et_al._2003_Instance_2","Paragraph":"Owing to their brightness at rest-frame FIR wavelengths, the ionized and neutral species of Carbon, Nitrogen, and Oxygen are powerful diagnostic lines for tracing the ISM of nearby and distant galaxies. When combined with photodissociation region (PDR) models (Tielens & Hollenbach 1985), measurement of the emission from different lines provide a means to constrain quantities such as the ionization rate and metallicity of the ISM of galaxies. Multilevel FIR transition lines have now been widely surveyed in local galaxies, originally by ISO (e.g. Luhman et al. 1998; Malhotra et al. 2001) and more recently using the PACS spectrometer (Poglitsch et al. 2010) on the ESA Herschel Space Observatory (Pilbratt et al. 2010). The [C\u2009ii]158\u2009\u03bcm line is typically the brightest in star-forming galaxies, arising from ionized, and even neutral gas where it is the main coolant (Wolfire et al. 2003). Another commonly observed FIR line tracer of the ionized ISM is [N\u2009ii]122\u2009\u03bcm. It has the advantage that it can be found associated with lower excitation gas, close to that observed in our own Galaxy (e.g. Goldsmith et al. 2015; Herrera-Camus et al. 2016). Another major coolant of the ISM is [O\u2009i]63\u2009\u03bcm (Wolfire et al. 2003). Owing to its high excitation temperature and critical density, it can dominate the cooling in regions of starburst activity (Kaufman et al. 1999; Kaufman, Wolfire & Hollenbach 2006; Brauher, Dale & Helou 2008; Croxall et al. 2012; Narayanan & Krumholz 2017). When combined with measurements of the [C\u2009ii]158\u2009\u03bcm line intensity and FIR luminosity, the [O\u2009i]63\u2009\u03bcm line intensity can constrain the FUV field, G, and the gas density using PDR models. The luminosity in these FIR lines generally exhibits a deficit in the most FIR luminous galaxies compared to the trend expected from lower luminosity galaxies (e.g. Malhotra et al. 2001; Graci\u00e1-Carpio et al. 2011; D\u00edaz-Santos et al. 2017). This has made the emission from lines like [O\u2009i]63\u2009\u03bcm more challenging to detect at high-redshifts.","Citation Text":["Wolfire et al. 2003"],"Functions Text":["Another major coolant of the ISM is [O\u2009i]63\u2009\u03bcm"],"Functions Label":["Background"],"Citation Start End":[[1199,1218]],"Functions Start End":[[1151,1197]]}
{"Identifier":"2021MNRAS.503..324M__Zhao_et_al._2019_Instance_1","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)"],"Functions Text":["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"],"Functions Label":["Similarities"],"Citation Start End":[[580,598]],"Functions Start End":[[419,579]]}
{"Identifier":"2020MNRAS.494.2851C__Dijkstra_et_al._2008_Instance_1","Paragraph":"Previously it has been postulated that the SMS formation by DC is only possible in atomic cooling haloes both with intense FUV irradiation and with primordial gas composition. Here we have shown that SMSs can form also in slightly metal-enriched cases as long as FUV irradiation is intense enough. This relaxation of the condition will increase the expected number density of massive seed BHs. Several authors have estimated the number density of such seeds formed in the usual DC scenario, which requires metal-free gas composition. Using the critical intensity advocated by recent studies (Agarwal et al. 2014; Sugimura et al. 2014; Inayoshi & Tanaka 2015), this ranges from a few Gpc\u22123 (e.g. Dijkstra et al. 2008; Dijkstra, Ferrara & Mesinger 2014), to 10\u22126 to 10\u22124 Mpc\u22123 (e.g. Agarwal et al. 2012; Visbal, Haiman & Bryan 2014a; Chon et al. 2016; Habouzit et al. 2016; Valiante et al. 2016). The dynamical heating model by Wise et al. (2019) predicts the number density of 10\u22127 to 10\u22126 Mpc\u22123. Those seeds can account for the origin of the rare BHs in the high-z universe. If all the SMBHs down to z = 0 shares the universal origin, the seed density should be comparable to that of the SMBHs ubiquitously residing in massive galaxies, \u223c0.01\u20130.1 Mpc\u22123 (Aller & Richstone 2002; Davis et al. 2014), which is much larger than the seed density predicted by the DC model. The DC in the primordial environment is terminated around z \u223c 10, as the metal enrichment proceeds (e.g. Trenti & Stiavelli 2009; Chon et al. 2016): once a Pop III star ends its life as a core-collapse SN, metallicity inside the host halo jumps up to 10\u22124 to 10\u22123 Z\u2299 (e.g. Maio et al. 2010; Ritter et al. 2015; Sluder et al. 2016; Chiaki, Susa & Hirano 2018). According to our result, the SMS formation can continue even in such second-generation haloes. Since the second-generation haloes tend to be under stronger radiation field than the pristine ones (e.g. Dijkstra et al. 2014), we expect significant increase in the number density of the SMSs. For example, the first episode of star formation delays another star formation for a few hundred million years by the radiative and SNe mechanical feedback, which ejects a gas from the halo. If the haloes approach close enough to a luminous galaxy before another episode of star formation, they can be ideal sites for SMS formation. Not only the radiation sources outside the halo, but also those inside the same halo can trigger SMS formation in an irradiated massive cloud as long as the metallicity is low enough. With such new varieties of SMS formation sites, the expected seed BH number will be largely enhanced. We will pursue the validity of this scenario using samples from the cosmological simulations in the future studies.","Citation Text":["Dijkstra et al. 2008"],"Functions Text":["Using the critical intensity advocated by recent studies","this ranges from a few Gpc\u22123 (e.g.","to 10\u22126 to 10\u22124 Mpc\u22123"],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[695,715]],"Functions Start End":[[534,590],[660,694],[753,774]]}
{"Identifier":"2020MNRAS.492.3509B__Garsden_et_al._2015_Instance_1","Paragraph":"In the context of RI imaging techniques, the standard clean algorithm implements a non-linear, greedy approach to perform deconvolution in an iterative manner (H\u00f6gbom 1974). Working pixelwise, clean implicitly assumes sparsity of the sought image, removing at each iteration, a fraction of the maximum intensity pixel convolved with the dirty beam from the computed residual image. Over the past years, many variants of this celebrated algorithm have been developed, for instance, its multiscale version (Cornwell 2008). In addition to these clean-based approaches, recent developments in the field of compressive sensing (CS) applied for astronomical imaging have given birth to many imaging algorithms, particularly for RI (Wiaux et al. 2009; Carrillo, McEwen & Wiaux 2012; Garsden et al. 2015; Onose et al. 2016). Leveraging optimization frameworks, these methods aim to solve the underlying image recovery problem by enforcing sparsity of the image of interest in some suitable domain. These techniques have shown the potential to surpass the image reconstruction quality obtained by clean -based approaches (Carrillo, McEwen & Wiaux 2014; Onose et al. 2016; Pratley & Johnston-Hollitt 2016; Onose, Dabbech & Wiaux 2017; Dabbech et al. 2018). While the aforementioned methods have been developed mainly for Stokes I imaging, as previously mentioned these can be extended to polarimetric imaging by following the same imaging approach for all the Stokes parameters. In the context of sparsity regularized approaches, one such extension for polarimetric imaging has been presented in Akiyama et al. (2017), promoting sparsity of each of the Stokes parameters using \u21131 norm combined with total variation (TV) regularization (Rudin, Osher & Fatemi 1992). A point worth noting here is that the above-mentioned approaches, whether clean and its variants or the sparsity regularized method, solve for the Stokes parameters totally independently. However, these images are physically linked via polarization constraint, i.e. the polarized intensity cannot be greater than the total intensity. Within the CS framework, this constraint has been exploited by a recently proposed approach, namely Polarized SARA, estimating jointly the Stokes parameters (Birdi, Repetti & Wiaux 2018a, b). This approach has been shown to provide better reconstruction quality in comparison with the case when the constraint is not accounted for.","Citation Text":["Garsden et al. 2015"],"Functions Text":["In addition to these clean-based approaches, recent developments in the field of compressive sensing (CS) applied for astronomical imaging have given birth to many imaging algorithms, particularly for RI","Leveraging optimization frameworks, these methods aim to solve the underlying image recovery problem by enforcing sparsity of the image of interest in some suitable domain."],"Functions Label":["Background","Background"],"Citation Start End":[[776,795]],"Functions Start End":[[521,724],[817,989]]}
{"Identifier":"2022MNRAS.517.5473G__Kapanadze_et_al._2018_Instance_1","Paragraph":"Various outbursts in GeV\u2013TeV energies from 1ES 1959+650 have been reported by Fermi-LAT and several ground-based Cherenkov experiments like VERITAS, MAGIC, FACT (Albert et al. 2006; Aliu et al. 2014; Biland et al. 2016) including a new highest historical \u03b3-ray activity, detected by MAGIC collaboration on 2016 June 13\u201314, when the VHE flux above 100 GeV reached up to 3 Crab unit (C.U) (Bhattacharyya et al. 2019). Many multiwavelength campaign from radio to TeV \u03b3-ray energies have been conducted since its discovery (Krawczynski et al. 2004; Biland et al. 2016) and this source has exhibited a number of outbursts across all energies. One interesting VHE outburst was detected on 2002 June 4 by the Whipple telescope without any activity in contemporaneous X-ray energies, called \u2018orphan flare\u2019, during a multiwavelength campaign (Krawczynski et al. 2004). In the recent past, 1ES 1959+650 has undergone a few major outbursts in all wavelengths specially in X-ray energies since 2015 (Kapanadze et al. 2016). Unprecedented flux and spectral variability during these episodes have been observed (Kapanadze 2015; Kaur et al. 2017; Kapanadze et al. 2018; Patel et al. 2018). Such flux and spectral variations in different energy bands over time provide an opportunity to study underlying physical processes in terms of particle acceleration processes, evolution of the accelerated particles (leptons\/hadrons) within the source, and their radiative cooling and escape processes. Many authors explained the broad-band emission from this source with leptonic one-zone SSC process (Krawczynski et al. 2004; Tagliaferri et al. 2008). But the lack of correlation between X-ray and VHE \u03b3-rays several times and also during the orphan flare, mentioned earlier, challenges the leptonic one-zone SSC model of emission. Alternatively, EC, lepto-hadronic and two-zone SSC models have been also tried to explain the broad-band SED of the source (Backes et al. 2012; Aliu et al. 2014; Asano & Hayashida 2018; Patel et al. 2018). Recently, the MAGIC collaboration again introspected the VHE flaring episodes in 2016, mentioned earlier, to study the feasibility of neutrino emission from 1ES 1959+650 as well under the framework of mixed lepto-hadronic model (MAGIC Collaboration 2020a). Although they (MAGIC collaboration) have concluded that it is difficult to detect neutrinos from the same source during extreme \u03b3-ray flares, it is customary to mention that more study with well-sampled data from multiple experiments are inevitable to understand the true physical processes inside this type of astrophysical sources in a better way.","Citation Text":["Kapanadze et al. 2018"],"Functions Text":["Unprecedented flux and spectral variability during these episodes have been observed"],"Functions Label":["Background"],"Citation Start End":[[1132,1153]],"Functions Start End":[[1012,1096]]}
{"Identifier":"2019MNRAS.488.5671I__Inoue_&_Iwata_2008_Instance_1","Paragraph":"While both Suprime-Cam\/NB359 and HST\/F336W trace LyC for sources at z > 3.06, NB359 has a narrower filter bandpass, and it preferentially captures ionizing photons closer to the Lyman limit (see Fig. 1). Because IGM attenuation varies as a function of wavelength, expected transmissions of photons through the IGM (here we denote the ratio of flux density with and without IGM attenuation as \u2018IGM transmission\u2019) are different for different filters. In Fig. 2 we show the median and 68 percentile IGM transmission values of NB359\u2009andF336W for a source at a redshift range of 3.1 \u2264 z \u2264 3.8. These IGM transmissions are estimated from Monte Carlo simulations (Inoue & Iwata 2008) which generate 10\u2009000 sightlines for redshifts consistent with the H\u2009i cloud distribution defined analytically by Inoue et al. (2014). For each sightline we calculate IGM transmission through NB359 or F336W for an object with a flat SED in f\u03bd (i.e. f\u03bd = constant) and obtain median and 68 percentile values from the 10\u2009000 instances. These calculations are repeated from z = 3.1 to 3.8 with a redshift step of 0.1. The median values of IGM transmission for NB359 are higher than those for F336W, especially at the lower redshift range. This is because NB359\u2009 traces a rest-frame wavelength range close to the Lyman limit where IGM transmission is higher than that for photons with shorter wavelengths. If the same limiting magnitude is achieved, NB359 is more sensitive for detecting LyC photons from z > 3 galaxies than F336W. Median values of IGM transmission for both NB359 and F336W rapidly decrease as the source redshift increases, which indicates that the detection of LyC photons becomes increasingly difficult due to increasing IGM opacity. However, higher-side 68 percentile IGM transmission values at z = 3.5 are 0.33 for NB359 and 0.20 for F336W, respectively. The fluctuation of IGM transmission among different sightlines is large, and there are still a reasonable number of sightlines with significant IGM transparency for LyC in the redshift range studied in this paper.","Citation Text":["Inoue & Iwata 2008"],"Functions Text":["These IGM transmissions are estimated from Monte Carlo simulations (",") which generate 10\u2009000 sightlines for redshifts consistent with the H\u2009i cloud distribution defined analytically by Inoue et al. (2014)."],"Functions Label":["Uses","Uses"],"Citation Start End":[[657,675]],"Functions Start End":[[589,657],[675,811]]}
{"Identifier":"2018MNRAS.477.4308R__Novak,_Ostriker_&_Ciotti_2011_Instance_1","Paragraph":"It is well known that the accretion on to compact objects may influence the nearby ambient around SMBHs in the centre of galaxies (e.g. Salpeter 1964; Fabian 1999; Barai 2008; Germain, Barai & Martel 2009). Together with the outflow phenomena, it is believed to play a major role in the feedback processes invoked by modern cosmological models (i.e. \u039b-Cold Dark Matter) to explain the possible relationship between the SMBH and its host galaxy (e.g. Magorrian et al. 1998; Gebhardt et al. 2000) as well as in the self-regulating growth of the SMBH. The problem of accretion on to an SMBH has been studied via hydrodynamical simulations (e.g. Ciotti & Ostriker 2001; Di Matteo, Springel & Hernquist 2005; Li et al. 2007; Ostriker et al. 2010; Novak, Ostriker & Ciotti 2011). In numerical studies of galaxy formation, spatial resolution permits resolving scales from the kpc to the pc, while the sub-parsec scales of the Bondi radius are not resolved. This is why a prescribed sub-grid physics is employed to solve this lack of resolution. With sufficiently high X-ray luminosities, the falling material will have the correct opacity, developing outflows that originate at sub-parsec scales. Therefore, calculations of the processes involving accretion on to SMBH have become of primary importance (e.g. Proga, Stone & Kallman 2000; Proga 2000, 2003; Proga & Kallman 2004; Proga 2007; Ostriker et al. 2010). Numerical calculations of the accretion of matter on to SMBHs, including the radiative-outflow component, have been mostly performed using Eulerian finite-difference methods (Mo\u015bcibrodzka & Proga 2013) [see also the overviews by Edgar (2004) and Foglizzo et al. (2005) and references therein for earlier work], while only a few calculations have been reported using smoothed particle hydrodynamics (SPH) techniques (Barai 2008; Barai, Proga & Nagamine 2011, 2012; Nagamine, Barai & Proga 2012), where results for the accretion rates, outflow rates, thermal instabilities, and impact of the thermal, mechanical, and X-ray feedbacks have been obtained for evolutions up to 5\u2009Myr and scales from 0.1 to 200\u2009pc.","Citation Text":["Novak, Ostriker & Ciotti 2011"],"Functions Text":["The problem of accretion on to an SMBH has been studied via hydrodynamical simulations (e.g."],"Functions Label":["Background"],"Citation Start End":[[742,771]],"Functions Start End":[[549,641]]}
{"Identifier":"2019MNRAS.485.5453S__Morlino_et_al._2012_Instance_1","Paragraph":"As well as being directly excited from the ground state, Balmer lines (intensity, line profile, polarization and so on) are affected by the conversion of Lyman lines to Balmer lines. For example, the absorption of Ly\u2009\u03b2 by a hydrogen atom results in radiative excitation from 1s to 3p, and the excited atom can emit H\u2009\u03b1 by the spontaneous transition from 3p to 2s. Simultaneously, the conversion yields the 2s-state hydrogen atom, which creates the two-photon continuum by the spontaneous transition from 2s to 1s. Thus, the Ly\u2009\u03b2 to H\u2009\u03b1 conversion impacts the total intensity, line profile, and net polarization of H\u2009\u03b1. Moreover, an adequate density of 2s-sate atoms can further scatter H\u2009\u03b1 photons. Although such fundamental physics is well known, it has not been well studied in SNR shocks. In fact, it is usually assumed that the Ly\u2009\u03b2 photons are at the limits of either completely optically thick or optically thin at SNR shocks, i.e. they are completely converted to H\u2009\u03b1 photons or not at all (e.g. Heng & McCray 2007; van Adelsberg et al. 2008; Morlino et al. 2012, 2013b; Shimoda et al. 2018). Contrary to this, Ghavamian et al. (2001) studied the conversion of Ly\u2009\u03b2 and Ly\u2009\u03b3 to H\u2009\u03b1 and H\u2009\u03b2 by Monte Carlo simulations and claimed that intermediate conversion occurs. However, they and previous studies did not consider the population of 2s-state hydrogen atoms. In this paper, we provide a formulation of the radiative line transfer with the rate equation of atomic population and study the nature of Balmer line emissions from SNR shocks. In this paper we do not consider the polarization, deferring that instead to a later work. Note that as a first step, our model makes several simplifications in the treatment of the SNR shock, handling the hydrogen atoms as fluids and supposing no particles leaking back upstream (e.g. cosmic rays). Our calculation of radiative transfer is based on so-called the ray-tracing method and uses updated atomic data from the literature (e.g. Heng & Sunyaev 2008; Tseliakhovich, Hirata & Heng 2012). Moreover, we consider only hydrogen line emissions and ignore bremsstrahlung radiation, thermal emissions from the SNR ejecta, and external radiation sources. Thus, our model possibly predicts somewhat smaller population of 2s-state hydrogen atoms than would be the case in a realistic SNR shock.","Citation Text":["Morlino et al. 2012"],"Functions Text":["Although such fundamental physics is well known, it has not been well studied in SNR shocks. In fact, it is usually assumed that the Ly\u2009\u03b2 photons are at the limits of either completely optically thick or optically thin at SNR shocks, i.e. they are completely converted to H\u2009\u03b1 photons or not at all (e.g."],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1050,1069]],"Functions Start End":[[699,1002]]}
{"Identifier":"2020ApJ...899..120T__Chou_et_al._2008_Instance_1","Paragraph":"The energy dependent pulse arrival time can usually be observed in AMXPs. Soft lags are most often seen in AMXPs, where the pulse arrival times of the pulsations from softer energy bands lag compared to the ones from the harder energy bands. This phenomenon was first discovered by Cui et al. (1998) during the 1998 outburst of the SAX J1808.4\u20133658. The soft lag tendency can clearly be detected from 2 keV extended up to 10 keV (hereafter, the break point) for about 200 \u03bcs (\u223c0.08 cycles) but saturated for harder X-ray bands. Similar energy dependent pulse arrival time behavior can also be found in XTE J1751\u2013305 (Gierli\u0144ski & Poutanen 2005), XTE J1814\u2013338 (Watts & Strohmayer 2006), HETE J1900.1\u20132455 (Galloway et al. 2007), and XTE J1807\u2013294 (Chou et al. 2008) with different break point energies. However, the energy dependent pulse arrival time acts differently in IGR J00291+5934. Falanga et al. (2005) found that the soft lags can extend to the break point energy at \u223c6 keV but the tendency reverses instead of saturation for the pulse of energy bands harder than the break point (hereafter hard lags) during its 2005 outburst. Similar behaviors are also observed during its 2015 outburst (Sanna et al. 2017) with a different break point energy at \u223c8 keV. Such hard lags are also marginally detected in IGR J17511\u20133057 (Falanga et al. 2011) and IGR J17498\u20132921 (Falanga et al. 2012). On the other hand, the energy dependent pulse arrival times for two AMXPs are rather peculiar. The pulse arrival times seem independent of energy bands below \u223c17 keV for SAX J1748.9\u20132021 (Patruno et al. 2009a) and only hard lags are detected in IGR J18245\u20132452 (De Falco et al. 2017). The newly discovered AMXP IGR J17379-3747 also shows unusual energy dependent pulse arrival times. Regular soft lags can be observed between 0.5 and 6 keV with no break point being detected; however, the phase difference between the softest and hardest bands can be as high as \u223c0.3 cycles (or \u223c110\u00b0; Bult et al. 2019).","Citation Text":["Chou et al. 2008"],"Functions Text":["Similar energy dependent pulse arrival time behavior can also be found in","and XTE J1807\u2013294","with different break point energies."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[748,764]],"Functions Start End":[[528,601],[729,746],[766,802]]}
{"Identifier":"2015ApJ...805....4Z__Terradas_et_al._2015_Instance_1","Paragraph":"Solar prominences or filaments are cool, dense plasmas embedded in the million-Kelvin corona (Mackay et al. 2010). The plasmas originate from the direct injection of chromospheric materials into a pre-existing filament channel, levitation of chromospheric mass into the corona, or condensation of hot plasmas from the chromospheric evaporation due to the thermal instability (Xia et al. 2011, 2012; Keppens & Xia 2014; Zhou et al. 2014). Prominences are generally believed to be supported by the magnetic tension force of the dips in sheared arcades (Guo et al. 2010b; Terradas et al. 2015) or twisted magnetic flux ropes (MFRs; Cheng et al. 2012, 2014a; Su & van Ballegooijen 2012; Sun et al. 2012a; Zhang et al. 2012a; Xia et al. 2014a, 2014b). They can keep stable for several weeks or even months, but may get unstable after being disturbed. Large-amplitude and long-term filament oscillations before eruption have been observed by spaceborne telescopes (Chen et al. 2008; Li & Zhang 2012; Zhang et al. 2012b; Bi et al. 2014; Shen et al. 2014) and reproduced by numerical simulations (Zhang et al. 2013), which makes filament oscillation another precursor for coronal mass ejections (CMEs; Chen 2011) and the accompanying flares. When the twist of a flux rope supporting a filament exceeds the threshold value (2.5\u03c0\u20133.5\u03c0), it will also become unstable and erupt due to the ideal kink instability (KI; Hood & Priest 1981; Kliem et al. 2004; T\u00f6r\u00f6k et al. 2004, 2010; Fan 2005; Srivastava et al. 2010; Aschwanden 2011; Kumar et al. 2012). However, whether the eruption of the kink-unstable flux rope is failed or ejective depends on how fast the overlying magnetic field declines with height (T\u00f6r\u00f6k & Kliem 2005; Liu 2008; Kumar et al. 2010). When the decay rate of the background field exceeds a critical value, the flux rope loses equilibrium and erupts via the so-called torus instability (TI; Kliem & T\u00f6r\u00f6k 2006; Amari et al. 2014; Jiang et al. 2014). On the other hand, if the confinement from the background field is strong enough, the filament will decelerate to reach the maximum height before falling back to the solar surface, which means the eruption has failed (Ji et al. 2003; Liu et al. 2009; Guo et al. 2010a; Kumar et al. 2011; Joshi et al. 2013, 2014; Song et al. 2014).","Citation Text":["Terradas et al. 2015"],"Functions Text":["Prominences are generally believed to be supported by the magnetic tension force of the dips in sheared arcades"],"Functions Label":["Background"],"Citation Start End":[[569,589]],"Functions Start End":[[438,549]]}
{"Identifier":"2019AandA...631A..88Y__Bohren_&_Huffman_(1998)_Instance_2","Paragraph":"Starting from the four aforementioned materials, we consider several composition mixtures and grain structures. For the sake of comparison, we first consider compact grains of purely a-Sil, a-C, or a-C:H. Subsequently, according to K\u00f6hler et al. (2015), we consider compact grains made of two thirds a-Sil and one third a-C (Mix 1) or one third a-C:H (Mix 2), in terms of volume fractions. These allow reproduction of the mass fractions derived by Jones et al. (2013) for the diffuse ISM. The effect of porosity is tested for the Mix 1 mixture, with a porosity degree of 50% (Mix 1:50). We also evaluate theimpact of the presence of a water ice mantle on compact Mix 1 grains (Mix 1:ice). We further consider two material compositions defined in Pollack et al. (1994) based on depletion measurements: (i) 21% a-Sil and 79% a-C (Mix 3); and (ii) 8% a-Sil, 30% a-C, and 62% water ice (Mix 3:ice). The various grain compositions are summarised in Table 1. For each grain composition, we derive the absorption and scattering efficiencies Qabs and Qsca, respectively, and the asymmetry factor of the phase function g = \u27e8cos\u03b8\u27e9. To allow fast calculations, we make the major assumption that the grains are spherical and compute their optical properties using the Mie theory (Mie 1908; Bohren & Huffman 1983) with the Fortran 90 version of the BHMIE routine given in Bohren & Huffman (1998). For grains consisting of two or three materials, we first derive effective optical constants following the Maxwell Garnett mixing rule (Maxwell Garnett 1904; Bohren & Huffman 1998). Indeed, we assume that in Mix 1 grains, for example, carbon appears as proper inclusions in the silicate matrix rather than assuming a completely random inhomogeneous medium. Mishchenko et al. (2016a,b) performed exhaustive studies of the applicability of the Maxwell Garnett mixing rule to heterogeneous particles. These latter authors showed that this rule can provide accurate estimates of the scattering matrix and absorption cross-section of heterogeneous grains at short wavelengths (typically up to the visible for a 0.1 \u03bcm grain and to the mid-infrared(MIR) for a 10 \u03bcm grain) if twocriteria are met: both the size parameter of the inclusions and the refractive index contrast between the host material and the inclusions have to be small. Moreover, Mishchenko et al. (2016a) demonstrated that the extinction and asymmetry-parameter errors of the Maxwell Garnett mixing rule are significantly smaller than the scattering-matrix errors, remaining small enough for most typical applications and in particular the kind of applications we perform here. It is however well known that this kind of mixing rule systematically underestimates the absorption efficiency in the FIR to millimetre wavelength range, the implications of which are discussed in Sect. 3.2. We perform our computations with the emc routine of V. Ossenkopf3. For Mix 1 and Mix 2, we assume a matrix of a-Sil with inclusions of a-C or a-C:H, and for Mix 3 a matrix of a-C with inclusions of a-Sil. For grains surrounded by an ice mantle, the optical properties are derived with the core-mantle Mie theory using the BHCOAT routine given in Bohren & Huffman (1998).","Citation Text":["Bohren & Huffman 1998"],"Functions Text":["For grains consisting of two or three materials, we first derive effective optical constants following the Maxwell Garnett mixing rule"],"Functions Label":["Uses"],"Citation Start End":[[1542,1563]],"Functions Start End":[[1384,1518]]}
{"Identifier":"2022AandA...663A..15W__Widmark_et_al._2021a_Instance_1","Paragraph":"The dynamics of stars can be related to the gravitational potential that they inhabit via the collisionless Boltzmann equation. For systems in a steady state with certain symmetry properties (typically spherical or axisymmetric) it is possible to find solutions to the phase-space distribution of a stellar tracer population, either through distribution function modelling or via the moments of the Boltzmann equation (Kapteyn 1922; Oort 1932; Bahcall 1984a,b; Kuijken & Gilmore 1989; Cr\u00e9z\u00e9 et al. 1998; Holmberg & Flynn 2000; Bienayme et al. 2006; Binney & Tremaine 2008; Garbari et al. 2012; Bovy & Rix 2013; Salomon et al. 2020; Guo et al. 2020; Widmark et al. 2021a). Given the relatively quiet conditions necessary to form disk galaxies, the assumption of equilibrium for near-equilibrium systems has been widely and favourably applied to the Milky Way and other galaxies (McMillan 2011; Binney & McMillan 2011). Our place in the Milky Way makes it ideal to accurately measure its gravitational potential and mass distribution, since it is the only system where we have access to full six-dimensional phase space information, from its inner regions all the way to its outermost edge (e.g. Deason et al. 2021). A precise and robust measurement of the gravitational potential is crucial for our general understanding of the Milky Way (Dehnen & Binney 1998; Klypin et al. 2002; Widrow et al. 2008; Weber & de Boer 2010; McMillan 2011, 2017; Kafle et al. 2014; Cole & Binney 2017), and also for probing its distribution of dark matter (Read 2014; Nitschai et al. 2020; Cautun et al. 2020; Li et al. 2020; de Salas & Widmark 2021). The local dark matter density is of fundamental importance for direct and indirect dark matter detection experiments (Jungman et al. 1996; Klasen et al. 2015). In a broader sense, dark matter\u2019s gravitational signatures, studied via stellar dynamics and gravitational lensing, is one of the most competitive avenues for constraining its thus far elusive particle nature (Bertone & Tait 2018; Ferreira 2021). The Gaia satellite has been instrumental to this field, pushing the size of the astrometric sample from a few hundred thousand stars (Perryman et al. 1997) to roughly two billion (Gaia Collaboration 2018a).","Citation Text":["Widmark et al. 2021a"],"Functions Text":["The dynamics of stars can be related to the gravitational potential that they inhabit via the collisionless Boltzmann equation. For systems in a steady state with certain symmetry properties (typically spherical or axisymmetric) it is possible to find solutions to the phase-space distribution of a stellar tracer population, either through distribution function modelling or via the moments of the Boltzmann equation"],"Functions Label":["Background"],"Citation Start End":[[649,669]],"Functions Start End":[[0,417]]}
{"Identifier":"2018AandA...618A.128C__Kenney_et_al._2004_Instance_1","Paragraph":"In addition to the galaxy mergers, the ram pressure may also induce AGN activity together to star formation in cluster galaxies (e.g., Poggianti et al. 2017; Marshall et al. 2018; Ramos-Mart\u00ednez et al. 2018). The ram pressure, acting on galaxies in a galaxy cluster moving through the intra-cluster medium (ICM), can strip (ram pressure stripping) gas out of the galaxy where the gas is gravitationally bound to the galaxy less strongly than the force from the ICM medium wind due to the ram pressure. A typical effect of the ram pressure stripping is the formation of long tails of stripped gas behind galaxies. A good example of this effect is the Virgo cluster, the closest rich galaxy cluster, considered a ram pressure stripping laboratory since a number of clearly ram pressure stripped galaxies have been observed (e.g., Kenney et al. 2004; Crowl et al. 2005; Chung et al. 2007). Another example of long tails of stripped gas is the Norma cluster (Abell 3627), containing a very extended multi-phase gas stripped tail of a late-type galaxy (ESO137-001 Sun et al. 2007, 2010). In the literature, there is an open debate on the role of the ram pressure stripping on the AGN activity. Observations suggest that ram-pressure stripping tends to produce a decreasing of the radiative-mode AGN activity in the centres of clusters with respect to low density regions (e.g., Ellison et al. 2011; Ehlert et al. 2014; Khabiboulline et al. 2014). This happens because the ram pressure has depleted the gas supply of the central galaxies. However, models and hydrodynamical simulations show that a moderate ram pressure is also able to compress the gas in the galaxy and trigger star formation (e.g., Fujita & Nagashima 1999; Kronberger et al. 2008; Kapferer et al. 2009; Tonnesen & Bryan 2009; Bekki 2014). This effect, produced by a low level of ram pressure, is observationally supported (e.g., Lee et al. 2017). Moderate ram pressure can also produce the removal of angular momentum from the disk gas (Tonnesen & Bryan 2009), producing gravitational instability in the galactic disk (Schulz & Struck 2001) and consequentially the gas fueling towards the central AGN.","Citation Text":["Kenney et al. 2004"],"Functions Text":["A good example of this effect is the Virgo cluster, the closest rich galaxy cluster, considered a ram pressure stripping laboratory since a number of clearly ram pressure stripped galaxies have been observed (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[828,846]],"Functions Start End":[[613,827]]}
{"Identifier":"2016MNRAS.463.4121T__Brouillet_et_al._2005_Instance_2","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"],"Functions Text":["The HCN(1\u20130)\/HCO+ ratios for M31 GMCs","are indicated by the turquoise shaded region in panels c and d."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1631,1652]],"Functions Start End":[[1592,1629],[1654,1717]]}
{"Identifier":"2022MNRAS.516..731B__Hattori,_Valluri_&_Vasiliev_2021_Instance_1","Paragraph":"Now with the recent availability of high quality full 6D phase-space information for large numbers of sources, much effort has been made to decrease the uncertainties in the Milky Way mass estimate. Recent works using a tracer mass estimator with 6D phase-space information include Sohn et al. (2018, globular clusters), Watkins et al. (2019, globular clusters), and Fritz et al. (2020, satellites). The most recent work using the spherical Jeans equation by Zhai et al. (2018) is very similar to our current investigation in method and data (LAMOST K giants) although only line-of-sight velocities were included, whereas we additionally make use of proper motions from Gaia to obtain the stellar tangential velocities. Using Bayesian analysis to fit a distribution function to full 6D phase-space data (globular clusters, satellites, and halo stars) has been a recent popular choice among many works (Eadie & Juri\u0107 2019; Posti & Helmi 2019; Vasiliev 2019; Deason et al. 2021; Correa Magnus & Vasiliev 2022; Shen et al. 2022; Slizewski et al. 2022; Wang, Hammer & Yang 2022) and a similar distribution function analysis using 5D phase-space data from Gaia (Hattori, Valluri & Vasiliev 2021). In addition to fitting the observational data with a distribution function, several works have incorporated into the fitting a comparison of the observed data with Milky Way-type galaxies from cosmological simulations (Callingham et al. 2019; Li et al. 2020). Newly discovered high velocity stars with full 6D phase-space information have been used to estimate the mass of the Milky Way (Hattori et al. 2018; Monari et al. 2018; Deason et al. 2019; Grand et al. 2019; Koppelman & Helmi 2021; Necib & Lin 2022). Vasiliev, Belokurov & Erkal (2021) and Craig et al. (2021) have estimated the Milky Way mass by fitting models for the Sagittarius and Magellanic Steams, respectively. Several recent studies have estimated the Milky Way mass using measurements of the rotation curve (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). Other works have used 6D satellite phenomenology, characterizing simulated Milky Way-type satellite populations and comparing to the observations of satellites in the Milky Way, to estimate the mass of the Milky Way (Patel et al. 2018; Villanueva-Domingo et al. 2021; Rodriguez Wimberly et al. 2022). Zaritsky et al. (2020) apply the timing argument to distant Milky Way halo stars to derive a lower limit to the Milky Way mass.","Citation Text":["Hattori, Valluri & Vasiliev 2021"],"Functions Text":["Using Bayesian analysis to fit","and a similar distribution function analysis using 5D phase-space data from Gaia"],"Functions Label":["Background","Background"],"Citation Start End":[[1157,1189]],"Functions Start End":[[720,750],[1075,1155]]}
{"Identifier":"2018ApJ...866L...1S__Pecharrom\u00e1n_et_al._1999_Instance_2","Paragraph":"It was found that the complex dielectric function from Pecharrom\u00e1n et al. (1999) for the sample obtained by heating bayerite at 1273 K, assuming a spheroid with depolarization parameters of (0.35, 0.003), produced an opacity with 11, 20, 28, and 32 \u03bcm features, so this component was included in the models. However, with only this component, the observed 20 \u03bcm features in the residual spectra were found to be wider than those in the models. By adding the opacity of the sample obtained by heating boehmite at 1173 K, the width of the 20 \u03bcm feature could be matched. This was done using the complex dielectric function for the sample obtained by heating boehmite at 1173 K from Pecharrom\u00e1n et al. (1999), assuming a spheroid with depolarization parameters of (0.35, 0.035). The complex dielectric functions of the samples obtained by heating bayerite and boehmite to various temperatures (Pecharrom\u00e1n et al. 1999) were derived by modeling the reflectance spectra of pellets obtained by pressing powders of these materials under great pressure. This method required Pecharrom\u00e1n et al. (1999) to assume an effective medium theory, such that a pellet is a mixture of one of their samples with a matrix of air. Pecharrom\u00e1n et al. (1999) noted that heating bayerite at 500\u00b0C eliminates the XRD pattern of bayerite, and they note that at 700\u00b0C, the infrared reflectance spectrum of the boehmite sample no longer shows OH\u2212 stretching bands. This must mean that the samples obtained from heating bayerite at 1273 K and from heating boehmite at 1173 K are no longer bayerite or boehmite, respectively. XRD performed by Pecharrom\u00e1n et al. (1999) of the sample of bayerite prepared at 1273 K suggests only \u03b8-alumina was present, and their infrared and NMR spectroscopy confirms this. XRD of their sample obtained from heating boehmite to 1173 K (Pecharrom\u00e1n et al. 1999) suggests \u03b4-alumina to be present, though some amounts of \u03b8-alumina and \u03b1-alumina are present, as they deduce from XRD and infrared and NMR spectroscopy.","Citation Text":["Pecharrom\u00e1n et al. (1999)"],"Functions Text":["This was done using the complex dielectric function for the sample obtained by heating boehmite at 1173 K from","assuming a spheroid with depolarization parameters of (0.35, 0.035)."],"Functions Label":["Uses","Uses"],"Citation Start End":[[680,705]],"Functions Start End":[[569,679],[707,775]]}
{"Identifier":"2018ApJ...854...63O__Nakamura_et_al._2015_Instance_1","Paragraph":"Many of the earliest simulations of core-collapse supernovae used or were concerned with a general relativistic (GR) approach (Misner & Sharp 1964; Colgate & White 1966; May & White 1966; Wilson 1971; Bruenn 1985; Burrows 1988), and rightly so since neutron stars are sufficiently compact that GR dramatically effects their equilibrium structure. While the importance of including GR gravity in simulations of core-collapse supernovae has always persisted in the literature and is in use in many current and state-of-the-art multidimensional, core-collapse supernova calculations (Marek & Janka 2009; Kuroda et al. 2012; M\u00fcller et al. 2012a; Bruenn et al. 2013, 2016; Ott et al. 2013; Lentz et al. 2015; Skinner et al. 2016; Burrows et al. 2018), many modern simulations have used a purely Newtonian approximation for gravity (Takiwaki et al. 2014; Couch & O\u2019Connor 2014; Handy et al. 2014; Couch & Ott 2015; Dolence et al. 2015; Nakamura et al. 2015; Pan et al. 2016; Suwa et al. 2016). Liebend\u00f6rfer et al. (2001) extensively compared GR gravity and Newtonian gravity in spherical symmetry with a full Boltzmann neutrino transport solver. While their baseline simulations in both prescriptions of gravity fail to explode in 1D, their conclusion is that, overall, GR is helpful for the development of the core-collapse supernova explosion. This conclusion comes out of serendipitous simulations where incorrect nucleon isoenergetic scattering cross sections were used. In these simulations, when GR gravity was used, the simulation predicted an explosion, but when Newtonian gravity was used, the simulations failed to achieve explosions. Bruenn et al. (2001), Buras et al. (2006b), and Lentz et al. (2012) also compared Newtonian and GR simulations in spherical symmetry with energy-dependent neutrino transport. Like Liebend\u00f6rfer et al. (2001), they observe higher neutrino luminosities and energies but do not extensively study the differences, in part because both simulations still fail to explode owing to the 1D nature of the simulations. Kuroda et al. (2012) examined the difference between special relativistic hydrodynamics and GR hydrodynamics in both 1D and 3D core-collapse simulations using approximate neutrino transport (a gray M1 scheme). However, while the 3D GR simulations show signs of increased susceptibility to explosion, the simulated time was not long enough to observe an explosion. M\u00fcller et al. (2012a) also investigated the influence of Newtonian gravity, GR gravity, and GR effective potential gravity in 2D core-collapse simulations using a variable-Eddington-factor, two-moment, energy-dependent neutrino transport scheme and the ray-by-ray+ approximation (where the neutrino transport is done along radial rays assuming spherical symmetry and with minimal coupling between neighboring rays). For the classic 15 M\u2299 model from Woosley & Weaver (1995), M\u00fcller et al. (2012a) find a late and asymmetric explosion when using fully GR gravity, but not with Newtonian gravity, GR effective potential gravity, or reduced set neutrino opacities. This was, and still is, the most direct evidence that GR aids in the explosion mechanism. It is important to validate the hypothesis that GR gravity does indeed lead to more favorable conditions for explosion in modern-day simulations, with updated progenitor models, and loosening the assumption of spherically symmetric.","Citation Text":["Nakamura et al. 2015"],"Functions Text":["While the importance of including GR gravity in simulations of core-collapse supernovae has always persisted in the literature and is in use in many current and state-of-the-art multidimensional, core-collapse supernova calculations","many modern simulations have used a purely Newtonian approximation for gravity"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[930,950]],"Functions Start End":[[347,579],[747,825]]}
{"Identifier":"2017MNRAS.464..635M__Ceverino_et_al._2012_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":["Ceverino et al. 2012"],"Functions Text":["This has been shown using","as well as cosmological simulations"],"Functions Label":["Background","Background"],"Citation Start End":[[803,823]],"Functions Start End":[[467,492],[713,748]]}
{"Identifier":"2021MNRAS.504.1473L__Mu\u00f1oz,_Miranda_&_Lai_2019_Instance_1","Paragraph":"Large numbers of works have extensively studied the binary-disk system in the Newtonian regime with analytic methods (Paczynski 1977; Papaloizou & Pringle 1977; Kocsis, Haiman & Loeb 2012a, b) and hydrodynamical simulations (Lin & Papaloizou 1979; Artymowicz & Lubow 1994; MacFadyen & Milosavljevi\u0107 2008; Mayama et al. 2010; de Val-Borro et al. 2011; Shi et al. 2012; D\u2019Orazio et al. 2013; Farris et al. 2014; Ju, Stone & Zhu 2016; Mu\u00f1oz & Lai 2016; Nelson A. & Marzari 2016; Miranda, Mu\u00f1oz & Lai 2017; Tang, MacFadyen & Haiman 2017; Tang, Haiman & MacFadyen 2018; Mu\u00f1oz, Miranda & Lai 2019). Simulations within the Newtonian regime during the inspiral of the binary black hole have also been investigated (Baruteau, Ramirez-Ruiz & Masset 2012; Farris et al. 2015; Cerioli, Lodato & Price 2016; Tang et al. 2018). Two-dimensional simulations by Noble et al. (2012) and three-dimensional ones by Zilh\u00e3o et al. (2015) conducted the fully relativistic magnetohydrodynamic evolution of a circumbinary disk surrounding two non-spinning black holes with equal mass using the near zone metric. Later on, the fully relativistic magnetohydrodynamic simulations of SMBBH\u2013disk interaction with both the near zone metric and the inner zone metric have been performed (Bowen et al. 2017, 2018, 2019; d\u2019Ascoli et al. 2018) and these simulations are limited to several binary orbits. In particular, Bowen et al. (2017) showed that the gravitational potential of the binary black hole in the post-Newtonian (PN) regime is shallower than that in the Newtonian regime, and that such shallow potential in the PN regime has significant effects on the dynamics of the mini-disk around each SMBH. Simulations of SMBBH\u2013disk interaction with full numerical relativity are so computationally expensive that the evolution of disk around SMBBH is constrained within the stage near merger (Farris et al. 2012; Giacomazzo 2012). In this work, we present hydrodynamical simulations of accretion of equal mass SMBBH by using PN hydrodynamics in which the near zone metric of SMBBH is used.","Citation Text":["Mu\u00f1oz, Miranda & Lai 2019"],"Functions Text":["Large numbers of works have extensively studied the binary-disk system in the Newtonian regime with","and hydrodynamical simulations"],"Functions Label":["Background","Background"],"Citation Start End":[[565,590]],"Functions Start End":[[0,99],[193,223]]}
{"Identifier":"2021ApJ...910...82L__Schrijver_2001_Instance_1","Paragraph":"Similar results of the on-disk condensations are obtained to those of the on-disk quiescent coronal rain in AR closed loops in H\u03b1 (see Section 3.2). If we considered the AIA 304 \u212b observations alone, the on-disk condensation, and hence coronal rain events we report here would resemble those occurring in magnetically closed field lines (Antolin & Rouppe van der Voort 2012; Antolin et al. 2012; Ahn et al. 2014), which are generally interpreted as a manifestation of the heating-condensation cycles due to thermal nonequilibrium (Schrijver 2001; M\u00fcller et al. 2003, 2004; Antolin 2020). However, combining the observations of SDO and STEREO A and B at different viewing angles, we find that the on-disk coronal rain in this study corresponds to the downflows of condensations facilitated by reconnection between open and closed structures. Similarly, if we analyze only the on-disk observations as recorded by the AIA as shown in Figure 7, it appears that the structures gradually brighten first in the 171 \u212b channel and then in the 131 \u212b channel. Such sequential appearance of coronal structures in the AIA channels sensitive to emission from progressively cooler plasma could be interpreted as a case of loop cooling after nanoflare heating (e.g., Viall & Klimchuk 2012; Li et al. 2015), for instance. In contrast, by employing observations from multiple vantage points, we show that the on-disk structure brightening in the AIA images is actually due to cooling and condensation of plasma facilitated by reconnection in the high corona that is not necessarily regulated by any heating mechanism. To search for the origination of structures harboring on-disk flow, e.g., coronal rain, in the transition region or chromospheric lines, observations from different viewing angles are thus quite important. If there is no observation from other viewing angles, evolution of the associated structures, that may be difficult to observe on the disk, in multiwavelength images during the cooling and condensation process needs to be examined because the condensation facilitated by reconnection cools down from \u223c0.9 MK, the characteristic temperature of the 171 \u212b channel, rather than from the higher temperatures.","Citation Text":["Schrijver 2001"],"Functions Text":["If we considered the AIA 304 \u212b observations alone, the on-disk condensation, and hence coronal rain events we report here would resemble those occurring in magnetically closed field lines","which are generally interpreted as a manifestation of the heating-condensation cycles due to thermal nonequilibrium","However, combining the observations of SDO and STEREO A and B at different viewing angles, we find that the on-disk coronal rain in this study corresponds to the downflows of condensations facilitated by reconnection between open and closed structures."],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[531,545]],"Functions Start End":[[149,336],[414,529],[588,840]]}
{"Identifier":"2022MNRAS.514.5178D__Ruderman_&_Roberts_2002_Instance_1","Paragraph":"There are several natural scenarios for the initial flare-caused impulsive perturbation of an active region to develop into a quasi-periodic response, which includes the effects of resonance in closed coronal plasma structures (acting as resonators), dispersion of a wave-guide, and non-linearity\/self-organization (McLaughlin et al. 2018; Zimovets et al. 2021). The observed doubling of the QPP and shorter lifetime of shorter-period QPP modes strongly indicate in favour of resonant dynamics of magnetohydrodynamic waves in a coronal loop, for which the oscillation period is prescribed by the local plasma conditions (i.e. the Alfv\u00e9n and sound speeds) and the loop length (Nakariakov et al. 2021; Wang et al. 2021). More specifically, fast- and slow-mode magnetoacoustic waves in coronal plasma structures are well-known to be subject to a frequency-dependent damping by e.g. resonant absorption and thermal conduction (e.g. Ofman & Aschwanden 2002; Ruderman & Roberts 2002; De Moortel & Hood 2003). Likewise, excitation of even and\/or uneven parallel harmonics of fast- and slow-mode standing waves in a plasma loop is known to be highly sensitive to the location of the initial perturbation along the loop. For example, Tsiklauri et al. (2004) and Selwa, Murawski & Solanki (2005) theoretically demonstrated that the second parallel harmonic of a slow standing wave can be effectively excited if the impulsive energy release occurs near the apex of the loop. Observations of higher harmonics of fast-mode oscillations in coronal loops were also shown to be subject to the excitation mechanism and location of the initial displacement of the loop (e.g. De Moortel & Brady 2007; Srivastava et al. 2008; Yuan & Van Doorsselaere 2016; Pascoe et al. 2017; Duckenfield et al. 2019). Thus, taking cs = 600\u2009km\u2009s\u22121 for the sound speed in a hot flaring loop, the Alfv\u00e9n speed cA = 1200\u2009km\u2009s\u22121 (e.g. Mathioudakis et al. 2006), and treating the observed QPP periods as characteristic acoustic or Alfv\u00e9n transit times along the loop, we can estimate the corresponding loop lengths as 80\u2013160\u2009Mm for the 2018 flare and 50\u2013100\u2009Mm for the 2019 flare, i.e. 0.2\u20130.7 R\u22c6. These estimations agree with a typical length of solar coronal loops, and, are approximately an order of magnitude larger than those derived by Mathioudakis et al. (2006). However, the above authors ruled out these small loop lengths suggesting instead that they may be due to a fast-MHD wave, with the modulation of the emission being due to the magnetic field. The present observation in the doubling of the QPP in both YZ CMi flares presents rare and compelling evidence for the presence of compact plasma loops in a stellar corona.","Citation Text":["Ruderman & Roberts 2002"],"Functions Text":["More specifically, fast- and slow-mode magnetoacoustic waves in coronal plasma structures are well-known to be subject to a frequency-dependent damping by e.g. resonant absorption and thermal conduction (e.g."],"Functions Label":["Uses"],"Citation Start End":[[953,976]],"Functions Start End":[[719,927]]}
{"Identifier":"2018AandA...615A..77L__Tacconi_et_al._2010_Instance_1","Paragraph":"A different approach is to count the amount of baryonic material within the proto-cluster bounds associated with the member galaxies and attempt to relate that back to the overall mass of the structure, an approach which has been employed successfully at low redshift when the galactic baryonic content of cluster member galaxies is dominated by stars (Andreon 2012). The approach we take here is similar to that of Lemaux et al. (2014a). Briefly, the total amount of stellar matter of the zspec members is counted and a completeness correction is made to this value for galaxies at stellar masses \n\n$\\log(\\mathcal{M}_{\\ast}\/M_{\\odot})\\ge 9.5$log(M\u2217\/M\u2299)\u22659.5\n (see Sect. 4.2.1 for the reasoning behind this cut) based on the number of zphot members and non-members without secure spectral redshifts within the bounds of the proto-cluster and the likelihood of their being true members. An additional correction is made to correct for galaxies in the stellar mass range \n\n$8.0 < \\log(\\mathcal{M}_{\\ast}\/M_{\\odot})< 9.5$8.0log(M\u2217\/M\u2299)9.5\n by integrating the stellar mass function of Davidzon et al. (2017) appropriate for this redshift. Here we additionally make the assumption that stellar mass comprises 50% of the baryonic content of galaxies by mass at these redshifts, a value broadly consistent with the few measurements made at or near these redshifts (e.g., Tacconi et al. 2010; Capak et al. 2011b; Schinnerer et al. 2016; Scoville et al. 2016). We assume for the purposes of this calculation that the proto-cluster is a closed system, with all gas being converted to stars by z = 0 and that the completeness-corrected galaxy population which lies within Rproj \u2264 3 Mpc at z ~ 4.57 comprises the entirety of the galaxy population which will eventually be contained within the cluster virial radius at z = 0. The latter assumption is broadly consistent with simulations (Muldrew et al. 2015) to ~10% accuracy, a factor which we account for in the calculation below. Any total to stellar mass conversion then provides a z = 0 total mass, which we de-evolve to z ~ 4.57 using the correction factors of Muldrew et al. (2015) appropriate for z ~ 4.57 (0.20 \u00b1 0.03), a correctionwhich is appropriate for descendents of all masses. Into this formalism we input the resulting completeness-corrected baryonic content of \n\n$\\log(\\mathrm{\\Sigma}\\mathcal{M}_{\\ast}\/M_{\\odot})=12.39^{0.05}_{-0.07}$log(\u03a3M\u2217\/M\u2299)=12.39\u22120.07+0.05\n to the r200 stellar mass to M500 total mass relation of Andreon (2012) and scale the resulting M500 to the virial radius (Rvir = 0.33 Mpc) using the methods presented in Lemaux et al. (2014a) giving:\n\n(2)\n\n\\begin{equation*}\n\\log(\\mathcal{M}\/\\mathcal{M}_{\\odot})_{z=4.57, \\, \\mathrm{\\Sigma}\\mathcal{M}_{\\ast, corr}} = 13.31^{0.23}_{-0.27}.\\end{equation*}log(M\/M\u2299)z=4.57,\u200a\u03a3M\u2217,corr=13.31\u22120.27+0.23.\n","Citation Text":["Tacconi et al. 2010"],"Functions Text":["Here we additionally make the assumption that stellar mass comprises 50% of the baryonic content of galaxies by mass at these redshifts, a value broadly consistent with the few measurements made at or near these redshifts (e.g.,"],"Functions Label":["Uses"],"Citation Start End":[[1362,1381]],"Functions Start End":[[1133,1361]]}
{"Identifier":"2020ApJ...904..185O__Takakuwa_et_al._2014_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":["Takakuwa et al. 2014"],"Functions Text":["Moreover, the disk formation process has been revealed to be much more complicated for binary and multiple cases, both in observations"],"Functions Label":["Background"],"Citation Start End":[[617,637]],"Functions Start End":[[441,575]]}
{"Identifier":"2019ApJ...871..191O__Vaucouleurs_et_al._1991_Instance_1","Paragraph":"NGC 4418 is one of the closest and best-studied galaxies hosting CONs. It is an LIRG with its 8\u20131000 \u03bcm infrared luminosity (LIR) of 1.4 \u00d7 1011 L (Sanders et al. 2003)4\n\n4\nWith our assumed distance, the luminosity is larger than the threshold for the LIRG classification of LIR = 1 \u00d7 1011 L (Sanders & Mirabel 1996). We note, however, that NGC 4418 is sometimes classified as a non-LIRG if one adopts smaller distance to the galaxy. For example, with its recession velocity (cz) and an assumed Hubble constant of H0 = 75 km s\u22121 Mpc\u22121, the distance is estimated to be 28.32 Mpc and LIR = 9. 4 \u00d7 1010 L.\n at a distance of 34 Mpc (1\u2033 = 165 pc; Sakamoto et al. 2013).5\n\n5\nSakamoto et al. (2013) adopted the distance from Sanders et al. (2003), in which distance is calculated with cz using the cosmic attractor model outlined in Appendix A of Mould et al. 2000, using H0 = 75 km s\u22121 Mpc\u22121 and adopting a flat cosmology in which \u03a9M = 0.3 and \u03a9\u03bb = 0.7. We adopt this distance for consistency with recent closely related studies of this galaxy (Aalto et al. 2012; Costagliola et al. 2013; Varenius et al. 2014).\n In spite of its huge infrared luminosity, the host galaxy is classified from the optical morphology as an ordinary Sa-type galaxy ([R\u2019]SAB[s]a; de Vaucouleurs et al. 1991) with moderate inclination (i = 62\u00b0; taken from Sakamoto et al. 2013 for consistency; based on Jarrett et al. 2000). It has a neighbor galaxy about 32 kpc (32) away (VV 655 = MCG+00- 32-013 = SDSS J122704.47-005420.6; \n\n\n\n\n\n km s\u22121; Evans et al. 2003; Varenius et al. 2017). Varenius et al. (2017) found a bridge of H i 21 cm emission connecting these two galaxies and argued that their strong tidal interaction occurred about 190 Myr ago. The optical nuclear emission line spectrum of NGC 4418 is classified as LINER (Armus et al. 1989; Shi & Gu 2005) or Seyfert 2 (Baan et al. 1998), but no direct signature of AGNs such as very broad permitted lines is known (see also Lehnert & Heckman 1995). NGC 4418 is deficient of X-ray emission for its infrared luminosity, and its interpretation has been a matter of controversy. Maiolino et al. (2003) tentatively identified, with low statistics, a Compton-thick (reflection-dominated) AGN on the basis of its small photon index (or flat X-ray spectrum; \n\n\n\n\n\n, where \n\n\n\n\n\n, E is the photon energy, and N(E) is the photon number density). On the other hand, Lehmer et al. (2010) found the opposite (soft) spectral index (\n\n\n\n\n\n), which is even softer than the one of Compton-thin AGNs (\u0393 \u2243 1.7; Mushotzky et al. 1993). We note that both analyses are based on the Chandra AICS-S data, and Lehmer et al. (2010) added \u224330% more integration onto the earlier data used by Maiolino et al. (2003).","Citation Text":["de Vaucouleurs et al. 1991"],"Functions Text":["In spite of its huge infrared luminosity, the host galaxy is classified from the optical morphology as an ordinary Sa-type galaxy ([R\u2019]SAB[s]a;"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1250,1276]],"Functions Start End":[[1106,1249]]}
{"Identifier":"2017ApJ...849..149S__Lee_2004_Instance_1","Paragraph":"There is a similarity between the ISM and the young stellar structures identified here. First, similar to the young stellar structures, the ISM also displays irregular morphologies and contains large amounts of substructures (clouds, clumps, cores, filaments, etc.) that are hierarchically organized (Rosolowsky et al. 2008). Second, the ISM substructures also follow a power-law size distribution, which indicates a scale-free behavior (e.g., Elmegreen & Falgarone 1996). The third aspect of their similarity comes from the fractal dimension. The projected fractal dimension of the ISM has been investigated based on the perimeter\u2013area relation of its projected boundaries. Typical values are close to D2 = 1.4\u20131.5 (e.g., Beech 1987; Scalo 1990; Falgarone et al. 1991; Vogelaar & Wakker 1994; Lee 2004; Lee et al. 2016; although smaller values have also been reported by, e.g., Dickman et al. 1990; Hetem & Lepine 1993), which are consistent with the fractal dimension as derived for the young stellar structures. Using power-spectrum analysis, Stanimirovic et al. (1999, 2000) reported D2 = 1.4 or 1.5 for the ISM in the SMC, also close to that of the young stellar structures. On the other hand, it is possible to measure the volume fractal dimension of the ISM, since clouds along the line of sight can be distinguished by their velocities. For instance, Elmegreen & Falgarone (1996) reported D3 = 2.3 \u00b1 0.3 based on the size distribution for a number of Galactic molecular clouds, and Roman-Duval et al. (2010) found D3 = 2.36 \u00b1 0.04 using the mass\u2013size relation. If the relation D3 = D2 + 1 holds for the ISM, these results would not be far from the fractal dimension of the young stellar structures. Unfortunately, there is no reported measurement of the fractal dimension of the ISM in the LMC bar region. However, it has been suggested that, despite a few exceptions, the fractal dimension is invariant from cloud to cloud, regardless of their nature as star-forming or quiescent, whether gravitationally bound or unbound (e.g., Williams et al. 2000).","Citation Text":["Lee 2004"],"Functions Text":["Typical values are close to D2 = 1.4\u20131.5","which are consistent with the fractal dimension as derived for the young stellar structures."],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[794,802]],"Functions Start End":[[675,715],[922,1014]]}
{"Identifier":"2017MNRAS.470..612F__Feng_etal._2016_Instance_1","Paragraph":"The millimetre bump in M87 as recently observed by the Atacama Large Millimeter\/submillimeter Array can be naturally modelled by the synchrotron emission of the thermal electrons in the ADAF, which is different from the prediction of the jet model. Therefore, it provides an opportunity to explore the accretion process near the BH horizon. In particular, Feng etal. (2016) and Li etal. (2016) both found that the rotation measure predicted from the ADAF is roughly consistent with the observational values. It is still difficult to constrain the BH spin parameter from the modelling of the SED of M87 due to some degeneracy in model parameters (wind parameter, s, magnetic parameter ; Feng etal. 2016). The spin parameter can be better constrained from the jet model if the relativistic jet is indeed powered by the rotating BHs as suggested by MHD simulations and some observations. We find that the dimensionless BH spin parameter should be larger than 0.96 for the lower limit of jet power derived from the X-ray cavities (e.g. Rafferty et al. 2006; Russell etal. 2013b). In this work, we adopt several typical values of parameters (e.g. 0.1, 0.3 and 0.5). The larger value of  will lead to a lower accretion rate near the horizon to explain the observed millimetre bump, and the BHs need to rotate faster to reproduce the observed jet power. The peak of synchrotron emission from the thermal electrons of ADAF will move to the submillimetre waveband if  is too small, which is different from the observed millimetre bump. We adopt the equipartition case of 0.5 in our calculations, where magnetic energy will become dominant if the BH is fast spinning, considering the possible amplification of the magnetic field by the frame dragging effect. For the weaker magnetic case (e.g. 0.5), the BH needs to rotate faster to explain the observed SED and jet power. We find that our results are not sensitive to the viscosity parameter . Therefore, we suggest that the BH should be fast rotating in M87 even after considering the possible uncertainties.","Citation Text":["Feng etal. (2016)"],"Functions Text":["In particular,","and Li etal. (2016) both found that the rotation measure predicted from the ADAF is roughly consistent with the observational values."],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[356,373]],"Functions Start End":[[341,355],[374,507]]}
{"Identifier":"2022MNRAS.515.2188H__Rorai_et_al._2018_Instance_1","Paragraph":"In this work, we follow the method for measuring the IGM thermal state based on Voigt profile decomposition of the Ly\u2009\u03b1 forest (Schaye et al. 1999; Ricotti et al. 2000; McDonald et al. 2001). In this approach, a transmission spectrum is treated as a superposition of multiple discrete Voigt profiles, with each line described by three parameters: redshift zabs, Doppler broadening b, and neutral hydrogen column density $N_{\\rm H\\, {\\small I}}$. By studying the statistical properties of these parameters, i.e. the $b-N_{{{{\\rm H\\, {\\small I}}}}{}}$ distribution, one can recover the thermal information encoded in the absorption profiles. The majority of past applications of this method constrained the IGM thermal state by fitting the low-b-$N_{\\rm H\\, {\\small I}}$ cutoff of the $b-N_{{{{\\rm H\\, {\\small I}}}}{}}$ distribution (Schaye et al. 1999, 2000; Ricotti et al. 2000; McDonald et al. 2001; Rudie, Steidel & Pettini 2012; Boera et al. 2014; Bolton et al. 2014; Garzilli, Theuns & Schaye 2015, 2020; Hiss et al. 2018; Rorai et al. 2018). The motivation for this approach is that the Ly\u2009\u03b1 lines are broadened by both thermal motion and non-thermal broadening resulting from combinations of Hubble flow, peculiar velocities, and turbulence. By isolating the narrow lines in the Ly\u2009\u03b1 forest that constitutes the lower cutoff in $b-N_{{{{\\rm H\\, {\\small I}}}}{}}$ distributions, of which the line-of-sight component of non-thermal broadening is expected to be zero, the broadening should be purely thermal, thus allowing one to constrain the IGM thermal state. However, this method has three crucial drawbacks. First, the IGM thermal state actually impacts all the lines besides just the narrowest lines. Therefore, by restricting attention to data in the distribution outskirts, this approach throws away information and reduces the sensitivity to the IGM thermal state significantly (Rorai et al. 2018; Hiss et al. 2019). Secondly, in practice, determining the location of the cutoff is vulnerable to systematic effects, such as contamination from the narrow metal lines (Hiss et al. 2018; Rorai et al. 2018). Lastly, the results from this approach critically depend on the choice of low-b cutoff fitting techniques, where different techniques might result in inconsistent T0 and \u03b3 measurements (Hiss et al. 2018; Rorai et al. 2018).","Citation Text":["Rorai et al. 2018"],"Functions Text":["The majority of past applications of this method constrained the IGM thermal state by fitting the low-b-$N_{\\rm H\\, {\\small I}}$ cutoff of the $b-N_{{{{\\rm H\\, {\\small I}}}}{}}$ distribution"],"Functions Label":["Background"],"Citation Start End":[[1027,1044]],"Functions Start End":[[640,830]]}
{"Identifier":"2016ApJ...826..117Y__Roux_&_Webb_2009_Instance_2","Paragraph":"Traditionally, the Parker transport equation (Parker 1965) was used to model pickup ion acceleration at the SWTS when using a transport theory approach. However, interesting Voyager results, such as strongly fluctuating pitch-angle anisotropies upstream, the detection of highly anisotropic intensity spikes at the SWTS, the average upstream anisotropy peaking at a surprisingly high energy far above the upstream flow energy, and energetic particle spectra with multiple power-law slopes with breaking points in between that are highly variable upstream (Decker et al. 2005, 2008b; Cummings et al. 2006), suggest that transport modeling should be modified in at least two ways. First, the turbulent nature of the magnetic field conditions at the SWTS should be taken into account, and second, a transport theory that applies when particle distributions are highly anisotropic is needed, given that the Parker transport equation only holds for nearly isotropic particle distributions. In response, shock acceleration transport models were developed in recent years based on the numerical solutions of the focused transport equation (K\u00f3ta & Jokipii 2004; le Roux et al. 2007; Florinski et al. 2008a, 2008b; le Roux & Webb 2009) to take advantage of the fact that focused transport is not restricted to small pitch-angle anisotropies. This is especially advantageous at lower suprathermal particle energies upstream, where particle distributions can be sporadically strongly anisotropic (Decker et al. 2006), allowing one to model particle injection into diffusive shock acceleration (DSA) naturally at those energies. Furthermore, statistical variations in the observed magnetic field direction near the SWTS were included as a time series to model time variations in the injection and DSA of pickup ions to simulate the highly volatile nature of actual DSA at the SWTS (Florinski et al. 2008a, 2008b; le Roux & Webb 2009; Arthur & le Roux 2013). This focused transport approach should be seen as complementary to more sophisticated self-consistent shock acceleration models based on hybrid codes (Kucharek & Scholer 1995; Giacalone 2005) and particle-in-the-cell models (Scholer et al. 2003; Lembege et al. 2004), but it has the virtue of relative simplicity because different statistical plasma parameters can easily be studied separately and in combination at the SWTS to gain a more clear conceptional understanding of the role of such statistics on pickup acceleration at the SWTS.","Citation Text":["le Roux & Webb 2009"],"Functions Text":["Furthermore, statistical variations in the observed magnetic field direction near the SWTS were included as a time series to model time variations in the injection and DSA of pickup ions to simulate the highly volatile nature of actual DSA at the SWTS"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1901,1920]],"Functions Start End":[[1617,1868]]}
{"Identifier":"2022MNRAS.515.5629D__Dubber_et_al._2021_Instance_1","Paragraph":"In the previous sections, we have shown the difficulty in balancing follow-up time with large numbers of unconfirmed candidate companions from single-epoch photometry. This follow-up dilemma was the motivation for designing the K-peak filter: can we optimize the observing time required to determine the nature of a target? One answer is to use a carefully chosen combination of photometric filters. Using only photometry, we can calculate colours that contain information about the type of object being observed, allowing approximate characterization with single-epoch photometry. Past works (e.g. Najita, Tiede & Carr 2000; Allers & Liu 2020) have shown that custom photometric filters can be used to greatly improve the confirmation rate of photometrically selected candidate low-mass brown dwarfs. In previous work (Allers & Liu 2020; Jose et al. 2020; Dubber et al. 2021), we used a custom filter centred on the deep 1.45 \u00b5m feature present in YPMOs to distinguish between them and background sources. In the 2\u20135 \u00b5m range covered by NIX, such water features are far less dominant, and there are strong telluric features across some of this range that would make a similar \u2018water\u2019 technique difficult to use. Instead, we use the differing spectral shape in K-band of very low mass brown dwarfs when compared to earlier spectral type stars. This can be seen in the sample of spectra shown in Fig. 3. By locating filters at key spectral points for defining the overall shape of the spectra, the extracted colour information can be used for direct characterization. Also shown in Fig. 3 are spectra of well-studied brown dwarfs and exoplanets, discovered via direct imaging: 51 Eri b (Macintosh et al. 2015), \u03b2 Pictoris b (Lagrange et al. 2010), PSOJ-318 (Liu et al. 2013), G196-3B (Rebolo et al. 1998), and HR 8799d and e (Marois et al. 2008, 2010). References for the spectral data plotted are detailed in the caption of Fig. 3. These spectra demonstrate the general variety in the spectral shapes of objects that have been detected via direct imaging previously, but also the similar features in the highlighted filter windows.","Citation Text":["Dubber et al. 2021"],"Functions Text":["In previous work","we used a custom filter centred on the deep 1.45 \u00b5m feature present in YPMOs to distinguish between them and background sources. In the 2\u20135 \u00b5m range covered by NIX, such water features are far less dominant, and there are strong telluric features across some of this range that would make a similar \u2018water\u2019 technique difficult to use. Instead, we use the differing spectral shape in K-band of very low mass brown dwarfs when compared to earlier spectral type stars."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[857,875]],"Functions Start End":[[802,818],[878,1343]]}
{"Identifier":"2021MNRAS.507.5882S__Mackereth_et_al._2018_Instance_2","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)"],"Functions Text":["Clarke et al. (2019) and Buck (2020) suggest that bimodality should be common in disc galaxies, whereas","suggest that it is rare."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1680,1703]],"Functions Start End":[[1576,1679],[1704,1728]]}
{"Identifier":"2021ApJ...910...78Z__Condon_et_al._2017_Instance_2","Paragraph":"The multiwavelength spectral data of 13 SNRs with hard \u03b3-ray spectra. The \u03b3-ray spectra are fitted with a hadronic model with the normalization of the individual spectrum as free parameters. The model assumes that protons have a single power-law energy distribution with an exponential high-energy cutoff. Note that the TeV spectra of G78.2+2.1 (HAWC) and N132D (HESS) cut off at relatively lower energies, and the soft spectral component of GeV of HESS 1912+101 may be from other contributors and are not considered in SED fitting. The best-fit model parameters are indicated in the figure. References for the observational data are as follows: RX J0852.0\u22124622: radio (Duncan & Green 2000), GeV (Tanaka et al. 2011), X-ray (Aharonian et al. 2007), TeV (H.E.S.S. Collaboration et al. 2018c); RX J1713.7\u22123946: radio (Lazendic et al. 2004), X-ray (Tanaka et al. 2008), GeV and TeV (H.E.S.S. Collaboration et al. 2018a); HESS J1731\u2212347: radio (Tian et al. 2008), GeV (Condon et al. 2017; Guo et al. 2018), X-ray (Doroshenko et al. 2017), TeV (H.E.S.S. Collaboration et al. 2011); RCW 86: radio (Clark et al. 1975; Lemoine-Goumard et al. 2012), X-ray (Lemoine-Goumard et al. 2012), GeV (Ajello et al. 2016), TeV (H.E.S.S. Collaboration et al. 2018d); SN 1006: radio Dyer et al. 2009, X-ray (Bamba et al. 2008), GeV (Condon et al. 2017), TeV (Acero et al. 2010); G150.3+4.5: radio (Gerbrandt et al. 2014), X-ray and GeV (Devin et al. 2020); G296.5 + 10.0: radio (Milne & Haynes 1994), GeV (this work), HESS J1534\u2212571: radio (Maxted et al. 2018), GeV (Araya 2017), X-ray and TeV (H.E.S.S. Collaboration et al. 2018b); RCW 103: radio (Dickel et al. 1996), GeV (Xing et al. 2014); G78.2+2.1: radio (Wendker et al. 1991; Zhang et al. 1997; Kothes et al. 2006; Gao et al. 2011), X-ray (Leahy et al. 2013), GeV (Abeysekara et al. 2018) and TeV (Fleischhack 2019); G279.0+1.1: radio (Woermann & Jonas 1988; Duncan et al. 1995), GeV Araya (2020); N132D: radio (Dickel & Milne 1995), X-ray (Hughes et al. 1998; Bamba et al. 2018), GeV (Y. L. Xin et al. 2020, in preparation), and TeV (H.E.S.S. Collaboration et al. 2015).","Citation Text":["Condon et al. 2017"],"Functions Text":["RCW 86:","GeV"],"Functions Label":["Uses","Uses"],"Citation Start End":[[1312,1330]],"Functions Start End":[[1077,1084],[1307,1310]]}
{"Identifier":"2022AandA...668A..50S__Ginski_et_al._(2016)_Instance_2","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)"],"Functions Text":["For our disk model, we are using the power-law profile for the scattering surface height H and the separation r found by","for the disk around HD 97048. This profile describes the flaring discovered in this source reasonably well up to a separation of ~270 au."],"Functions Label":["Uses","Uses"],"Citation Start End":[[851,871]],"Functions Start End":[[730,850],[872,1009]]}
{"Identifier":"2021MNRAS.507.6012Z__Kendrick_2018_Instance_1","Paragraph":"Being a benchmark system H + H2, H + HD, and their isotopic counterparts have received much attention over the last several decades (Marinero et al. 1984; Zhang & Miller 1989; D\u2019Mello et al. 1991; Harich et al. 2002; Gao et al. 2015; Yuan et al. 2018a, b, 2020). Most early experimental and theoretical investigations were centered around benchmarking theory against experiments and providing improved descriptions of the H3 potential energy surfaces (PES; Boothroyd et al. 1996; Mielke, Garrett & Peterson 2002; Yuan et al. 2018a, b). Among the available PESs for the H3 system, the one by Boothroyd et al. (1996) referred to as the BKMP2 PES and by Mielke et al. (2002) referred to as the CCI PES, nearly equally well account for most experimental data for H + H2, H + HD, and D + HD collisions. These PESs have also been able to account for even subtle effects such as the GP (Kendrick et al. 2015; Hazra et al. 2016; Croft et al. 2017; Kendrick 2018, 2019). The GP effect, while not significant at temperatures relevant to astrophysics, is important below 1 K as illustrated in a series of calculations on H + HD (v, j) collisions for vibrational levels v = 4 \u2212 9 (Kendrick et al. 2015; Hazra et al. 2016; Croft et al. 2017; Kendrick 2018, 2019). Though several prior studies of Flower and co-workers  (Flower 1999, 2000; Flower & Roueff 1999; Wrathmall et al. 2007) have reported rate coefficients for H + HD collisions, due to the approximations involved in the scattering calculations (e.g. neglect of hydrogen atom-exchange), the reliability of the available rate coefficients has been a source of debate (Desrousseaux et al. 2018). Recently, Desrousseaux et al. (2018) reported rate coefficients for pure rotational transitions for j \u2264 10 within the v = 0 vibrational level using accurate quantum calculations that include the exchange channel. In this paper, we report rate coefficients for state-to-state rovibrational transitions in HD induced by H atoms between and within the v = 0 and 1 vibrational levels and for temperatures ranging from T = 1\u20131000 K.","Citation Text":["Kendrick 2018"],"Functions Text":["These PESs have also been able to account for even subtle effects such as the GP"],"Functions Label":["Background"],"Citation Start End":[[940,953]],"Functions Start End":[[798,878]]}
{"Identifier":"2016MNRAS.461.4406A__Santos_et_al._2015_Instance_1","Paragraph":"However, we note that the quoted eccentricity for the inner planet in HD 155358 is 0.17 \u00b1 0.03 (Robertson et al. 2012a). The mean value is exceeded for the standard run at 4000 orbits after going into resonance with the implication that the planets could not have been in resonance longer than this time. If this is the case, the above discussion would have to be modified to allow the planets to migrate independently from larger radii before converging on to resonance close to their final locations. This is likely to need to be considered for different possible exterior disc models and in addition the planets may have built up their masses as they went (see e.g. Tadeu dos Santos et al. 2015). These considerations are beyond the scope of this paper. None the less, because the migration rates for single planets and the resonantly coupled planets are in general similar, the estimated starting radii would also be similar for disc models that are similar to those we considered. But note that the attained eccentricities depend on the eccentricity damping rates which depend on the details of the disc model (see Crida et al. 2008). For example, we found that for the same amount of relative resonant migration, the entirely inactive disc model led to smaller eccentricities while the 3D layered model led to larger eccentricities. Thus, it is important to note that there is uncertainty as to how long the planets could have been in resonance. In the same context, we comment that migration in the completely active disc model was slower by a factor of \u223c1.6 compared to the standard case on account of its lower mass, that being determined so as to maintain the same steady-state accretion rate as in the standard case. Furthermore, the potential importance of a residual inner gaseous disc for damping the eccentricity of the inner planet and so preventing the eccentricities of both planets from continuing to increase in the later stages of the orbital evolution has been stressed by Crida et al. (2008). In addition, Murray, Paskowitz & Holman (2002) indicate that a residual disc of planetesimals could produce a similar effect.","Citation Text":["Tadeu dos Santos et al. 2015"],"Functions Text":["This is likely to need to be considered for different possible exterior disc models and in addition the planets may have built up their masses as they went (see e.g.","These considerations are beyond the scope of this paper."],"Functions Label":["Future Work","Future Work"],"Citation Start End":[[669,697]],"Functions Start End":[[503,668],[700,756]]}
{"Identifier":"2016ApJ...827...75L__Nishizawa_&_Nakamura_2014_Instance_1","Paragraph":"According to general relativity, in the limit in which the wavelength of gravitational waves is small compared to the radius of curvature of the background spacetime, the waves propagate with the velocity of light, i.e., c (see Will 1998, and references therein). In other theories, the speed \n\n\n\n\n\n could differ from c. Let us define the parameter\n4\n\n\n\n\n\nIf the gravitational wave velocity is subluminal (i.e., \n\n\n\n\n\n), then cosmic rays lose their energy via gravitational Cerenkov radiation significantly. The detection of ultrahigh-energy cosmic rays thus imposes a stringent constraint \n\n\n\n\n\n (i.e., the \u201csubluminal constraint\u201d), depending on the Galactic or extragalactic origin of such particles (Caves 1980; Moore & Nelson 2001). However, there is no theoretical argument (or pathology) against GWs propagating faster than light (see Nishizawa & Nakamura 2014; Blas et al. 2016, and references therein) and the weak bounds from radiation damping in binary systems are \n\n\n\n\n\n (Yagi et al. 2014). The time lag of arrival times between the GW and the simultaneously radiated photons is\n5\n\n\n\n\n\nwhere \n\n\n\n\n\n is the differential distance the photons have traveled. Note that in this work we take the flat cosmological model (i.e., \n\n\n\n\n\n), \n\n\n\n\n\n, and \n\n\n\n\n\n is Hubble\u2019s constant (Riess et al. 2011; Ade et al. 2014). In general, \u03c2 may be a function of the GW frequency (f) and especially when graviton mass is non-zero (i.e., \n\n\n\n\n\n), which gives \n\n\n\n\n\n, where h is Planck\u2019s constant. For simplicity we assume a constant \u03c2 and focus on the association of GWs with electromagnetic counterparts at redshifts \n\n\n\n\n\n. Hence Equation (5) yields (see also Will 1998; Nishizawa & Nakamura 2014)\n6\n\n\n\n\n\nIn reality the photons and the coalescence are not usually simultaneous and we have \n\n\n\n\n\n, where \n\n\n\n\n\n and \n\n\n\n\n\n are the differences in arrival time and emission time, respectively, of the GW and the photons. In most cases, it is rather hard to get an a priori value for \n\n\n\n\n\n. Assuming \n\n\n\n\n\n (i.e., the GW and the electromagnetic counterparts were emitted simultaneously; see Nishizawa (2016) for a more general discussion), we constrain the absolute amplitude of \u03c2 as\n7\n\n\n\n\n\nWe call the above process the \u201ccanonical approach\u201d to measuring the GW velocity directly, in which the graviton and photon are assumed to make the same journey (i.e., EEP is guaranteed). The advantage is that as long as a GW\/GRB association is established one can constrain \n\n\n\n\n\n directly. In Section 4.2 we outline an approach to measure the GW velocity with a simultaneous test of EEP.","Citation Text":["Nishizawa & Nakamura 2014","Nishizawa & Nakamura 2014"],"Functions Text":["However, there is no theoretical argument (or pathology) against GWs propagating faster than light (see","and references therein)"],"Functions Label":["Uses","Uses"],"Citation Start End":[[841,866],[1663,1688]],"Functions Start End":[[737,840],[886,909]]}
{"Identifier":"2018MNRAS.476.1835F__Danforth_et_al._2010_Instance_1","Paragraph":"In the classical unified model, blazars constitute a class of active galactic nuclei (AGNs) viewed at small angles from the jet axis (Blandford & Rees 1978; Antonucci 1993; Urry & Padovani 1995). Traditionally, blazars have been further splitted into two subclasses based on the strength of the features present in their optical spectra. While flat-spectrum radio quasars (FSRQs) show emission lines with equivalent width \u22735\u2009\u00c5, in BL Lacertae objects (BL Lacs) the non-thermal synchrotron radiation of the jet completely dominates the optical\/UV emission, ending up in a typical featureless power-law spectrum. This makes the determination of their redshift via the detection of absorption\/emission lines from the nuclear emission and\/or from the host galaxy particularly challenging, even with 8\u201310 m class telescopes (e.g. Sbarufatti et al. 2005a, 2006, 2009; Landoni et al. 2013; Sandrinelli et al. 2013; Shaw et al. 2013; Falomo, Pian & Treves 2014; Pita et al. 2014; Paiano et al. 2016; Rosa-Gonzalez et al. 2017 for a review). In past years, several alternatives have been proposed to constrain the redshift of BL Lac objects, including the detection of intervening absorption features either from the halo of lower redshift galaxies (e.g. Shaw et al. 2013; Landoni et al. 2014) or from the neutral hydrogen in the intergalactic medium (e.g. Danforth et al. 2010; Furniss et al. 2013); the spectroscopy of galaxies in the environment where the blazars are embedded (e.g. Muriel et al. 2015; Farina et al. 2016); the detection of molecular emission lines from the host galaxy (e.g. Fumagalli et al. 2012); the study of the effect of the interaction with the extragalactic background light (EBL) in the blazar emission in the GeV and TeV domain (e.g. Prandini et al. 2010; Prandini, Bonnoli & Tavecchio 2012). In particular, the narrow distribution in luminosity of BL Lac host galaxies (Urry et al. 2000; Sbarufatti, Treves & Falomo 2005b) opened the possibility to use them as standard candles, and thus to measure their distance via broad-band imaging (e.g. Nilsson et al. 2008; Meisner & Romani 2010; Kotilainen et al. 2011). The main challenge of this approach is to accurately remove the bright central emission that typically outshine the host galaxy. Given the average 1.0\u2009arcsec (or 3.2\u2009kpc in the considered cosmology) effective radius calculated from the collection of z \u2272 0.6 BL Lac hosts observed with HST by Scarpa et al. (2000), it is clear that images with an exquisite spatial resolution and high contrast are necessary to unveil the faint and diffuse starlight emission around the bright, point-like emission from the active nucleus. In this paper, we exploit the capabilities of the new Advanced Rayleigh guided Ground layer adaptive Optics System (ARGOS; Rabien et al. 2010) mounted on the Large Binocular Telescope (LBT; Hill & Salinari 2004; Hill et al. 2012) to collect high-resolution near-infrared (NIR) LUCI\u20091 (i.e. LBT Utility Camera in the Infrared; Seifert et al. 2003; Ageorges et al. 2010) observations of HESS J1943+213. This blazar was detected by HESS in the very high energy (VHE) domain (i.e. at E >100\u2009GeV) during a VHE galactic survey (H.E.S.S. Collaboration et al. 2011), making it the only BL Lac object known located in the Galactic plane. Broad Ks-band images gathered with the 3.5\u2009m CAHA telescope revealed the presence of an extended emission that has been attributed to the host galaxy of HESS J1943+21 (Peter et al. 2014). A comparison with the typical size of blazar host derived by Cheung et al. (2003) allowed Peter et al. (2014) to set a lower limit on the redshift of z > 0.03. This is consistent with the z > 0.14 derived from the fit of the spectral energy distribution of HESS J1943+21 (Cerruti 2011; H.E.S.S. Collaboration et al. 2011) and with the z  0.45 limit obtained via modelling the attenuation of the VHE emission by the EBL (Peter et al. 2014). A tighter constraint on the redshift is however necessary to understand the nature of the VHE emission and to derive the EBL properties.","Citation Text":["Danforth et al. 2010"],"Functions Text":["In past years, several alternatives have been proposed to constrain the redshift of BL Lac objects, including","or from the neutral hydrogen in the intergalactic medium (e.g."],"Functions Label":["Background","Background"],"Citation Start End":[[1348,1368]],"Functions Start End":[[1033,1142],[1285,1347]]}
{"Identifier":"2022AandA...663A..50B__Seibert_et_al._2005_Instance_1","Paragraph":"In the absence of dust, the spectral emission of a normal star-forming galaxy is dominated by stellar populations of different ages with superimposed nebular emission, mainly in the form of recombination lines as well as continuum. The interaction with dust has a dramatic effect, both dimming and reddening the emission from stars and ionized gas. This negatively impacts our ability to measure star formation as energetic photons produced by massive young stars are far more easily attenuated than longer wavelength photons, and even a small quantity of dust can lead to a significant attenuation in the ultraviolet (UV). In the case of particularly dust-rich galaxies, it can render their detection in the rest-frame UV especially difficult. However, as the FUV emission vanishes due to dust attenuation, this dust re-emits the absorbed energy in the mid-infrared (MIR) and far-infrared (FIR), which can in turn be exploited to trace star formation. Except for the most extreme cases (e.g., when the dust content is negligible or, conversely, when almost all of the UV photons are absorbed by dust), an attenuation correction must be carried out to retrieve the star formation rate (SFR). One of the most direct ways is to simply apply a hybrid SFR estimator combining the rest-frame UV with the IR (e.g., Hao et al. 2011; Boquien et al. 2016). The obvious downside is that this requires observations of the dust emission that are costly and difficult to obtain, and even more so at increasing redshifts, where they tend to be limited to vanishingly small samples. With the rest-frame UV emission being relatively easy to obtain from the ground from z\u2004\u223c\u20042 and beyond, techniques have been developed to relate the UV slope (\u03b2) to the UV attenuation (the IRX-\u03b2 relation). While this approach initially appeared to work remarkably well in the case of starburst galaxies (Meurer et al. 1999), there is now ample evidence that there is no tight universal relation between the UV slope and the attenuation (e.g., Buat et al. 2005; Seibert et al. 2005; Howell et al. 2010; Casey et al. 2014). In fact, this relation relies on two strong underlying assumptions: the intrinsic UV slope of the stellar populations in the absence of dust and the exact shape of the attenuation curve. Numerous studies have analyzed their respective impact in an attempt to understand why and when such relations fail and build more reliable ones (e.g., Kong et al. 2004; Boquien et al. 2009, 2012; Popping et al. 2017, and many others). In particular, the recent study of Salim & Boquien (2019) found that the diversity of attenuation curves is a strong driver of the scatter around the IRX-\u03b2 relation. This finding, which is consistent with simulations (Narayanan et al. 2018b; Liang et al. 2021), is especially important in that we can observe a broad variety of attenuation curves at all redshifts (e.g., Salmon et al. 2016; Buat et al. 2018; Salim et al. 2018). With the shape of the attenuation curve being strongly dependent on the relative geometry of stars, ionized gas, and dust (Salim & Narayanan 2020), from the disturbed morphologies observed at higher redshifts, we can only expect important variations there as well (e.g., Faisst et al. 2017). However, due to the great difficulty in measuring them and given the sparsity of the data available, our knowledge of attenuation curves beyond z\u2004=\u20044 remains limited. In effect, most observational studies on the attenuation properties of distant galaxies tend to concentrate on redshifts between 2 and 4 (e.g., Noll et al. 2009b; Buat et al. 2012, 2019; Reddy et al. 2012, 2015; Shivaei et al. 2015; \u00c1lvarez-M\u00e1rquez et al. 2016; Salmon et al. 2016; Fudamoto et al. 2017, 2020b; Lo Faro et al. 2017; \u00c1lvarez-M\u00e1rquez et al. 2019; Reddy et al. 2018; Koprowski et al. 2020). There is only a handful of examples at higher redshift (Capak et al. 2015; Scoville et al. 2015; Bouwens et al. 2016; Barisic et al. 2017; Koprowski et al. 2018). Because of the inherent limits of the observations, studies based on numerical simulations of galaxies at very high redshift (e.g., Mancini et al. 2016; Cullen et al. 2017; Di Mascia et al. 2021) are an important source of information. However, they lead to contrasted results, finding both flat (Cullen et al. 2017) and steep (Mancini et al. 2016) attenuation curves.","Citation Text":["Seibert et al. 2005"],"Functions Text":["While this approach initially appeared to work remarkably well in the case of starburst galaxies","there is now ample evidence that there is no tight universal relation between the UV slope and the"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[2028,2047]],"Functions Start End":[[1773,1869],[1892,1990]]}
{"Identifier":"2021MNRAS.501.4148L__Hatzes_2014_Instance_1","Paragraph":"If the separation between the period of the planet and all the other periodic signals is large enough, and the RV signal has a similar or larger semi-amplitude, it is possible to determine the RV semi-amplitude for a USP planet without any assumptions about the number of planets in the system or the activity of the host star. Under such conditions, during a single night, the influence of any other signal is much smaller than the measurement error and, thus, it can be neglected. If two or more observations are gathered during the same night and they span a large fraction of the orbital phase, the RV semi-amplitude of the USP planet can be precisely measured by just applying nightly offsets to remove all the other signals (e.g. Hatzes et al. 2010; Howard et al. 2013; Pepe et al. 2013; Frustagli et al. 2020 for a recent example). Such an approach, also known as floating chunk offset method (FCO; Hatzes 2014), has proven extremely reliable even in the presence of complex activity signals, as shown by Malavolta et al. (2018). In our case, the shortest, next periodic signal (i.\u2009e. TOI-561 c at 10.78\u2009d) is \u224324 times the period of TOI-561 b (i.\u2009e. the USP planet at 0.45\u2009d), with similar predicted RV semi-amplitude, making this target suitable for the FCO approach. Thanks to our observational strategy (see Section 2.2), we could use 10 different nights for this analysis. Most notably, during two nights, we managed to gather six observations spanning nearly 5\u2009h, i.\u2009e. more than 40\u2009per\u2009cent of the orbital period of TOI-561 b, at opposite orbital phases, thus, providing a good coverage in phase of the RV curve. We did not include RV measurements with an associated error greater than 2.5\u2009m\u2009s\u22121 (see Appendix B1). We performed the analysis with PyORBIT as specified in Section 5, assuming a circular orbit for the USP planet and including a RV jitter as a free parameter to take into account possible short-term stellar variability and any underestimation of the error bars. From our analysis, we obtained a RV semi-amplitude of Kp = 1.80 \u00b1 0.38\u2009m\u2009s\u22121, corresponding to a mass of Mp = 1.83 \u00b1 0.39\u2009M\u2295. The resulting RV jitter is j  0.9\u2009m\u2009s\u22121(84.13th percentile of the posterior). We show the phase folded RVs of the USP planet in Fig. 9. Since the greater reliability of this method over a full fit of the RV data set is counter-balanced by the smaller number of RVs, we decided not privilege one over the other. Therefore, we assumed as final semi-amplitude and mass of TOI-561 b the weighted mean of the values obtained from the two methods (FCO approach and joint photometric and RV fit), i.\u2009e. Kb = 1.56 \u00b1 0.35\u2009m\u2009s\u22121, corresponding to a mass of Mb = 1.59 \u00b1 0.36\u2009M\u2295. Table 5 lists the above-mentioned values for TOI-561 b.","Citation Text":["Hatzes 2014"],"Functions Text":["Such an approach, also known as floating chunk offset method (FCO;","has proven extremely reliable even in the presence of complex activity signals, as shown by Malavolta et al. (2018)."],"Functions Label":["Uses","Compare\/Contrast"],"Citation Start End":[[906,917]],"Functions Start End":[[839,905],[920,1036]]}
{"Identifier":"2016ApJ...832...57P__Vasquez_&_Markovskii_2012_Instance_1","Paragraph":"We employ two types of kinetic codes, hybrid particle-in-cell (PIC) and full PIC simulations. Both types make use of the P3D family of codes (Zeiler et al. 2002), in hybrid PIC (e.g., Parashar et al. 2011) mode, and fully kinetic PIC mode (e.g., Wu et al. 2013b). All simulations discussed here are performed in the 2.5D geometry (two-dimensional (2D) grid and all three components of field vectors). The hybrid simulation has \n\n\n\n\n\n (where \n\n\n\n\n\n is the ion inertial length, with c the speed of light and \n\n\n\n\n\n the proton plasma frequency), \n\n\n\n\n\n, 200 particles per cell, \n\n\n\n\n\n, cold isothermal electrons with \n\n\n\n\n\n. The simulation is initialized with energy only in wavevectors \n\n\n\n\n\n that have \n\n\n\n\n\n. v and b fluctuations are chosen with a specified initial spectral shape, Gaussian random phases, and only in essentially incompressive modes of the system. This simulation was also used in a recent study of variance anisotropy in kinetic plasmas (Parashar et al. 2016). The first full PIC simulation has \n\n\n\n\n\n, \n\n\n\n\n\n, 200 particles per cell, \n\n\n\n\n\n, \n\n\n\n\n\n. The initial condition is the Orszag\u2013Tang vortex (OTV) (e.g., Orszag & Tang 1979; Dahlburg & Picone 1989; Parashar et al. 2009; Vasquez & Markovskii 2012). This simulation was performed for a recent study of transition from kinetic to MHD-like behavior (Parashar et al. 2015). The final PIC simulation (Turb812) has \n\n\n\n\n\n, \n\n\n\n\n\n, 400 particles per cell, \n\n\n\n\n\n, \n\n\n\n\n\n. The initial condition is MHD-like, and more \u201cturbulent,\u201d with v and b fluctuations excited in a band of wave-vectors with \n\n\n\n\n\n with a specified initial spectrum. This simulation was done as part of a recent study that discussed the relation of timescales at the proton gyroscale and their relation to relative proton\u2013electron heating (Matthaeus et al. 2016). PIC codes have an inherent noise associated with them due to the finite number of particles per cell. While performing these simulations, the two most important numerical criteria that we paid attention to were: (i) excellent conservation of total energy (less than a few percent change in any fluctuation energy), and (ii) the particle noise in the spectrum was significant only at scales much smaller than the scales of interest (Debye length \n\n\n\n\n\n for PIC and di for hybrid PIC). On this basis, the modest number of particles employed here was considered adequate. As an additional measure, we employed filtering (e.g., Wan et al. 2012) to remove particle noise at grid scales prior to computing gradients (e.g., vorticity).","Citation Text":["Vasquez & Markovskii 2012"],"Functions Text":["The initial condition is the Orszag\u2013Tang vortex (OTV) (e.g.,"],"Functions Label":["Uses"],"Citation Start End":[[1196,1221]],"Functions Start End":[[1069,1129]]}
{"Identifier":"2016ApJ...819...51L__Svensson_et_al._2010_Instance_1","Paragraph":"Further diagnostics are clearly needed to form firm conclusions. There are likely to be three routes through which this can come. The first is via spectroscopy of the bumps in any further examples. High quality spectroscopy, allied to detailed modeling can yield diagnostics even in the case of relatively weak or featureless spectra, as recently demonstrated in the case of the ultra-long and luminous supernovae pairing GRB 111209A\/SN2011kl (Greiner et al. 2015). The unique identification of features expected in luminous SNe (e.g., turn-off due to line blanketing, absorption lines seen in SLSNe) or TDFs (e.g., blueshifted narrow lines from streams (Strubbe & Quataert 2011)) would then provide a clinching argument as to the origin of the bumps in the longest high-energy transients. A second route arises through studying the locations of the transients within their hosts. Swift J1644+57 clearly arises very close to the galactic nucleus, and Swift J2058+0516 is also consistent with the nucleus of a much fainter galaxy (Pasham et al. 2015). In the case of GRBs, approximately 1\/6 of examples are consistent with a galaxy nucleus (Fruchter et al. 2006; Svensson et al. 2010); this number may be lower for SLSNe (Lunnan et al. 2015) although the origin of SLSNe in the nuclei of galaxies may be ambiguous (e.g., Dong et al. 2016). Further examples, all in the nuclei of their hosts, would rapidly remove any SNe model from consideration. Finally, we can also consider the host galaxy more globally. TDFs can be observed in quiescent, non-star-forming galaxies while SLSNe are thought to arise from massive star collapse (Gal-Yam 2012) and in principle should occur only in star forming galaxies. A prime model for SLSNe is that they arise from supernovae in which the shock wave is re-energized by the spin-down energy of a recently formed magnetar (Kasen & Bildsten 2010). While magnetars similar to those suggested to power SLSNe can be formed via accretion induced collapse of two merging white dwarfs (Usov 1992; Levan et al. 2006), and may provide a similar energy input, in the case of a white dwarf merger there would be minimal remnant to re-energize, and hence no luminous SNe. This means that the presence of an extremely long event within an quiescent elliptical galaxy would rule out SNe models, and strongly favor an origin as a relativistic tidal flare. Since a reasonable fraction (\u223c50%) of candidate tidal disruptions arise from passive systems, (i.e., those with little sign of star formation) (e.g., Arcavi et al. 2014) such a test should be possible with only a handful of additional examples since we would expect to observe an example in a system without star formation in the near future.","Citation Text":["Svensson et al. 2010"],"Functions Text":["In the case of GRBs, approximately 1\/6 of examples are consistent with a galaxy nucleus"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1162,1182]],"Functions Start End":[[1051,1138]]}
{"Identifier":"2017MNRAS.468.2206M__Ceplecha_(1951)_Instance_1","Paragraph":"The first identification of a possible meteor shower with a radiant point in the constellation of Ursa minor was by Denning (1912), occurring around the winter solstice. However, very few observations were subsequently made. Hoffmeister recorded its activity in 1914, but its existence was not confirmed until 1945 (Becvr 1946). There are two reasons for the lack of records of the Ursids. First, the Ursids have very low activity levels in most years, its Zenith hourly rate (the number of meteors that would be observed under good observing conditions in one hour if the radiant was at the zenith) is usually ZHR 10 (Jenniskens 1994). Secondly, the weather can be bad in mid-December and many observers choose to only observe the more predictable Geminids. Ceplecha (1951) showed that the annual Ursid shower was related to comet 8P\/Tuttle, a Jupiter family comet with a period of 13.6yr (JPL-HORIZONS, ssd.jpl.nasa.gov) and a perihelion distance slightly greater than 1au. This means that the comet is usually brightest (close to perihelion) and close to the Earth at roughly the same time. Hence, it has been observed at all perihelion passages since discovery apart from 1953, when observing conditions were poor throughout. These perihelion passages were in 1858, 1872, 1885, 1899, 1913, 1926, 1940, 1953, 1967, 1980, 1994 and 2008. As it can take a considerable time after being ejected from the nucleus for meteoroids to disperse away from the nucleus locality, an enhancement is generally to be expected in stream activity at the time when the comet is close to perihelion and new meteoroids are injected (see Williams, Johnson  Fox 1986 for an early discussion and mathematical formulation of this). An enhancement at some such times has also been observed in the Ursids, where the ZHR reaches around three times the normal rate (e.g. 1900, 1914, 1953, 1981 and 1996, data taken from www.meteorshowersonline.com). However, in the case of the Ursids, several perihelion passage years do not show any significant enhancement and so this is unlikely to be the explanation. Instead, Jenniskens etal. (2007) suggested that they are caused by cometary material released at very old comet perihelion passages between ad300ad1400 and appear as a wide and stretched stream (which we will call and refer to as a filament).","Citation Text":["Ceplecha (1951)"],"Functions Text":["showed that the annual Ursid shower was related to comet 8P\/Tuttle, a Jupiter family comet with a period of 13.6yr (JPL-HORIZONS, ssd.jpl.nasa.gov) and a perihelion distance slightly greater than 1au."],"Functions Label":["Background"],"Citation Start End":[[759,774]],"Functions Start End":[[775,975]]}
{"Identifier":"2019ApJ...879...52S__Wei_et_al._2010_Instance_1","Paragraph":"The lack of data at complementary wavelengths also makes resolved multiwavelength analyses applied to low-redshift galaxies, such as the Schmidt\u2013Kennicutt relation (the correlation between galaxies\u2019 SFR and gas mass surface densities; e.g., Schmidt 1959; Kennicutt 1989, 1998), significantly less common at high redshift. High-resolution CO observations are critical for evaluating where high-redshift galaxies fall on the true surface density version of the Schmidt\u2013Kennicutt relation, where \u03a3SFR and \u03a3gas can be compared on a pixel-by-pixel basis within individual galaxies (as done for local galaxies; e.g., Kennicutt et al. 2007; Bigiel et al. 2008, 2011; Wei et al. 2010; Leroy et al. 2013). Many high-redshift analyses use star formation and gas properties averaged over the entire galaxy (e.g., Buat et al. 1989; Kennicutt 1989, 1998; Daddi et al. 2010b; Genzel et al. 2010; Tacconi et al. 2013) or avoid the additional uncertainties in source size and scaling factors by using the total luminosities of the star formation and gas tracers (e.g., Young et al. 1986; Solomon & Sage 1988; Gao & Solomon 2004). These different methods for determining SFRs and gas masses make it difficult to compare studies that focus on different galaxy populations, leading to significant uncertainties in the power-law index of the Schmidt\u2013Kennicutt relation and the relative placement of different galaxy types in the \u03a3SFR\u2013\u03a3gas plane. Accurately characterizing the Schmidt\u2013Kennicutt relation is important, since offsets imply a difference in star formation efficiency (SFE), and the power-law index probes the underlying physical processes of star formation (for example, a linear correlation would imply supply-limited star formation, whereas superlinear correlations occur if star formation depends on cloud\u2013cloud collisions or total gas freefall collapse times; e.g., Tan 2000; Krumholz & McKee 2005; Ostriker & Shetty 2011). Systematic differences in the Schmidt\u2013Kennicutt relation between different galaxy populations would imply important differences in their star formation processes.","Citation Text":["Wei et al. 2010"],"Functions Text":["High-resolution CO observations are critical for evaluating where high-redshift galaxies fall on the true surface density version of the Schmidt\u2013Kennicutt relation, where \u03a3SFR and \u03a3gas can be compared on a pixel-by-pixel basis within individual galaxies (as done for local galaxies; e.g.,"],"Functions Label":["Motivation"],"Citation Start End":[[660,675]],"Functions Start End":[[322,610]]}
{"Identifier":"2022AandA...665A..68M__Kempen_et_al._2009_Instance_1","Paragraph":"Within a star-forming cloud core, the protostar is the main source of luminosity and heat due to the release of gravitational energy from contraction and material accretion. The amount of protostellar heating dictates the temperature structure of the cloud core. In an idealized spherical scenario, the temperature alone would dictate the snowline locations within the cloud core. However, star-forming cloud cores are not spherically symmetric. The outflow cavity, flattened structures around the protostar (e.g., pseudo-disks and rotationally supported disks; hereafter referred to as disks for simplicity), and variations within the envelope density can all impact how heat is distributed within the cloud core. Studies have shown that heating mainly escapes through the outflow cavity in deeply embedded sources (van Kempen et al. 2009; Y\u0131ld\u0131z et al. 2015; Murillo et al. 2018a). Thus the extent of chemical richness in the outflow cavity provides insight into the luminosity of the protostar and the physical conditions of the envelope (e.g., Drozdovskaya et al. 2015; Murillo et al. 2018b; Tychoniec et al. 2019, 2020). Observations of embedded protostars (so-called Class 0 and I systems) have shown the presence of disks, both as flattened dust continuum structures (e.g., J\u2205rgensen et al. 2009; Enoch et al. 2011; Persson et al. 2016; Segura-Cox et al. 2018; Tobin et al. 2020) and rotationally supported disks traced in molecular gas (e.g., Murillo et al. 2013; Harsono et al. 2014; Yen et al. 2015, 2017; Maret et al. 2020). These disks show a wide range of geometries and radii ranging from a few 10 AU up to ~200 AU. Additional studies have shown that the presence of a disk can alter the temperature profile along the equator (disk mid-plane) of the cloud core (e.g., Murillo et al. 2015, 2018b; van\u2019t Hoff et al. 2018b; Hsieh et al. 2019a). Multiplicity, that is two or more protostars within a single cloud core, can produce further asymmetries due to differences in luminosity from the multiple components and their locations with respect to each other (e.g., Chen et al. 2009; Koumpia et al. 2016; Murillo et al. 2016, 2018b). At an early evolutionary stage, protostellar luminosity is dominated by accretion, that is accretion luminosity (Hartmann & Kenyon 1996). Variability in protostellar luminosity has been detected toward several targets (V1647 Ori: \u00c1brah\u00e1m et al. 2004; Andrews et al. 2004; Acosta-Pulido et al. 2007; Fedele et al. 2007; Aspin & Reipurth 2009; OO Serpentis: K\u00f3sp\u00e1l et al. 2007; CTF93 216-2 Caratti o Garatti et al. 2011; VSX J205126.1: Covey et al. 2011; K\u00f3sp\u00e1l et al. 2011; HOPS383 Safron et al. 2015; S255IR-SMA1 Safron et al. 2015; Liu et al. 2018; EC53 Herczeg et al. 2017; Yoo et al. 2017). Such variability is considered to be a product of the nonuniform accretion of material with a variable amount and frequency onto the protostar, that is to say episodic accretion (Audard et al. 2014).","Citation Text":["van Kempen et al. 2009"],"Functions Text":["Studies have shown that heating mainly escapes through the outflow cavity in deeply embedded sources"],"Functions Label":["Background"],"Citation Start End":[[817,839]],"Functions Start End":[[715,815]]}
{"Identifier":"2022MNRAS.515.2188H__Rorai_et_al._2018_Instance_3","Paragraph":"In this work, we follow the method for measuring the IGM thermal state based on Voigt profile decomposition of the Ly\u2009\u03b1 forest (Schaye et al. 1999; Ricotti et al. 2000; McDonald et al. 2001). In this approach, a transmission spectrum is treated as a superposition of multiple discrete Voigt profiles, with each line described by three parameters: redshift zabs, Doppler broadening b, and neutral hydrogen column density $N_{\\rm H\\, {\\small I}}$. By studying the statistical properties of these parameters, i.e. the $b-N_{{{{\\rm H\\, {\\small I}}}}{}}$ distribution, one can recover the thermal information encoded in the absorption profiles. The majority of past applications of this method constrained the IGM thermal state by fitting the low-b-$N_{\\rm H\\, {\\small I}}$ cutoff of the $b-N_{{{{\\rm H\\, {\\small I}}}}{}}$ distribution (Schaye et al. 1999, 2000; Ricotti et al. 2000; McDonald et al. 2001; Rudie, Steidel & Pettini 2012; Boera et al. 2014; Bolton et al. 2014; Garzilli, Theuns & Schaye 2015, 2020; Hiss et al. 2018; Rorai et al. 2018). The motivation for this approach is that the Ly\u2009\u03b1 lines are broadened by both thermal motion and non-thermal broadening resulting from combinations of Hubble flow, peculiar velocities, and turbulence. By isolating the narrow lines in the Ly\u2009\u03b1 forest that constitutes the lower cutoff in $b-N_{{{{\\rm H\\, {\\small I}}}}{}}$ distributions, of which the line-of-sight component of non-thermal broadening is expected to be zero, the broadening should be purely thermal, thus allowing one to constrain the IGM thermal state. However, this method has three crucial drawbacks. First, the IGM thermal state actually impacts all the lines besides just the narrowest lines. Therefore, by restricting attention to data in the distribution outskirts, this approach throws away information and reduces the sensitivity to the IGM thermal state significantly (Rorai et al. 2018; Hiss et al. 2019). Secondly, in practice, determining the location of the cutoff is vulnerable to systematic effects, such as contamination from the narrow metal lines (Hiss et al. 2018; Rorai et al. 2018). Lastly, the results from this approach critically depend on the choice of low-b cutoff fitting techniques, where different techniques might result in inconsistent T0 and \u03b3 measurements (Hiss et al. 2018; Rorai et al. 2018).","Citation Text":["Rorai et al. 2018"],"Functions Text":["Secondly, in practice, determining the location of the cutoff is vulnerable to systematic effects, such as contamination from the narrow metal lines"],"Functions Label":["Background"],"Citation Start End":[[2097,2114]],"Functions Start End":[[1929,2077]]}
{"Identifier":"2022AandA...666A..44K__Martig_et_al._2009_Instance_1","Paragraph":"Meanwhile, the quiescence of the stellar disks (albeit possibly not entirely passive), is the first indication of the presence of an apparent \u201coutside-in\u201d quenching mode in high-redshift massive galaxy populations, in contradiction to the classic \u201cinside-out\u201d configuration (Lang et al. 2014; Tacchella et al. 2015; Breda & Papaderos 2018). We consider various possible modes of quenching. Since we observe the galaxies to be undergoing an outside-in quenching, this could not have primarily occurred due to feedback from active galactic nuclei (AGN; Alatalo et al. 2015), known to quench galaxies inside-out (Tacchella et al. 2018). We also find no evidence of AGN activity from X-ray observations (Daddi et al. 2021), although a radio excess in Galaxy-C is indicative of weak past AGN activity in that galaxy. The process of morphological quenching (Martig et al. 2009), where the formation of a stellar bulge stabilizes the disk against further star-formation, is also improbable. The galaxies are far from being bulge-dominated, essential for this mode of quenching to be applicable, based on their stellar mass distributions. Also unlikely is cosmological starvation (Feldmann & Mayer 2015) since the availability of gas has already been established in RO-1001. Finally, ram-pressure stripping (RPS; Gunn & Gott 1972) could be regarded as a possible contributor as it is known to remove gas from external regions of galaxies (Bravo-Alfaro et al. 2000; Vollmer et al. 2001; Fumagalli et al. 2009; Boselli et al. 2014; Loni et al. 2021), therefore resulting in an apparent outside-in quenching (for a review, Boselli et al. 2022). It could also lead to compression of the gas in the galaxy which could result in the lopsidedness (although the stellar morphology would not be affected). However, as shown in Appendix B, a conservative lower-limit of the radius up to which RPS can remove gas from the sample in RO-1001 would be 10\u201315 kpc. This is more than an order of magnitude higher than the kpc-scale cores beyond which the galaxies have suppressed star-formation. Hence we do not expect RPS to be playing a major role.","Citation Text":["Martig et al. 2009"],"Functions Text":["The process of morphological quenching","where the formation of a stellar bulge stabilizes the disk against further star-formation, is also improbable."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[852,870]],"Functions Start End":[[812,850],[873,983]]}
{"Identifier":"2021ApJ...921...19W__Arcones_&_Montes_2011_Instance_1","Paragraph":"The abundances for cold and hot tracers with different Ye treatments are shown in Figure 12 for model s15_F120_R006. In general, the abundances for cold and hot tracers show a clear iron peak. The cold component is characterized by the odd\u2013even distribution from carbon to calcium, following the progenitor composition. The hot component reaches heavier elements than the cold one with a slight dependency of the abundances on the exact electron fraction. The different assumptions for the initial Ye lead to some variability for the abundances beyond iron covering all expected possibilities for yields from neutrino-driven supernovae. The production of elements in the region of Sr, Y, and Zr is more efficient for slightly neutron-rich conditions, corresponding to RLn = 1.35 (Arcones & Montes 2011; Arcones & Bliss 2014). We observe a very similar behavior and dependency of abundances on the Ye distribution for all exploding models of the 15 M\u2299 progenitor. As a summary, we show in Figure 13 the final abundances only for the narrow Ye distribution with RLn = 1 for all models. We do not find any clear correlation between the abundances and the final explosion energy, although we see slight dependency on the explosion energy at \u223c1 s after bounce. The lack of correlation is probably due to our simple treatment of the neutrinos and the correction of the electron fraction, as well as to the fact that we are calculating abundances only for the 15 M\u2299 progenitor. Simulations with detailed neutrino transport and for several progenitors are necessary to narrow the uncertainties in the abundances and to link those to other astrophysical conditions. Even if our results are not completely conclusive for abundances beyond iron, they indicate that there is not a strong variability for iron group elements. This may be important to estimate uncertainties in the production of 56Ni and 44Ti as shown in Figure 14 where we present an overview of all models including cold and hot components and variations of Ye.","Citation Text":["Arcones & Montes 2011"],"Functions Text":["The production of elements in the region of Sr, Y, and Zr is more efficient for slightly neutron-rich conditions, corresponding to RLn = 1.35","We observe a very similar behavior and dependency of abundances on the Ye distribution for all exploding models of the 15 M\u2299 progenitor."],"Functions Label":["Uses","Similarities"],"Citation Start End":[[780,801]],"Functions Start End":[[637,778],[826,962]]}
{"Identifier":"2022MNRAS.513L..78D__Marcos_2020_Instance_1","Paragraph":"The fields targeted by our proof-of-concept mini-survey were observed at relatively high airmasses (see Section 2). Therefore, it can be argued that our null result could be due to enhanced extinction caused by observing at such low elevations. However, this is likely not the case because the same EURONEAR collaboration carried out a nearly concurrent mini-survey looking for Atira and Vatira asteroids (see e.g. Greenstreet, Ngo & Gladman 2012) using the same instrumental setup and similar automated reduction pipeline. Atiras or Interior Earth Objects (IEOs) have aphelion distances 0.983 au and can only be observed at low solar elongations (often 70\u25cb); Vatiras have aphelia 0.718 au, are observed at very low solar elongations (typically 40\u25cb\u201345\u25cb), and just one such object is currently known (Bolin et al. 2020; de la Fuente Marcos & de la Fuente Marcos 2020; Greenstreet 2020), 594913 \u2018Ayloi\u2019chaxnim (2020 AV2). Although the nearly concurrent EURONEAR small-scale survey (12 nights with an average time of 1.5 h per night devoted to searching for Vatira and Atira asteroids towards the evening and early morning skies) did not find any new Vatiras or IEOs, it did observe at elevations in the range 15\u25cb\u201330\u25cb and recovered known objects at apparent magnitudes close to or above 23 \u2013 for example 2017 HO20, a main belt asteroid that was observed near conjunction. Furthermore, Fig. 2, top panel, shows the r\u2032-magnitude distribution of the detected stars from one representative pointing out of the 77 (\u00d75) pointings intended to recover the TESS candidate in Cepheus. The histogram was produced using the matplotlib library (Hunter 2007) with sets of bins computed using numpy (van der Walt, Colbert & Varoquaux 2011; Harris et al. 2020) by applying the Freedman and Diaconis rule (Freedman & Diaconis 1981). Fig. 2, bottom panel, shows the expected S\/N for WHT prime-focus imaging as a function of the exposure time from the exposure time calculator (ETC)11 for a similar camera (Red + 4, PFIP) for observations carried out under a seeing of 0${_{.}^{\\prime\\prime}}$9 at an airmass of 2.3 in dark time and R = 23.2 mag. TNOs have colour index V \u2212 R in the range 0.2\u20131.2 (see e.g. Peixinho, Delsanti & Doressoundiram 2015). Consistently, we are confident that we reached a limiting magnitude r\u2032= 23.0 mag, which actually corresponds to $V\\, \\gt $23.0 mag (in fact, it could be 23.2\u201324.2 mag). Therefore, our distant TNO candidate survey could have detected candidates 9 and 11 in table 2 of Rice & Laughlin (2020) unless both candidates are much fainter than predicted or they are false positives in TESS data. In any case, Rice & Laughlin (2020) stated that many if not most of the high signal significances reported in their table 2 could be the result of unmodelled systematic errors. On the other hand, in the fields observed towards Cepheus, there are a few bright stars that might have hidden a putative moving object with the properties of candidate 9; we estimate that the probability of having missed the candidate as a result of it moving projected towards one of those bright stars is 1 per cent (considering the area affected by the bright stars and associated diffraction spikes).","Citation Text":["de la Fuente Marcos & de la Fuente Marcos 2020"],"Functions Text":["Atiras or Interior Earth Objects (IEOs) have aphelion distances 0.983 au and can only be observed at low solar elongations (often 70\u25cb); Vatiras have aphelia 0.718 au, are observed at very low solar elongations (typically 40\u25cb\u201345\u25cb), and just one such object is currently known"],"Functions Label":["Background"],"Citation Start End":[[819,865]],"Functions Start End":[[524,798]]}
{"Identifier":"2017MNRAS.469.4933L__Friedmann_et_al._1996_Instance_1","Paragraph":"IC 59 has a lower surface brightness than IC 63, which may be the result of both lower column densities as well as a larger physical distance from the illuminating star (Blouin et al. 1997; Karr, Noriega-Crespo & Martin 2005) compared to IC 63. Because of their rare illumination geometry and proximity to Earth, these two nebulae have been studied intensively at wavelengths ranging from the Lyman limit of hydrogen to the hydrogen line and continuum near 21 cm as ideal test cases for PDR models (Perrin & Sivan 1992; Jansen, van Dishoeck & Black 1994; Blouin et al. 1997; Luhman et al. 1997; Hurwitz 1998; Habart et al. 2004; France et al. 2005; Karr et al. 2005; Thi et al. 2009; Fleming et al. 2010; Miao et al. 2010). As a result of these studies, the physical conditions in these nebulae are exceptionally well constrained. The distances of the nebulae from the illuminating star are about 1 pc, assuming that they are located near the plane of the sky occupied by \u03b3 Cas (Friedmann et al. 1996). Thus, the incident radiation field at locations within the two nebulae is essentially unidirectional, with the contribution from \u03b3 Cas about 650 times (Habart et al. 2004) more intense than the average ISRF in the solar vicinity at the front PDR in IC 63 facing \u03b3 Cas. This leads to well-defined PDRs that are viewed essentially edge-on. This is the case particularly in IC 63. Thus, it is possible to observe radiation from different emission processes requiring photons of different energies for their excitation, such as recombination lines from ionized hydrogen and sulphur, fluorescence from photoexcited molecular hydrogen, emissions from ionized and neutral polycyclic aromatic hydrocarbons (PAHs), and dust-scattered continuum at different depths behind the front of the PDR. In this regard, IC 63 is a more ideal geometry for the morphological study of ERE than NGC 7023, where the illuminating star, HD 200775, is centrally embedded and where projection effects make it more difficult to spatially separate different emission processes (Witt et al. 2006; Bern\u00e9 et al. 2008). One of the prime objectives of this study is to determine the spatial distribution of the ERE in IC 59 and IC 63 in relation to other emissions with known excitation requirements. For this analysis, we take advantage of the wavelength dependence of the opacity of the ISM consisting of atomic and molecular gas and interstellar dust in the UV\u2013optical spectral range (Ryter 1996). An added benefit of the special geometry of IC 59 and IC 63 with respect to \u03b3 Cas is that scattering by dust occurring here at large scattering angles near 90\u00b0 is rather inefficient, which makes it easier to separate other UV\/optical emission components from the scattered light, including the ERE. The subsequent chapters of this paper are organized as follows: In Section 2, we present new optical observations of IC 59 and IC 63 and discuss their reductions and absolute intensity calibrations. Additional archival observations obtained from the WISE, Spitzer and Herschel Space Observatories will be introduced as well. In Section 3, we present the analysis of the ERE morphology in the nebulae based on digital image subtraction and division. We will estimate the ERE intensity using the colour-difference technique and determine the wavelength region of ERE excitation. A discussion of the implications for the ERE emission process and specific ERE carrier models will follow in Section 4, with conclusions to be presented in Section 5. Data tables with the results of measurements of the relative surface brightnesses of IC59 and IC63 are presented in the appendices.","Citation Text":["Friedmann et al. 1996"],"Functions Text":["The distances of the nebulae from the illuminating star are about 1 pc, assuming that they are located near the plane of the sky occupied by \u03b3 Cas"],"Functions Label":["Uses"],"Citation Start End":[[979,1000]],"Functions Start End":[[831,977]]}
{"Identifier":"2021MNRAS.501.3046F__Thorstensen_et_al._1991_Instance_1","Paragraph":"Cataclysmic variables (CVs) are strong interacting binaries. They are systems with an accreting white dwarf (WD, the primary) and a red dwarf (the donor). The donor star is on a late-type main sequence, or sometimes in a slightly evolved state. An accretion disc forms and accreted material reaches the WD in the case in which the magnetic field of the primary is too weak to interrupt the accretion flow (B  0.01 MG): such systems are referred to as non-magnetic CVs, characterized by their eruptive behaviour (Warner 1995). Mass and angular momentum transfer from the donor star to the WD drive a rich variety of variability, as well as the evolution of these systems (Knigge, Baraffe & Patterson 2011). Nova-like variables (NLs) and dwarf novae (DNs) are subtypes of non-magnetic CVs: in addition to the more frequent eruption of DNs, there are many differences between them. NLs have high accretion rates (\u223c10\u22128 M\u2299 yr\u22121), so are in a persistent high state, and they gather above the period gap in the period distribution of CVs (Thorstensen et al. 1991); while the accretion rates of DNs are lower than those of NLs by one or two orders of magnitude (\u226410\u22129 M\u2299 yr\u22121), their luminosities only match those of NLs during DN outbursts and their periods are more commonly below the period gap. The possible evolutionary links between the two classes were discussed by Shara et al. (1986) as part of the hibernation scenario, in which systems are NLs immediately after a nova burst and then metamorphose into DNs if their luminosity becomes sufficiently low within several centuries. The hibernation hypothesis also predicts that CVs will experience novae, NLs, and DNs within 104\u2212105 yr and repeat the cycle during their lifetimes. The NL system V1315 Aql was estimated to experience a nova outburst about 500\u20131200 yr earlier (Sahman et al. 2018): this is consistent with nova-induced cycles in hibernation theory. Kovetz, Prialnik & Shara (1988) estimated that the mass-transfer rate declines at a rate of 0.012 mag yr\u22121 for the century following outburst. Despite enhanced interest in long-term observations of post-novae and their final decline inspired by the hibernation hypothesis, analysis of the secular evolution of NLs is not yet well known.","Citation Text":["Thorstensen et al. 1991"],"Functions Text":["NLs have high accretion rates (\u223c10\u22128 M\u2299 yr\u22121), so are in a persistent high state, and they gather above the period gap in the period distribution of CVs"],"Functions Label":["Background"],"Citation Start End":[[1033,1056]],"Functions Start End":[[879,1031]]}
{"Identifier":"2018ApJ...856..136P__Pingel_et_al._2013_Instance_1","Paragraph":"Depending on the specific driver, the characteristics of turbulence will then be imprinted within the ISM mainly as three-dimensional density and velocity fluctuations, and these fluctuations have been traditionally studied via correlation functions such as the spatial power spectrum (SPS) (e.g., Crovisier & Dickey 1983), \u0394-variance (e.g., Stutzki et al. 1998), and structure function (e.g., Padoan et al. 2002; Burkhart et al. 2015b). In particular, the SPS approach has been applied to observations of various Galactic and extragalactic environments (e.g., Plume et al. 2000; Dickey et al. 2001; Elmegreen et al. 2001; Burkhart et al. 2010; Combes et al. 2012; Zhang et al. 2012; Pingel et al. 2013), showing power spectral slopes \u03b2 roughly ranging from \u22122.7 to \u22123.7 depending on the tracers used (e.g., H i, carbon monoxide (CO), and dust). These slopes essentially provide information on the relative amount of structure as a function of spatial scale and can be compared with theoretical models of turbulence (mainly numerical simulations) to characterize turbulence cascade (e.g., Burkhart et al. 2010), to determine the influence of shocks (e.g., Beresnyak et al. 2005), to reveal the injection and dissipation scales of turbulent energy (e.g., Kowal & Lazarian 2007; Federrath & Klessen 2013; Chen et al. 2015), and to trace the evolution of MCs (e.g., Burkhart et al. 2015a). The proximity and abundance of multi-wavelength observations make MCs in the solar neighborhood an ideal laboratory for probing the impact of turbulence on their formation and evolution. In this paper, we focus on the Perseus MC, which is a nearby (\u223c300 pc; e.g., Herbig & Jones 1983; \u010cernis 1990), low-mass (\u223c2 \u00d7 104 M\u2299; e.g., Sancisi et al. 1974; Lada et al. 2010) cloud. Its star formation activities, as well as atomic and molecular gas content, have been extensively examined over the past decade (e.g., Ridge et al. 2006; J\u00f8rgensen et al. 2007; Pineda et al. 2008; Lee et al. 2012, 2014, 2015; Mercimek et al. 2017), revealing that the cloud consists of several individual dark and star-forming regions (e.g., B5, B1, B1E, IC 348, and NGC 1333) and is actively forming low- to intermediate-mass stars (see Bally et al. 2008 for a review).","Citation Text":["Pingel et al. 2013"],"Functions Text":["In particular, the SPS approach has been applied to observations of various Galactic and extragalactic environments (e.g.,","showing power spectral slopes \u03b2 roughly ranging from \u22122.7 to \u22123.7 depending on the tracers used (e.g., H i, carbon monoxide (CO), and dust)."],"Functions Label":["Background","Compare\/Contrast"],"Citation Start End":[[684,702]],"Functions Start End":[[438,560],[705,845]]}
{"Identifier":"2022MNRAS.517..130P__Navarro,_Frenk_&_White_1996_Instance_1","Paragraph":"The catalogue represents a luminosity-complete sample of galaxies with r-band absolute magnitude Mr \u2264 \u221219 in a (300h\u22121Mpc)3 comoving volume at z = 0. The mock contains both central and satellite galaxies, populated in dark matter haloes identified in an N-body simulation having 10243 particles with a flat \u039bCDM WMAP7 cosmology (Komatsu et al. 2011; \u03a9m = 0.276, h = 0.7). To achieve the luminosity completeness threshold of Mr \u2264 \u221219, haloes containing \u226540 particles are considered. Since the halo concentration cannot be reliably measured for haloes with fewer than about \u223c400 particles, we use the method presented by Ramakrishnan, Paranjape & Sheth (2021) to assign concentrations cvir (assuming Navarro, Frenk & White 1996, NFW profiles) conditioned on the mass and local tidal environment of individual haloes.2 The galaxies were populated using a halo occupation distribution (HOD) model and H\u2009i-optical scaling relations calibrated by Paul, Choudhury & Paranjape (2018) and Paul, Pahwa & Paranjape (2019) using luminosity- and colour-dependent clustering measurements from the Sloan Digital Sky Survey (SDSS; York et al. 2000; Zehavi et al. 2011) and H\u2009i-dependent clustering measurements from the Arecibo Legacy Fast ALFA survey (ALFALFA; Giovanelli et al. 2005; Guo et al. 2017). Each galaxy in the mock is assigned absolute magnitudes in SDSS u, g, and r bands (with a threshold Mr \u2264 \u221219 imposed by the SDSS clustering measurements) and a stellar mass m* using a colour-dependent mass-to-light ratio. The H\u2009i-optical scaling relation additionally leads to a fraction ${\\sim}60{{\\ \\rm per\\ cent}}$ of galaxies to be assigned an H\u2009i mass $m_{\\rm{H}\\,{\\small I}}$. PCS21 presented extensive tests of this algorithm. We focus in this work on the population of mock central galaxies containing massive H\u2009i discs, with $m_{\\rm{H}\\,{\\small I}}\\ge 10^{9.7}h^{-2}M_{\\odot }$; the resulting \u223c50\u2009000 such objects in our catalogue form a volume-complete sample of H\u2009i-selected galaxies.","Citation Text":["Navarro, Frenk & White 1996"],"Functions Text":["Since the halo concentration cannot be reliably measured for haloes with fewer than about \u223c400 particles, we use the method presented by Ramakrishnan, Paranjape & Sheth (2021) to assign concentrations cvir (assuming","NFW profiles) conditioned on the mass and local tidal environment of individual haloes."],"Functions Label":["Uses","Uses"],"Citation Start End":[[698,725]],"Functions Start End":[[482,697],[727,814]]}
{"Identifier":"2015ApJ...805..163D__Liang_et_al._2010_Instance_1","Paragraph":"In the field of GRBs, evidence of PFD jets has been collected independently in several directions. First, a prominent thermal emission component as expected in the fireball-internal-shock model (e.g., M\u00e9sz\u00e1ros & Rees 2000) has been seen only in a small fraction of GRBs (e.g., GRB 090902B, Ryde et al. 2010; Zhang et al. 2011). The majority of GRBs either show no evidence of a thermal component or a weak, sub-dominant thermal component (e.g., Abdo et al. 2009; Guiriec et al. 2011; Axelsson et al. 2012). These GRBs require that the GRB central engine is highly magnetized, and jet is still PFD at the emission site (Zhang & Pe\u2019er 2009; Gao & Zhang 2015). Next, strong linear polarization was discovered during the prompt gamma-ray emission phase for some GRBs (Yonetoku et al. 2011, 2012), and during the reverse-shock-dominated early optical afterglow emission phase for some others (Steele et al. 2009; Mundell et al. 2013), which hint at the existence of globally ordered magnetic fields in the jet. Furthermore, strong PeV neutrino emission as predicted by the MFD models has not been observed from GRBs so far (Abbasi et al. 2012), which is consistent with the expectation of the PFD models (Zhang & Kumar 2013). Finally, the MFD internal shock (IS) model for GRBs also suffers some criticisms, such as low energy dissipation efficiency (Kumar 1999; Panaitescu et al. 1999), electron fast cooling (Ghisellini et al. 2000), the electron number excess (Bykov & M\u00e9sz\u00e1ros 1996; Daigne & Mochkovitch 1998; Shen & Zhang 2009), and inconsistency with some empirical (Amati\/Yonetoku) relations (Zhang & M\u00e9sz\u00e1ros 2002; Liang et al. 2010). Zhang & Yan (2011) proposed a novel PFD outflow model named as \u201cthe Internal-Collision-induced MAgnetic Reconnection and Turbulence (ICMART),\u201d which can potentially keep the merits of the IS model but alleviate the criticisms faced by the IS model mentioned above. The main idea of the ICMART model is that the GRB jets are PFD. The Poynting flux is catastrophically discharged at a relatively large distance (e.g., 1015 cm) from the central engine through collision-induced magnetic reconnection. The magnetic energy is converted to particle energy and radiation efficiently, leading to a very high radiation efficiency as demanded by the GRB data (Panaitescu & Kumar 2002; Zhang et al. 2007). A PFD jet has less leptons than the MFD model so that the electron excess problem is avoided. A large emission radius favors a moderately fast cooling, which can account for the right low-energy spectral index observed in GRBs (Uhm & Zhang 2014). It also gives a natural explanation of the seconds-duration of \u201cslow variability component\u201d observed in GRBs (Gao et al. 2012). The rapid \u201cfast variability component\u201d can be interpreted within this scenario as mini jets due to locally Lorentz boosted regions (see also Lyutikov & Blandford 2003; Narayan & Kumar 2009)3\n\n3\nLyutikov & Blandford (2003) and Narayan & Kumar (2009) proposed that GRB variability is a consequence of mini jets due to relativistic outflow from reconnection or relativitic turbulence. There is no simple explanation to the observed slow variability component in these models. Zhang & Yan (2011) attributed the two variability components (slow and fast) as due to central engine activity and mini jets, respectively. Monte Carlo simulations by Zhang & Zhang (2014) showed that the ICMART model can indeed reproduce the observed GRB light curves.\n. It is speculated that turbulent reconnection in a moderately high-\u03c3 flow can give rise to relativistic motion of mini jets within the bulk relativistic motion of the jets.","Citation Text":["Liang et al. 2010"],"Functions Text":["Finally, the MFD internal shock (IS) model for GRBs also suffers some criticisms, such as","and inconsistency with some empirical (Amati\/Yonetoku) relations"],"Functions Label":["Differences","Differences"],"Citation Start End":[[1618,1635]],"Functions Start End":[[1221,1310],[1529,1593]]}
{"Identifier":"2021AandA...655A..72S___2019_Instance_3","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"],"Functions Text":["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;"],"Functions Label":["Background"],"Citation Start End":[[1699,1711],[1718,1722]],"Functions Start End":[[1496,1698]]}
{"Identifier":"2018MNRAS.477.3520L__Abolfathi_et_al._2018_Instance_3","Paragraph":"Over time, the data releases have treated the Balmer line regions in different ways. The presence of the artificial curvature was first reported by Busca et al. (2013) in the context of the DR9 data release. To minimize this effect, a different scheme was used in DR12 (Alam et al. 2015, see their table 2) by using a linear function (instead of an iterative b-spline procedure) to interpolate the flux over the masked regions. Surprisingly, we observe that this data reduction change was only applied to the Balmer \u03b2, \u03b3, and \u03b4 lines but not applied to the Balmer \u03b1 line. For the latter one, the original iterative spline interpolation scheme has been used all along. As a result, an absorption-like feature at the location of Balmer \u03b1 is found in SDSS data releases 9 up to now, i.e. the latest data release 14 (Abolfathi et al. 2018). To illustrate this, we show examples of calibration vectors for SDSS BOSS DR9 (Ahn et al. 2012; Dawson et al. 2013), DR12 (Alam et al. 2015), eBOSS DR14 (Dawson et al. 2016; Abolfathi et al. 2018) data release as well as calibration vectors for the MaNGA survey (Bundy et al. 2015) DR14 data release (Abolfathi et al. 2018) in Fig. 6. In the upper panel, we show small-scale features in calibration vectors from the same plate (number 3647) processed by the DR9 (blue), DR12 (green), and DR14 (red) pipelines. The spectra are obtained by normalizing the large-scale features in calibration vectors with a cubic b-spline with break points separated by 50 \u00c5. The three spectra are mostly identical while the smooth curving features at the wavelengths of Balmer series in DR9 disappear in DR12 and DR14. From DR12, the pipeline corrects the features with linear interpolation across wavelength regions with Balmer \u03b2, \u03b3, and \u03b4 lines (Alam et al. 2015). One can also observe the change in the median residual spectra of DR9 and DR12 in Fig. 1. However, the H\u03b1 feature remains uncorrected from DR9 to DR14. We also show a DR7 calibration vector (black) in which most of the wiggles are absent. As pointed previously, this is due to the fact that the DR7 pipeline interpolates the calibration vectors using an effective scale larger than that used in subsequent data releases.","Citation Text":["Abolfathi et al. 2018"],"Functions Text":["data release as well as calibration vectors for the MaNGA survey","DR14 data release","in Fig. 6.","In the upper panel, we show small-scale features in calibration vectors from the same plate (number 3647) processed by the DR9 (blue), DR12 (green), and DR14 (red) pipelines. The spectra are obtained by normalizing the large-scale features in calibration vectors with a cubic b-spline with break points separated by 50 \u00c5. The three spectra are mostly identical while the smooth curving features at the wavelengths of Balmer series in DR9 disappear in DR12 and DR14. From DR12, the pipeline corrects the features with linear interpolation across wavelength regions with Balmer \u03b2, \u03b3, and \u03b4 lines","One can also observe the change in the median residual spectra of DR9 and DR12 in Fig. 1. However, the H\u03b1 feature remains uncorrected from DR9 to DR14."],"Functions Label":["Uses","Uses","Uses","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1138,1159]],"Functions Start End":[[1034,1098],[1119,1136],[1161,1171],[1172,1765],[1786,1937]]}
{"Identifier":"2020AandA...640L..11B__Segretain_1996_Instance_2","Paragraph":"Another possibly important cooling delay may arise from the phase separation of 22Ne during crystallization (Isern et al. 1991; Althaus et al. 2010). Our current best understanding is that at the small 22Ne concentrations typical of C\/O white dwarfs (\u223c1% by number), the presence of 22Ne should not affect the phase diagram, except near the azeotropic point of the C\/O\/Ne phase diagram. Thus, the crystallization of the C\/O core initially proceeds as in the case without 22Ne with no redistribution of neon ions between the solid and liquid phases. After a significant fraction of the core has crystallized, the temperature approaches the azeotropic point and the existing calculations indicate that the liquid phase is enriched in 22Ne relative to the solid (Segretain 1996; Garc\u00eda-Berro et al. 2008). The 22Ne-poor solid is lighter than the surrounding liquid and floats upward where it eventually melts. This gradually displaces the 22Ne-rich liquid downward toward the solid\u2013liquid interface until the azeotropic composition is reached, thereby releasing a considerable amount of gravitational energy. Given our very limited knowledge of the ternary C\/O\/Ne phase diagram (Segretain 1996; Hughto et al. 2012), this effect cannot be quantitatively implemented in our evolution models. However, we note that our current understanding of 22Ne phase separation is remarkably consistent with the missing cooling delay. In Fig. 2 we show the luminosity function obtained by adding an artificial 0.6 Gyr delay when 60% of the core is crystallized. These parameters are entirely consistent with those found in preliminary studies (Segretain 1996; Garc\u00eda-Berro et al. 2008) and yield an excellent fit to the crystallization pile-up3. Based on the current (albeit limited) knowledge of the C\/O\/Ne phase diagram, we propose that the phase separation of 22Ne in the advanced stage of crystallization significantly contributes to the pile-up in the luminosity function of 0.9\u22121.1\u2006M\u2299 white dwarfs (Fig. 2).","Citation Text":["Segretain 1996"],"Functions Text":["Given our very limited knowledge of the ternary C\/O\/Ne phase diagram","this effect cannot be quantitatively implemented in our evolution models.","However, we note that our current understanding of 22Ne phase separation is remarkably consistent with the missing cooling delay."],"Functions Label":["Differences","Differences","Similarities"],"Citation Start End":[[1176,1190]],"Functions Start End":[[1106,1174],[1213,1286],[1287,1416]]}
{"Identifier":"2016MNRAS.462.1508G__Gaur_et_al._2012c_Instance_1","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. 2012","c"],"Functions Text":["Interestingly, we found that, four HSPs did not show any IDV"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[497,513],[515,516]],"Functions Start End":[[408,468]]}
{"Identifier":"2018AandA...616L...2K__Frew_et_al._(2016)_Instance_2","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)"],"Functions Text":["The newest model developed by","is based on similar ideas, but makes use of the optical H\u03b1 surface brightness and a wide set of various calibrators."],"Functions Label":["Background","Background"],"Citation Start End":[[1523,1541]],"Functions Start End":[[1493,1522],[1542,1658]]}
{"Identifier":"2020MNRAS.498.1801K__Hosler,_Jensen_&_Goldshlak_1957_Instance_1","Paragraph":"Water ice is ubiquitous in the cold regions of the Universe, owing to the fact that hydrogen and oxygen are the two most abundant elements to form a solid such as icy dust particles and comets. It is, therefore, commonly accepted that the essential component of dust particles and planetesimals in protoplanetary discs is water ice beyond the so-called snow line, at which the temperature of gas is low enough for water vapour to condense into ices (e.g. Cyr, Sears & Lunine 1998). Reactive accretion of water ice from hydrogen and oxygen atoms on the surface of dust particles takes place in the dense core of molecular clouds where the growth of dust particles has been observed by scattering of stellar radiation (Steinacker et al. 2010). It is worthwhile noting that laboratory experiments on the coagulation growth of water-ice particles have a long history outside astronomy and planetary science, since coagulation is observed in daily life and is a plausible route to the formation of snowflakes (e.g. Faraday 1860; Hosler, Jensen & Goldshlak 1957). Recent works on laboratory measurements of cohesion between crystalline water-ice particles at vacuum conditions provided encouraging results that dust particles composed of water ice might be much more cohesive than previously believed (Gundlach et al. 2011; Gundlach & Blum 2015; Jongmanns et al. 2017). Form a theoretical point of view, Chokshi, Tielens & Hollenbach (1993) demonstrated that the JKR theory of elastic contact formulated by Johnson, Kendall & Roberts (1971) is a powerful tool for better understanding of dust coagulation. Numerical simulations incorporating the JKR theory have shown that dust aggregates consisting of submicrometre-sized water-ice particles proceed with coagulation growth even at a collision velocity of 50 m s\u22121 (Wada et al. 2009, 2013). As a result, the majority of recent studies on dust coagulation and planetesimal formation assume that silicate aggregates are disrupted by mutual collision at a velocity of vdisrupt \u223c 1 m s\u22121, but icy aggregates at vdisrupt \u223c 10 m s\u22121 (e.g. Birnstiel, Dullemond & Brauer 2010; Vericel & Gonzalez 2019). Such a trendy assumption led Dr\u0105\u017ckowska & Alibert (2017) to propose planetesimal formation by the \u2018traffic jam\u2019 effect at the snow line, provided that sticky water-ice particles grow faster and thus drift toward the central star faster than less-sticky bare silicate particles, implying that aggregates of the former catch up the latter at the snow line, which results in a traffic jam. However, we argue that the importance of water ice to dust coagulation is still open to debate, since water ice is not necessarily stickier than other materials such as silicates and complex organic matter (Kimura et al. 2015, 2020a; Musiolik & Wurm 2019).","Citation Text":["Hosler, Jensen & Goldshlak 1957"],"Functions Text":["It is worthwhile noting that laboratory experiments on the coagulation growth of water-ice particles have a long history outside astronomy and planetary science, since coagulation is observed in daily life and is a plausible route to the formation of snowflakes (e.g."],"Functions Label":["Background"],"Citation Start End":[[1024,1055]],"Functions Start End":[[742,1009]]}
{"Identifier":"2017ApJ...835...79S__Wilson_et_al._1974_Instance_1","Paragraph":"The PN central white dwarf generates a significant ultraviolet (UV) radiation field during its transition from the proto-PN stage to the PN stage, typically \u223c105 times that of the general interstellar medium (ISM), but decreasing to 10\u2013100 as the remnant envelope expands and the luminosity of the white dwarf declines (Cox 1997). The copious amount of UV photons penetrating the nebula is predicted to photodissociate the remnant molecular material from the AGB stage within a thousand years (e.g., Redman et al. 2003). Indeed, the first molecular searches of PNe turned up negative (e.g., Penzias et al. 1971; Wilson et al. 1974). Subsequently, both CO and H2 have been detected in numerous PNe (e.g., Huggins & Healy 1989; Huggins et al. 1996, 2005; Hora et al. 1999; Likkel et al. 2006). More recent studies have identified other molecules in planetary nebulae with ever increasing complexity. Toward the young PNe NGC 7027 and NGC 6537, for example, HCN, HNC, CCH, CS, SO, H2CO, HCO+, and N2H+ have been identified (Zhang et al. 2008; Edwards & Ziurys 2013). In more evolved nebulae, species such as HCN, HNC, HCO+, CS, and CN have been found (e.g., Cox et al. 1992; Bachiller et al. 1997; Edwards et al. 2014). The very old Helix Nebula (age \u223c12,000 years.) has an assortment of interesting molecules, including HCO+, CN, HCN, H2CO, CCH, c-C3H2, and HNC (Bachiller et al. 1997; Tenenbaum et al. 2009), and mapping data have shown that the molecular material is distributed throughout the PN (Zack & Ziurys 2013; Zeigler et al. 2013). SiO, SO, and SO2 have recently been found in M2-48 (age \u223c5000 years.) as well (Edwards & Ziurys 2014). Perhaps even more striking is the identification of C60 in several PNe with low central star temperatures (Cami et al. 2010; Garc\u00eda-Hern\u00e1ndez et al. 2010, 2011, 2012; Otsuka et al. 2013). The presence of long-lived molecular material is believed to arise from shielding by high-density clumps (Howe et al. 1994), as seen in images of the Helix and the Ring (e.g., Meaburn et al. 1992; Speck et al. 2003; Meixner et al. 2005).","Citation Text":["Wilson et al. 1974"],"Functions Text":["Indeed, the first molecular searches of PNe turned up negative (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[612,630]],"Functions Start End":[[521,590]]}
{"Identifier":"2019MNRAS.485.4343C__its_2018_Instance_1","Paragraph":"As of 2018, three multimeasurement catalogues including a substantial amount of redshift-independent extragalactic distance measurements have been released: HyperLEDA (Makarov et al. 2014), NED-D (Mazzarella & Team 2007; Steer et al. 2017), and Cosmicflows-3 (Tully, Courtois & Sorce 2016). HyperLEDA includes a homogenized catalogue for extragalactic distances in the nearby universe, with 12\u2009866 distance measurements for 518 galaxies to date. NED-D is the NASA\/IPAC Extragalactic Distance catalogue of redshift-independent distances, which compiles 326\u2009850 distance measurements for 183\u2009062 galaxies in its 2018 version. Here, \u223c1800 galaxies (\u223c1 per cent) have more than 13 distance measurements, and 180 galaxies (\u223c0.1 per cent) have distance measurements using more than six different methods. Cosmicflows-3 is the most up-to-date catalogue, which reports distance measurements for 10\u2009616 galaxies (all of which include errors) using up to four distance determination methods, calibrated with supernova luminosities. However, unlike HyperLEDA or NED-D, Cosmicflows-3 only reports the latest distance measurement for each method. In HyperLEDA, NED-D, and Cosmicflows-3 errors are reported as one standard deviation from the reported distance modulus. Treatment of errors for combining distance moduli across methods or across measurements is suggested by Mazzarella & Team (2007) and Tully et al. (2016) to be based on weighted estimates such as the uncertainty of the weighted mean, albeit with caution partly due to the heterogeneous origin of the compiled data and partly due to Malmquist bias. In the case of NED-D, this is additionally complicated by the fact that many errors are not reported or are reported as zero. In fact, the TF relation method has the largest number of galaxies with non-reported distance modulus errors (884 to date). Even though extragalactic distances measured using the TF relation were originally reported to have a relative error in distance modulus of 10\u201320 per cent (Tully & Fisher 1977), we consider that this conservative estimate can be improved upon by using a predictive model based on the distance error of galaxies that use the same distance determination method. This requires a robust estimation of the variance of extragalactic distances based on the available data.","Citation Text":["Mazzarella & Team (2007)"],"Functions Text":["Treatment of errors for combining distance moduli across methods or across measurements is suggested by","and Tully et al. (2016) to be based on weighted estimates such as the uncertainty of the weighted mean, albeit with caution partly due to the heterogeneous origin of the compiled data and partly due to Malmquist bias."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1359,1383]],"Functions Start End":[[1255,1358],[1384,1601]]}
{"Identifier":"2019MNRAS.486.5239D__Galli_et_al._2014_Instance_1","Paragraph":"The emission from both types of sources contaminates measurements of the Cosmic Microwave Background (CMB), which contains information on cosmological parameters. The measurement of the CMB polarization, particularly the E-mode damping tail, can help break some of the degeneracies between cosmological parameters. Measurements of the polarization damping tail are expected to become foreground-limited at a smaller angular scale (higher \u2113) than the temperature damping tail, because of the expected low polarization of dusty point sources (see below). Further, the higher contrast of the acoustic features in EE power spectrum compared to astrophysical foregrounds (Calabrese et al. 2014; Galli et al. 2014) will ultimately provide independent and tighter constraints on the standard cosmological parameters, such as the scalar spectral index ns, than those from the temperature data alone. High-resolution measurements of the E-mode polarization will improve the delensing of the primordial B-modes (Seljak & Hirata 2004), ultimately tightening the constraint on the tensor-to-scalar ratio r. As measurements of the small angular-scale fluctuations in the CMB are attaining higher sensitivity and finer resolution, ongoing and planned ground-based CMB surveys, such as Advanced ACTPol (Henderson et al. 2016), SPT-3G (Benson et al. 2014), Simons Observatory (Ade et al. 2019 (SO)), CCAT-prime (Stacey et al. 2018), and CMB Stage 4 (Abazajian et al. 2016; Abitbol et al. 2017) will be capable of extracting information from the E-mode damping tail out to \u2113 \u2248 9000. However, the contribution of the extragalactic point sources to the CMB power spectrum increases towards smaller angular scales, and it is expected to be a significant fraction of the CMB polarization power. For example, extragalactic foreground sources are expected to be the predominant contaminant for angular scales smaller than 30\u2032 (\u2113\u2273 400) in the 70\u2013100 GHz frequency range (Toffolatti et al. 1998). Hence, characterization of these sources in terms of their spectral and spatial distributions is essential for separating foregrounds from the CMB.","Citation Text":["Galli et al. 2014"],"Functions Text":["Further, the higher contrast of the acoustic features in EE power spectrum compared to astrophysical foregrounds","will ultimately provide independent and tighter constraints on the standard cosmological parameters, such as the scalar spectral index ns, than those from the temperature data alone."],"Functions Label":["Future Work","Future Work"],"Citation Start End":[[690,707]],"Functions Start End":[[553,665],[709,891]]}
{"Identifier":"2022MNRAS.511.3477E__Carter_&_Erd\u00e9lyi_2008_Instance_1","Paragraph":"It has been cleared out in the observations that a coronal flux tube could consist of a twisted magnetic field (e.g. Kwon & Chae 2008; Aschwanden et al. 2012; Thalmann et al. 2014; Wang et al. 2015). In a magnetically twisted flux tube, the background magnetic field has an azimuthal component which causes the magnetic field lines to wind around the axis of the tube. The winding number (the number of twist turns around the tube axis) has been reported to have values in the range (0.11, 0.87) (Kwon & Chae 2008). The role of twisted magnetic field in the MHD waves has been studied in many theoretical works (e.g. Bennett et al. 1999; Erd\u00e9lyi & Carter 2006; Erd\u00e9lyi & Fedun 2006, 2007, 2010; Ruderman 2007, 2015; Carter & Erd\u00e9lyi 2008; Karami & Barin 2009; Karami & Bahari 2010, 2012; Terradas & Goossens 2012; Ruderman & Terradas 2015; Ebrahimi & Karami 2016; Ebrahimi et al. 2017; Terradas et al. 2018). Sakurai et al. (1991a) studied resonant absorption of MHD waves in presence of a twisted magnetic field and obtained jumps of the perturbations around the resonance point (see also Sakurai et al. 1991b; Goossens et al. 1992). Ebrahimi et al. (2017) investigated resonant absorption of kink MHD waves and phase mixing of the perturbations in twisted magnetic flux tubes. They showed that depending on the profile of the azimuthal component of the magnetic field and the azimuthal mode number of the global kink wave (m = \u00b11) magnetic twist can enhance or suppress the decay rate of the kink wave and the rate of phase mixing of the perturbations. Ebrahimi et al. (2017) concluded that even a small amount of twist can have a large impact on the generation of small spatial scales and transferring the wave energy from the global mode to the local Alfv\u00e9nic oscillations. However, Ebrahimi et al. (2017) neglected the dissipation effects in their analysis. It is known that dissipation could affect the generation of small scales due to phase mixing (McLaughlin et al. 2011; Ebrahimi et al. 2020; Howson et al. 2020).","Citation Text":["Carter & Erd\u00e9lyi 2008"],"Functions Text":["The role of twisted magnetic field in the MHD waves has been studied in many theoretical works (e.g."],"Functions Label":["Background"],"Citation Start End":[[716,737]],"Functions Start End":[[516,616]]}
{"Identifier":"2021MNRAS.506.5015H__Weinberg,_Miller_&_Lamb_2001_Instance_1","Paragraph":"Based purely on quality of fits of atmosphere model spectra to observed spectra obtained here, one composition is not preferred over another for the three older CCOs. The primary differences in fit results when the atmosphere is composed of carbon instead of hydrogen are a \u223c40 per\u2009cent lower temperature and \u223c3 times larger emission radius (to maintain nearly constant $R_{\\rm em}^2T^4$). The larger Rem could be an argument in favour of carbon because it is closer to the neutron star radius R, which would imply that the entire surface is essentially at a single temperature and would explain non-detection thus far of pulsations from each of these three CCOs. However, a hot region with radius 3\u20134 times smaller than R can still produce a pulsed fraction below current limits of 20\u201340 per\u2009cent in the spin period range 0.1\u20130.4 s of known CCOs (see Section 1; for pulsed fraction dependence on spot size, see, e.g. Psaltis, \u00d6zel & DeDeo 2000; DeDeo, Psaltis & Narayan 2001; Weinberg, Miller & Lamb 2001; Bogdanov, Grindlay & Rybicki 2008; Lamb et al. 2009; Baub\u00f6ck, Psaltis & \u00d6zel 2015; Elshamouty et al. 2016). For example, Gotthelf, Perna & Halpern (2010) find that a model which includes a hotspot with Rem\/R \u223c 0.4 (and a second smaller spot) is able to produce a pulsed fraction that matches the 11 per\u2009cent measured for the CCO in Puppis A. Also of note is the stronger limit on the Cassiopeia A CCO pulsed fraction of 12 per\u2009cent for P > 0.01s compared to the other CCOs (see Table 1). Meanwhile, a hydrogen atmosphere is a natural consequence of even a very low-level of accretion from the interstellar medium on to a relatively cool neutron star surface several hundred years after neutron star formation (see Section 5.1 for further discussion). Thus from an evolution standpoint, a hydrogen atmosphere for the older CCOs studied here might be preferred. Future measurements of pulsations or improvements to pulsation constraints could provide stronger indications of their atmosphere composition.","Citation Text":["Weinberg, Miller & Lamb 2001"],"Functions Text":["However, a hot region with radius 3\u20134 times smaller than R can still produce a pulsed fraction below current limits of 20\u201340 per\u2009cent in the spin period range 0.1\u20130.4 s of known CCOs (see Section 1; for pulsed fraction dependence on spot size, see, e.g."],"Functions Label":["Uses"],"Citation Start End":[[977,1005]],"Functions Start End":[[664,917]]}
{"Identifier":"2017MNRAS.469.2720G__Hern\u00e1ndez-Garc\u00eda_et_al._2015_Instance_1","Paragraph":"With all this in mind, our last question is: What powers soft X-rays and [O\u2009III] in LINERs? This has no clear answer, but both are not tracing the same mechanism, since none of them match in morphologies. This is a clear difference between type-2 Seyferts and LINERs. In favour of the soft X-ray emission being originated by AGN photoionization, the RGS spectra studied by Gonzalez-Martin et al. (2010b) showed that in at least 30 per cent of their sample, a contribution of photoionization by the AGNs is required due to the presence of radiative recombination continua (RRC) from CV emission line. However, this does not guarantee a dominance of this emission mechanism. Moreover, cone-like morphologies at soft X-rays in some objects in this study point out again to the photoionization by the AGNs being responsible for the soft X-ray emission. Nevertheless, this assumes that we are seing LINERs with an LOS perpendicular to the accretion disc, which might not be the case. Indeed, the ultraviolet and X-ray variability detected for many of these LINERs (Maoz et al. 2005; Hern\u00e1ndez-Garc\u00eda et al. 2015) is in favour of a direct view of the AGNs (i.e. perpendicular to the disc under the UM). This is consistent with the fact that most of the [O\u2009III] morphologies found for LINERs are spheroids, if we assume that the [O\u2009III] traces the NLR. In addition, the fact that we detected a clear correspondence between soft X-ray and [O\u2009III] morphologies only in objects with log\u2009(LHX)>40, and also that all the objects where soft X-rays and [O\u2009III] match their morphologies seem to better follow the previously found relation between the size of the region and the hard X-ray luminosity (see Fig. 2 and Section 5.2), may argue in favour of the scenario in which the AGNs do not have enough thrust to ionize in the low-luminosity regime (Elitzur & Shlosman 2006; Elitzur & Ho 2009), ruling out photoionization by the AGNs at both soft X-ray and [O\u2009III] emissions. In this case, the most reasonable explanation for the [O\u2009III] is the host galaxy emission, which, anyhow, could also be on top of the AGNs, preventing its detection and erasing the connection (Gonz\u00e1lez-Mart\u00edn et al. 2014). The host galaxy can contribute either as star formation or shocks to the total [O\u2009III] emission. Regarding the soft X-ray origin, Mingo et al. (2014) confirmed that jets are the main responsible for soft X-ray emission from their sources. In our sample, jets are identified in NGC 1052 (Kadler et al. 2004), where the jet position angle would be consistent with the extended soft X-ray emission shown here.","Citation Text":["Hern\u00e1ndez-Garc\u00eda et al. 2015"],"Functions Text":["Indeed, the ultraviolet and X-ray variability detected for many of these LINERs","is in favour of a direct view of the AGNs (i.e. perpendicular to the disc under the UM). This is consistent with the fact that most of the [O\u2009III] morphologies found for LINERs are spheroids, if we assume that the [O\u2009III] traces the NLR."],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[1078,1106]],"Functions Start End":[[979,1058],[1108,1345]]}
{"Identifier":"2020AandA...642A.140G__Guilera_et_al._(2019)_Instance_1","Paragraph":"In addition, we also perform two new simulations changing the type I migration recipes from those of Tanaka et al. (2002) to those of Paardekooper et al. (2011) and Jim\u00e9nez & Masset (2017). With this, we want to analyze whether or not the planet migration trap could be a robust result even using more sophisticated type I migration recipes, which were derived for nonisothermal disks. We remark that despite our simplifications in the computation of the thermodynamics of the disk, we compute the time evolution of the radial profiles of all the quantities needed to calculate the migration recipes mentioned above, that is, we compute the time evolution of the radial profiles of the density, pressure, opacity, and so on while the temperature radial profile remains fixed in time. Explicitly, the simplification lies in the fact that the evolution of the temperature is not linked self-consistently with the evolution of the gas surface density. We can see in Fig. 8 that the density and pressure maximum also acts as a planet migration trap using the type I migration recipes from Jim\u00e9nez & Masset (2017). However, using the type I migration prescriptions from Paardekooper et al. (2011), the migration trap is broken close to the mass needed for the planet to open a gap in the disk. This happens because now the normalized torque is not only a function of the local gradient of the gas surface density (as in the isothermal case) but also of the local gradient of the temperature, the viscosity, and the mass of the planet (these last two dependencies through the corotation torque). Thus, the combination of such quantities breaks the migration trap at some moment. Guilera et al. (2019) showed that the main differences between the migration recipes from Jim\u00e9nez & Masset (2017) and Paardekooper et al. (2011) lie in the computation on the corotation torque. Thus, in this case, when the planet reaches a mass of a several tens of Earth masses, the migration trap is broken. At this moment, the planet quickly migrates inwards (due to the fact that it has a mass of ~ 80 M\u2295) until it opens a gap at ~1.15 au. We note that the time evolution of the mass of the core and the mass of the envelope for the simulations using the type I migration recipes from Paardekooper et al. (2011) and Jim\u00e9nez & Masset (2017) are very similar to those of the fiducial simulation. In Fig. 9, we emphasize this by plotting the planet formation tracks as a function of time. Simulations also stopped when the planets opened a gap in the disk. This happens when the total mass of the planet is about 110 M\u2295 for the case where we used the migration recipes from Paardekooper et al. (2011) and about 95 M\u2295 for the case where we adopt the migration recipes from Jim\u00e9nez & Masset (2017). The differences in the total mass at which the planet opens the gap are due to the fact that this mass depends on the viscosity of the disk. Inside the ice line, the viscosity is larger (because of the larger \u03b1-parameter), and therefore the total mass of the planet has to be higher to open a gap. For the second case, the planet opens the gap in the \u03b1 transition region between 10\u22125 (outside the ice line) and 10\u22123 (inside the ice line). We note that if the thermodynamics of the disk are computed self-consistently, the migration trap can break at a different planet mass depending mainly on the viscosity and thermal diffusivity of the disk (see Morbidelli 2020).","Citation Text":["Guilera et al. (2019)"],"Functions Text":["showed that the main differences between the migration recipes from Jim\u00e9nez & Masset (2017) and Paardekooper et al. (2011) lie in the computation on the corotation torque."],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1673,1694]],"Functions Start End":[[1695,1866]]}
{"Identifier":"2022AandARv..30....6M__Ellison_et_al._2015_Instance_1","Paragraph":"In particular, two main scenarios have emerged: the first one attributes differences in radio morphology to the large-scale (i.e., host galaxy and environment, see also below in this Section) properties of radio-AGN which influence the interaction between the radio jet and the external medium (e.g., Kaiser and Best 2007; Wing and Blanton 2011; Miraghaei and Best 2017; Mingo et al. 2019, 2022), with those between HERGs and LERGs being instead dictated by different fuelling mechanisms as discussed earlier. According to this framework, HERGs are powered by accretion of cold gas, provided by e.g., a recent merger with a gas-rich galaxy, while LERGs accrete hot intergalactic gas from dense environments at a low rate (e.g., Best and Heckman 2012), with fuelling from major mergers strongly disfavoured by recent observations (e.g., Ellison et al. 2015). On the other hand, other works do not observe any difference in the host and\/or environmental properties of FRI and FRII galaxies, except in the rare cases of FRII HERGs (e.g., Lin et al. 2010; Capetti et al. 2017b; Jimenez-Gallardo et al. 2019; Massaro et al. 2019, 2020; Vardoulaki et al. 2021) or even between those of HERGs and LERGs (e.g., Fernandes et al. 2015), so that an alternative mechanism for their large-scale radio behaviour has to be invoked. In this case, ageing processes can be thought as the main driver for the observed morphological differences, with an evolutionary pattern that proceeds from FRII HERGs that switch from efficient to inefficient accretion due to gas starvation and transform themselves into FRII LERGs, sources that still maintain their large radio structures thanks to the past nuclear activity at high efficiency (e.g., Ghisellini and Celotti 2001; Tadhunter 2016; Macconi et al. 2020; Grandi et al. 2021). The switch off\/change in accretion mode will eventually show in the radio morphology with the delay needed to reach Kpc-to-Mpc distances, and the source will ultimately turn into an FRI galaxy.","Citation Text":["Ellison et al. 2015"],"Functions Text":["According to this framework, HERGs are powered by accretion of cold gas, provided by e.g., a recent merger with a gas-rich galaxy, while LERGs accrete hot intergalactic gas from dense environments at a low rate","with fuelling from major mergers strongly disfavoured by recent observations (e.g.,"],"Functions Label":["Background","Background"],"Citation Start End":[[836,855]],"Functions Start End":[[510,720],[752,835]]}
{"Identifier":"2020MNRAS.494.1045B__Raymond_et_al._2009_Instance_1","Paragraph":"In our simulations, we assumed that the growth of dust grains and planetesimals occurred during the first stages of planetary accretion. Thus, we used a bimodal disc composed of embryos (60 per cent of the disc\u2019s mass) and planetesimals (40 per cent of the disc\u2019s mass; Izidoro et al. 2015). It was assumed that embryos are formed by oligarchic growth and are thus spaced randomly by 5\u201310 mutual Hill radii (Kokubo & Ida 1998, 2000) with a density of 3\u2009g\u2009cm\u22123. The mass of each planetesimals is \u2248 0.002\u2009M\u2295. In the numerical integrations, the planetesimals do not have gravitational interactions with themselves but only with stars, planets, and protoplanetary embryos. The protoplanetary embryos masses scale as M \u2248 r3(2 \u2212 x)\/2\u03943\/2 (Kokubo & Ida 2002; Raymond, Quinn & Lunine 2005; Raymond et al. 2009; Izidoro et al. 2015), where \u0394 is the mutual Hill radii separations between embryos orbits. As we are using distinct systems with different parameters, the number of embryos and planetesimals are not the same among the systems; see Table 4. Fig. 5 shows the initial conditions of x = 1.5 and 2.5, for all our simulations. Discs with x = 2.5 have embryos more massives in the inner region of the disc, and with 1.5 in the outer region. We chose these two values of the parameter x, besides being used frequently in some works (Raymond, Quinn & Lunine 2004; Raymond et al. 2005; Izidoro et al. 2014a,2015; Izidoro, Morbidelli & Raymond 2014b; Izidoro & Raymond 2018), to study the dynamic evolution in discs that have mass growing as a function of the orbital radius of the bodies and in cases where the mass decreases as a function of the orbital radius (see Fig. 5). This is an important point because we have some systems where the giant planet is in the inner region of the disc and others in the outer region. So we needed to consider the two cases for all systems. The orbital inclination of the embryos and planetesimals is chosen randomly between $10^{-4}\\, ^{\\circ }$ and $10^{-3}\\, ^{\\circ }$ with respect to the binary plane, and the eccentricity is chosen between 0 and 0.01.","Citation Text":["Raymond et al. 2009"],"Functions Text":["he protoplanetary embryos masses scale as M \u2248 r3(2 \u2212 x)\/2\u03943\/2","where \u0394 is the mutual Hill radii separations between embryos orbits."],"Functions Label":["Uses","Uses"],"Citation Start End":[[782,801]],"Functions Start End":[[670,731],[825,893]]}
{"Identifier":"2017MNRAS.469S.731L__Kolokolova_&_Kimura_2010_Instance_1","Paragraph":"The specific radiance received from the optically thin dust plume can be expressed as\n(1)\r\n\\begin{equation}\r\nL_\\lambda =f_{\\rm plume}\\frac{p}{\\pi } \\frac{\\phi (\\alpha )}{\\phi (0)} \\frac{f_{\\lambda }}{r_{\\rm h}{^2}} \r\n\\end{equation}\r\nwhere p is the geometric albedo of the dust particles at wavelength \u03bb, \u03d5(\u03b1) is the phase function at phase angle \u03b1, f\u03bb is the solar irradiance (in W m\u22122 nm\u22121 at 1 au), rh is the heliocentric distance in au and fplume is the dust filling factor (e.g. the fraction of a pixel covered by dust). Assuming that the geometric albedo of the dust particles is similar to that of the nucleus, the geometric albedo and specific solar flux are 0.068 and 1.513 W \u22122 nm\u22121 for the orange (648.6 nm) filter, and 0.027 and 0.23 W \u22122 nm\u22121 for the UV (270.7nm) filter (Fornasier et al. 2015), respectively. The flux ratio \u03d5(\u03b1)\/\u03d5(0) can be derived by using the dust phase function (Kolokolova & Kimura 2010). By integrating the radiance over the image area (e.g. within the boxes in Fig. 7) containing the dust from a difference outburst plume, the total cross-section from difference images of the dust jet from within a given time interval can be determined. Then, we used a power-law index g = 3.7 for the size distribution with a constant bulk density of 1000 kg m\u22123 for all ejected particles. The ejected mass in size interval of a1  a  a2 (i.e. 1 \u03bcm to 1 mm) is given by\n(2)\r\n\\begin{equation}\r\nM=\\frac{({4}\/{3})\\pi \\rho N}{4-g} \\left(a_2^{4-g}-a_1^{4-g}\\right), \r\n\\end{equation}\r\nwhere N is the total number of dust particles in the size interval 1 \u03bcm to 1 mm. The estimated dust cross-section and ejected mass are given in Table 2. The uncertainty of the ejected mass is typically of the order of 5\u201310 per cent, depending on the uncertainty on the brightness integration of the selected box. The averaged mass ejection rate for the outbursts can be estimated from the ejected mass (M) by dividing by the time interval between the difference images. Note that our calculation is based on the chosen time intervals and hence produces lower limits on the mass ejection rates. For example, the estimation of the ejected mass for a major outbursts on July 29 is about 4550 kg. Given a time interval of 18 min, the average mass ejection rate would be of the order of 4 kg s\u22121. However, if this outburst lasted for a time interval of 5 min only, which is the cadence for the outburst sequence, the corresponding mass ejection rate would be as much as 15 kg s\u22121.","Citation Text":["Kolokolova & Kimura 2010"],"Functions Text":["The flux ratio \u03d5(\u03b1)\/\u03d5(0) can be derived by using the dust phase function"],"Functions Label":["Uses"],"Citation Start End":[[896,920]],"Functions Start End":[[822,894]]}
{"Identifier":"2016MNRAS.455..552B__Riess_et_al._1998_Instance_1","Paragraph":"The plethora of cosmological observations has turned cosmology into a quantitative science. From combining several probes that observe the Universe at different epochs and have different systematics and statistics, emerged the \u2018concordance model\u2019 of cosmology, a six parameter model, most of them measured to the accuracy of a per cent. Among these probes, Type Ia supernovae (SNe) are a powerful cosmological tool to directly measure the expansion history of the Universe (see e.g. a recent review of Weinberg et al. 2013). The type Ia SNe are known to be standard candles and one can measure the luminosity distances to them accurately. Several SNe surveys have set strong constraints on cosmological models from the distance\u2013redshift relation (e.g. Riess et al. 1998; Perlmutter et al. 1999; Riess et al. 2007; Sullivan et al. 2011; Campbell et al. 2013). However the measured apparent magnitudes, have a residual scatter arising from its intrinsic scatter and effects due to line-of-sight (LOS) structures. The intrinsic scatter (\u223c0.4 mag) can be significantly reduced to \u223c0.1 mag by empirically calibrating the luminosity curves. The scatter due to photon deflection along the LOS is composed of many different physical effects. The dominant effects are peculiar velocities at z  0.1 and gravitational lensing at z \u2273 0.3. Within the realm of cosmological perturbation theory and the stochastic nature of the LOS structures, all effects are expressed as an integral over the power spectrum with appropriate kinematical factors (Ben-Dayan et al. 2012). The lensing effects create residuals from the best-fitting curve in the magnitude\u2013redshift relation. The lensing magnifications of SNe can be extracted by correlating the residuals with the surface densities of nearby foreground galaxies. The lensing dispersion is roughly proportional to the SNe redshift z, (e.g. Holz & Linder 2005; Ben-Dayan et al. 2013), and specifically for the concordance model, the predicted value is \u223c0.06z mag.","Citation Text":["Riess et al. 1998"],"Functions Text":["Several SNe surveys have set strong constraints on cosmological models from the distance\u2013redshift relation (e.g.","However the measured apparent magnitudes, have a residual scatter arising from its intrinsic scatter and effects due to line-of-sight (LOS) structures."],"Functions Label":["Background","Background"],"Citation Start End":[[752,769]],"Functions Start End":[[639,751],[859,1010]]}
{"Identifier":"2015MNRAS.451.2123T__King_&_Begelman_1999_Instance_1","Paragraph":"The results of Ivanova et al. (2003) are somewhat closer to ours. A remaining key difference is that Ivanova et al. (2003) find that their calculated mass-transfer rates for systems with initial orbital periods, Porb, i  0.4\u2009d sometime exceed a critical mass-transfer rate, $\\dot{M}_{\\rm crit}$ related to the location of the so-called trapping radius (e.g. Begelman 1979; Chevalier 1993; MacLeod & Ramirez-Ruiz 2015b). In systems with very large super-Eddington mass-transfer rates, matter presumably piles up around the NS and forms a growing, bloated cloud engulfing a large fraction of the accretion disc. A system will only avoid a CE if it manages to evaporate the bulk of the transferred matter via the liberated accretion energy at a distance from the NS larger than the trapping radius. Otherwise, the incoming material has too much negative binding energy to be ejected. At the same time, this trapping radius must be located inside the Roche lobe of the NS in order to avoid a CE (King & Begelman 1999). The exact location of the trapping radius, and thus the value of $\\dot{M}_{\\rm crit}$, is difficult to calculate because it also depends on the cooling processes of the infalling gas (Narayan & Yi 1995; Blandford & Begelman 1999). Nevertheless, in all of our models presented in Table 1 with Porb, i > 0.06\u2009d we find that $|\\dot{M}_{\\rm He}^{\\rm max}|\\le \\dot{M}_{\\rm crit}$, where $\\dot{M}_{\\rm crit}$ is calculated from equation 15 in Ivanova et al. (2003), and varies from about 3.7 \u00d7 10\u22123\u2009M\u2299\u2009yr\u22121 (for Porb, i = 2.0\u2009d) to 4.3 \u00d7 10\u22124\u2009M\u2299\u2009yr\u22121 (for Porb, i = 0.08\u2009d). Only for Porb, i = 0.06\u2009d, we find a couple of systems where $|\\dot{M}_{\\rm He}^{\\rm max}|> \\dot{M}_{\\rm crit}$. However, these systems are anyway found to result in a runaway mass transfer and thus formation of a CE. Therefore, our calculated mass-transfer rates from the helium stars donors must be slightly smaller than those calculated by Ivanova et al. (2003). The reason for this, besides from using different stellar evolution codes, is perhaps that our mass-transfer rates are calculated using the prescription by Ritter (1988) whereas Ivanova et al. (2003) adopted the prescription by Tout & Eggleton (1988). To summarize, we caution that all numerical calculations of Case BB RLO could potentially be affected by uncertain accretion processes at high mass-transfer rates exceeding the Eddington limit by \u223c4 orders of magnitude.","Citation Text":["King & Begelman 1999"],"Functions Text":["In systems with very large super-Eddington mass-transfer rates, matter presumably piles up around the NS and forms a growing, bloated cloud engulfing a large fraction of the accretion disc. A system will only avoid a CE if it manages to evaporate the bulk of the transferred matter via the liberated accretion energy at a distance from the NS larger than the trapping radius. Otherwise, the incoming material has too much negative binding energy to be ejected. At the same time, this trapping radius must be located inside the Roche lobe of the NS in order to avoid a CE"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[992,1012]],"Functions Start End":[[420,990]]}
{"Identifier":"2018MNRAS.478.2541F__Smith_&_Tombleson_2015_Instance_1","Paragraph":"Utilizing the full range of peak absolute magnitudes observed in LBV stars [M \u2243 \u2009(\u221213)\u2013(\u22129)\u2009mag; e.g. Smith et al. 2011b] provides a range of peak apparent magnitudes of m \u2243 \u200916.5\u201320.5\u2009mag, for a distance of D \u2243 8\u2009Mpc (Table 1). Hence, if the transient source (m \u2248 21.0\u2009mag; Section 2.1.2) is an LBV star, it must have been observed a short time after its peak (Fig. 2; left-hand column; rows 4\u20135); quiescent LBV stars can have absolute magnitudes as low as M \u2243 \u22126\u2009mag (e.g. Smith et al. 2011b), or m \u2243 23.5\u2009mag (for a distance of D \u2243 8\u2009Mpc; Table 1). An LBV star provides an adequate explanation for the transient time-scale, the isolation (e.g. Smith & Tombleson 2015; Smith 2016), the lack of a host H\u2009ii region (Fig. 2 and 3; see below, however), and the transient\/main source offset (d \u2248 0.3\u2009kpc, for D \u2243 8\u2009Mpc; Table 1; see below, however). In addition, (net) fading of the LBV peak event by \u0394m \u2272 3\u2009mag (over a period of approximately 40 yr) may possibly produce an apparent brightness centroid shift\/morphology variation in the main source (Sections 2.1.2 and 2.1.3). However, this scenario has its caveats. First, it is difficult to interpret the flux variability of the main source (Sections 2.1.4 and 2.3.1) within such a context. Secondly, although a fraction of LBV stars are isolated (e.g. Smith & Tombleson 2015; Smith 2016), in the only XMP with a documented LBV star, the LBV star appears embedded within an H\u2009ii region (DDO 68; Pustilnik et al. 2017); no clear H\u2009ii region at the location of the transient source is discernible in the HST images (Fig. 3; top), although it should be detectable given the typical lifetime of an H\u2009ii region (few Myr; e.g. Alvarez, Bromm & Shapiro 2006). Thirdly, a (net) magnitude variation of \u0394m \u22723\u2009mag of the transient over a period of approximately 60 yr would result in a quiescent LBV source (m \u2272 23.5\u2009mag) that should have been detectable in the HST images (Fig. 3; top). Lastly, as LBV stars brighten, they become redder (e.g. Sterken 2003), which appears to contradict the POSS data (Table 2; Fig. 2; rows 2\u20135). Consequently, as the LBV scenario explains some of the observables challenged by other scenarios (e.g. transient timeline), it remains a contender for the observed phenomenon.","Citation Text":["Smith & Tombleson 2015"],"Functions Text":["An LBV star provides an adequate explanation for the transient time-scale, the isolation (e.g.","the lack of a host H\u2009ii region (Fig. 2 and 3; see below, however), and the transient\/main source offset (d \u2248 0.3\u2009kpc, for D \u2243 8\u2009Mpc; Table 1; see below, however)."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[647,669]],"Functions Start End":[[552,646],[684,846]]}
{"Identifier":"2020MNRAS.494.5270Z__Pesnell,_Thompson_&_Chamberlin_2012_Instance_1","Paragraph":"In particular, the presence of slow magneto-acoustic-gravity waves (MAG waves, hereafter) guided by the magnetic field is supported by increasingly observational evidence in different atmospheric structures, e.g. photospheric flux tubes (Roberts & Ulmschneider 1997), sunspots (Jess et al. 2013; Freij et al. 2014; Khomenko & Collados 2015; Madsen, Tian & DeLuca 2015), coronal loops (King et al. 2003), and coronal plumes (Nakariakov 2006). These oscillations are also used to estimate the formation heights of different emission spectral lines. As a result of the analysis of 3 min oscillations detected in observations from the Atmospheric Imaging Assembly instrument (AIA; Lemen et al. 2012) on-board the Solar Dynamic Observatory (SDO; Pesnell, Thompson & Chamberlin 2012), Deres & Anfinogentov (2015) found that the formation heights of the corresponding spectral lines are consistent with models of the sunspots umbra involving strong temperature gradients of the type used in this paper (e.g. Fontenla et al. 2009). Meanwhile the lower atmospheric magnetic structure is frequently believed to be mostly formed by small magnetic flux tubes of circular cross-section emerging from the photosphere and expanding upwardly in the corona (Solanki, Inhester & Sch\u00fcssler 2006). The photosphere and chromosphere are known to be dominated by acoustic-gravity waves with observed short periods of \u223c[3\u20135] min. Whereas in the lower corona a wider range of periods are present (up to about 80 min \u2013 e.g. Sakurai et al. 2002\u2013), short periods ([3\u20135] min) have also been detected in certain coronal structures, e.g. coronal loops and intense magnetic flux tubes (Srivastava & Dwivedi 2010; Jess et al. 2012; Reznikova et al. 2012). A rich range of periods are found in sunspots, the periods becoming larger as the distance from the sunspot umbra increases; a likely result of the influence of the magnetic field (Bogdan & Judge 2006; Jess et al. 2013; Yuan et al. 2014; Madsen et al. 2015).","Citation Text":["Pesnell, Thompson & Chamberlin 2012"],"Functions Text":["As a result of the analysis of 3 min oscillations detected in observations from the Atmospheric Imaging Assembly instrument","on-board the Solar Dynamic Observatory (SDO;"],"Functions Label":["Background","Background"],"Citation Start End":[[741,776]],"Functions Start End":[[547,670],[696,740]]}
{"Identifier":"2021AandA...647A.144A__Bretherton_1966_Instance_2","Paragraph":"Understanding how inertial waves interact with co-rotation resonances is thus a key issue in quantifying tidal dissipation, especially since waves may deeply interact with the background flow at this particular location, which in turn may alter the background flow (as was proposed first by Eliassen & Palm 1961, for terrestrial mountain waves). In binary systems and for late-type stars, Goldreich & Nicholson (1989) showed that the angular momentum transported by gravity waves and exchanged at co-rotation can lead to the successive synchronisation of the layers, from the base to the top of the radiative envelope. More generally, a body of work in various domains, from astrophysical disks (e.g. Goldreich & Tremaine 1979; Baruteau & Masset 2008; Latter & Balbus 2009; Tsang & Lai 2009) to geophysical fluid dynamics (e.g. Bretherton 1966; Yamanaka & Tanaka 1984), has tried to understand the properties of wave propagation and dissipation around co-rotation and, more generally, at all special locations in fluids that correspond to singularities in the linear wave propagation equation. We will refer to them as \u2018critical levels\u2019 in the following (Maslowe 1986), or as \u2018critical layers\u2019 in the case of a viscous medium. This distinction is analogous to the distinction between shear layers and attractors of characteristics that are kinds of singularities for the governing equation of inertial waves in a spherical shell. The aforementioned singularities can act very differently, with either severe absorption at the critical level (as in Booker & Bretherton 1967, for stratified vertical shear flows) or no attenuation if the wave propagates in a peculiar direction (Jones 1967; Acheson 1972; Grimshaw 1975a, for stratifiedvertical shear flows with rotation and magnetism). In other cases, a critical level may even give rise to wave amplification under certain conditions related to the first and second derivatives of the mean flow velocity (Lindzen & Tung 1978; Lindzen & Barker 1985, for barotropic and stratified shear flows, respectively). These studies all used an invariant quantity (the Reynolds stress or thewave action for rotating or magnetic flows) as a diagnostic tool to interpret the role of the critical level in terms of energy transmission and to quantify exchanges between the wave and the mean flow (Eliassen & Palm 1961; Bretherton 1966).","Citation Text":["Bretherton 1966"],"Functions Text":["These studies all used an invariant quantity (the Reynolds stress or thewave action for rotating or magnetic flows) as a diagnostic tool to interpret the role of the critical level in terms of energy transmission and to quantify exchanges between the wave and the mean flow"],"Functions Label":["Background"],"Citation Start End":[[2353,2368]],"Functions Start End":[[2056,2329]]}
{"Identifier":"2018AandA...618A.145O__Codella_et_al._2016_Instance_1","Paragraph":"The chemical composition of protostellar envelopes and their properties along the evolutionary stage of protostars is an important topic in astrochemistry. Since the pioneering work by Cazaux et al. (2003) and Sakai et al. (2008), systematic chemical studies of solar-type protostars (see Ceccarelli et al. 2007; Caselli & Ceccarelli 2012 for a review; also Lefloch et al. 2018) have identified two classes of objects. The first class corresponds to the so-called \u201chot corinos\u201d, that is, sources which display a rich content in complex organic molecules (COMs) in the central inner regions of the protostellar envelope (see Ceccarelli et al. 2007 for a review; also Taquet et al. 2015). Only a few hot corinos have been identified so far either with single dish or interferometric observations: IRAS16293-2422 (Cazaux et al. 2003; Bottinelli et al. 2004b; J\u00f8rgensen et al. 2011, 2016), IRAS2, IRAS4B (Bottinelli et al. 2007), IRAS4A (Bottinelli et al. 2004a; Taquet et al. 2015), HH212 (Codella et al. 2016), L483 (Oya et al. 2017), B335 (Imai et al. 2016), SVS13A (Bianchi et al. 2017), Serpens SMM1, and SMM4 (\u00d6berg et al. 2011). We note that very few sources were investigated in a systematic manner meaning that the COM budget in hot corino sources is very inhomogeneous, making a general picture difficult to come by. Hot corinos share some similarities with the hot cores observed around high-mass stars but they are not scaled-down versions of these. Bottinelli et al. (2007) showed that the abundances of O-bearing species scaled to methanol are higher than those measured in hot cores by one to two orders of magnitude or more. The second chemical class of protostars corresponds to the so-called Warm Carbon Chain Chemistry (WCCC) sources, which have a rich content in C-chains but are poor in COMs. A recent survey of a sample of 36 Class 0\/I protostars of the Perseus molecular cloud complex by Higuchi et al. (2018) shows that the majority of the sources observed have intermediate characters between these two distinct chemistry types.","Citation Text":["Codella et al. 2016"],"Functions Text":["Only a few hot corinos have been identified so far either with single dish or interferometric observations:","HH212"],"Functions Label":["Background","Background"],"Citation Start End":[[987,1006]],"Functions Start End":[[687,794],[980,985]]}
{"Identifier":"2022AandA...664A.117D__Peters_et_al._2015_Instance_1","Paragraph":"We mentioned that our light curves are characterized by two large gaps of one year and seven months plus eight months, as shown in Table 1. From the table, it is apparent that mean and median observed baseline values, computed for the sources in the main sample for each season, are very close to the maximum observed baseline (when not exactly coincident with it) for the corresponding season. Similarly, the mean and median number of visits for each season are very close to the total number of visits (when not exactly coincident with it) for the corresponding season. This shows how, for individual seasons, our dataset can take advantage of a dense sampling, which plays a key role in the context of AGN detection efficiency, as shown in De Cicco et al. (2019). Nevertheless, the two gaps affect the shape of our SF with inadequate sampling in correspondence with some timescales. Sparse and\/or irregular sampling is a very common issue in SF analysis (e.g., de Vries et al. 2003; Peters et al. 2015; Simm et al. 2016; Sartori et al. 2019). Indeed, there are works from the literature where the cadence is low but the sampling is regular: as an example, Hawkins (2002) uses quasar light curves from a long-term monitoring program with 24 yearly observations per source to investigate the origin of the emission mechanism in AGN. Nonetheless, one of the advantages in the use of the SF is its relative insensitivity to irregular sampling when sources are considered as an ensemble rather than individually (e.g., Hawkins 2007; Koz\u0142owski 2016; Sartori et al. 2019). As mentioned in Sect. 3.2, Bauer et al. (2009) analyze the effect of irregular sampling by means of simulations, and conclude that the turnover that is observed in the light curve is an effect of sparse sampling at longer timescales, and not a real feature in the SF of AGN. Emmanoulopoulos et al. (2010) also resort to simulations in order to assess whether and to what extent the SF is robust against the presence of gaps in the light curves. They simulated a single light curve and then investigated the effect of three different gaps in the data, representing three different situations: almost periodic data gaps, dense and sparse sampling, and purely sparsely sampled data, corresponding to 57%, 83%, and 92% of the data being removed from a single simulated light curve that is 2000 time units long, respectively. They then used bootstrapping to extract 1000 light curves from each of the obtained light curves with gaps, then compared the results obtained with and without gaps. While they find that the presence of gaps is responsible for the presence of wiggles and bends in the SF, from the right panels of their Fig. 12 we can infer that these wiggles and bends do not alter the slope of the SF obtained from the light curve with no gaps. In this work we are not investigating the turnover as our baseline is not long enough (Sect. 3.2); our analysis is instead focused on the linear region of the SF (and the possible dependence on physical quantities of interest), where the irregularity of the sampling does not constitute a major issue.","Citation Text":["Peters et al. 2015"],"Functions Text":["Sparse and\/or irregular sampling is a very common issue in SF analysis"],"Functions Label":["Motivation"],"Citation Start End":[[986,1004]],"Functions Start End":[[886,956]]}
{"Identifier":"2019AandA...631A.106S__Bord\u00e9_&_Traub_2006_Instance_1","Paragraph":"With a priori knowledge of speckle evolution lifetime (Milli et al. 2016), more evolved a posteriori algorithms may well calibrate the speckle pattern. However, any such method can directly benefit from an active technique that minimizes the static or quasi-static speckles in each science image during an observation. Active suppression of these speckles requires measurement of the electric field associated with the speckles directly from a coronagraphic image using a focal plane wavefront sensor (FPWFS). Several FPWFSs have been proposed such as phase diversity (Bord\u00e9 & Traub 2006; Give\u2019on et al. 2007; Sauvage et al. 2012) and the self-coherent camera (SCC; Baudoz et al. 2006, 2012). Once the electric field is measured, one or several deformable mirrors (DM) can then be used to minimize the speckle intensity in a region of the image called the dark hole (Malbet et al. 1995). Very encouraging laboratory results have been obtained on the coronagraphic testbeds simulating the space-related environment. The stellar speckle intensity is shown to be reduced by a factor of up to 105 (Belikov et al. 2007; Trauger et al. 2011, 2012; Mazoyer et al. 2013, 2014) which would enable the detection of planets 1010 times fainter than their host star. Several attempts on ground-based Extreme-AO instruments have been performed with moderate results. The stellar speckle intensity has been suppressed by only a factor of up to ten (Savransky et al. 2012; Martinache et al. 2014; Bottom et al. 2017; Matthews et al. 2017; Wilby et al. 2017; Vigan et al. 2019; Galicher et al. 2019) reaching contrast levels of roughly 10\u22126. Most of these techniques temporally modulate the speckle intensity to measure their phase and at least three images are needed for each estimation of the electric field. The quasi-static speckles that evolve faster than every four images cannot be correctly estimated and therefore set a limitation for the speckle estimation. Addressing this concern, our team proposed the SCC that spatially modulates the speckle intensity so that the associated complex electric field can be measured in every science image. The drawback is that a finer sampling of the coronagraphic image is required as compared to the temporal modulation techniques. The SCC has been developed and rigorously tested in space-related environments on the THD2 bench at the Paris Observatory (Baudoz et al. 2018).","Citation Text":["Bord\u00e9 & Traub 2006"],"Functions Text":["Active suppression of these speckles requires measurement of the electric field associated with the speckles directly from a coronagraphic image using a focal plane wavefront sensor (FPWFS). Several FPWFSs have been proposed such as phase diversity"],"Functions Label":["Background"],"Citation Start End":[[569,587]],"Functions Start End":[[319,567]]}
{"Identifier":"2015MNRAS.447.3243M__Middleton_et_al._2011b_Instance_1","Paragraph":"Ultraluminous X-ray sources (ULXs) have been widely observed in the local Universe, with inferred isotropic luminosities above 1039 erg s\u22121 (Roberts 2007; Feng & Soria 2011). Those below \u223c3 \u00d7 1039 erg s\u22121 can be readily associated with accretion on to stellar mass black holes (BHs) (\u223c10 M\u2299) accreting close to or at their Eddington limit (see Sutton, Roberts & Middleton 2013, and references therein). There is now strong evidence to support this assertion, with the discovery of extremely bright ballistic jets from a ULX in M31 (Middleton et al. 2013; Middleton, Miller-Jones & Fender 2014b), which unambiguously links the flow with Eddington rate accretion (Fender, Belloni & Gallo 2004), and the first dynamical mass measurement of the compact object in a ULX, from M101 ULX-1 (Liu et al. 2013). Observations of such \u2018low-luminosity\u2019 ULXs (Middleton et al. 2011b, 2012; Kaur et al. 2012; Soria et al. 2012) have revealed changes in the disc emission that may imply the creation of a radiation pressure supported, larger scaleheight flow in the inner regions (Middleton et al. 2012) or magnetic pressure support (Straub, Done & Middleton 2013). Although emission below \u223c2 keV is generally heavily photoelectrically absorbed by material in the Galactic plane (e.g. Zimmermann et al. 2001), similar spectral behaviour may also be seen in a small number of Galactic BH X-ray binaries (BHBs) at high rates of accretion (e.g. Ueda, Yamaoka & Remillard 2009; Uttley & Klein-Wolt, in preparation). Such \u2018extreme\u2019 high state BHBs probably dominate the ULX population (Walton et al. 2011), yet a significant number of ULXs can still be found at higher luminosities. Those above 1041 erg s\u22121 are dubbed hyperluminous X-ray sources (HLXs; Gao et al. 2003) and provide the best evidence (Farrell et al. 2009; Davis et al. 2011; Servillat et al. 2011; Webb et al. 2012) for a population of intermediate-mass BHs (IMBHs; Colbert & Mushotzky 1999). Such IMBHs (with masses above those expected from direct stellar collapse: >100s of M\u2299) could potentially be formed in globular clusters (Miller & Hamilton 2002, but see Maccarone et al. 2007), through capturing and tidally stripping a dwarf galaxy (King & Dehnen 2005) or mergers in young super star clusters (Portegies-Zwart, McMillan & Gerhard 2003; Portegies-Zwart et al. 2004).","Citation Text":["Middleton et al. 2011b"],"Functions Text":["Observations of such \u2018low-luminosity\u2019 ULXs","have revealed changes in the disc emission that may imply the creation of a radiation pressure supported, larger scaleheight flow in the inner regions","or magnetic pressure support"],"Functions Label":["Background","Background","Background"],"Citation Start End":[[845,867]],"Functions Start End":[[801,843],[912,1062],[1087,1115]]}
{"Identifier":"2018ApJ...856...19N__Isern_et_al._1991_Instance_1","Paragraph":"Models of the ECSN progenitor cores suggest the onset of the electron-capture instability occurs at a unique ONeMg core mass in the mass range of 1.366\u20131.377 M\u2299. (Miyaji et al. 1980; Nomoto 1984, 1987; Podsiadlowski et al. 2005; Takahashi et al. 2013). Electron captures cause the core to contract, and O and Ne burning is ignited in the central regions and propagates outwards in a deflagration front (Schwab et al. 2015), processing material to nuclear statistical equilibrium, where further electron captures and photdissociation accelerates the collapse (Miyaji et al. 1980; Nomoto 1987; Takahashi et al. 2013). Whether the core collapses or the deflagration disrupts the core depends sensitively on the ignition density (Isern et al. 1991; Jones et al. 2016). If the core does collapse, the explosion proceeds via delayed explosion on short timescales (Mayle & Wilson 1988; Kitaura et al. 2006; Fischer et al. 2010), and 2D simulations suggest the explosion occurs before significant convection has had time to develop (Wanajo et al. 2011) and hence a symmetric explosion results. This, coupled with the steep density gradient at the core surface, leads to very little mass loss from the core; estimates of mass loss include of order 10\u22123 M\u2299 (Podsiadlowski et al. 2005), 10\u22122 M\u2299 (Kitaura et al. 2006), and 1.39 \u00d7 10\u22122 M\u2299 (1.14 \u00d7 10\u22122 M\u2299) for the 1D (2D) models of Wanajo et al. (2009, 2011). Therefore the ONeMg progenitor core mass is a good estimate of the baryon mass MB of the resulting NS (Podsiadlowski et al. 2005). Indeed, PSR J0737-3039A and the companion to PSR J1756-2251 have gravitational masses consistent with baryon masses \u223c1.37 M\u2299 when their gravitational binding energies are taken into account (Lattimer & Yahil 1989). Population synthesis calculations incorporating the various binary evolution channels that might lead to production of NSs via ECSNe show that J0737-3039B most likely formed in an ECSN, and the companion to PSR J1756-2251 is consistent with such a formation scenario (Andrews et al. 2015). Other systems with candidates for ECSNe formation also exist (Keith et al. 2009; Chen et al. 2011).","Citation Text":["Isern et al. 1991"],"Functions Text":["Whether the core collapses or the deflagration disrupts the core depends sensitively on the ignition density"],"Functions Label":["Uses"],"Citation Start End":[[726,743]],"Functions Start End":[[616,724]]}
{"Identifier":"2019MNRAS.482..560M__Kong_et_al._2004_Instance_1","Paragraph":"An empirical expression for FUV extinction based on UV reddening of a diverse, UV-selected sample of 200 galaxies (Seibert et al. 2005) is\n(10)\r\n\\begin{eqnarray*}\r\nA_{\\rm FUV}(\\beta)=3.978(m_{\\rm FUV}-m_{\\rm NUV})+0.143,\r\n\\end{eqnarray*}\r\nwhere mFUV and mNUV are the respective GALEX AB magnitudes.13 The equation is similar to relations derived by others (e.g. Hao et al. 2011). Fig. 9 shows the SFRS galaxies in the IRX\u2013(mFUV \u2212 mNUV) (or equivalently IRX\u2013\u03b2) space. There is a correlation between AFUV(\u03b2) and AFUV(IRX) with Pearson correlation coefficient r = 0.71 and mean \u2329AFUV(\u03b2) \u2212 AFUV(IRX)\u232a = 0.33\u2009mag, but the rms scatter in AFUV(IRX) as derived from AFUV(\u03b2) is 0.44\u2009dex. Galaxies with $A_{\\rm FUV}(\\beta)\\lesssim 2$ can have bolometric extinctions as high as 6\u2009mag, and AFUV(\u03b2) applied to LFUV greatly underestimates their FIR luminosity and therefore SFR. This is consistent with other results (e.g. Kong et al. 2004; Johnson et al. 2006, 2007), which have shown that galaxies having higher current SFR relative to their past averaged SFR are likely to deviate above the IRX\u2013\u03b2 relation, i.e. have larger AFUV(IRX) for a given AFUV(\u03b2). Despite this qualitative agreement, the Kong et al. mean numerical relation for their UV-selected sample of 50 local starbursts is not a good fit to the FIR-selected SFRS data as shown in Fig. 9. Regardless of numerical values, all these studies agree that galaxies with higher SFR are more obscured at fixed \u03b2 (also see Iglesias-P\u00e1ramo et al. 2004; Cortese et al. 2006; Moore et al. 2010). At the low SFR end, galaxies with L60  109.3\u2009L\u2299, which at $z$ \u2248 0 are mostly early-type cluster galaxies, form two groups. Around 75 per\u2009cent of them are near the mean IRX\u2013\u03b2 relation, but the rest show AFUV(IRX) \u226a AFUV(\u03b2). One possibility is that these galaxies have older stellar populations with intrinsically high values of \u03b2. In the middle range 109.3 L60  1010.7\u2009L\u2299, there is a general trend for AFUV(IRX) to follow AFUV(\u03b2) but with rms scatter \u223c0.34\u2009dex. At L60 > 1010.7\u2009L\u2299, the scatter is \u223c0.56\u2009dex.","Citation Text":["Kong et al. 2004"],"Functions Text":["This is consistent with other results (e.g.","which have shown that galaxies having higher current SFR relative to their past averaged SFR are likely to deviate above the IRX\u2013\u03b2 relation, i.e. have larger AFUV(IRX) for a given AFUV(\u03b2).","Despite this qualitative agreement, the Kong et al. mean numerical relation for their UV-selected sample of 50 local starbursts is not a good fit to the FIR-selected SFRS data as shown in Fig. 9."],"Functions Label":["Similarities","Similarities","Differences"],"Citation Start End":[[909,925]],"Functions Start End":[[865,908],[955,1143],[1144,1339]]}
{"Identifier":"2020ApJ...888..118M__Soto_et_al._2013_Instance_1","Paragraph":"To determine foreground and background contamination by star-forming galaxies, AGNs, shock emission, and extended PAH emission, as well as field stars, we performed the infrared color selection method described in Gutermuth et al. (2009). In Phase I of this method, we used only GLIMPSE sources that have photometry in all four IRAC bands, and have photometric uncertainties \u03c3  0.2 mag in all four bands, which corresponds to a total of 2723 sources. In Figure 4 we show CCDs from the first step of the Phase I selection method, that allowed us to identify contamination from star-forming galaxies and AGNs. Then, we proceeded to the elimination of shock emission and extended PAH emission contamination, as well as YSO class selection as shown on Figure 5. This procedure resulted in a total of 2702 Spitzer\/GLIMPSE sources without contamination. We obtained a total of 42 sources classified as Class I, and 177 sources classified as Class II. We continued with a process similar to the Phase II selection method described by Gutermuth et al. (2009), which we slightly modified due to the use of VISTA\/VVV data instead of 2MASS data used in the Gutermuth et al. (2009) paper. In this step, both Spitzer\/GLIMPSE and VISTA\/VVV data were used. Specifically, the Phase II selection method is applied to Spitzer\/GLIMPSE sources that lack [5.8] and\/or [8.0] detections. The Spitzer\/GLIMPSE sources were first matched with their VISTA\/VVV counterparts. We only selected high-quality VISTA\/VVV detections with \u03c3  0.1 mag, and Spitzer\/GLIMPSE detections with photometric uncertainties \u03c3  0.2 mag in the detected bands for this analysis. A cross-match of both catalogs was done using a search radius of 1\u2033, to create a matched list of 13,309 sources. A further selection on these sources was done using their VISTA\/VVV magnitudes. We excluded stars with saturated photometry (Soto et al. 2013), by limiting the magnitudes of the detected VISTA\/VVV sources to 13.8, 12.8, and 12.8, for the J, H, and KS bands, respectively. A total of 9974 non-saturated sources were selected. Whenever possible we checked for contamination using similar color\u2013color criteria as described for Phase I in Gutermuth et al. (2009), and ended up with a total of 9844 sources without contamination. We selected this sample of 9844 stars where objects with infrared excesses are expected to be found. To select these infrared excess sources, we used the dereddened color selection criteria detailed in Gutermuth et al. (2009) with [K\u2013[3.6]]0 and [[3.6]\u2013[4.5]]0 colors. Displaying the position of all the new infrared excess candidates showed that they are uniformly distributed in the region, thus we consider that our selected data set might still suffer contamination by foreground sources. To eliminate any possible chance of selecting any contaminant, we therefore applied additional selection criteria to these infrared excess sources candidates, as described in Winston et al. (2011) and Megeath et al. (2012), by plotting two CCDs, as shown in Figure 6. The first selection was done using the CCD J\u2013H versus H\u2013[4.5], with the slope of the reddening band \n\n\n\n\n\n (Chen et al. 2013). The second color selection was done using the CCD H\u2013KS versus KS\u2013[4.5] with the slope of the reddening band \n\n\n\n\n\n (Chen et al. 2013). We defined as YSOs all sources passing the selection criteria of the first or of the second CCD, and found a total of 1566 YSOs. Initial inspection of this selection indicated that there was still contamination by normally reddened stars, and the limits were moved by 0.2 mag toward the red compared with the formulas given by Winston et al. (2007) and and Jose et al. (2016). To then isolate Class I and Class II YSOs from these sources, we used the Phase II color selection from Gutermuth et al. (2009), with the CCD [KS\u2013[3.6]]0 versus [[3.6]\u2013[4.5]]0 as shown in Figure 7. We obtained a total of 32 Class I YSOs and a total of 385 Class II YSOs.","Citation Text":["Soto et al. 2013"],"Functions Text":["We excluded stars with saturated photometry","by limiting the magnitudes of the detected VISTA\/VVV sources to 13.8, 12.8, and 12.8, for the J, H, and KS bands, respectively."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1867,1883]],"Functions Start End":[[1822,1865],[1886,2013]]}
{"Identifier":"2019ApJ...875..129K__Kovtyukh_&_Andrievsky_1999_Instance_1","Paragraph":"Finally, the final estimates depend slightly on whether very strong lines with X > \u22126 are used or not. In Figure 7, very strong lines clearly show a systematic tilt. These strong lines have an impact on the slopes, e.g., seen in Figure 4. The lower \n\n\n\n\n\n values of the stronger lines in MB99 would give higher \n\n\n\n\n\n values with a fixed \u03be, but this would also cause a tilt in Figure 4. A larger \u03be is therefore required so that \n\n\n\n\n\n values of strong and weak lines get balanced. While this is an important difference between the two line lists, generally speaking, it is suggested that using very strong lines often introduces complications such as non-LTE effects into a chemical abundance analysis (e.g., Kovtyukh & Andrievsky 1999; Gratton et al. 2006; Takeda et al. 2013). Based on synthetic spectra, we found that, in the case of lines with X \u2273 \u22126, the line core does not grow any more with increasing metallicity and the damping wing starts to contribute to the EW at around the solar metallicity. If we run the bootstrap method with the same lines but including those with X > \u22126, we obtain moderately different results for the MB99 list, as illustrated in Figure 5. Four lines from MB99 have X > \u22126, and including them leads to higher \u03be and lower \n\n\n\n\n\n values: \n\n\n\n\n\n = (1.47 \u00b1 0.18, 6.94 \u2213 0.05) for Arcturus and (1.61 \u00b1 0.16, 7.71 \u2213 0.06) for \u03bc Leo. The changes caused by including the strongest lines are marginally significant, 1\u20132\u03c3, for the former but are negligible for the latter. Figure 6 shows that one line, Fe i \u03bb 11973.04, with the largest \n\n\n\n\n\n has a particularly strong impact on the slope in the X versus \n\n\n\n\n\n diagram for Arcturus with MB99. The same line gives \n\n\n\n\n\n \u223c 8.10 dex, which is also higher than the average, for \u03bc Leo. However, the scatter of \n\n\n\n\n\n from lines within the low-X range is large, which explains the relatively small effect of including the high-X lines for \u03bc Leo. In contrast, six VALD3 lines that we selected have X > \u22126, but including them has a negligible impact on the \n\n\n\n\n\n measurements. For VALD3, the Fe i \u03bb 11973.046 line leads to \n\n\n\n\n\n values that are very close to the average abundances from other lines for both Arcturus and \u03bc Leo. This line corresponds to the rightmost point in Figure 7 and has a very large difference, 0.8 dex, between the \n\n\n\n\n\n values in the two line lists. Considering these complications, we decided to adopt the \n\n\n\n\n\n values obtained without the lines at X > \u22126 as our best estimates. Although the \n\n\n\n\n\n from individual lines depend on \u03be as described above, we found that the [Fe\/H] obtained in different works are not correlated with \u03be (Figure 8). This is probably because systematic differences in previous works, such as differences in line lists and atmosphere models, introduced a scatter larger than the expected correlation between the two parameters.","Citation Text":["Kovtyukh & Andrievsky 1999"],"Functions Text":["While this is an important difference between the two line lists, generally speaking, it is suggested that using very strong lines often introduces complications such as non-LTE effects into a chemical abundance analysis (e.g.,"],"Functions Label":["Uses"],"Citation Start End":[[709,735]],"Functions Start End":[[481,708]]}
{"Identifier":"2021AandA...647A.137J__Beccari_et_al._2017_Instance_1","Paragraph":"The advent of the ESA Gaia satellite for the first time allows the study of young stellar populations on large scales (> 100 pc) in six-dimensional (6D) position and velocity space. For example, using Gaia Data Release 2 (DR2) (Gaia Collaboration 2018a), Kounkel et al. (2018) and Zari et al. (2019) demonstrated that the Orion star-forming region is composed of a variety of populations with different spatial and kinematic properties each that are all likely generated in multiple events instead of in a progressive star formation history (see also Beccari et al. 2017; Jerabkova et al. 2019b; Kroupa et al. 2018, discussing three bursts of star formation in the Orion Nebula cluster). Zari et al. (2018) described the 6D properties and age structures in which young stars are found within a region of 500 pc around the Sun, and a number of other studies focusing on individual star-forming regions have been reported. One example is the large-scale picture of the Gamma Velorum region (Beccari et al. 2018; Cantat-Gaudin et al. 2019). The Gaia DR2 catalogue also led to the discovery of stellar relic filaments. These are spatial structures of a few pc that are more than 90 pc long and consist of stars of equal ages that are younger than a few hundred million years (Jerabkova et al. 2019a; Beccari et al. 2020). A Galaxy-wide survey using the Gaia DR2 catalogue reveals many elongated structures made of coeval stars, many of which are tidal tails of dissolving star clusters (Kounkel & Covey 2019). Meingast et al. (2021) described the nearby extended spatial distribution of stars that co-move with their clusters but are not bound to them. After the release of the Gaia DR2 catalogue, tidal tails have been found around four nearby ( 300\u2006pc) OCs, namely Blanco 1 (\u2248100\u2006 Myr, Zhang et al. 2020), the Hyades (\u2248600\u2005\u2212\u2005700\u2006Myr; Meingast & Alves 2019; R\u00f6ser et al. 2019; Douglas et al. 2019; Gossage et al. 2018; Reino et al. 2018; Lodieu 2020; Gaia Collaboration 2018b), Coma Berenices (\u2248750\u2006Myr; Tang et al. 2019; F\u00fcrnkranz et al. 2019), and Praesepe (\u2248800\u2006Myr; R\u00f6ser & Schilbach 2019).","Citation Text":["Beccari et al. 2017"],"Functions Text":["For example, using Gaia Data Release 2 (DR2)","Kounkel et al. (2018) and Zari et al. (2019) demonstrated that the Orion star-forming region is composed of a variety of populations with different spatial and kinematic properties each that are all likely generated in multiple events instead of in a progressive star formation history (see also","discussing three bursts of star formation in the Orion Nebula cluster)."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[551,570]],"Functions Start End":[[182,226],[255,550],[616,687]]}
{"Identifier":"2020MNRAS.498.6013A__Tonry_1998_Instance_1","Paragraph":"In this section, we describe our new compilation of SLS. To construct Dobs, we have chosen only systems with spectroscopically data well measured from different surveys. We have considered 19 SLS from the CASTLES, 107 from SLACS, 38 from BELLS, 4 from LSD, 35 from SL2S, and 1 system from the DES survey. The final list has a total of 204 systems, being the largest sample of SLS to date. We use spectroscopy to select those lenses with lenticular (S0) or elliptical (E) morphologies that have been modelled assuming a SIS ($\\sim 3{{\\ \\rm per\\ cent}}$) or SIE ($\\sim 97{{\\ \\rm per\\ cent}}$) lens model. Many systems have not been taken into account due to several issues. For instance, the system PG1115+080 (Tonry 1998) from the CASTLES survey has been discarded because the lens mass model is steeper than isothermal. In addition, the system MGJ0751+2716 (Spingola et al. 2018) was also discarded because the main lens belongs to a group of galaxies. From the SLACS survey (Bolton et al. 2008; Auger et al. 2009), we remove the systems SDSSJ1251\u22120208, SDSSJ1432+6317, SDSSJ1032+5322, and SDSSJ0955+0101 since the lens galaxies are late-type. The same reason is applied to the systems SDSSJ1611+1705 and SDSSJ1637+1439 from the BELLS survey (Shu et al. 2016). We have also discarded the systems SDSSJ2347\u22120005 and SDSSJ0935\u22120003 from the SLACS survey and the system SDSSJ111040.42+364924.4 from the BELLS survey because they have large measured velocity dispersions (\u223c400 km s\u22121 or bigger values), suggesting the lens might be part of a group of galaxies or that there is substructure in the line of sight. For those systems without reported velocity dispersion error, we assumed the average error of the measurements in the survey subsample as follows. For the nine systems from CASTLES, we consider the average error on \u03c3 for this survey, i.e. 14  per\u2009cent. In the case of the system DES J2146\u22120047 (Agnello et al. 2015), we have assumed a 10 per\u2009cent error on \u03c3, which is the average error of the entire sample. The LSD survey (Koopmans & Treu 2003; Treu & Koopmans 2004) reports \u03c3 corrected by circular aperture using the expression obtained by Jorgensen, Franx & Kjaergaard (1995a,b). A close inspection of the \u03c3spec values, with and without aperture correction, presented by Cao et al. (2015) show the differences is smaller than the reported error. Thus, we decided to use the observed values (\u03c3) and the reported error for the sample without the aperture correction. On the other hand, in those systems in which the Einstein radius error was not reported, we followed Cao et al. (2015) and assumed an error of \u03b4\u03b8E = 0.05, which is the average value of the systems with reported errors in this sample.","Citation Text":["Tonry 1998"],"Functions Text":["Many systems have not been taken into account due to several issues. For instance, the system PG1115+080","from the CASTLES survey has been discarded because the lens mass model is steeper than isothermal."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[709,719]],"Functions Start End":[[603,707],[721,819]]}
{"Identifier":"2019AandA...629A.134G__Groh_et_al._2008_Instance_1","Paragraph":"The shape of the spectral energy distribution and the emission rates of ionizing photons (see Table 1 of Paper II) depend on the assumed wind mass-loss rates, wind speeds, and wind clumping. These parameters are uncertain. Theoretical predictions are now available (e.g., Krti\u010dka et al. 2016; Vink 2017), but they have not yet been thoroughly tested against observations, because only very few stripped stars with sufficiently strong wind mass-loss have been identified and studied in detail so far (e.g., Groh et al. 2008). In Paper I, we showed that variations in wind mass-loss rate primarily affect the predicted emission rate of He\u202fII-ionizing photons, while the emission rates of H\u202fI- and He\u202fI-ionizing photons are not significantly affected. The mass-loss rates assumed in our models were chosen to smoothly connect the mass-loss rates of subdwarfs (Krti\u010dka et al. 2016) with the observed mass-loss rates of WR stars (Nugis & Lamers 2000). Our assumed mass-loss rates also match well with the observed mass-loss rate of the stripped star in the binary system HD 45166 (Groh et al. 2008). The recent theoretical predictions by Vink (2017) suggest that the mass-loss rates of stripped stars may be ten times lower than what we assume in this paper. The winds of stripped stars are likely not reaching close to the Eddington limit, in contrary to massive main-sequence and WR stars (cf. Bestenlehner et al. 2014). This suggests that the wind mass-loss rate from stripped stars is lower than that from WR stars and thus not well-described by the recipe for WR stars of Nugis & Lamers (2000). To establish which are the wind mass-loss rates from stripped stars, observations of a sample of stripped stars are necessary. If, as suggested by Vink (2017), the mass-loss rates from stripped stars indeed are lower than what the recipe from Nugis & Lamers (2000) predicts, it would likely imply an increase of the emission rates of He\u202fII-ionizing photons presented in this work. The emission rates of H\u202fI- and He\u202fI-ionizing photons are robust against wind uncertainties.","Citation Text":["Groh et al. 2008","Groh et al. 2008"],"Functions Text":["Theoretical predictions are now available","but they have not yet been thoroughly tested against observations, because only very few stripped stars with sufficiently strong wind mass-loss have been identified and studied in detail so far (e.g.,"],"Functions Label":["Background","Background"],"Citation Start End":[[506,522],[1076,1092]],"Functions Start End":[[223,264],[305,505]]}
{"Identifier":"2021AandA...647A.144A__Bretherton_1966_Instance_1","Paragraph":"Understanding how inertial waves interact with co-rotation resonances is thus a key issue in quantifying tidal dissipation, especially since waves may deeply interact with the background flow at this particular location, which in turn may alter the background flow (as was proposed first by Eliassen & Palm 1961, for terrestrial mountain waves). In binary systems and for late-type stars, Goldreich & Nicholson (1989) showed that the angular momentum transported by gravity waves and exchanged at co-rotation can lead to the successive synchronisation of the layers, from the base to the top of the radiative envelope. More generally, a body of work in various domains, from astrophysical disks (e.g. Goldreich & Tremaine 1979; Baruteau & Masset 2008; Latter & Balbus 2009; Tsang & Lai 2009) to geophysical fluid dynamics (e.g. Bretherton 1966; Yamanaka & Tanaka 1984), has tried to understand the properties of wave propagation and dissipation around co-rotation and, more generally, at all special locations in fluids that correspond to singularities in the linear wave propagation equation. We will refer to them as \u2018critical levels\u2019 in the following (Maslowe 1986), or as \u2018critical layers\u2019 in the case of a viscous medium. This distinction is analogous to the distinction between shear layers and attractors of characteristics that are kinds of singularities for the governing equation of inertial waves in a spherical shell. The aforementioned singularities can act very differently, with either severe absorption at the critical level (as in Booker & Bretherton 1967, for stratified vertical shear flows) or no attenuation if the wave propagates in a peculiar direction (Jones 1967; Acheson 1972; Grimshaw 1975a, for stratifiedvertical shear flows with rotation and magnetism). In other cases, a critical level may even give rise to wave amplification under certain conditions related to the first and second derivatives of the mean flow velocity (Lindzen & Tung 1978; Lindzen & Barker 1985, for barotropic and stratified shear flows, respectively). These studies all used an invariant quantity (the Reynolds stress or thewave action for rotating or magnetic flows) as a diagnostic tool to interpret the role of the critical level in terms of energy transmission and to quantify exchanges between the wave and the mean flow (Eliassen & Palm 1961; Bretherton 1966).","Citation Text":["Bretherton 1966"],"Functions Text":["More generally, a body of work in various domains,","to geophysical fluid dynamics (e.g.","has tried to understand the properties of wave propagation and dissipation around co-rotation and, more generally, at all special locations in fluids that correspond to singularities in the linear wave propagation equation."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[828,843]],"Functions Start End":[[619,669],[792,827],[870,1093]]}
{"Identifier":"2018MNRAS.473.2000T__Noutsios_et_al._2011_Instance_1","Paragraph":"The launch of the Fermi Gamma-ray Space Telescope has spurred on the search for pulsars in \u03b3-rays (Grenier & Harding 2015), yielding over 2001 detections and triggering multiwavelength observations. While pulsars are common targets in the X-rays, they are very challenging targets in the optical and very few of them have been identified (see Mignani et al. 2016, and references therein). Here, we report on Large Binocular Telescope (LBT) observations of an isolated pulsar, PSR\u2009J2043+2740 (Taylor, Manchester & Lyne 1993), detected by both AGILE (Pellizzoni et al. 2009) and Fermi (Abdo et al. 2010; Noutsios et al. 2011). It was discovered as a radio pulsar (Ray et al. 1996) and later on as an X-ray source by XMM\u2013Newton (Becker et al. 2004), although X-ray pulsations have not yet been found. PSR\u2009J2043+2740 is one of the very few non-recycled pulsars older than 1\u2009Myr detected in \u03b3-rays, with a characteristic age \u03c4c = 1.2\u2009Myr, inferred from the values of its spin period Ps = 0.096\u2009s and its derivative $\\dot{P}_{\\rm s} = 1.27 \\times 10^{-15}$\u2009s\u2009s\u22121 (Ray et al. 1996). This also yields a rotational energy loss rate $\\dot{E}_{\\rm rot} = 5.6 \\times 10^{34}$\u2009erg\u2009s\u22121 and a surface dipolar magnetic field Bs = 3.54 \u00d7 1011 G.2 Although PSR\u2009J2043+2740 does not have a very large spin-down power compared to young (\u223c1\u201310 kyr) pulsars (\u223c1036\u20131038\u2009erg\u2009s\u22121), it is still a factor of 2 larger than that of middle aged \u03b3-ray pulsars (\u223c0.1\u20130.5\u2009Myr), such as Geminga, PSR\u2009B0656+14 and PSR\u2009B1055\u221252, all detected in the optical, thanks to their distances \u2272 500\u2009pc (Abdo et al. 2013). The distance to PSR\u2009J2043+2740 is uncertain owing to the lack of a radio parallax measurement. The radio dispersion measure (DM = 21.0\u2009\u00b1\u20090.1\u2009pc cm\u22123; Ray et al. 1996) gives a distance of 1.8\u2009\u00b1\u20090.3\u2009kpc from the NE2001 model of the Galactic-free electron density (Cordes & Lazio 2002). A slightly smaller distance (1.48\u2009kpc) is inferred from the model of Yao, Manchester & Wang (2017). The hydrogen column density towards the pulsar obtained from the X-ray spectral fits (NH \u2272 3.6 \u00d7 1020\u2009cm\u22122; Abdo et al. 2013) suggests a distance of a few hundred pc (He, Ng & Kaspi 2013), although these estimates depend on the model X-ray spectrum. Such a distance would make PSR\u2009J2043+2740 a viable target for deep optical observations, never carried out until now, and might be compatible with the debated association (Noutsios et al. 2011) with the Cygnus Loop supernova remnant (SNR) at $540^{+100}_{-80}$\u2009pc (Blair, Sankrit & Raymond 2005).","Citation Text":["Noutsios et al. 2011"],"Functions Text":["Here, we report on Large Binocular Telescope (LBT) observations of an isolated pulsar, PSR\u2009J2043+2740","detected by","and Fermi"],"Functions Label":["Background","Background","Background"],"Citation Start End":[[602,622]],"Functions Start End":[[389,490],[525,536],[573,582]]}
{"Identifier":"2015AandA...582A..88W__Lamberts_et_al._2013_Instance_1","Paragraph":"For completeness, we also supplement our reaction scheme with grain-surface association reactions extracted from the publicly available Ohio State University (OSU) network3 (Garrod et al. 2008). For those species important in grain-surface chemical reaction schemes, e.g., the CH3O radical, which are not included in Rate12, we also extract the corresponding gas-phase formation and destruction reactions from the OSU network. The grain-surface network has been further updated to include all studied routes to water formation under interstellar and circumstellar conditions (Cuppen et al. 2010a; Lamberts et al. 2013). The grain-surface reaction rates are calculated assuming the Langmuir-Hinshelwood mechanism only, and using the rate-equation method as described in Hasegawa et al. (1992). We limit the chemically \u201cactive\u201d zone to the top two monolayers of the ice mantle. We assume the size of the barrier to surface diffusion is 0.3\u00d7 the binding energy; in this way, volatile species diffuse at a faster rate than strongly bound species. This value lies at the optimistic end of the range determined by recent off-lattice kinetic Monte Carlo simulations of surface diffusion of CO and CO2 on crystalline water ice (Karssemeijer & Cuppen 2014). This allows the efficient formation of complex organic molecules via radical-radical association reactions at \u227320 K (see, e.g., Vasyunin & Herbst 2013; Walsh et al. 2014). For the lightest reactants, H and H2, we use either the classical diffusion rate or the quantum tunnelling rate depending on which is fastest (see, e.g., Tielens & Hagen 1982; Hasegawa et al. 1992). For the latter rates, we follow Garrod & Pauly (2011) and adopt a rectangular barrier of width 1.5 \u00c5. We also include reaction-diffusion competition in which the reaction probability is determined by the relative rates between the barrier-mediated reaction and thermal diffusion (see, e.g., Chang et al. 2007; Garrod & Pauly 2011). Although still relatively simplistic, this takes into account the increased probability of reaction in the limit where the thermal diffusion of the reactants away from a common binding site is slow compared with the barrier-mediated reaction rate. ","Citation Text":["Lamberts et al. 2013"],"Functions Text":["The grain-surface network has been further updated to include all studied routes to water formation under interstellar and circumstellar conditions"],"Functions Label":["Uses"],"Citation Start End":[[597,617]],"Functions Start End":[[427,574]]}
{"Identifier":"2018AandA...612A..77M__Gromadzki_&_Miko\u0142ajewska_(2009)_Instance_3","Paragraph":"\u201cWiggling\u201d outflows are often observed among young stellar jets and protostellar molecular outflows (Eisloffel et al. 1996; Terquem et al. 1999). Terquem et al. (1999) investigated such binary systems where the accretion disk, from which the jet originates, is inclined to the binary orbital plane. They concluded that the observed jet \u201cwiggling\u201d is a consequence of the jet precession caused by tidal interactions in such non-coplanar binary systems. Nichols & Slavin (2009) as well as Hollis & Michalitsianos (1993) suggested that the precession of the accretion disk around the WD may be responsible for the bending of the wide-angle outflow found in the previous studies. In analogy to these observations of young stellar jets, we suggest that the \u201cwiggling\u201d that we also find here for the R Aqr jet may result from disk precession as well. We estimated the precession period using the theory from Terquem et al. (1999) and the binary parameters taken from Gromadzki & Miko\u0142ajewska (2009) \u2013 Mh = 0.8M\u2299 (the mass of the hot WD companion), Mp\u2215Mh = 1.65 (where Mp is the mass of the primary companion), D = 15.5 AU (the mean semi-major axis of the system), and e = 0.25 (eccentricity). The value of Rd is unknown in our case and we used Rd = 5 AU giving D\u2215Rd \u2248 3 which corresponds to the average value of 2 \u2264 D\u2215Rd \u2264 4 (the range taken from Terquem et al. 1999). For this calculation, we assumed that the angle \u03b4 between the disk plane and that of the binary orbit is small enough (10\u00b0) and we adopted cos\u03b4 = 1. Using Eq. (1) from Gromadzki & Miko\u0142ajewska (2009), we derived the precession time of T \u2248 530 yr. This value is quite large for the wiggling waves that we see. We estimated the projected spatial wavelength \u03bbproj of the \u201cwiggling\u201d wave according to \u03bb = \u03bbproj\u2215sin i, where i is the angle between the jet symmetry axis and the line of sight, and T = \u03bb\u2215\u03c5, where \u03c5 is the jet velocity, from Gromadzki & Miko\u0142ajewska (2009). Using i = 72\u00b0 and \u03c5 ~ 100 km\/s, we derive \u03bbproj \u2248 10 500 AU which is more than 20 times larger than the projected length of the observed wiggling outflow (2\u2032\u2032 \u2248 440 AU). However, we should note that the precessing time strongly depends on the D\u2215Rd; the T decreases significantly with increasing R. It may also be the case that the \u201cwiggling\u201d model developed for YSO jets is not fitting for R Aqr which consists of evolved objects, and both the WD and the disk where the jet probably forms are much hotter than YSO systems. Furthermore, we cannot exclude that the steady wiggling might be a sequence of dynamical interactions of the two collimated flows tilted to each other.","Citation Text":["Gromadzki & Miko\u0142ajewska (2009)"],"Functions Text":["We estimated the projected spatial wavelength \u03bbproj of the \u201cwiggling\u201d wave according to \u03bb = \u03bbproj\u2215sin i, where i is the angle between the jet symmetry axis and the line of sight, and T = \u03bb\u2215\u03c5, where \u03c5 is the jet velocity, from"],"Functions Label":["Uses"],"Citation Start End":[[1898,1929]],"Functions Start End":[[1672,1897]]}
{"Identifier":"2016MNRAS.463.3783B__Reid_&_White_2011_Instance_1","Paragraph":"It is well known that a per cent level understanding of the anisotropy of the redshift-space galaxy clustering is needed to accurately recover cosmological information from the RSD signal in order to shed light on the issue of dark energy versus modified gravity. From a statistical point of view, the source of the anisotropy is the galaxy line-of-sight pairwise velocity distribution. It is therefore important to adopt a realistic functional form for this velocity PDF when fitting models to the data. To this purpose, in Paper I we introduced the GG prescription for the velocity PDF. In this work, we have continued the development of this model by making explicit the dependence of the GG distribution on quantities predictable by theory, namely its first three moments, and extending it to the more general concept of GQG. To keep the model as simple as possible, we have proposed an ansatz with two free dimensionless parameters that describe how infall velocity and velocity dispersion vary when moving from one place to another in our Universe. Since their interpretation is clear, these parameters can be theoretically predicted or, assuming a more pragmatic approach, tuned to simulations or used as nuisance parameters. State-of-the-art PT has proven successful in predicting the large-scale behaviour of the velocity PDF and the correspondent monopole and quadrupole of the redshift-space correlation function (e.g. Reid & White 2011; Wang, Reid & White 2014), at least for massive haloes, M \u223c 1013\u2009M\u2299. Unfortunately, by definition, any PT breaks down for small separations. Consequently, alternative approaches have been suggested in the literature, spanning from purely theoretical (e.g. Sheth 1996) to hybrid techniques in which N-body simulations plus an HOD are employed to deal with the issue of non-linearities (e.g. Tinker 2007; Reid et al. 2014). One of the main results from our work is to provide a framework in which perturbation and small-scale theories are smoothly joined, so that all available RSD information can be coherently extracted from redshift surveys. A fundamental requirement for a redshift-space model is that it must be precise on all scales interest, and it should inform the user of the scales on which the model can be trusted. We have compared to N-body simulations the well-known GSM (Reid & White 2011), the more recent ESM (Uhlemann et al. 2015) and the GQG prescription over a broad range of separations, from 0 to 80\u2009h\u22121 Mpc. Different redshifts, from z = 0 to 1, and different tracers, namely DM particles and two mass-selected catalogues of DM haloes, have been considered. We have concluded that, among the three, QGQ is the only model capable of providing a precise redshift-space correlation function on scales down to \u223c5\u2009h\u22121 Mpc over the range of redshifts covered by future surveys. Keeping in mind that the range of validity of the models depends on tracer, redshift and order of the Legendre multipoles we are interested in, for finiteness, we can say that all the models converge to the expected amplitude on scales \u227330\u2009h\u22121 Mpc, at least for multipole and quadrupole. Since these scales roughly coincide with the range of validity of state-of-the-art PTs, if we rely only on PT and if we are not interested in higher order multipoles, the most natural choice is the simplest model among the three, i.e. the GSM. As for the ESM, we have found it to be unbiased down to smaller scales and for higher order multipoles than the GSM, thus confirming the results by Uhlemann et al. (2015), but, on the other hand, it seems to behave even worse than the GSM on the smallest scales. We can therefore think of it as a natural extension of the GSM in the perspective of further PT developments. In particular, a better prediction of the third moment of the velocity PDF is required before the ESM can be applied to data on smaller scales. Formally, the same argument holds for the GQG model, none the less, since this latter is meant to include non-linear scales, it could be possible to obtain a prediction for the third moment by interpolating between (very) small and (very) large scales. More precisely, as shown in the lower-right panel of Fig. A1, the functions $c^{(3)}_t$ and $c^{(3)}_r$, which fully characterize the third moment, are peaked at r \u2272 10\u2009h\u22121 Mpc. By adopting a model for the small-scale limit that includes those separation, most likely using simulations in a similar way to that proposed in Reid et al. (2014), we would then be able to interpolate between these peaks and their large-scale limit, which is trivially 0.","Citation Text":["Reid & White 2011"],"Functions Text":["State-of-the-art PT has proven successful in predicting the large-scale behaviour of the velocity PDF and the correspondent monopole and quadrupole of the redshift-space correlation function (e.g",", at least for massive haloes, M \u223c 1013\u2009M\u2299."],"Functions Label":["Background","Background"],"Citation Start End":[[1430,1447]],"Functions Start End":[[1233,1428],[1473,1516]]}
{"Identifier":"2021MNRAS.507.4389G__Masters_et_al._2011_Instance_3","Paragraph":"Erwin (2018) showed that, in a sample drawn from the Spitzer Survey of Stellar Structure in Galaxies (S4G), the bar fraction is constant over a range of (g \u2212r) colours and gas fractions. Their bar fraction does not increase, but rather decreases for stellar masses higher than \u223c 109.7M\u2299. These results are in contrast to many SDSS-based studies cited above. Erwin (2018) argues that this apparent contradiction can be explained if SDSS-based studies miss bars in low-mass blue galaxies. In Figs 5 and 6, we showed that the newly detected bars in GZD (compared to GZ2) are weak bars in low-mass blue galaxies. Nevertheless, the \u2018combined\u2019 bar fraction in Fig. 6 is not constant over (g \u2212r) colour and agrees well with Masters et al. (2011) for redder colours [(g \u2212r) colour > 0.5]. Additionally, our \u2018combined\u2019 bar fraction remains roughly constant over stellar mass. As mentioned before, we conclude that strong bars drive the trends of bar fraction with (g \u2212r) colour, stellar mass, and SFR observed in other studies (Nair & Abraham 2010b; Masters et al. 2011, 2012; Vera et al. 2016; Cervantes Sodi 2017). However, the addition of weak bars in low-mass blue galaxies is insufficient to resolve the apparent disagreement between Erwin (2018) and many SDSS-based studies (Masters et al. 2011, 2012; Vera et al. 2016; Cervantes Sodi 2017; Kruk et al. 2018), which instead seems likely to be due to the very different sample selection of the S4G and SDSS galaxy samples. For example, the median stellar mass of the sample used in Erwin (2018) is \u223c109.6M\u2299 (based on their Fig. 4 and the bins in the top left-hand panel of their Fig. 5). However, the median stellar mass of our sample is 1010.6M\u2299. As stellar mass correlates with many parameters (including bar length), this can have major consequences. Additionally, as Erwin (2018) notes, there is also the issue of resolution to consider. With an r-band FWHM of 1.18 arcsec from DECaLS (Dey et al. 2019) and a mean redshift of 0.036, the mean linear resolution of our sample is approximately 834 pc, which is higher than the 165 pc of Erwin (2018). This explains why they observe many sub-kpc bars, while we do not. These differences in stellar mass and resolution will manifest themselves in the conclusions, so a more detailed analysis is needed for a proper comparison with Erwin (2018).","Citation Text":["Masters et al. 2011"],"Functions Text":["However, the addition of weak bars in low-mass blue galaxies is insufficient to resolve the apparent disagreement between Erwin (2018) and many SDSS-based studies","which instead seems likely to be due to the very different sample selection of the S4G and SDSS galaxy samples."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1272,1291]],"Functions Start End":[[1108,1270],[1357,1468]]}
{"Identifier":"2020AandA...644A..59K__Heays_et_al._2017_Instance_1","Paragraph":"Analyzing optical emission lines, emanating from within the northern lobe, Tylenda et al. (2019) found a reddening with EB\u2005\u2212\u2005V\u2004\u2248\u20040.9 mag or AV\u2004\u2248\u20042.8 mag, which we assume is mainly circumstellar in origin. Hajduk et al. (2013) observed two stars shining through the southern lobe and found AV\u2004=\u20043.3\u2005\u2212\u20054.4 mag with unknown contribution from the interstellar component. We assume here that those observations quantify the amount of circumstellar dust that is the main actor in shielding molecules from the central source. We recalculated the lifetimes of molecules assuming AV\u2004=\u20043 mag, and with (1) standard dust properties (i.e. with composition and size distribution as of interstellar dust) or (2) with larger and less opaque grains, at the gas-to-dust mass ratio of 124 (see Heays et al. 2017, for more details on the assumed dust properties). We used shielding functions from Heays et al., which include effects in lines. Results are shown in Cols. (3) and (4) of Table 3. The presence of ISM grains makes it possible for the observed molecules to survive for a very long time, longer than 350 yr. The lifetimes in the presence of the large grains considered by Heays et al. are typically a few times shorter than the age of the remnant. It is uncertain what kind of grains populate the dusty remnant of CK Vul, but given its anomalous elemental and molecular compositions and eruptive history, dust may have a peculiar chemical composition and size distribution. In such a case, the total to selective extinction law would also be different and the assumed AV may not be adequate. Nevertheless, if the molecules formed 350 yr ago and are shielded by big grains, with the calculated lifetimes a considerable fraction of molecular species would survive, except perhaps for a few most fragile ones which indeed are almost absent in the lobes. We conclude that the lifetimes in Table 3 that were calculated with an attenuated ISRF are consistent with the molecule formation 350 yr ago or more recently.","Citation Text":["Heays et al. 2017"],"Functions Text":["We recalculated the lifetimes of molecules assuming AV\u2004=\u20043 mag, and with (1) standard dust properties (i.e. with composition and size distribution as of interstellar dust) or (2) with larger and less opaque grains, at the gas-to-dust mass ratio of 124 (see","for more details on the assumed dust properties)."],"Functions Label":["Uses","Uses"],"Citation Start End":[[776,793]],"Functions Start End":[[519,775],[795,844]]}
{"Identifier":"2020MNRAS.493.3045B__Jaisawal_&_Naik_2015a_Instance_3","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"],"Functions Text":["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"],"Functions Label":["Differences"],"Citation Start End":[[2007,2028]],"Functions Start End":[[1825,2005]]}
{"Identifier":"2019AandA...627A.114K__Robitaille_2010_Instance_1","Paragraph":"In the main simulation, we modeled a Cartesian cube 2000 AU in size with 150 grid points in each direction. The cell size of 13.3 AU is comparable to the size of the star, which was represented by a point source. Clumps A to D were simulated as 3D structures of a Gaussian density profile. The choice of a Gaussian distribution was dictated only by the simplicity of the implementation. Whereas observations constrain the clump sizes and overall relative location in two spatial dimensions, the extent of a given feature along the line of sight was a free parameter. We included scattering in the radiative transfer calculations. It turned out early on in the simulations that the dusty medium is very optically thick for visual and ultraviolet (UV) photons making the calculations particularly time-consuming. To speed them up, we used the Modified Random Walk method implemented in RADMC-3D (Robitaille 2010). The number of photons used in the Monte Carlo simulations varied from 105 to 109. For each spatial configuration of dust distribution, first a thermal structure was calculated. Based on this calculation, regions where the dust temperature exceeded the sublimation temperature were removed from the simulation and the thermal structure was calculated once again. In this way we reconstructed images that were further processed in CASA for a direct comparison to ALMA continuum maps. Given a very high number of free parameters (over 40 in the basic version of the simulation), the best configurations were searched through iterative trial-and-error modifications of the models. We aimed to construct only a general 3D model of the dusty environment requiring a minimum dust mass possible. We attempted to reproduce the peak and integrated ALMA fluxes of each feature within the order of magnitude. This aim, however unambitious it may appear, was challenging without increasing further the number of free parameters. For example, better fits to total flux and its distribution would require considering density distributions other than Gaussian.","Citation Text":["Robitaille 2010"],"Functions Text":["To speed them up, we used the Modified Random Walk method implemented in RADMC-3D"],"Functions Label":["Uses"],"Citation Start End":[[894,909]],"Functions Start End":[[811,892]]}
{"Identifier":"2015ApJ...800...24K___2012_Instance_1","Paragraph":"Analyzing the fraction of quenched galaxies for centrals and satellites as a function of stellar mass and environmental parameters has emerged as a fruitful way of gaining insights into the phenomenology and the physical processes of quenching. However, one of the difficulties in interpreting this fraction is that many of the environmental parameters are statistically correlated both with each other or with the stellar mass of the galaxy. For example, there is good evidence that halo mass correlates with the mass of the central galaxy (e.g., Yang et\u00c2 al. 2009), that the local overdensity correlates with the group-centric distance (e.g., Peng et\u00c2 al. 2012; Woo et\u00c2 al. 2013), and that the local overdensity also correlates with halo mass (e.g., Haas et\u00c2 al. 2012; Carollo et\u00c2 al. 2013b). These correlations among the parameters can introduce spurious dependencies, i.e., dependencies which have no direct causal relation to quenching, if the quenched fraction is regarded as a function of certain parameters in isolation. Therefore, in principle one has to study the quenched fraction in the full parameter space and to vary only one parameter at a time, while keeping the others fixed (see, e.g., the discussion in Mo et\u00c2 al. 2010; Section 15.5). However, this approach does also not necessarily guarantee that the measured dependence of the quenched fraction on a certain parameter is directly related to a physical quenching process, since the red fraction may also depends on the history of the galaxy, i.e., how the galaxy is moving through the parameter space. Therefore, interpreting trends of the quenched fraction even in the full parameter space has turned out to be quite difficult and has led to some confusing and apparently inconsistent statements in the literature. In this paper, we use the Sloan Digital Sky Survey (SDSS; York et\u00c2 al. 2000) and the group catalog of Yang et\u00c2 al. (2012) to take a new look at some of the issues that have been recently raised.","Citation Text":["Peng et\u00c2 al. 2012"],"Functions Text":["For example, there is good evidence","that the local overdensity correlates with the group-centric distance (e.g.,"],"Functions Label":["Background","Background"],"Citation Start End":[[645,662]],"Functions Start End":[[443,478],[568,644]]}
{"Identifier":"2015ApJ...813...20B__Kim_2001_Instance_1","Paragraph":"In our investigation we were looking for alignment of galaxy clusters in a sample of 1056 low redshift ACO clusters with known Bautz\u2013Morgan type. We found a statistically significant effect for structures with separation distance 30  R \u2264 45 h\u22121 Mpc. This allowed us to conclude that, for the analyzed clusters, the effect has a range to about 45 h\u22121 Mpc. We found a stronger alignment of elongated clusters of BM type I (an excess of small values of the \u0394\u03b8 angles is observed), having a range to about 45 h\u22121 Mpc. The alignment for BM I is probably connected with the origin of a supergiant galaxy. One should note that during studies of the isolated Abell groups (Flin & Olowin 1991; Trevese et al. 1992; Kim 2001; Niederste-Ostholt et al. 2010), only a rudimentary alignment was found and related only to the brightest cluster members. Gonzalez-Sanchez & Teodoro (2010) interpreted the alignment of just the brightest galaxies within a cluster as the effect of gravitational tidal forces. Correlation between the orientation of the brightest galaxy within a cluster and the cluster\u2019s large axis was also found by Sastry (1968), Carter & Metcalfe (1980), Binggeli (1982), Struble & Peebles (1985), Rhee & Katgert (1987), West (1989, 1994), van Kampen & Rhee (1990), Plionis (1994), Fuller et al. (1999), and Kim et al. (2002). The role of a much more massive central galaxy is related to the problem of galaxy mergers in a cluster. This, in turn, is due to the fact of (generally speaking, random) interaction between cluster members and the dynamic evolution of structures near to specific cluster members. The problem of dynamic evolution can be studied by combined studies of the Binggeli effect, i.e., the relation between the positions of major axes in groups or clusters of galaxies and the directions toward their neighbors, and of the mutual orientation of the brightest galaxy (and other bright galaxies) in a structure relative to the position of cluster major axes or even examination of structure ellipticity redshift dependence, especially in the enlarged observational samples.","Citation Text":["Kim 2001"],"Functions Text":["One should note that during studies of the isolated Abell groups","only a rudimentary alignment was found and related only to the brightest cluster members."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[706,714]],"Functions Start End":[[599,663],[748,837]]}
{"Identifier":"2018AandA...617A..86L__Tian_2017_Instance_1","Paragraph":"The IRIS spectra measure the flare in a \u201csit-and-stare\u201d mode with a roll angle of 45\u2218. The spectral scale is \u223c25.6 m\u00c5 per pixel in the far-ultraviolet (FUV) wavelengths. The IRIS slit crosses the flaring loop and one ribbon (Fig. 1). Two red bars enclose the flaring loop region used to study the quasi-periodic oscillations in this work. IRIS spectrum was pre-processed with the SSW routines of \u201ciris_orbitval_corr_l2.pro\u201d (Tian et al. 2014; Cheng et al. 2015) and \u201ciris_prep_despike.pro\u201d (De Pontieu et al. 2014). To improve the signal-to-noise ratio, we apply a running average over five pixels to the IRIS spectra along the slit (Tian et al. 2012, 2016). We also manually perform the absolute wavelength calibration using a relatively strong neutral line, O\u202fI 1355.60 \u00c5 (see De Pontieu et al. 2014; Tian et al. 2015; Tian 2017). IRIS observations show that Fe\u202fXXI 1354.08 \u00c5 is a hot (\u223c11 MK) and broad emission line and is always blended with many narrow chromospheric lines at the flaring ribbons (Li et al. 2015b, 2016b; Tian et al. 2015; Young et al. 2015; Brosius et al. 2016; Polito et al. 2016). However, the Fe\u202fXXI 1354.08 \u00c5 line is much stronger than those blended emission lines at the flaring loops (Tian et al. 2016). Figure 2a gives the time evolution of the line profiles of Fe\u202fXXI 1354.08 \u00c5, averaged over the slit positions between \u223c18.3\u2033 and 21.6\u2033. Figure 2 panels b\u2212f show the spectral line profiles at the time indicated by the yellow lines in panel a. We can see that only the cool line of C\u202fI 1354.29 \u00c5 is blended with the hot line of Fe\u202fXXI 1354.08 \u00c5, but its contribution is negligible. Therefore, double Gaussian functions superimposed on a linear background are used to fit the IRIS spectra at \u201cO\u202f\u202fI\u201d window (Tian et al. 2016). Next, we can extract the hot line of Fe\u202fXXI 1354.08 \u00c5, as shown by the turquoise profile. The purple profile is the cool line of C\u202fI 1354.29 \u00c5. Two orange peaks represent the cool lines of O\u202fI 1354.60 \u00c5 and C\u202fI 1354.84 \u00c5 (Tian 2017), which are far away from the flaring line of Fe\u202fXXI 1354.08 \u00c5. Finally, the line properties of Fe\u202fXXI 1354.08 \u00c5 are extracted from the fitting results, that is, Doppler velocity, peak intensity, and line width (Li et al. 2016b; Tian et al. 2016; Tian & Chen 2018).","Citation Text":["Tian 2017"],"Functions Text":["We also manually perform the absolute wavelength calibration using a relatively strong neutral line, O\u202fI 1355.60 \u00c5 (see"],"Functions Label":["Uses"],"Citation Start End":[[821,830]],"Functions Start End":[[659,778]]}
{"Identifier":"2018ApJ...853...50F__Bernard_et_al._2015b_Instance_3","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"],"Functions Text":["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"],"Functions Label":["Uses"],"Citation Start End":[[1570,1590]],"Functions Start End":[[1372,1541]]}
{"Identifier":"2020AandA...641A.151S__Spinoglio_et_al._(2002)_Instance_1","Paragraph":"The galaxy NGC 7213 benefits from a detailed SED decomposition performed by G16, which allowed us to disentangle the relative contributions of AGN and SF activity to the global IR outcome of the source, providing a characterisation of the host galaxy in terms of stellar and dust content (M\u22c6 and Mdust, respectively), and ongoing SF (SFR). Here, we briefly introduce the photometric data collected from the archive and the SED decomposition procedure adopted by G16. The homogenised catalogue of total fluxes, from the UV to the FIR, is presented by G16 (see also their Table A.1: the flux densities are corrected for the aperture and magnitude zero point). In the case of NGC 7213, the photometric data included in the analysis are: the U, B, V, and R bands from de Vaucouleurs et al. (1991); the near-infrared measurements from the catalogue by Jarrett et al. (2000); the IRS spectrum re-binned by Gruppioni et al. (2016) and the photometry by Gallimore et al. (2010) and Moshir et al. (1990) in the MIR; and the FIR photometry by Spinoglio et al. (2002). The adopted SED fitting code was SED3FIT5 (Berta et al. 2013), which reproduces the stellar emission and the emission of the dust heated by the stars and the torus emission simultaneously. The code used the library by Bruzual & Charlot (2003) for the stellar contribution, that of da Cunha et al. (2008) for the IR dust-emission, and the library of smooth AGN tori by Fritz et al. (2006), updated by Feltre et al. (2012). In order to limit the degeneracy among the torus parameters, in G16 the AGN configurations of obscured sources were excluded for NGC 7213 (as supported by optical observations of the source and by the X-ray spectral properties presented in Sect. 3.1.2). The best fit model and the decomposition in the different components is presented in Fig. 5. The host-stellar contribution and the dusty SF dominate over the AGN in the optical bands and in the entire IR band, respectively. While this could appear to be in contrast with the type 1\/broad-line nature of the AGN, it is in agreement with the relatively weak nuclear activity observed in NGC 7213 (revealed also through the X-ray spectral analysis reported Sect. 3.1.2).","Citation Text":["Spinoglio et al. (2002)"],"Functions Text":["In the case of NGC 7213, the photometric data included in the analysis are:","and the FIR photometry by"],"Functions Label":["Uses","Uses"],"Citation Start End":[[1033,1056]],"Functions Start End":[[658,733],[1007,1032]]}
{"Identifier":"2020ApJ...897...84T__Zhang_et_al._2019_Instance_1","Paragraph":"A comparison between the constraints in Figure 2 and those found from the analysis of the reflection spectrum of the disk of other sources in previous studies is not straightforward because these measurements are quite sensitive to the specific source and the quality of the data, so it may be dangerous to generalize the results found from our analysis of 17 RXTE observations of LMC X-1. In general, a correlation between the estimates of a* and \u03b113 is common even when we analyze the reflection spectrum. However, such a degeneracy can be broken when the inner edge of the disk is very close to the compact object (Tripathi et al. 2018, 2019b; Zhang et al. 2019). For example, Tripathi et al. (2019a) analyzed simultaneous XMM-Newton and NuSTAR observations of MCG\u20136\u201330\u201315 obtaining \n\n\n\n\n\n and \n\n\n\n\n\n (90% confidence level for one relevant parameter). Even if there is a correlation between these measurements of a* and \u03b113, see Figure 6 in Tripathi et al. (2019a), we can get quite stringent constraints on both parameters. This is not the case with the analysis of the thermal component presented in this paper. When we assume the Kerr metric, observations 7, 12, and 14 of LMC X-1 give quite high spin values (see left column of Table 1), comparable to the XMM-Newton and NuSTAR observations of MCG\u20136\u201330\u201315. However, when we leave \u03b113 free it is not easy to constrain a* and \u03b113 at the same time. While the limited energy resolution of RXTE with respect to XMM-Newton may have some effect, the key-point is in the difference between the reflection spectrum and the thermal one. The former is characterized by many features, notably, but not only, the iron K\u03b1 complex around 6\u20137 keV. Such features help to break the parameter degeneracy, even if the reflection spectrum has several parameters to fit. The thermal spectrum, on the contrary, has quite a simple shape and there is an intrinsic degeneracy among the model parameters. This is true even when we assume the Kerr metric: it is possible to measure the black-hole spin only when we have independent estimates of the black-hole mass, distance, and inclination angle of the disk. If we want to use the continuum-fitting method to test the Kerr metric and we add a deformation parameter, the problem of degeneracy between a* and \u03b113 should not be a surprise.","Citation Text":["Zhang et al. 2019"],"Functions Text":["However, such a degeneracy can be broken when the inner edge of the disk is very close to the compact object"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[647,664]],"Functions Start End":[[508,616]]}
{"Identifier":"2016MNRAS.462.3441D__Namouni_1999_Instance_4","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)"],"Functions Text":["Asteroid 469219 librates around $\\omega _{\\rm r}=-90{^\\circ }$ for Venus, the Earth, and Jupiter. This behaviour corresponds to domain III in","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."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1491,1505]],"Functions Start End":[[1349,1490],[1507,1760]]}
{"Identifier":"2016MNRAS.458...84A__Larsson_et_al._2007_Instance_1","Paragraph":"While the properties of the SLSN hosts themselves are of interest, they are most diagnostic when compared to other classes of extragalactic transient whose progenitors are better understood. To this end, we employ a comparison sample of LGRB and CCSN3 host galaxies. In principle, CCSNe should trace all core-collapse events, although the mass function means they will be dominated by stars at the lower mass end (\u223c8 M\u2299 to \u223c25 M\u2299). There also remains a possibility that some very massive stars can undergo core collapse without yielding a LSNe (e.g. Smartt 2009; Ugliano et al. 2012; Kochanek 2014) such that CCSNe samples might only provide a census of lower mass core collapsing stars (e.g. 8  M*  20 M\u2299). Indeed, constraints from explosion parameters have shown the majority of CCSNe to be consistent with lower mass progenitors, as opposed to more massive Wolf\u2013Rayet stars (Cano 2013; Lyman et al. 2016) GRBs likely represent a population with rather larger initial masses (Larsson et al. 2007; Raskin et al. 2008). LGRBs are now known to be associated with the core collapse of massive stars, and broad line SN Ic are near ubiquitously associated with low-z events (where such signatures can be seen; Hjorth et al. 2012). When compared to the hosts of CCSNe they are generally smaller and of lower luminosity, consistent with an origin in galaxies of lower metallicity (Fruchter et al. 2006; Svensson et al. 2010). In relatively local examples, where spatially resolved gas phase metallicities can be obtained, these indeed appear to be lower for GRBs than for CCSNe, even in cases where the luminosity of the galaxy is relatively high (i.e. the GRB host galaxies lie off the mass\u2013metallicity relation; Modjaz et al. 2008; Graham & Fruchter 2013). Hence, comparing the hosts of SLSNe to these events allows us to test the large-scale environments of SLSNe against those of the bulk core-collapse population and a subset which appears to derive largely from massive stars at lower metallicity, although we note that agreement on this matter is not complete (e.g. Podsiadlowski, Joss & Hsu 1992; Eldridge, Izzard & Tout 2008; Smartt 2009; Drout et al. 2011). By exploiting both LGRB and CCSN host samples we can ascertain if there is a strong metallicity dependence in SLSN production, and if this is more or less extreme than that observed in GRB hosts.","Citation Text":["Larsson et al. 2007"],"Functions Text":["Indeed, constraints from explosion parameters have shown the majority of CCSNe to be consistent with lower mass progenitors, as opposed to more massive Wolf\u2013Rayet stars","GRBs likely represent a population with rather larger initial masses"],"Functions Label":["Background","Background"],"Citation Start End":[[978,997]],"Functions Start End":[[708,876],[908,976]]}
{"Identifier":"2015AandA...582A..41H__Davis_et_al._2014_Instance_1","Paragraph":"The snowlines of various volatiles (sublimation temperature Tsub \u2272 160 K) play a major role for planet formation. Beyond the snowline, the high abundances of solids allow for efficient sticking to form larger bodies, which is further enhanced by the presence of ices (e.g., Stevenson & Lunine 1988; Ros & Johansen 2013). Extensive studies have investigated the snowline in protoplanetary disks around pre-main-sequence stars similar to the nebula out of which the solar system is assumed to have formed (e.g., Lissauer 1987; Pollack et al. 1996). In such models, the water snowline is located at a radius of a few AU. It is thought that the early pre-solar nebula was hot (>1500 K), such that both volatiles and refractories (Tsub \u2273 1400 K) are in the gas phase out to larger distances (Cassen 2001; Scott 2007; Davis et al. 2014; Marboeuf et al. 2014). The evidence of such a hot solar nebula comes from the history of the refractories, but the volatile content of comets seems to indicate that a part of the disk remains cold (Bockel\u00e9e-Morvan et al. 2000; Mumma & Charnley 2011; Pontoppidan et al. 2014). The evolution of the snowline due to disk and star evolution and its accretion rate clearly affects the chemical composition in the region relevant to planet formation (e.g., Lodders 2004; Davis 2005; \u00d6berg et al. 2011b). The most relevant volatiles are the known major ice species: H2O, CO2, and CO. Observations (Meijerink et al. 2009; Zhang et al. 2013) and models (e.g., D\u2019Alessio et al. 1998; Dullemond et al. 2007) of protoplanetary disks around pre-main sequence T Tauri stars indicate that such disks are not warm enough to have gas-phase volatiles in the midplane beyond 30 AU, as claimed in some early solar nebula models, and, for the case of H2O, a snowline of only a few AU is typically found. Higher temperatures at large radii might be achieved, but only during the deeply embedded phase of star formation when the accretion rate is high. The question remains how hot an embedded accreting disk can be when the accretion rate is high (\u226510-6M\u2299 yr-1, see Dunham et al. 2014, for a recent review). ","Citation Text":["Davis et al. 2014"],"Functions Text":["It is thought that the early pre-solar nebula was hot (>1500 K), such that both volatiles and refractories (Tsub \u2273 1400 K) are in the gas phase out to larger distances"],"Functions Label":["Background"],"Citation Start End":[[812,829]],"Functions Start End":[[618,785]]}
{"Identifier":"2021MNRAS.506.2181L__Qi_et_al._2019_Instance_1","Paragraph":"There are still several sources of systematics we do not consider in this paper. For instance, whether the use of a different mass distribution models for these lenses could significantly affect the final result. Therefore, we performed a sensitivity analysis and repeated the above calculation using the extend power law (EPL) lens model, in which the luminosity density profile [v(r) \u223c r\u2212\u03b4] is different from the total mass (luminous plus dark matter) density profile \u03c1(r) \u223c r\u2212\u03b1. Such lens model has found widespread astrophysical applications in the literature (Cao et al. 2016a; Xia et al. 2017; Qi et al. 2019), considering the anisotropic distribution of stellar velocity dispersion \u03b2 (Koopmans 2005; Cao et al. 2017b; Chen et al. 2019). With the EPL lens parameters (\u03b1, \u03b2, \u03b4) modelled by Gaussian distributions \u03b1 = 2.00 \u00b1 0.08, \u03b4 = 2.40 \u00b1 0.11, and \u03b2 = 0.18 \u00b1 0.13 (Gerhard et al. 2001; Bolton, Rappaport & Burles 2006; Graur et al. 2014; Schwab, Bolton & Rappaport 2010), the scatter plot of the deviation Tzs in EPL model is shown in Fig. 4. Our results provide the deviation Tzs = 0.974 \u00b1 0.017 [corresponding to c(zs) = 2.922(\u00b1 0.051) \u00d7 105 km\u2009s\u22121], and the median value Med(Tzs) = 0.983 with the median absolute deviation MAD(Tzs) = 0.259 for the full lens sample. Therefore, our results show that the assumed lens model has a slight impact on the SOL constraint, which highlights the importance of auxiliary data (such as more high quality integral field unit) in improving constraints on the density profile of gravitational lenses. More detailed models of mass distribution, such as Navarro\u2013Frenk\u2013White density profile (suitable for dark matter distribution; Navarro, Frenk & White 1997), Sersic-like profile (suitable for stellar light distribution) (S\u00e9rsic 1968), pseudo-isothermal elliptical mass distribution, could also be considered in this context (Kassiola & Kovner 1993). However, strong lensing observables we used are determined by the total mass inside the Einstein radius. Hence, they are not so sensitive to the details of the very central distribution like cusps (besides the extremal cases). Moreover, the Einstein rings of the lenses we used corresponded to less than 10 kpc hence the NFW profile of the dark halo would not likely be manifested.","Citation Text":["Qi et al. 2019"],"Functions Text":["Such lens model has found widespread astrophysical applications in the literature"],"Functions Label":["Background"],"Citation Start End":[[600,614]],"Functions Start End":[[482,563]]}
{"Identifier":"2019AandA...629L...4A__Arzoumanian_et_al._2011_Instance_1","Paragraph":"Herschel imaging observations have shown that filamentary structures are truly ubiquitous in the cold interstellar medium (ISM) of the Milky Way (Molinari et al. 2010), dominate the mass budget of Galactic molecular clouds at high densities (\u2273104 cm\u22123) (Schisano et al. 2014; K\u00f6nyves et al. 2015), and feature a high degree of universality in their properties. In particular, detailed analysis of the radial column density profiles indicates that, at least in the nearby clouds of the Gould Belt, molecular filaments are characterized by a narrow distribution of crest-averaged inner widths with a typical full width at half maximum (FWHM) value Wfil\u200b\u200b\u2004\u223c\u2004\u200b\u200b0.1 pc and a dispersion of less than a factor of \u223c2 (Arzoumanian et al. 2011, 2019; Koch & Rosolowsky 2015). Another major result from Herschel (e.g., Andr\u00e9 et al. 2010; K\u00f6nyves et al. 2015; Marsh et al. 2016) is that the vast majority (> 75%) of prestellar cores are found in dense \u201ctranscritical\u201d or \u201csupercritical\u201d filaments for which the mass per unit length, Mline, is close to or exceeds the critical line mass of nearly isothermal, long cylinders (e.g., Inutsuka & Miyama 1997), \n\n\n\n\nM\n\nline,crit\n\n\u200b\n=\n\u2009\u200b\n2\n\nc\ns\n2\n\n\/G~16\nM\n\u2299\n\n\n\n$ {M_{{\\rm{line}},{\\rm{crit}}}} = \\,2{\\mkern 1mu} c_{\\rm{s}}^2\/G{\\mkern 1mu} \\sim {\\mkern 1mu} 16{\\mkern 1mu} {M_ \\odot } $\n\n\n pc\u22121, where cs\u200b\u2004\u223c\u2004\u200b0.2 km s\u22121 is the isothermal sound speed for molecular gas at T\u200b\u2004\u223c\u2004\u200b10 K. Moreover, most prestellar cores lie very close to the crests (i.e., within the inner 0.1 pc portion) of their parent filaments (e.g., K\u00f6nyves et al. 2019; Ladjelate et al. 2019). These findings support a filamentary paradigm in which low-mass star formation occurs in two main steps (Andr\u00e9 et al. 2014; Inutsuka et al. 2015): (1) multiple large-scale compressions of cold interstellar material in supersonic magneto-hydrodynamic (MHD) flows generate a cobweb of \u223c0.1 pc-wide filaments within sheet-like or shell-like molecular gas layers in the ISM and (2) the densest molecular filaments fragment into prestellar cores (and then protostars) by gravitational instability near or above the critical line mass, Mline,\u2006crit, corresponding to \n\n\n\n\n\u03a3\n\ngas\n\ncrit\n\n\u223c\n\nM\n\nline\n,\ncrit\n\n\n\/\n\nW\nfil\n\n\u223c\n160\n\n\nM\n\u2299\n\n\n\n$ \\Sigma_{\\mathrm{gas}}^{\\mathrm{crit}} \\sim M_{\\mathrm{line, crit}}\/W_{\\mathrm{fil}} \\sim 160\\,M_\\odot $\n\n\n pc\u22122 in gas surface density (AV\u2004\u223c\u20047.5) or nH2\u2004\u223c\u20042\u2005\u00d7\u2005104 cm\u22123 in volume density. This paradigm differs from the classical gravo-turbulent picture (Mac Low & Klessen 2004) in that it relies on the anisotropic formation of dense structures (such as shells, filaments, cores) in the cold ISM and the unique properties of filamentary geometry (see Larson 2005).","Citation Text":["Arzoumanian et al. 2011"],"Functions Text":["In particular, detailed analysis of the radial column density profiles indicates that, at least in the nearby clouds of the Gould Belt, molecular filaments are characterized by a narrow distribution of crest-averaged inner widths with a typical full width at half maximum (FWHM) value Wfil\u200b\u200b\u2004\u223c\u2004\u200b\u200b0.1 pc and a dispersion of less than a factor of \u223c2"],"Functions Label":["Background"],"Citation Start End":[[710,733]],"Functions Start End":[[361,708]]}
{"Identifier":"2016MNRAS.461..248S__Munari_et_al._2013_Instance_4","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"],"Functions Text":["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","\u03c38 = 0.8) may also play a role."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[2611,2629]],"Functions Start End":[[2458,2609],[2631,2662]]}
{"Identifier":"2021AandA...655A.109B__Foreman-Mackey_et_al._2013_Instance_1","Paragraph":"When the size of the logarithmic redshift bin is small enough, we can use fluxes in place of luminosities, performing a test on the (non-)evolution with redshift that is completely independent from any assumption on cosmology. Risaliti & Lusso (2019) analysed in detail the choice of the bin size and verified that, as long as \u0394log(z)\u22640.1, the slope in the relation does not depend on it. Thanks to the statistics available, we chose \u0394log(z)=0.06 and we limited our analysis to the redshift bins with more than five objects. We performed the same analysis for bins of size \u0394log(z)=0.05 and \u0394log(z)=0.07, finding no significant difference (see Fig. 9 for a comparison of the results for different sizes of the redshift bins). For our selection in \u0393 and the choice of the threshold for the Eddington bias, the division yielded 17 redshift bins. To perform the fitting to the data, we adopted the Python package emcee (Foreman-Mackey et al. 2013), a pure-Python implementation of Goodman & Weare\u2019s Affine Invariant Markov chain Monte Carlo (MCMC) Ensemble sampler. To check that the results were independent from the employed method, we also performed the analysis using the Linmix package (Kelly 2007), an algorithm that makes use of a Bayesian approach to linear regression and takes into account the errors in both the x and y variable. We performed the fit a first time, then applied a 3\u03c3 clipping to the data, repeating this sequence for a total of three times. This yielded no significant difference with respect to the analysis without \u03c3 clipping. The results are shown in Figs. 7 and 9, and summarised in Table 3. In Fig. 7, red points indicate when the observations are characterised by an SN  5 in the soft band. Most of them are (X-ray) fainter objects at intermediate redshifts. This confirms that data points that passed our selection criteria, even if with a low SN, follow the relation and are representative of the population of blue quasars.","Citation Text":["Foreman-Mackey et al. 2013"],"Functions Text":["To perform the fitting to the data, we adopted the Python package emcee",", a pure-Python implementation of Goodman & Weare\u2019s Affine Invariant Markov chain Monte Carlo (MCMC) Ensemble sampler."],"Functions Label":["Uses","Uses"],"Citation Start End":[[916,942]],"Functions Start End":[[843,914],[943,1061]]}
{"Identifier":"2019ApJ...887...75T__Stritzinger_et_al._2012_Instance_1","Paragraph":"Figure 6 compares the IR SED of SN 2014C with the SEDs of all other interacting SNe for which observations beyond 5 \u03bcm are available in the literature. Apart from SN 2014C, only four other strongly interacting SNe IIn have been observed in the mid-IR. Three of these were observed at epochs comparable to the epoch at which SN 2014C was observed in the mid-IR. SN 2005ip was observed at 936 days post-explosion during the cryogenic phase of Spitzer. It was the only interacting SN with an InfraRed Spectrograph (IRS; Houck et al. 2004) mid-IR spectrum from 5 to 12 \u03bcm and IRAC photometry at 5.8 and 8 \u03bcm (Fox et al. 2010). SN 2006jd was observed at 1638 days, an epoch very similar to that of SN 2014C, with Spitzer\/IRAC and WISE (Stritzinger et al. 2012). SN 2010jl was observed at 1279 days with Spitzer\/IRAC and SOFIA\/FORCAST (Herter et al. 2018) at 11.1 \u03bcm, resulting in a deep upper limit after 6400 s of total integration time (Williams & Fox 2015). In all three cases, the photometry and spectra from 1 to 10 \u03bcm are well fitted by purely carbonaceous dust models, which we overplot in Figure 6 using dust parameters from the literature. Lastly, SN 1995N was observed at more than 10 yr post-explosion by Spitzer\/IRAC and WISE (Van Dyk 2013). Its SED shape differs markedly from those of the other three SNe observed at earlier epochs. The carbonaceous dust model that Van Dyk (2013) fitted to the data is shown in Figure 6. However, we note that the shallow slope from 3 to 10 \u03bcm cannot be fitted with a single-temperature dust model, regardless of composition, and would require a range of dust temperatures. Given the late epoch of the observation, one might also consider a nonthermal origin for the IR emission from SN 1995N, as the SED can also be described with a broken power law (F\u03bd \u221d \u03bd\u22123; also overplotted) with a knee at around 12 \u03bcm. In addition to these H-rich interacting SNe, the H-poor interacting SN 2006jc (Ibn) was also observed beyond 5 \u03bcm with AKARI  (Sakon et al. 2009). Its SED was also best fitted with a two-temperature amorphous carbon dust model. This comparison highlights SN 2014C\u2019s unique SED shape among other interacting SNe for which data are available beyond 5 \u03bcm at comparable epochs, showing for the first time an evidence for a silicate dust feature in the IR SED of an interacting SN. It also accentuates the need for observations of interacting SNe at late times, out to decades post-explosion, in the near- to mid-IR, which will be enabled by the upcoming James Webb Space Telescope (JWST).","Citation Text":["Stritzinger et al. 2012"],"Functions Text":["SN 2006jd was observed at 1638 days, an epoch very similar to that of SN 2014C, with Spitzer\/IRAC and WISE","In all three cases, the photometry and spectra from 1 to 10 \u03bcm are well fitted by purely carbonaceous dust models, which we overplot in Figure 6 using dust parameters from the literature."],"Functions Label":["Uses","Similarities"],"Citation Start End":[[731,754]],"Functions Start End":[[623,729],[956,1143]]}
{"Identifier":"2019ApJ...887..137S__Vekstein_2017_Instance_2","Paragraph":"As mentioned above, the magnetic reconnection is introduced as breaking and reconfiguration of the oppositely directed magnetic field lines in highly conducting plasma. The magnetic field lines collapse near the X-point and form an extended magnetic singularities known as a current sheet. There are two mechanism of the current-sheet formation. The first kind of current-sheet formation is associated with the MHD instabilities (e.g., resistive tearing mode and ideal kink mode) known as spontaneous magnetic reconnection (e.g., White 1984; Baty 2000; Vekstein 2017). The second kind of current sheet can be formed in the MHD stable configuration, where some external perturbations trigger the forced magnetic reconnection (Hahm & Kulsrud 1985). The forced magnetic reconnection may be activated by nonlinear MHD waves, which may be caused by explosive solar activities (e.g., Sakai et al. 1984; Dewar et al. 2013; Beidler et al. 2017). The forced magnetic reconnection may be developed due to boundary perturbations, which induce a surface current in such a way that it opposes the progress of the reconnection (Ishizawa & Tokuda 2000, 2001; Fitzpatrick 2003). The multimode simulation approach has been adopted to investigate the thinning of the current sheet induced by forced magnetic reconnection (Birn et al. 2005). The motion of the photospheric footpoints of the coronal magnetic field may also trigger the forced magnetic reconnection, which may be caused by the explosive solar coronal events (e.g., Vekstein & Jain 1998; Jain et al. 2005; Vekstein 2017). Although there is a remarkable development in the theory of the forced magnetic reconnection, Jess et al. (2010) have suggested that there is no observational evidence of explosive flare or coronal activities triggered by forced magnetic reconnection. They have observed a microflare activity driven by forced magnetic reconnection. The lower solar atmosphere (photosphere & chromosphere) is dominated by cool, partially ionized and collision dominated plasma. Most of the energy releases during the forced magnetic reconnection may be consumed by such plasma systems (e.g., Litvinenko 1999; Chen et al. 2001; Chen & Ding 2006; Litvinenko et al. 2007).","Citation Text":["Vekstein 2017"],"Functions Text":["The motion of the photospheric footpoints of the coronal magnetic field may also trigger the forced magnetic reconnection, which may be caused by the explosive solar coronal events (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[1551,1564]],"Functions Start End":[[1323,1510]]}
{"Identifier":"2018ApJ...866...48U__Korngut_et_al._2011_Instance_1","Paragraph":"RX J1347.5\u20131145 is one of the most luminous X-ray galaxy clusters and is located at a redshift of z = 0.451. It was thought to be a relaxed cluster when it was discovered in the ROSAT all sky survey (Schindler et al. 1997). Komatsu et al. (1999) made the first measurements of the Sunyaev\u2013Zel\u2019dovich effect (SZE: Sunyaev & Zeldovich 1972) toward this cluster with the James Clerk Maxwell Telescope at 350 GHz as well as with the 45 m Nobeyama Radio Telescope at 21 and 43 GHz. A higher angular resolution observation of the SZE was performed by Komatsu et al. (2001) using the Nobeyama Bolometer Array and they found a prominent substructure which has no counterpart in the soft X-ray image from ROSAT. The presence of the substructure has been confirmed by Chandra and XMM-Newton (e.g., Allen et al. 2002; Gitti & Schindler 2004) as well as by more recent SZE measurements (Mason et al. 2010; Korngut et al. 2011; Plagge et al. 2013; Adam et al. 2014; Kitayama et al. 2016). Allen et al. (2002) measured the mean temperature of the ICM to be over 10 keV, which is relatively high compared to other typical clusters. Kitayama et al. (2004) and Ota et al. (2008) found a very hot (>20 keV) component of the ICM in this cluster. In addition, the radial profile and spatial distribution of the ICM temperature indicate that the temperature drops to \u223c6 keV toward the cluster center so that the cool core is formed (e.g., Allen et al. 2002; Ota et al. 2008; Kreisch et al. 2016). A disturbed morphology is further supported by radio synchrotron observations (e.g., Ferrari et al. 2011) and gravitational lensing maps (e.g., K\u00f6hlinger & Schmidt 2014). The total mass of RX J1347.5\u20131145 within r200 is estimated to be \u223c1.5 \u00d7 1015 h\u22121 \n\n\n\n\n\n  using weak-lensing analysis, where r200, the radius within which the mean mass density is 200 times the critical density of the universe, is 1.85 h\u22121 Mpc (Lu et al. 2010) for this galaxy cluster.18\n\n18\nThey adopted the Hubble constant of 70 km s\u22121 Mpc\u22121.\n\n","Citation Text":["Korngut et al. 2011"],"Functions Text":["The presence of the substructure has been confirmed by","as well as by more recent SZE measurements"],"Functions Label":["Background","Background"],"Citation Start End":[[894,913]],"Functions Start End":[[703,757],[831,873]]}
{"Identifier":"2021MNRAS.505.4289P__Vila-Costas_&_Edmunds_1993_Instance_1","Paragraph":"For decades, the metallicity and the above chemical abundance ratios have been estimated in samples of star-forming galaxies (SFGs), by using their emission line ratios to calculate the abundances of the different elements that originate them. Many techniques were developed for such studies, such as the use of the Te-method, photoionization models, or optical calibrations based on strong emission lines (see the review by Maiolino & Mannucci 2019 for a more detailed discussion). By using these techniques, some studies show a correlation between both 12 + log\u2009(O\/H) and log\u2009(N\/O) (Vila-Costas & Edmunds 1993; Masegosa, Moles & Campos-Aguilar 1994; Andrews & Martini 2013). On the other hand, it is also important how the different properties of the host galaxies correlate with their chemical abundances. For instance, preliminary studies of the Local Group of galaxies (McClure & van den Bergh 1968; Lequeux et al. 1979) revealed the existence of a relation between the luminosity of an SFG and its metallicity, the so-called luminosity\u2013metallicity relation (the most luminous galaxies have the higher chemical abundances), which has been later studied in further detail (Garnett & Shields 1987; Skillman, Kennicutt & Hodge 1989; Brodie & Huchra 1991; Vilchez 1995; Mateo 1998). Later on it was probed that the luminosity\u2013metallicity relation was in fact the result of a more fundamental mass\u2013metallicity relation (Garnett 2002; P\u00e9rez-Gonz\u00e1lez et al. 2003; Pilyugin, V\u00edlchez & Contini 2004). This relation has been probed in large sample of SFG both at low-redshift (Contini et al. 2002; Melbourne & Salzer 2002; Lamareille et al. 2004; Tremonti et al. 2004) and at high-redshift (Erb et al. 2006a; P\u00e9rez-Montero et al. 2009, 2013; Izotov et al. 2015; Gao et al. 2018; Torrey et al. 2019). In addition, it has been reported that the star-formation rate is also related to gas-phase Z in SFG (Lara-L\u00f3pez et al. 2010; Mannucci et al. 2010; Yates, Kauffmann & Guo 2012). Morphological type may also be related with the chemical abundances, as early-type galaxies, despite some of them can present episodes of on-going star-formation (Zhu, Blanton & Moustakas 2010), tend to have on average older stellar populations than late-type galaxies (Kennicutt 1998; Trager et al. 2000; Gebhardt et al. 2003; Schiavon 2007), being thus more chemically evolved.","Citation Text":["Vila-Costas & Edmunds 1993"],"Functions Text":["By using these techniques, some studies show a correlation between both 12 + log\u2009(O\/H) and log\u2009(N\/O)"],"Functions Label":["Background"],"Citation Start End":[[585,611]],"Functions Start End":[[483,583]]}
{"Identifier":"2018MNRAS.478.4336M__Porter_&_Raychaudhury_2007_Instance_1","Paragraph":"Recently filaments, in particular the warm hot intergalactic medium (WHIM) in filaments, have been the pivot of several studies. Using emission in the soft X-ray bands, such studies estimate the WHIM temperature in filaments to be \u223c3\u20138 keV (Eckert et al. 2015; Akamatsu et al. 2017; Parekh et al. 2017; Tanimura et al. 2017). Furthermore, the cosmic microwave background map from the Planck together with the Canada\u2013France\u2013Hawaii Telescope Lensing Survey, as well as the Two-Micron All-Sky Redshift Survey of galaxies suggest that at least half of the missing baryons in the Universe may reside as WHIM in large-scale filaments tracing the dark matter distribution (Van Waerbeke, Hinshaw & Murray 2014; G\u00e9nova-Santos et al. 2015). Therefore, undeniably it is crucial to characterize the large-scale structure (LSS) of the Universe and comprehend the impact of the cosmic-web on properties of galaxies. The Coma supercluster is one of the richest large-scale structures (Chincarini & Rood 1976) in the nearby Universe comprising two clusters of galaxies, connected by a web of large-scale filaments around $30\\,h_{70}^{-1}$ Mpc long (e.g. Fontanelli 1984). The two clusters, Coma (Abell 1656) and Abell 1367, along with the filaments of galaxies dispersed with several small galaxy groups span \u223c500 deg2 on the sky (Mahajan et al. 2010). The large-scale filaments in the Coma supercluster have not just been observed by means of the galaxy distribution in the optical wavebands (Gregory & Thompson 1978; Mahajan et al. 2010), but also diffuse emission in the radio continuum (Kim et al. 1989). Studies of clusters and groups at $z$ \u223c 0 have evidently shown that outskirts of groups and clusters (Zabludoff & Mulchaey 1998; Wang, Owen & Ledlow 2004; Rines et al. 2005; Cortese et al. 2007; Tran et al. 2009; Gavazzi et al. 2010; Smith et al. 2010; Sun et al. 2010; Coppin et al. 2011; Mahajan, Raychaudhury & Pimbblet 2012; Verdugo et al. 2012; Mahajan 2013) and filaments of galaxies (Porter & Raychaudhury 2007; Bou\u00e9 et al. 2008; Fadda et al. 2008; Porter et al. 2008; Edwards et al. 2010a; Biviano et al. 2011) are favourable sites for galaxy transformations. Based on a study using optical data from the Sloan Digital Sky Survey (SDSS) data release (DR) 7, Mahajan et al. (2010) found that the star formation (SF)\u2013density relation in the Coma supercluster for the giant galaxies is much weaker than their dwarf counterparts. However, the fraction of star-forming galaxies for both declines to \u223c0 at the centre of the clusters (also see Mahajan, Haines & Raychaudhury 2011). Cybulski et al. (2014) further study the star formation in the Coma supercluster by combining a complementary optical data set from SDSS DR 9, with IR data from the Wide-Field Infrared Survey Explorer (Wright et al. 2010) and UV data from the Galaxy Evolution Explorer (GALEX; Martin et al. 2005). Cybulski et al. (2014) corroborated the results of Mahajan et al. (2010, 2011) by probing both obscured and unobscured star formation down to \u223c0.02 M\u2299 yr\u22121, in order to quantify the effect of different types of large-scale environments: groups, clusters, filaments, and voids, on quenching SF in galaxies. In the absence of dust in star-forming galaxies, the UV emission is a good tracer of massive (>10 M\u2299) star formation. On the other hand, optical emission lines such as H \u03b1 probe instantaneous star formation over a time-scale of \u227220 Myr (Kennicutt 1998). Assuming that the UV luminosity is not overwhelmed by contribution from the old stellar populations due to the UV upturn such as in massive early-type galaxies (O\u2019Connell 1999), the UV luminosity measures star formation over a time-scale of \u223c100 Myr (Kennicutt 1998). Hence, the star formation rate (SFR) estimated from optical emission lines delineates the continuous SF in a galaxy, while the SFR determined from the UV is representative of its recent SF activity. But even though GALEXand its predecessor UV imagers have been used to investigate individual galaxies within clusters and groups (e.g. Hicks & Mushotzky 2005), or galaxy populations therein (e.g. Donas et al. 1990; Donas, Milliard & Laget 1995; Cornett et al. 1998; Boselli et al. 2005b), limited work has been done to analyse the UV properties of galaxies in the large-scale cosmic-web. Since the Coma supercluster is one of the most well studied regions in the nearby Universe, many other authors (e.g. Bernstein et al. 1995; Mobasher et al. 2003; Hammer et al. 2012; Smith et al. 2010; Smith, Lucey & Carter 2012b) have made use of optical and UV data to study the Coma and Abell 1367 clusters and their surroundings. With the advent of large redshift surveys, several studies (Gavazzi et al. 2010; Mahajan et al. 2010, 2011; Gavazzi et al. 2013; Cybulski et al. 2014 ) have also used multiwavelength data at optical, UV, and 21 cm continuum to study the properties of galaxies in the entire supercluster region. In this paper, we make use of similar data sets: UV data derived from GALEX and optical spectroscopic and photometric data from the SDSS for the entire Coma supercluster to further explore the impact of environment on the properties of galaxies. Conventionally, the \u2018environment\u2019 of galaxies is quantified as the projected density of galaxies in a fixed 2D or 3D region of the sky (e.g. Dressler 1980). Muldrew et al. (2012) combined 20 published methods of defining environment into two methods: (i) method that uses nearest-neighbours to probe the underlying galaxy density and (ii) fixed aperture methods. Muldrew et al. (2012) found that while the former are better suited for quantifying internal density within massive haloes, the latter fixed-aperture methods are better for probing the large-scale environment. Therefore, in order to characterize the large-scale cosmic-web, a combination of these methodologies is required to quantify the environment on different scales. In this work, we implement their result by making use of two different algorithms to define the large-scale filaments, and high density nodes of the cosmic-web characterized as clusters and groups. We also present a catalogue of all spectroscopically confirmed galaxy members of the Coma supercluster detected in the UV. This paper is organized as follows: in the next section we describe our data sets, followed by definition of environment in Section 3. In Section 4, we analyse the broad-band colours of galaxies as a function of environment, while in Section 5 we study the impact of the large-scale filaments on the properties of galaxies. Finally, we discuss our results in the context of the existing literature in Section 6 and summarize our results in Section 7. Throughout this paper, we use concordance \u039b cold dark matter cosmological model with H0 = 70 \u2009km\u2009s\u22121 Mpc\u22121, $\\Omega _\\Lambda =0.7$, and \u03a9m = 0.3 to calculate distances and magnitudes. We note that at the redshift of the Coma cluster ($z$ = 0.023) our results are independent of the cosmological model used.","Citation Text":["Porter & Raychaudhury 2007"],"Functions Text":["Studies of clusters and groups at $z$ \u223c 0 have evidently shown that outskirts of groups and clusters","and filaments of galaxies","are favourable sites for galaxy transformations."],"Functions Label":["Motivation","Motivation","Motivation"],"Citation Start End":[[1984,2010]],"Functions Start End":[[1593,1693],[1957,1982],[2112,2160]]}
{"Identifier":"2020AandA...643A..58B__Trujillo_et_al._2001_Instance_1","Paragraph":"Adaptive optics (AO, Roddier 1999) is a game changer in the quest for high-angular resolution, especially for ground-based astronomical observations that face the presence of wavefront aberrations introduced by the atmosphere (Roddier 1981). Thanks to AO, the point spread function (PSF) delivered by an optical instrument is much narrower, by a factor up to 50 on the full width at half maximum (FWHM), than the seeing-limited scenario (Roddier 1981). Still, some correction residuals persist and render the AO PSF shape complex to model. Consequently, standard parametric models that reliably reproduce the seeing-limited PSFs, such as a Moffat function (Trujillo et al. 2001; Moffat 1969), become inefficient at describing the AO PSF. Moreover, contrary to seeing-limited observations, AO-corrected images suffer from the anisoplanatism effect (Fried 1982) that strengthens the spatial variations of the PSF on top of instrument defects. Determining the AO PSF is necessary for two major reasons. Firstly, understanding and accurately modeling the PSF morphology is key to diagnosing AO performance. From the PSF, we can identify the major contributors to the AO residual error (Beltramo-Martin et al. 2019; Ferreira et al. 2018; Martin et al. 2017). Secondly, the image delivered by an optical instrument depends on both the science object we want to characterize and the PSF. In order to estimate the interesting astrophysical quantities, one can either use a deconvolution technique (F\u00e9tick et al. 2020, 2019a; Benfenati et al. 2016; Flicker & Rigaut 2005; Mugnier et al. 2004; Fusco et al. 2003; Drummond 1998) or include the PSF as part of a model, as is performed in PSF-fitting astrometry\/photometry retrieval techniques (Beltramo-Martin et al. 2019; Witzel et al. 2016; Schreiber et al. 2012; Diolaiti et al. 2000; Bertin & Arnouts 1996; Stetson 1987) and galaxy kinematics estimation (Puech et al. 2018; Bouch\u00e9 et al. 2015; Epinat et al. 2010) for instance.","Citation Text":["Trujillo et al. 2001"],"Functions Text":["Still, some correction residuals persist and render the AO PSF shape complex to model. Consequently, standard parametric models that reliably reproduce the seeing-limited PSFs, such as a Moffat function","become inefficient at describing the AO PSF."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[657,677]],"Functions Start End":[[453,655],[693,737]]}
{"Identifier":"2020AandA...642A...2R__Hyder_&_Lites_1970_Instance_1","Paragraph":"Investigation of solar wind outflow velocities is generally performed by the analysis of UV spectrometric observations. The UVCS\/SOHO instrument has provided H\u202fI Ly\u03b1 spectral line data over a longer time than a whole solar activity cycle (1996\u20132012), giving the possibility of studying coronal dynamics in different activity phases. The analysis of UVCS daily Ly\u03b1 synoptic data, in combination with electron densities derived from the white-light RS observations, has allowed to derive H\u202fI outflow speed maps (see e.g. Dolei et al. 2018, 2019). One of the methods based on the synergy between UV and WL observations is the Doppler dimming technique (Hyder & Lites 1970; Noci et al. 1987; Withbroe et al. 1982). It exploits the progressive UV intensity reduction of the coronal resonantly scattered component of the coronal H\u202fI Ly\u03b1 line emission with increasing outflow velocities. The line emission depends on the physical quantities involved in the Ly\u03b1 resonant scattering process, such as, for instance, the coronal electron density and temperature, and the chromospheric Ly\u03b1 radiation that excites the coronal H\u202fI atoms. The Ly\u03b1 intensity is also sensitive to the speed of the outflowing plasma from about 50\u2013500 km s\u22121, that are the typical velocities for neutral hydrogen atoms in the inner corona. Following the approach of Withbroe et al. (1982) and Noci et al. (1987), the intensity of the resonantly scattered Ly\u03b1 line can be also numerically computed by iteratively tuning the plasma speed value in order to reproduce the observed UV line intensity. The best match between computed and observed Ly\u03b1 intensity provides an estimate of the solar wind H\u202fI outflow velocity. Figure 9 shows an example of 2D outflow velocity map, in the range of heliocentric distances between 1.5 and 4.0 R\u2299, obtained via Doppler dimming technique (Dolei et al. 2018). The speed values radially increase with altitude up to about 150\u2013200 km s\u22121 in the equatorial regions and 400 km s\u22121 in the polar regions. These values are in agreement with the expected latitudinal distribution of slow and fast solar wind components, corresponding to equatorial regions, and mid-latitude and polar regions, respectively. The methodology put in place throughout this project will be applied later to the data acquired by the Solar Orbiter\u2019s Metis instrument, giving an unprecedented daily picture of the coronal dynamics, see Antonucci et al. (2020) for a detailed description of diagnostic techniques for Metis data.","Citation Text":["Hyder & Lites 1970"],"Functions Text":["One of the methods based on the synergy between UV and WL observations is the Doppler dimming technique"],"Functions Label":["Uses"],"Citation Start End":[[650,668]],"Functions Start End":[[545,648]]}
{"Identifier":"2021AandA...647A..67B__Cao_2011_Instance_1","Paragraph":"Keeping in mind the caveats discussed in Sect. 4.2, a relation between the properties of the jet collimation region and the properties of the accretion disk is suggested by Figs. 5 and 6. According to theoretical models and simulations, both thin disks (e.g., Blandford & Payne 1982; Fendt 2006; Liska et al. 2019) and geometrically-thick hot disks (e.g., Blandford & Begelman 1999; McKinney 2006; Begelman 2012; Mo\u015bcibrodzka & Falcke 2013; Mo\u015bcibrodzka et al. 2016) can launch collimated outflows. Due to the higher mass loading and lower speed, the disk-driven jet is expected to dominate the emission in radio galaxies with respect to the de-boosted black hole-launched jet. This is confirmed in observations by the direct imaging of limb-brightened jet structures (e.g., Boccardi et al. 2016a; Mertens et al. 2016; Giovannini et al. 2018) as well as by kinematic studies of radio galaxies, which generally show much lower intrinsic speeds than measured in blazars (Lister et al. 2019). As these properties are observed in high-luminosity and low-luminosity radio galaxies alike, a jet sheath must be produced from disks spanning different accretion regimes. Our results, however, indicate that the disk-driven jet in LERG originates at small disk radii (few RS, as measured in M 87), and indeed the expansion profiles of most of the LERG are well aligned with those of BL Lacs, which are expected to be dominated by the black hole-launched spine (see e.g., Ghisellini et al. 2014). This result is in broad agreement with models of jet launching from ADAFs (e.g., Cao 2011; Yuan & Narayan 2014, and references therein), which predict the formation of a thin and mildly-relativistic outer layer. ADAF models also predict the launch of a nonrelativistic disk-wind component carrying the bulk of the disk mass outflow and spanning a large solid angle. There is no evidence for such a component based on the analyzed VLBI images, at least in the considered frequency regime. The jet profiles in HEG, on the other hand, are all shifted upward, and a back-extrapolation down to the jet base suggests that the jet sheath is launched at larger disk radii. Taking as a reference Cygnus A, which shows the thinnest jet among HEG and for which an initial jet width of \u223c200 RS was measured based on GMVA observations (Boccardi et al. 2016b), the present data suggest that thin disks could launch collimated winds with an initial outer radius \u2273100 RS. This possible difference in the outer radius of the jet sheath is accompanied by a different extent of the collimation region in HEG and LEG (Fig. 6). Modeling of jet collimation by disk winds, presented by Globus & Levinson (2016), revealed a direct link between the wind outer radius and the collimation radius: for a given wind power, larger wind radii correspond to more extended collimation zones. A sufficiently high ratio (> 0.1) of wind power to jet power is required for this process to be efficient. When this condition is verified in reality, is a matter of debate. In recent simulations presented by Hervet et al. (2017), the diverse kinematic behavior of VLBI knots in blazars of different powers could be well explained by varying this ratio. Except for the least powerful class among BL Lacs (that of the High-frequency peaked BL Lacs, HBLs), whose properties could be reproduced by assuming an absent or very weak wind, ratios larger than 0.3 were suggested for blazars. A question remains concerning the portion of these winds which is actually detected in VLBI observations. When attempting to model the M 87 jet collimation profile, Globus & Levinson (2016) have suggested that the radio emission is produced in the shocked interface between the relativistic jet and the outer wind, which is undetected. Observational constraints on extended disk winds may be provided through other methods. For instance, we note that for all the HERG in our sample (except PKS 1514+00) the detection of ultra-fast outflows was reported based on X-ray observations (Tombesi et al. 2010, 2014; Reynolds et al. 2015). These outflows, whose launching mechanism is unclear, are suggested to be characterized by mildly relativistic speeds, to originate at disk radii of 102\u2005\u2212\u2005104RS (in agreement with our findings), and to carry a significant fraction of the jet kinetic power. Thus collimation via the action of disk winds, where by disk winds we mean a mildly relativistic jet sheath plus possible broader outflows, appears to be a viable mechanism, especially for high-luminosity sources.","Citation Text":["Cao 2011"],"Functions Text":["This result is in broad agreement with models of jet launching from ADAFs (e.g.,","which predict the formation of a thin and mildly-relativistic outer layer."],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[1567,1575]],"Functions Start End":[[1486,1566],[1623,1697]]}
{"Identifier":"2022AandA...659A..85F__Jeffries_et_al._2014_Instance_2","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"],"Functions Text":["Both clusters host two kinematically distinct populations"],"Functions Label":["Background"],"Citation Start End":[[1085,1105]],"Functions Start End":[[1026,1083]]}
{"Identifier":"2018ApJ...854...26L___2015a_Instance_3","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"],"Functions Text":["More details can be found in our previous papers"],"Functions Label":["Background"],"Citation Start End":[[1823,1838]],"Functions Start End":[[1773,1821]]}
{"Identifier":"2020AandA...642A..24W__Willingale_&_M\u00e9sz\u00e1ros_2017_Instance_1","Paragraph":"Selected as part of ESA\u2019s Cosmic Vision programme, Athena will study the evolution of large scale structure through the detection of WHIM filaments to trace the missing baryons in the local universe. Athena aims to measure the local cosmological baryon density in the WHIM to better than 10% and to constrain structure formation models in the low-density regime by measuring the redshift distribution and the physical parameters of 200 filaments against bright background sources. To achieve this, Athena will detect 200 filaments in the WHIM through absorption, 100 towards active galactic nuclei (AGN), and 100 towards bright gamma-ray burst (GRB) afterglows, up to redshifts of z = 1 (Kaastra et al. 2013). The WHIM can be detected by observing the absorption of highly ionised elements such as C, N, O, Ne, and possibly Fe (e.g. O\u202fVII, O\u202fVIII, Ne\u202fIX, Fe\u202fXVII). The strongest lines expected correspond to the H-like and He-like oxygen ions of the O\u202fVII 1s\u22122p X-ray resonance line (574 eV) and the unresolved O\u202fVIII 1s\u22122p X-ray doublet (653.5 eV, 653.7 eV). A bright X-ray background source such as an AGN (e.g. Padovani et al. 2017) or a GRB (e.g. Willingale & M\u00e9sz\u00e1ros 2017) can be used to detect the WHIM by producing absorption features in its energy spectra (Perna & Loeb 1998; Hellsten et al. 1998; Fiore et al. 2000). These sources provide a high probability of detection for the WHIM because they are sufficiently bright and distant to obtain a large statistical sample of lines. Bright AGN are common but are typically nearby, having an average redshift of z \u2248 0.8 for flat-spectrum radio quasars (FRSQ) and z \u2248 0.3 for BL-Lacs (Ackermann et al. 2011) and so, they probe relatively short lines of sight. Gamma-ray bursts occur at an average redshift of z \u2248 2 (Evans et al. 2009) and have been detected out to redshifts of 9.4 (Cucchiara et al. 2011). This allows for the probing of long lines of sight and can potentially provide multi-filament detections in a single observation. In addition, GRBs occur at an approximate rate of 1 GRB per day, with Fermi-GBM detecting \u2248250 per year with a \u224870% sky coverage (Bhat et al. 2016). Both sources provide similar fluences in the 0.3\u221210 keV energy range, but GRBs can provide this fluence in much shorter integration times. The challenge associated with the use of GRBs as background sources is their transient nature, emitting a considerable percentage of the soft X-ray photons within the first hour of their afterglow phase. Therefore, an instrument capable of having a high efficiency to react on target of opportunity (ToO) events is required to probe the WHIM with GRBs. In spite of this, GRBs should allow Athena to perform its science mission of tracing the missing baryons in GRB afterglow spectra throughout its four-year mission lifetime.","Citation Text":["Willingale & M\u00e9sz\u00e1ros 2017"],"Functions Text":["A bright X-ray background source such as","or a GRB (e.g.","can be used to detect the WHIM by producing absorption features in its energy spectra"],"Functions Label":["Background","Background","Background"],"Citation Start End":[[1151,1177]],"Functions Start End":[[1060,1100],[1136,1150],[1179,1264]]}
{"Identifier":"2021ApJ...910...78Z__Condon_et_al._2017_Instance_1","Paragraph":"The multiwavelength spectral data of 13 SNRs with hard \u03b3-ray spectra. The \u03b3-ray spectra are fitted with a hadronic model with the normalization of the individual spectrum as free parameters. The model assumes that protons have a single power-law energy distribution with an exponential high-energy cutoff. Note that the TeV spectra of G78.2+2.1 (HAWC) and N132D (HESS) cut off at relatively lower energies, and the soft spectral component of GeV of HESS 1912+101 may be from other contributors and are not considered in SED fitting. The best-fit model parameters are indicated in the figure. References for the observational data are as follows: RX J0852.0\u22124622: radio (Duncan & Green 2000), GeV (Tanaka et al. 2011), X-ray (Aharonian et al. 2007), TeV (H.E.S.S. Collaboration et al. 2018c); RX J1713.7\u22123946: radio (Lazendic et al. 2004), X-ray (Tanaka et al. 2008), GeV and TeV (H.E.S.S. Collaboration et al. 2018a); HESS J1731\u2212347: radio (Tian et al. 2008), GeV (Condon et al. 2017; Guo et al. 2018), X-ray (Doroshenko et al. 2017), TeV (H.E.S.S. Collaboration et al. 2011); RCW 86: radio (Clark et al. 1975; Lemoine-Goumard et al. 2012), X-ray (Lemoine-Goumard et al. 2012), GeV (Ajello et al. 2016), TeV (H.E.S.S. Collaboration et al. 2018d); SN 1006: radio Dyer et al. 2009, X-ray (Bamba et al. 2008), GeV (Condon et al. 2017), TeV (Acero et al. 2010); G150.3+4.5: radio (Gerbrandt et al. 2014), X-ray and GeV (Devin et al. 2020); G296.5 + 10.0: radio (Milne & Haynes 1994), GeV (this work), HESS J1534\u2212571: radio (Maxted et al. 2018), GeV (Araya 2017), X-ray and TeV (H.E.S.S. Collaboration et al. 2018b); RCW 103: radio (Dickel et al. 1996), GeV (Xing et al. 2014); G78.2+2.1: radio (Wendker et al. 1991; Zhang et al. 1997; Kothes et al. 2006; Gao et al. 2011), X-ray (Leahy et al. 2013), GeV (Abeysekara et al. 2018) and TeV (Fleischhack 2019); G279.0+1.1: radio (Woermann & Jonas 1988; Duncan et al. 1995), GeV Araya (2020); N132D: radio (Dickel & Milne 1995), X-ray (Hughes et al. 1998; Bamba et al. 2018), GeV (Y. L. Xin et al. 2020, in preparation), and TeV (H.E.S.S. Collaboration et al. 2015).","Citation Text":["Condon et al. 2017"],"Functions Text":["References for the observational data are as follows: RX J0852.0\u22124622:","GeV"],"Functions Label":["Uses","Uses"],"Citation Start End":[[965,983]],"Functions Start End":[[592,662],[960,963]]}
{"Identifier":"2015ApJ...803...17Z__Campbell_&_Lattanzio_2008_Instance_1","Paragraph":"The metallicity calculated for HD 212869 is higher than that estimated prior to using high-resolution spectra. The iron abundance is found to be [Fe\/H] = \n\n\n\n\n\n 0.2 dex on the basis of ionized lines that are almost free from NLTE effects. The carbon abundance was found to be high in the atmosphere of HD 112869, \n\n\n\n\n\n = 8.3 \u00b1  0.1 dex. The nitrogen abundance is \n\n\n\n\n\n = 6.55 \u00b1 0.2 dex. With the obtained abundances [C\/Fe] = +2.2 dex and [C\/N] = +1.15, HD 112869 occupies the region of CEMP-s stars on the plots [C\/Fe] versus [Fe\/H] and [C\/N] versus [Fe\/H] (see Figures 5 and 6 in Campbell & Lattanzio 2008). However, the s-process elements Sr, Y, and Ba are not enhanced significantly, thus confirming by definition (Beers & Christlieb 2005) the CEMP-no status of HD 112869. However, the Nd abundance seems to be enhanced relative to iron, and a similar overabundance was recognized for lanthanum and samarium. From inspection of the Eu ii line at 6645.127 \u212b, the upper limit was set for the r-process element europium, [Eu\/Fe] \u2272 +0.8 dex. According to calculations carried out by Bisterzo et al. (2011), the abundances of three s-peaks are strongly dependent on the choice of the \n\n\n\n\n\n-pocket, as well as on the initial mass and the metallicity. [ls\/Fe], [hs\/Fe] and [Pb\/Fe] do not follow a linear behavior with decreasing metallicity and can cover a large range of values. For the low-mass models the enhancement of the first s-process peak is low or absent. A spectrum of very high resolution is needed to estimate the abundances for a large number of the second s-process peaks and to recognize an overabundance of the third s-process peak. With the adopted oxgen abundance, [O\/Fe] = +0.8 dex, the carbon-to-oxygen ratio was found to be very high for HD 112869, C\/O \n\n\n\n\n\n 12.6. The isotopic lines of C2 and CN are too week to be detected in the crowded spectrum, and the lower limit of isotopic ratio was found to be extremely high, \n\n\n\n\n\n 1500. A large isotopic ratio is not typical of CEMP stars. However, for low-mass AGB stars the CN processing is not expected after the second dredge-up and a total amount of \n\n\n\n\n\n dredged-up during the AGB phase leads to a high [C\/N] ratio and a high \n\n\n\n\n\n ratio observed for HD 112869. On the contrary, intermediate-mass AGB stars with a hot bottom burning should have low both [C\/N] and \n\n\n\n\n\n ratios.","Citation Text":["Campbell & Lattanzio 2008"],"Functions Text":["With the obtained abundances [C\/Fe] = +2.2 dex and [C\/N] = +1.15, HD 112869 occupies the region of CEMP-s stars on the plots [C\/Fe] versus [Fe\/H] and [C\/N] versus [Fe\/H] (see Figures 5 and 6 in"],"Functions Label":["Uses"],"Citation Start End":[[583,608]],"Functions Start End":[[389,582]]}
{"Identifier":"2021AandA...646A..96C__Brusa_et_al._2018_Instance_1","Paragraph":"AGN-driven outflows. Another possible effect of the AGN activity on the molecular gas is through outflows. This possibility is supported by observations of individual objects: For example, Carniani et al. (2017), Brusa et al. (2018) and Loiacono et al. (2019) find low gas fractions in powerful AGN at cosmic noon hosting high-velocity molecular and ionized outflows (but see also Herrera-Camus et al. 2019). AGN feedback in action in these targets could be depleting the molecular gas reservoir (Brusa et al. 2015). F\u00f6rster Schreiber et al. (2019), studying outflows in a large sample of 0.6\u2004\u2004z\u2004\u20042.7 galaxies through integral field spectroscopy of the H\u03b1 emission line, find that incidence, strength, and velocity of AGN-driven winds are strongly correlated with the stellar mass. In particular, they find that high-velocity (\u223c1000\u20132000 km s\u22121) AGN-driven outflows are commonly detected at masses above log(M*\/M\u2299) = 10.7, and present in up to 75% of the population for log(M*\/M\u2299) > 11.2. Interestingly, above this stellar-mass threshold we find a significant CO luminosity deficit in our AGN sample with respect to inactive galaxies (Fig. 3, bottom). Moreover, our AGN show on average gas fractions 0.57 dex (by using uniform assumptions, Sect. 4) lower than inactive galaxies at the 2.2\u03c3 level. Quantitatively, this translates into Mgas,\u2006mol\/M*\u2004\u2248\u20040.3 for AGN (0.16 if we use r31\u2004=\u20040.92; Kirkpatrick et al. 2019) and \u22481 in inactive galaxies. This representative value for our AGN is in line but not as low as previous work targeting extremely powerful sources (e.g., Mgas,\u2006mol\/M*\u2004\u20040.05 in Brusa et al. 2018). Our team is performing a systematic investigation of ionized gas outflows with SINFONI as part of the SUPER survey, and 11 targets of our ALMA sample have complementary good quality SINFONI data (Kakkad et al. 2020; Perna et al., in prep.). For some of them we measured [O\u202fIII] line widths larger than 600 km s\u22121, interpreted as a clear signature of the presence of an AGN-driven outflow in these objects (Kakkad et al. 2020). A detailed comparison between outflow and CO properties for these targets will be presented in a future work. Distinguishing among the scenarios described above is challenging with the current dataset. AGN feedback could proceed in different ways and different mechanisms likely overlap in shaping the properties of the molecular gas reservoir. For example, AGN radiation could both heat and\/or dissociate CO molecules. In this case, AGN would produce a feedback mechanism that does not require outflows but would potentially work toward inhibiting further star formation. As for AGN-driven outflows, they could impact the gas content by ejecting material out of the galaxy (e.g., Travascio et al. 2020), or they could produce CO heating or dissociation due to shocks. Additionally, numerical simulations predict that AGN-driven outflows may heat via shocks a significant quantity of the gas in the ISM, reaching the high temperatures required for the excitation of high-J CO transitions (Costa et al. 2018). To reach a deeper understanding of the impact of AGN on the molecular gas reservoir, also on longer timescales, predictions from simulations providing the spatial scales and effects of AGN activity on CO properties as a function of cosmic time are needed.","Citation Text":["Brusa et al. (2018)"],"Functions Text":["Another possible effect of the AGN activity on the molecular gas is through outflows. This possibility is supported by observations of individual objects: For example, Carniani et al. (2017),","and Loiacono et al. (2019) find low gas fractions in powerful AGN at cosmic noon hosting high-velocity molecular and ionized outflows"],"Functions Label":["Background","Background"],"Citation Start End":[[213,232]],"Functions Start End":[[21,212],[233,366]]}
{"Identifier":"2022MNRAS.511.1121M__Reig_&_Nespoli_2013_Instance_3","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"],"Functions Text":["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)"],"Functions Label":["Background"],"Citation Start End":[[1545,1564]],"Functions Start End":[[1327,1543]]}
{"Identifier":"2019ApJ...885...50W__Falgarone_et_al._2009_Instance_1","Paragraph":"The physical conditions within giant molecular clouds (GMCs) establish the initial conditions for star formation, thus understanding the factors that determine molecular cloud properties is of major interest. A correlation between the size and line width, of the form \u03c3v \u221d R\u03b1 with \u03b1 \u2248 0.5, has long been noted in samples of nearby molecular clouds (Larson 1981; Solomon et al. 1987, hereafter S87). This correlation, hereafter referred to as the R\u2013\u03c3v relation, is usually interpreted as the result of turbulent motions in the interstellar medium on all scales (Mac Low & Klessen 2004; Falgarone et al. 2009). It closely resembles the turbulent cascade with a power-law slope falling between the Kolmogorov (1941) and Burgers (1939) values for incompressible and highly supersonic turbulence, respectively (see also Falgarone et al. 1994; Brunt & Heyer 2002; Federrath 2013; Kritsuk et al. 2013). At the same time, a study of 13CO emission in the Boston University\u2013FCRAO Galactic Ring Survey (GRS) by Heyer et al. (2009) showed that the normalization of the relation, \n\n\n\n\n\n, exhibits a linear correlation with the mass surface density, \u03a3 = M\/\u03c0R2, across more than an order of magnitude in \u03a3. The data are consistent with a state of virial balance between gravity and turbulent motions, except that the observed normalization, v0, is about a factor of 2 too large. Subsequent work by Field et al. (2011) has suggested that the larger than expected v0 may result from external pressure confinement, although to a lesser extent than has been inferred for clouds in the outer Galaxy (Heyer et al. 2001) or near the Galactic Center (Oka et al. 2001; Shetty et al. 2012). On the other hand, Ballesteros-Paredes et al. (2011) have interpreted the Heyer et al. (2009) result in terms of gravitational collapse near freefall, which differs from the virial equilibrium prediction by a factor of \n\n\n\n\n\n in v0 and is thus roughly consistent with the GRS data. A third possibility is that errors in the measured or inferred cloud properties create the appearance of excess kinetic energy when in fact clouds are close to being virialized.","Citation Text":["Falgarone et al. 2009"],"Functions Text":["This correlation, hereafter referred to as the R\u2013\u03c3v relation, is usually interpreted as the result of turbulent motions in the interstellar medium on all scales"],"Functions Label":["Background"],"Citation Start End":[[585,606]],"Functions Start End":[[399,559]]}
{"Identifier":"2018ApJ...856...94Z__Bieber_et_al._1991_Instance_1","Paragraph":"Figures 8 and 9 show the effects of solar activity on the CR parallel \u03bb\u2225 (blue line), perpendicular \u03bb\u22a5 (red line), and radial mean free path \u03bbrr (gray line) for a proton with rigidity 445 MV (corresponding to a 100 MeV proton) for the inwardly and outwardly directed IMF, respectively. As described in Zank et al. (1998), the parallel mean free path (mfp) based on standard QLT and assuming magnetostatic turbulence is approximated by\n13\n\n\n\n\n\nwhere \n\n\n\n\n\n, \n\n\n\n\n\n, and \n\n\n\n\n\n. RL is the particle Larmor radius, P is the particle rigidity, and B0 is the mean magnetic field strength. The analytic form of the perpendicular mfp based on NLGC theory is given by (Zank et al. 2004; Shalchi et al. 2010)\n14\n\n\n\n\n\nwhere a2 = 1\/3 is a factor related to the gyrocenter velocity. \n\n\n\n\n\n is a constant such that \u03bd = 5\/6 yields a Kolmogorov (1941) spectrum. Note that Equation (14) was derived under the assumption of a specific form of 2D wave spectrum, which is a constant at large turbulence scales. It means that the 2D turbulence spectrum is independent of wavenumber in the energy range in Equation (14). Observation of magnetic fluctuations in the SW indicates that omnidirectional power spectra approach a k\u22121 wavenumber dependence and that low-frequency turbulence exhibits some sunspot cycle variability (Bieber et al. 1991). Based on this, Engelbrecht & Burger (2015) derived the perpendicular mfp by specifying the energy range spectral index of 2D turbulence power spectra as \u22121. A more general form of the 2D power spectrum with an energy range spectral index q was proposed by Shalchi et al. (2010). They show that the spectral index has a strong influence on the perpendicular diffusion coefficient. In their model, negative values of q correspond to a decreasing spectrum in the energy range, q = 0 corresponds to the constant spectrum we use here, and positive values of q correspond to an increasing spectrum. Matthaeus et al. (2007) presented a similar spectrum in different regimes: energy range, inertial range, and intermediate regime where the spectrum is proportional to k\u22121 to coincide with observations (Bieber et al. 1991; Goldstein & Roberts 1999). However, Shalchi (2013) argues that a spectrum that behaves like k\u22121 does not provide a different perpendicular diffusion coefficient (see also Shalchi et al. 2010), since the field lines for such length scales behave superdiffusively as in the inertial range (Shalchi & Kourakis 2007). In view of this uncertainty, we do not take into account a more elaborate spectrum in the present paper. The behavior of the 2D wave spectrum in the energy range, which may also be correlated with the sunspot cycle, is an important factor in deriving the CR perpendicular mfp. A general form (e.g., Shalchi et al. 2010; Shalchi 2013) should be employed in future studies of CR diffusion.","Citation Text":["Bieber et al. 1991"],"Functions Text":["Observation of magnetic fluctuations in the SW indicates that omnidirectional power spectra approach a k\u22121 wavenumber dependence and that low-frequency turbulence exhibits some sunspot cycle variability"],"Functions Label":["Background"],"Citation Start End":[[1303,1321]],"Functions Start End":[[1099,1301]]}
{"Identifier":"2020ApJ...892L..10Y__Macchi_2013_Instance_1","Paragraph":"In this section, we consider the plasma properties under the propagation of strong waves. In strong waves, the motion of electrons in the plasma becomes relativistic. However, different from free electrons that have a relativistic drift velocity in the direction of the incident electromagnetic wave (see Section 3.1), in plasma the space-charge potential is important in preventing the drift of electrons (Waltz & Manley 1978). For nonrelativistic electrons in plasma, if the wave duration \u03c4 is much larger than c\/\u03c9p, where \n\n\n\n\n\n is the plasma frequency, the drift velocity would be close to zero (Waltz & Manley 1978; Sprangle et al. 1990b). In this case, electrons in plasma under a strong wave would have a typical Lorentz factor (\n\n\n\n\n\n) similar to that (\u03b3) in the laboratory frame, so that \n\n\n\n\n\n is satisfied. Due to the relativistic and magnetic force effects, the propagation and dispersion properties of an electromagnetic wave depend on its amplitude. For a circular polarized wave, the dispersion relation in the laboratory frame is given by (e.g., Gibbon 2005; Macchi 2013; Macchi et al. 2013; see the Appendix)\n11\n\n\n\n\n\nThe dispersion relation of strong electromagnetic waves is altered due to the effective electron mass increased by the relativistic effect (e.g., Sarachik & Schappert 1970; Gibbon 2005; Macchi 2013). One can define the effective plasma frequency as\n12\n\n\n\n\n\nso that the wave can propagate in the region where \n\n\n\n\n\n. With respect to the nonrelativistic linear case, this is known as relativistically self-induced transparency. We note that since the dispersion depends on the electromagnetic field amplitude in the nonlinear case, the dispersion relation must be taken with care. The propagation of a pulse will be affected by the complicated effects of nonlinear propagation and dispersion, and finally the spatial and temporal shape of the pulse itself would also be modified. In particular, for linear polarization, the relativistic factor \u03b3 is not a constant (see Section 3.1). The propagation of the linearly polarized wave with a relativistic amplitude would lead to generation of the higher-order harmonics. Sprangle et al. (1990b) proved that the propagation of the first harmonic component, i.e., of the \u201cmain\u201d wave, is still reasonably described by Equation (11) with \n\n\n\n\n\n. Thus, we will directly adopt Equation (11) in the following discussion.","Citation Text":["Macchi 2013"],"Functions Text":["For a circular polarized wave, the dispersion relation in the laboratory frame is given by (e.g.,"],"Functions Label":["Uses"],"Citation Start End":[[1075,1086]],"Functions Start End":[[964,1061]]}
{"Identifier":"2016ApJ...822...37S__Freeman_et_al._2001_Instance_1","Paragraph":"Each pulsar system in this analysis was observed with XMM-Newton (Jansen et al. 2001), using the European Photon Imaging Camera (EPIC) with the pn detector in full frame mode with thin filters (the data from the MOS and RGS detectors had insufficient counts for analysis). PSR J0337 was observed on 2013 August 1 (observation number 0722920101) for 16.2 ks. An X-ray source was detected 16 away from the radio position of Ransom et al. (2014), consistent with the 2\u2033 astrometric precision of XMM;12\n\n12\nSee xmm.vilspa.esa.es\/docs\/documents\/CAL-TN-0018.pdf.\n we show an image of the detection in Figure 1. We measured 164 \u00b1 13 background-subtracted counts between 0.2 and 2.0 keV, as determined using calc_data_sum in Sherpa (Freeman et al. 2001; Doe et al. 2007) and uncertainty given by a Poisson distribution.13\n\n13\nNote that our count-rate for PSR J0337 is below the 2\u03c3 upper limit from Prinz & Becker (2015), who analyzed the same data set. Nonetheless, we are confident in our detection (Figure 1) and do not know the reason for the discrepancy.\n The chance coincidence probability, given the number of sources in the field with similar or higher count rates, is approximately 8 \u00d7 10\u22125. PSR J0636 was observed on 2013 October 13 (observation number 0722920201) for 15.0 ks, and we found an X-ray source within 03 of the radio position of Stovall et al. (2014); see Figure 1. The chance coincidence probability for PSR J0636 is also approximately 8 \u00d7 10\u22125. We measured 170 \u00b1 13 counts between 0.2 and 2.0 keV. Finally, PSR J0645 was observed on 2014 March 29 (observation number 0722920301) for 34.9 ks, but removing a flare from the data reduced the effective observation length to 23 ks. No source was found by the XMM pipeline near the radio position of Stovall et al. (2014, see Figure 1), and we measured only 18 \u00b1 9 source counts between 0.2 and 2.0 keV. The time resolution of 73.4 ms was too coarse to detect pulsations at the rotational periods of the pulsars (2.73, 2.87, and 8.85 ms; Ransom et al. 2014; Stovall et al. 2014), but the observed flux can guide future searches for pulsed X-rays. We reprocessed the data using SAS v13.0.1, specifically epchain. Using HEAsoft v6.14 and CIAO v4.6, and some custom scripts, we extracted the source counts from within a radius of 25\u2033, and background counts from an annular region with radii of 50\u2033 and 125\u2033, restricted to the same CCD chip with other sources removed. We limited the data to events with PATTERN \u2264 4 (singles and doubles), but also experimented with using PATTERN \u2264 12 (singles, doubles, and triples). We found that the change in the results when including triple events was negligible. Because of the high background rate at low energies and the expected softness of the source spectra, we limited our analysis to energies between 0.2 and 2.0 keV. We grouped the counts such that each energy bin had at least 15 events in it and subtracted the background from the source.","Citation Text":["Freeman et al. 2001"],"Functions Text":["We measured 164 \u00b1 13 background-subtracted counts between 0.2 and 2.0 keV, as determined using calc_data_sum in Sherpa"],"Functions Label":["Uses"],"Citation Start End":[[725,744]],"Functions Start End":[[605,723]]}
{"Identifier":"2022AandA...667A..69S__Hut_1985_Instance_1","Paragraph":"Finally, it is worth noting that two-body relaxation and tidal interactions affect the IMF in different ways. The former causes mass segregation with higher-mass stars moving inward and lower-mass stars outward, which enables the lower-mass stars to evaporate from the cluster (see also, e.g., Chandrasekhar 1942; Spitzer 1969; Binney & Tremaine 2008). The cluster then forms a dense core of high-mass stars and once it starts to collapse, binaries form in the centre. These perturb their neighbours, and we begin to observe massive stars escaping from the mass segregated core (see also Hills 1975; Hut 1985). On the other hand, if the relaxation timescales are long, the cluster does not segregate as quickly and the tidal harassment is the dominant reason for mass-loss. Stars are peeled-off from the outer regions independent of their masses. Consequently, the relative contributions of these two processes determine the mass distribution of the escapers. This is shown in Fig. 5, where the mass distribution of escaping stars divided by the original IMF in models R07_Sal13 and R07_Sal13_Nb. When the environment is present, the tidal shocks enhance fraction of escaping stars at intermediate masses. Instead, the loss of massive stars, similar in the two models, can be linked to the relaxation process ongoing in the core that is composed almost exclusively of massive stars and binaries at later times. In our model, the percentage of mass loss is small so this does not significantly affect the mass function of the remaining members. However, the same process happens in protoclusters still embedded in the parent cloud. Their gas density is orders of magnitude higher than our n\u2004=\u200410\u2006cm\u22123, and structures are closer to the cluster (Kruijssen et al. 2012; Kruijssen 2012). At the same time, the IMF of massive protoclusters is much wider, since massive stars are still on the main sequence, so that the evolutionary timescale is also reduced (Allison et al. 2009; Yu et al. 2011). In the end, the effects we described could indeed play a crucial role when inferring the IMF of young clusters.","Citation Text":["Hut 1985"],"Functions Text":["The cluster then forms a dense core of high-mass stars and once it starts to collapse, binaries form in the centre. These perturb their neighbours, and we begin to observe massive stars escaping from the mass segregated core"],"Functions Label":["Background"],"Citation Start End":[[600,608]],"Functions Start End":[[353,577]]}
{"Identifier":"2019AandA...628A.110M__Kryukova_et_al._(2012)_Instance_3","Paragraph":"Deriving the completeness limits of the WISE photometry is mandatory to assess the reliability of our catalogue of starless cores. We examined the histograms of the number of mid-infrared (MIR) sources versus magnitude; taking into account the effects of the cuts required to fulfil the criteria of Koenig et al. (2012), rough completeness limits are [3.6] ~ 14, [4.6] ~ 12, [12] ~ 9 and [22] ~ 7. These values are 1\u20133 mag brighter than the sensitivity limits quoted in the WISE Explanatory Supplement3 for the relevant sky region. Once converted into flux units and, for example, compared with the models of Class I and Class 0 sources of 0.5 M\u2299 by Whitney et al. (2004), it can be seen that the completeness limits at 3.6 and 4.6\u03bcm are faint enough to detect such objects taking into account a distance of 700 pc and a further foreground reddening up to AV = 20. Even in the worst case of edge-on discs, these objects would be detectable at 3.6 and 22\u03bcm. Furthermore, the completeness limit at 22 \u03bcm is faint enough to allow detection of Class I and Class 0 sources of even-lower-mass central objects. Alternatively, one can compute the bolometric luminosity following Kryukova et al. (2012). Starting from our completeness limit at 22 \u03bcm, after conservatively dereddening it by AV = 20, we assumed a spectral index \u03b3 = \u22122 (see Table 3 for definition) to compute the MIR luminosity from Eq. (6) of Kryukova et al. (2012). Equation (7) of Kryukova et al. (2012) then yields Lbol ~ 1.7\u20132.8 L\u2299, depending on whether the NIR flux is neglected (which may be the case) or extrapolated from \u03b3 = \u22122. A comparison with the birthline of Palla & Stahler (1993) indicates a mass of ~ 0.4\u20130.5 M\u2299 for the central protostar. For the sake of comparison, we can roughly estimate the completeness limit in central masses of the Herschel protostellar cores in Giannini et al. (2012) using their quoted completeness limit at 70 \u03bcm of 0.21 Jy and following Dunham et al. (2008). By using Eq. (2) of Dunham et al. (2008), scaled to a distance of 700 pc, we found that the flux density at 70 \u03bcm translates into a bolometric luminosity of the central (proto)star Lbol ~ 0.28 L\u2299 (we note that Dunham et al. 2008 indicate this luminosity as Lint). We highlight the fact that the 70 \u03bcm emission is in principle a more sensitive protostellar tracer than WISE. However, this contrasts with the much lower number of protostellar cores found by Giannini et al. (2012), which may be due to a poorer effective sensitivity because of their selection criteria.","Citation Text":["Kryukova et al. (2012)"],"Functions Text":["Equation (7) of","then yields Lbol ~ 1.7\u20132.8 L\u2299, depending on whether the NIR flux is neglected (which may be the case) or extrapolated from \u03b3 = \u22122."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1440,1462]],"Functions Start End":[[1424,1439],[1463,1593]]}
{"Identifier":"2020MNRAS.496.3582C__Komacek_&_Showman_2016_Instance_1","Paragraph":"3D General Circulation Models (GCMs) with simplified thermal forcing are one possible intermediate step between 3D GCMs with full coupling between radiation and dynamics like those used by Showman et al. (2009) and Amundsen et al. (2016), and shallow water models, i.e. atmosphere models with one atmosphere layer comprising vertically averaged flow (Showman, Cho & Menou 2010). Fully coupled GCMs have the highest accuracy in stellar radiation and flow coupling and thus the highest predictive power. They are, however, computationally much more expensive and their complexity makes it more difficult to test underlying modelling assumptions compared to GCMs with simplified thermal forcing. The latter are thus better suited to run simulations for various scenarios, to understand large-scale flow and circulation properties in 3D climate models under different conditions (Liu & Showman 2013; Mayne et al. 2014; Tsai et al. 2014; Carone et al. 2015, 2016; Komacek & Showman 2016; Hammond & Pierrehumbert 2017; Mayne et al. 2017). Such models have been proven to be very useful: superrotation in hot Jupiters was first inferred by Showman & Guillot (2002) in a 3D GCM with Newtonian cooling. Recently, Showman, Tan & Zhang (2019) used Newtonian cooling to establish a clean, simple environment to diagnose flow dynamics in brown dwarfs, Jupiter, and Saturn-like planets. Shallow water models present an even simpler model framework and represent 3D flow patterns in an atmosphere depth-dependent (2D) formalism (Showman & Polvani 2010, 2011; Penn & Vallis 2017). There are other useful radiative forcing parametrizations such as those using the dual-band radiative scheme, which can also explore a large parameter space and basic assumptions (see e.g. the model used by Komacek et al. 2017). Generally, a hierarchy of models with various levels of complexity has proven to be extremely beneficial to understand complex flow patterns in full 3D climate simulations. Here, we establish a clean, simple environment to understand possible dynamical feedback between the lower boundary and observational flow via Newtonian cooling.","Citation Text":["Komacek & Showman 2016"],"Functions Text":["They are, however, computationally much more expensive and their complexity makes it more difficult to test underlying modelling assumptions compared to GCMs with simplified thermal forcing. The latter are thus better suited to run simulations for various scenarios, to understand large-scale flow and circulation properties in 3D climate models under different conditions"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[959,981]],"Functions Start End":[[502,874]]}
{"Identifier":"2021ApJ...913...55H__Goldman_et_al._2017_Instance_2","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"],"Functions Text":["In addition, recent observational and theoretical studies on RSGs and SNe II indicate that RSG wind mass-loss rates may be independent from metallicity"],"Functions Label":["Background"],"Citation Start End":[[988,1007]],"Functions Start End":[[835,986]]}
{"Identifier":"2015ApJ...804..130C___2013_Instance_1","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"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[869,890]],"Functions Start End":[[678,867]]}
{"Identifier":"2019ApJ...887..137S__Vekstein_2017_Instance_1","Paragraph":"As mentioned above, the magnetic reconnection is introduced as breaking and reconfiguration of the oppositely directed magnetic field lines in highly conducting plasma. The magnetic field lines collapse near the X-point and form an extended magnetic singularities known as a current sheet. There are two mechanism of the current-sheet formation. The first kind of current-sheet formation is associated with the MHD instabilities (e.g., resistive tearing mode and ideal kink mode) known as spontaneous magnetic reconnection (e.g., White 1984; Baty 2000; Vekstein 2017). The second kind of current sheet can be formed in the MHD stable configuration, where some external perturbations trigger the forced magnetic reconnection (Hahm & Kulsrud 1985). The forced magnetic reconnection may be activated by nonlinear MHD waves, which may be caused by explosive solar activities (e.g., Sakai et al. 1984; Dewar et al. 2013; Beidler et al. 2017). The forced magnetic reconnection may be developed due to boundary perturbations, which induce a surface current in such a way that it opposes the progress of the reconnection (Ishizawa & Tokuda 2000, 2001; Fitzpatrick 2003). The multimode simulation approach has been adopted to investigate the thinning of the current sheet induced by forced magnetic reconnection (Birn et al. 2005). The motion of the photospheric footpoints of the coronal magnetic field may also trigger the forced magnetic reconnection, which may be caused by the explosive solar coronal events (e.g., Vekstein & Jain 1998; Jain et al. 2005; Vekstein 2017). Although there is a remarkable development in the theory of the forced magnetic reconnection, Jess et al. (2010) have suggested that there is no observational evidence of explosive flare or coronal activities triggered by forced magnetic reconnection. They have observed a microflare activity driven by forced magnetic reconnection. The lower solar atmosphere (photosphere & chromosphere) is dominated by cool, partially ionized and collision dominated plasma. Most of the energy releases during the forced magnetic reconnection may be consumed by such plasma systems (e.g., Litvinenko 1999; Chen et al. 2001; Chen & Ding 2006; Litvinenko et al. 2007).","Citation Text":["Vekstein 2017"],"Functions Text":["The first kind of current-sheet formation is associated with the MHD instabilities (e.g., resistive tearing mode and ideal kink mode) known as spontaneous magnetic reconnection (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[553,566]],"Functions Start End":[[346,529]]}
{"Identifier":"2017AandA...606A..17M__Kennicutt_(1998)_Instance_1","Paragraph":"The SFR reported in Table C.1 refers to a stellar mass range from Mlow = 0.1M\u2299 to Mup = 100M\u2299, is averaged over the past \u0394t = 100 Myr, and was calculated using the standard SFR(LIR) relationship from Kennicutt (1998; here scaled to a Chabrier 2003, IMF) (1)\\begin{equation} \\label{eq:sfr} \\textit{SFR}=10^{-10}\\times L_{\\rm IR}[{L}_{\\sun}]\\, {M}_{\\sun}~{\\rm yr}^{-1}. \\end{equation}SFR=10-10\u00d7LIR[L\u2299]\u2009M\u2299yr-1.This calibration relies on the starburst synthesis models of Leitherer & Heckman (1995), and it is based on the assumption of solar metallicity, and an optically thick (\u03c4dust \u226b 1) starburst region, in which case LIR is a good proxy of the system\u2019s bolometric luminosity (LIR \u2243 Lbol), and hence a sound, calorimetric probe of the obscured, current stellar birth rate. A possible caveat is that the contribution to the dust heating by more evolved stellar populations (the cirrus component; e.g. Helou 1986; Lonsdale Persson & Helou 1987; Walterbos & Greenawalt 1996) is not taken into account. If the cirrus ISM component heated by the more general galactic UV radiation field contributes to LIR, then the Kennicutt (1998) relationship overestimates the SFR. Another issue is the fact that some percentage of the UV photons can escape the starburst region without being absorbed, and hence are not reprocessed into IR photons (indeed, some of our SMGs are visible in the rest-frame UV images; Miettinen et al. 2017b). The MAGPHYS code also gives the SFR as an output, and contrary to the aforementioned LIR diagnostic, the model permits for the heating of the dust by older and longer-lasting stellar populations. We found that the SFR(LIR) is somewhat higher on average than SFRMAGPHYS: the SFR(LIR) \/SFRMAGPHYS ratio was found to range from 0.47 to 6.92 with a median of \\hbox{$1.31^{+0.83}_{-0.17}$}1.31-0.17+0.83, where the \u00b1 errors represent the 16th\u201384th percentile range (see the corresponding panel in Fig. 2). If, instead of \u0394t = 100 Myr, the aforementioned comparison is done by using the SFRMAGPHYS values averaged over the past \u0394t = 10 Myr, the median SFR(LIR) \/SFRMAGPHYS ratio is found to be \\hbox{$1.15^{+0.38}_{-0.27}$}1.15-0.27+0.38, which is consistent with the results obtained by da Cunha et al. (2015). Unless otherwise stated, in our subsequent analysis we use the SFR averaged over the past 100 Myr as calculated using Eq. (1). ","Citation Text":["Kennicutt (1998"],"Functions Text":["he SFR reported in Table C.1 refers to a stellar mass range from Mlow = 0.1M\u2299 to Mup = 100M\u2299, is averaged over the past \u0394t = 100 Myr, and was calculated using the standard SFR(LIR) relationship from","here scaled to a Chabrier 2003, IMF) (1)\\begin{equation} \\label{eq:sfr} \\textit{SFR}=10^{-10}\\times L_{\\rm IR}[{L}_{\\sun}]\\, {M}_{\\sun}~{\\rm yr}^{-1}. \\end{equation}SFR=10-10\u00d7LIR[L\u2299]\u2009M\u2299yr-1."],"Functions Label":["Uses","Uses"],"Citation Start End":[[200,215]],"Functions Start End":[[1,199],[217,407]]}
{"Identifier":"2017ApJ...835....2X___2003_Instance_1","Paragraph":"On the other hand, a clear physical interpretation of the observed pulse broadening phenomenon requires a good understanding of the interstellar electron density structure. A power-law model of electron density fluctuations is commonly adopted in theoretical constructions on radio wave propagation (Lee & Jokipii 1976; Rickett 1977, 1990) and is compatible with observational indications (e.g., Armstrong et al. 1995). Recent advances in understanding the properties of magnetohydrodynamic (MHD) turbulence (Goldreich & Sridhar 1995; Lithwick & Goldreich 2001; Cho & Lazarian 2002, 2003) stimulate a renewed investigation on density statistics (Beresnyak et al. 2005; Kowal et al. 2007; Lazarian et al. 2008; Burkhart et al. 2009, 2010, 2015; Collins et al. 2012; Federrath & Klessen 2012), which provide important insight into key physical processes such as star formation in the turbulent and magnetized ISM (see reviews by, e.g., McKee & Ostriker 2007; Lazarian et al. 2015). The density spectrum in compressible astrophysical fluids was systematically studied in Kowal et al. (2007) by carrying out an extensive set of MHD numerical simulations with varying compressibility and magnetization. Instead of a single Kolmogorov slope with a power-law index of \n\n\n\n\n\n, significant variations in the spectral slope of density fluctuations are present. For supersonic turbulence, their results are consistent with earlier findings in both magnetized (Beresnyak et al. 2005) and nonmagnetized (Kim & Ryu 2005) fluids. It shows that the density power spectrum becomes shallower as the sonic Mach number (\n\n\n\n\n\n) increases, where VL is the turbulent velocity at the outer scale of turbulence and cs is the sound speed in the medium, and there is a significant excess of density structures at small scales in highly supersonic turbulence. This behavior is naturally expected as the gas is compressed in shocks by supersonic flows and the interacting shocks produce local density enhancements (Mac Low & Klessen 2004; Padoan et al. 2004b).","Citation Text":["Cho & Lazarian","2003"],"Functions Text":["Recent advances in understanding the properties of magnetohydrodynamic (MHD) turbulence","stimulate a renewed investigation on density statistics","which provide important insight into key physical processes such as star formation in the turbulent and magnetized ISM"],"Functions Label":["Motivation","Motivation","Motivation"],"Citation Start End":[[562,576],[583,587]],"Functions Start End":[[420,507],[589,644],[792,910]]}
{"Identifier":"2018AandA...614A..86G__Hamann_et_al._(2006)_Instance_1","Paragraph":"In the supersonic part of WR outflows, the presence of instabilities and inhomogeneities should be taken into account. Invoking density inhomogeneities, or \u201cclumping\u201d, in stellar models has an influence on the mean opacity, due to the enhanced density in clumps (Moffat et al. 1988; Hamann & Koesterke 1998). However, a porous structure could also imply lowered mean opacity, counteracting the effect of small-scale clumping (Shaviv 1998; Oskinova et al. 2007). Assuming that the material near the Fe-opacity peak is clumped, Gr\u00e4fener et al. (2012) were able to extend the surface radii of their hydrostatic helium star models, in practical terms, by significantly enhancing the iron bump opacity (see also Appendix A). The surface temperatures of such hydrostatic, strongly inflated models computed with plane parallel grey atmospheres are then compared with the fictitious effective temperatures at \u03c4 = 20 of the atmosphere calculations performed by Hamann et al. (2006). This effective temperature at \u03c4 = 20 should not be confused with the actual blanketed temperature at the base of WNE wind models, which is about a factor of 2 larger (see e.g. Appendix D). As we show in Sect. 4, strongly inflated hydrodynamic solutions imply supersonic flows already at the base of the inflated envelopes for the typical mass-loss rates of WNE stars (cf. the plane parallel atmosphere model in Fig. 4). Moreover, it would be inconsistent to think that the iron opacity bump is simultaneously responsible for both the inflation of the envelope and the acceleration of the flow. While the detailed effects of clumping and porosity remain a subject for future research, we concentrated in this work on the homogeneous case as those instabilities are expected to initiate above the sonic point (Sundqvist et al. 2013; Owocki 2015). If this is not the case and clumping is already present at the relatively high densities and optical depth of the sonic point of massive helium star models, it might affect to some extent the local opacity.","Citation Text":["Hamann et al. (2006)"],"Functions Text":["The surface temperatures of such hydrostatic, strongly inflated models computed with plane parallel grey atmospheres are then compared with the fictitious effective temperatures at \u03c4 = 20 of the atmosphere calculations performed by"],"Functions Label":["Uses"],"Citation Start End":[[952,972]],"Functions Start End":[[720,951]]}
{"Identifier":"2021MNRAS.508.4332M__Draine_1978_Instance_1","Paragraph":"As a first step, we need to specify the photoelectron sheath features. In this course, we first evaluate the steady state potential over the lunar surface (equation 4), and then after, we use this as a boundary condition to solve the Poisson equation (equation 2) and estimate the photoelectron sheath profile. In calculations, Lyman \u03b1 (\u03bb \u223c 121.57 nm, 10.29 eV, \u039b \u223c 3 \u00d7 1011 cm\u20132 s\u20131) spike of solar photon radiation (Bauer 1973) is considered as the dominant source for the generation of photoelectrons from the lunar surface. The work function of the regolith material is taken from Grobman & Blank (1969), where it is suggested to vary in the range \u03d5r \u223c (4\u20136) V for the region across the subsolar point and limb. Moreover, Draine\u2019s formulation is accounted to determine the lunar surface\u2019s photoelectric efficiency (Draine 1978; Draine & Salpeter 1979) \u2013 its spectral dependence can be represented as ${\\chi _{\\nu r}} = {\\chi _o}[1 - ({\\phi _r}\/{E_\\nu })]$. For instance, for the Lyman \u03b1 radiation \u03c7\u03bdr = 0.042 for optimum efficiency \u03c7o = 0.1 (Sickafoose et al. 2001) and \u03d5r = 6 V (Grobman & Blank 1969). Another significant parameter is the surface temperature which describes the electron population within the lattice available for the photoemission. Lunar Reconnaissance Orbiter based measurements (Williams et al. 2017) suggest that the surface temperature may vary from the equator (\u223c400 K) to the terminator (poles, \u223c150 K). In order to take this account, we use the latitude (\u03b8) dependent empirical relation ${T_\\theta } = {T_0}[1 - (5\/4\\pi )\\theta ]$; for instance, at \u03b8 = 70\u00b0, and To \u2248 205 K. These three parameters, viz., \u03d5r, To, \u03c7\u03bd, and \u039b drive the photoemission current from the lunar regolith. The nominal solar wind plasma parameters are considered for calculating collection current over lunar regolith; the constituents are considered as40-41nes \u2248 nis = 8.7 cm\u20133 and Tes \u2248 Tis = 1.4 \u00d7 105 K (Mann et al. 2011; Kureshi et al. 2020). These solar radiation and wind plasma parameters might vary widely during active solar events and alter surface charging and sheath features. Popel et al. (2018) suggest the dust number density may also vary in a wide range depending on lunar altitude and particle size; for instance, nd \u223c 800 cm\u20133 for the particles of size 100 nm \u2264 ao \u2264 200 nm and \u03b8 = 77\u00b0. Note that the secondary electron emission (Seitz 1940; Misra, Mishra & Sodha 2013) from the lunar regolith (and floating dust) is ignored, as it minimally contributes to the charging of sunlit surfaces (Mishra & Bhardwaj 2020). These parameters, along with equation (4), yield steady state potential over the sunlit locations. This estimate of the surface potential is used as a boundary condition (i.e. at l = 0, \u03c5 = \u03c5o) alongwith \u03c5\u2019 = \u03c5 = 0 as $l \\to \\infty $ to solve the Poisson equation (equation 2) numerically \u2013 using this framework, the sheath structure is derived in terms of electric potential (\u03c5), electric field (Es), and photoelectron population density (npe).","Citation Text":["Draine 1978"],"Functions Text":["Moreover, Draine\u2019s formulation is accounted to determine the lunar surface\u2019s photoelectric efficiency (",") \u2013 its spectral dependence can be represented as ${\\chi _{\\nu r}} = {\\chi _o}[1 - ({\\phi _r}\/{E_\\nu })]$."],"Functions Label":["Uses","Uses"],"Citation Start End":[[819,830]],"Functions Start End":[[716,819],[854,960]]}
{"Identifier":"2015ApJ...815...15W__Dib_&_Kaspi_2014_Instance_1","Paragraph":"1E 1841\u2013045, located at the center of supernova remnant (SNR) Kes 73 (Vasisht & Gotthelf 1997), is another steady X-ray emitter. The source was monitored with NuSTAR between 2013 September 5 and 21, during which time six bursts were detected (An et al. 2015). In order to eliminate the contamination from the SNR, an annulus region with inner and outer radii of 60\u2033 and 100\u2033 around the source position was selected for background spectral extraction, as was also done in An et al. (2013). We searched for soft X-ray data from 2013 August 1 to October 30, and identified 11 Swift XRT observations, 10 executed in WT mode and 1 short observation in Photon Counting (PC) mode. Note that in WT mode, 10 rows are compressed into a single row, only the central 200 columns are read, and only 1D imaging is preserved. Consequently, it is impossible to resolve the SNR contribution and three SNR related emission lines (Mg, Si, and S) clearly appear in the stacked spectrum of 10 WT data. Since the source emission has been stable for more than 15 years (Dib & Kaspi 2014), we performed the joint fit with the longest, burst-free NuSTAR observations (ObsID = 30001025012) and XMM-Newton observations that were performed \u223c11 years earlier. For the XMM-Newton observation, since the MOS1\/MOS2 data were performed in the full frame mode and were seriously affected by pile up, we only use the pn data in the following spectral fitting. The joint spectra are fit by our model (STEMS3D+PL) and a 2% systematic error is added to the XMM-Newton data to account for calibration uncertainties. The fit results yielded that 1E 1841\u2013045 has the lowest twist angle (\u0394\u03d5 \u223c 0.83) and the highest electron velocities in the magnetosphere (\u03b2 \u223c 0.28) among the four sources studied here. The resulting photon index of the PL component (\u0393 \u223c 1.23) is consistent with the value reported based on INTEGRAL data (Kuiper et al. 2006). However, due to the short exposure of the XMM-Newton data, in addition to the high interstellar absorption, the parameter values have large uncertainties (see Table 2 and Figure 2). Note that 1E 1841\u2013045 is the only source whose non-thermal component (FP \u223c 5.1 \u00d7 10\u221211 erg s\u22121 cm\u22122) dominates the STEMS3D component (FS \u223c 3.6 \u00d7 10\u221211 erg s\u22121 cm\u22122) in the range of 1\u201379 keV.","Citation Text":["Dib & Kaspi 2014"],"Functions Text":["Since the source emission has been stable for more than 15 years","we performed the joint fit with the longest, burst-free NuSTAR observations (ObsID = 30001025012) and XMM-Newton observations that were performed \u223c11 years earlier."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1047,1063]],"Functions Start End":[[981,1045],[1066,1230]]}
{"Identifier":"2019MNRAS.488.5029H__Malhotra_et_al._2001_Instance_1","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":["Malhotra et al. 2001"],"Functions Text":["However, LC\u2009II\/LFIR is not constant, but declines with increasing LFIR, known as the \u2018[C\u2009ii] deficit\u2019"],"Functions Label":["Background"],"Citation Start End":[[943,963]],"Functions Start End":[[809,910]]}
{"Identifier":"2016AandA...591A..38C__Roediger_et_al._(2011a)_Instance_1","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"],"Functions Text":["Despite the considerable scatter in both colors and color gradients","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."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[115,136]],"Functions Start End":[[0,67],[139,291]]}
{"Identifier":"2015AandA...581A..31S__Giodini_et_al._2013_Instance_1","Paragraph":"Although the origin and evolution of linear-scale clustering is well described by the concordance model (Spergel et al. 2007), gravitational clustering of matter on smaller scales (galaxy clusters and groups) belongs to a non-linear regime of structure formation. This regime is more difficult to understand and to simulate because its evolution must include the role of baryons, which are driven by complex physics. Clusters of galaxies that are the most massive gravitationally bounded structures have been widely used over the past years to probe the cosmic evolution of the large-scale structures in the Universe (Voit 2005; Allen et al. 2011). In the standard model of structure formation driven by gravitation alone, clusters form a self-similar population that is only characterized by their mass and redshift. Including baryon physics introduces some distortions in the scaling relations between the mass and other physical quantities such as temperature, X-ray, or optical luminosity (Kaiser 1986; Giodini et al. 2013). Most recent research works have focused on the relationship between the dominant dark matter and the baryonic matter that forms gas and stars (Lin et al. 2003; Giodini et al. 2009). Both the mass-to-light (M\/L) ratio of structures and the halo occupation number (HON, or the number of satellite galaxies per halo) correspond to observables that are easy to compare to predictions from numerical simulations (Cooray & Sheth 2002; Tinker et al. 2005). They are both representative of the way stellar formation occurred in the early stages of halo formation (Marinoni & Hudson 2002; Borgani & Kravtsov 2011). Recent progress on numerical simulations (Murante et al. 2007; Conroy et al. 2007; Aghanim et al. 2009) has also stressed the role of hierarchical building of structures in enriching the intra-cluster medium (ICM) with stars in a consistent way with the observed amount of ICM globular clusters and ICM light. This ICM light, although hardly detectable, can be considered as the extension of the diffuse envelope often seen in the central galaxy in rich clusters of galaxies. It is an important component, although not the only one, that explains the formation of the brightest cluster galaxies (BCG) in the centre of clusters of galaxies (Dubinski 1998; Presotto et al. 2014). ","Citation Text":["Giodini et al. 2013"],"Functions Text":["Including baryon physics introduces some distortions in the scaling relations between the mass and other physical quantities such as temperature, X-ray, or optical luminosity"],"Functions Label":["Background"],"Citation Start End":[[1007,1026]],"Functions Start End":[[818,992]]}
{"Identifier":"2018ApJ...867..123M__Helling_&_Fomins_2013_Instance_1","Paragraph":"Regardless of their exact composition, particles suspended in these exoplanet environments likely undergo repeated particle\u2013particle collisions in response to atmospheric circulation. Such dynamics have been inferred to drive efficient triboelectrification (e.g., Helling et al. 2013), resulting in electrified cloudy or hazy environments at elevation. As happens within Earth\u2019s clouds, exoplanet clouds are likely gravitationally stratified, meaning that smaller, lighter grains become concentrated at the top of the clouds, while larger, heavier grains remain at lower elevation (see Figures 2 and 8 in Helling et al. 2008; also Helling & Fomins 2013; Helling et al. 2016). Because, as discussed above, the polarity of charge collected by particles from triboelectric processes depends on grain size, this stratification (smaller particles at elevation; larger particles on the bottom) has the ability to set up coherent electric fields. Such charge separation occurs in both thunderstorms and volcanic plumes and ultimately drives spark discharges\u2013lightning\u2013either through conventional breakdown of the gas or via runaway electron avalanche (Kikuchi & Endoh 1982; Gurevich et al. 1992; James et al. 1998; Dwyer 2005; Cimarelli et al. 2014; Dwyer & Uman 2014; Aizawa et al. 2016). On Earth, these discharges support a global electric circuit and have the ability to modulate chemical and physical reactions in the atmosphere (Price 1993; Rakov & Uman 2007; Siingh et al. 2007; Genareau et al. 2015; Wadsworth et al. 2017; Mueller et al. 2018). Indeed, lightning may have been involved in the production of prebiotic molecules in an early-Earth environment (Miller & Urey 1959; Navarro-Gonz\u00e1lez et al. 1998) and such lightning has been hypothesized to have been associated within dusty flows (namely, volcanic plumes) rather than hydrometeor clouds (Navarro-Gonz\u00e1lez et al. 1998; Segura & Navarro-Gonz\u00e1lez 2005; Johnson et al. 2008). If the mineral clouds inferred to exist on extrasolar worlds can be considered analogs to the dusty environments in our own solar systems, charging within these systems may also stimulate a wide array of electrostatic phenomena and help catalyze prebiotic chemistry (Hodos\u00e1n et al. 2016).","Citation Text":["Helling & Fomins 2013"],"Functions Text":["As happens within Earth\u2019s clouds, exoplanet clouds are likely gravitationally stratified, meaning that smaller, lighter grains become concentrated at the top of the clouds, while larger, heavier grains remain at lower elevation"],"Functions Label":["Uses"],"Citation Start End":[[631,652]],"Functions Start End":[[353,580]]}
{"Identifier":"2022AandA...663A...5M__Andrae_et_al._2018_Instance_1","Paragraph":"With the real-time pipeline, we also frequently detect eclipsing binaries, due to their significant changes in brightness. Often, when we trigger these sources, PV light curves reveal potentially eclipsing behaviour. In these cases, we can attempt to verify the nature of the variable source by creating Lomb-Scargle periodograms for the light curves38 and inspecting, by eye, the folded light curves that we create using the best periods found. We note that this is done independently of the rest of pipeline; it is not part of the automatic software. An example of an eclipsing binary we found using TUVOpipe is shown in Fig. 32. The strong variability (by over 1 optical magnitude) led us to create a long-term light curve; the PV light curves of all the used UVOT filters are shown in Fig. 33. The source is a known source that appears in many sky survey catalogues, including Gaia, where its G-band magnitude is listed as 14.1 and its inferred temperature is ~5000 K (see Andrae et al. 2018 for how the listed temperature we extract from the Gaia catalogue is inferred from the Gaia data). This G-band magnitude is brighter than the U-UV bands seen in the UVOT light curve, indicating (along with the temperature) that it is a relatively red source. The source is also present in the ATLAS catalogue of variable stars Heinze et al. (2018), where it is denoted as AT J098.1962+05.6837. The ATLAS catalogue gives a period of 1.985 days. However, no classification of the source was so far known, so the cause of the found periodicity is not clear. Our method gave a best period of 1.9899 days, so we used this period to fold the UVOT U-band light curve, which is shown in Fig. 34. The light curve clearly reveals an eclipse in which the source becomes dimmer by around 1.5 magnitudes, and the shape of the light curve confirms it as an eclipsing binary. There may also be additional very weak minima seen halfway between the deep eclipses (see around phase \u22120.5 and 0.5 in the plot shown). The fact that any potential secondary minima are very weak may suggest that this is an Algol-type system (these are known as EA binaries; see Carmo et al. 2020 for a review of these systems and example light curves; the inferred Gaia temperature is also consistent with this being an EA system).","Citation Text":["Andrae et al. 2018"],"Functions Text":["The source is a known source that appears in many sky survey catalogues, including Gaia, where its G-band magnitude is listed as 14.1 and its inferred temperature is ~5000 K (see","for how the listed temperature we extract from the Gaia catalogue is inferred from the Gaia data)."],"Functions Label":["Uses","Uses"],"Citation Start End":[[977,995]],"Functions Start End":[[798,976],[996,1094]]}
{"Identifier":"2015ApJ...815...39W__Frank_et_al._1996_Instance_1","Paragraph":"With improved resolution, in this work we have resolved corrugated swept-up shells as well as small blobs of mixed material in the wind region that were not seen in Paper I. The swept-up shells are much smoother in the simulations including poloidal field, which is related to the stabilization of shear instability by the magnetic field. To check how an even higher resolution can affect the results, we have run simulations for case b2 at a resolution of 3200 \u00d7 990. In the upper panel of Figure 9 we present the density map at t = 1000 year for the run without ambient poloidal field. A close comparison to its ordinary resolution counterpart (see the upper left panel of Figure 5) shows that the details of the mixing structures are quite different in the two runs. More complex and fragmented mixing structures are generated in the higher resolution simulation, which indicates that the process producing these structures is not completely resolved at the current resolution (see, e.g., Frank et al. 1996). Despite the different fine structures, the overall outflow shape and the jet structure in the two simulations are in good agreement. In the lower panel of Figure 9 we present the same maps for the run with ambient poloidal field. The smoother and thicker shell seen here is almost identical to its ordinary resolution counterpart (see the lower left panel of Figure 5), and similar feather-like structures are also present near the wind boundary. There is also no sign of intense stochastic behavior like that seen in the upper panel of Figure 9. This suggests that the stabilization of the wind-ambient interface by the poloidal field is quite robust, and this conclusion holds well at least for the resolution we have used. Finally, it should be noted that the wind in the current simulations fans out in all radial directions even at the equator where an accreting inner envelope or disk is instead expected. Whether this setup could have anything to do with the growth of stochastic mixing structures near the inner toroid will need to be resolved in future simulations with more realistic setup of boundary conditions.","Citation Text":["Frank et al. 1996"],"Functions Text":["More complex and fragmented mixing structures are generated in the higher resolution simulation, which indicates that the process producing these structures is not completely resolved at the current resolution (see, e.g.,"],"Functions Label":["Uses"],"Citation Start End":[[992,1009]],"Functions Start End":[[770,991]]}
{"Identifier":"2021ApJ...908...95H__Sanders_&_Mirabel_1996_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":["Sanders & Mirabel 1996"],"Functions Text":["These often-called starburst galaxies, with an IR luminosity LIR \u223c (0.1\u20135) \u00d7 1012 L\u2299 (e.g.,","become increasingly more common at high z."],"Functions Label":["Background","Background"],"Citation Start End":[[825,847]],"Functions Start End":[[733,824],[873,915]]}
{"Identifier":"2020ApJ...895..128M__Zaldarriaga_et_al._2018_Instance_1","Paragraph":"We analyze the 10 BBH mergers reported by LIGO and Virgo in their O1 and O2 observing runs (Abbott et al. 2019a; LIGO Scientific Collaboration & Virgo Collaboration 2019). Before discussing results, it is useful to review expectations from the literature for the spin distributions resulting from different formation scenarios. Isolated binary evolution is predicted to yield black holes with spins preferentially aligned with their orbit. Although spin misalignments may be introduced by natal supernova kicks, episodes of mass transfer and tidal torques serve to realign component spins before the formation of the final black hole binary (Rodriguez et al. 2016; Zevin et al. 2017; Gerosa et al. 2018; Qin et al. 2018; Zaldarriaga et al. 2018; Bavera et al. 2020). The black holes\u2019 spin magnitudes in this scenario are much more uncertain. Recent work indicates that angular momentum is efficiently transported away from stellar cores, leaving black holes with natal spins as low as a \u223c 10\u22122 (Qin et al. 2018; Fuller & Ma 2019). While tides on the progenitor of the second-born black hole can spin up the progenitor star (Zaldarriaga et al. 2018), this effect can be counteracted by mass loss in stellar winds, and more detailed simulations find only low or moderate spin increases due to tides (Qin et al. 2018; Bavera et al. 2020). Meanwhile, dynamically formed systems in dense stellar clusters have no a priori preferred axis, and so are likely to have random spin configurations (Rodriguez et al. 2016, 2018, 2019; Doctor et al. 2020). Once again, however, the expected spin magnitudes are largely unknown, subject to the same uncertainties mentioned above regarding natal black hole spins. One firm prediction of the dynamical scenario concerns the spins of second-generation binaries, whose components were themselves formed from previous mergers. Regardless of their component spins, black hole mergers generally yield remnants with a \u223c 0.7; thus the effective spin of two such second-generation binaries may be large (Fishbach et al. 2017; Gerosa & Berti 2017; Rodriguez et al. 2018, 2019; Doctor et al. 2020).","Citation Text":["Zaldarriaga et al. 2018"],"Functions Text":["Although spin misalignments may be introduced by natal supernova kicks, episodes of mass transfer and tidal torques serve to realign component spins before the formation of the final black hole binary"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[721,744]],"Functions Start End":[[440,640]]}
{"Identifier":"2019MNRAS.483.2362R__Calderone_et_al._2013_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":["Calderone et al. 2013"],"Functions Text":["However, from spectro-polarimetric observations of a \u03b3-ray emitting NLSy1 galaxy, PKS 2004 \u2212 447","and accretion disc modelling of a sample of 23 radio-detected NLSy1 galaxies","indicate that they have masses similar to the blazar class of AGNs."],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[813,834]],"Functions Start End":[[618,714],[735,811],[836,903]]}
{"Identifier":"2021MNRAS.507.6012Z__Kendrick_2018_Instance_2","Paragraph":"Being a benchmark system H + H2, H + HD, and their isotopic counterparts have received much attention over the last several decades (Marinero et al. 1984; Zhang & Miller 1989; D\u2019Mello et al. 1991; Harich et al. 2002; Gao et al. 2015; Yuan et al. 2018a, b, 2020). Most early experimental and theoretical investigations were centered around benchmarking theory against experiments and providing improved descriptions of the H3 potential energy surfaces (PES; Boothroyd et al. 1996; Mielke, Garrett & Peterson 2002; Yuan et al. 2018a, b). Among the available PESs for the H3 system, the one by Boothroyd et al. (1996) referred to as the BKMP2 PES and by Mielke et al. (2002) referred to as the CCI PES, nearly equally well account for most experimental data for H + H2, H + HD, and D + HD collisions. These PESs have also been able to account for even subtle effects such as the GP (Kendrick et al. 2015; Hazra et al. 2016; Croft et al. 2017; Kendrick 2018, 2019). The GP effect, while not significant at temperatures relevant to astrophysics, is important below 1 K as illustrated in a series of calculations on H + HD (v, j) collisions for vibrational levels v = 4 \u2212 9 (Kendrick et al. 2015; Hazra et al. 2016; Croft et al. 2017; Kendrick 2018, 2019). Though several prior studies of Flower and co-workers  (Flower 1999, 2000; Flower & Roueff 1999; Wrathmall et al. 2007) have reported rate coefficients for H + HD collisions, due to the approximations involved in the scattering calculations (e.g. neglect of hydrogen atom-exchange), the reliability of the available rate coefficients has been a source of debate (Desrousseaux et al. 2018). Recently, Desrousseaux et al. (2018) reported rate coefficients for pure rotational transitions for j \u2264 10 within the v = 0 vibrational level using accurate quantum calculations that include the exchange channel. In this paper, we report rate coefficients for state-to-state rovibrational transitions in HD induced by H atoms between and within the v = 0 and 1 vibrational levels and for temperatures ranging from T = 1\u20131000 K.","Citation Text":["Kendrick 2018"],"Functions Text":["The GP effect, while not significant at temperatures relevant to astrophysics, is important below 1 K as illustrated in a series of calculations on H + HD (v, j) collisions for vibrational levels v = 4 \u2212 9"],"Functions Label":["Motivation"],"Citation Start End":[[1229,1242]],"Functions Start End":[[962,1167]]}
{"Identifier":"2015MNRAS.453.2747M__Shaw_et_al._2008_Instance_1","Paragraph":"\u0393CR can be expressed as the product of three factors \u2013 the total cosmic ray ionization rate per H nucleus \u03b6H (including both primary and secondary ionizations), the average energy deposited into the medium per ionization \u0394Q, and nH. Both \u03b6H and \u0394Q are rather uncertain and can vary considerably over different Galactic environments. Typical values of \u03b6H in dense gas are \u223c1\u20135 \u00d7 10\u221217 s\u22121 (Dalgarno 2006), but there is evidence from H3+ observations that \u03b6H is considerably higher in the diffuse gas under consideration here (Dalgarno 2006; Indriolo & McCall 2012). Indriolo & McCall (2012) find a mean \u03b6H of 1.8 \u00d7 10\u221216 s\u22121 in their sample of diffuse molecular sight lines, and values as large as \u223c1 \u00d7 10\u221215 s\u22121 have been reported in the literature (Snow & McCall 2006; Shaw et al. 2008). In this paper, we adopt the value 1.8 \u00d7 10\u221216 s\u22121. For \u0394Q, we use 10 eV, as estimated for diffuse molecular gas from table 6 of Glassgold, Galli & Padovani (2012), although it is important to note that this value can vary by several eV depending on the precise physical and chemical conditions in the cloud. Combining these factors, the cosmic ray heating rate is\n\n(11)\n\n\\begin{eqnarray}\n\\Gamma _{\\rm {CR}} &amp;=&amp; \\zeta \\Delta Q n_{\\rm {H}} \\nonumber \\\\\n&amp;\\approx&amp; 1.9 \\times 10^{-25} \\left( \\frac{n_{\\rm {H}}}{30 \\hspace{3.0pt} \\rm {cm}^{-3}} \\right) \\hspace{3.0pt} \\rm {ergs} \\hspace{3.0pt} \\rm {cm}^{-3} \\hspace{3.0pt} \\rm {s}^{-1}.\n\\end{eqnarray}\n\nFor \u0393PE, we adopt the expression:\n\n(12)\n\n\\begin{equation}\n\\Gamma _{\\rm {PE}} = 1.3 \\times 10^{-24} \\hspace{3.0pt} n_{\\rm {H}} \\epsilon G_0 \\hspace{3.0pt} \\rm {ergs} \\hspace{3.0pt} \\rm {cm}^{-3} \\hspace{3.0pt} \\rm {s}^{-1}\n\\end{equation}\n\nfrom Wolfire et al. (2003), where G0 is the intensity of FUV light in units of the Habing (1968) field and \u03f5 is the heating efficiency factor given by equation 20 of Wolfire et al. (2003). For nH = 30\u2009cm\u22123, T = 100\u2009K, an electron fraction of 1.6 \u00d7 10\u22124, and a FUV field of G0 = 1.1 (Mathis, Mezger & Panagia 1983), \u03f5 evaluates to 1.8 \u00d7 10\u22122, yielding\n\n(13)\n\n\\begin{equation}\n\\Gamma _{\\rm {PE}} = 7.6 \\times 10^{-25} \\left( \\frac{n_{\\rm {H}}}{30 \\hspace{3.0pt} \\rm {cm}^{-3}} \\right) \\hspace{3.0pt} \\rm {ergs} \\hspace{3.0pt} \\rm {cm}^{-3} \\hspace{3.0pt} \\rm {s}^{-1}.\n\\end{equation}\n\n","Citation Text":["Shaw et al. 2008"],"Functions Text":["Indriolo & McCall (2012) find a mean \u03b6H of 1.8 \u00d7 10\u221216 s\u22121 in their sample of diffuse molecular sight lines, and values as large as \u223c1 \u00d7 10\u221215 s\u22121 have been reported in the literature"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[770,786]],"Functions Start End":[[565,748]]}
{"Identifier":"2022AandA...658A.188S__Liu_et_al._(2013)_Instance_1","Paragraph":"The trends between the LF slope \u03b1 and the aforementioned parameters, with the addition of the morphological T type, are shown in Fig. 4 together with their Spearman correlation coefficient (\u03c1) and their p value, indicating the probability that the two sets of data are uncorrelated. We summarize the properties for which we look for a correlation in Table 4. We define a correlation to be negligible when |\u03c1|=[0\u2005\u2212\u20050.2], weak when |\u03c1|=[0.2\u2005\u2212\u20050.4], moderate when |\u03c1|=[0.4\u2005\u2212\u20050.6], strong when |\u03c1|=[0.6\u2005\u2212\u20050.8], and very strong when |\u03c1|=[0.8\u2005\u2212\u20051]; using the p value to evaluate the probability that, despite showing a correlation, two variables may be uncorrelated. It should be noted that only a handful of studies so far have looked at the correlation between \u03b1 and global galaxy properties: Kennicutt et al. (1989), Elmegreen & Salzer (1999), Youngblood & Hunter (1999), van Zee (2000), and Thilker et al. (2002) investigated nebular LFs as in this paper, while Liu et al. (2013) identified H\u202fII regions via Pa\u03b1, and Cook et al. (2016) studied the GALEX far-ultraviolet (FUV) LFs of H\u202fII regions. While the sample of Cook et al. (2016) includes a few hundred galaxies, the other studies are based on samples ranging from 10 to 35 galaxies, similar to our study. In this section and in Sect. 6.1, we compare our results to those studies that, as in our case, applied a uniform analysis methodology on galaxy samples. It should be noted that using different tracers means probing different source ages and, as reported by Oey & Clarke (1998), older H\u202fII regions tend to have steeper LF slopes, mainly due to the short main-sequence lifetimes of the more massive stars constituting the brighter H\u202fII regions. This is the reason why, for example, FUV observations, probing H\u202fII regions with ages less than 100 Myr, are expected to deliver a steeper LF compared to H\u03b1 observations, typically probing H\u202fII regions younger than 10 Myr, and our comparison remains qualitative.","Citation Text":["Liu et al. (2013)"],"Functions Text":["It should be noted that only a handful of studies so far have looked at the correlation between \u03b1 and global galaxy properties:","while","identified H\u202fII regions via Pa\u03b1"],"Functions Label":["Background","Background","Background"],"Citation Start End":[[960,977]],"Functions Start End":[[661,788],[954,959],[978,1009]]}
{"Identifier":"2019MNRAS.490.5478W__Winter_et_al._2018b_Instance_1","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"],"Functions Text":["However, the stellar number densities required for tidal truncation are high, and in practice few observed regions satisfy this condition"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[735,754]],"Functions Start End":[[596,733]]}
{"Identifier":"2018MNRAS.478.4357S__Barr_&_Hochberg_1988_Instance_1","Paragraph":"Since long cosmologists have felt motivated to look for alternative explanations for the DE beyond a rigid cosmological constant \u039b. The scalar field paradigm was then profusely used also to make the cosmic vacuum dynamical: \u039b = \u039b(\u03d5). In the old days, the main aim was to adjust the large value of \u039b typically predicted in QFT to be zero. There were many early proposals (see e.g. Endo & Fukui 1977, 1982; Fujii 1982; Dolgov 1983; Abbott 1985; Zee 1985; Barr 1987; Ford 1987; Peccei, Sol\u00e0 & Wetterich 1987; Weiss 1987; Barr & Hochberg 1988). In spite of the hopes raised by these works at solving the \u2018old CC Problem\u2019, it was later shown in Weinberg (1989) through the so-called no-go theorem that most if not all the dynamical adjustment mechanisms existing in the literature to date were plagued by more or less obvious forms of subtly hidden fine tuning. For this reason, the subsequent use of scalar fields in cosmology was mostly focused on trying to explain another aspect of the CCP: the cosmic coincidence problem (viz. the fact that $\\rho _\\Lambda$ happens to be so close to the matter density \u03c1m right now; see e.g. Peebles & Ratra 2003). The new wave of dynamical scalar fields in cosmology crystalized in the notions of quintessence, phantom fields and the like, which have had a tremendous influence in cosmology till our days (see e.g. Peebles & Ratra 1988; Ratra & Peebles 1988; Wetterich 1988; Wetterich 1995; Caldwell, Dave & Steinhardt 1998; Zlatev, Wang & Steinhardt 1999; Amendola 2000; Caldwell, Kamionkowski & Weinberg 2003), the reviews (Sahni & Starobinsky 2000; Padmanabhan 2003; Peebles & Ratra 2003; Copeland, Sami & Tsujikawa 2006), and the many references therein. At the same time a blooming crest of models based on ascribing a direct phenomenological time-dependence to the CC term, \u039b = \u039b(t), broke with impetus into the market. For an account of some of the old attempts, see Overduin & Cooperstock (1998, and references therein).","Citation Text":["Barr & Hochberg 1988"],"Functions Text":["Since long cosmologists have felt motivated to look for alternative explanations for the DE beyond a rigid cosmological constant \u039b. The scalar field paradigm was then profusely used also to make the cosmic vacuum dynamical: \u039b = \u039b(\u03d5). In the old days, the main aim was to adjust the large value of \u039b typically predicted in QFT to be zero. There were many early proposals (see e.g."],"Functions Label":["Background"],"Citation Start End":[[518,538]],"Functions Start End":[[0,379]]}
{"Identifier":"2016ApJ...817..173L__Takahashi_2004_Instance_1","Paragraph":"Horizon scale imaging promises to test basic predictions of GR and improves our understanding of the physics responsible for accretion and emission in a strong gravitational field. In particular, imaging a black hole shadow has been a long-standing goal of black hole astronomy. However, imaging the black hole shadow feature in Sgr A* has been inherently challenged by two known effects. First, the scattering by interstellar medium blurs the strong GR features near the black hole. In a recent work, it has been shown that this effect can be mitigated based on the fact that the scattering is well understood over the relative range of baseline lengths provided by the EHT (Fish et al. 2014). Second, while the predicted shadow feature is nearly independent of the spin or orientation of the black hole to within 10% (Bardeen 1973; Takahashi 2004), the emission region surrounding the black hole depends on the details of the underlying accretion process and is intrinsically time variable primarily due to the stochastic nature of magnetorotational-instability-driven turbulence and magnetic reconnection in the accretion flow. Magnetorotational instability (MRI, Balbus & Hawley 1991, 1998) is believed to be the leading mechanism driving turbulence in accretion disks and develops on orbital timescales. The timescale for the Keplerian motion at the innermost stable circular orbit around the black hole in Sgr A* ranges from 30 minutes for a non-rotating black hole to 4 minutes for prograde orbits around a maximally rotating black hole (Doeleman et al. 2009b). These timescales are much less than the typical duration of a Very Long Baseline Interferometry (VLBI) experiment, which violates one of the basic requirements for VLBI Earth-rotation aperture synthesis imaging. In contrast, the corresponding timescales in the nearby giant elliptical galaxy M87, which has the second largest apparent event horizon, are much larger (a minimal timescale of a few days).","Citation Text":["Takahashi 2004"],"Functions Text":["Second, while the predicted shadow feature is nearly independent of the spin or orientation of the black hole to within 10%","the emission region surrounding the black hole depends on the details of the underlying accretion process and is intrinsically time variable primarily due to the stochastic nature of magnetorotational-instability-driven turbulence and magnetic reconnection in the accretion flow."],"Functions Label":["Background","Background"],"Citation Start End":[[834,848]],"Functions Start End":[[695,818],[851,1130]]}
{"Identifier":"2015ApJ...801..112L__Kulsrud_1983_Instance_1","Paragraph":"The main goal of this paper was to further expand the recent transport theory of Zank et\u00c2 al. (2014) and thus, by default, earlier attempts by Drake et\u00c2 al. (2006, 2013) and Bian & Kontar (2013). In our approach, similar to Zank et\u00c2 al. (2014), the focus was on inertial-scale flux ropes that were modeled as quasi-2D magnetic islands superposed transversely on a strong large-scale magnetic field imbedded in the large-scale plasma flow, as suggested by the 3D MHD turbulence simulations with a strong guide field of Dmitruk et\u00c2 al. (2004). We adopt a different perspective in proceeding from the basic guiding center kinetic equation (e.g., Kulsrud 1983) instead of the transformed Vlasov equation (Skilling 1975) used by Zank et\u00c2 al. (2014). This offers the advantage of explicitly identifying drift acceleration for various guiding center drift mechanisms and the betatron acceleration that energetic particles would be subject to in both the large-scale plasma flow and magnetic field, and in the inertial-scale plasma flow and magnetic fields of contracting and merging flux ropes. Our approach allows us to distinguish between particle drift and betatron energization when magnetic flux ropes behave in an incompressible fashion (later stage contraction or during merging) or when they are compressible (early stage contraction and during island collisions). The transformation of the guiding center kinetic equation for nearly gyrotropic particle distributions into a focused transport equation clarifies the role of drift and betatron acceleration in focused transport. In the process, a more general focused transport formalism of particle transport and momentum change in contracting and merging flux ropes was derived, whereas Zank et\u00c2 al. (2014) followed a more targeted focused transport approach based on specific conservation laws (magnetic moment conservation, conservation of parallel action, and conservation of magnetic flux) that are thought to apply to the above-mentioned three acceleration mechanisms (see also Drake et\u00c2 al. 2006, 2013).","Citation Text":["Kulsrud 1983"],"Functions Text":["We adopt a different perspective in proceeding from the basic guiding center kinetic equation (e.g.,","instead of the transformed Vlasov equation (Skilling 1975) used by Zank et\u00c2 al. (2014)."],"Functions Label":["Uses","Compare\/Contrast"],"Citation Start End":[[643,655]],"Functions Start End":[[542,642],[657,744]]}
{"Identifier":"2021MNRAS.508.4512L__Fishbach_et_al._2019_Instance_1","Paragraph":"An EM counterpart is not strictly necessary to use compact binary mergers such as BBHs and BNSs as standard sirens. By matching the sky localization region of GW sources \u2013 which can be inferred from the GW measurements \u2013 with galaxy catalogues, one might in fact be able to extract complementary information on the redshift of the sources, without the need of spotting an EM counterpart. The idea was originally proposed by Schutz (1986) and it has subsequently been used and developed in different analyses (Holz & Hughes 2005; MacLeod & Hogan 2008; Petiteau, Babak & Sesana 2011; Del Pozzo 2012; Chen et al. 2018; Gray et al. 2020). It has already been tested with real data collected by the LIGO and Virgo detectors (Fishbach et al. 2019; Soares-Santos et al. 2019; Palmese et al. 2020; Abbott et al. 2020d; Finke et al. 2021), though the constraints obtained so far with this \u2018statistical\u2019 method are not competitive with the ones derived from GW170817 and its EM counterpart, mainly because of the poor spatial resolution of the current network of ground-based interferometers. Future observations, taken with an enlarged network of ground-based GW detectors, will allow for better cosmological measurements (Chen et al. 2018), mainly thanks to the improved sky localization accuracy. Other complementary methods, which analogously do not require the identification of an EM counterpart, might yield interesting results as well (Taylor & Gair 2012; Oguri 2016; Del Pozzo, Li & Messenger 2017; Mukherjee & Wandelt 2018; Farr et al. 2019; Mukherjee et al. 2020a; Mukherjee, Wandelt & Silk 2020b, c, d; Ezquiaga & Holz 2021). The era of precise cosmological measurements with GWs will however start only with next-generation interferometers, such as the Einstein Telescope (ET) (Punturo et al. 2010; Sathyaprakash, Schutz & Van Den Broeck 2010; Belgacem et al. 2019a; Maggiore et al. 2020) and the Cosmic Explorer (Abbott et al. 2017a; Reitze et al. 2019a, b) on the Earth, or TianQin (Mei et al. 2020), Taiji (Luo et al. 2020), and the Laser Interferometer Space Antenna (LISA) (Amaro-Seoane et al. 2017) in space. The latter instrument is the focus of the present investigation. In what follows, we will briefly introduce LISA and review previous studies of LISA\u2019s capability to do cosmological analyses using standard sirens. More details on how to extract cosmology from GWs by statistically matching with galaxy catalogues will be given in Section 2.","Citation Text":["Fishbach et al. 2019"],"Functions Text":["It has already been tested with real data collected by the LIGO and Virgo detectors","though the constraints obtained so far with this \u2018statistical\u2019 method are not competitive with the ones derived from GW170817 and its EM counterpart, mainly because of the poor spatial resolution of the current network of ground-based interferometers."],"Functions Label":["Background","Differences"],"Citation Start End":[[720,740]],"Functions Start End":[[635,718],[831,1082]]}
{"Identifier":"2019MNRAS.487.3776P__Moriondo,_Giovanardi_&_Hunt_1998_Instance_1","Paragraph":"Driven by the influx of spatially resolved observations coming from integral field units (IFU), more recent investigations have focused on attempting to infer structural, dynamical, and\/or chemical properties for localized regions of galaxies, by decomposing them into physically motivated components. In fact, this concept pre-dates the large-scale use of IFU with works dealing with, for instance, photometric disc\/bulge decompositions of surface brightness profiles (e.g. see Kent 1985; Cinzano & van der Marel 1993; Scorza & Bender 1995; Moriondo, Giovanardi & Hunt 1998; Krajnovi\u0107 et al. 2013), including using multiple filters (Dimauro et al. 2018). Aside from surface brightness, radial profiles of other parameters have also been subject to analogous decompositions. For instance, the decomposition of gas (usually H\u2009i) circular velocity profiles into contributions from different galaxy sub-components is a well-established practise (as in van Albada et al. 1985; Carignan, Sancisi & van Albada 1988; Battaglia et al. 2006; Noordermeer et al. 2007; Swaters et al. 2012; Aniyan et al. 2016, 2018; Sofue 2017), whilst decompositions of mass profiles have attempted to infer the contributions from dark matter (DM) and baryons (stars and globular clusters, gas, et cetera; Annunziatella et al. 2017; Poci, Cappellari & McDermid 2017; Bellstedt et al. 2018). These concepts have been extended to two dimensions, including multiband photometric disc\/bulge decompositions of images, rather than profiles (for instance, see Scorza et al. 1998; de Souza, Gadotti & dos Anjos 2004; Norris, Sharples & Kuntschner 2006; Simard et al. 2011; M\u00e9ndez-Abreu et al. 2017; Dalla Bont\u00e0 et al. 2018). Moreover, there have been recent efforts to conduct the decomposition directly on an observed spectrum (Johnston et al. 2012; Coccato et al. 2015; Tabor et al. 2017; Coccato et al. 2018), to similarly determine the contributions to various spectral features coming from \u2018distinct\u2019 galaxy sub-components. This type of component-based approach attempts to isolate the distinct contributions to observed galaxies from regions that may or may not have had different origins and\/or formation paths; however, they have thus far dealt with the problem from only one perspective \u2013 dynamics or stellar populations.","Citation Text":["Moriondo, Giovanardi & Hunt 1998"],"Functions Text":["In fact, this concept pre-dates the large-scale use of IFU with works dealing with, for instance, photometric disc\/bulge decompositions of surface brightness profiles (e.g. see"],"Functions Label":["Background"],"Citation Start End":[[542,574]],"Functions Start End":[[302,478]]}
{"Identifier":"2017AandA...601A.109W__Hoyt_et_al._1994_Instance_1","Paragraph":"The WSN\/ISN series is based on the counts of both sunspot groups and individual sunspots, with the former being weighted with a factor of ten: (1)\\begin{equation} R = k\\cdot(10\\cdot G + S), \\end{equation}R=k\u00b7(10\u00b7G+S),where G and S are the numbers of sunspot groups and individual sunspots, respectively, and k is a correction factor, characterizing each observer. However, resolving individual spots may be imprecise with poor instrumentations, and a new series, based only on sunspot groups, was proposed, called the group sunspot number, GSN (Hoyt et al. 1994; Hoyt & Schatten 1998). The GSN is more robust than WSN regarding observational conditions (e.g., Usoskin 2017). There is still a potential problem related to the grouping of individual spots, which may have been considered by earlier observers in a different manner to that currently accepted (Clette et al. 2014). This uncertainty is related to both WSN\/ISN and GSN but can be fixed by redefining groups in historical sunspot drawings (Arlt et al. 2013). The GSN series produced by Hoyt & Schatten (1998) also uses the linear scaling and daisy-chaining method to reduce different data to the same reference observer, for which the RGO was chosen. The GSN is constantly scaled up by a factor 12.08 to make it comparable with the WSN series. The main advantage of the GSN series is that Hoyt & Schatten (1998) had collected and published the original database of raw data, including all the records of individual observers. This makes it possible to revise the entire series if needed. Since some corrections and additions have been recently made to this dataset, a revised database of the sunspot group numbers, separate for each observer, is published (Vaquero et al. 2016, referred to as V16 hereafter). The GSN series was revised by Svalgaard & Schatten (2016) who performed a full re-calibration of the observers using a modified daisy-chaining method with a reduced number of links (the \u201cbackbone\u201d method). The revised backbone GSN series suggests that the level of solar activity was relatively high in the 18th and 19th centuries, much higher than that implied by the original GSN series by Hoyt & Schatten (1998) and by WSN. ","Citation Text":["Hoyt et al. 1994"],"Functions Text":["However, resolving individual spots may be imprecise with poor instrumentations, and a new series, based only on sunspot groups, was proposed, called the group sunspot number, GSN"],"Functions Label":["Background"],"Citation Start End":[[545,561]],"Functions Start End":[[364,543]]}
{"Identifier":"2015ApJ...800...62B___2014b_Instance_1","Paragraph":"We find that, even though the physical parameters such as kTe and \u00cf\u0084e are very well constrained by the data, it is still impossible to formally distinguish the geometry. The slab (disk-like) and the spherical geometries, as parameterized by the compTT and compPS models used here, both describe the MCG\u00e2\u0080\u009305-23-016 spectrum equally well. We note that a similar result was found in observations of the Seyfert 1.2 IC\u00c2 4329a and the narrow-line Seyfert 1 SWIFT J2127.4+5654 with NuSTAR (Brenneman et\u00c2 al. 2014a, 2014b; Marinucci et\u00c2 al. 2014). Both of these AGNs and MCG\u00e2\u0080\u009305-23-016 are radio-quiet, however, they differ in other properties. With a mass of the super-massive black hole of \u00e2\u0088\u00bc5 \u00c3\u0097 107\u00c2 M (Wandel & Mushotzky 1986), the mean intrinsic 2\u00e2\u0080\u009310\u00e2\u0080\u0089keV luminosity of 1.66 \u00c3\u0097 1043 erg\u00c2 s\u00e2\u0088\u00921 (see Table\u00c2 1) and a bolometric correction from Marconi et\u00c2 al. (2004), MCG\u00e2\u0080\u009305-23-016 is accreting at approximately 5% of the Eddington rate. This is almost an order of magnitude less than the key other two AGNs. Interestingly, SWIFT J2127.4+5654 has the lowest black hole mass and the lowest cut-off, followed by MCG\u00e2\u0080\u009305-23-016 in the middle, and IC\u00c2 4329a with highest mass and cut-off energy. In a number of other AGNs, a stringent lower limit on the cut-off energy was placed using the NuSTAR data, indicating a generally higher coronal temperature and lower optical depths, e.g., Ecut > 190\u00e2\u0080\u0089keV in 3C 382 (Ballantyne et\u00c2 al. 2014) and in Ark 120 (Matt et\u00c2 al. 2014), and Ecut > 210\u00e2\u0080\u0089keV in NGC\u00c2 2110 (Marinucci et\u00c2 al. 2015). Using long-term averaged data from INTEGRAL, Malizia et\u00c2 al. (2014) constrained cut-off energies for 26 AGNs in the range between 50 and 200\u00e2\u0080\u0089keV, some of which have been or will be observed with NuSTAR. With more high-quality measurements in the near future, covering a wide range of physical properties, it will be possible to directly probe the physics of the AGN corona. In order to distinguish the fine differences due to the coronal geometry, longer observations of sources with a weaker reflection continuum will be needed.","Citation Text":["Brenneman et\u00c2 al.","2014b"],"Functions Text":["We note that a similar result was found in observations of the Seyfert 1.2 IC\u00c2 4329a and the narrow-line Seyfert 1 SWIFT J2127.4+5654 with NuSTAR"],"Functions Label":["Similarities"],"Citation Start End":[[485,502],[510,515]],"Functions Start End":[[338,483]]}
{"Identifier":"2021AandA...656A..15G___1974_Instance_1","Paragraph":"The comparison between Ulysses and Pamela overlapping measurements revealed that the proton flux in the rigidity interval 1.6\u20131.8 GV (0.92\u20131.09 GeV, corresponding approximately to the median energy of the GCR spectrum at solar minimum) has a radial intensity variation of 2.7 \u00b1 0.2% AU\u22121, and a latitudinal gradient of \u22120.024 \u00b1 0.005% degree\u22121 (De Simone et al. 2011). Positive (negative) latitudinal gradients are observed during positive (negative) polarity periods. In addition, Experiment 6 (E6) on board Helios-A and Helios-B provided ion data from four to several hundreds of MeV n\u22121 (Winkler 1976; Marquardt & Heber 2019). The Helios-A and Helios-B S\/C were launched on December 10, 1974 and January 15, 1976 during a positive polarity epoch and were sent into ecliptic orbits of 190-day and 185-day periods around the Sun. The orbits perihelia were 0.3095 AU and 0.290 AU, respectively. The aphelia were approximately 1 AU. As a result, the Helios data are representative of the cosmic-ray bulk variations that are to be experienced by Solar Orbiter, which will also reach maximum distances from the Sun of about 1 AU. In the recent paper by Marquardt & Heber (2019), the Helios proton data radial gradients of the GCR flux were found to be 6.6 \u00b1 4% above 50 MeV and 2 \u00b1 2.5% between 250 and 700 MeV between 0.4 and 1 AU. These results are in agreement with those from Pamela\/Ulysses (within the statistical and systematic errors). In conclusion, variations in the GCR proton-dominated flux along the Solar Orbiter orbit are expected to be of a few % at most; consequently, it is plausible to assume that models for cosmic-ray modulation developed on the basis of observations gathered near Earth will also apply to Solar Orbiter. On the other hand, the Metis data will allow us to verify this assumption in the unexplored region of tens of degrees above the solar equator. Analogously, even though no SEP data were gathered up to present time, it is likely that a study of the evolution of SEP events near the Sun above the solar equator will be possible for the first time also with Metis, in addition to the dedicated instruments flown on board Solar Orbiter.","Citation Text":["Marquardt & Heber (2019)"],"Functions Text":["In the recent paper by","the Helios proton data radial gradients of the GCR flux were found to be 6.6 \u00b1 4% above 50 MeV and 2 \u00b1 2.5% between 250 and 700 MeV between 0.4 and 1 AU. These results are in agreement with those from Pamela\/Ulysses (within the statistical and systematic errors)."],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[1150,1174]],"Functions Start End":[[1127,1149],[1176,1439]]}
{"Identifier":"2019AandA...631A..35B__Bridges_et_al._(1996)_Instance_1","Paragraph":"The collision velocity dependence of the coefficient of restitution between particles was observed in experiments (Bridges et al. 1996; Higa et al. 1996) and is discussed in the literature (e.g., Ram\u00edrez et al. 1999; Zhang & Vu-Quoc 2002). However, the experiments by Hei\u00dfelmann et al. (2010), used in the present paper to support our assumption of a constant coefficient of restitution, do not see a variation of the coefficient of restitution between particles at low collision velocities (\u2264 1 cm s\u22121). This discrepancy in results might originate in the nature of the collisions studied in these different experiments: Bridges et al. (1996) and Higa et al. (1996) performed collisions of a particle with a flat surface, while Hei\u00dfelmann et al. (2010) observed particle-particle collisions in a free-floating environment. The latter is an experimental environment very similar to NanoRocks. In such inter-particle collisions in free-floating environments, other physical effects lead to a different behavior of the energy dissipation during collisions. In particular, the damping behavior of a large plate or surface is expected to differ from that of a same-sized particle, so that the velocity dependence of the coefficient of restitution might be an effect of the experimental setup in Bridges et al. (1996) and Higa et al. (1996). Colwell et al. (2016) and Brisset et al. (2018) studied collisions between a round cm-sized particle and a flat surface of fine grains. They also observed an increase of the coefficient of restitution with decreasing collision velocity. While the composition of the target surface was different than in Bridges et al. (1996) and Higa et al. (1996) (fine granular material vs. solid ice), the similar behavior of the coefficient of restitution supports the fact that particle-surface collisions are very different from particle-particle collisions, and coefficients of restitution are only velocity dependent for collisions with particles with very different sizes (a much larger particle can be approximated as a target surface).","Citation Text":["Bridges et al. 1996"],"Functions Text":["The collision velocity dependence of the coefficient of restitution between particles was observed in experiments"],"Functions Label":["Background"],"Citation Start End":[[115,134]],"Functions Start End":[[0,113]]}
{"Identifier":"2016ApJ...818...38S__Vanzella_et_al._2011_Instance_1","Paragraph":"As described in the introduction, spectroscopically confirmed Ly\u03b1 emitters at high redshift (and the nonconfirmations from follow-up campaigns) have proven very valuable for studying the early universe and the environment at the epoch of reionization (Pentericci et al. 2011, 2014; Caruana et al. 2012, 2014; Treu et al. 2012, 2013; Faisst et al. 2014; Tilvi et al. 2014). In particular, confirmed Ly\u03b1 emitters fix the redshift of the object, resulting in improved prediction power from fitting stellar population synthesis models to the photometry. Assuming a set of stellar population models (e.g., Bruzual & Charlot 2003; Maraston 2005) to generate spectral energy distributions for galaxies at the emission-line redshift, and fitting them to the available photometry, can give estimates of physical quantities of the galaxies like total stellar mass (the normalization between the observed flux and best-fit model), the star formation rate, metallicity, and the age of the stellar populations, i.e., the galaxy (e.g., Labb\u00e9 et al. 2006; Vanzella et al. 2011; Coe et al. 2013, 2014; Finkelstein et al. 2013; Huang et al. 2015; Oesch et al. 2015; Zitrin et al. 2015b). When performing the spectral energy distribution fitting, assuming a dust law can furthermore predict the dust content of the galaxy. This can be directly compared to the measured UV spectral slope, if available from the data (e.g., Bouwens et al. 2015; Finkelstein et al. 2015b; Oesch et al. 2015). Another direct comparison can be obtained from independently determining the star formation rate from scaling relations with the UV photometry (Kennicutt 1998; Madau et al. 1998). As part of our study of IRAC-detected high-redshift galaxies presented by Huang et al. (2015), we estimated the physical properties of the confirmed Ly\u03b1 emitter presented in Section 7.1. An important aspect of this study was the availability of ancillary Spitzer photometry from SURFS-UP (Brada\u010d et al. 2014). Photometry in the rest-frame optical falling in the Spitzer IRAC infrared bands for high-redshift galaxies has proven to be an important part of reliably predicting the physical properties of (high-redshift) galaxies through spectral energy distribution fitting (e.g., Schaerer & de Barros 2010; Labb\u00e9 et al. 2013; Smit et al. 2014a, 2014b; Finkelstein et al. 2015a; Huang et al. 2015; Wilkins et al. 2015). Furthermore, fixing the redshift of the spectral energy distributions when fitting to photometry can also be used as a test of the validity of potential low-redshift contaminants. If, for instance, the best-fit low-redshift model predicts a dusty red and old stellar population, it would be very unlikely to see strong [O ii] emission, therefore making a high-redshift Ly\u03b1 scenario more likely (Coe et al. 2013; Finkelstein et al. 2013). Similar arguments can be used to rule out other line-emitting low-redshift contaminants. Fixing the redshift of high-redshift sources behind massive clusters, like the ones presented in the current study behind the GLASS clusters, is not only important for the study of individual sources and high-redshift galaxy populations. Knowing the redshift, i.e., the luminosity distance to any object, precisely, especially if it is multiply lensed, is also very valuable for lens modeling of the foreground clusters (e.g., Coe et al. 2013, 2014; Zitrin et al. 2014). Lastly, the sizes of high-redshift galaxies have also been shown to provide useful information about the environment and epoch they inhabit (Ono et al. 2013; Curtis-Lake et al. 2014; Holwerda et al. 2015).","Citation Text":["Vanzella et al. 2011"],"Functions Text":["Assuming a set of stellar population models","to generate spectral energy distributions for galaxies at the emission-line redshift, and fitting them to the available photometry, can give estimates of physical quantities of the galaxies like total stellar mass (the normalization between the observed flux and best-fit model), the star formation rate, metallicity, and the age of the stellar populations, i.e., the galaxy (e.g.,"],"Functions Label":["Uses","Uses"],"Citation Start End":[[1041,1061]],"Functions Start End":[[550,593],[640,1021]]}
{"Identifier":"2019MNRAS.485.3088C__Cheng_et_al._2018_Instance_1","Paragraph":"As displayed in Fig. 1 and Supplementary Movie 1, the control torque \u03c4, applied to the lander at rest, actuates it to pivot over the leading edge, and then the bottom side is dragged downwards in contact with the granular regolith, yielding a reaction force on to the lander (Allen et al. 2013). Given sufficient acceleration impulsive moment, the lander will leave the surface, hopping forwards in a ballistic trajectory. Hockman et al. (2017) stated that the key mechanism of this internally actuated mobility is the transmission of rotational energy to translation energy, indicating that the acceleration duration T of the control torque could determine the hopping outcomes. To elucidate the fundamental dynamics that governs surface locomotion on asteroids, we adopt 12 different acceleration durations ranging from 20 to 800 $\\rm {ms}$ in this section, in which \u03bc is set to 0.5, \u03b2 is set to 1.0, and c is set to 0 Pa. Two typical dynamics data, corresponding to 20 and 600 $\\rm {ms}$, are displayed in the insets of Figs 2(a) and (b), which shows instantaneous velocity $v$ and angle \u03b8 (relative to horizon), versus time t. Generally, during the hopping process, the acceleration first shows a gradual increase, and then decreases rapidly at the end of control torque, during which the reaction force is highly fluctuating due to intermittent transmissions of acoustic energy along force-chain-like networks (Clark, Kondic & Behringer 2012; Cheng et al. 2018), as illustrated in Fig 3. In this collisional scenario, the acceleration torque excites strong force chains which exert sporadic impacts on the bottom of the lander, while after the end of acceleration, the force networks become sparse and fragmented immediately, leading to the rapid decrease of the reaction force as evidenced by Fig. 2(c). The work done by the lander to impact with regolith surface generates the contact force, and then increase its velocity and height before the lander loses contact with the regolith at time tesc. After that, the lander behaves as free projectile motion under only gravity, which corresponds to a gradual decrease of the vertical velocity, forming a peak value in the velocity curve and a turning point in the angle curve. Hereafter, we define these two values as hopping velocity V and hopping angle \u03d1, respectively. Note that the lowercase $v$ and \u03b8 represent the instantaneous velocity and angle, while the uppercase V and \u03d1 represent hopping outcomes. The hopping distance can be deduced from V and \u03d1, given as 2V2cos\u2009(\u03d1)sin\u2009(\u03d1)\/g. Thus in the case of T = 600 ms, the lander hops about 57 m under the gravity of 1.0 \u00d7 10\u22124$\\rm {m\\, s^{-2}}$, corresponding to a large surface coverage on small asteroids. And in a lower acceleration duration like T = 40 ms, the lander only moves about 45 cm, which gives a much more precise form of surface relocation. The results presented here provide strong evidence on the feasibility of the hopping landers for the exploration of low-gravity bodies. Additionally, our simulations show the lander takes about 1\u201317 min to complete the hop, which is far less than the battery life of MASCOT (16 h, Ho et al. 2017). This information is crucial for the design of landers equipped without solar batteries.","Citation Text":["Cheng et al. 2018"],"Functions Text":["Generally, during the hopping process, the acceleration first shows a gradual increase, and then decreases rapidly at the end of control torque, during which the reaction force is highly fluctuating due to intermittent transmissions of acoustic energy along force-chain-like networks","as illustrated in Fig 3."],"Functions Label":["Background","Background"],"Citation Start End":[[1448,1465]],"Functions Start End":[[1131,1414],[1468,1492]]}
{"Identifier":"2022AandA...661A.129S__Rodr\u00edguez-Almeida_et_al._2021_Instance_2","Paragraph":"Radio astronomy is recognized as one of the most effective techniques to search for interstellar molecules. By comparing the spectra of candidate molecules in the laboratory with the spectra observed in astronomical surveys, we can determine whether these molecules exist in interstellar space. Therefore, it is necessary to provide rotational spectra of candidates for astronomical detection. Radio astronomy has helped to detect several sulfur-containing molecules in the ISM in recent years: in particular, thiols, the sulfur analogs of alcohols. Methanethiol (or methyl mercaptan, CH3SH) was detected in the Sagittarius B2 (Sgr B2) region of the center of our Galaxy (Linke et al. 1979; Gibb et al. 2000; M\u00fcller et al. 2016; Rodr\u00edguez-Almeida et al. 2021) and in the protostar IRAS 16293-2422 (Majumdar et al. 2016). Two groups reported to have detected several signs of ethanethiol (C2H5SH) in Sgr B2 (M\u00fcller et al. 2016) and Orion (Kolesnikov\u00e1 et al. 2014). Most recently, Rodriguez-Almeida reported the first unambiguous detection of ethanethiol in the ISM, toward the G+0.603-0.027 molecular cloud (Rodr\u00edguez-Almeida et al. 2021). Moreover, several sulfur-containing species have been observed in comets (Altwegg et al. 2017). Some recent efforts, both from spectroscopy and astronomical searches, to detect S-stitutes of other classes of compounds have also been reported. For instance, thioformic acid (HC(O)SH) was very recently detected in G+0.693\u20130.027. Its trans-isomer has an abundance of ~1 \u00d7 10\u201310 (Rodr\u00edguez-Almeida et al. 2021). Conversely, thioformamide (NH2CHS), the counterpart of for-mamide (NH2CHO), was characterized in the laboratory up to 660 GHz, and its transitions were searched for toward the hot cores Sgr B2(N1S) and Sgr B2(N2), but it was not detected (Motiyenko et al. 2020). The rotational spectrum of thioac-etamide was recently analyzed in the 59.6\u2013110.0 GHz frequency region (5.03\u20132.72 mm). Its emission was searched for in regions associated with star formation using the IRAM 30 m ASAI observations toward the prestellar core L1544 and the outflow shock L1157\u2013B1. The molecule was not detected, but the study allowed placing constraints on the thioacetamide abundances (Maris et al. 2019; Remijan et al. 2022).","Citation Text":["Rodr\u00edguez-Almeida et al. 2021"],"Functions Text":["Most recently, Rodriguez-Almeida reported the first unambiguous detection of ethanethiol in the ISM, toward the G+0.603-0.027 molecular cloud"],"Functions Label":["Background"],"Citation Start End":[[1107,1136]],"Functions Start End":[[964,1105]]}
{"Identifier":"2021AandA...655A..25Z__Garc\u00eda-Burillo_et_al._2014_Instance_1","Paragraph":"Outflows are ubiquitous in both luminous AGN and in local Seyfert galaxies, and occur on a wide range of physical scales, from highly ionised semi-relativistic winds and jets in the nuclear region at subparsec scales to galactic scale outflows seen in mildly ionised, molecular, and neutral gas (Morganti et al. 2016; Fiore et al. 2017; Fluetsch et al. 2019; Lutz et al. 2020; Veilleux et al. 2020, and references therein). In some cases molecular and ionised winds have similar velocities and are nearly co-spatial, suggesting a cooling sequence scenario where molecular gas forms from the cooling of the gas in the ionised wind (Richings & Faucher-Giguere 2017; Menci et al. 2019). Other AGN show ionised winds that are faster than the molecular winds, suggesting a different origin of the two phases (Veilleux et al. 2020, and references therein). The molecular phase is a crucial element of the feeding and feedback cycle of AGN because it constitutes the bulk of the total gas mass and it is the site of star formation activity. On galactic scales massive molecular winds are common in local Seyfert galaxies (e.g. Feruglio et al. 2010; Cicone et al. 2014; Dasyra et al. 2014; Morganti et al. 2015; Garc\u00eda-Burillo et al. 2014, 2017, 2019); these winds likely suppress star formation (i.e. negative feedback) as they reduce the molecular gas reservoir by heating or expelling gas from the host-galaxy ISM. In late-type AGN-host galaxies, the gas kinematics appears complex at all scales, showing several components such as bars, rings, and (warped) discs, with high velocity dispersion regions (e.g. Shimizu et al. 2019; Feruglio et al. 2020; Fern\u00e1ndez-Ontiveros et al. 2020; Alonso-Herrero et al. 2020; Aalto et al. 2020; Audibert et al. 2020). Accurate dynamical modelling of the molecular gas kinematics reveals kinematically decoupled nuclear structures, high velocity dispersion at nuclei, trailing spirals, and evidence of inflows and AGN-driven outflows. (e.g. Combes et al. 2019; Combes 2019, 2021). The outflow driving mechanism (wind shock, radiation pressure, or jet), their multiphase nature, and their relative weights and impact on the galaxy ISM are still open problems (Faucher-Gigu\u00e8re & Quataert 2012; Zubovas & King 2012; Richings & Faucher-Giguere 2017; Menci et al. 2019; Ishibashi et al. 2021). To date, far different outflow phases have been observed only for a handful of sources. Atomic, cold, and warm molecular outflows have been observed in radio galaxies (e.g. Morganti et al. 2007; Dasyra & Combes 2012; Dasyra et al. 2014; Tadhunter et al. 2014; Oosterloo et al. 2017). The nuclear semi-relativistic phase and the galaxy-scale molecular phase have been observed simultaneously in less than a dozen sources, with varied results: in some cases data suggest energy driven flows (Feruglio et al. 2015; Tombesi et al. 2015; Longinotti et al. 2018; Smith et al. 2019), in other cases data suggest momentum driven flows (e.g. Garc\u00eda-Burillo et al. 2014; Feruglio et al. 2017; Fluetsch et al. 2019; Bischetti et al. 2019; Marasco et al. 2020). Fiore et al. (2017), using a compilation of local and high redshift winds, showed that there is a broad distribution of the momentum boost, suggesting that both energy- and momentum-conserving expansion may occur. Enlarging the sample of local AGN-host galaxies with outflows detected in different gas phases is important to understand the nature and driving mechanisms of galaxy-scale outflows.","Citation Text":["Garc\u00eda-Burillo et al. 2014"],"Functions Text":["On galactic scales massive molecular winds are common in local Seyfert galaxies (e.g."],"Functions Label":["Background"],"Citation Start End":[[1204,1230]],"Functions Start End":[[1034,1119]]}
{"Identifier":"2019ApJ...872...52C__Linsky_2017_Instance_1","Paragraph":"Studies aiming at measuring and modeling solar radiation and its variability are strongly motivated by the impact that solar irradiance (that is, the electromagnetic energy emitted by the Sun received at the top of Earth\u2019s atmosphere in units of area and time), especially in the UV, has on the chemistry and physical properties of Earth\u2019s atmosphere and climate (e.g., Gray et al. 2010; Matthes et al. 2017). Studies of solar variability have been recently also driven by the necessity of improving our understanding of stellar variability (see Fabbian et al. 2017, for a recent review), which, in turn, is essential to characterize the habitable zones of stars and the atmospheres of their exoplanets. As for Earth, modeling of exoplanet atmospheres requires as fundamental input the spectral energy distribution of the hosting star, especially UV and shorter wavelengths (e.g., Tian et al. 2014; Ranjan et al. 2017; Rugheimer & Kaltenegger 2018). Unfortunately, measurements of UV radiation are strongly hampered by the interstellar medium absorption (up to 70%\u201390%), which is significant even for relatively close stars, so that estimates of stellar UV radiation strongly rely on modeling (see Linsky 2017, for a recent review). Moreover, because there is no mission scheduled in the near future to observe stellar spectra in the UV, after the Hubble Space Telescope ceases operations, the characterization of UV spectra of stars hosting exoplanets that will be discovered by current and future missions (e.g., TESS or James Webb Space Telescope) will necessarily rely on indirect estimates, performed, for instance, through the use of semi-empirical models (e.g., Mauas et al. 1997; Fontenla et al. 2016; Bus\u00e1 et al. 2017) or proxies (e.g., Stelzer et al. 2013; Shkolnik et al. 2014). Stellar irradiance variability also affects the detectability of exoplanets (see, e.g., the recent review by Oshagh 2018). The passage over the disk of spots and faculae may induce photometric variations of amplitude similar to or larger than photometric variations induced by planetary transits. Moreover, the presence of active regions may alter spectral line profiles, thus hindering exoplanet detections performed through radial velocity measurements. Similarly, spectroscopic techniques that allow us to estimate the physical properties of exoplanet atmospheres (see, e.g., Kreidberg 2017, for a recent view) require as fundamental input the spectra synthetized through models representing quiet and active regions (faculae and sunspots). Finally, stellar irradiance variability observed at different spectral ranges, especially in the UV, is a fundamental observable for the characterization of the magnetic activity of a star, and therefore for the understanding of dynamo processes in stellar objects (e.g., Reinhold et al. 2013; Basri 2016; Salabert et al. 2016). Because stellar photometric and spectral variability can be modeled using the semi-empirical approaches developed for the Sun described above (see, e.g., Shapiro et al. 2016; Witzke et al. 2018), understanding the limitations of current irradiance models is fundamental to improving our capability of modeling stellar variability.","Citation Text":["Linsky 2017"],"Functions Text":["Unfortunately, measurements of UV radiation are strongly hampered by the interstellar medium absorption (up to 70%\u201390%), which is significant even for relatively close stars, so that estimates of stellar UV radiation strongly rely on modeling (see",", for a recent review)."],"Functions Label":["Background","Background"],"Citation Start End":[[1198,1209]],"Functions Start End":[[950,1197],[1209,1232]]}
{"Identifier":"2015ApJ...806..118S__Moffatt_1978_Instance_1","Paragraph":"Magnetic fields observed in various astrophysical systems, such as the Earth, the Sun, disk galaxies, accretion disks, etc., possess large-scale magnetic fields in addition to a fluctuating component. The magnetic field survives for timescales much larger than the diffusion timescales in those systems, and therefore are thought to be self-sustained by turbulent dynamo action. The standard model of such a turbulent dynamo producing a large-scale magnetic field involves the amplification of seed magnetic fields due to the usual \u03b1 effect, where \u03b1 is a measure of the net kinetic helicity in the flow (see, e.g., Moffatt 1978; Parker 1979; Krause & R\u00e4dler 1980; Brandenburg & Subramanian 2005; Brandenburg et al. 2012). Since it is not necessary for the turbulent flow always to be helical, it is interesting to study the dynamo action in non-helically forced shear flows. Dynamo action due to shear and turbulence in the absence of the \u03b1 effect received some attention in the astrophysical contexts of accretion disks (Vishniac & Brandenburg 1997) and galactic disks (Blackman 1998; Sur & Subramanian 2009). The presence of large-scale shear in turbulent flows is expected to have significant effects on transport properties (R\u00e4dler & Stepanov 2006; R\u00fcdiger & Kitchatinov 2006; Leprovost & Kim 2009; Sridhar & Singh 2010; Singh & Sridhar 2011). It has also been demonstrated that the mean shear in conjunction with rotating turbulent convection gives rise to the growth of large-scale magnetic fields (K\u00e4pyl\u00e4 et al. 2008; Hughes & Proctor 2009). The problem we are interested in may be stated as follows: in the absence of the \u03b1 effect, will it be possible to generate a large-scale magnetic field solely through the action of non-helical turbulence in the background shear flow on the seed magnetic field? This question was recently numerically studied by Brandenburg et al. (2008), Yousef et al. (2008a, 2008b). These works clearly demonstrated the growth of large-scale magnetic fields due to non-helical stirring at small scales in the background linear shear flow.","Citation Text":["Moffatt 1978"],"Functions Text":["The standard model of such a turbulent dynamo producing a large-scale magnetic field involves the amplification of seed magnetic fields due to the usual \u03b1 effect, where \u03b1 is a measure of the net kinetic helicity in the flow (see, e.g.,"],"Functions Label":["Background"],"Citation Start End":[[615,627]],"Functions Start End":[[379,614]]}
{"Identifier":"2020ApJ...893...84B__Almeida_et_al._2014a_Instance_1","Paragraph":"Moreover, if mixing is effective over kiloparsec scales, then we would expect the neutral gas observed toward the QSO and the ionized gas seen toward the central H ii region to have the same metallicity. In which case, in order to find \n\n\n\n\n\n  \n\n\n\n\n\n, a significant fraction of the neutral gas must be many kiloparsecs away from the star-forming regions in order to not be contaminated. Such an explanation was adopted by Cannon et al. (2005), who postulated the existence of a low-metallicity halo beyond the inner ISM to explain the discrepancy between \n\n\n\n\n\n and \n\n\n\n\n\n toward star-forming regions in NGC 625. Our results are consistent with this idea and provide additional evidence for IGM gas feeding into galaxies through streams from the cosmic web (e.g., S\u00e1nchez Almeida et al. 2014a, and references therein). Dwarf galaxy pairs in particular are thought to have enhanced star formation because they are fed by significant reservoirs of neutral gas in which they reside, or because of their mutual interactions (Lelli et al. 2014; Stierwalt et al. 2015; Pearson et al. 2016). As noted in Section 2.4, UGC 5282 does lie near a galaxy of similar mass (UGC 5287 in Figure 3) and interactions between the two galaxies may be the source of gas flowing into UGC 5282. For example, Pearson et al. (2018) have suggested that multiple encounters between dwarf galaxies can \u201cpark\u201d gas at significant distances from the protagonists, which can then return over several Gyr. It remains unknown whether the metallicity of such returning debris would be low enough to cause the decrease in metallicity in either UGC 5282 or UGC 5287, but it may reflect whatever buildup of metals occurred in the dwarfs at much earlier times. Alternative tidal models that cause strong metallicity gradients and the removal of low-metallicity gas at the edges of dwarf galaxies (Williamson et al. 2016) seem less well supported by our results. Indeed, the situation for UGC 5282 may be even more complicated, depending on whether it has entered the halo of NGC 3003 (Figure 3) and has begun to feel any effects from ram pressure stripping by gas in the host\u2019s halo. The effects on the metallicity of the dwarf galaxy may, however, be less significant than any effects from tidal stripping (Williamson & Martel 2018).","Citation Text":["S\u00e1nchez Almeida et al. 2014a"],"Functions Text":["Our results are consistent with this idea and provide additional evidence for IGM gas feeding into galaxies through streams from the cosmic web (e.g.,"],"Functions Label":["Similarities"],"Citation Start End":[[764,792]],"Functions Start End":[[613,763]]}
{"Identifier":"2018AandA...610A..38F__Bisterzo_et_al._2017_Instance_2","Paragraph":"Similarly to the [\u03b1\/Fe] ratio, the ratio of the slow (s-) neutron capture process elements to iron can be regarded as a cosmic clock. Ba, Sr, La, and Y are mainly s-process elements produced on long timescales by low mass AGB stars (Matteucci 2012). Since a low mass star must evolve to the AGB phase before the s-process can occur, the s-process elements are characterized by a delay in the production, much like the delay of iron production by SNe Ia relative to the \u03b1 elements production by core collapse SNe. Among the four s-process elements mentioned above, GES provides the abundances of Y II (the first s-process peak element) and Ba II (the second s-process peak element) for all our sample stars. Their abundances behave differently in the Galactic thick and thin discs (Bensby et al. 2005, 2014; Israelian et al. 2014; Bisterzo et al. 2017; Delgado Mena et al. 2017). Unlike the Galactic thick disc stars, which show an almost constant [Ba\/Fe] abundance close to the solar value, the Galactic thin disc stars have their [Ba\/Fe] abundances increasing with [Fe\/H] and reaching their maximum values around solar metallicity, after which a clear decline is seen (see also Cristallo et al. 2015a,b, for the most recent s-process calculation in AGB yields). The same trend is observed in our sample. In Fig. 13 we display the Li-[Ba\/Fe], [Ba\/Fe] as a function of [Fe\/H], and the evolution of absolute Ba abundance A(Ba), as derived from Ba II lines. Similar figures are also plotted for yttrium (Y II). [Ba\/Fe] and [Y\/Fe] values here are derived from MCMC simulations, taking into account the measurement uncertainties of A(Ba II)\/A(Y II) and [Fe\/H]. By applying the same MCMC setups used for [\u03b1\/Fe] (see Sect. 3.1), we calculate the mean values of [Ba\/Fe] and [Y\/Fe] for each star. These values, together with their corresponding 1\u03c3 uncertainties, are listed in Table 1. In the literature there are several theoretical works on the evolution of [Ba\/Fe] and [Y\/Fe] in the Galactic thin disc (e.g. Pagel & Tautvaisiene 1997; Travaglio et al. 1999, 2004; Cescutti et al. 2006; Maiorca et al. 2012; Bisterzo et al. 2017). For comparison, we show in Fig. 13 the predictions of the most recent one (Bisterzo et al. 2017) where the updated nuclear reaction network was used.","Citation Text":["Bisterzo et al. 2017"],"Functions Text":["In the literature there are several theoretical works on the evolution of [Ba\/Fe] and [Y\/Fe] in the Galactic thin disc (e.g."],"Functions Label":["Background"],"Citation Start End":[[2101,2121]],"Functions Start End":[[1877,2001]]}
{"Identifier":"2018MNRAS.480.5113M__Murgia_et_al._2009_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":["Murgia et al. 2009"],"Functions Text":["The main and novel aspect of our work is the analysis of the diffuse radio emission resulting from radio haloes in galaxy clusters"],"Functions Label":["Extends"],"Citation Start End":[[539,557]],"Functions Start End":[[380,510]]}
{"Identifier":"2019ApJ...878...84M__Oppenheimer_et_al._2016_Instance_1","Paragraph":"These observations have identified a large reservoir of baryons surrounding galaxies. This circumgalactic gas extends to roughly the virial radius (Shull 2014) and contains a substantial fraction of the baryons (Werk et al. 2014; Bregman et al. 2018) and metals (Peeples et al. 2014) associated with the dark matter halo. The larger column densities of the O+5 ion in the CGM of star-forming galaxies compared to passive galaxies (Tumlinson et al. 2011) has generated great interest in the CGM, because the processes producing this dichotomy may explain why star formation is quenched in massive halos (Blanton et al. 2003; Kauffmann et al. 2003; Schawinski et al. 2014). In cosmological hydrodynamical simulations, the O vi absorbing gas lies behind the halo accretion shock and is maximal in L* galaxies because their virial temperature is close to the temperature T \u2248 105.5 K where the O+5 ionization fraction peaks (Oppenheimer et al. 2016). Feedback from supermassive black holes may suppress the O+5 fraction in the halos of red galaxies relative to the halos of blue galaxies of similar stellar mass (Nelson et al. 2018). The nucleus would not typically still be active by the time its outflow impacted the gas properties at half the virial radius, so differentiating between nuclear activity and halo mass is challenging observationally (Berg et al. 2018). Simulations that zoom in on individual galaxies include more physics than cosmological simulations (Hummels et al. 2013; Su et al. 2018). They qualitatively agree that enhanced star formation feedback increases the strength of high-ionization absorption lines, as Heckman et al. (2017) observed. Quantitatively, however, the star formation feedback does not produce enough O vi absorption, nor does it permanently quench star formation. A solution may require a completely different schema for the CGM. Stern et al. (2018) argue, for example, that the O vi absorption occurs beyond the accretion shock, where O vi would be photoionized by the UV background, a low-pressure scenario.","Citation Text":["Oppenheimer et al. 2016"],"Functions Text":["In cosmological hydrodynamical simulations, the O vi absorbing gas lies behind the halo accretion shock and is maximal in L* galaxies because their virial temperature is close to the temperature T \u2248 105.5 K where the O+5 ionization fraction peaks"],"Functions Label":["Uses"],"Citation Start End":[[920,943]],"Functions Start End":[[672,918]]}
{"Identifier":"2020AandA...644A..59K__Pastorello_et_al._2019_Instance_2","Paragraph":"The year 2020 marks the 350 yr anniversary of the discovery of the eruption of Nova 1670 (or CK Vul) made by European astronomers (Shara et al. 1985). Their observations, predominantly performed with a naked eye, traced the object\u2019s evolution on the sky in 1670\u22121672. From the archive records, we know that the eruption was rather unusual, in particular it was very much unlike classical novae. The light curve of CK Vul displayed three peaks and the star was described as reddish in the later stages of the eruption (Hevelius 1671; Shara et al. 1985). These characteristics resemble closest the behavior often observed in (luminous) red novae (Kato 2003; Tylenda et al. 2013), a modern category of eruptive stars known from our and other galaxies (e.g. Pastorello et al. 2019). Red novae are recognized as manifestations of on-going mergers of non-compact stars such as main-sequence dwarfs, sub-giants, or red giants (Soker & Tylenda 2003; Tylenda & Soker 2006; Tylenda et al. 2011; Pastorello et al. 2019). While the number of known red novae, mainly extragalactic ones, is quickly rising (e.g., Stritzinger et al. 2020), we know only a few red-nova remnants that are decades old (Kami\u0144ski et al. 2018a). The remnant of the 1670 eruption of CK Vul, as a candidate post-merger site, could be the oldest (counting from the onset of the eruption) known object of this type and as such offers the opportunity to investigate a merger aftermath centuries after the stellar coalescence. The nature of the progenitor system of CK Vul has been debated. Eyres et al. (2018) proposed that the seventeenth-century merger took place between a white dwarf and a brown dwarf, but there is little quantitative evidence to support this. Based on the analysis of the source\u2019s chemical and isotopic composition, including the unique presence of the radioactive isotope of 26Al, Kami\u0144ski et al. (2018b) found that the progenitor system of CK Vul included at least one red-giant-branch (RGB) star with a fully developed helium core.","Citation Text":["Pastorello et al. 2019"],"Functions Text":["Red novae are recognized as manifestations of on-going mergers of non-compact stars such as main-sequence dwarfs, sub-giants, or red giants"],"Functions Label":["Background"],"Citation Start End":[[985,1007]],"Functions Start End":[[779,918]]}
{"Identifier":"2022MNRAS.512..186K__Chakraborty_et_al._2019a_Instance_1","Paragraph":"A widely used statistical property of the sky brightness distribution is its power spectrum (Lazarian 1995; Bharadwaj & Sethi 2001, and others). As the redshifted 21-cm signal is expected to be faint and hard to detect with imaging, estimating its power spectrum or equivalently intensity mapping gives a possible probe of the evolution of the baryonic matter distribution over cosmic time. Bharadwaj & Sethi (2001) show that visibility correlation directly measures the power spectrum. This method and its variants (Datta, Choudhury & Bharadwaj 2007; Choudhuri et al. 2014; Choudhuri et al. 2016; Bharadwaj et al. 2019; Choudhuri et al.2019, and others) have been used to estimate the angular power spectrum of the diffused galactic foreground (Ghosh et al. 2012; Choudhuri et al. 2017b; Chakraborty et al. 2019a; Choudhuri et al. 2020) as well as the power spectrum of H\u2009i distribution in nearby galaxies (Dutta et al. 2009; Dutta & Bharadwaj 2013; Nandakumar & Dutta 2020). These works propagate the uncertainties in each visibility estimate and combine that with the sample variance error in measuring the power spectrum to quote uncertainties in the power spectrum estimates. In this work, we use the estimator discussed in Choudhuri et al. (2014), where visibilities are gridded before estimating the power spectrum. Given an angular field of view of \u03b80 to which the telescope is sensitive, it has been shown (Bharadwaj & Sethi 2001; Bharadwaj & Ali 2005; Choudhuri et al. 2014) that the visibilities in the nearby baselines remain correlated to a baseline separation of $\\Delta U \\lt \\frac{1}{\\pi \\theta _0}$. The size of the uv-grids is chosen such that they are large enough to include a sufficient number of baselines in a given uv-grid and small enough to have all visibilities in the uv-grid correlated. In each uv-grid, they estimate the power spectrum by correlating visibilities only in nearby baselines, omitting the visibility autocorrelations. This drastically reduces the noise bias in estimates of the power spectrum in uv-grids. The contribution from each uv-grid within a given annulus in $U = \\mid \\vec{U} \\mid$ is then combined, and the real part of it is used to quote the value of the isotropic power spectrum for the baseline separation U. We may schematically write it as\n(4)$$\\begin{eqnarray}\r\n\\mathcal {E} \\lbrace P(U)\\rbrace = \\mathcal {R} [\\langle \\tilde{V}(\\vec{U})^{*} \\tilde{V}(\\vec{U}+\\Delta \\vec{U}) \\rangle].\r\n\\end{eqnarray}$$Here, the average is taken over the uv-grid first and then within the annulus, as explained above. Note that the power spectrum estimator here assumes that a perfect calibration is done and the gains are all unity. In such a case, the power spectrum estimate has no bias arising from instrumental noise, and its uncertainties can be written as (Ali et al. 2008; Dutta 2011)\n(5)$$\\begin{eqnarray}\r\n\\sigma _P^2 = \\frac{P^2(U)}{N_\\mathrm{ G}} + 2\\frac{P(U)\\sigma _N^2}{N_\\mathrm{ B}} + 2\\frac{\\sigma _N^4}{N_\\mathrm{ B}},\r\n\\end{eqnarray}$$where NG is the number of independent estimates of the power spectrum in a given annulus bin at U, NB is the total number of visibility pairs in the bin.","Citation Text":["Chakraborty et al. 2019a"],"Functions Text":["This method and its variants","have been used to estimate the angular power spectrum of the diffused galactic foreground","These works propagate the uncertainties in each visibility estimate and combine that with the sample variance error in measuring the power spectrum to quote uncertainties in the power spectrum estimates."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[789,813]],"Functions Start End":[[487,515],[655,744],[977,1180]]}
{"Identifier":"2021ApJ...912..132K__Abramowicz_&_Zurek_1981_Instance_1","Paragraph":"The magnetic field assumed in our models is normalized either to \u03b2 = 100, which means thermally dominated accretion flows, or to \u03b2 = 1, which means actually quite strong magnetization and equilibrium ratio of the magnetic and gas pressures. The gas pressure is set by the solution of the Bondi transonic accretion flow, which is our initial condition for the simulation. Only if the magnetic pressure does not add a high contribution to this flow structure can we safely assume that the solution holds, and the gas falls supersonically to the black hole downstream of the sonic point, ultimately reaching the speed of light at the black hole horizon. We perturb this flow in time, by adding a slow rotation to the gas; hence, our sonic point is moving, and the shape of the sonic surface is disturbed. However, when the magnetic pressure is substantially large, the initial solution for the transonic accretion may not be a good approximation. Already at the beginning, multiple critical points might exist in the flow. The possible existence of shocks in low angular momentum flows connected with the presence of multiple critical points in the phase space has been studied from different points of view during the past 30 yr. However, the theoretical works that describe the fundamental properties of the low angular momentum accretion and that usually treat the steady solution of the equations have so far been carried out only for nonmagnetized and usually also nonviscous flows (Abramowicz & Zurek 1981; Abramowicz & Chakrabarti 1990; Das 2002; Das & Czerny 2012). Hydrodynamical models of the low angular momentum accretion flows have been studied already in two and three dimensions, e.g., by Proga & Begelman (2003) and Janiuk et al. (2008, 2009). In those simulations, a single, constant value of the specific angular momentum was assumed, while the variability of the flows occurred owing to, e.g., nonspherical or nonaxisymmetric distribution of the matter. The level of this variability was also dependent on the adiabatic index (see also Palit et al. 2019). However, these studies considered nonmagnetized flow only, which limits the applicability of these models to observational data and to the stellar evolution and presupernova studies (Heger et al. 2005). The range of strength of the magnetic field varies from very low values in the interstellar medium to extremely high values of \u223c 1015 G in the magnetars, which are extremely magnetized neutron stars. During the accretion process, the magnetic field in the star is likely to be amplified in the innermost region of the black hole, so a negligible amount of magnetic field in those regions, where the shock and the sonic point could be located, is thus not anticipated. The only method, however, to describe the effect of strong magnetic fields in the context of transonic accretion with low angular momentum is via numerical simulations, because no analytic solution of this problem exists, to the best of our knowledge. Hence, setting up the initial condition with strong magnetic fields overimposed on the slowly rotating transonic Bondi flow is a simple but working approach that we make here.","Citation Text":["Abramowicz & Zurek 1981"],"Functions Text":["However, the theoretical works that describe the fundamental properties of the low angular momentum accretion and that usually treat the steady solution of the equations have so far been carried out only for nonmagnetized and usually also nonviscous flows"],"Functions Label":["Background"],"Citation Start End":[[1485,1508]],"Functions Start End":[[1228,1483]]}
{"Identifier":"2021AandA...655A..12T__Tang_et_al._2017b_Instance_7","Paragraph":"Using the RADEX3 non local thermodynamic equilibrium (LTE) modeling program (van der Tak et al. 2007) with collisional rate coefficients from Wiesenfeld & Faure (2013), we modeled the relation between the gas kinetic temperature and the measured average of para-H2CO 0.5 \u00d7 [(322\u2013221 + 321\u2013220)\/303\u2013202] ratios, adopting a 2.73 K background temperature, an average observational linewidth of 4.0 km s\u22121, and column densities N(para-H2CO) = 2.7 \u00d7 1012 and 3.7 \u00d7 1012 cm\u22122 for N113 and N159W, respectively. The results are shown in Fig. 5. The values of the para-H2CO column density were obtained with APEX data (beam size ~30\u2033; Tang et al. 2017b), which cover similar regions. Different column densities of para-H2CO only weakly affect derived kinetic temperatures (see Fig. 3 in Tang et al. 2017b or Fig. 4 in Tang et al. 2018a; this was also shown in Fig. 13 and discussed in Sect. 4.3.1 of Mangum & Wootten 1993) as long as all lines are optically thin. Considering that the relation between the gas temperature and the para-H2CO line ratio may vary at different spatial densities (see Fig. 2 in Tang et al. 2017b), we modeled it at spatial densities 104, 105, and 106 cm\u22123 in Fig. 5. It appears that Tkin at n(H2) = 105 cm\u22123 is consistently lower than values at 104 and 106 cm\u22123 by \u227223% and \u227234%, respectively, for Tkin \u2272 100 K. Local thermodynamic equilibrium (LTE) is a good approximation for the H2CO level populations under optically thin and high-density conditions (Mangum & Wootten 1993; Tang et al. 2017a,b, 2018b). Following the method applied by Tang et al. (2017b) in their Eq. (2), we plot the relation between the LTE kinetic temperature, TLTE, and the para-H2CO (3\u20132) line ratio in Fig. 5. Apparently, TLTE agrees well with Tnon-LTE at volume densities n(H2) ~ 105 cm\u22123 as long as Tkin \u2272 100 K. Previous observations show that para-H2CO (3\u20132) is sensitive to gas temperature at density 105 cm\u22123 (Ginsburg et al. 2016; Immer et al. 2016; Tang et al. 2017b). The spatial density measured with para-H2CO (303\u2013202) and C18O (2\u20131) in N113 and N159W is n(H2) ~ 105 cm\u22123 on a size of ~30\u2033 (Tang et al. 2017b). Therefore, here we adopt 105 cm\u22123 as an averaged spatial gas density in the N113 and N159W regions.","Citation Text":["Tang et al. 2017b"],"Functions Text":["The spatial density measured with para-H2CO (303\u2013202) and C18O (2\u20131) in N113 and N159W is n(H2) ~ 105 cm\u22123 on a size of ~30\u2033"],"Functions Label":["Uses"],"Citation Start End":[[2099,2116]],"Functions Start End":[[1973,2097]]}
{"Identifier":"2015ApJ...810..107M__LBGs_and_Ouchi_et_al._2008_Instance_1","Paragraph":"Here, we present a revised calculation of the emissivity of ionizing photons at z = 2.85 based on the analysis of the HST data in the HS1549 field. We estimate the comoving specific emissivity as\n3\n\n\n\n\n\nfollowing the assumptions of M13 and Nestor et al. (2013). In this expression, L is the non-ionizing UV luminosity, \u03a6 is the non-ionizing UV luminosity function, and (FUV\/FLyC)corr is the average flux-density ratio of non-ionizing to ionizing UV radiation for the entire galaxy sample, corrected for the mean IGM attenuation in the LyC spectral region.19\n\n19\nTo correct for absorption of LyC photons by neutral hydrogen in the IGM, we use the sample-averaged transmission values calculated in M13: tLAE = 0.44 \u00b1 0.03 and tLBG = 0.35 \u00b1 0.04.\n We perform this emissivity calculation separately for the main sample of spectroscopically confirmed LBGs and LAEs from M13 (using the UV luminosity functions from Reddy et al. 2008 for LBGs and Ouchi et al. 2008 for LAEs), and combine the LBG and LAE emissivities to obtain a total emissivity for star-forming galaxies. As in M13, we use the LRIS V-band to represent non-ionizing UV flux and NB3420 to represent LyC flux. The difference between our calculation and that of M13 lies in our estimation of the average flux-density ratio. Rather than estimating the average amount of foreground contamination from simulations, we instead know exactly which galaxies are contaminated based on the HST data. There were only two NB3420-detected galaxies in the M13 spectroscopic sample for which we were unable to acquire U336V606J125H160 imaging (D24 and lae4680), and for these objects we could not evaluate whether or not the NB3420 detections are due to foreground contamination. We thus calculate the emissivity twice in order to quote the full range of possible values: in one calculation we assume that MD5 is the only true LyC detection, and in the other calculation we assume that MD5, D24, and lae4680 are all true LyC-emitters. In addition to using the HST data to remove the NB3420 flux of foreground contaminants, we also use these measurements to estimate the percentage of contaminated flux in the non-ionizing UV. All objects with foreground contaminants identified through the HST imaging are blended in the LRIS V imaging, and it is impossible to isolate the uncontaminated z \u223c 2.85 flux in the LRIS V image. Thus, for each contaminated object, we decrease its LRIS V-band flux to match the fraction of uncontaminated V606 flux in the HST imaging. For objects that do not have HST U336V606J125H160imaging and are undetected in NB3420, we decrease their LRIS V-band flux to match the average fraction of uncontaminated V606 flux in the full sample of HST-imaged galaxies without NB3420 detections (99% for LBGs, 91% for LAEs). Finally, we use the same sample-averaged IGM correction to compute (FUV\/FLyC)corr as described in M13, employing statistics of H i  absorbers from Rudie et al. (2013). We note that the clustering of Lyman limit systems is not taken into account in these absorber statistics, and thus the true mean IGM transmission may be slightly higher than the values presented in M13 (see, e.g., Prochaska et al. 2014).","Citation Text":["Ouchi et al. 2008"],"Functions Text":["We perform this emissivity calculation separately for the main sample of spectroscopically confirmed LBGs and LAEs from M13 (using the UV luminosity functions from","for LAEs),"],"Functions Label":["Uses","Uses"],"Citation Start End":[[940,957]],"Functions Start End":[[745,908],[958,968]]}
{"Identifier":"2015AandA...584A..32M__Smol\u010di\u0107_et_al._2015b_Instance_1","Paragraph":"Ikarashi et al. (2015) discussed that if both the radio and FIR continuum are tracers of star-forming regions, then the z \u2273 3 SMGs are more compact than the lower-redshift SMGs typically observed in radio continuum emission (e.g. Biggs & Ivison 2008). As shown in Fig. 7, our present VLA 3 GHz data do not suggest such a trend, and, as mentioned earlier, there is actually a hint of larger radio sizes at z ~ 2.5\u20135 compared to lower redshifts. However, the highest-redshift SMG in our sample, AzTEC3 at z \u2243 5.3, shows the most compact size among our sources, consistent with the rest-frame FIR sizes from Ikarashi et al. (2015). We note that Capak et al. (2011) found that AzTEC3 belongs to a spectroscopically confirmed protocluster containing eight galaxies within a 1 arcmin2 area, and therefore the environment might also play a role in the galaxy size evolution (see also Smol\u010di\u0107 et al. 2015b). However, it is currently unclear whether the environmental effects in a galaxy overdensity will lead to a more compact or more extended radio-emitting size than field galaxies. On one side, a protocluster environment is expected to show an elevated merger rate (e.g. Hine et al. 2015), and, as discussed above, mergers are expected to pull the galactic magnetic fields to larger spatial scales, and hence lead to a more extended radio synchrotron emission. On the other side, the ram and\/or thermal pressures of the intracluster medium could compress the ISM of the galaxy, increase the magnetic field strength, and hence cause an excess in radio emission (consistent with a low IR-radio q parameter of \u22722 for AzTEC3; Miettinen et al., in prep.). The aforementioned pressure forces can drive shock waves into the ISM, and hence accelerate the CR particles (Murphy 2009). Consequently, the cooling time and diffusion length-scale of CR electrons can decrease (see Appendix E), resulting in a compact radio-emitting area. More detailed environmental analysis of SMGs is needed to understand this further. ","Citation Text":["Smol\u010di\u0107 et al. 2015b"],"Functions Text":["We note that Capak et al. (2011) found that AzTEC3 belongs to a spectroscopically confirmed protocluster containing eight galaxies within a 1 arcmin2 area, and therefore the environment might also play a role in the galaxy size evolution (see also","However, it is currently unclear whether the environmental effects in a galaxy overdensity will lead to a more compact or more extended radio-emitting size than field galaxies."],"Functions Label":["Background","Compare\/Contrast"],"Citation Start End":[[877,897]],"Functions Start End":[[629,876],[900,1076]]}
{"Identifier":"2022MNRAS.511.1362T__Tanikawa_et_al._2021a_Instance_1","Paragraph":"Among the main proposed formation channels for merging binary COs, we find: pairing of primordial BHs (e.g. Carr & Hawking 1974; Bird et al. 2016; Carr et al. 2016; Scelfo et al. 2018; De Luca et al. 2021); isolated binary evolution via common envelope (e.g. Tutukov & Yungelson 1973; Bethe & Brown 1998; Portegies Zwart & Yungelson 1998; Belczynski, Kalogera & Bulik 2002; Belczynski et al. 2008; Dominik et al. 2013; Belczynski et al. 2016; Eldridge & Stanway 2016; Mapelli et al. 2017; Stevenson, Berry & Mandel 2017; Ablimit & Maeda 2018; Klencki et al. 2018; Kruckow et al. 2018; Mapelli & Giacobbo 2018; Eldridge, Stanway & Tang 2019; Mapelli et al. 2019; Neijssel et al. 2019; Spera et al. 2019), via stable mass transfer (e.g. Kinugawa et al. 2014; Inayoshi et al. 2017; van den Heuvel, Portegies Zwart & de Mink 2017; Kinugawa, Nakamura & Nakano 2020; Tanikawa et al. 2021a,b), or via chemically homogeneous mixing (e.g. de Mink & Mandel 2016; Mandel & de Mink 2016; Marchant et al. 2016; du Buisson et al. 2020); dynamical perturbations in the field (Michaely & Perets 2019, 2020); dynamical formation in young star clusters (YSCs, e.g. Banerjee, Baumgardt & Kroupa 2010; Ziosi et al. 2014; Mapelli 2016; Askar et al. 2017; Banerjee 2017; Banerjee 2018; Rastello et al. 2018; Di Carlo et al. 2019, 2020a,b; Kumamoto, Fujii & Tanikawa 2019, 2020; Rastello et al. 2020; Banerjee 2021; Rastello et al. 2021; Trani et al. 2021); globular clusters (GCs, e.g. Portegies Zwart & McMillan 2000; Downing et al. 2010; Tanikawa 2013; Samsing, MacLeod & Ramirez-Ruiz 2014; Rodriguez et al. 2015; Rodriguez, Chatterjee & Rasio 2016; Rodriguez & Antonini 2018; Samsing, Askar & Giersz 2018; Zevin et al. 2019; Antonini & Gieles 2020); nuclear star clusters (NSCs, e.g. Miller & Lauburg 2009; O\u2019Leary, Kocsis & Loeb 2009; Antonini & Perets 2012; Prodan, Antonini & Perets 2015; Antonini & Rasio 2016; Petrovich & Antonini 2017; Arca-Sedda & Gualandris 2018; Gond\u00e1n et al. 2018; Arca-Sedda & Capuzzo-Dolcetta 2019; Rasskazov & Kocsis 2019; Arca Sedda 2020; Arca Sedda et al. 2020); and AGN discs (e.g. McKernan et al. 2012; Bartos et al. 2017; Stone, K\u00fcpper & Ostriker 2017; McKernan et al. 2018; Yang et al. 2019; Tagawa, Haiman & Kocsis 2020).","Citation Text":["Tanikawa et al. 2021a"],"Functions Text":["Among the main proposed formation channels for merging binary COs, we find:","via stable mass transfer (e.g."],"Functions Label":["Background","Background"],"Citation Start End":[[861,882]],"Functions Start End":[[0,75],[704,734]]}
{"Identifier":"2020MNRAS.497.2651K__Teuff,_Millar_&_Markwick_2000_Instance_1","Paragraph":"As the H\u2009ii regions S235 A and S235 C are deeply embedded in the molecular cloud, we cannot use published predictions from standard plane-parallel PDR models without foreground absorption (e.g. Kaufman et al. 1999). In order to understand how the observational view by SOFIA is matched by a \u2019classical\u2019 model of expanding H\u2009ii regions, described analytically and numerically by Spitzer (1978), Elmegreen & Lada (1977), Hosokawa & Inutsuka (2006), Raga et al. (2012), Kirsanova et al. (2009), and Bisbas et al. (2015), we make simulations with the MARION model (Kirsanova et al. 2009; Akimkin et al. 2015). The gas-phase chemical network from R\u00f6llig et al. (2007) (mainly based on the UMIST99 ratefile; Le Teuff, Millar & Markwick 2000) allows us to reproduce the C+, CO, and HCO+ abundances in PDRs (see Kirsanova & Wiebe 2019), so we use this network together with the ionization of the atomic species and corresponding recombinations in the H\u2009ii regions. The cross-sections for most of photoreactions are taken from the Leiden data base of \u2018Photodissociation and photoionization of astrophysically relevant molecules\u2019 (Heays, Bosman & van Dishoeck 2017). We implement the formation of H2 on grain surfaces and accretion and desorption processes of other neutral species, but other chemical reactions on dust surfaces are not considered in the calculations to save computation time. The rates of accretion and desorption processes are based on the work by Hasegawa & Herbst (1993), with updated desorption energies from Garrod, Wakelam & Herbst (2007). We use the \u2018high-metallicity\u2019 initial elemental abundances based on the \u2018EA2\u2019 set from Wakelam & Herbst (2008). The initial conditions are cold, molecular, and solid. This means that we start with all carbon and oxygen in CO and H2O on dust surfaces. Fifty two chemical species are included: H, H+, H2, H$_2^+$, H$_3^+$, O, O+, O++, OH+, OH, O2, O2:d, O$_2^+$, H2O, H2O+, H3O+, C, C+, C++, CH, CH+, CH2, CH$_2^+$, CH3, CH$_3^+$, CH4, CH$_4^+$, CH$_5^+$, CO, CO:d, CO+, HCO+, He, He+, S, S+, S++, Si, Si+, H:d, H2:d, O:d, OH:d, H2O:d, C:d, CH:d, CH2:d, CH3:d, CH4:d, S:d, Si:d, and e\u2212, where the postfix \u2018:d\u2019 indicates species on dust grain ices. Chemical species containing Si and S are included only to obtain correct gas temperature in ionized region. We recognize that the chemical network is far from complete, but our choise was motivated by limited computation time. The heating and cooling processes included in the model are listed in Akimkin et al. (2015) and Kirsanova & Wiebe (2019).","Citation Text":["Le Teuff, Millar & Markwick 2000"],"Functions Text":["The gas-phase chemical network from R\u00f6llig et al. (2007) (mainly based on the UMIST99 ratefile;","allows us to reproduce the C+, CO, and HCO+ abundances in PDRs","so we use this network together with the ionization of the atomic species and corresponding recombinations in the H\u2009ii regions."],"Functions Label":["Uses","Uses","Uses"],"Citation Start End":[[702,734]],"Functions Start End":[[606,701],[736,798],[829,956]]}
{"Identifier":"2017ApJ...850...75S__Bugaev_et_al._2016_Instance_1","Paragraph":"A realistic EoS that is able to reproduce the properties of compact astrophysical objects has to fulfill several requirements. The possibility of including many particle species, which is known as multicomponent character, is of crucial importance for modeling the NS interiors, which in even the simplest treatment include neutrons, protons, and electrons, while more advanced descriptions have to account for the presence of hyperons (Schaffner-Bielich et al. 2002). Therefore, the grand canonical ensemble is the natural choice for the formulation of such an EoS. Another element of the realistic phenomenological hadronic EoS corresponds to the short-range repulsive interaction of the hard core nature between particles  (Andronic et al. 2006; Bugaev et al. 2016). Analysis of the particle yields produced in relativistic A+A collisions within statistical (thermal) models, i.e., the Hadron Resonance Gas (HRG) model (Bugaev et al. 2016), shows the importance of the particle hard core repulsion. In this approach every particle species is defined as a rigid sphere with a fixed radius estimated from experimental data analysis. These radii do not exceed 0.5 fm  (Andronic et al. 2017; Sagun et al. 2017a). Note that the hard core of hadrons in phenomenological EoSs is important in order to suppress thermal excitations of the hadronic spectrum and provide deconfinement of the color degrees of freedom expected at high temperatures\/densities  (Satz 2012). Another requirement to the phenomenological EoS is related to its causal behavior when the speed of sound cannot exceed the speed of light. At sufficiently high densities this condition is violated by the hard core repulsion. As was shown by Sagun et al. (2017a), introducing the induced surface tension (IST) of particles to the model with the hard core repulsion between an arbitrary number of hadron species makes the EoS significantly softer and extends its causality range up to 7.5 normal nuclear densities, where formation of the quark-gluon plasma is expected. The IST is the key element of this approach (Sagun et al. 2014), as it allows us to account for the hard core repulsion between constituents in the most accurate way, and to properly reproduce the virial expansion of the multicomponent EoS. Recently, the IST EoS was used to describe the experimental data of hadron multiplicities measured at AGS, SPS, RHIC, and LHC energies of nuclear collisions (Sagun et al. 2017b), as well as the nuclear matter properties (Sagun et al. 2014). In this work the focus is on the application of IST EoS to the study of NS properties.","Citation Text":["Bugaev et al. 2016"],"Functions Text":["Another element of the realistic phenomenological hadronic EoS corresponds to the short-range repulsive interaction of the hard core nature between particles"],"Functions Label":["Background"],"Citation Start End":[[749,767]],"Functions Start End":[[567,724]]}
{"Identifier":"2022ApJ...935..136O__Ioppolo_et_al._2011_Instance_1","Paragraph":"HCOOH is the simplest carboxylic acid and has been observed toward high-mass and low-mass star-forming regions (e.g., Woods et al. 1983; Liu et al. 2001, 2002; Bisschop et al. 2007; Lefloch et al. 2017; Oya et al. 2017; Csengeri et al. 2019), protoplanetary disks (e.g., Favre et al. 2018), pre-stellar sources (e.g., Irvine et al. 1990; Vastel et al. 2014; Jim\u00e9nez-Serra et al. 2016), and comets of the solar system (Biver et al. 2014). Although the production of HCOOH has also been investigated experimentally (Ioppolo et al. 2011) and theoretically (Tielens & Hagen 1982; Garrod & Herbst 2006; Aikawa et al. 2008; Garrod et al. 2008; Vasyunin et al. 2017), its formation process and chemical link to nitrogen-bearing species are puzzling. Indeed, the chemical network calculation tracing a pre-stellar\/protostellar core by Aikawa et al. (2020) indicates that nitrogen-bearing COMs and HCOOH appear in the same region (i.e., a hot core) as oxygen-bearing COMs: no differentiation is found. It should be noted that the differentiation is not ascribed to the desorption temperature of these molecules formed on grain mantles. As described in Section 4.2, the temperature of the outer envelope traced by CH3OH is 150\u2013165 K, which is higher than that of the inner envelope traced by HCOOH and NH2CHO, 75\u2013112 K. Hence, ice mantles containing various COMs should have already been liberated from dust grains in the outer envelope. However, the results of PCA-3D suggest that HCOOH and the nitrogen-bearing species, NH2CHO, HNCO, and HC3N, do not appear in the gas phase outside the radius of \u223c0.\u203306, even though these desorption temperatures are comparable to those of CH3OH (e.g., Oya et al. 2019). In other words, other factors rather than the desorption temperature should be responsible for these molecules to be observed in the gas phase. It may be the high-density condition, protostellar radiation, or both of them. The gas-phase production of nitrogen-bearing COMs and HCOOH should be considered seriously.","Citation Text":["Ioppolo et al. 2011"],"Functions Text":["Although the production of HCOOH has also been investigated experimentally","its formation process and chemical link to nitrogen-bearing species are puzzling."],"Functions Label":["Background","Motivation"],"Citation Start End":[[514,533]],"Functions Start End":[[438,512],[661,742]]}
{"Identifier":"2016MNRAS.457.3492H__Boylan-Kolchin_et_al._2009_Instance_1","Paragraph":"N-body simulations represent the most widely used and convenient method of exploring the highly non-linear regime of cosmic structure formation. Starting from a set of initial conditions, the numerical simulations follow the formation and evolution of structures from an early epoch down to present day. Motivated by the fact that DM represents most of the matter in the Universe and because of the relatively simple physics of collisionless DM particles, DM-only simulations represent the most widely used category of numerical simulations. When designing a cosmological N-body experiment, one is concerned by two major factors. Ideally, one would like to simulate a region of the universe that is as large as possible to get a representative census of the structures encompassed within it. On the other hand, one would also want very high mass resolution, to be able to resolve accurately even the smallest cosmologically relevant objects. Unfortunately, due to limited computational resources, these two requirements are in conflict, which implies that various compromises need to be made when designing a numerical simulation. So far, the biggest efforts were focused into two, somewhat complementary approaches. The first is represented by simulations like Millennium (ms; Springel et al. 2005), Millennium II (ms-II; Boylan-Kolchin et al. 2009), Millennium XXL (MXXL; Angulo et al. 2012), Bolshoi (Klypin et al. 2011), MultiDark (Prada et al. 2012), Horizon Run I-III (Kim et al. 2009, 2011), Horizon-4\u03c0 (Prunet et al. 2008; Teyssier et al. 2009), MareNostrum Universe (Gottloeber et al. 2006), Jubilee project (Watson et al. 2014), Coyote Universe (Heitmann et al. 2010), DEUS simulation (Alimi et al. 2012; Rasera et al. 2014) or MICE suite (Fosalba et al. 2015). These follow structure formation in a large cosmological volume at the expense of having a medium or a low-mass resolution. Such simulations provide the formation histories for a very large number of medium- and high-mass DM haloes, but do not necessary resolve all the details relevant for galaxy formation. On the other side we have N-body simulations like the aquarius project (Springel et al. 2008), the Via Lactea (Diemand, Kuhlen & Madau 2007), the Phoenix project (Gao et al. 2012), CLUES (Gottloeber, Hoffman & Yepes 2010) and the ELVIS suite (Garrison-Kimmel et al. 2014b) that are characterized by a very high mass and force resolution but are limited to very small cosmic volumes. These give a very detailed picture of galaxy- and cluster-size haloes, but do so only for a very limited number of objects, which makes their results sensitive to small number statistics, and are unable to capture the full interconnection between small (DM haloes) and large (the cosmic web) cosmic scales.","Citation Text":["Boylan-Kolchin et al. 2009"],"Functions Text":["So far, the biggest efforts were focused into two, somewhat complementary approaches. The first is represented by simulations like","Millennium II (ms-II;","These follow structure formation in a large cosmological volume at the expense of having a medium or a low-mass resolution. Such simulations provide the formation histories for a very large number of medium- and high-mass DM haloes, but do not necessary resolve all the details relevant for galaxy formation."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[1323,1349]],"Functions Start End":[[1131,1261],[1301,1322],[1772,2080]]}
{"Identifier":"2019AandA...632A..76N__Ni_2018_Instance_1","Paragraph":"The Juno spacecraft has measured Jupiter\u2019s gravitational field to high precision through precise Doppler tracking in its polar orbit around Jupiter, compared with the previous values detected by Pioneer 10 and 11 and by Voyager 1 and 2 (Folkner et al. 2017; Bolton et al. 2017; Iess et al. 2018). These new gravity data have improved our knowledge of Jupiter\u2019s interior. The even gravity harmonics are affected by the shape and internal structure of Jupiter in its hydrostatic equilibrium under the effect of rotational distortion. To accurately describe the shape and internal structure of Jupiter, various interior models with new ingredients, such as two-layer and three-layer structure models, have been established by several groups (Guillot & Morel 1995; Guillot 1999; Hubbard 1999, 2013; Anderson & Schubert 2007; Kaspi et al. 2010; Nettelmann et al. 2012; Helled et al. 2015; Vazan et al. 2015; Kong et al. 2016; Wahl et al. 2017; Ni 2018; Guillot et al. 2018; Debras & Chabrier 2019). Most of these models describe the internal structure of Jupiter via the physical laws of nature given an empirical or simulated equation of state (EOS), such as polytropic EOSs, EOSs obtained using the free-energy minimization method (Saumon et al. 1995), and ab initio EOSs (Nettelmann et al. 2008, 2012; French et al. 2012; Becker et al. 2014; Militzer et al. 2008; Militzer & Hubbard 2013; Chabrier et al. 2019). The EOSs describe the microscopic properties of planetary matter and play an important role in calculating the structure and evolution of planets, in spite of uncertainties in the hydrogen\u2013helium phase separation (Hubbard et al. 2002; Saumon & Guillot 2004; Fortney & Nettelmann 2010; Miguel et al. 2016, 2018; Militzer et al. 2016; Debras & Chabrier 2019). Juno\u2019s gravity measurements have demonstrated small values for high-order even gravitational harmonics with respect to the heavy-element abundance measured by the Galileo entry probe (Wahl et al. 2017; Debras & Chabrier 2019). In order to reconcile the calculated gravitational harmonics with these small values of J4 to J8, Wahl et al. (2017) proposed a dilute core model where heavy elements are dissolved in a hydrogen\u2013helium mixture and expanded outward through a portion of Jupiter. Alternatively, these latter authors modified the abundances of helium and heavy elements in the outer molecular envelope to be lower than those measured by the Galileo entry probe. Stimulated by dilute cores, the three-layer structure models for Jupiter have been generalized into four-layer structure models by introducing a dilute core region above central compact cores. Guillot et al. (2018) compared the effective even harmonics \n\n$J_{2i}^{\\mathrm{eff}}(H,m)\\;{=}\\;J_{2i}^{\\mathrm{Juno}}-\\bigtriangleup J_{2i}(H,m)$\nJ2ieff(H,m)\u2009=\u2009J2iJuno\u2212\u25b3J2i(H,m)\n with the ones obtained from interior models assuming rigid rotation to explore the rotation of Jupiter\u2019s deep interior. It is found that Jupiter\u2019s deep interior exhibits an almost rigid-body rotation and Jupiter\u2019s atmospheric zonal flow extends to a depth H of 2000\u22123500 km. Debras & Chabrier (2019) established new models of Jupiter which satisfy both Juno\u2019s gravity data and the outer helium and heavy element abundances from Galileo, showing a considerable entropy increase between the outer and inner envelopes and an inward decreasing abundance of heavy elements in the inner envelope.","Citation Text":["Ni 2018"],"Functions Text":["To accurately describe the shape and internal structure of Jupiter, various interior models with new ingredients, such as two-layer and three-layer structure models, have been established by several groups"],"Functions Label":["Background"],"Citation Start End":[[939,946]],"Functions Start End":[[532,737]]}
{"Identifier":"2018MNRAS.478...95K__Caselli_et_al._2002_Instance_1","Paragraph":"Strong controversies are however still present regarding the way stars form in gravitationally unstable cores, in particular high-mass stars. Whether it happens on approximately a free-fall time, as suggested in the competitive accretion scenario by Bonnell et al. (2001), or rather slowly, implying at least several free-fall times, in the core accretion model by McKee & Tan (2003), which assumes that turbulence and\/or magnetic fields provide a major stabilizing contribution. It is therefore of central importance to obtain independent constraints on the time-scale of star formation. A potentially powerful tool that was suggested in the literature is the measurement of deuteration fractions, which may be translated into time-scales via chemical models (Caselli et al. 2002; Fontani et al. 2011; Pagani et al. 2011a; Kong et al. 2015; Barnes et al. 2016; Lackington et al. 2016). In relation to filaments, a time-scale estimate has recently been achieved by Lackington et al. (2016) for the infrared dark cloud L332. They found deuteration ratios N2D+\/N2H+ in the range 0.003\u20130.14. Based on the chemical models by Kong et al. (2015), the authors deduced time-scales for various cores within L332 of the order of several free-fall times to match the observed deuteration fractions, indicating rather old cores. The free-fall time of their best-fitting model was tff = 1.39 \u00d7 105\u2009yr at a density of nH = 105\u2009cm\u22123. However, the authors emphasize that a change in the CO depletion factor reduces the time-scale to only one free-fall time, which points towards (dynamically) young objects. In addition, Barnes et al. (2016) deduce a chemical age of about eight free-fall times (defined for spherical systems at a density of n = 104\u2009cm\u22123, giving tff = 3.4 \u00d7 105\u2009yr) for the dark cloud G035.39-00.33 to fit the observed average deuteration fraction of 0.04 \u00b1 0.01, which they attribute to support of the filament by magnetic fields and turbulence. Furthermore, these authors find that deuteration is widespread over the filament rather than concentrated in individual cores, in agreement with recent findings by Pillai et al. (2012) who report widespread H2D+ emission in Cygnus X.","Citation Text":["Caselli et al. 2002"],"Functions Text":["A potentially powerful tool that was suggested in the literature is the measurement of deuteration fractions, which may be translated into time-scales via chemical models"],"Functions Label":["Background"],"Citation Start End":[[761,780]],"Functions Start End":[[589,759]]}
{"Identifier":"2016AandA...588A..44Y__Jones_et_al._2014_Instance_1","Paragraph":"As described in K\u00f6hler et al. (2015)2, we assume that the dust properties change with increasing local density through accretion and coagulation. First, in the transition between the diffuse ISM and dense clouds, a second mantle can form on the surface of the CM grains that is due to the coagulation of the small aromatic-rich carbon grains on top of the bigger grains, which might be subsequently hydrogenated (a-C \u2192 a-C:H) and\/or to the accretion of a-C:H material from the gas phase. Such carbonaceous mantles are efficiently processed by UV photons (Alata et al. 2014), but can stay H-rich as long as the radiation field is attenuated (Jones et al. 2014) or if the rehydrogenation process is efficient enough. This leads to grains with two mantles (core\/mantle\/mantle grains, CMM). Second, inside dense clouds, the CMM grains can coagulate into aggregates (AMM). On average, regarding the material abundance, one aggregate is composed of three CMM grains with amorphous silicate cores and one CMM grain with an amorphous carbon core. The formation of ice mantles on the surface of the aggregates (AMMI) can also occur in the densest regions, where the shielding from energetic photons is efficient enough to allow either gas molecules to form and to freeze out on the grains or surface chemistry to proceed effectively. The three types of evolved grains (CMM, AMM, and AMMI) contain 406 ppm of C in agreement with Parvathi et al. (2012), who found that 355 \u00b1 64 ppm of C are enclosed within grains for lines of sights where NH \u2273 2 \u00d7 1021 H\/cm2 (equivalently where E(B \u2212 V) \u2273 0.4 or AV \u2273 2). For each grain type, a size distribution consistent with grain growth is considered as explained in K\u00f6hler et al. (2015) and shown in Fig. 2. For the CMM grains, ~65% of the dust mass is in grains with \\hbox{$40~{\\rm nm} \\leqslant a \\leqslant 0.25$}40 nm \u2a7d a\u2a7d 0.25\u03bcm, ~25% in grains with \\hbox{$0.25 \\leqslant a \\leqslant 0.5$}0.25 \u2a7d a\u2a7d 0.5\u03bcm, and ~10% in grains with \\hbox{$0.5 \\leqslant a \\leqslant 0.7$}0.5 \u2a7d a\u2a7d 0.7\u03bcm. For the AMM (AMMI) aggregates, ~50% of the dust mass is in grains with \\hbox{$48~(91)~{\\rm nm} \\leqslant a \\leqslant 0.25$}48 (91) nm \u2a7d a\u2a7d 0.25\u03bcm, ~40% in grains with \\hbox{$0.25 \\leqslant a \\leqslant 0.5$}0.25 \u2a7d a\u2a7d 0.5\u03bcm, and ~10% in grains with \\hbox{$0.5 \\leqslant a \\leqslant 0.7$}0.5 \u2a7d a\u2a7d 0.7\u03bcm. The optical properties of all grains and a detailed description of the calculation method can also be found in K\u00f6hler et al. (2015), while a more specific description of the grain scattering properties and efficiencies are given in Paper I. A schematic view of the dust composition and stratification from the diffuse ISM to dense molecular clouds is shown in Fig. 1. ","Citation Text":["Jones et al. 2014"],"Functions Text":["Such carbonaceous mantles are efficiently processed by UV photons","but can stay H-rich as long as the radiation field is attenuated"],"Functions Label":["Background","Background"],"Citation Start End":[[641,658]],"Functions Start End":[[488,553],[575,639]]}
{"Identifier":"2020AandA...644A..59K__Heays_et_al._2017_Instance_2","Paragraph":"Analyzing optical emission lines, emanating from within the northern lobe, Tylenda et al. (2019) found a reddening with EB\u2005\u2212\u2005V\u2004\u2248\u20040.9 mag or AV\u2004\u2248\u20042.8 mag, which we assume is mainly circumstellar in origin. Hajduk et al. (2013) observed two stars shining through the southern lobe and found AV\u2004=\u20043.3\u2005\u2212\u20054.4 mag with unknown contribution from the interstellar component. We assume here that those observations quantify the amount of circumstellar dust that is the main actor in shielding molecules from the central source. We recalculated the lifetimes of molecules assuming AV\u2004=\u20043 mag, and with (1) standard dust properties (i.e. with composition and size distribution as of interstellar dust) or (2) with larger and less opaque grains, at the gas-to-dust mass ratio of 124 (see Heays et al. 2017, for more details on the assumed dust properties). We used shielding functions from Heays et al., which include effects in lines. Results are shown in Cols. (3) and (4) of Table 3. The presence of ISM grains makes it possible for the observed molecules to survive for a very long time, longer than 350 yr. The lifetimes in the presence of the large grains considered by Heays et al. are typically a few times shorter than the age of the remnant. It is uncertain what kind of grains populate the dusty remnant of CK Vul, but given its anomalous elemental and molecular compositions and eruptive history, dust may have a peculiar chemical composition and size distribution. In such a case, the total to selective extinction law would also be different and the assumed AV may not be adequate. Nevertheless, if the molecules formed 350 yr ago and are shielded by big grains, with the calculated lifetimes a considerable fraction of molecular species would survive, except perhaps for a few most fragile ones which indeed are almost absent in the lobes. We conclude that the lifetimes in Table 3 that were calculated with an attenuated ISRF are consistent with the molecule formation 350 yr ago or more recently.","Citation Text":["Heays et al."],"Functions Text":["We used shielding functions from","which include effects in lines."],"Functions Label":["Uses","Uses"],"Citation Start End":[[878,890]],"Functions Start End":[[845,877],[892,923]]}
{"Identifier":"2015AandA...576A...5C__J\u00f8rgensen_et_al._2012_Instance_2","Paragraph":"The relative abundances of the three species are derived from the column densities in Table 2 and are compared with other star-forming regions and comets in Table 3. The (CH2OH)2\/CH2OHCHO abundance ratio of ~0.3\u20130.5 previously derived in IRAS 16293 by J\u00f8rgensen et al. (2012) was revised. Indeed, the assignment in J\u00f8rgensen et al. (2012) was based on only one line of the gGg\u2032 conformer of ethylene glycol about 200 cm-1 (~290 K, M\u00fcller & Christen 2004) above the lowest-energy aGg\u2032 conformer \u2013 and thus tentative. An analysis from observations of six transitions of the lower energy conformer from ALMA Cycle 1 observations at 3 mm (four spectral windows at 89.48\u201389.73, 92.77\u201393.03, 102.48\u2013102.73, and 103.18\u2013103.42 GHz; J\u00f8rgensen et al., in prep.) results in a higher ethylene glycol-to-glycolaldehyde abundance ratio of 1.0\u2009\u00b1\u20090.3. This new estimate is consistent with the ratio expected between the aGg\u2032 and gGg\u2032 conformers under thermal equilibrium conditions at 300\u2009K, the excitation temperature of glycolaldehyde derived in IRAS 16293 (J\u00f8rgensen et al. 2012). The (CH2OH)2\/CH2OHCHO abundance ratio in IRAS2A is estimated at 5.5 \u00b1 1.0 if we consider the column densities derived from the rotational diagrams. It is slightly lower (4.6), however, if we use the column density of ethylene glycol of 1.1 \u00d7 1016 cm-2 that does not overproduce the peak intensities of a few lines (see Fig. 3). The (CH2OH)2\/CH2OHCHO abundance ratio consequently is a factor ~5 higher than in the Class 0 protostar IRAS 16293. It is also higher than in the other star-forming regions (see Table 3), but similar to the lower limits derived in comets (\u22733\u20136). This indicates that the glycolaldehyde chemistry may in general vary among hot corinos. It is possible that like IRAS2A, other very young low-mass protostars show high (CH2OH)2\/CH2OHCHO abundance ratios, in agreement with the cometary values. The CH3OCHO\/CH2OHCHO column density ratio found in IRAS2A (~20) ranges between the values derived in the molecular clouds from the Galactic center (~3.3\u20135.2) and the high-mass star-forming regions (~40\u201352). A lower limit of 2 was derived for comet Hale-Bopp. ","Citation Text":["J\u00f8rgensen et al. (2012)"],"Functions Text":["Indeed, the assignment in","was based on only one line of the gGg\u2032 conformer of ethylene glycol about 200 cm-1","above the lowest-energy aGg\u2032 conformer \u2013 and thus tentative."],"Functions Label":["Differences","Differences","Differences"],"Citation Start End":[[315,338]],"Functions Start End":[[289,314],[339,421],[455,515]]}
{"Identifier":"2021ApJ...914...52F__Dere_et_al._1997_Instance_1","Paragraph":"In Equations (24) and (25) the cumulative effect of all ions with Zi \u2265 2 is included in the term \u03b6(T, f), where a (rather weak) dependence on frequency f appears due to the Gaunt factor. The \u03b6(T, f) function depends both on the elemental abundances and on each element\u2019s ionization state, which are different in equilibrium and nonequilibrium conditions. In this work, we will provide the \u03b6(T, f) function calculated under the assumption of ionization equilibrium using the CHIANTI spectral code (Dere et al. 1997, 2019), as detailed below. In this framework it is straightforward to users to replace the CHIANTI values for the equilibrium ion fractions by their own values calculated under nonequilibrium conditions, depending on the physical process they model. We carried out the calculations for the plasma element abundances typical of the solar corona (Feldman 1992), as well as for two abundance models obtained for the solar photosphere (Caffau et al. 2011; Scott et al. 2015). The main difference between these models is that in the solar corona all elements with low first ionization potential (FIP) are overabundant by a factor of about four over their photospheric abundances. However, this overabundance factor is somewhat uncertain\u2014while it has been extensively used in the field, its value has been suggested to vary with time (Widing & Feldman 2001) and also within the same active region; also, the absolute correction of the abundances has been disputed, as some studies suggest that the FIP effect affects both low-FIP and high-FIP ions (Schmelz et al. 2012). Furthermore, solar abundances can be different from those of other stars, both in the photosphere and in the corona, even to the point that some stars exhibit an inverse FIP effect (Laming 2015). So, the software we implemented to carry out the calculations allows the user to interactively select an abundance data set that is best suited for the star or the solar region under consideration.","Citation Text":["Dere et al. 1997"],"Functions Text":["In this work, we will provide the \u03b6(T, f) function calculated under the assumption of ionization equilibrium using the CHIANTI spectral code","as detailed below. In this framework it is straightforward to users to replace the CHIANTI values for the equilibrium ion fractions by their own values calculated under nonequilibrium conditions, depending on the physical process they model."],"Functions Label":["Uses","Uses"],"Citation Start End":[[497,513]],"Functions Start End":[[355,495],[522,763]]}
{"Identifier":"2020ApJ...903L..28S__Shivaei_et_al._2015_Instance_1","Paragraph":"The effectiveness of an IRX-\u03b2 relation depends on its validity for different types of galaxies across cosmic time. While the original MHC99 IRX-\u03b2 relation applies to majority of galaxies, large scatters around this relation have been observed. Theoretical studies have shown that the IRX-\u03b2 scatter may depend on the type of dust, gas metallicity, star formation history (SFH), dust-star geometry, and stellar population age (Popping et al. 2017; Safarzadeh et al. 2017; Narayanan et al. 2018; Schulz et al. 2020). Observations have shown that the IRX-\u03b2 relation varies with stellar mass (Bouwens et al. 2016, 2020; Reddy et al. 2018; Fudamoto et al. 2020), infrared (IR) luminosity (Buat et al. 2012; Casey et al. 2014), age (Siana et al. 2009; Reddy et al. 2010; Shivaei et al. 2015), redshift (Capak et al. 2015; Pannella et al. 2015), and intrinsic \u03b20 (\u03b2 for a dust-free system; Boquien et al. 2012; Reddy et al. 2018; Schulz et al. 2020). These variations are linked to the diversity of galaxies\u2019 attenuation curves, seen in previous studies (e.g., Kriek & Conroy 2013; Scoville et al. 2015; Battisti et al. 2020). The attenuation curve variations stem from two main sources: (a) different geometrical distributions of dust with respect to stars (and different dust optical depths), and (b) different dust grain properties, which affect the shape of the underlying extinction curve irrespective of the dust-star geometry (see the review by Salim & Narayanan 2020). While scenario (a) is extensively studied by comparing the attenuation curves of galaxies with different dust optical depths (references above), scenario (b) is less explored. Dust compositions are related to gas-phase element depletions and abundances (Jenkins 2009). Following this relation, Shivaei et al. (2020, hereafter S20) studied scenario (b) by deriving the attenuation curve of z \u223c 2 galaxies in two different metallicity (\n\n\n\n\n\n) bins, while the dust optical depth distributions were kept the same. They found a steep SMC-like curve for galaxies with \n\n\n\n\n\n, and a shallower Calzetti et al. (2000, hereafter C00)-like curve at \n\n\n\n\n\n.","Citation Text":["Shivaei et al. 2015"],"Functions Text":["Observations have shown that the IRX-\u03b2 relation varies with","age"],"Functions Label":["Background","Background"],"Citation Start End":[[764,783]],"Functions Start End":[[514,573],[721,724]]}
{"Identifier":"2018MNRAS.479.3254V___2000_Instance_1","Paragraph":"The lifetime of molecular clouds (MCs) remains an active research topic in the study of the interstellar medium and star formation, and most recent studies, both observational and theoretical, place this lifetime at a few times 107 yr for clouds in the 105\u2013106M\u2299 mass range (e.g. Blitz & Shu 1980; Kawamura et al. 2009; Zamora-Avil\u00e9s, V\u00e1zquez-Semadeni & Col\u00edn 2012; Zamora-Avil\u00e9s & V\u00e1zquez-Semadeni 2014; Lee, Miville-Desch\u00eanes & Murray 2016). In addition, several observational studies have suggested that the star formation rate (SFR) of the clouds appears to increase over their lifetimes. For example, studies of young clusters embedded in moderate-mass MCs (\u223c104M\u2299) (e.g. Palla & Stahler 1999, 2000; Da Rio et al. 2010) have shown that their age histograms contain a large majority of young (1\u20132 Myr) objects, but also a tail of older (up to several Myr) ones suggesting an accelerating star formation activity, sometimes followed by a subsequent decline (see also Povich et al. 2016; Schneider et al. 2018). In addition, Kawamura et al. (2009) reported a clear evolutionary process over the lifetime of giant molecular clouds (GMCs; of masses \u223c105\u2013106M\u2299) in the Large Magellanic Cloud, evidenced by the increasing number of massive stars across the sequence of GMC \u2018classes\u2019 proposed by those authors. Finally, on the basis of the large scatter in the observed star formation efficiency (SFE) in Milky Way GMCs, Lee et al. (2016) have concluded that the SFR in those clouds must also be time variable. Numerical simulations of MC formation and evolution also exhibit time varying, increasing SFRs during their early stages (e.g. V\u00e1zquez-Semadeni et al. 2007; Hartmann, Ballesteros-Paredes & Heitsch 2012). Also, in the presence of stellar feedback, at late times the SFRs reach a maximum and begin to decrease again (e.g. V\u00e1zquez-Semadeni et al. 2010; Col\u00edn, V\u00e1zquez-Semadeni & G\u00f3mez 2013). V\u00e1zquez-Semadeni, Gonz\u00e1lez-Samaniego & Col\u00edn (2017) have recently shown that the simulations of Col\u00edn et al. (2013) in fact produce stellar age histograms highly resemblant of the observed ones (Palla & Stahler 1999, 2000; Da Rio et al. 2010), and reproduce observed radial age gradients in clusters (Getman et al. 2014) as well as bottom-heavy stellar initial mass functions (IMFs) in scattered regions of massive star formation (Povich et al. 2016).","Citation Text":["Palla & Stahler","2000"],"Functions Text":["In addition, several observational studies have suggested that the star formation rate (SFR) of the clouds appears to increase over their lifetimes. For example, studies of young clusters embedded in moderate-mass MCs (\u223c104M\u2299) (e.g.","have shown that their age histograms contain a large majority of young (1\u20132 Myr) objects, but also a tail of older (up to several Myr) ones suggesting an accelerating star formation activity, sometimes followed by a subsequent decline"],"Functions Label":["Background","Background"],"Citation Start End":[[677,692],[699,703]],"Functions Start End":[[444,676],[725,959]]}
{"Identifier":"2022MNRAS.517.1058H__Navarro-Gonz\u00e1lez_et_al._1989_Instance_1","Paragraph":"It is likely that some of these products are also formed in our experiments, but except for CO2, HNCO, and NCO\u2212, they cannot be unequivocally identified with IR spectroscopy alone. Some of them have the same functional groups as urea and their IR spectra present absorptions in the same spectral regions, albeit with different band shapes and intensities. They could contribute somehow to the changes, both in shape and intensity, observed in the absorption bands during processing and thus introduce uncertainty, but by analogy with the literature results just mentioned (Renoult et al. 1969; Navarro-Gonz\u00e1lez et al. 1989; Duvernay et al. 2005; Poch et al. 2014) we expect that the prevalent dissociation channels will lead to small molecules that will not interfere much with the absorptions of urea. It is worth noting that urea fits in an isoelectronic sequence with carbonic acid and acetone: O = C(OH)2, O = C(NH2)2, and O = C(CH3)2. Under energetic processing, carbonic acid ice has been shown to decompose into H2O and CO2 (Peeters et al. 2010), and acetone ice into CH4 and H2C2O (ketene) among other products (Hudson 2018). In a similar way, urea, the middle term in the sequence, should decompose into NH3 and HNCO. Acid base reactions could transform HNCO and NH3 into OCN\u2212 and NH4+. As indicated repeatedly in the text, HNCO and OCN\u2212 were found in our measurements, but we did not find NH3 (expected peak at 1070 cm\u22121, A = 1.4 \u00d7 10\u221217 cm molecule\u22121). The presence of NH4+, which has been reported by Poch et al (2014) in the photolysis of crystalline urea, is not evident in our measurements, but cannot be completely excluded. The \u03bd4 bending mode of ammonium (A = 4.1 \u00d7 10\u221217 cm molecule\u22121; van Broekuizen et al. 2004) appears at 1440\u20131480 cm\u22121, but its exact location and bandwidth are strongly dependent on the ice environment (see discussion in Mat\u00e9 et al. 2009; G\u00e1lvez et al. 2010). In our case, the corresponding ammonium band, which would be the most obvious counterion for OCN\u2212, could be masked by the \u03bdCN band of urea. Compounds of very high molecular weight do not seem to form in appreciable amounts in our experiments, since all products evaporate without leaving any residue upon warming of the substrate to 300 K during a few minutes.","Citation Text":["Navarro-Gonz\u00e1lez et al. 1989"],"Functions Text":["It is likely that some of these products are also formed in our experiments, but except for CO2, HNCO, and NCO\u2212, they cannot be unequivocally identified with IR spectroscopy alone. Some of them have the same functional groups as urea and their IR spectra present absorptions in the same spectral regions, albeit with different band shapes and intensities. They could contribute somehow to the changes, both in shape and intensity, observed in the absorption bands during processing and thus introduce uncertainty, but by analogy with the literature results just mentioned","we expect that the prevalent dissociation channels will lead to small molecules that will not interfere much with the absorptions of urea."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[594,622]],"Functions Start End":[[0,571],[664,802]]}
{"Identifier":"2018AandA...617A..86L__Li_et_al._2015b_Instance_1","Paragraph":"The IRIS spectra measure the flare in a \u201csit-and-stare\u201d mode with a roll angle of 45\u2218. The spectral scale is \u223c25.6 m\u00c5 per pixel in the far-ultraviolet (FUV) wavelengths. The IRIS slit crosses the flaring loop and one ribbon (Fig. 1). Two red bars enclose the flaring loop region used to study the quasi-periodic oscillations in this work. IRIS spectrum was pre-processed with the SSW routines of \u201ciris_orbitval_corr_l2.pro\u201d (Tian et al. 2014; Cheng et al. 2015) and \u201ciris_prep_despike.pro\u201d (De Pontieu et al. 2014). To improve the signal-to-noise ratio, we apply a running average over five pixels to the IRIS spectra along the slit (Tian et al. 2012, 2016). We also manually perform the absolute wavelength calibration using a relatively strong neutral line, O\u202fI 1355.60 \u00c5 (see De Pontieu et al. 2014; Tian et al. 2015; Tian 2017). IRIS observations show that Fe\u202fXXI 1354.08 \u00c5 is a hot (\u223c11 MK) and broad emission line and is always blended with many narrow chromospheric lines at the flaring ribbons (Li et al. 2015b, 2016b; Tian et al. 2015; Young et al. 2015; Brosius et al. 2016; Polito et al. 2016). However, the Fe\u202fXXI 1354.08 \u00c5 line is much stronger than those blended emission lines at the flaring loops (Tian et al. 2016). Figure 2a gives the time evolution of the line profiles of Fe\u202fXXI 1354.08 \u00c5, averaged over the slit positions between \u223c18.3\u2033 and 21.6\u2033. Figure 2 panels b\u2212f show the spectral line profiles at the time indicated by the yellow lines in panel a. We can see that only the cool line of C\u202fI 1354.29 \u00c5 is blended with the hot line of Fe\u202fXXI 1354.08 \u00c5, but its contribution is negligible. Therefore, double Gaussian functions superimposed on a linear background are used to fit the IRIS spectra at \u201cO\u202f\u202fI\u201d window (Tian et al. 2016). Next, we can extract the hot line of Fe\u202fXXI 1354.08 \u00c5, as shown by the turquoise profile. The purple profile is the cool line of C\u202fI 1354.29 \u00c5. Two orange peaks represent the cool lines of O\u202fI 1354.60 \u00c5 and C\u202fI 1354.84 \u00c5 (Tian 2017), which are far away from the flaring line of Fe\u202fXXI 1354.08 \u00c5. Finally, the line properties of Fe\u202fXXI 1354.08 \u00c5 are extracted from the fitting results, that is, Doppler velocity, peak intensity, and line width (Li et al. 2016b; Tian et al. 2016; Tian & Chen 2018).","Citation Text":["Li et al. 2015b"],"Functions Text":["IRIS observations show that Fe\u202fXXI 1354.08 \u00c5 is a hot (\u223c11 MK) and broad emission line and is always blended with many narrow chromospheric lines at the flaring ribbons"],"Functions Label":["Uses"],"Citation Start End":[[1003,1018]],"Functions Start End":[[833,1001]]}
{"Identifier":"2018MNRAS.477.1664G__Behroozi,_Wechsler_&_Wu_2013_Instance_1","Paragraph":"We test our calibration procedure on the buzzard-v1.1 simulation, a mock DES Y1 survey created from a set of dark-matter-only simulations. The simulation and creation of the mock survey data are detailed in DeRose et al. (in preparation), Wechsler et al. (in preparation), and MacCrann et al. (2018), so we provide only a brief summary of both. buzzard-v1.1 is constructed from a set of three N-body simulations run using l-gadget2, a version of gadget2 (Springel 2005) modified for memory efficiency. The simulation boxes ranged from 1 to 4\u2009Gpc\u2009h\u22121. Light cones from each box were constructed on the fly. Haloes were identified using rockstar (Behroozi, Wechsler & Wu 2013), and galaxies were added to the simulations using the Adding Density Dependent GAlaxies to Light-cone Simulations (addgals) algorithm (Wechsler et al., in preparation). addgals uses the large-scale dark matter density field to place galaxies in the simulation based on the probabilistic relation between density and galaxy magnitude. The latter is calibrated from subhalo abundance matching in high-resolution N-body simulations. SEDs are assigned to the galaxies from a training set of spectroscopic data from SDSS Data Release 7 (DR7; Cooper et al. 2011) based on local environmental density. The SEDs are integrated in the DES pass bands to generate griz magnitudes. Galaxy sizes and ellipticities are drawn from distributions fit to deep SuprimeCam $i^{{\\prime }}$-band data. Galaxies are added to the simulation to the projected apparent magnitude limit of DES Year 5 (Y5) data out to redshift z = 2. The galaxy positions, shapes and magnitudes are then lensed using the multiple-plane ray-tracing code Curved-sky grAvitational Lensing for Cosmological Light conE simulatioNS (calclens; Becker 2013). Finally, the catalogue is cut to the DES Y1 footprint with RA > 0 using the footprint and bad region masking including bright stars, regions of high extinction, etc., used in the actual Y1 data, and photometric errors are added using the DES Y1 depth map (Rykoff, Rozo & Keisler 2015). This yields a total masked area of 1108.13\u2009deg2, 12 million WL source galaxies, and 102\u2009120 galaxies in the higher luminosity redMaGiC sample used in this paper, as will be discussed in Sections 3.2 and 3.3.","Citation Text":["Behroozi, Wechsler & Wu 2013"],"Functions Text":["Haloes were identified using rockstar"],"Functions Label":["Uses"],"Citation Start End":[[645,673]],"Functions Start End":[[606,643]]}
{"Identifier":"2016AandA...593A..22R__Shibuya_et_al._2015_Instance_1","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"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[673,692]],"Functions Start End":[[119,281]]}
{"Identifier":"2017AandA...602A..75R__Kaneko_&_Yokoyama_(2015)_Instance_1","Paragraph":"We can see how this is consistent with the development of small scales via phase mixing by introducing local wavenumbers for the variation with \u03b1 and \u03b2, (3)\\begin{eqnarray} \\xi \\propto \\exp {\\rm i}\\left[ \\int \\kappa_\\alpha {\\rm d}\\alpha + \\int \\kappa_\\beta {\\rm d}\\beta \\right]\\!. \\label{number3} \\end{eqnarray}\u03be\u221dexpi\u222b\u03ba\u03b1d\u03b1+\u222b\u03ba\u03b2d\u03b2.Here \u03ba\u03b1 and \u03ba\u03b2 are the wavenumbers in \u03b1 and \u03b2 and have units that are the inverse of the units of their respective coordinates. These wavenumbers should be distinguished from the perpendicular components of the usual wave vector k, which has units of 1\/length. The different wavenumbers may be related through the scale factors (h) that relate elemental coordinate increments to physical distances: dr = e\u03b1h\u03b1d\u03b1 + e\u03b2h\u03b2d\u03b2 + e\u03b3h\u03b3d\u03b3, where e\u03b1 is a unit vector in the \u03b1 direction, etc. In this notation \u2207\u22a5 = (e\u03b1\/h\u03b1)\u2202\/\u2202\u03b1 + (e\u03b2\/h\u03b2)\u2202\/\u2202\u03b2, and noting that \u2207\u22a5\u03be \u2248 ik\u22a5\u03be, Eq. (3) yields \\begin{eqnarray} \\bdelp\\xi &amp;=&amp; {\\rm i} \\vec{k}_\\perp \\xi = {\\rm i} \\left( \\frac{\\vec{e}_\\alpha}{h_\\alpha}\\frac{ \\partial}{\\partial\\alpha} +\\frac{\\vec{e}_\\beta}{h_ \\beta}\\frac{ \\partial}{\\partial \\beta}\\right)\\xi \\\\ &amp;\\equiv&amp; {\\rm i}\\left( \\frac{\\vec{e}_\\alpha}{h_\\alpha}\\kappa_\\alpha +\\frac{\\vec{e}_\\beta}{h_ \\beta}\\kappa_\\beta\\right)\\xi. \\label{number4} \\end{eqnarray}\u2207\u22a5\u03be=ik\u22a5\u03be=ie\u03b1h\u03b1\u2202\u2202\u03b1+e\u03b2h\u03b2\u2202\u2202\u03b2\u03be\u2261Equating components of the second and fourth expressions in Eq. (5) gives the expected relations between the various wavenumbers, (6)\\begin{eqnarray} k_\\alpha = \\kappa_\\alpha \/h_\\alpha, \\qquad k_\\beta = \\kappa_ \\beta \/h_ \\beta. \\label{number5} \\end{eqnarray}k\u03b1=\u03ba\u03b1\/h\u03b1,\u2001k\u03b2=\u03ba\u03b2\/h\u03b2.Equations (2) and (5) give a direct and elegant expression for the perpendicular wave vector as (7)\\begin{eqnarray} \\vec{k}_\\perp \\approx -(\\bdel \\omega_{\\rm c}) t, \\label{number6} \\end{eqnarray}k\u22a5\u2248\u2212(\u2207\u03c9c)t,which is a generalisation to three dimensions of the results of Mann et al. (1995), (Wright et al. 1999) and Kaneko & Yokoyama (2015) for lower dimensional systems, which developed phase mixing in only one perpendicular coordinate. The above expression allows phase mixing in both perpendicular directions, giving physical phase mixing lengths (or wavelengths) in the \u03b1 and \u03b2 directions of (8)\\begin{eqnarray} L_{{\\rm ph}\\alpha} = \\frac{2\\pi}{|k_\\alpha |} \\equiv \\frac{2\\pi h_\\alpha}{|\\partial \\omega_{\\rm c}\/\\partial\\alpha |t}, \\qquad L_{{\\rm ph}\\beta} = \\frac{2\\pi}{|k_ \\beta |} \\equiv \\frac{2\\pi h_\\beta}{|\\partial \\omega_{\\rm c}\/\\partial \\beta |t}\\cdot \\label{number7} \\end{eqnarray}Lph\u03b1=2\u03c0|k\u03b1|\u22612\u03c0h\u03b1|\u2202\u03c9c\/\u2202\u03b1|t,\u2001Lph\u03b2=2\u03c0|k\u03b2|\u22612\u03c0h\u03b2|\u2202\u03c9c\/\u2202\u03b2|t\u00b7If the phase mixing lengths are expressed in the same units as \u03b1 and \u03b2, rather than physical length as in Eq. (8), slightly simpler expressions are found, i.e. (9)\\begin{eqnarray} \\ell_{{\\rm ph}\\alpha} = \\frac{2\\pi}{|\\kappa_\\alpha |} \\equiv \\frac{2\\pi}{|\\partial \\omega_{\\rm c}\/\\partial\\alpha |t}, \\qquad \\ell_{{\\rm ph}\\beta} = \\frac{2\\pi}{|\\kappa_ \\beta |} \\equiv \\frac{2\\pi}{|\\partial \\omega_{\\rm c}\/\\partial \\beta |t}\\cdot \\label{number8} \\end{eqnarray}\u2113ph\u03b1=2\u03c0|\u03ba\u03b1|\u22612\u03c0|\u2202\u03c9c\/\u2202\u03b1|t,\u2001\u2113ph\u03b2=2\u03c0|\u03ba\u03b2|\u22612\u03c0|\u2202\u03c9c\/\u2202\u03b2|t\u00b7The development of the phase mixing length can be pictured simply as the tendency for each field line to oscillate with its own natural frequency. Even if all the field lines start to oscillate with the same phase, they soon drift out of phase with one another as time passes. Not only does the phase mixing process generate perpendicular scales, but points of constant phase can be seen to move across field lines. This phase motion has been seen in magnetospheric data of Alfv\u00e9n waves (see the review by Wright & Mann 2006) and the simulations of coronal oscillations by Kaneko & Yokoyama (2015). These studies note that the direction of motion is related to the spatial variation of \u03c9c. The results of these papers for the perpendicular phase velocity in physical space generalise to Vph = \u03c9c\/k\u22a5, giving the components (10)\\begin{eqnarray} V_{{\\rm ph}\\alpha} = \\frac{-\\omega_{\\rm c} h_\\alpha}{(\\partial \\omega_{\\rm c}\/\\partial\\alpha )t}, \\qquad V_{{\\rm ph}\\beta} = \\frac{-\\omega_{\\rm c} h_ \\beta}{(\\partial \\omega_{\\rm c}\/\\partial \\beta )t}, \\qquad \\label{number9} \\end{eqnarray}Vph\u03b1=\u2212\u03c9ch\u03b1(\u2202\u03c9c\/\u2202\u03b1)t,\u2001Vph\u03b2=\u2212\u03c9ch\u03b2(\u2202\u03c9c\/\u2202\u03b2)t,\u2001If the excitation occurred at a time ti, the subsequent properties are found by replacing t with t\u2212ti in the above formulae. ","Citation Text":["Kaneko & Yokoyama (2015)"],"Functions Text":["\\label{number5} \\end{eqnarray}k\u03b1=\u03ba\u03b1\/h\u03b1,\u2001k\u03b2=\u03ba\u03b2\/h\u03b2.Equations (2) and (5) give a direct and elegant expression for the perpendicular wave vector as (7)\\begin{eqnarray} \\vec{k}_\\perp \\approx -(\\bdel \\omega_{\\rm c}) t, \\label{number6} \\end{eqnarray}k\u22a5\u2248\u2212(\u2207\u03c9c)t,which is a generalisation to three dimensions of the results of Mann et al. (1995),","and","for lower dimensional systems, which developed phase mixing in only one perpendicular coordinate."],"Functions Label":["Uses","Uses","Uses"],"Citation Start End":[[1905,1929]],"Functions Start End":[[1541,1879],[1901,1904],[1930,2027]]}
{"Identifier":"2018MNRAS.480.2881M__Berkley,_Kazanas_&_Ozik_2000_Instance_1","Paragraph":"There has additionally been the concern that the optical light curves do not look as expected if they arise from reprocessing of X-ray emission from a small central corona, e.g. of size similar to that which we measure from microlensing observations, i.e. $\\rm{\\, \\stackrel{\\lt }{_\\sim }\\,}10 R_{G}$ (Dai et al. 2010; Mosquera et al. 2013), or from X-ray low\/high energy reverberation, i.e. \u223c4Rg (Cackett et al. 2014; Emmanoulopoulos et al. 2014). The observed optical light curves are smoother than expected and an insufficient fraction of the X-ray emission hits the disc to power the optical variability (e.g. Berkley, Kazanas & Ozik 2000; Ar\u00e9valo et al. 2008). Larger coronal sizes are required. Gaskell (2008) also notes the energetics problem and proposes variations originating independently in different parts of the disc. However, although such a model is useful for explaining the uncorrelated variations between bands which are sometimes seen, it cannot explain the well correlated multiwavelength variations seen in the Swift observations. Gardner & Done (2017) proposed an alternative model in which the X-ray emission does not directly impact on the outer disc but mainly heats up the very inner edge of the disc, which then inflates and re-radiates at hard UV wavelengths on to the outer disc. In this model there should be an additional lag between the X-ray and UVW2 emission, over and above that expected from an extrapolation of the longer wavelength lags down to the X-ray waveband. This additional lag would correspond to the thermal time-scale for the incident X-ray heating to pass through the inner disc to the re-radiation surface. Gardner & Done (2017) note the existence of such a lag when the unfiltered X-ray and UVW2 observations of NGC 5548 are compared. However, McHardy et al. (2014) do not see any additional lag in NGC 5548 if those light curves are filtered to remove variations on time-scales longer than 20 d. In NGC 4151, Edelson et al. (2017) see a very large excess lag between the X-ray and UVW2 bands. Unlike in NGC 5548, the excess lag in NGC 4151 is strongly energy dependent, with the highest energy X-rays having the largest lag. NGC 4151 is the most absorbed of the few AGN whose lags have been well studied so far and so the energy dependence may be a function of scattering in the absorbing medium.","Citation Text":["Berkley, Kazanas & Ozik 2000"],"Functions Text":["The observed optical light curves are smoother than expected and an insufficient fraction of the X-ray emission hits the disc to power the optical variability (e.g."],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[613,641]],"Functions Start End":[[448,612]]}
{"Identifier":"2022MNRAS.511.4333K__Gibbons_et_al._2014_Instance_1","Paragraph":"The first structures we consider are cosmic voids. Cosmic voids are defined as large underdense regions of the cosmic web, they are the largest structures in the Universe and make up most of its volume (Cautun et al. 2014; Falck & Neyrinck 2015). Historically, their existence was one of the earliest predictions of the concordance cosmological model (Hausman, Olson & Roth 1983), and their observational detection goes back to roughly 40 yr ago (Gregory, Thompson & Tifft 1978; Kirshner et al. 1981). Voids are in particular extremely underdense near their centres, and their spherically averaged density profile shows a characteristic shape (Colberg et al. 2005; Ricciardelli, Quilis & Planelles 2013; Hamaus, Sutter & Wandelt 2014a; Nadathur et al. 2014b; Ricciardelli, Quilis & Varela 2014). Recently, cosmic voids are becoming a promising cosmological probes: first they could represent a population of statistically ideal spheres with a homogeneous distribution at different redshifts which size evolution could be used to probe the expansion of the Universe using Alcock & Paczynski tests (Alcock & Paczynski 1979; Lavaux & Wandelt 2012; Sutter et al. 2012; Sutter et al. 2014b; Hamaus et al. 2015; Hamaus et al. 2016; Mao et al. 2017; Hamaus et al. 2022). Moreover, due to their low density, voids are naturally sensitive to dark energy and thus the interest to use them as probe of alternative Dark Energy models and modified gravity scenarios is increasing (Odrzywo\u0142ek 2009; Lavaux & Wandelt 2010; D\u2019Amico et al. 2011; Li 2011; Bos et al. 2012; Clampitt, Cai & Li 2013; Gibbons et al. 2014; Barreira et al. 2015; Cai, Padilla & Li 2015; Pisani et al. 2015; Zivick et al. 2015; Pollina et al. 2016; Baldi & Villaescusa-Navarro 2018), as well as the possibility of using them to put constraints on neutrinos masses (Massara et al. 2015; Kreisch et al. 2019; Contarini et al. 2021). Their imprint on the observed Cosmic Microwave Background (CMB) is also becoming an encouraging new cosmological probe, either through their Integrated Sachs-Wolfe (ISW) imprint (Baccigalupi, Amendola & Occhionero 1997; Baccigalupi 1999; Granett, Neyrinck & Szapudi 2008; Cai et al. 2014; Granett, Kov\u00e1cs & Hawken 2015; Hotchkiss et al. 2015; Ade et al. 2016; Nadathur & Crittenden 2016; Kov\u00e1cs et al. 2017; Kov\u00e1cs et al. 2019; Hang et al. 2021), or their lensing imprint (Cai et al. 2017; Raghunathan et al. 2020; Vielzeuf et al. 2021). Furthermore, the observed cold spot of the CMB could be explained as the imprint of the ISW sourced by very large voids along the line of sight (Rees, Sciama & Stobbs 1968; Kovac et al. 2013; Finelli et al. 2014; Nadathur et al. 2014a). Moreover, some works such as Jamieson & Loverde (2019) studied the properties of the voids via the separate universe technique. Finally, some studies tried to link high redshift intergalactic voids in the transmitted Lyman-\u03b1 flux to the gas density (Viel, Colberg & Kim 2008). Because they are almost empty regions, their evolution during cosmic history is at most weakly non-linear and their properties could possibly be impacted by the primordial density fields from which they formed. This fact motivates us to investigate the effects of baryon-CDM relative perturbations on these objects and their statistics.","Citation Text":["Gibbons et al. 2014"],"Functions Text":["Moreover, due to their low density, voids are naturally sensitive to dark energy and thus the interest to use them as probe of alternative Dark Energy models and modified gravity scenarios is increasing"],"Functions Label":["Motivation"],"Citation Start End":[[1580,1599]],"Functions Start End":[[1264,1466]]}
{"Identifier":"2021AandA...656A.137G__White_et_al._1997_Instance_1","Paragraph":"To determine the median spectral index and the variation within the sample we perform a bootstrap analysis, randomly drawing from our sample (with replacement) and median stacking both the LoTSS and FIRST cutouts independently to obtain a median peak pixel flux density measurement in each survey. This is then repeated 10 000 times to produce a distribution that also incorporates uncertainties due to sample variation. The 16th and 84th percentile of the resulting peak flux density distributions, together with the systematic flux scale density uncertainty of 5% (White et al. 1997) and 10% (Shimwell et al., in prep.) for the FIRST and LoTSS survey, respectively, are used to determine the error on the spectral index. This stacking procedure is done for all 93 quasars that overlap with FIRST (red) and for only the 35 sources detected by LoTSS with S\/N\u2004>\u20042 that overlap with FIRST (grey), resulting in spectral indices of \n\n\n\n\u2212\n0\n.\n\n29\n\n\u2212\n0.09\n\n\n+\n0.10\n\n\n\n\n$ -0.29^{+0.10}_{-0.09} $\n\n\n and \n\n\n\n\u2212\n0\n.\n\n24\n\n\u2212\n0.22\n\n\n+\n0.09\n\n\n\n\n$ -0.24^{+0.09}_{-0.22} $\n\n\n, respectively. Six of the FIRST detected quasars have spectral indices higher than these stacked values, which is as expected because these flat or positive spectral index quasars are naturally selected because of the FIRST flux limit of \u223c144 \u03bcJy, which is shallower than the LOFAR flux limit assuming a negative spectral index. In general, LOFAR detected quasars are brighter at 144 MHz than the non-detected ones, therefore one might expect the stacked S\/N\u2004>\u20042 sample to result in a steeper spectral index than when stacking all quasars. However, given the uncertainties on the median spectral index there is no significant deviation found between the two. In this work, we therefore assume a spectral index of \u22120.29\u00b10.10 for the quasars not detected by FIRST when converting radio fluxes to luminosities. This spectral index is in line with the previous work of G\u00fcrkan et al. (2019), who found a median spectral index of \u22120.26\u2005\u00b1\u20050.02 for a large sample of optically selected quasars at z\u2004\u2272\u20043.","Citation Text":["White et al. 1997"],"Functions Text":["The 16th and 84th percentile of the resulting peak flux density distributions, together with the systematic flux scale density uncertainty of 5%","for the FIRST and LoTSS survey, respectively, are used to determine the error on the spectral index."],"Functions Label":["Uses","Uses"],"Citation Start End":[[567,584]],"Functions Start End":[[421,565],[622,722]]}
{"Identifier":"2018MNRAS.473.3810Y__Mitrushchenkov_et_al._2017_Instance_1","Paragraph":"The lack of data on inelastic processes due to collisions with neutral hydrogen atoms has been a major limitation on modelling of F-, G- and K-star spectra in statistical equilibrium, and thus to reliably proceeding beyond the assumption of local thermodynamic equilibrium (LTE) in analysis of stellar spectra and the determination of elemental abundances. This problem has been well documented, e.g. see Lambert (1993); Barklem (2016a) and references therein. Significant progress has been made in recent times through detailed full-quantum scattering calculations, based on quantum chemical data, for the cases of simple atoms such as Li, Na, Mg and Ca (Belyaev & Barklem 2003; Barklem, Belyaev & Asplund 2003; Belyaev et al. 2010; Barklem et al. 2010; Belyaev et al. 2012; Barklem et al. 2012; Mitrushchenkov et al. 2017). These calculations have demonstrated the importance of the ionic-covalent curve crossing mechanism leading naturally to charge transfer processes (mutual neutralization and ion-pair production), in addition to excitation and de-excitation processes. The importance of this mechanism has allowed various simplified model approaches to be developed, which may be used in cases where suitable quantum chemistry data are not been available. In particular a semi-empirical model has been employed for Al, Si, Be and Ca (Belyaev 2013a,b; Belyaev, Yakovleva & Barklem 2014b; Yakovleva, Voronov & Belyaev 2016; Belyaev et al. 2016), and a theoretical model based on a two-electron asymptotic linear combinations of atomic orbitals (LCAO) approach, has also been employed for Ca (Barklem 2016b, 2017). Comparisons of the two methods show quite good agreement and reasonable agreement with the full quantum calculations is found, particularly for the most important processes with the largest rates (Barklem 2016b, 2017; Mashonkina, Sitnova & Belyaev 2017; Mitrushchenkov et al. 2017). Thus, the model approaches provide a useful route for obtaining estimates of the rates for these processes for many elements of astrophysical interest.","Citation Text":["Mitrushchenkov et al. 2017"],"Functions Text":["Significant progress has been made in recent times through detailed full-quantum scattering calculations, based on quantum chemical data, for the cases of simple atoms such as Li, Na, Mg and Ca"],"Functions Label":["Background"],"Citation Start End":[[797,823]],"Functions Start End":[[461,654]]}
{"Identifier":"2019MNRAS.490.1714P__Dolag_et_al._2009_Instance_1","Paragraph":"The E-MOSAICS project (Pfeffer et al. 2018; Kruijssen et al. 2019a) is a suite of cosmological, hydrodynamical simulations of galaxy formation in the \u039b cold dark matter cosmogony that couples the MOSAICS model for star cluster formation and evolution (Kruijssen et al. 2011; Pfeffer et al. 2018) to the EAGLE model of galaxy formation and evolution (Crain et al. 2015; Schaye et al. 2015). The simulations are run with a highly modified version of the N-body, smoothed particle hydrodynamics code gadget3 (last described by Springel 2005). Bound galaxies (subhaloes) were identified within the simulations using the subfind algorithm (Springel et al. 2001; Dolag et al. 2009), in the same manner as in the EAGLE simulations (for details see Schaye et al. 2015). EAGLE includes subgrid routines describing radiative cooling (Wiersma, Schaye & Smith 2009a), star formation (Schaye & Dalla Vecchia 2008), stellar evolution, and mass-loss (Wiersma et al. 2009b), the seeding and growth of black holes (BHs) via gas accretion and BH\u2013BH mergers (Rosas-Guevara et al. 2015), and feedback associated with star formation and BH growth (Booth & Schaye 2009). As current cosmological simulations lack the resolution and physics necessary to compute the feedback efficiencies from first principles, the stellar and active galactic nuclei feedback parameters are calibrated such that the simulations of cosmologically representative volumes reproduce the galaxy stellar mass function, galaxy sizes, and BH masses at z \u2248 0. The EAGLE simulations successfully reproduce a range of galaxy properties, including the stellar masses (Furlong et al. 2015) and sizes (Furlong et al. 2017) of galaxies, their luminosities and colours (Trayford et al. 2015), their cold gas properties (Lagos et al. 2015, 2016; Bah\u00e9 et al. 2016; Marasco et al. 2016; Crain et al. 2017), and the properties of circumgalactic and intergalactic absorption systems (Rahmati et al. 2015, 2016; Oppenheimer et al. 2016; Turner et al. 2016, 2017). The simulations also largely reproduce the cosmic SFR density and relation between specific SFR and galaxy mass (Furlong et al. 2015). The simulations are therefore ideal for comparisons with observed galaxy populations.","Citation Text":["Dolag et al. 2009"],"Functions Text":["Bound galaxies (subhaloes) were identified within the simulations using the subfind algorithm","in the same manner as in the EAGLE simulations"],"Functions Label":["Uses","Similarities"],"Citation Start End":[[657,674]],"Functions Start End":[[540,633],[677,723]]}
{"Identifier":"2017ApJ...847...42D__Purcell_et_al._2011_Instance_1","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"],"Functions Text":["As a result, recent studies have shifted to using dark-matter halo masses as large as 1011 M\u2299","based on halo abundance matching arguments."],"Functions Label":["Background","Background"],"Citation Start End":[[639,658]],"Functions Start End":[[544,637],[679,722]]}
{"Identifier":"2018MNRAS.478.3890B__Heckman_et_al._2017_Instance_4","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)"],"Functions Text":["However, the results of","show strong offsets in the kinematics of the gas from the host galaxy (\u2248100\u2009km s\u22121;","whereas the COS-AGN sightlines do not (bottom panel of Fig. 6)"],"Functions Label":["Differences","Differences","Differences"],"Citation Start End":[[1090,1111]],"Functions Start End":[[1066,1089],[1112,1195],[1234,1296]]}
{"Identifier":"2016ApJ...825...47P__Krimm_et_al._2011a_Instance_1","Paragraph":"A tidal disruption event (TDE) is an astronomical phenomenon that occurs when a star gets too close to a supermassive black hole in the galaxy center and is disrupted by the tidal force of the black hole. Part of stellar material is bound and accreted by the central black hole, resulting in bright optical, UV, and soft X-ray emission (Rees 1988; Lodato et al. 2015 and references therein). There are a growing number of candidate TDEs being discovered in soft X-ray, ultraviolet, and optical surveys; see Komossa (2015) for a recent review. Recently, three unusual TDE candidates have been discovered by Swift, i.e., Swift J164449.3+573451, Swift J2058.4+0516, and Swift J1112.2-8238 (hereafter Sw J1644+57, Sw J2058+05, and Sw J1112-82 for short, respectively), which have very bright nonthermal hard X-ray and radio emissions (Bloom et al. 2011; Burrows et al. 2011; Krimm et al. 2011a, 2011b; Zauderer et al. 2011; Cenko et al. 2012; Brown et al. 2015). The luminous nonthermal X-ray and radio emissions are thought to be produced by relativistic jets (Bloom et al. 2011; Burrows et al. 2011; Levan et al. 2011; Zauderer et al. 2011; Cao & Wang 2012; Metzger et al. 2012; Wang et al. 2014; Liu et al. 2015). Sw J1644+57 shows a highly variable light curve in X-rays, as observed by the X-ray Telescope on board Swift. At redshift \n\n\n\n\n\n, the isotropic luminosity of the X-ray emission is as high as \n\n\n\n\n\n erg s\u22121. Sw J2058+05 exhibits a luminous, long-lived X-ray outburst with an isotropic peak luminosity of \n\n\n\n\n\n (at redshift \n\n\n\n\n\n). Its total isotropic energy (0.3\u201310 keV) on a timescale of the first 2 months amounts to 1054 erg. Sw J1112-82 was initially also discovered by Swift\/BAT (Burst Alert Telescope) in 2011 June as an unknown, long-lived (order of days) \u03b3-ray transient source. It exhibits a similar bright X-ray flare, and its position is consistent with the nucleus of a faint galaxy at \n\n\n\n\n\n (Brown et al. 2015). The peak X\/\u03b3-ray luminosity of Sw J1112-82 exceeds \n\n\n\n\n\n.","Citation Text":["Krimm et al. 2011a"],"Functions Text":["Recently, three unusual TDE candidates have been discovered by Swift, i.e., Swift J164449.3+573451, Swift J2058.4+0516, and Swift J1112.2-8238 (hereafter Sw J1644+57, Sw J2058+05, and Sw J1112-82 for short, respectively), which have very bright nonthermal hard X-ray and radio emissions"],"Functions Label":["Motivation"],"Citation Start End":[[871,889]],"Functions Start End":[[543,829]]}
{"Identifier":"2019ApJ...879...52S___2010_Instance_1","Paragraph":"Kennicutt & Evans (2012) present a compilation of disk-averaged SFR and gas mass surface densities whose values have been calculated in a uniform manner across different galaxy types (including normal disk galaxies and dusty starburst galaxies selected in the IR) and find a power-law index of n \u223c 1.4. However, this result may be an artifact of combining galaxies of different interaction states. For a sample of z \u223c 1\u20133 MS galaxies, Tacconi et al. (2013) find an index consistent with unity and only a slight offset between their high-redshift sample and a low-redshift sample with similar masses. However, SMGs and other ultra-\/luminous infrared galaxies (U\/LIRGs) are further offset above the correlation for star-forming disk galaxies even when similar CO-to-H2 conversion factors are used for all galaxy populations (see also Bigiel et al. 2008; Daddi et al. 2010b; Genzel et al. 2010, 2015; Tacconi et al. 2018). In analyses of the resolved star formation properties of nearby disks, a near-unity index for the Schmidt\u2013Kennicutt relation is also found in regimes where the molecular gas dominates the total gas mass surface density (\u03a3gas >9 M\u2299 pc\u22122; e.g., Bigiel et al. 2008, 2010; Schruba et al. 2011). The surface density version of the Schmidt\u2013Kennicutt relation has been evaluated within only eight high-redshift galaxies: SMM J14011+0252 at z = 2.56 (Sharon et al. 2013), EGS 13011166 at z = 1.53 (Genzel et al. 2013), HLS 0918 at z = 5.24 (Rawle et al. 2014), GN20 at z = 4.05 (Hodge et al. 2015), PLCK G244.8+54.9 at z = 3.00 (Ca\u00f1ameras et al. 2017), AzTEC-1 at z = 4.34 (Tadaki et al. 2018), and the two components of HATLAS J084933 at z = 2.41 (G\u00f3mez et al. 2018).9\n\n9\nFreundlich et al. (2013) and Sharda et al. (2018) also examine the Schmidt\u2013Kennicutt relation at z > 1, but they analyze individually resolved clumps within high-redshift galaxies rather than performing full pixel-by-pixel comparisons.\n These studies find a range of Schmidt\u2013Kennicutt relation indices (n = 1\u20132). It is particularly worth noting that Genzel et al. (2013) find that their measured index depends strongly on which spatially resolved extinction correction they apply to their H\u03b1 measurements.","Citation Text":["Bigiel et al.","2010"],"Functions Text":["In analyses of the resolved star formation properties of nearby disks, a near-unity index for the Schmidt\u2013Kennicutt relation is also found in regimes where the molecular gas dominates the total gas mass surface density (\u03a3gas >9 M\u2299 pc\u22122; e.g.,"],"Functions Label":["Similarities"],"Citation Start End":[[1163,1176],[1183,1187]],"Functions Start End":[[920,1162]]}
{"Identifier":"2022AandA...666A..83N__Chan_et_al._(2018)_Instance_1","Paragraph":"The detection criterion of having the full interval \u03b1\u2005\u00b1\u2005\u03c3\u03b1 inside the prior range [0.8,\u20061.2] (e.g., Ata et al. 2018; Chan et al. 2018) has commonly been employed in the literature. However, we can observe from these analyses that the distribution of the \u03b1 parameter in general seems to be more Gaussian than what we find in our analyses (left panels in Figs. 3 and 4). A possible reason for this is the lower redshift we considered here, where the contribution from nonlinear effects is stronger. However, results from Villaescusa-Navarro et al. (2017), who investigated the BAO detection from 21 cm signal for the SKA case for the redshift range 0.35\u2004\u2004z\u2004\u20043.05, also showed \u03b1 distributions with clear deviations from a perfect Gaussian (but note the smaller number of simulations employed there, namely, 100). As pointed out by Chan et al. (2018) and Abbott et al. (2022), a natural consequence from having a (approximate) Gaussian distribution is a reasonable concordance among the three different error measurements, that is, \u27e8\u03c3\u03b1\u27e9\u223c\u03c368\u2004\u223c\u2004\u03c3std. This is not our case, as can be seen from Tables 2 and 3 (as well as from Tables 4 and 5), indicating that \u27e8\u03c3\u03b1\u27e9 is not meaningful or representative of the error in the \u03b1 measurements for individual mock realizations (see also Figs. 3 and 4). Our results show that compared to \u27e8\u03c3\u03b1\u27e9, the errors given by \u03c368 are overestimated by \u223c18% to 33%, from the smallest to the highest z-bins when the ACF is used, but is underestimated by \u223c50% for APS. Moreover, comparing the expected \n\n\n\n\n\u03c3\n\u03b1\nm\n\n\n\n$ \\sigma^m_\\alpha $\n\n\n\n obtained fitting the mean C\u2113 and \u03c9(\u03b8) to the average \u27e8\u03c3\u03b1\u27e9, we find a reasonable agreement for intermediate and higher z-bins, but not for the lower z-bins, in particular, for the ACF estimator. In addition, although \u03c3std agrees better with \u03c368, we still find non-negligible differences among them, which confirms our \u03b1 distributions as non-Gaussian (a larger \u03c3std indicates non-Gaussian tails). These reasons motivated our choice of using the \u03b1 values instead of \u03b1\u2005\u00b1\u2005\u03c3\u03b1 ranges, which belong to the interval [0.8,\u20061.2] as a criterion for a BAO detection (defining the Nd fraction), as well as our choice of using the 68% spread of the \u03b1 distributions, \u03c368, as the representative error in our measurements.","Citation Text":["Chan et al. (2018)"],"Functions Text":["As pointed out by","and Abbott et al. (2022), a natural consequence from having a (approximate) Gaussian distribution is a reasonable concordance among the three different error measurements, that is, \u27e8\u03c3\u03b1\u27e9\u223c\u03c368\u2004\u223c\u2004\u03c3std."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[828,846]],"Functions Start End":[[810,827],[847,1044]]}
{"Identifier":"2019MNRAS.486.1608N__Cicone_et_al._2014_Instance_1","Paragraph":"The clearest examples of host\u2013AGN interaction are arguably found in nearby brightest cluster galaxies. The AGN in these systems have been shown to deposit vast amounts of energy into the surrounding intracluster medium via heating and (mega-parsec scale) jets both observationally and by means of modelling (e.g. Binney 2004; Scannapieco, Silk & Bouwens 2005; Gitti, Brighenti & McNamara 2012, for a review; English, Hardcastle & Krause 2016), which maintain the hot-gas reservoirs in these systems, prevent cooling flows, and thus suppress star formation (e.g. Binney & Tabor 1995; Li et al. 2015). In lower density environments (where the majority of galaxies live), the empirical picture is much less clear. Direct observational evidence for AGN feedback on galactic scales in such environments remains sparse. More specifically, while outflows have been observationally detected in a number of instances around AGN in various gas phases, most of these detections have been made in ultra-luminous infrared galaxies or some of the closest quasar\u2013host galaxies (e.g. Rupke, Veilleux & Sanders 2005b; Nyland et al. 2013; Harrison et al. 2014), with only a few examples where the AGN have been shown to couple with kpc-scale outflows that are capable of impacting star-formation on galactic scales (for a review Cicone et al. 2014; Harrison 2017). Thus, it is still unclear to what extent the general AGN population could drive kpc-scale galactic outflows capable of possibly limiting or quenching star-formation in the local Universe. Indeed, recent observational work has cast doubt on the ability of AGN to directly regulate star formation in the nearby Universe. For example, Schawinski et al. (2014) found that black hole accretion occurs preferentially in quenched galaxies that experienced a rapid decay of their star-formation rates. Based on stellar population analysis, such a time delay between the peak of star formation and the onset of AGN activity has been reported to be of at least several dynamical time-scales (e.g. Kaviraj et al. 2015; Shabala et al. 2017, for radio and e.g. Schawinski et al. 2007; Kaviraj 2009; Wild, Heckman & Charlot 2010 for optical). AGN may therefore not play a significant direct role in regulating their associated star formation episodes as they would not couple directly to the cold-gas reservoir (see, e.g. the models of Kaviraj et al. 2011). Using observations for cold-gas outflows, this has been directly confirmed for radio AGN, which are found to couple mainly to residual gas in galaxies where the gas reservoir is already significantly depleted (e.g. Sarzi et al. 2016). A fuller understanding of the role of AGN in regulating star formation demands a direct study of whether outflows of neutral material (which are ultimately responsible for quenching star formation) are more likely launched in AGN hosts. Most importantly, a quantitative statement about the putative role of AGN in influencing the evolution of their host galaxy requires a study that employs a complete sample of such AGN in the local Universe. Performing such an analysis is the purpose of this paper.","Citation Text":["Cicone et al. 2014"],"Functions Text":["Direct observational evidence for AGN feedback on galactic scales in such environments remains sparse. More specifically, while outflows have been observationally detected in a number of instances around AGN in various gas phases, most of these detections have been made in ultra-luminous infrared galaxies or some of the closest quasar\u2013host galaxies","with only a few examples where the AGN have been shown to couple with kpc-scale outflows that are capable of impacting star-formation on galactic scales (for a review","Thus, it is still unclear to what extent the general AGN population could drive kpc-scale galactic outflows capable of possibly limiting or quenching star-formation in the local Universe."],"Functions Label":["Background","Background","Motivation"],"Citation Start End":[[1311,1329]],"Functions Start End":[[711,1061],[1144,1310],[1347,1534]]}
{"Identifier":"2017ApJ...834..178Y__Tachihara_et_al._2007_Instance_1","Paragraph":"In order to investigate the gas kinematics at an early evolutionary stage and the formation of Keplerian disks, we conduct ALMA observations toward three candidate young protostars, Lupus 3 MMS, IRAS 15398\u22123559, and IRAS 16253\u22122429. They are selected from our SMA sample (Yen et al. 2015a). These three protostars all have relatively low protostellar masses (0.1 M\u2299), inferred from the infalling motions in their protostellar envelopes, and they do not show clear signs of a spin-up rotation on a 1000 au scale; i.e., no signatures of Keplerian disks are seen in our SMA observations (Yen et al. 2015a). Lupus 3 MMS is a Class 0 protostar with a bolometric luminosity (Lbol) of 0.41 L\u2299 and a bolometric temperature (Tbol) of 39 K in the Lupus 3 cloud at a distance of 200 pc (Tachihara et al. 2007; Comer\u00f3n 2008; Dunham et al. 2013). Our SMA results suggest that the protostellar mass in Lupus 3 MMS can be as low as 0.1 M\u2299 (Yen et al. 2015a). IRAS 15398\u22123559 is a Class 0\/I protostar with an Lbol of 1.2 L\u2299 and a Tbol of 61 K in the Lupus 1 cloud at a distance of 150 pc (Froebrich 2005; Comer\u00f3n 2008). Early single-dish observations of its CO outflow suggest that IRAS 15398\u22123559 is close to face on (van Kempen et al. 2009). Recent SMA and ALMA observations show that it is actually closer to edge on (Oya et al. 2014; Bjerkeli et al. 2016). With this new estimated inclination angle (\u223c70\u00b0), our SMA data suggest a low protostellar mass (0.1 M\u2299) and a low specific angular momentum in the protostellar envelope (\u223c1 \u00d7 10\u22124 km s\u22121 pc; Yen et al. 2015a). IRAS 16253\u22122429 is a Class 0 protostar with an Lbol of 0.24 L\u2299 and a Tbol of 36 K in the \u03c1 Ophiuchus star-forming region at a distance of 125 pc (Dunham et al. 2013). Both CARMA and our SMA results suggest that its protostellar mass is 0.1 M\u2299 (Tobin et al. 2012a; Yen et al. 2015a). These three protostars are all embedded in dense cores with masses \u22730.5 M\u2299 (Froebrich 2005; Tachihara et al. 2007; Enoch et al. 2009). Therefore, they are excellent targets by which to study the gas motions on a 100 au scale at an early evolutionary stage.","Citation Text":["Tachihara et al. 2007"],"Functions Text":["Lupus 3 MMS is a Class 0 protostar with a bolometric luminosity (Lbol) of 0.41 L\u2299 and a bolometric temperature (Tbol) of 39 K in the Lupus 3 cloud at a distance of 200 pc"],"Functions Label":["Background"],"Citation Start End":[[776,797]],"Functions Start End":[[604,774]]}
{"Identifier":"2022MNRAS.517.5032D__Dihingia_et_al._2021_Instance_1","Paragraph":"In our study, we consider that the accretion disc is threaded by the poloidal magnetic field lines. The initial poloidal field lines are prescribed by implementing a vector potential A\u03d5 following Zanni et al. (2007) and Vourellis et al. (2019). The functional form of the vector potential is given by\n(6)$$\\begin{eqnarray}\r\nA_\\phi \\propto \\left(r \\sin \\theta \\right)^{3\/4} \\frac{m^{5\/4}}{\\left(m^2 + \\tan ^{-2}(\\theta -\\pi \/2)\\right)^{5\/8}}.\r\n\\end{eqnarray}$$The parameter m(= 0.1) is related to the initial inclination of the field lines and it also determines the magnetic flux of the system. The parameter m play crucial role in the launching of Blandford\u2013Payne type wind from the accretion disc (Blandford & Payne 1982; Dihingia et al. 2021). The initial strength of the poloidal magnetic field is determined by the choice of the plasma-\u03b2 parameter at the truncation radius rtr on the equatorial plane as, $\\beta _{\\rm tr} = p_{\\rm gas}^{\\rm tr}\/p_{\\rm mag}^{\\rm tr}$. Here, superscript \u2018tr\u2019 denotes quantities calculated at the r = rtr. Following our motivation, we carry out eight axisymmetric simulation models, by choosing different truncation radius (rtr), initial plasma-\u03b2 (\u03b2tr), effective resolution (with AMR), and magnetic resistivity. For axisymmetric models, we consider the highest effective resolution to be 2048 \u00d7 1024 (with three refinement levels). The list of input parameters, effective resolution, and the final simulation time (tfinal) for all the simulation models are shown in the Table 1. Out of all these models, we consider 2D40AH to be our reference model for the sake of explanation and comparison. In Fig. 1, we show the initial density distribution (log\u2009(\u03c1\/\u03c1tr)) and the initial gas pressure distribution $(\\log (p_{\\rm gas}\/p_{\\rm gas}^{\\rm tr}))$ for the reference model at panels (a) and (b), respectively. In panel Fig. 1(a), we also show the initial magnetic field lines in terms of grey lines. In Fig. 1(b), the white line represent the boundary of plasma-\u03b2 = 1. In the figure, the density distribution follows equation (2) outside the truncation radius (rtr = 40). Near the equatorial plane, the matter distribution is gas pressure dominated, and far from the equatorial plane and inside the truncation radius, the matter distribution is magnetic pressure dominated.","Citation Text":["Dihingia et al. 2021"],"Functions Text":["The parameter m play crucial role in the launching of Blandford\u2013Payne type wind from the accretion disc"],"Functions Label":["Background"],"Citation Start End":[[724,744]],"Functions Start End":[[595,698]]}
{"Identifier":"2019MNRAS.482.3950S__Shultz_et_al._2015_Instance_1","Paragraph":"As a first step to analysis of NU\u2009Ori\u2019s magnetic field, least-squares deconvolution (LSD) profiles were extracted using a line mask developed from an extract stellar request from the Vienna Atomic Line Database 3 (VALD3; Piskunov et al. 1995; Ryabchikova et al. 1997, 2015; Kupka et al. 1999, 2000) using the stellar parameters determined for NU\u2009Ori\u2009Aa (Teff\u2009=\u200930.5 \u00b1 0.5\u2009kK, log\u2009g = 4.2 \u00b1 0.1) by Sim\u00f3n-D\u00edaz et al. (2011). These parameters were selected since Petit et al. (2008) identified the Aa component as the magnetic star, given that the Stokes V signature is much wider than the $v$sin\u2009i of the secondary component. The line mask was cleaned of all H lines, as well as lines strongly blended with H line wings, lines in spectral regions strongly affected by telluric contamination, lines blended with nebular or interstellar features, and lines in spectral regions affected by ripples. While He lines are often removed due to substantial differences between magnetometry results obtained from He versus metallic lines (e.g. Shultz et al. 2015, 2018b; Yakunin et al. 2015), in this case He lines were left in the mask since the majority of the Stokes V line flux comes from these lines, and the Stokes V profiles extracted using a line mask with He lines excluded did not result in detectable Zeeman signatures. Because of the very high $v$sin\u2009i of the Aa component, LSD profiles were extracted using a velocity range of \u00b1600\u2009km\u2009s\u22121 (in order to include enough continuum for normalization) and a velocity pixel size of 7.2\u2009km\u2009s\u22121, or four times the average ESPaDOnS velocity pixel (thus raising the per pixel S\/N by about a factor of 2). The significance of the signal in Stokes V was evaluated using false alarm probabilities (FAPs), with observations classified as definite detections (DDs), marginal detections (MDs), or non-detections (NDs) according to the criteria described by Donati, Semel & Rees (1992) and Donati et al. (1997). Since FAPs essentially evaluate the statistical significance of the Stokes V signal inside the stellar line by comparing it to the noise level, they are primarily sensitive to the amplitude of Stokes V, which unlike \u2329B$z$\u232a is not strongly dependent on rotational phase. FAPs are thus a complementary means of checking for the presence of a polarization signal, the principal advantage being that they can detect a magnetic field even at magnetic nulls, i.e. \u2329B$z$\u232a\u2009=\u20090.","Citation Text":["Shultz et al. 2015"],"Functions Text":["While He lines are often removed due to substantial differences between magnetometry results obtained from He versus metallic lines (e.g.","in this case He lines were left in the mask since the majority of the Stokes V line flux comes from these lines, and the Stokes V profiles extracted using a line mask with He lines excluded did not result in detectable Zeeman signatures."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1033,1051]],"Functions Start End":[[895,1032],[1082,1319]]}
{"Identifier":"2020ApJ...891...28T__Lind_et_al._2015_Instance_1","Paragraph":"Most GCs are now found to host multiple populations through photometry and spectroscopy (e.g., Carretta et al. 2010b; M\u00e9sz\u00e1ros et al. 2015; Milone et al. 2015; Piotto et al. 2015; Tang et al. 2017, 2018). Chemical abundances from spectroscopic data suggest that GCs have a group of so-called \u201csecond generation\u201d (SG) stars with enhanced N and Na (sometimes He and Al), but depleted C and O (sometimes Mg). These kinds of stars presumably are only formed in the dense environments of GCs. Therefore, identifying field stars with an SG-like chemical pattern is a feasible way to find a link between field stars and GC ejection\/dissolution. Thanks to large spectroscopic surveys, the search for these chemically peculiar stars is becoming more efficient. Using high spectral resolution surveys, multiple elements, like C, N, O, Na, Mg, and Al, can be measured, depending on the wavelength range and signal-to-noise ratio. Toward this, the Apache Point Observatory Galactic Evolution Experiment (APOGEE; Majewski et al. 2017) and Gaia-ESO survey have led to the discovery of a large group of N-rich field stars (Lind et al. 2015; Fern\u00e1ndez-Trincado et al. 2016, 2017, 2019a; Martell et al. 2016; Schiavon et al. 2017). While high-resolution spectra give more elements for a detailed investigation of their chemical history, low-resolution spectra can supposedly extend the search for N-rich field stars toward fainter and more numerous samples (Martell & Grebel 2010; Martell et al. 2011; Koch et al. 2019). Simultaneously observing 4000 stars with fibers makes LAMOST an unprecedented machine in collecting low-resolution stellar spectra. Using the CN\u2013CH band features around 4000 \u212b, we have identified \u223c40 N-rich field stars8\n\n8\nAlso called CN-strong, CH-normal stars in Paper I.\n in LAMOST DR3 (Tang et al. 2019, hereafter Paper I). The derived N abundances of these stars are clearly higher than those of the metal-poor field stars, indicating that (1) our sample is a bona fide sample of N-rich field stars and (2) the classical extra-mixing theory may not work for these stars. Moreover, a substantial fraction of retrograding N-rich field stars suggest that some N-rich field stars may be accreted. In this work, we expand our sample to \u223c100 N-rich field stars in LAMOST DR5 (Section 2), making it more robust for drawing statistical conclusions, especially for the GC origin of these field stars. We put forward a detailed analysis of high-resolution chemical abundances and kinematics (Sections 3 and 4) to discuss the origins of these N-rich field stars (Section 5). As the second paper of this series, we will also call this present work Paper II.","Citation Text":["Lind et al. 2015"],"Functions Text":["Thanks to large spectroscopic surveys, the search for these chemically peculiar stars is becoming more efficient. Using high spectral resolution surveys, multiple elements, like C, N, O, Na, Mg, and Al, can be measured, depending on the wavelength range and signal-to-noise ratio. Toward this,","and Gaia-ESO survey have led to the discovery of a large group of N-rich field stars"],"Functions Label":["Background","Background"],"Citation Start End":[[1108,1124]],"Functions Start End":[[638,931],[1022,1106]]}
{"Identifier":"2017AandA...601A..64Z__Milosavljevi\u0107_&_Nakar_2006_Instance_1","Paragraph":"The distribution function of the electrons inside the precursor can be formally determined through a Lorentz transform of the corresponding distribution function in the shock frame, and the latter can be obtained by solving a stationary transport equation with a loss term corresponding to the finite residence time spent in the precursor. Here, we approximate this distribution function in the precursor frame as follows: (3)\\begin{equation} \\frac{{\\rm d}n_{\\rm e\\vert\\rm p}}{{\\rm  d}\\gamma_{\\rm e}}\\,\\simeq\\,\\frac{R^2}{r^2}\\, \\frac{\\vert  s-1\\vert}{\\gamma_{\\rm m\\vert p}}\\left(\\frac{\\gamma_{\\rm e}}{\\gamma_{\\rm  m\\vert p}}\\right)^{-s}\\, 2\\Gamma_{\\rm sh}^2n_{\\rm  u}\\,\\Theta\\left[r_{\\rm p}(\\gamma_{\\rm e})-r\\right].  \\label{eq:dndg1} \\end{equation}dne|pd\u03b3e\u2009\u2243\u2009R2r2\u2009|s\u22121|\u03b3m|p\u03b3e\u03b3m|p\u2212s\u20092\u0393sh2nu\u2009\u0398[rp(\u03b3e)\u2212r].The Heaviside function models the finite length scale up to which particles of a given Lorentz factor can travel. This precursor length scale is related to the (upstream frame) residence time tres through \\hbox{$r_{\\rm p}\\,=\\,R+c t_{\\rm res}\\left(1-\\beta_{\\rm sh}\\right)\\,\\simeq\\,R+ct_{\\rm res}\/(2\\Gamma_{\\rm sh}^2)$}rp\u2009=\u2009R+ctres1\u2212\u03b2sh\u2009\u2243\u2009R+ctres\/(2\u0393sh2). The residence time can be calculated as the time it takes to deflect the accelerated electron by an angle ~1\/\u0393sh (Achterberg et al. 2001; Milosavljevi\u0107 & Nakar 2006; Pelletier et al. 2009). If the background magnetic field controls the transport of the accelerated electrons, this residence time is \\hbox{$t_{\\rm res}^{(1)}\\,\\simeq\\,\\Gamma_{\\rm sh}^{-1}\\gamma_{\\rm e} m_{\\rm e} c\/(e B_{\\rm u})$}tres(1)\u2009\u2243\u2009\u0393sh-1\u03b3emec\/(eBu). In contrast, if pitch angle scattering in a micro-turbulence of length scale \u03bb\u03b4B governs the return of particles to the shock, \\hbox{$t_{\\rm res}^{(2)}\\,\\simeq\\,\\Gamma_{\\rm sh}^{-2}\\gamma_{\\rm e}^2m_{\\rm e}^2c^3\/(\\lambda_{\\delta B}e^2\\delta B_{\\rm p}^2)$}tres(2)\u2009\u2243\u2009\u0393sh-2\u03b3e2me2c3\/(\u03bb\u03b4Be2\u03b4Bp2). Whether one or the other occurs depends on \u03b3e and the hierarchy between Bu and \u03b4Bp. As the scattering frequency in a micro-turbulence scales as \\hbox{$\\gamma_{\\rm e}^{-2}$}\u03b3e-2, while the gyrofrequency scales with \\hbox{$\\gamma_{\\rm e}^{-1}$}\u03b3e-1, we expect the background field to control the residence time at higher energies. In order to bracket this realistic scenario, we consider in the following the above two extremes: (1) where regular deflection in a background field dominates at all energies and (2) where stochastic deflection in a small scale micro-turbulence of uniform energy density dominates at all energies. Correspondingly, we write \\hbox{$r_{\\rm p}(\\gamma_{\\rm e})=R+\\Delta_i\\gamma_{\\rm e}^i$}rp(\u03b3e)=R+\u0394i\u03b3ei, with i = 1,2 for models (1) or (2); \\hbox{$\\Delta_i=ct_{\\rm res}^{(i)}\/(2\\Gamma_{\\rm sh}^2)$}\u0394i=ctres(i)\/(2\u0393sh2). ","Citation Text":["Milosavljevi\u0107 & Nakar 2006"],"Functions Text":["The residence time can be calculated as the time it takes to deflect the accelerated electron by an angle ~1\/\u0393sh"],"Functions Label":["Uses"],"Citation Start End":[[1295,1321]],"Functions Start End":[[1157,1269]]}
{"Identifier":"2021AandA...650A.155Z__Oh_et_al._2012_Instance_2","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"],"Functions Text":["In addition, there is evidence demonstrating that bars can enhance star formation in the central regions of galaxies (e.g."],"Functions Label":["Background"],"Citation Start End":[[720,734]],"Functions Start End":[[597,719]]}
{"Identifier":"2015AandA...584A..32M__Sargent_et_al._2010_Instance_1","Paragraph":"Nearly all of the cm-wavelength radio emission from star-forming galaxies, such as SMGs, is non-thermal synchrotron radiation from relativistic electrons accelerated in supernova (SN) remnants produced by the short-lived, high-mass OB-type stars (M \u2273 8 M\u2299; main-sequence lifetime \u03c4MS \u2272 30 Myr). Because SNe trace the recent\/on-going star formation, the radio synchrotron emission has the potential to trace the spatial scales on which star formation is occurring. This connection between radio emission and star formation is strongly supported by the close infrared (IR)-radio correlation observed in galaxies (e.g. Helou et al. 1985; Beck & Golla 1988; Xu et al. 1992; Condon 1992; Yun et al. 2001; Bell 2003; Tabatabaei et al. 2007; Murphy et al. 2008; Sargent et al. 2010; Mori\u0107 et al. 2010; Dumas et al. 2011). On the basis of this correlation, the IR-emitting region of a star-forming galaxy is expected to be comparable in size to that of radio continuum emission. However, the most recent studies of the sizes of IR-emitting regions of SMGs based on continuum imaging observations with the Atacama Large Millimetre\/submillimetre Array (ALMA) show that these are significantly smaller than SMG radio sizes presented in the literature (Simpson et al. 2015a; Ikarashi et al. 2015). A possible explanation for this discrepancy, as suggested by Simpson et al. (2015a), is cosmic ray (CR) diffusion in the galactic magnetic field away from their acceleration site, which would render larger radio sizes. To test this further here we present a study of radio sizes of SMGs from a well selected sample of SMGs in the Cosmic Evolution Survey (COSMOS; Scoville et al. 2007) deep field using radio data from the Karl G. Jansky Very Large Array (VLA)-COSMOS 3 GHz Large Project (1\u03c3 noise of 2.3 \u03bcJy beam-1, angular resolution \\hbox{$0\\farcs75$}0 .\u030b 75; Smol\u010di\u0107 et al., in prep.). We describe the SMG sample and the employed VLA data in detail in Sect. 2. The 3 GHz images are presented in Sect. 3, and the analysis (size measurements and radio spectral indices) are presented in Sect. 4. We compare our results with literature studies in Sect. 5, discuss the results in Sect. 6, and summarise the main results of the paper in Sect. 7. ","Citation Text":["Sargent et al. 2010"],"Functions Text":["This connection between radio emission and star formation is strongly supported by the close infrared (IR)-radio correlation observed in galaxies (e.g."],"Functions Label":["Background"],"Citation Start End":[[755,774]],"Functions Start End":[[464,615]]}
{"Identifier":"2022ApJ...926...85S__Ehrenreich_et_al._2020_Instance_2","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)"],"Functions Text":["We use the quadratic limb-darkening coefficients reported by","to establish the stellar center-to-limb intensity profile,"],"Functions Label":["Uses","Uses"],"Citation Start End":[[919,943]],"Functions Start End":[[858,918],[944,1002]]}
{"Identifier":"2022MNRAS.513.4464T__Hopkins_et_al._2020_Instance_2","Paragraph":"\nGalactic winds: Galactic winds driven by CRs have often been simulated in two limits: a diffusion-dominated regime, due possibly to \u2018extrinsic confinement\u2019, where CRs are scattered by extrinsic turbulence, and\/or due to various wave damping mechanisms (e.g. ion neutral damping) and streaming-dominated \u2018self confinement\u2019, where CRs are confined by Alfven waves they produce via the gyroresonant streaming instability. In the diffusive \u2018extrinsic confinement\u2019 case, CRs do not heat the gas.19 In the streaming dominated \u2018self confinement\u2019 case, CR transport heats gas at a rate vA \u00b7 \u2207Pc. The diffusive case fits \u03b3 ray observations better, because CRs can propagate out of the galaxy faster (Chan et al. 2019). It is also much better at driving winds, because the CRs do not suffer strong energy losses via Alfven wave heating (Wiener et al. 2017b; Hopkins et al. 2020). However, we expect self-confinement to be very strong at the \u223cGeV energies where CR energy peaks (Kulsrud & Pearce 1969; Farmer & Goldreich 2004; Wiener et al. 2013), while extrinsic compressible turbulence is strongly damped at small scales, and unlikely to efficiently scatter \u223cGeV CRs (Yan & Lazarian 2002). Thus, CR winds should be streaming dominated and relatively inefficient. The CR staircase changes these dichotomies by changing the structure of the wind. We have seen how CR pressure can build up in streaming dominated simulations, due to trapping at bottlenecks. This increases mass outflow rates, similar to the effect of increased opacity in radiative outflows. In CR streaming simulations of isothermal winds where the CR acoustic instability arose, Quataert et al. (2022a) found an increase in wind mass loss rates by an order of magnitude, compared to analytic models without a CR staircase, illustrating the potential impact of CR staircases. High-resolution cosmological zoom simulations of CR staircases are actually well within reach. As seen in Appendix Section B, all that is required is that the diffusion length $l_{\\rm diff} \\sim \\kappa \/c_{\\rm s} \\sim 2 \\, {\\rm kpc} \\, \\left(\\frac{\\kappa }{10^{29} {\\rm cm^2 s^{-1}}} \\right)\\left(\\frac{c_{\\rm s}}{150 \\, {\\rm km \\, s^{-1}}} \\right)^{-1}$ is resolved. However, to date only the FIRE collaboration has implemented the two moment method (capable of dealing with CR streaming) in such simulations, and \u2013 in contrast to, for instance, van de Voort et al. (2021) \u2013 the plasma \u03b2 in their winds is too high for the acoustic instability to develop (Hopkins et al. 2020). But alternate setups where CR staircases appear are certainly numerically feasible.","Citation Text":["Hopkins et al. 2020"],"Functions Text":["However, to date only the FIRE collaboration has implemented the two moment method (capable of dealing with CR streaming) in such simulations, and \u2013 in contrast to, for instance, van de Voort et al. (2021) \u2013 the plasma \u03b2 in their winds is too high for the acoustic instability to develop"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[2490,2509]],"Functions Start End":[[2201,2488]]}
{"Identifier":"2016ApJ...826..168X__Bai_2014_Instance_1","Paragraph":"MRI is considered to be the most promising mechanism driving angular-momentum transport in protoplanetary disks (Balbus & Hawley 1991; Brandenburg et al. 1995; Hawley et al. 1995; Balbus et al. 1996; Balbus & Hawley 1998). However, protoplanetary disks are cold, dense, and, therefore, poorly ionized. The low level of ionization tends to decouple the disk gas from magnetic fields, which generates non-ideal MHD effects: Ohmic dissipation, ambipolar diffusion (AD), and the Hall effect (e.g., Armitage 2011; Turner et al. 2014). These effects quench MRI in different ways: Ohmic dissipation originates from collisions between electrons and neutrals, AD from collisions between ions and neutrals, and the Hall effect from drift between electrons and ions (Fleming et al. 2000; Sano & Stone 2002; Bai & Stone 2011). Ohmic dissipation operates in high-density regions with weak field, AD dominates in highly ionized and low-density regions, and the Hall effect lies in between (Fleming & Stone 2003; Bai & Stone 2011; Bai 2014). So far, the effect of Ohmic dissipation has been best studied. Investigations show the layered accretion in the inner disk, where the midplane region is \u201cdead\u201d due to low ionization while the surface layer is \u201cactive\u201d due to sufficient ionization (Gammie 1996; Jin 1996; Fleming et al. 2000; Fleming & Stone 2003; Turner et al. 2007; Ilgner & Nelson 2008; Oishi & Mac Low 2009; Okuzumi & Hirose 2011). Recent works that take into account both Ohmic dissipation and AD show that AD may render the surface layer and portions of the outer disk inactive (Bai & Stone 2011; Landry et al. 2013; Kalyaan et al. 2015). Bai & Stone (2013) find that MRI is completely suppressed in the inner disk and a strong magnetocentrifugal wind is launched. Three-dimensional simulations that include all three non-ideal MHD effects are also performed (Bai 2014, 2015; Lesur et al. 2014; Simon et al. 2015). In the inner disk, the influence of the Hall effect on midplane angular-momentum transport depends on the orientation of the vertical magnetic field with the disk rotation axis. When the field is aligned with the axis, the enhanced Maxwell stress promotes angular-momentum transport. When the field is anti-aligned with the axis, the midplane remains quiescent. In the outer disk, the Hall effect has little influence on the disk turbulence. Although the inclusion of AD and the Hall effect substantially changes the level of turbulence in the protoplanetary disks, the feature that the viscosity is low in the inner disk and high in the outer disk is still valid. In this study, we assume that gas giant planets form in situ via the core accretion scenario, which implies that their formation locations are always in the low-viscosity region. Since in this study we focus on the relation between photoevaporation and planet formation and gap opening by planets in the disk, we adopt Ohmic dissipation to represent the non-ideal MHD effects on the MRI. We consider that this simplification has little influence on our main calculation results.","Citation Text":["Bai 2014"],"Functions Text":["Ohmic dissipation operates in high-density regions with weak field, AD dominates in highly ionized and low-density regions, and the Hall effect lies in between"],"Functions Label":["Background"],"Citation Start End":[[1016,1024]],"Functions Start End":[[815,974]]}
{"Identifier":"2020AandA...643A.149S__Nilsson_et_al._2011_Instance_1","Paragraph":"An increasing number of recent works focus on the study of high-redshift Lyman-\u03b1 emitters (LAEs), objects showing prominent rest-frame Ly\u03b1 emission within a spectrum (usually) devoid of other line features (e.g. Cassata et al. 2011; Nakajima et al. 2018). The spectral properties of LAEs are usually interpreted as coming from young (\u227250 Myr) and low-mass (M*\u200b\u2004 \u2004\u200b1010\u2006M\u2299) galaxies (e.g. Wilkins et al. 2011; Amor\u00edn et al. 2017; Hao et al. 2018; Santos et al. 2020) with small rest-frame UV half-light radii (R\u200b\u2004\u2272\u2004\u200b1\u2005\u2212\u20052 Kpc, as in e.g. M\u00f8ller & Warren 1998; Lai et al. 2008; Bond et al. 2012; Guaita et al. 2015; Kobayashi et al. 2016; Ribeiro et al. 2016; Bouwens et al. 2017a; Paulino-Afonso et al. 2018) which are actively star forming (SFR\u200b\u2004\u223c\u2004\u200b1\u2005\u2212\u2005100\u2006M\u2299\u2006yr\u22121) and dust poor (dust attenuation AV\u200b\u2004 \u2004\u200b0.2, see e.g. Gawiser et al. 2006, 2007; Guaita et al. 2011; Nilsson et al. 2011; Bouwens et al. 2017b; Arrabal Haro et al. 2020). When observed at high redshift, isolated and grouped LAEs appear to represent the progenitors of present-day galaxies and clusters, respectively, providing extremely valuable insights into structure formation (e.g. Matsuda et al. 2004, 2005; Venemans et al. 2005; Gawiser et al. 2007; Overzier et al. 2008; Guaita et al. 2010; Mei et al. 2015; Bouwens et al. 2017b; Khostovan et al. 2019). A basic statistical tool to study the population of high-z LAEs is the description of their number density at a given redshift as a function of line luminosity (LLy\u03b1), namely the Ly\u03b1 luminosity function (LF; see e.g. Gronke et al. 2015, for a theoretical approach). Several recent works have focused on the construction of the Ly\u03b1 LF at z\u200b\u2004\u2265\u2004\u200b2 (Gronwall et al. 2007; Ouchi et al. 2008; Blanc et al. 2011; Cl\u00e9ment et al. 2012; Konno et al. 2016; Sobral et al. 2017, 2018a) by making use of deep observations of narrow sky regions (up to few squared degrees, as in e.g. Matthee et al. 2014, 2017b; Cassata et al. 2015; Ono et al. 2018). Their findings describe a Ly\u03b1 LF that follows a Schechter function (Schechter 1976) at relatively faint line luminosity (i.e. LLy\u03b1\u2004\u2272\u20041042.5, see e.g. Ouchi et al. 2008; Konno et al. 2016; Sobral et al. 2016; Matthee et al. 2017a), a regime mostly occupied by low-mass star-forming galaxies (e.g. Hu et al. 1998; Kudritzki et al. 2000; Stiavelli et al. 2001; Santos et al. 2004; van Breukelen et al. 2005; Gawiser et al. 2007; Rauch et al. 2008; Guaita et al. 2011).","Citation Text":["Nilsson et al. 2011"],"Functions Text":["The spectral properties of LAEs are usually interpreted as coming from young (\u227250 Myr) and low-mass (M*\u200b\u2004 \u2004\u200b1010\u2006M\u2299) galaxies","which are actively star forming (SFR\u200b\u2004\u223c\u2004\u200b1\u2005\u2212\u2005100\u2006M\u2299\u2006yr\u22121) and dust poor (dust attenuation AV\u200b\u2004 \u2004\u200b0.2, see e.g."],"Functions Label":["Background","Background"],"Citation Start End":[[866,885]],"Functions Start End":[[256,381],[708,818]]}
{"Identifier":"2019AandA...631A.109W__Rivera_et_al._(2017)_Instance_1","Paragraph":"With the advent of the LOw Frequency ARray (LOFAR; R\u00f6ttgering et al. 2011; van Haarlem et al. 2013) which combines a large field of view with high sensitivity on both small and large angular scales, we can now study the FIRC at lower frequencies where the contribution from thermal free-free emission is even less important than at 1.4 GHz. Operating between 30 and 230 MHz, LOFAR offers complementary information to the wealth of data collected at higher frequencies. Using deep LOFAR 150 MHz observations in the 7 deg2 Bo\u00f6tes field (Williams et al. 2016), Calistro Rivera et al. (2017) studied the FIRC at 150 MHz from z\u2004\u223c\u20040.05 out to z\u2004\u223c\u20042.5. They found fairly mild redshift evolution in the logarithmic IR to radio luminosity ratio in the form of qIR\u2004\u223c\u2004(1\u2005+\u2005z)\u22120.22\u2005\u00b1\u20050.05. However, if the FIRC is non-linear (i.e. the logarithmic slope is different from one), then it implies that the qIR parameter would depend on luminosity. Therefore the reported redshift dependence of qIR may simply be a consequence of the non-linearity of the FIRC (Basu et al. 2015) as the mean SFR of galaxies is generally larger at higher redshifts (e.g., Hopkins & Beacom 2006; Madau & Dickinson 2014; Pearson et al. 2018; Liu et al. 2018; Wang et al. 2019). Based on LOFAR observations of the Herschel Astrophysical Terahertz Large Area Survey (H-ATLAS; Eales et al. 2010) 142 deg2 North Galactic Pole (NGP) field (Hardcastle et al. 2016), G\u00fcrkan et al. (2018) found that a broken power-law (with a break around SFR \u223c1\u2006M\u2299 yr\u22121) compared to a single power law is a better calibrator for the relationship between RC luminosity and SFR, possibly implying additional mechanisms for generating cosmic rays and\/or magnetic fields. Also using LOFAR data in the NGP field, Read et al. (2018) found evidence for redshift evolution of the FIRC at 150 MHz. Heesen et al. (2019) studied the relation between radio emission and SFR surface density using spatially resolved LOFAR data of a few nearby spiral galaxies. They found a sublinear relation between the resolved RC emission and the SFR surface densities based on GALEX UV and Spitzer 24 \u03bcm data.","Citation Text":["Calistro Rivera et al. (2017)"],"Functions Text":["Using deep LOFAR 150 MHz observations in the 7 deg2 Bo\u00f6tes field","studied the FIRC at 150 MHz from z\u2004\u223c\u20040.05 out to z\u2004\u223c\u20042.5. They found fairly mild redshift evolution in the logarithmic IR to radio luminosity ratio in the form of qIR\u2004\u223c\u2004(1\u2005+\u2005z)\u22120.22\u2005\u00b1\u20050.05.","However, if the FIRC is non-linear (i.e. the logarithmic slope is different from one), then it implies that the qIR parameter would depend on luminosity."],"Functions Label":["Background","Background","Compare\/Contrast"],"Citation Start End":[[558,587]],"Functions Start End":[[469,533],[588,777],[778,931]]}
{"Identifier":"2019AandA...632A.129W__Shodhan_et_al._2000_Instance_1","Paragraph":"Coronal mass ejections (CMEs) are intense solar explosive eruptions during which large amounts of plasma and magnetic field from the solar atmosphere are ejected into interplanetary space. The interplanetary manifestations of CMEs (ICMEs; Kilpua et al. 2017) can be measured by a spacecraft at about 1 AU and exhibit the following characteristics: increase in total magnetic magnitude (Cane & Richardson 2003), helium abundance (Hirshberg et al. 1972; Zwickl et al. 1982; Richardson & Cane 2004), average iron ionization (Lepri et al. 2001; Lepri & Zurbuchen 2004), and O7+ abundance (Richardson & Cane 2004; Wang & Feng 2016); decrease in proton temperatures and proton densities (Gosling et al. 2001; Zhang et al. 2013); counterstreaming suprathermal electron (CSE) strahls and declining speed (Zwickl et al. 1982; Gosling et al. 1987; Shodhan et al. 2000; Burlaga et al. 2001). A subset of ICMEs was defined as magnetic cloud (MC) by Burlaga et al. (1981) empirically using the following properties: (1) the magnetic field strength is higher than average, (2) a smooth change in field direction as observed by a spacecraft passing through the cloud, and (3) low proton temperature compared to that of the ambient proton. Magnetic clouds usually have magnetic flux rope structures, and they are the main source of major geomagnetic storms (Burlaga et al. 1981; Webb et al. 2000; Huttunen et al. 2002; Zhang et al. 2007). Observations at 1 AU show that 30\u2005\u2212\u200540% of ICMEs are MCs, and this percentage depends on the solar cycle (Richardson & Cane 2004). However, CMEs are usually assumed to have magnetic flux rope structures near the Sun because of their helical shapes (Canfield et al. 1999; Liu et al. 2010; Rust & Kumar 1996). This begs the question of whether or not nonMC ICMEs also have flux rope structures. The journal of Solar Physics once made a special issue to address this question (Gopalswamy et al. 2013a). A comparative study of 23 MCs and 31 nonMC ICMEs was completed, and the source regions of the 54 ICMEs were located within \u00b115\u00b0 longitude from the disk center. Yashiro et al. (2013) found no significant difference between the structures of the post-eruption arcades of MCs and nonMC ICMEs during launch. Gopalswamy et al. (2013b) observed that MCs and nonMC ICMEs have significant enhancement in Fe and O charge states, and Fe and O charge-state measurements are positively correlated with flare properties, including flare size and soft X-ray flare intensity. Their observations suggest that these CMEs have similar explosive environment and flux rope structures near the Sun. Furthermore, some studies indicate that CMEs associated with MCs tend to propagate along the Sun\u2013Earth line, whereas nonMC events are deflected away from the Sun\u2013Earth line (Kim et al. 2013; M\u00e4kel\u00e4 et al. 2013; Zhang et al. 2013). Therefore, many researchers believe that all ICMEs have magnetic flux rope structures and that the nonMC events are due to observational limitations, that is, that the observing spacecraft crosses the flanks of the ropes and therefore the ICMEs appear as nonMCs. This has been shown by some multi-satellite-observed ICMEs, namely spacecraft farther from the axis detect less clear flux rope signatures than centrally crossing spacecraft for the same event (Cane et al. 1997; Kilpua et al. 2011).","Citation Text":["Shodhan et al. 2000"],"Functions Text":["The interplanetary manifestations of CMEs","can be measured by a spacecraft at about 1 AU and exhibit the following characteristics:","counterstreaming suprathermal electron (CSE) strahls and declining speed"],"Functions Label":["Background","Background","Background"],"Citation Start End":[[838,857]],"Functions Start End":[[189,230],[259,347],[723,795]]}
{"Identifier":"2019ApJ...878..108L___1991_Instance_1","Paragraph":"A charged particle moving in the electromagnetic field can feel Landau\u2013Lifshitz radiation reaction force due to synchrotron radiation, thus modifying the motion of the charged particle significantly when taking into account the radiation reaction force. Recent observations have demonstrated that there is strong evidence that a magnetic field of several hundred Gauss exists in the vicinity of the supermassive black hole at the center of the Milky Way (Eatough et al. 2013). A dynamo mechanism from an accretion disk around a black hole accounts for the appearance of such a magnetic field (Punsly 2001; Brandenburg et al. 1995; Hawley et al. 1996). There is also clear evidence that the magnetic field on the surface of the neutron star can be up to 1014 G (Duncan & Thompson 1992; Paczy\u0144ski 1992; Usov 1992; Thompson & Duncan 1995, 1996; Vasisht & Gotthelf 1997). The Landau\u2013Lifshitz radiation reaction has been investigated in detail in Landau & Lifshitz (1975) in flat space, while, in curved space, the radiation reaction is described in Sokolov et al. (1978, 1983), DeWitt & Brehme (1960), and Tursunov et al. (2018). Thus, the radiation reaction force can affect the charged particles from the accretion disk around the black hole and the neutron star significantly. Radiation reactions have been an important element of pair creation scenarios in positron\u2013electron plasma just above the pole of the event horizon (Hirotani & Pu 2016). Curvature radiation in black hole magnetosphere pair creation schemes is radiation resistance limited (Broderick & Tchekhovskoy 2015). Recently, a radiation reaction has also been applied to protonic acceleration in the vortex above the pole of the black hole (Ruffini et al. 2018). Radiation reactions are of fundamental importance in the evacuated vortex of black hole magnetospheres (Punsly 2001). In particular, if the field line angular velocity is set much less than the horizon angular velocity by distant plasma and a very tenuous plasma exists in the event horizon magnetosphere, then radiation resistance will determine the flow dynamics of accretion as well as the rotational energy extraction by a putative jet (Punsly 2001, 1991). The dynamics mentioned above cannot be revealed by MHD simulations with mass floors. The ad hoc injection of mass will damp any large waves that can break ideal MHD and prevent the associated large local electromagnetic forces from being achieved. Therefore, all existing numerical simulations of the black hole magnetosphere bypass the radiation reaction dominated dynamics as a consequence of numerical dissipation of waves and numerical diffusion in the MHD system in the evacuated vortex above the event horizon. Thus, a proper treatment of radiation reaction is critical for assessing the time evolution of these types of astrophysical systems. Numerical simulations using general relativistic magnetohydrodynamics (GRMHDs) are applied to investigate the physical process in accretion disks around neutron stars, as well as in microquasars (\u03bcQSOs), gamma-ray bursts, and active galactic nuclei. In curved spacetime, the dynamical evolution of the high energy disks made of ion\u2013electron plasma in simulated GRMHD is usually performed without the contribution from Landau\u2013Lifshitz radiation reaction force, which may play a crucial role. The method of particle-in-cell (Chen & Beloborodov 2014; Philippov & Spitkovsky 2014; Belyaev 2015; Cerutti et al. 2015, 2016; Philippov et al. 2015b) containing the radiation reaction forces has been investigated in flat space, while, in curved space, the same method with a radiation reaction has been achieved in Philippov et al. (2015a). Incorporating the radiation reaction into the relativistic magnetohydrodynamic equations governing the dynamics of plasma has been studied by Tam & Kiang (1979) and Berezhiani et al. (2004, 2008). In a recent study, Liu et al. (2018) achieved the one-fluid relativistic magnetohydrodynamics description of two-fluid plasma in which the Landau\u2013Lifshitz radiation reaction is incorporated. However, numerical simulation with the radiation reaction from Liu et al. (2018) is not practical due to its highly nonlinear form at present, we expect that analytical investigations are essential and provide more motivation for future work. Thus, in this work, similar to Liu et al. (2018), we get the GRMHD equation for the one-fluid description of two-fluid plasma containing a Landau\u2013Lifshitz radiation reaction in curved space, and these results could be applied to both positron\u2013electron and proton\u2013electron plasma.","Citation Text":["Punsly","1991"],"Functions Text":["In particular, if the field line angular velocity is set much less than the horizon angular velocity by distant plasma and a very tenuous plasma exists in the event horizon magnetosphere, then radiation resistance will determine the flow dynamics of accretion as well as the rotational energy extraction by a putative jet"],"Functions Label":["Background"],"Citation Start End":[[2169,2175],[2182,2186]],"Functions Start End":[[1846,2167]]}
{"Identifier":"2020MNRAS.496.4127V__Homan_et_al._2010_Instance_1","Paragraph":"The recent report by Buisson et al. (2020b) of the detection of Type-I bursts in Sw J1858 shows the presence of a neutron star primary. Therefore, here, we first consider whether Eddington-limited accretion in such a neutron star LMXB might explain the observed radio behaviour in Sw J1858. The neutron star LMXBs with the highest mass accretion rates are the Z-sources, which are thought to accrete near or at the Eddington luminosity, tracing out Z-shaped tracks in their X-ray colour\u2013colour diagrams (Hasinger & van der Klis 1989; Homan et al. 2010). Z-sources can show strong changes in radio brightness, related to the branch they are positioned on in their colour\u2013colour diagram track; they are radio brighter and more variable in the Horizontal Branch than in the Flaring and Normal Branches (Penninx et al. 1988; Hjellming et al. 1990a,b; Spencer et al. 2013; Motta & Fender 2019). Time-resolved radio studies of the different branches in Sco X-1 by Hjellming et al. (1990b) and Cyg X-2 by Hjellming et al. (1990a) show that these sources have similar levels of radio variability and luminosity to Sw J1858 during their radio-faint Flaring and lower Normal Branches. However, in the X-ray \u2013 radio luminosity plane, these sources are located to the right of Sw J1858, around the neutron star Eddington limit of 2 \u00d7 1038\u2009erg\u2009s\u22121. For Sw J1858 to be similar to these sources, it would therefore have to be viewed a high inclination, reducing its observed X-ray flux and masking the Z-source variability properties through obscuration. However, a clear radio difference between Sw J1858 and the Z-sources is the radio spectral index: contrary to Sw J1858, Z-sources typically show steep spectra and are indeed associated with the launch of (resolved) discrete ejecta (Motta & Fender 2019). Also, in this scenario, Sw J1858 should not have resided in the much more radio bright Horizontal Branch during any of the observations, which is unlikely given the time-scales and commonness of transitions between the branches (Homan et al. 2010).","Citation Text":["Homan et al. 2010","Homan et al. 2010"],"Functions Text":["Therefore, here, we first consider whether Eddington-limited accretion in such a neutron star LMXB might explain the observed radio behaviour in Sw J1858. The neutron star LMXBs with the highest mass accretion rates are the Z-sources, which are thought to accrete near or at the Eddington luminosity, tracing out Z-shaped tracks in their X-ray colour\u2013colour diagrams","Also, in this scenario, Sw J1858 should not have resided in the much more radio bright Horizontal Branch during any of the observations, which is unlikely given the time-scales and commonness of transitions between the branches"],"Functions Label":["Uses","Compare\/Contrast"],"Citation Start End":[[534,551],[2023,2040]],"Functions Start End":[[136,502],[1794,2021]]}
{"Identifier":"2017MNRAS.465..383R__Tong_et_al._2013_Instance_1","Paragraph":"In general, the Hall time-scale for magnetic field evolution depends on the strength of the magnetic field, as seen in equation (4). For young NSs with fields below 1014 G, this time-scale may be longer than the observed SNR age. Therefore, magnetic field growth does not have a dramatic effect on these young NSs. However, many of these systems are observed to have a braking index n  3 (Espinoza 2012; Archibald et al. 2016). A possible explanation for these low braking indices may be through the emission of a relativistic particle wind (Thompson & Blaes 1998; Harding, Contopoulos & Kazanas 1999; Tong et al. 2013). However, the conclusive detection of wind nebulae around magnetars in particular is challenging due to the presence of dust-scattering haloes that accompany these X-ray bright, heavily absorbed objects (Esposito et al. 2013; Safi-Harb 2013). Only a handful of such nebulae have been proposed to be associated with highly magnetized NSs. For example, a wind nebula has been proposed to surround the magnetar Swift J1834.9\u20130846 in W41 (Younes et al. 2016), and the luminosity of a particle wind was estimated for SGR 1806\u201320 based on the X-ray and radio observations of the wind-powered nebula G10.0\u20130.3 (Thompson & Duncan 1996; Marsden, Rothschild & Lingenfelter 1999; Gaensler et al. 2005). In the pulsar wind model, relativistic particles load the magnetosphere with charge and distort the dipole field at large scales outside of the light cylinder. Besides affecting the NS spin-down, the emission of a relativistic wind can also offer an explanation for the significant timing noise that generally affects magnetar observations (Tsang & Konstantinos 2013). The HBPs J1119\u20136127 and J1846\u20130258 are clearly associated with pulsar wind nebulae (Gavriil et al. 2008; Kumar & Safi-Harb 2008; Ng et al. 2008; Safi-Harb & Kumar 2008; Safi-Harb 2013), suggesting that particle wind emission should play an important role in their evolution. We also expect that AXP and SGR evolution may be affected by wind emission due to candidate wind nebulae, but do not consider these models for the CCOs that do not show any evidence of PWN. However, we note that braking exclusively due to a steady particle wind produces a torque with a braking index n = 1, too low for the NSs with secure SNR associations in Table 1.","Citation Text":["Tong et al. 2013"],"Functions Text":["A possible explanation for these low braking indices may be through the emission of a relativistic particle wind"],"Functions Label":["Background"],"Citation Start End":[[602,618]],"Functions Start End":[[428,540]]}
{"Identifier":"2017MNRAS.470.2959K___2013_Instance_1","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)"],"Functions Text":["For example, Sirko et al. (2004), Kafle et al. (2012),","Kafle et al. (2014) and King et al. (2015) fit an ellipsoidal distribution of velocities"],"Functions Label":["Background","Background"],"Citation Start End":[[1263,1282]],"Functions Start End":[[1208,1262],[1284,1372]]}
{"Identifier":"2021AandA...645A..95H__Boese_2000_Instance_2","Paragraph":"Next we chose an optimum cut-off radius for the detector FOV. The PSPC has a circular FOV with a radius 57\u2032. The PSPC entrance window has a rib support structure with an inner ring at a radius corresponding to 20\u2032 (Pfeffermann et al. 1987; Hasinger & Zamorani 2000). Both the ROSAT telescope angular resolution and its vignetting function are roughly constant within the inner 20\u2032 ring, but degrade significantly towards larger off-axis angles. The combined detector and telescope PSFs are described in detail in Boese (2000). To the first order, the PSF at each off-axis angle can be approximated by a Gaussian function with a half power radius (HPR) of 13, 22, 52, 93, 130, and 180\u2033, at off-axis angles of 0, 12, 24, 36, 48, and 57\u2032, respectively (at 1 keV). The vignetting function at 1 keV drops almost linearly to about 50% at an off-axis angle of 50\u2032. Taking into account all these effects, the HPR of the overall RASS PSF is 84\u2033 (Boese 2000). This means that the classical confusion limit (40 beams per source) is reached at a source density of about 15 sources deg\u22122, which is exceeded in the high-exposure areas of our survey. In addition, we need to optimally discriminate between extended and point-like X-ray sources, calling for an angular resolution that is as high as possible. We therefore have to reduce the detector FOV. The sharpest imaging is achieved within the inner 20\u2032 of the PSPC FOV, corresponding to the inner ring-like rib of the PSPC support structure (see Fig. 1). However, there is a trade-off between image sharpness and the number of photons required for detection and image characterization. In particular in the outer areas of our survey, where the RASS exposure times drop significantly, a 20\u2032 FOV radius does not provide sufficient exposure time. Taking into account the various competing factors in this trade-off, we made a few tests varying the FOV cut-off radius, and finally decided on an optimum FOV radius of 30\u2032. The PSPC detector coordinates have a pixel size of 0.934\u2033. We thus removed all X-ray events from the dataset, which are further than 1925 pixels from the PSPC centre pixel coordinate [4119,3929]. A similar cut had to be applied to the modified PSPC instrument map (MOIMP), which is used later for the construction of the survey exposure map.","Citation Text":["Boese 2000"],"Functions Text":["Taking into account all these effects, the HPR of the overall RASS PSF is 84\u2033"],"Functions Label":["Background"],"Citation Start End":[[937,947]],"Functions Start End":[[858,935]]}
{"Identifier":"2021MNRAS.506.3313G__Gao_&_Ho_2017_Instance_1","Paragraph":"The B\/T ratio of galaxies has been traditionally measured by modelling their surface brightness profile. Such modelling allows one to find the light-profiles of individual photometric components of galaxies and then measure their properties such as luminosity, shape, size etc. The earliest such works aimed at photometric decomposition of galaxies involved fitting for the azimuthally averaged 1D surface brightness profile (Kent 1985). However, this method leads to systematic errors as creation of 1D light profile by either averaging over the galaxy image or by taking a single cut along the major axis do not properly take into account non-axisymmetric features like a bar or isophotal twist (for more details see Gao & Ho 2017). Due to these shortcomings, techniques to fit the whole 2D image of the galaxy, pixel-by-pixel, were developed (Wadadekar, Robbason & Kembhavi 1999). There are now a number of implementations of 2D bulge disc decomposition available \u2013 gim2d: Simard (2010), galfit: Peng et al. (2002), imfit: Erwin (2015), profit: Robotham et al. (2016), budda: de Souza, Gadotti & dos Anjos (2004), gasp2d: M\u00e9ndez-Abreu et al. (2008) \u2013 which fit the 2D images of galaxies. They usually employ different algorithms and methods to find the best-fitting model for a galaxy image. In the last two decades, some pipeline codes have been developed (pymorph: Vikram et al. (2010), galapagos: Barden et al. (2012)) to carry out bulge disc decomposition for large galaxy samples, in a mostly automated fashion. This 2D bulge-disc decomposition approach, while being powerful, is time consuming. The employed fitting algorithm and code carries out a search in parameter space defined by various model parameters and then finds the best fit. Most of the time, one also needs to carry out additional steps (masking other objects in the field and providing PSF images) before the actual fitting of galaxies can commence. Thus estimating the B\/T ratio for a single galaxy involves a significant amount of time and computational cost. More importantly, the fitting procedure scales only linearly with the size of the galaxy sample. This limitation will be a major problem in the era of next generation sky surveys (e.g. LSST, Euclid sky surveys) which will observe several billion galaxies.","Citation Text":["Gao & Ho 2017"],"Functions Text":["However, this method leads to systematic errors as creation of 1D light profile by either averaging over the galaxy image or by taking a single cut along the major axis do not properly take into account non-axisymmetric features like a bar or isophotal twist (for more details see"],"Functions Label":["Background"],"Citation Start End":[[719,732]],"Functions Start End":[[438,718]]}
{"Identifier":"2021MNRAS.500.1772N__Shibata,_Kiuchi_&_Sekiguchi_2017_Instance_1","Paragraph":"While these early studies demonstrated the viability of neutron star mergers as a major r-process site, they identified only one ejection channel: \u2018dynamical ejecta\u2019 that are tidally flung out by gravitational torques. Since they are never substantially heated, these ejecta carry their original \u03b2 \u2212equilibrium electron fraction from the original neutron star, Ye \u2248 0.05, and this enormous neutron-richness allows them to undergo a \u2018fission cycling\u2019 process (Goriely, Bauswein & Janka 2011; Korobkin et al. 2012), which produces a very robust r-process abundance distribution close to the solar pattern for A \u2265 130, but hardly any lighter r-process elements. Oechslin, Janka & Marek (2007) pointed out that there is a second channel of mass ejection that also happens on a dynamical time-scale: shock-heated matter from the interface where the stars come into contact. As of today, many more mass ejection channels have been discussed: matter that becomes unbound on secular time-scales (\u223c1 s) from the post-merger accretion torus (Beloborodov 2008; Metzger, Piro & Quataert 2008; Fernandez & Metzger 2013; Fernandez et al. 2015; Just et al. 2015; Siegel & Metzger 2017, 2018; Fernandez et al. 2019; Miller et al. 2019a), as MHD-driven winds (Siegel & Ciolfi 2015) and by viscous effects (Shibata, Kiuchi & Sekiguchi 2017; Radice et al. 2018a; Shibata & Hotokezaka 2019) from a long-lived neutron star merger remnant. Similar to the case of proto-neutron stars, the enormous neutrino luminosities (>1053 erg s\u22121) after a neutron star merger can also drive substantial matter outflows (Ruffert et al. 1997; Rosswog & Ramirez-Ruiz 2002; Dessart et al. 2009; Perego et al. 2014; Martin et al. 2015; Radice et al. 2018b). The secular torus ejecta contain approximately 40 per cent of the initial torus mass and, although the latter may vary substantially from case to case, they likely contribute the lion\u2019s share to the total ejecta mass. Due to their different thermal histories and exposure times to neutrinos, the ejecta channels can have different electron fractions Ye and therefore different nucleosynthesis yields.1 For electron fractions below a critical value, $Y_{\\rm e}^{\\rm crit}\\approx 0.25$ (Korobkin et al. 2012; Lippuner & Roberts 2015), lanthanides and actinides are efficiently produced, which, due to their open f-shells, have particularly high bound\u2013bound opacities (Barnes & Kasen 2013; Kasen, Badnell & Barnes 2013; Tanaka & Hotokezaka 2013; Tanaka et al. 2020) and therefore lead to red transients that peak days after the merger. Ejecta with electron fractions above $Y_{\\rm e}^{\\rm crit}$, in contrast, only produce \u2018lighter\u2019 elements with lower opacities and thus result in bluer transients that peak after about 1 d. Opaque, low-Ye ejecta blocking the view on high-Ye ejecta can lead to a \u2018lanthanide curtaining\u2019 effect (Kasen, Fern\u00e1ndez & Metzger 2015; Wollaeger et al. 2018), which will efficiently block blue light. Therefore, it is important to understand the layering, dynamics, interaction and potential mixing of different ejecta channels.","Citation Text":["Shibata, Kiuchi & Sekiguchi 2017"],"Functions Text":["As of today, many more mass ejection channels have been discussed:","and by viscous effects"],"Functions Label":["Background","Background"],"Citation Start End":[[1289,1321]],"Functions Start End":[[869,935],[1265,1287]]}
{"Identifier":"2019AandA...623A..75V__Tetarenko_et_al._2018a_Instance_1","Paragraph":"The X-ray transient MAXI J1820+070 was first detected on 2018 March 11 (Kawamuro et al. 2018) by the Monitor of All-sky X-ray Image (MAXI, Matsuoka et al. 2009) and was associated with the optical transient ASASSN-18ey (Denisenko 2018; Tucker et al. 2018). In the X-rays, the source flux exceeded 3 Crabs (Bozzo et al. 2018; Mereminskiy et al. 2018), and in the optical, the source reached a magnitude of mV\u2004=\u200412\u2005\u2212\u200513 (Littlefield 2018; Russell et al. 2018). The parallax of the source \u03c0\u2004=\u20040.3\u2005\u00b1\u20050.1 mas was presented in the Gaia DR2 catalogue (Gaia Collaboration 2018). This corresponds to a distance of \n\n\n\n3\n.\n\n9\n\n\u2212\n1.3\n\n\n+\n3.3\n\n\n\n\n$ 3.9^{+3.3}_{-1.3} $\n\n\n kpc (Gandhi et al. 2018a). This unusually bright event allows a detailed investigation of multi-wavelength spectral and timing properties. The course of the outburst was monitored in radio (Trushkin et al. 2018; Polisensky et al. 2018), sub-millimeter (Tetarenko et al. 2018a), optical (Baglio et al. 2018), X-rays (Uttley et al. 2018), and \u03b3-rays (Bozzo et al. 2018; Kuulkers et al. 2018). Because the object was sufficiently bright even for small telescopes, the target was almost continuously monitored, and a rich variety of phenomena was observed. Fast variability and powerful flares in the optical and infrared (Littlefield 2018; Sako et al. 2018; Gandhi et al. 2018b; Casella et al. 2018), optical, and X-ray quasi-periodic oscillations (Mereminskiy et al. 2018; Yu et al. 2018a,b; Buisson et al. 2018; Zampieri et al. 2018) as well as low linear polarisation (Berdyugin et al. 2018) were detected in the source. A 17 h photometric period was recently reported (Patterson et al. 2018) and was tentatively associated with the orbital or superhump period (previously, the source showed a 3.4 h periodicity, Richmond 2018). The X-ray spectral and timing properties as well as the optical-to-X-ray flux ratio suggests that the source is a black hole binary (Baglio et al. 2018; Mereminskiy et al. 2018).","Citation Text":["Tetarenko et al. 2018a"],"Functions Text":["The course of the outburst was monitored in","sub-millimeter"],"Functions Label":["Background","Background"],"Citation Start End":[[913,935]],"Functions Start End":[[799,842],[897,911]]}
{"Identifier":"2019MNRAS.490.5739X__Zahn_et_al._2011_Instance_1","Paragraph":"It is generally believed that reionization started first in high-density regions, where the first luminous objects formed first. In the \u2018bubble model\u2019 of reionization (Furlanetto, Zaldarriaga & Hernquist 2004), the amount of star formation and the resulting ionizing photons are estimated using the excursion set model. Based on this idea, the so-called \u2018seminumerical simulations\u2019 have been developed. For example, Mesinger, Furlanetto & Cen (2011) developed the 21cmFAST,1 to simulate the evolution of the 3D density, ionization, and 21cm brightness temperature fields efficiently. The \u2018bubble model\u2019 considers spherical regions of increasingly smaller scales and identifies ionized bubbles by comparing the cumulative number of ionizing photons produced within the region with the number consumed in reionization process. It has been demonstrated that the statistical predictions of the \u2018bubble model\u2019 and the 21cmFAST agree fairly well with radiative\u2013hydrodynamic simulations (Zahn et al. 2007; Mesinger et al. 2011; Zahn et al. 2011), at least when recombination, feedback, etc. are ignored. However, during most epochs of reionization, the topology of the ionization field is much more complicated than the isolated bubbles configuration. The bubbles start to connect with each other as early as when the global ionized fraction is just about 10\u2009per\u2009cent, and the Universe starts the percolation process when the ionized fraction gets \u223c30\u2009per\u2009cent (see e.g. Furlanetto & Oh 2016; Chen et al. 2018). Inspired by the bubble model, the 21cmFAST determines the ionization state of each point by comparing the expected ionizing photon production in the surrounding region with the required number, but it allows non-spherical geometry for the ionized regions. In order to give a better description of the evolution of neutral regions after percolation, Xu et al. (2014) developed the so-called \u2018island model\u2019, assuming isolated neutral islands topology. The island model also takes into account an ionizing background that is inevitable during the late EoR (Furlanetto & Oh 2005; McQuinn, Oh & Faucher-Gigu\u00e8re 2011; Emberson, Thomas & Alvarez 2013). Based on the \u2018island model\u2019, a seminumerical code named islandFAST was developed to mimic the islands evolution during the last stage of reionization (Xu et al. 2017). In the islandFAST, the effect of small-scale absorbers is taken into account empirically by adopting a fitting formula for the evolution of mean free path (MFP) of ionizing photons (Songaila & Cowie 2010), based on the observed number density of Lyman limit systems up to redshift 6. Before the completion of reionization, the MFP is limited by both the underdense islands and the overdense absorbers. The evolution of the ionization field and the intensity of the ionizing background are derived self-consistently by an iterative procedure to ensure convergence in the total effective MFP of the ionizing photons.","Citation Text":["Zahn et al. 2011"],"Functions Text":["It has been demonstrated that the statistical predictions of the \u2018bubble model\u2019 and the 21cmFAST agree fairly well with radiative\u2013hydrodynamic simulations","at least when recombination, feedback, etc. are ignored."],"Functions Label":["Similarities","Compare\/Contrast"],"Citation Start End":[[1021,1037]],"Functions Start End":[[825,979],[1040,1096]]}
{"Identifier":"2021AandA...655A.111K__Bovy_et_al._2012_Instance_1","Paragraph":"Over the last decade, the radial and vertical dependences of the metallicity-alpha-element distribution have been studied in more and more detail with increasingly larger samples (e.g., Bensby et al. 2011; Anders et al. 2014; Nidever et al. 2014; Hayden et al. 2015; Queiroz et al. 2020). Figure 6 is mostly consistent with similar plots shown in the above papers. In the inner 10 kpc, it displays two over-densities, a high alpha-element (here [Mg\/Fe]), and a low one. Between Rg\u2004=\u20046 and 10 kpc, the two over-densities define two different sequences. In Appendix E, we note that when the sample is restricted to a \u00b1500 pc layer around the Galactic plane, two close but separated sequences are observed in the Rg\u2004\u2208\u2004[4,\u20066] kpc interval. Because of their scale height (Bovy et al. 2012), kinematics (Bensby et al. 2003), and age properties (Haywood et al. 2013), these two sequences are associated with the thick disc (high-alpha) and thin disc (low-alpha), respectively. Moving inward of Rg\u2004=\u20044\u2005\u2212\u20056 kpc, Fig. 6 shows that the two over-densities connect through a zone of lower density to form a single sequence. This is in agreement with the observations of Hayden et al. (2015), Bensby et al. (2017), Zasowski et al. (2019), Bovy et al. (2019), and Lian et al. (2020a,b), who also report a single sequence in the inner disc and\/or in the bulge\/bar area. Conversely, Rojas-Arriagada et al. (2019) and Queiroz et al. (2020) observe two sequences in the inner regions. In Appendix F, we compare the distributions of different APOGEE DR16 alpha elements in the ([Fe\/H], [\u03b1\/Fe]) plane (restricting the sample to the stars contained in the Rg\u2004\u2208\u2004[0,\u20062] kpc interval). The different elements produce different patterns: the global alpha-element abundance5 and oxygen show a double sequence, while magnesium, silicon, and calcium present a single sequence. This could explain, at least partly, why Queiroz et al. (2020), who use a combined \u03b1-element abundance, observe a double sequence, while we see a single one with magnesium. However, this does not explain the discrepancy with Rojas-Arriagada et al. (2019), who also used magnesium. Beyond Rg\u2004=\u200410 kpc, the high-alpha sequence gradually vanishes. This is in agreement with the finding that the thick disc has a shorter scale length than the thin disc (Bensby et al. 2011; Cheng et al. 2012; Bovy et al. 2012). It should be emphasised that in this paragraph the term \u2018sequence\u2019 is used in the geometrical sense. It does not presuppose the number of chemical tracks that form the sequence or sequences. In particular, based on Fig. 6, it can not be excluded that the single geometrical sequence observed in the inner disc be made of two chemical tracks, with the low-alpha one restricted to a narrow metallicity range. We discuss and propose an interpretation of the inner disc sequence in Sect. 5.","Citation Text":["Bovy et al. 2012","Bovy et al. 2012"],"Functions Text":["Because of their scale height","these two sequences are associated with the thick disc (high-alpha) and thin disc (low-alpha), respectively.","This is in agreement with the finding that the thick disc has a shorter scale length than the thin disc"],"Functions Label":["Uses","Uses","Similarities"],"Citation Start End":[[767,783],[2337,2353]],"Functions Start End":[[736,765],[861,969],[2193,2296]]}
{"Identifier":"2017ApJ...845..160P__Lyubarsky_2009_Instance_1","Paragraph":"It is a pressing question as to how the radiation that is observed in relativistic jets in active galactic nuclei (AGNs) is generated (e.g., Blandford & K\u00f6nigl 1979; Marscher 1980; Zensus 1997; Laing & Bridle 2002; Honda 2010; Levinson & Rieger 2011; Mo\u015bcibrodzka et al. 2011; Ito et al. 2013; Mason et al. 2013; Potter & Cotter 2013; Hovatta et al. 2014; Scott & Stewart 2014; Shih & Stockton 2014; Wang et al. 2014; Turner & Shabala 2015; Asada et al. 2016; Hirotani et al. 2016; Koay et al. 2016; Khabibullin et al. 2016; Prieto et al. 2016). Although there is a common consensus that the emitters are energetic particles, how these particles are accelerated to such high energies, how they dissipate their energy, and how they are transported with the jets themselves are still the subject of feverish investigation (e.g., Blandford & Eichler 1987). It is argued that relativistic outflows from black holes are associated with accretion flows (Blandford 1976; Fender et al. 2004; Meier 2005; Ferreira et al. 2006; Trump et al. 2011; Pu et al. 2012; Wu et al. 2013; Ishibashi et al. 2014; Sbarrato et al. 2014). In the case of collimated relativistic jets, magnetic fields must play an important role (Camenzind 1986a, 1986b, 1987; Fendt & Greiner 2001; Vlahakis & K\u00f6nigl 2004; Komissarov et al. 2007; Lyubarsky 2009; Nakamura & Asada 2013; Homan et al. 2015), and it is argued that jets are powered at the expense of the black hole, wherein energy is extracted from a reservoir of rotational energy from the black hole itself, either by electromagnetic means (Blandford & Znajek 1977; Komissarov 2004, 2005; Toma & Takahara 2014) or through magnetohydrodynamical processes (Phinney 1983; Takahashi et al. 1990; Koide et al. 2002; McKinney & Gammie 2004; Hawley & Krolik 2006). Models have been proposed for both of these cases, and they both in principle possess certain testable predictions. In particular, for the latter, numerical GRMHD simulations (e.g., McKinney 2006) and analytical GRMHD studies (e.g., Takahashi et al. 1990; Pu et al. 2016) consistently show the presence of a stagnation or separation surface (a separatrix). This surface separates the (inner) inflow region from the (outer) outflow region, both of which follow the same global, black-hole-threading magnetic field lines. The relatively slow radial velocities near the stagnation surface imply a high concentration of fluid particles. If energetic particles are injected in the vicinity of the stagnation surface or near the black hole event horizon, they must accumulate in high concentrations near the stagnation surface, provided that the cooling timescale is not significantly shorter than the dynamical timescale of the jet fluid flow. This surface is a unique feature of relativistic GRMHD jets and in contrast to an ideal force-free magnetic jet (e.g., McKinney & Narayan 2007; Tchekhovskoy et al. 2008; Broderick & Loeb 2009).","Citation Text":["Lyubarsky 2009"],"Functions Text":["In the case of collimated relativistic jets, magnetic fields must play an important role"],"Functions Label":["Motivation"],"Citation Start End":[[1305,1319]],"Functions Start End":[[1115,1203]]}
{"Identifier":"2019ApJ...875L...7D__Saikia_et_al._2018_Instance_1","Paragraph":"The first transit of HD 21749b was identified by both the MIT Quick Look Pipeline (which searches for planet candidates in the 30 minutes Full Frame Images) and the Science Processing Operations Center (SPOC) pipeline based at the NASA Ames Research Center (Jenkins et al. 2016). No other matching transits were found in the publicly released data from sectors 1 and 2. After TOI 186.01 was alerted, we searched for archival spectroscopy of this very bright star and found 59 High-accuracy Radial-velocity Planet Searcher (HARPS) radial velocities (RVs) in the European Southern Observatory (ESO) archive (see Section 2.4). A periodogram of these RVs showed a clear signal at 35.57 days, but the TESS photometry and the RHK index (Boro Saikia et al. 2018) indicate a stellar rotation period of around 35 days, calling for caution. If the strongest period in the RVs did correspond to the planet, then we expected to see additional transits in sectors 3 and 4. Once the sector 3 data were released, we discovered that a momentum dump29\n\n29\n\u201cMomentum dumps\u201d consist of resetting the momentum wheel speed every 2.5\u20133 days and are used to mitigate the noisier-than-expected measurements of the spacecraft momentum wheel speeds at higher wheel speeds (see https:\/\/archive.stsci.edu\/files\/live\/sites\/mast\/files\/home\/missions-and-data\/active-missions\/tess\/_documents\/TESS_Instrument_Handbook_v0.1.pdf for details). Momentum dumps require brief interruptions to Fine Pointing mode, during which an increase in the flux dispersion is noticeable in the science data, so data acquired during these intervals are excluded from our analysis.\n occurred approximately 35.6 days after the sector 1 transit (see Figure 1). We did not let this unexpected turn of events foil our search efforts, and upon close inspection of the light curve we succeeded in recovering a partial transit (including egress) immediately following the momentum dump. Finally, we observed a third transit in sector 4, thus allowing for a robust ephemeris determination (see Section 3.3). Serendipitously, when applied to the first three sectors of the HD 21749 light curve, the SPOC pipeline yielded an additional planet candidate with a period of 7.9 days (TOI 186.02).","Citation Text":["Boro Saikia et al. 2018"],"Functions Text":["A periodogram of these RVs showed a clear signal at 35.57 days, but the TESS photometry and the RHK index","indicate a stellar rotation period of around 35 days, calling for caution."],"Functions Label":["Differences","Differences"],"Citation Start End":[[731,754]],"Functions Start End":[[624,729],[756,830]]}
{"Identifier":"2019MNRAS.487.4473M__Hosokawa_&_Omukai_2009_Instance_1","Paragraph":"The simulation model Run1-hr of Meyer et al. (2018) begins with the gravitational collapse of $100\\, \\rm M_{\\odot }$ of pre-stellar rotating molecular material on to the stellar embryo. At the end of the free-fall collapse phase, the infalling material ends on a centrifugally-balanced disc, from which the gas is subsequently transferred to the growing protostar. Fig. 1 plots the accretion rate history on to the central massive protostar (solid blue line, in $\\rm M_{\\odot }\\, \\rm yr^{-1}$) together with the evolution of the protostellar mass (dashed red line, in $\\rm M_{\\odot }$) that is calculated as the integrated disc-to-star mass transfer rate through the sink cell. The vertical thin black line indicates the onset of disc formation, when the free-fall collapse of the envelope material on to the protostar stops and the star begins to gain its mass exclusively via accretion from its surrounding disc. After the initial infall of material, the collapse of the parent pre-stellar core material generates an initial increase of the mass flux through the inner boundary at a time ${\\approx } 2\\, \\rm kyr$. The accretion rate then reaches the standard value predicted for MYSOs of $10^{-3}\\, \\rm M_{\\odot }\\, \\rm yr^{-1}$ (Hosokawa & Omukai 2009) up to the onset of the disc formation happening at ${\\approx } 12\\, \\rm kyr$. Variabilities in the accretion flow begin right after the disc formation and the accretion rate history exhibit numerous peaks of growing intensity as the MYSO becomes heavier. This is not caused by different inner boundary conditions, but simply reflects the time-dependent azimuthal anisotropies induced in the accretion flow by the disc evolution. The efficient gravitational instabilities produce complex substructures in the disc, such as overdense spiral arms in which gaseous clumps of various morphologies form at radii of ${\\sim } 100\\, \\rm au$ and inward-migrate down to the central protostar. This produces luminosity outbursts via the mechanism revealed in Meyer et al. (2017) for massive stars. These outbursts are responsible for step-like increases in the stellar mass evolution (see thick dotted red line) due to the fast accretion of dense circumstellar material inside the sink cell. A more detailed description of the evolution of circumstellar discs around young massive stars irradiating their self-gravitating discs can be found in our precedent study (Meyer et al. 2018). In Fig. 1, several magenta dots mark the time instances of the chosen simulation snapshots considered in this study. Note that the young star becomes, by definition, a massive object when $M_{\\star }=8\\, \\rm M_{\\odot }$. Hence, our selected disc models are all in the high-mass regime.","Citation Text":["Hosokawa & Omukai 2009"],"Functions Text":["The accretion rate then reaches the standard value predicted for MYSOs of $10^{-3}\\, \\rm M_{\\odot }\\, \\rm yr^{-1}$","up to the onset of the disc formation happening at ${\\approx } 12\\, \\rm kyr$."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1232,1254]],"Functions Start End":[[1116,1230],[1256,1333]]}
{"Identifier":"2019AandA...621A.124B__Leconte_et_al._2010_Instance_1","Paragraph":"A large proportion of systems where one planet or more is orbiting closer to its host star than Mercury to the Sun have been observed. Tidal interactions play a key role in the orbital configuration of these very compact systems since they are likely to circularize orbits, align spins, and synchronize periods (Zahn 1977; Mathis & Remus 2013; Ogilvie 2014). These interactions consists in an exchange of angular momentum between the orbit and the spins of the celestial bodies. This exchange is the consequence of the dissipation of tidal flows. The kinetic energy of these tidal flows is converted into heat through tidal dissipation. Since the planet is synchronized within a timescale of a few thousand years, the stellar tide drives the secular orbital evolution (Guillot et al. 1996; Rasio et al. 1996; Leconte et al. 2010). In this work, we neglect the impact of the dissipation in the radiative zone. In stellar convection zones, there are two kinds of tides and both are dissipated by the turbulent friction applied by convective eddies. On the one hand, the equilibrium tide is the large-scale velocity field associated with tidal deformation, the so-called tidal bulge. This nonwave-like entity corresponds to the hydrostatic adjustment of the star to the gravitational perturbation (Zahn 1966; Remus et al. 2012). The friction applied by convective motions delays the response of the star to the perturbation (e.g., Zahn 1989; Ogilvie & Lesur 2012; Mathis et al. 2016). This results in a lag angle between the axes of the tidal bulge and the line of centers. This angle increases with dissipation magnitude. Hansen (2012) calibrated its value for several stellar masses by constraining the dissipation using observations of planetary systems. Since lower-mass stars have deeper convective envelopes, they dissipate more energy than higher-mass stars. On the other hand, in rotating bodies such as stars, at low tidal frequencies, the Coriolis acceleration acting on this equlibrium tide excites inertial modes (Ogilvie & Lin 2007). Their ensemble, the dynamical tide, constitutes the wavelike part of the tidal response. Its dissipation strongly depends on internal structure since it arises from the reflection of inertial modes on the radiative, stably stratified core (Ogilvie 2013; Mathis 2015). The dynamical tide may also vary over several orders of magnitude with rotation since inertial waves are restored by the Coriolis force. At low frequencies, dissipation of the dynamical tide is several orders of magnitude higher than the dissipation of the equilibrium tide (Ogilvie & Lin 2007).","Citation Text":["Leconte et al. 2010"],"Functions Text":["Since the planet is synchronized within a timescale of a few thousand years, the stellar tide drives the secular orbital evolution"],"Functions Label":["Background"],"Citation Start End":[[809,828]],"Functions Start End":[[637,767]]}
{"Identifier":"2018ApJ...866L...1S__Pecharrom\u00e1n_et_al._1999_Instance_4","Paragraph":"It was found that the complex dielectric function from Pecharrom\u00e1n et al. (1999) for the sample obtained by heating bayerite at 1273 K, assuming a spheroid with depolarization parameters of (0.35, 0.003), produced an opacity with 11, 20, 28, and 32 \u03bcm features, so this component was included in the models. However, with only this component, the observed 20 \u03bcm features in the residual spectra were found to be wider than those in the models. By adding the opacity of the sample obtained by heating boehmite at 1173 K, the width of the 20 \u03bcm feature could be matched. This was done using the complex dielectric function for the sample obtained by heating boehmite at 1173 K from Pecharrom\u00e1n et al. (1999), assuming a spheroid with depolarization parameters of (0.35, 0.035). The complex dielectric functions of the samples obtained by heating bayerite and boehmite to various temperatures (Pecharrom\u00e1n et al. 1999) were derived by modeling the reflectance spectra of pellets obtained by pressing powders of these materials under great pressure. This method required Pecharrom\u00e1n et al. (1999) to assume an effective medium theory, such that a pellet is a mixture of one of their samples with a matrix of air. Pecharrom\u00e1n et al. (1999) noted that heating bayerite at 500\u00b0C eliminates the XRD pattern of bayerite, and they note that at 700\u00b0C, the infrared reflectance spectrum of the boehmite sample no longer shows OH\u2212 stretching bands. This must mean that the samples obtained from heating bayerite at 1273 K and from heating boehmite at 1173 K are no longer bayerite or boehmite, respectively. XRD performed by Pecharrom\u00e1n et al. (1999) of the sample of bayerite prepared at 1273 K suggests only \u03b8-alumina was present, and their infrared and NMR spectroscopy confirms this. XRD of their sample obtained from heating boehmite to 1173 K (Pecharrom\u00e1n et al. 1999) suggests \u03b4-alumina to be present, though some amounts of \u03b8-alumina and \u03b1-alumina are present, as they deduce from XRD and infrared and NMR spectroscopy.","Citation Text":["Pecharrom\u00e1n et al. (1999)"],"Functions Text":["This method required","to assume an effective medium theory, such that a pellet is a mixture of one of their samples with a matrix of air."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1067,1092]],"Functions Start End":[[1046,1066],[1093,1208]]}
{"Identifier":"2022MNRAS.517.5541T__Mooley_et_al._2018a_Instance_1","Paragraph":"GRB 170817A, the short GRB directly associated with the first gravitational wave signal GW170817 from a binary neutron star merger, is a great example presenting afterglow spectra with a beautiful single power-law extending from radio to X-ray (Abbott et al. 2017a, b; Goldstein et al. 2017; Savchenko et al. 2017). A successful launch of a relativistic jet in GRB 170817A is supported by the detection of superluminal motion of a compact radio source (Mooley et al. 2018b; Ghirlanda et al. 2019). The jet was observed from off-axis (i.e. the jet axis and the line of sight was misaligned), leading to a faint gamma-ray emission (Abbott et al. 2017c; Ioka & Nakamura 2018, 2019; Salafia & Ghirlanda 2022, and references therein) and early rising afterglow light curves (Alexander et al. 2017; Hallinan et al. 2017; Margutti et al. 2017; Troja et al. 2017). The slow rising afterglow light curves also rejected a uniform top-hat jet (Mooley et al. 2018a; Troja et al. 2018) and revealed a launch of a structured jet, while various jet structures can explain the observations (see, however, Lamb, Levan & Tanvir (2020) for an alternative scenario with a refreshed shock in an off-axis uniform jet). Often assumed structured jets are a Gaussian jet or a power-law jet (D\u2019Avanzo et al. 2018; Gill & Granot 2018; Lamb & Kobayashi 2018; Lyman et al. 2018; Resmi et al. 2018; Ghirlanda et al. 2019; Lamb et al. 2019; Troja et al. 2019, 2020; Ryan et al. 2020). Other non-trivial jet structures such as hollow-cone jets and spindle jets are also possible candidates as recently discovered by Takahashi & Ioka (2020, 2021), while the consistency of a hollow-cone jet was also confirmed by Nathanail et al. (2020, 2021). As predicted by Troja et al. (2018), the rapid decline of the afterglow after the peak (Mooley et al. 2018c; Lamb et al. 2019; Hajela et al. 2020; Makhathini et al. 2021; Troja et al. 2020; Balasubramanian et al. 2021) revealed a distinctive signature of a successfully launched structured jet, setting it apart from a chocked jet scenario. The observed electron power-law index derived from the afterglow spectrum falls within the range predicted in a theory of particle acceleration at trans-relativistic shock (Keshet & Waxman 2005) as mentioned by Margutti et al. (2018) and is consistent with being constant in time within observational errors (D\u2019Avanzo et al. 2018; Dobie et al. 2018; Margutti et al. 2018; Fong et al. 2019; Kilpatrick et al. 2022).","Citation Text":["Mooley et al. 2018a"],"Functions Text":["The slow rising afterglow light curves also rejected a uniform top-hat jet"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[933,952]],"Functions Start End":[[857,931]]}
{"Identifier":"2016MNRAS.456.1901H__Yoo_&_Miralda-Escud\u00e9_2004_Instance_1","Paragraph":"The growth of supermassive BHs is also altered when considering non-Gaussianities. After deriving the merger history of the most massive haloes at z = 6.5 in both the Gaussian and non-Gaussian simulations, we study the evolution of BHs in massive haloes down to z = 6.5. To investigate the cumulative effect over cosmic times on the BHs assembly, we model the growth of BHs in three different ways. Different probabilities for a halo of hosting a seed BH, and different accretion models (either each BH accretes at the Eddington limit for a dynamical time after a major merger or using an accretion rate based on a distribution probability derived from a large-scale hydrodynamical simulation) are adopted. We have not included in our models the effects of \u2018kicks\u2019 caused by asymmetric emission of gravitational waves, which have been proposed to be possibly responsible for ejecting BHs from haloes with shallow potential wells, thus halting or reducing the growth of high-redshift BHs hosted in small haloes (e.g. Yoo & Miralda-Escud\u00e9 2004; Volonteri & Rees 2006; Tanaka & Haiman 2009). This effect, however, seems to affect less than 10 per cent of binaries and it becomes negligible for BH mergers at z  10 (Volonteri & Rees 2006). We find that non-Gaussianities imply a larger number of massive BHs and also an increase in the mean BH mass (up to 0.36 in the most favourable experiment). A population of supermassive BHs will then grow faster and to higher masses in a universe with scale-dependent non-Gaussian primordial density fluctuations. If the seed masses are similar to those of Pop III star remnants, BHs will not be able to grow above few \u00d7 105 M\u2299 by z = 6. However, our simulations do not resolve mini-haloes, and we may underestimate the growth of seeds at earlier times. We argue that, in a simulation resolving mini-haloes, BHs would have formed earlier through the Pop III remnant scenario, leading to a longer time for them to grow in mass. If we assumed that Pop III remnant seeds with mass 100\u2009M\u2299 form at z \u223c 30 in haloes unresolved in our simulations, they would have grown, assuming, optimistically, constant growth at the Eddington rate (but see Johnson & Bromm 2007; Alvarez, Wise & Abel 2009; Milosavljevi\u0107, Couch & Bromm 2009; Park & Ricotti 2011) to \u223c103 M\u2299 by z = 18, where we start our analysis. The final BH mass at z = 6 would then be \u223c one order of magnitude larger, a few \u00d7 106\u2009M\u2299, still short of the \u223c109 M\u2299 required. The very limited growth obtained for the Pop III remnant case suggests that large seeds or super-Eddington accretion (see Volonteri, Silk & Dubus 2015, and references therein) may be necessary for successful BH growth. We have done the same experiments on BH growth starting with initial 105 M\u2299 BH masses (not shown in the paper, but see Section 4). In this case we found that it is much easier for BHs to grow to higher BH masses, but still only to several 108\u2009M\u2299. This is not unexpected, because our simulation box does not contain the very rare and biased DM haloes with masses \u223c1013 M\u2299 believed to be hosting these extreme BHs.","Citation Text":["Yoo & Miralda-Escud\u00e9 2004"],"Functions Text":["We have not included in our models the effects of \u2018kicks\u2019 caused by asymmetric emission of gravitational waves, which have been proposed to be possibly responsible for ejecting BHs from haloes with shallow potential wells, thus halting or reducing the growth of high-redshift BHs hosted in small haloes (e.g.","This effect, however, seems to affect less than 10 per cent of binaries"],"Functions Label":["Differences","Compare\/Contrast"],"Citation Start End":[[1016,1041]],"Functions Start End":[[707,1015],[1089,1160]]}
{"Identifier":"2016ApJ...827...93N__Neuhauser_et_al._2007_Instance_1","Paragraph":"The results presented in this work answer a number of questions regarding the character of the GG Tau A system, while raising others and leaving others untouched. First among the questions raised by our results are the questions of whether the detailed morphology of features in the circumbinary torus as seen in our simulations are actually present in the GG Tau A system. The current best resolution observations of the torus provide tantalizing hints that such features do exist, but a definitive statement that such features are present must await higher resolution observations of quality similar to those described by the ALMA Partnership (2015), for the HL Tau circumstellar disk. Another important question concerns planet formation in multiple systems. Whether or not any theoretical mechanism presently exists to explain the formation of giant planets in binary systems, the fact remains that at least a few binary systems, such as \u03b3 Cephei (Neuhauser et al. 2007), GI86 (Queloz et al. 2000), and HD 41004 (Zucker et al. 2004), do harbor planets. Therefore, some formation mechanism does in fact exist. While we find that accretion into the circumstellar disks occurs rapidly enough so that they can survive for the comparatively long timescales needed to form planets, other conditions, such as temperatures, remain quite unfavorable. What are the mechanisms still missing from our models that permit such objects to form? Finally, our simulations model the evolution of the \u201cA\u201d component of the full GG Tau system over a time span extending over only a tiny fraction of its formation timescale and in only two spatial dimensions. We neglected the full dynamical effects expected to be present in the system, insofar as our results include neither the distant binary \u201cB\u201d component of the GG Tau system, nor the newly discovered tight binary nature of the GG Tau Ab component. Even so, we find large scale morphological changes even over this short time span, and restricted dimensionality and physical system. Given the vigor of the activity over such a short timescale, we would expect activity of similar scale to occur over longer time spans as well, with correspondingly large consequences on the system morphology. To what extent will 3D effects also play a role in the evolution? What will be the end state configuration of the GG Tau system as a whole? Will the components eventually break apart? Merge? Future investigations extending the work presented here will be required in order to answer these questions.","Citation Text":["Neuhauser et al. 2007"],"Functions Text":["Whether or not any theoretical mechanism presently exists to explain the formation of giant planets in binary systems, the fact remains that at least a few binary systems, such as \u03b3 Cephei","do harbor planets. Therefore, some formation mechanism does in fact exist."],"Functions Label":["Uses","Uses"],"Citation Start End":[[952,973]],"Functions Start End":[[762,950],[1038,1112]]}
{"Identifier":"2018ApJ...855...48Q__Nagy_et_al._2017_Instance_1","Paragraph":"The Orion Bar is probably the best studied PDR in our Galaxy. It is located between the Orion Molecular Cloud 1 and the H ii region excited by the Trapezium cluster, and is exposed to an FUV field a few 104 times the mean interstellar radiation field. Owing to its proximity (417 pc, Menten et al. 2007) and nearly edge-on orientation, the Bar provides an ideal laboratory for testing PDR models (e.g., Jansen et al. 1995; Gorti & Hollenbach 2002; Andree-Labsch et al. 2017) and a primary target for observational studies of physical and chemical structures of PDRs (e.g., Tielens et al. 1993; Walmsley et al. 2000; van der Wiel et al. 2009; Arab et al. 2012; Peng et al. 2012; Goicoechea et al. 2016; Nagy et al. 2017). Observations of various molecular spectral lines have shown that the emissions could be better interpreted with an inhomogeneous density structure containing an extended and relatively low density (\n\n\n\n\n\n cm\u22123) medium and a compact and high-density (\n\n\n\n\n\n cm\u22123) component (e.g., Hogerheijde et al. 1995; Young Owl et al. 2000; Leurini et al. 2006, 2010; Goicoechea et al. 2016). However, due to the scarcity of high-resolution observations capable of spatially resolving the density structure, the nature of the high-density clumps or condensations is still not well understood. Lis & Schilke (2003) mapped the Bar in H13CN (1\u20130) with the Plateau de Bure Interferometer (PdBI) at an angular resolution of about 5\u2033, and detected 10 dense clumps. They proposed that the H13CN clumps are in virial equilibrium and may be collapsing to form stars. Goicoechea et al. (2016) performed Atacama Large Millimeter\/submillimeter Array HCO+ (4\u20133) observations of the Bar and detected over-dense substructures close to the cloud edge, and found that the substructures have masses much lower than the mass needed to make them gravitationally unstable. These two interferometric observations both target molecular spectral lines. A high-resolution map of the dust continuum emission of the Bar, which is highly desirable in constraining the mass and density of the dense condensations, is still lacking. Here we report our Submillimeter Array (SMA) observations of the dust continuum and molecular spectral line observations of the Bar.","Citation Text":["Nagy et al. 2017"],"Functions Text":["Owing to its proximity","and nearly edge-on orientation, the Bar provides an ideal laboratory for testing PDR models","and a primary target for observational studies of physical and chemical structures of PDRs (e.g.,"],"Functions Label":["Background","Background","Background"],"Citation Start End":[[702,718]],"Functions Start End":[[252,274],[304,395],[475,572]]}
{"Identifier":"2015AandA...584A.103S__G\u00f6gelein_&_M\u00fcther_2007_Instance_1","Paragraph":"Before leaving this section, in Fig. 8 we display the spatial dependence of the self-consistent neutron and proton density profiles for the optimal solutions in spherical WS cells with average baryon densities nb = 0.0475 fm-3, 0.065 fm-3, and 0.076 fm-3. It is observed that in denser matter the size of the WS cell decreases, as we discussed previously, and that the amount of free neutrons in the gas increases, as expected. It can be seen that the nuclear surface is progressively washed out with increasing average baryon density as the nucleon distributions become more uniform. At high nb the density profile inside the WS cell extends towards the edge of the cell, pointing out that the WS approximation may be close to its limits of validity (Negele & Vautherin 1973; Chamel et al. 2007; Baldo et al. 2007; Pastore et al. 2011; G\u00f6gelein & M\u00fcther 2007; Newton & Stone 2009). Although the proton number Z is similar for the three average baryon densities of Fig. 8, the local distribution of the protons is very different in the three cases. In Fig. 8c the proton density profile extends more than 3 fm farther from the origin than in Fig. 8a, while the central value of the proton density has decreased by more than a factor 2, hinting at the fact that the neutrons have a strong drag effect on the protons. Figure 9 presents the nucleon density profiles obtained for cylindrical and planar geometries at the same average density nb = 0.076 fm-3 as in Fig. 8c. From Figs. 8c (droplets),  9a (rods), and 9b (slabs) we see that the size of the WS cells decreases with decreasing dimensionality, i.e. Rc,droplet>Rc,rod>Rc,slab. At high average densities near the crust-core transition, nucleons inside the WS cell can arrange themselves in such a way that the region of higher density is concentrated at the edge of the cell, leaving the uniform region of lower density in the inner part of the cell. This distribution of nucleons corresponds to the cylindrical tube and spherical bubble configurations. In Figs. 9c and d, we plot the neutron and proton density profiles of the optimal solution for tubes and bubbles at nb = 0.076 fm-3. At equal average density, the size of the cells containing tubes and bubbles is larger than the size of the cells accommodating rods and droplets, respectively, as can be appreciated by comparing Fig. 9a for rods with Fig. 9c for tubes, and Fig. 8c for droplets with Fig. 9d for bubbles. As a consequence of this fact and of the effectively larger value of the integration factors 2\u03c0r and 4\u03c0r2 when the densities are accumulated near the edge of the cell, the total number of nucleons and the atomic number in the tube and bubble cells is about 1.5\u22122 times larger than in their rod and droplet counterparts. The proton fraction xp = Z\/A is, however, practically the same for all geometries. ","Citation Text":["G\u00f6gelein & M\u00fcther 2007"],"Functions Text":["At high nb the density profile inside the WS cell extends towards the edge of the cell, pointing out that the WS approximation may be close to its limits of validity"],"Functions Label":["Background"],"Citation Start End":[[837,859]],"Functions Start End":[[585,750]]}
{"Identifier":"2020AandA...641A.155V__Ceverino_et_al._2010_Instance_1","Paragraph":"It has also become evident that the normalization of the MS rapidly increases with redshift: distant galaxies form stars at higher paces than in the local Universe, at fixed stellar mass (e.g., Daddi et al. 2007; Elbaz et al. 2007; Whitaker et al. 2012; Speagle et al. 2014; Schreiber et al. 2015). This trend could be explained by the availability of copious molecular gas at high redshift (Daddi et al. 2010a; Tacconi et al. 2010, 2018; Scoville et al. 2017a; Riechers et al. 2019; Decarli et al. 2019; Liu et al. 2019a), ultimately regulated by the larger accretion rates from the cosmic web (Kere\u0161 et al. 2005; Dekel et al. 2009a). Moreover, higher SFRs could be induced by an increased efficiency of star formation due to the enhanced fragmentation in gas-rich, turbulent, and gravitationally unstable high-redshift disks (Bournaud et al. 2007, 2010; Dekel et al. 2009b; Ceverino et al. 2010; Dekel & Burkert 2014), reflected on their clumpy morphologies (Elmegreen et al. 2007; F\u00f6rster Schreiber et al. 2011; Genzel et al. 2011; Guo et al. 2012, 2015; Zanella et al. 2019). IR-bright galaxies with prodigious SFRs well above the level of the MS are observed also in the distant Universe, but their main physical driver is a matter of debate. While a star formation efficiency (SFE\u2004=\u2004SFR\/Mgas) monotonically increasing with the distance from the main sequence (\u0394MS\u2004=\u2004SFR\/SFRMS, Genzel et al. 2010, 2015; Magdis et al. 2012; Tacconi et al. 2018, 2020) could naturally explain the existence of these outliers, recent works suggest the concomitant increase of gas masses as the main driver of the starbursting events (Scoville et al. 2016; Elbaz et al. 2018). In addition, if many bright starbursting (sub)millimeter galaxies (SMGs, Smail et al. 1997) are indeed merging systems as in the local Universe (G\u00f3mez-Guijarro et al. 2018, and references therein), there are several well documented cases of SMGs hosting orderly rotating disks at high redshift (e.g., Hodge et al. 2016, 2019; Drew et al. 2020), disputing the pure merger scenario. The same definition of starbursts as galaxies deviating from the main sequence has been recently questioned with the advent of high spatial resolution measurements of their dust and gas emission. Compact galaxies with short depletion timescales typical of SBs are now routinely found on the MS, being possibly on their way to leave the sequence (Barro et al. 2017a; Popping et al. 2017; Elbaz et al. 2018; G\u00f3mez-Guijarro et al. 2019; Puglisi et al. 2019; Jim\u00e9nez-Andrade et al. 2019); or galaxies moving within the MS scatter, due to mergers unable to efficiently boost the star formation (Fensch et al. 2017) or owing to gravitational instabilities and gas radial redistribution (Tacchella et al. 2016).","Citation Text":["Ceverino et al. 2010"],"Functions Text":["Moreover, higher SFRs could be induced by an increased efficiency of star formation due to the enhanced fragmentation in gas-rich, turbulent, and gravitationally unstable high-redshift disks"],"Functions Label":["Background"],"Citation Start End":[[876,896]],"Functions Start End":[[636,826]]}
{"Identifier":"2016AandA...596A.113B__Perryman_et_al._1997_Instance_1","Paragraph":"The Pleiades open cluster does not only offer us a beautiful spectacle during the fall, it is one of the best-studied stellar associations and one of the cornerstones to understand stellar properties and evolution. In fact, the literature includes more than one thousand refereed papers dealing with the Pleiades, most of which use the Pleiades as a reference, just in the last ten years. In spite of this, the Pleiades cluster still has many secret and basic parameters, such as its distance and age, that are not clearly established. Even the census of this cluster is incomplete, although the recent works by Stauffer et al. (2007), Lodieu et al. (2012), and Bouy et al. (2013) have considerably improved the membership list, In fact, Bouy et al. (2015) has increased the number of known members by 50% with public archival data, very accurate proper motions, and multiwavelength photometry (see additional details in Sarro et al. 2014). Regarding its distance, there are currently two different methodologies based on parallaxes from Hipparcos (Perryman et al. 1997) and isochrone fitting, respectively. Pre-Hipparcos distances for the Pleiades range between 125 and 130 pc (see, for instance, Soderblom et al. 1993d), whereas the initial distance derived by Hipparcos is much closer, about 119 pc (van Leeuwen 1999). This last value is significantly different from the distance derived by Pinsonneault et al. (1998), who used color-magnitude diagrams and fitting isochrones and obtained 133.5 \u00b1 1.2 pc. More recently, van Leeuwen (2009), by reanalyzing Hipparcos data, derived a distance of 120.2 \u00b1 1.9 pc. These values should correspond to the distance to the cluster center, whose core radius should be around 3 degrees, which corresponds to 5\u20136 pc. Another trigonometric parallax, based on Hubble Space Telescope data and three members, was found by Soderblom et al. (2005), who obtained 134.6 \u00b1 3.1 pc. Recently, Melis et al. (2014) derived a distance of 136.2 \u00b1 1.2 pc based on an accurate parallax for four bona fide members obtained with the VLBI. This value also agrees with that derived by Galli et al. (in prep.) using accurate proper motions and the convergence point method (137.7 \u00b1 2.5 pc). ","Citation Text":["Perryman et al. 1997"],"Functions Text":["Regarding its distance, there are currently two different methodologies based on parallaxes from Hipparcos","and isochrone fitting, respectively."],"Functions Label":["Background","Background"],"Citation Start End":[[1049,1069]],"Functions Start End":[[941,1047],[1071,1107]]}
{"Identifier":"2018AandA...614A..48B__Keselman_&_Nusser_2012_Instance_1","Paragraph":"The driving mechanisms and chronology of the buildup of bulges in late-type galaxies (LTGs) is an issue of key relevance to our understanding of galaxy evolution. According to our current knowledge on bulge demographics in the local universe, a large fraction of LTGs host pseudo-bulges (PBs; e.g., Gadotti 2009; Fisher & Drory 2011; Fern\u00e1ndez Lorenzo et al. 2014) that substantially differ from classical bulges (CBs) in their spectrophotometric and kinematical characteristics. The latter resemble in many respects \u201cold and dead\u201d elliptical galaxies, lacking ongoing star-formation (SF), exhibit a spheroidal shape with inwardly steeply increasing surface brightness profiles (SBPs) being well approximated by the S\u00e9rsic (1963) fitting law with a high (\u22733) exponent \u03b7, show stellar kinematics dominated by velocity dispersion (\u03c3\u22c6) and obey the Kormendy (1977) scaling relations for normal elliptical galaxies (Fisher & Drory 2010). It is observationally established that CBs contain a super-massive black hole (SMBH) with a mass M\u2219 tightly correlating with their stellar mass \n\n${\\cal M}_{\\star,\\textrm{B}}, \\sigma_{*}$M\u22c6,B,\u03c3*\n and optical luminosity (Ho 2008; Kormendy & Ho 2013; see also Ferrarese & Merritt 2000). Traditionally, bulges were thought to invariably form early-on via violent quasi-monolithic gas collapse (Larson 1974) or mergers (Bender et al. 1992; Aguerri et al. 2001; Keselman & Nusser 2012) associated with vigorous nuclear starbursts (Okamoto 2012), with the disk gradually building up around them. Whereas this inside-out galaxy formation scenario appears consistent with important integral characteristics of CBs (e.g., their red colors), it does not offer a plausible explanation for the presence of PBs in present-day LTGs. These generally show ongoing SF, a significant degree of rotational support (Kormendy & Kennicutt 2004, for a review) and flatter\/ellipsoidal shapes with nearly exponential SBPs (\u03b7\u22722; e.g., Drory & Fisher 2007; Fisher & Drory 2010). Even though there is observational evidence that PBs also contain a SMBH (Kormendy et al. 2011; Kormendy & Ho 2013), in some cases revealing itself as an active galactic nucleus (AGN; e.g., Kotilainen et al. 2016; see Kormendy & Ho 2013 for a review), these do not follow the M\u2219 \u2013\u03c3* correlation for CBs, which appears to be consistent with a different formation route. Indeed, the prevailing concept on PB formation is that these entities emerge gradually out of galactic disks through gentle gas inflow spawning quasi-continuous SF and the emergence of a central bulge-like luminosity excess at their centers (e.g., Courteau et al. 1996; Carollo et al. 2001; Kormendy & Kennicutt 2004). Besides bar-driven gas inflow (e.g., Springel & Hernquist 2005), various other mechanisms, such as inward stellar migration, minor mergers with low-mass satellites, or a purely dynamical re-arrangement of the disk (Scannapieco et al. 2010; Guedes et al. 2013; Bird et al. 2012; Roskar et al. 2012; Grand et al. 2014; Halle et al. 2015) have been proposed as further contributors to PB growth along the Gyr-long secular evolution of LTGs.","Citation Text":["Keselman & Nusser 2012"],"Functions Text":["Traditionally, bulges were thought to invariably form early-on via","or mergers","with the disk gradually building up around them.","Whereas this inside-out galaxy formation scenario appears consistent with important integral characteristics of CBs (e.g., their red colors), it does not offer a plausible explanation for the presence of PBs in present-day LTGs."],"Functions Label":["Background","Background","Background","Compare\/Contrast"],"Citation Start End":[[1391,1413]],"Functions Start End":[[1219,1285],[1338,1348],[1475,1523],[1524,1752]]}
{"Identifier":"2019AandA...630A..26M__Wozniakiewicz_et_al._(2012)_Instance_2","Paragraph":"The majority of the particles collected by Stardust are olivine and pyroxene silicates with solar isotopic compositions, which suggests an origin in our solar system rather than an interstellar provenance. These polymineralic particles dominate those made of a single mineral even down to sizes smaller than 100 nm, indicating that the dust composition is surprisingly consistent at different scales and that the smallest subunits of the dust may be as small as tens of nanometers (H\u00f6rz et al. 2006; Zolensky et al. 2006). The sizes of these smallest single mineral impactors are similar to those of the nanocrystals determined by Rietmeijer (1993). As discussed above, they might also be existing in MIDAS dust particles and might be fused into the 100 nm features. Price et al. (2010) and Wozniakiewicz et al. (2012) investigated the sizes of particles smaller than 10 \u03bcm that impacted the aluminum foils of the Stardust probe. The distribution peaks at about 175 nm, but if we assume that the particles areagglomerates of smaller subunits, as indicated by their common polymineralic nature, then the subunit size distribution would peak at sizes below 100 nm (Price et al. 2010). A study of over 450 particles that do not seem to be agglomerates, that is, those that show single mineral impactors of silicate or sulfide, found geometric mean sizes of \n\n$532^{741}_{-310}$\n\n\n\n\n532\n\n\u2212310\n\n+741\n\n\n\n nm for the silicate particles and \n\n$406^{491}_{-222}$\n\n\n\n\n406\n\n\u2212222\n\n+491\n\n\n\n nm for the sulfides (Wozniakiewicz et al. 2013). These sizes are notably larger than the 175 nm (or less) found for the whole dataset. This large spread of subunit sizes could indicate a size distribution with a large width. No fits of these sizedistributions are available, but the figures in Wozniakiewicz et al. (2012) and Price et al. (2010) indicate that the differential sizes may follow a log-normal distribution. When we assume that the smallest subunit sizes are possibly between tens and hundreds of nanometers, the subunit size range found for MIDAS smallest features would be encompassed. The determination of the size distributions for the small Stardust particles and a detailed comparison to the distributions obtained for comet 67P could be the work of an interesting future project.","Citation Text":["Wozniakiewicz et al. (2012)"],"Functions Text":["No fits of these sizedistributions are available, but the figures in","and Price et al. (2010) indicate that the differential sizes may follow a log-normal distribution."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1772,1799]],"Functions Start End":[[1703,1771],[1800,1898]]}
{"Identifier":"2021MNRAS.507..524M__Steidel_et_al._2016_Instance_1","Paragraph":"In particular, stellar wind P-Cygni profiles and photospheric absorption lines are detected with high significance (green and yellow dashed lines in Fig. 2). The detection of photospheric lines in J0121+0025 indicates unambiguously that the UV luminosity is dominated by stellar emission, rather than an AGN. We identify more than ten photospheric features. Some of them are resolved and detected with high significance (e.g. C\u2009ii 1324 \u00c5, O\u2009iv 1343 \u00c5, and S\u2009v 1501 \u00c5; see Fig. 3). We use these to determine the systemic redshift zsys = 3.244 \u00b1 0.001 of J0121+0025. Others are seen in blends from multiple transitions (e.g. Si\u2009iii 1417 \u00c5, C\u2009iii 1427 \u00c5, and Fe\u2009v 1430 \u00c5 at \u03bb0 \u2243 1415\u20131435 \u00c5). These stellar absorption lines are intrinsically weak in star-forming galaxies, with EW0 typically well bellow 1 \u00c5 (e.g. Shapley et al. 2003; Steidel et al. 2016; Rigby et al. 2018). As they are formed in the photospheres of hot stars and are seen in absorption, the background radiation should be dominated by the starlight, otherwise they would not be detected. Even a small contribution of an AGN to the UV continuum (\u2272 25 per\u2009cent), that is featureless in these spectral regions, would make these lines disappear at the SNR of our spectrum. In addition, the observed P-Cygni profiles in N\u2009v 1240 \u00c5 and C\u2009iv 1550 \u00c5 can be also well explained\/modelled by stellar models with a very young age (\u22433 Myr burst; see Fig. 3 and Section 3.2 for details), similar to those seen in other very young starbursts (e.g. Rivera-Thorsen et al. 2019; Vanzella et al. ), some of them also very\/extremely luminous (Vanzella et al. 2018; Marques-Chaves et al. 2020b). While some rare AGNs, such as broad or narrow absorption line QSOs (BAL\/NAL QSOs), can show N\u2009v and C\u2009iv profiles mimic those of stellar P-Cygni, from the combination of a broad emission and a redshifted absorption (see for example Bentz, Osmer & Weinberg 2004; Appenzeller et al. 2005), photospheric lines are not present in the spectra of AGNs.","Citation Text":["Steidel et al. 2016"],"Functions Text":["These stellar absorption lines are intrinsically weak in star-forming galaxies, with EW0 typically well bellow 1 \u00c5 (e.g."],"Functions Label":["Background"],"Citation Start End":[[832,851]],"Functions Start End":[[690,810]]}
{"Identifier":"2018ApJ...863..162M__Liu_et_al._2013_Instance_1","Paragraph":"NLFFF extrapolation provides the reconstructed coronal magnetic field for AR 11158 from 2011 February 13 \u2212 2011 February 15 (Figures 1(d)\u2013(f)). The field lines (yellow lines) within the core of the AR have arcade-like structure with a relatively strong twist mainly near the PIL. These figures show that the magnetic field evolved during this period. Although we did not quantitatively compare the field lines with the observation, in general, the reconstructed coronal field morphologies match with the observations in Figures 1(a)\u2013(c). The general morphologies and the locations of the high-twist fields are also in agreement with many previous studies (Jing et al. 2012; Sun et al. 2012; Dalmasse et al. 2013; Inoue et al. 2013, 2014a; Liu et al. 2013; Wang et al. 2013; Aschwanden et al. 2014; Malanushenko et al. 2014; Zhao et al. 2014). Unlike Zhao et al. (2014) who could identify the twisted flux rope from the topology of the reconstructed coronal field, we could not find an obvious topological signature of a flux rope existing in our NLFFF during our analysis time window. It might be due to the fact that there was little magnetic flux with twist higher than one turn in our NLFFF and it is difficult to topologically define it as a flux rope. However, our result is consistent with other NLFFF results (Jing et al. 2012; Sun et al. 2012; Liu et al. 2013; Wang et al. 2013; Inoue et al. 2014a; Malanushenko et al. 2014). The high-twist region in our result is also in agreement with the region with high helicity flux (Dalmasse et al. 2013) and the location of the flare ribbons (Bamba et al. 2013; Liu et al. 2013), as well as the high current density region (Janvier et al. 2014). Figures 1(g)\u2013(i) show the evolution of the twist distribution map, with the magnetic twist of the field lines plotted at the footpoints of field lines according to a color scale. This shows that the high-twist (strongly right-handed twist corresponding to Tw > 0.5) areas are concentrated in only a limited part of the AR. The high-twist area grew and became even more twisted just before the X2.2 flare (Figure 1(i)). Most parts of the AR have twist values less than 0.25, but near the PIL the twist can reach more than 0.5, even up to about a full turn. This is consistent with the results of Sun et al. (2012) and Inoue et al. (2014a). A high-twist (strong negative\/left-handed twist) area also developed in the eastern part of the AR, which did not exist initially on February 13. Both of these high-twist areas produced several flares. However, here we focus on the flares that resulted from the high-twist core region near the center of the AR, where the M6.6 and X2.2 flares occurred.","Citation Text":["Liu et al. 2013"],"Functions Text":["The general morphologies and the locations of the high-twist fields are also in agreement with many previous studies"],"Functions Label":["Similarities"],"Citation Start End":[[739,754]],"Functions Start End":[[538,654]]}
{"Identifier":"2022AandA...663L...4L___2020_Instance_1","Paragraph":"While previous studies mainly focused on the optical-UV properties of a MAD in RLAGNs, for the first time we try to investigate their X-ray properties in this work. The origin of X-ray emission in RLAGNs is still under debate, which may come from a corona, jet, or both. In observations, there is a big difference between the X-ray properties of radio-quiet AGNs (RQAGNs) and RLAGNs. Firstly, the average X-ray flux in RLAGNs is found to be 2\u20133 times higher than that in RQAGNs (e.g., Zamorani et al. 1981; Wilkes & Elvis 1987; Li & Gu 2021). Secondly, Laor et al. (1997) reported that RLAGNs have harder 2\u201310 kev X-ray spectra than RQAGNs by compiling a sample of 23 quasars observed with ROSAT, which was subsequently confirmed by Shang et al. (2011) with a larger sample. Comparing the X-ray spectrum of 3CRR quasars and that of radio-quiet quasars, Zhou & Gu (2020) also gave a similar result. In addition, the X-ray reflection features of RLAGNs are weaker than those of RQAGNs (Wozniak et al. 1998). All of these results seem to indicate that the contribution of a jet to X-ray spectra cannot be neglected. However, several recent works suggested a totally different result. First, the slope of LUV\u2005\u2212\u2005LX is found to be consistent for RLAGNs and RQAGNs (Zhu et al. 2020, 2021; Li & Gu 2021). Second, Gupta et al. (2018, 2020) discovered that the distributions of X-ray photon spectral indices between RLAGNs and their radio-quiet counterpart are very similar (see Zhu et al. 2021 either). This opposite conclusion may be due to the effect of sample selection. The sample of Gupta et al. (2018, 2020) was X-ray selected (and optically selected for the sample of Zhu et al. 2021), which may lead to the radio jet power being very feeble compared to the bolometric luminosity in most of the RLAGNs. These weakly jetted RLAGNs can therefore have different X-ray photon indices compared to the strong jetted RLAGNs, such as the 3CRR quasars of Zhou & Gu (2020). P15 also indicated that the weakly jetted RLAGNs have similar \u03b1EUV as RQAGNs. However, interestingly, Markoff et al. (2005) demonstrated that both the corona model and the jet model can fit the X-ray data of some Galactic X-ray binaries well and that the jet base may be subsumed to corona in some ways. The 3CRR quasars are low frequency radio selected and have a strong jet on a large scale. However, it is still unclear whether all the objects with a strong jet harbor a MAD, or just containing MAD when jet is firstly launching millions of years ago. We focus on the RLAGNs with an EUV deficit in this work, which should possess a MAD in the inner disk region as suggested by P15. The presence of a MAD surrounding the black hole may bring a remarkable difference to the X-ray emissions since the structure of disk-corona greatly changes in the case of MAD (e.g., Tchekhovskoy et al. 2011; McKinney et al. 2012; White et al. 2019). In theory, it has been suggested that X-ray emission increases when an advection-dominated accretion flow (ADAF) becomes a MAD in its inner region (Xie & Zdziarski 2019). Nevertheless, how MAD affects the disk-corona corresponding to the X-ray emission of quasars is still an open issue. This work can constrain a future theoretical model for MAD in RLAGNs.","Citation Text":["Gupta et al.","2020"],"Functions Text":["Second,","discovered that the distributions of X-ray photon spectral indices between RLAGNs and their radio-quiet counterpart are very similar"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1305,1317],[1325,1329]],"Functions Start End":[[1297,1304],[1331,1463]]}
{"Identifier":"2016MNRAS.462.3945A__Ebbets_&_Savage_1982_Instance_1","Paragraph":"A comparison of the He and O abundances in RWT 152 with evolutionary models of stellar yields by Marigo (2001) suggests a progenitor star with an initial mass of \u223c1.3 M\u2299 and very low metallicity Z = 0.004. The stellar mass is compatible with that expected for sdOs (Heber 2009) and the low metallicity indicates that the progenitor was formed in a poor-metal environment. In contrast, similar models of initial mass \u223c1.25 M\u2299 and Z = 0.004 by Karakas (2010) predict abundances substantially different from those found in RWT 152. Even for the lowest metallicity in the Karakas models (Z = 0.0001), we do not recover the chemical abundances of RWT 152. Therefore, it is clear that drawing conclusions about the progenitor star of RWT 152 from the chemical abundances of the nebula may be a bit misleading. For this reason, we have obtained information about the progenitor from the current status of its CS. Fig. 6 shows the position of RWT 152 in the Hertzsprung-Russell diagram log\u2009g\u2013Teff with the post-AGB tracks by Bloecker (1995) and Schoenberner (1983). The location of the star (Teff \u2243 45 000 K, log\u2009g \u2243 4.5; Ebbets & Savage 1982) is consistent with a current mass of M\u223c0.55 M\u2299 which implies an initial mass in the main sequence of \u223c1 M\u2299. For such a low-mass star, the ejected mass during the AGB evolution is expected to be small. In fact, with the electron density derived above and the size of the nebula, we obtain values of 1.3 \u00d7 10\u22122\u20131.6 \u00d7 10\u22121 M\u2299 for the ionized nebular mass (assuming 2.4 and 6.5 kpc, respectively, and a filling factor of 0.6). These values are much smaller than ionized masses usually obtained for PNe (see e.g. Hua & Kwok 1999), further supporting a low-mass progenitor for RWT 152. Moreover, taking into account that 0.1\u20130.3 M\u2299 are lost in the red giant branch phase of a low-mass star (Dorman, Rood & O'Connell 1993), the current mass of the CS and the obtained ionized mass, we recover a progenitor star with a \u223c0.8\u20131.0 M\u2299, in agreement with the value obtained from the position of the CS in the log\u2009g\u2013Teff diagram. Such a low ionized mass, combined with a relatively high kinematical age, may explain the low surface brightness of RWT 152. If a low-mass progenitor is involved in the evolution of other PNe+sdO systems, it is not surprising that these PNe are very faint and, in some cases, may have faded beyond detection.","Citation Text":["Ebbets & Savage 1982"],"Functions Text":["The location of the star (Teff \u2243 45 000 K, log\u2009g \u2243 4.5;","is consistent with a current mass of M\u223c0.55 M\u2299 which implies an initial mass in the main sequence of \u223c1 M\u2299."],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[1114,1134]],"Functions Start End":[[1058,1113],[1136,1243]]}
{"Identifier":"2021AandA...653A.111R__Jones_et_al._(2021)_Instance_1","Paragraph":"As done by Le F\u00e8vre et al. (2020), we visually inspect the ancillary data, the intensity maps, the velocity and velocity dispersion fields presented in Sect. 3.1 to search for the presence of multiple components or disturbed morphology near the position of the targets. The channel maps, the spectra and the PVDs are checked together searching for consistent emission features. By taking into account the results of the initial qualitative classification by Le F\u00e8vre et al. (2020) and of the more recent quantitative analysis of a subsample of the ALPINE targets by Jones et al. (2021), we proceed with a more in-depth characterization of the [CII]-detected galaxies aimed at obtaining a robust merger fraction at z\u2004\u223c\u20045. Adopting the same criteria described in Sect. 2 to differentiate the targets and considering the S\/N of the minor merger component as described in Sect. 3.1, we find a slightly lower fraction (\u223c31%, 23 out of 75 [CII]-detected sources) of mergers7 if compared to the 40% found by Le F\u00e8vre et al. (2020), with 12, 20 and 7% of the sample made by rotating, extended and compact dispersion dominated sources, respectively. To be more conservative in the classification of the galaxies (especially for obtaining a more robust merger statistics), we define the remaining 30% of the sample as \u2018uncertain\u2019, a new category that includes the weak galaxies (as described in Le F\u00e8vre et al. 2020) and also objects that, by visual inspection, present features that are intermediate to those of various classes. This category is similar to the \u2018uncertain\u2019 (UNC) class introduced in Jones et al. (2021) that, based on the results of the 3DBarolo fits, contains sources they are unable to classify because of the low S\/N and\/or spectral resolution, or contrasting evidence in their classification criteria. Although the uncertain category is populated by a significantly larger fraction of sources with respect to the weak class (\u223c16%) in Le F\u00e8vre et al. (2020), we recover the same qualitative morpho-kinematic distribution of the previous analysis, confirming the high fraction of rotators and mergers at these early epochs.","Citation Text":["Jones et al. (2021)"],"Functions Text":["By taking into account the results of the initial qualitative classification by Le F\u00e8vre et al. (2020) and of the more recent quantitative analysis of a subsample of the ALPINE targets by","we proceed with a more in-depth characterization of the [CII]-detected galaxies aimed at obtaining a robust merger fraction at z\u2004\u223c\u20045."],"Functions Label":["Extends","Extends"],"Citation Start End":[[566,585]],"Functions Start End":[[378,565],[587,720]]}
{"Identifier":"2017AandA...605A..88L__Cordiner_et_al._2015_Instance_1","Paragraph":"Altogether, the approximately thirty molecules recently detected have confirmed the chemical complexity in the nebula, and generated our interest for the present study. Of these species, we will focus our attention on the seventeen species listed by molecular families in Table 1. As can be seen in this table, the WHISPER survey allowed the detection of some organic molecules in the Horsehead nebula, such as formaldehyde (H2CO) and methanol (CH3OH), which constitute key species in the likely synthesis of more complex organic molecules such as some prebiotic molecules (Bernstein et al. 2002; Mu\u00f1oz Caro et al. 2002; Garrod et al. 2008). Because they are detected in a wide variety of interstellar sources \u2013 in hot cores (Sutton et al. 1995; Ceccarelli et al. 2000), dark clouds (Bergman et al. 2011), shocked regions (e.g. Sakai et al. 2012; Codella et al. 2012; Tafalla et al. 2010) and even in comets (Mumma & Charnley 2011; Cordiner et al. 2015) \u2013 it is of prime importance to understand well how these precursor molecules form. H2CO is commonly thought to form both in the gas-phase and on grain surfaces, while CH3OH is believed to be only formed on grain surfaces (Garrod et al. 2006; Geppert et al. 2006). Guzman et al. (2013) reported the observations of these two molecules toward the Horsehead nebula in both the PDR and Core positions. Unable to reproduce the observed abundances of either H2CO or CH3OH at the PDR position with only pure gas-phase models, they concluded that, for this region, both species are formed on grain surfaces and then photodesorbed into the gas phase. On the other hand, at the Core position, a pure gas-phase model can reproduce the observed H2CO abundance, while photodesorption of ices is still needed to explain the observed abundance of CH3OH. Other organic molecules were reported in the Horsehead nebula as first detections in a PDR environment, including HCOOH (formic acid), CH2CO (ketene), CH3CHO (acetaldehyde), and CH3CCH (propyne) (Guzman et al. 2014). Their abundances were found to be higher at the PDR position than at the Core, revealing that complex organic chemistry is also occurring in UV-illuminated neutral gas (Guzman et al. 2014). Of these molecules, some \u2013 HCOOH, CH2CO, and CH3CHO \u2013 have now also been detected in the Orion bar PDR (Cuadrado et al. 2016, 2017). ","Citation Text":["Cordiner et al. 2015"],"Functions Text":["Because they are detected in a wide variety of interstellar sources","and even in comets","it is of prime importance to understand well how these precursor molecules form."],"Functions Label":["Background","Motivation","Motivation"],"Citation Start End":[[932,952]],"Functions Start End":[[642,709],[889,907],[956,1036]]}
{"Identifier":"2018MNRAS.473.1512A__Granot_&_Sari_2002_Instance_1","Paragraph":"The duration of the manually scheduled follow-up observations were also increased to improve the likelihood of a radio detection. The recent investigation of the entire sample of radio-detected GRBs before 2011 April by Chandra & Frail (2012) demonstrated that the majority of GRBs detected in the radio band at 8.5\u2009GHz had a peak flux between 0.1 and 0.2 mJy\u2009beam\u22121 at 5\u201310 d post-burst (see fig. 4 of Chandra & Frail 2012). A 4 h AMI observation is therefore required to reach an rms noise of \u223c0.03\u20130.04 mJy\u2009beam\u22121 that will allow the reliable detection of >0.1\u20130.2 mJy\u2009beam\u22121 sources. However, it is worth noting that since GRB relativistic blast waves generate synchrotron radiation as they expand into the circumstellar (wind generated) medium (Granot & Sari 2002), we expect the forward-shock of the afterglow to peak more brightly at 15.7\u2009GHz and at earlier times than the peaks recorded by Chandra & Frail (2012). We therefore require a higher monitoring cadence at early times (within 5\u2009d post-burst) to detect similar radio peaks. As the range of radio peaks observed by Chandra & Frail (2012) will be brighter at 15.7\u2009GHz, the rms achieved by a 4\u2009h AMI observation will be sufficient for detecting events similar to those seen in their sample. The follow-up observations are manually scheduled to occur near transit approximately 24 h, 3, 7, and 10\u2009d post-burst, with this temporal spacing designed to catch the peak of the forward- or reverse-shock at 15.7 GHz at a range of redshifts (z \u2272 5; e.g. see figs 22 and 23 of Chandra & Frail 2012). In the event that a GRB radio counterpart was detected, the AMI observing cadence was increased to a 4 h observation every 1 or 2 d. As part of the AMI GRB observing programme, we also obtained manually scheduled observations of GRBs that were detected with the Fermi Large Area Telescope (LAT; Atwood et al. 2009), the Fermi Gamma-ray Burst Monitor (GBM; Meegan et al. 2009) and the International Gamma-Ray Astrophysics Laboratory (INTEGRAL; Winkler et al. 2003), whose positions had been more precisely localized through the identification of X-ray and\/or optical counterparts, usually by the Swift-XRT, Swift-UVOT, or one of the ground-based GRB follow-up programmes.","Citation Text":["Granot & Sari 2002"],"Functions Text":["However, it is worth noting that since GRB relativistic blast waves generate synchrotron radiation as they expand into the circumstellar (wind generated) medium","we expect the forward-shock of the afterglow to peak more brightly at 15.7\u2009GHz and at earlier times than the peaks recorded by Chandra & Frail (2012).","We therefore require a higher monitoring cadence at early times (within 5\u2009d post-burst) to detect similar radio peaks."],"Functions Label":["Uses","Compare\/Contrast","Uses"],"Citation Start End":[[750,768]],"Functions Start End":[[588,748],[771,921],[922,1040]]}
{"Identifier":"2018ApJ...863..162M__Liu_et_al._2013_Instance_3","Paragraph":"NLFFF extrapolation provides the reconstructed coronal magnetic field for AR 11158 from 2011 February 13 \u2212 2011 February 15 (Figures 1(d)\u2013(f)). The field lines (yellow lines) within the core of the AR have arcade-like structure with a relatively strong twist mainly near the PIL. These figures show that the magnetic field evolved during this period. Although we did not quantitatively compare the field lines with the observation, in general, the reconstructed coronal field morphologies match with the observations in Figures 1(a)\u2013(c). The general morphologies and the locations of the high-twist fields are also in agreement with many previous studies (Jing et al. 2012; Sun et al. 2012; Dalmasse et al. 2013; Inoue et al. 2013, 2014a; Liu et al. 2013; Wang et al. 2013; Aschwanden et al. 2014; Malanushenko et al. 2014; Zhao et al. 2014). Unlike Zhao et al. (2014) who could identify the twisted flux rope from the topology of the reconstructed coronal field, we could not find an obvious topological signature of a flux rope existing in our NLFFF during our analysis time window. It might be due to the fact that there was little magnetic flux with twist higher than one turn in our NLFFF and it is difficult to topologically define it as a flux rope. However, our result is consistent with other NLFFF results (Jing et al. 2012; Sun et al. 2012; Liu et al. 2013; Wang et al. 2013; Inoue et al. 2014a; Malanushenko et al. 2014). The high-twist region in our result is also in agreement with the region with high helicity flux (Dalmasse et al. 2013) and the location of the flare ribbons (Bamba et al. 2013; Liu et al. 2013), as well as the high current density region (Janvier et al. 2014). Figures 1(g)\u2013(i) show the evolution of the twist distribution map, with the magnetic twist of the field lines plotted at the footpoints of field lines according to a color scale. This shows that the high-twist (strongly right-handed twist corresponding to Tw > 0.5) areas are concentrated in only a limited part of the AR. The high-twist area grew and became even more twisted just before the X2.2 flare (Figure 1(i)). Most parts of the AR have twist values less than 0.25, but near the PIL the twist can reach more than 0.5, even up to about a full turn. This is consistent with the results of Sun et al. (2012) and Inoue et al. (2014a). A high-twist (strong negative\/left-handed twist) area also developed in the eastern part of the AR, which did not exist initially on February 13. Both of these high-twist areas produced several flares. However, here we focus on the flares that resulted from the high-twist core region near the center of the AR, where the M6.6 and X2.2 flares occurred.","Citation Text":["Liu et al. 2013"],"Functions Text":["The high-twist region in our result is also in agreement with","the location of the flare ribbons"],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[1612,1627]],"Functions Start End":[[1434,1495],[1558,1591]]}
{"Identifier":"2021AandA...645A..96P__Marcantonio_et_al._(2018)_Instance_2","Paragraph":"The ESPRESSO DFS concept (Di Marcantonio et al. 2018) was conceived during its preliminary design phases with the goal of maximizing operational efficiency, flexibility, and scientific output while complying with the standard Paranal Observatory operational scheme. The main challenge derives from the requirement to operate ESPRESSO in a seamless way with any of the UT\u2019s or with all four UT\u2019s simultaneously. This must be possible not only with a predetermined schedule, but also \u201con the fly\u201d. The flexibility in ESPRESSO\u2019s operations has been tackled by adopting a new DFS deployment plan described in Di Marcantonio et al. (2018) that is exceptional under various aspects because it has to cope with various telescope and instrument configurations while remaining operationally simple. Figure 8 shows the main ESPRESSO DFS elements and their final deployment. Besides the software packages already described, part of the software for the control of the CT devices has been incorporated into the VLT telescope CS to allow CT operations even when ESPRESSO is offline (thus avoiding conflicts, e.g., with instruments of the VLT Interferometer operations). In addition to the standard DFS software packages, ESPRESSO is the first instrument to also provide a data analysis package that is able to extract relevant astronomical observables from the reduced data. The following DFS subsystems are specific to ESPRESSO: (1) the ETC hosted on the ESO web page9; (2) the CS with the full suite of acquisition, observation, and calibration templates that are able to control all vital parts of the instrument and the CT (Calderone et al. 2018); (3) the data reduction software (DRS) package (or \u201cpipeline\u201d) capable of providing \u201cscience-ready\u201d reduced data only minutes after the end of the individual observation; (4) the data analysis software (DAS) package that produces higher-level astronomical observables with no or limited supervision; (5) the DRS and DAS are distributed to the community10.","Citation Text":["Di Marcantonio et al. (2018)"],"Functions Text":["The flexibility in ESPRESSO\u2019s operations has been tackled by adopting a new DFS deployment plan described in","that is exceptional under various aspects because it has to cope with various telescope and instrument configurations while remaining operationally simple."],"Functions Label":["Background","Background"],"Citation Start End":[[605,633]],"Functions Start End":[[496,604],[634,789]]}
{"Identifier":"2022MNRAS.513.2349C__Brada\u010d_et_al._2002_Instance_1","Paragraph":"Eigenbrod et al. (2006) first modelled this system with HST imaging and suspected that the second set of bluer arcs in F814W band (see Fig. 1) inside and outside the area delimited by the red arcs in F160W band could be either a second source at a different redshift or a star-forming region in the source galaxy. We examine the possibility of a second source plane existing at a lower redshift than the source (z = 1.52) due to the bluer colour of the arc and find that the scenario is very unlikely, as the macro model determined by the red arc cannot reproduce a reasonable source for the blue arcs given a possible range of the source redshift from z = 0.5 to z  1.52. In contrast, we do find that a star-forming region can be reconstructed at the same source redshift. Faure et al. (2011) modelled the lens with high-resolution H and Ks imaging obtained using the European Southern Obseratory (ESO) Very Large Telescope (VLT) with AO and the laser guide star system. They identified a luminous object, located \u223c0.3 arcsec to the north of the lens galaxy, but showed that it cannot be responsible for the anomalous flux ratios. Many studies (e.g. Metcalf & Madau 2001; Brada\u010d et al. 2002; Dalal & Kochanek 2002; Pooley et al. 2012; Schechter et al. 2014; Glikman et al. 2018; Badole et al. 2020) have shown that the macro model cannot explain the flux ratios in this lens, which suggested the presence of microlensing or dark matter substructures. Thus, to avoid possible biases caused by flux ratios, we only use the lensed quasar positions and the extended arc to constrain the mass model, which is also the standard procedure for H0 measurements in TDCOSMO Collaboration. Ignoring the fluxes of the lensed quasar images will not affect the constraining power of the imaging data, since the lensed arc emission is much more constraining than the lensed quasar fluxes. In addition, For the error budget of H0 contributed from the uncertainties of the lensed quasar positions, Chen et al. (2021a) have showed that with high-resolution AO imaging it is a subdominate term given configuration of the J0924+0219 lens system. Gilman et al. (2020) also show that the presence of substructures do not bias H0 above the per cent level. We use glee, a strong lens modelling code to model (Suyu & Halkola 2010; Suyu et al. 2012) the three HST bands and one Keck AO band simultaneously. We describe the models in the following for fitting the high-resolution imaging data. We show the imaging, models, normalized residuals, and reconstructed sources in Fig. 4. Note that since the source in F555W band has more clumpy star-forming regions, the reconstructed source is less regular, with small-scale structures and more noise.5 In addition, the noise-overfitting problem is due to the fact that the outer region of the source plane is under-regularized, but this effect will not affect the uncertainty because the uncertainty will be dominated by the time delay and velocity dispersion measurements. Besides, we model the imaging with different source resolutions and marginalize over them to control the systematics.","Citation Text":["Brada\u010d et al. 2002"],"Functions Text":["Many studies (e.g.","have shown that the macro model cannot explain the flux ratios in this lens, which suggested the presence of microlensing or dark matter substructures. Thus, to avoid possible biases caused by flux ratios, we only use the lensed quasar positions and the extended arc to constrain the mass model, which is also the standard procedure for H0 measurements in TDCOSMO Collaboration. Ignoring the fluxes of the lensed quasar images will not affect the constraining power of the imaging data, since the lensed arc emission is much more constraining than the lensed quasar fluxes."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1173,1191]],"Functions Start End":[[1132,1150],[1300,1873]]}
{"Identifier":"2021ApJ...913..115A__Linden_et_al._2012_Instance_1","Paragraph":"The origin of the GC VHE emission remains undetermined, due in part to source confusion and the limitations of current instruments. The source of VER J1745\u2013290 may be Sgr A* (Atoyan & Dermer 2004; Aharonian & Neronov 2005; Ballantyne et al. 2011; Chernyakova et al. 2011; Fatuzzo & Melia 2012; Kusunose & Takahara 2012; Fujita et al. 2017; Rodr\u00edguez-Ram\u00edrez et al. 2019) or PWN G359.95-0.04 (Wang et al. 2006; Hinton & Aharonian 2007), with which it is spatially coincident (Acero et al. 2010). Other possible origins include the annihilation of dark matter particles (Bergstr\u00f6m et al. 2005a, 2005b; Horns 2005; Profumo 2005; Aharonian et al. 2006c; Belikov et al. 2012; Cembranos et al. 2012, 2013; Gammaldi et al. 2016) or a population of millisecond pulsars (Bednarek & Sobczak 2013; Bartels et al. 2016; Gu\u00e9pin et al. 2018). The mechanism of gamma-ray emission may be predominantly due to hadronic processes, where relativistic protons interact with gas and subsequently produce gamma rays through neutral pion decay (Aharonian & Neronov 2005; Ballantyne et al. 2011; Chernyakova et al. 2011; Fatuzzo & Melia 2012; Linden et al. 2012; Gu\u00e9pin et al. 2018), leptonic processes where gamma rays are produced when electrons and positrons undergo inverse Compton scattering off a radiation field (Atoyan & Dermer 2004; Hinton & Aharonian 2007; Kusunose & Takahara 2012; Lacroix et al. 2016), or a combination of processes (hybrid scenario), where leptons produce high-energy, but not VHE, gamma rays (Guo et al. 2013). Both the correlation of VHE emission with the CMZ and the lack of a cutoff in the diffuse spectrum support a hadronic scenario, capable of explaining both VER J1745\u2013290 and the diffuse emission (Aharonian et al. 2006b; Linden et al. 2012; Abramowski et al. 2016). Measurement of the diffuse spectrum by H.E.S.S. up to energies of tens of TeV with no evidence of a cutoff has also been interpreted as evidence for the existence of PeV protons within the central 10 parsecs of the GC, accelerated by Sgr A* (Abramowski et al. 2016). While cosmic rays are known to extend up to PeV energies (e.g., H\u00f6randel 2003), few, if any, accelerators of PeV cosmic rays, or \u201cPeVatrons,\u201d have been clearly established (e.g., Abramowski et al. 2016; Abeysekara et al. 2020). Discovering the nature of PeVatrons in our Galaxy is thus a particularly important step in understanding the origins of cosmic rays.","Citation Text":["Linden et al. 2012"],"Functions Text":["The mechanism of gamma-ray emission may be predominantly due to hadronic processes, where relativistic protons interact with gas and subsequently produce gamma rays through neutral pion decay"],"Functions Label":["Background"],"Citation Start End":[[1119,1137]],"Functions Start End":[[829,1020]]}
{"Identifier":"2015MNRAS.451.2174T__Yoshida_et_al._2006_Instance_1","Paragraph":"In the last two decades, it has been demonstrated that massive galaxies harbour supermassive black holes (SMBHs) in the centres of their bulge components (Kormendy & Ho 2013, and references therein). Also, at redshifts higher than 6, quasars are found that possess SMBHs with the mass higher than 109\u2009M\u2299 (Fan et al. 2001; Kurk et al. 2007). The formation history of these SMBHs is among the most significant unsolved issues in astrophysics (Volonteri & Bellovary 2012; Haiman 2013). Two recently discovered high-redshift quasars, ULAS J112010+641 with the mass of MBH = 2 \u00d7 109\u2009M\u2299 at redshift z = 7.085 (Mortlock et al. 2011) and SDSS J01001+2802 with MBH = 1.2 \u00d7 1010\u2009M\u2299 at z = 6.30 (Wu et al. 2015) have raised a serious problem for the formation of SMBHs. Possible building blocks for such high-redshift SMBHs are the remnants of first stars. The initial mass function of first stars is thought to be more or less top-heavy (Abel, Bryan & Norman 2000; Nakamura & Umemura 2001; Bromm, Coppi & Larson 2002; Yoshida et al. 2006; Greif et al. 2011; Hirano et al. 2014; Susa, Hasegawa & Tominaga 2014). First stars of several tens M\u2299 undergo supernovae, leaving black holes (BHs) of few tens M\u2299 (Heger & Woosley 2002). For SMBHs to grow from such first star remnants through mass accretion at z \u2273 6, a super-Eddington accretion rate is requisite. If SMBHs grow continuously by mass accretion from BH remnants of \u223c20\u2009M\u2299, the Eddington ratio (\u03bb) is required to be \u03bb = 1.4 for ULAS J112010+641, or \u03bb = 1.3 for SDSS J01001+2802. However, the continuous accretion is unlikely to be sustained due to feedbacks, and thus the average mass accretion rates should be lower than the Eddington rate (Alvarez, Wise & Abel 2009; Milosavljevic, Couch & Bromm 2009). On the other hand, seed BHs may stem from supermassive stars of 104\u20136\u2009M\u2299 as a result of the direct collapse of primordial density fluctuations (Umemura, Loeb & Turner 1993; Bromm & Loeb 2003; Inayoshi & Omukai 2012). These BHs are thought to be incorporated into a primordial galaxy of \u223c108\u2013109\u2009M\u2299 (Greene 2012). If an SMBH grows via gas accretion from such a massive BH (MBH), the constraint on the accretion rate can be alleviated.","Citation Text":["Yoshida et al. 2006"],"Functions Text":["The initial mass function of first stars is thought to be more or less top-heavy"],"Functions Label":["Background"],"Citation Start End":[[1008,1027]],"Functions Start End":[[846,926]]}
{"Identifier":"2018ApJ...864..158L__Zank_&_Matthaeus_1992_Instance_1","Paragraph":"So far, the discussion has emphasized particle energization by local plasma regions of contracting and fast reconnecting (merging) small-scale flux ropes generated in the vicinity of large-scale primary current sheets through turbulent current sheet reconnection. However, one can also approach this topic from the perspective of MHD turbulence theory, simulations, and related solar wind observations. Theoretical considerations and simulations of MHD turbulence in the presence of a significant background\/guide magnetic field suggest that solar wind turbulence can to lowest order be modeled in terms of a combination of a dominant quasi-2D turbulence component of coherent structures (small-scale magnetic islands) perpendicular to the background\/guide field and a minor parallel-propagating Alfv\u00e9n wave turbulence component (Shebalin et al. 1983; Zank & Matthaeus 1992, 1993; Dmitruk et al. 2004; Zank et al. 2017), a view that is consistent with analysis of solar wind observations (Matthaeus et al. 1990; Bieber et al. 1996) and with the finding that quasi-2D turbulence alone is not sufficient to explain observed solar wind turbulence (Turner et al. 2012). A recent analysis of Wind data to identify inertial-scale flux ropes indicates that these structures are much more commonplace in the solar wind near 1 au than previously thought. Zheng (2017), Zheng et al. (2017), and Zheng & Hu (2018) identified an unprecedented number of small-scale flux ropes at 1 au with scales in the inertial range using the Grad\u2013Shafranov reconstruction approach (\u223c3500 per year on average) with a clear solar cycle dependence, a number that is expected to grow when the data analysis shifts to shorter timescales. Furthermore, an axial (out-of-plane) current density distribution constructed from the Grad\u2013Shafranov-based data analysis yielded a non-Gaussian probability density function (pdf) entirely consistent with the out-of-plane current density pdf produced from compressible 2D MHD turbulence simulations with a strong out-of-plane guide field, in which merging magnetic island structures are a common occurrence (Greco et al. 2009). This result, combined with the sheer number of small-scale flux ropes being identified, suggests that the common occurrence of small-scale flux ropes in the low-latitude solar wind near 1 au is a natural consequence of local MHD turbulence evolution in a highly conductive plasma with a strong guide field, independent of additional flux-rope production at primary current sheets. Furthermore, observational evidence of merging (magnetic reconnection) of neighboring small-scale flux ropes at Earth (Khabarova et al. 2015, 2016), including evidence on the basis of Grad\u2013Shafranov reconstruction of small-scale flux ropes (Zheng & Hu 2016; Zheng 2017; Zheng et al. 2017), is consistent with the concept of quasi-2D turbulence theory of an inverse cascade of magnetic island energy to smaller wavenumbers.","Citation Text":["Zank & Matthaeus 1992"],"Functions Text":["Theoretical considerations and simulations of MHD turbulence in the presence of a significant background\/guide magnetic field suggest that solar wind turbulence can to lowest order be modeled in terms of a combination of a dominant quasi-2D turbulence component of coherent structures (small-scale magnetic islands) perpendicular to the background\/guide field and a minor parallel-propagating Alfv\u00e9n wave turbulence component","a view that is consistent with analysis of solar wind observations","and with the finding that quasi-2D turbulence alone is not sufficient to explain observed solar wind turbulence"],"Functions Label":["Background","Similarities","Similarities"],"Citation Start End":[[852,873]],"Functions Start End":[[403,828],[921,987],[1032,1143]]}
{"Identifier":"2019ApJ...874..166C__Dyks_et_al._2004_Instance_1","Paragraph":"The realistic structures of the pulsar magnetosphere still remain uncertain. Knowledge about the pulsar magnetosphere structures can be used to identify the potential sites of particle acceleration and gamma-ray emission. A vacuum dipole field is generally adopted in the early study of pulsar emission, because it has an exact analytical solution given by Deutsch (1955). Based on this field structure, different theoretical models have been developed to explain the observed pulsar emission. In these models, it is widely believed that particles are accelerated in the gap region where an accelerating electric field is created because of the deficit of charges. Gamma-ray emission is produced by the curvature or inverse-Compton radiation from high-energy particles accelerated in these gaps. Due to different emission zone locations, standard pulsar radiation models include the polar cap (PC; e.g., Ruderman & Sutherland 1975; Daugherty & Harding 1982), the slot-gap (SG) (e.g., Dyks & Rudak 2003; Dyks et al. 2004; Muslimov & Harding 2004), and the outer gap (OG; e.g., Cheng et al. 1986, 2000; Zhang & Cheng 1997, 2001; Zhang et al. 2004) models. These gap models have achieved great successes in explaining pulsar high-energy emissions and light curves (e.g., Watters et al. 2009; Romani & Watters 2010). However, the vacuum solution has no plasma; it is not able to reproduce any pulsar phenomena. It is well known that the pulsar magnetosphere should be filled with plasma (Goldreich & Julian 1969). In the presence of abundant plasma, all accelerating electric fields can be efficiently screened to form a force-free (FF) magnetosphere. The FF solution for an aligned rotator was first obtained by Contopoulos et al. (1999). The CKF solution consists of a closed field line region extending to the light cylinder (LC), an open field line region, and an equatorial current sheet beyond the LC. Moreover, the time-dependent simulations for the FF axisymmetric rotator also confirmed the closed\u2013open CKF solution (e.g., Komissarov 2006; McKinney 2006; Timokhin 2006; Yu 2011; Parfrey et al. 2012; Cao et al. 2016a; Etienne et al. 2017).","Citation Text":["Dyks et al. 2004"],"Functions Text":["Due to different emission zone locations, standard pulsar radiation models include the","the slot-gap (SG) (e.g.,","models."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[1003,1019]],"Functions Start End":[[796,882],[959,983],[1146,1153]]}
{"Identifier":"2018AandA...620A.122S__\u0160vanda_et_al._(2014)_Instance_1","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)"],"Functions Text":["By averaging over 222 976 supergranular cells and applying time-distant helioseismic inversions,","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."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1107,1127]],"Functions Start End":[[1010,1106],[1128,1492]]}
{"Identifier":"2020AandA...639A..46B__\u0160tver\u00e1k_et_al._(2009)_Instance_4","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)"],"Functions Text":["Scudder & Olbert (1979) also predict that the halo Ebp\/kBTc ratio remains constant with heliocentric distance, whereas","find that the halo Ebp\/kBTc ratio decreases with heliocentric distance."],"Functions Label":["Differences","Differences"],"Citation Start End":[[1376,1397]],"Functions Start End":[[1257,1375],[1398,1469]]}
{"Identifier":"2021ApJ...916...64T__Lundstedt_et_al._2002_Instance_1","Paragraph":"The above results, when considered as relating to different aspects of the same physical process, suggest the probable existence of correlations between periods of significant Alfv\u00e9nic turbulence, higher time-integrated AE and HILDCAA periods, and longer storm recovery phases. However, a statistical demonstration of the coupling between the solar wind Alfv\u00e9nicity and the ring current dynamics, aimed at providing evidence in support of this possible scenario and distinguishing between a temporal coincidence and a causal linkage, is still missing in the literature and motivates the present paper, which deals with the correlation between the presence of Alfv\u00e9nic streams and long-living periods of geomagnetic activity at low latitudes, after the main phases of storms induced by recurrent or non-recurrent solar transients. Specifically, by exploiting a 16 yr survey of interplanetary and geomagnetic data, an attempt is here sought to provide an empirical law that quantitatively links the duration of the recovery phase (as inferred from the SYM-H index-related magnetospheric output at low latitudes) with the Alfv\u00e9nic content of the solar wind fluctuations. The goal is twofold. First, this would allow a robust establishment of the extent to which Alfv\u00e9nic intervals contribute to slow recovery phases. Second, interesting predictions of the duration of the total geomagnetic storm (main plus recovery phase) would be available once the empirical law between the two periods was established. This is crucial since, although space weather forecasting often focuses on estimating the onset and intensity of geomagnetic storms (Joselyn 1995; Lundstedt et al. 2002; Abunina et al. 2013), the impacts on ground infrastructure are also related to the whole storm duration (closely linked to the storm intensity; e.g., Haines et al. 2019) through time-integrated effects (Balan et al. 2016; Lockwood et al. 2016; Mourenas et al. 2018). Indeed, the investigation of the different storm indices for low, intermediate, and high latitudes has shown that infrastructures on Earth, namely, the equipment of electric power transmission networks, are affected by the geomagnetically induced currents. In particular, transformers and some electrical substations have been observed to be very sensitive to prolonged periods of substorms, with delayed increases in anomalies (\u0160vanda et al. 2020). Moreover, highly energetic charged particles in near-Earth space represent a risk for satellite operations. A strong relationship exists between relativistic electrons and HILDCAA events: the longer the HILDCAAs last, the higher the energy reached by relativistic electrons (Thorne et al. 2013). Hajra et al. (2015), taking into account solar cycle 23, found a flux enhancement of relativistic electrons in the outer radiation belt during HILDCAAs, corroborating the results of Meredith et al. (2002, 2003). Furthermore, there, large events are always followed by high flux peaks of 2 MeV electrons (Mourenas et al. 2019). It is worth pointing out that the investigation of the particle flux for electron energy larger than 2 MeV is extremely important for the space weather forecasting as a key indicator of serious hazard for damage of spacecraft in low, medium, and geosynchronous Earth orbits (Forsyth et al. 2020). Therefore, prolonged periods of strong Alfv\u00e9nic turbulence in the solar wind could be precursors of many adverse effects in both ground installations and space satellites.","Citation Text":["Lundstedt et al. 2002"],"Functions Text":["This is crucial since, although space weather forecasting often focuses on estimating the onset and intensity of geomagnetic storms"],"Functions Label":["Background"],"Citation Start End":[[1650,1671]],"Functions Start End":[[1503,1634]]}
{"Identifier":"2021AandA...655A..99D__Bensby_&_Feltzing_2006_Instance_1","Paragraph":"In Fig. 6, we plot the [C\/O] ratios dependence on [O\/H] for the different stellar populations. This figure serves to evaluate the balance between the two different elements directly with the evolution of one of them, which in this case has a well-known, single production site. Here, we chose to only show the ratios with the oxygen abundances from the 6158 \u212b indicator as it was shown to present less dispersion and be moretrustworthy (Bertran de Lis et al. 2015). Nevertheless, for completeness, we present the figure with the forbidden oxygen line in the appendix (Fig. A.1). Previous works in the literature have shown an increase of [C\/O] ratios as [O\/H] increases (e.g. Bensby & Feltzing 2006; Nissen et al. 2014; Amarsi et al. 2019), but our ratios present a quite large dispersion that together with the shorter range in [O\/H] prevents us from seeing a clear behaviour. If we focus on the results with the 6158 \u212b line, the [C\/O] ratios present a general flat trend (see the upper panel of Fig. 6 with the running mean for each population), as [O\/H] increases to then clearly decrease at [O\/H] \u2273 0 dex, both for thin-disk and h\u03b1mr stars. To better evaluate the significance of this apparent trend, we applied a weighted least squares fit to the [C\/O] values of thin disk and h\u03b1mr stars at [O\/H] \u2273 0 dex. We then used the values of the slopes and the associated uncertainties to assess the significance of the fits. The p-values come from the F-statistics that tests the null hypothesis that the data can be modelled accurately by setting the regression coefficients to zero. The resulting p-values are 3.3 e\u22122 and 1.6 e\u22122 for the thin disk and h\u03b1mr stars, meaning that the correlations are significant. This trend is in agreement with the previously mentioned turning point of [C\/O] at [O\/H] ~ 0.0 dex (Carigi et al. 2005). However, this is in contrast with the flattening at super-solar [O\/H] presented by Nissen et al. (2014) or the increasing trend found by Franchini et al. (2021) which might be caused by the use of different oxygen indicators. On the other hand, the [C\/O] ratios for thick-disk stars present no trend but show a clear offset in its running mean with respect to the thin disk as also reported by Amarsi et al. (2019). This separation between the thick and thin disk, also observed for other elements, supports the different formation episodes of both populations (Chiappini et al. 2001; Amarsi et al. 2019).","Citation Text":["Bensby & Feltzing 2006"],"Functions Text":["Previous works in the literature have shown an increase of [C\/O] ratios as [O\/H] increases (e.g.","but our ratios present a quite large dispersion that together with the shorter range in [O\/H] prevents us from seeing a clear behaviour."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[676,698]],"Functions Start End":[[579,675],[741,877]]}
{"Identifier":"2016MNRAS.461.1719C__Fu_et_al._2012_Instance_2","Paragraph":"HATLAS12-00 had already been identified as a candidate gravitationally lensed galaxy as a result of its high submm flux (i.e. F500 > 100 mJy), red Herschel colours and the lack of a bright optical or radio counterpart (see e.g. Negrello et al. 2010 for a discussion of the selection of lens candidates in H-ATLAS and other Herschel surveys). This source was therefore observed spectroscopically in the submm. A CO spectroscopic redshift of 3.26 was first suggested by Z-spec (Bradford et al. 2004) observations, then subsequently confirmed by observations by the CARMA interferometer (Leeuw et al., in preparation) and the Zpectrometer instrument (Harris et al. 2007) on the Greenbank Telescope (Harris et al. 2012; see also Fu et al. 2012). Additional followup observations in the optical, near-IR, submm and other wavelengths were targeted at the lensed z = 3.26 source and the foreground objects responsible for the lensing, resulting in detailed analyses of this lensing system by Fu et al. (2012) and Bussmann et al. (2013). Their conclusions are that the z = 3.26 source HATLAS12-00 is subject to gravitational lensing, with a magnification of 9.6 \u00b1 0.5 in both the submm continuum and CO, and 16.7 \u00b1 0.8 in the K\u2032 band, by two foreground galaxies, one at a spectroscopically determined redshift of 1.22, and another with photometry suggesting that it lies at a similar redshift. The submm photometry of HATLAS12-00 at 890\u2009\u03bcm acquired with the Submillimeter Array (SMA) as part of this programme (Fu et al. 2012; Bussmann et al. 2013) is fully consistent with the 870\u2009\u03bcm and 850\u2009\u03bcm fluxes derived for this source from the LABOCA and SCUBA2 data to be presented here. The spectral energy distribution (SED) of the lensed source, after correcting for the lensing amplification, is well matched by the optically thick SED model for Arp220 from Rangwala et al. (2011), with a lensing-corrected far-IR luminosity of 1.2 \u00b1 0.2 \u00d7 1013\u2009L\u2299, and an implied star formation rate of 1400 \u00b1 300 \u2009M\u2299 yr\u22121. In many ways the unlensed properties of this object match those of the broader population of bright submm selected galaxies first discovered by the SCUBA submm imager (see e.g. Chapman et al. 2005; Clements et al. 2008; Micha\u0142owski, Hjorth & Watson 2010). The unlensed 870\u2009\u03bcm flux of this object would be \u223c7.7 mJy.","Citation Text":["Fu et al. (2012)"],"Functions Text":["Additional followup observations in the optical, near-IR, submm and other wavelengths were targeted at the lensed z = 3.26 source and the foreground objects responsible for the lensing, resulting in detailed analyses of this lensing system by","and Bussmann et al. (2013)."],"Functions Label":["Background","Background"],"Citation Start End":[[985,1001]],"Functions Start End":[[742,984],[1002,1029]]}
{"Identifier":"2022MNRAS.517.1313M__Meidt_et_al._2018_Instance_1","Paragraph":"Star formation is an inefficient process, as evidenced by observed gas depletion times,1 which are two orders of magnitude above the dynamical time, both in galaxies (e.g. Leroy et al. 2017; Utomo et al. 2018), and in individual giant molecular clouds (GMCs) (e.g. Krumholz & Tan 2007; Evans, Heiderman & Vutisalchavakul 2014; Heyer et al. 2016; Pokhrel et al. 2020; Hu et al. 2022). Theoretical models explain this inefficiency through a combination of mechanisms that provide support against gravitational collapse, including turbulence, magnetic fields, stellar feedback, and dynamical stabilization (Krumholz & McKee 2005; Ostriker, McKee & Leroy 2010; Federrath & Klessen 2012; Krumholz, Klein & McKee 2012b; Federrath 2013b; Padoan et al. 2014; Federrath 2015; Burkhart 2018; Meidt et al. 2018; Krumholz & Federrath 2019; Evans, Kim & Ostriker 2022). Recent progress in both theory and observations have highlighted the pivotal role that feedback, especially due to massive (main-sequence) stars, plays in star\/star-cluster formation (Krumholz et al. 2014; Krumholz, McKee & Bland-Hawthorn 2019), and the lifecycle of GMCs (see Chevance et al. 2020, 2022a for reviews). This massive-star feedback has been suggested to be largely responsible for limiting the integrated star formation efficiency (\u03f5*) to low values in typical environments, where \u03f5* is given by\n(1)$$\\begin{eqnarray}\r\n\\epsilon _* = \\frac{M_{*}}{M_{\\mathrm{gas}}},\r\n\\end{eqnarray}$$which quantifies the net efficiency of star formation over the lifetime of a GMC, i.e. the ratio of the final stellar mass M* and the available gas mass in the parent molecular cloud Mgas. Feedback achieves this by (i) disrupting GMCs in order \u223c unity dynamical time-scales, through the momentum and energy carried by feedback processes (e.g. Grudi\u0107 et al. 2018), and (ii) driving turbulent motions that could further provide support against collapse (e.g. Mac Low & Klessen 2004; Krumholz, Matzner & McKee 2006; Elmegreen 2009; Gritschneder et al. 2009; Federrath et al. 2010; Wibking, Thompson & Krumholz 2018; Gallegos-Garcia et al. 2020; Menon, Federrath & Kuiper 2020; Menon et al. 2021).","Citation Text":["Meidt et al. 2018"],"Functions Text":["Theoretical models explain this inefficiency through a combination of mechanisms that provide support against gravitational collapse, including turbulence, magnetic fields, stellar feedback, and dynamical stabilization"],"Functions Label":["Background"],"Citation Start End":[[782,799]],"Functions Start End":[[384,602]]}
{"Identifier":"2015ApJ...799..149J___2014_Instance_1","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)"],"Functions Text":["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"],"Functions Label":["Uses"],"Citation Start End":[[304,324]],"Functions Start End":[[139,303]]}
{"Identifier":"2020MNRAS.499.2575E__Madau,_Shen_&_Governato_2014_Instance_1","Paragraph":"We note in passing that recent studies address improved satellite modellimg that ameliorates many of these issues, including the core\u2013cusp issue via non-sphericity of the stellar velocity distribution (Hayashi, Chiba & Ishiyama 2020) and the detectability of MWG satellites (Nadler et al. 2020). Other proposed solutions include those invoking baryonic physics, ranging from inclusion of baryon-contraction-induced diversity (Lazar et al. 2020), through dynamical friction-mediated coupling with baryonic clumps (El-Zant, Shlosman & Hoffman 2001; El-Zant et al. 2004; Tonini, Lapi & Salucci 2006; Romano-D\u00edaz et al. 2008; Goerdt et al. 2010; Cole, Dehnen & Wilkinson 2011; Del Popolo et al. 2014; Nipoti & Binney 2015), or through dynamical feedback driven by starbursts or active galactic nuclei (AGNs; Read & Gilmore 2005; Mashchenko, Couchman & Wadsley 2006; Mashchenko, Wadsley & Couchman 2008; Peirani, Kay & Silk 2008; Governato et al. 2012; Pontzen & Governato 2012; Zolotov et al. 2012; Martizzi, Teyssier & Moore 2013; Teyssier et al. 2013; Madau, Shen & Governato 2014; Ogiya & Mori 2014; Pontzen & Governato 2014; El-Zant, Freundlich & Combes 2016; Silk 2017; Freundlich et al. 2020). Alternatively, modifications to the particle physics model of the dark matter have been proposed. Such proposals include \u2018pre-heated\u2019 warm dark matter (e.g. Col\u00edn, Avila-Reese & Valenzuela 2000; Bode, Ostriker & Turok 2001; Macci\u00f2 et al. 2012; Schneider et al. 2012; Shao et al. 2013; Lovell et al. 2014; El-Zant, Khalil & Sil 2015) and self-interacting dark matter, whereby energy flows into the central cores of haloes through conduction (e.g. Burkert 2000; Kochanek & White 2000; Spergel & Steinhardt 2000; Miralda-Escud\u00e9 2002; Peter et al. 2013; Zavala, Vogelsberger & Walker 2013; Elbert et al. 2015). Ultralight axions, with boson mass \u223c10\u221222 eV, have also been considered as dark matter candidates in connection with these same small (sub)galactic scale problems (e.g. Peebles 2000; Hu, Barkana & Gruzinov 2000; Peebles 2000; Marsh & Silk 2014; Schive et al. 2014b; Hui et al. 2017; Mocz et al. 2019; Nori et al. 2019; see Niemeyer 2019 for recent review). Here the zero-point momentum associated with a long de Broglie wavelength corresponding to the small mass comes along with \u2018fuzziness\u2019 in particle positions. This in turn leads to a hotter halo core with non-diverging central density and a cut-off in halo mass. Such axion fields can also be relevant for inflationary scenarios or late dark energy models. The non-thermal production implies that the axions are present with the required abundance for dark matter; they behave as cold dark matter on larger scales despite the tiny masses (Marsh 2016, 2017).","Citation Text":["Madau, Shen & Governato 2014"],"Functions Text":["Other proposed solutions include those","or through dynamical feedback driven by starbursts or active galactic nuclei (AGNs;"],"Functions Label":["Background","Background"],"Citation Start End":[[1050,1078]],"Functions Start End":[[296,334],[720,803]]}
{"Identifier":"2019ApJ...887..185S__Pushkarev_et_al._2010_Instance_1","Paragraph":"The radio versus optical DCF peaks at non-zero lags (\u03c4delay) with the optical emissions leading the radio emissions, implying that the radio and optical emission regions are not cospatial with the optical\/IR emission region being closer to the base of the jet. This is consistent with the observation that optical emission regions are smaller than radio emission regions (Table 8). These observations are in accordance with standard shock-in-jet models where higher frequencies are emitted closer to the shock front while lower frequencies are produced from larger volumes that extend farther away from the shock (e.g., Marscher & Gear 1985; Marscher et al. 2008). This could also be understood as a manifestation of the position offset of optically thick features that can be interpreted as a frequency-dependent shift of the self-absorbed core of the jet (e.g., Lobanov 1998; Pushkarev et al. 2012). The linear separation of the V and 37 GHz emission region can be estimated using the relation (Pushkarev et al. 2010; Lisakov et al. 2017):\n10\n\n\n\n\n\nwhere \u03b2app is the apparent jet speed and \u03b8 is the viewing angle. Using a range of 15\u2013100 days for \u03c4delay from Table 6, the maximum linear separation between the emission regions (\n\n\n\n\n\n) is estimated to be in the range of 4.06 (for Segment 6) to 27.04 pc (for Segment 1) and the corresponding projected separation varies from 0.18 to 1.20 pc using a viewing angle of \u03b8 = 13. The resulting angular separation is 0.022\u20130.151 mas. The simplest jet geometry is that of a conical jet. It cannot, however, explain the change in the separation between emission regions. In a conical jet geometry, the distance of an emission region from the central engine can be calculated using \n\n\n\n\n\n (Abdo et al. 2011) assuming the emission region fills the cross-section of the jet and the opening angle \n\n\n\n\n\n. The obtained dce is given in Table 8 and the separation comes out to be 1.4 pc. Thus, the conical jet model also severely underestimates the separation between the emission regions, as it does not take into account the jet collimation. An alternative model is that of an inhomogeneous curved jet, where synchrotron radiation of decreasing frequency is produced in an outer and wider jet region that changes orientation with time. It is possible that the long-term variability behavior of 3C 454.3 during our extended observation is dominated by geometrical effects that also lead to temporal delays between the radio and optical bands.","Citation Text":["Pushkarev et al. 2010"],"Functions Text":["The linear separation of the V and 37 GHz emission region can be estimated using the relation"],"Functions Label":["Uses"],"Citation Start End":[[997,1018]],"Functions Start End":[[902,995]]}
{"Identifier":"2015MNRAS.448..666S__Steidel_et_al._2014_Instance_1","Paragraph":"In local galaxies, the ISM conditions are often described by some physical quantities such as ionization parameter (q), gaseous metallicity (Z) and electron density (ne). At high redshift, the ionization parameter is raised by a large flux of ionizing photons in ISM originated from hot O, B stars due to intensive star formation in relatively small galaxies. Previous studies suggest a high ionization parameter of SF galaxies at z > 2 compared to that of local galaxies (Erb et al. 2010; Nakajima et al. 2013; Masters et al. 2014; Nakajima & Ouchi 2014). Secondly, the chemical abundance of SF galaxies at z \u223c 2 is lower by 0.1\u20130.3\u2009dex for a given stellar mass compared to those at low-z (Erb et al. 2006a; Sanders et al. 2014; Steidel et al. 2014). This leads to more compact and hotter O, B stars due to lower opacity (Ezer & Cameron 1971; Maeder 1987), and thus UV radiation becomes harder and produces more ionizing photons. Thirdly, the strength of collisionally excited emission lines (e.g. [O\u2009iii], [N\u2009ii]) strongly depends on the electron density. It is closely related to the number of electrons to collide since the excitation potential of this line is \u223c1 eV, which is nearly the same as the energy of electrons at the virial temperature (\u223c104 K) (Dyson & Williams 1980). Due to low excitation potential, the transition of collisionally excited line is reliant on electron density compared to gaseous metallicity. Recent observations have suggested a high electron density (ne > 100 cm\u22123) in SF galaxies at z \u223c 2 (Newman et al. 2012; Masters et al. 2014; Shirazi, Brinchmann & Rahmati 2014; Wuyts et al. 2014). This value is larger than that of normal SF galaxies at low-z by an order of magnitude, and close to that of interacting galaxies which are seen as (ultra)luminous infrared (IR) galaxies in the present-day Universe (Krabbe et al. 2014). Such a large electron density contributes to the offset of galaxy distributions on the BPT diagram together with other physical parameters (Brinchmann, Pettini & Charlot 2008a). In this way, the cosmic dependence of the BPT diagram can be attributed to such physical parameters which determine ISM conditions.","Citation Text":["Steidel et al. 2014"],"Functions Text":["Secondly, the chemical abundance of SF galaxies at z \u223c 2 is lower by 0.1\u20130.3\u2009dex for a given stellar mass compared to those at low-z"],"Functions Label":["Background"],"Citation Start End":[[730,749]],"Functions Start End":[[557,689]]}
{"Identifier":"2022MNRAS.512.3243S__Mason_&_Gronke_2020_Instance_1","Paragraph":"A more equitable view of the catalogue is given by considering the distributions of band values across all haloes. In Fig. 14, we quantify the relative probability for a given integrated ($\\mathcal {T}_\\rm{IGM}^\\rm{\\,int}$, bottom panels) and maximum ($\\mathcal {T}_\\rm{IGM}^\\rm{\\,max}$, top panels) band transmission at each redshift. For reference, we also show the cumulative distribution functions as dashed curves and the median and 1\u03c3 summary statistics in the middle panels. It is extremely unlikely to have non-negligible transmission spikes in the blue bands at z \u2273 8, although this is certainly allowed (int) or even common (max) at z \u2272 6. When also considering the central band in rare cases it is plausible to witness multiple-peaked or otherwise complex spectral line profiles (e.g. see the discussions by Byrohl & Gronke 2020; Mason & Gronke 2020; Gronke et al. 2021; Park et al. 2021). This also stresses the need for systemic tracers beyond \u2009Ly\u03b1 to distinguish IGM signatures and pinpoint the origins of various spectral features. Perhaps more significant is the broad distribution in the red band. The transmission is relatively high after the midpoint of reionization z \u2272 7.67, and roughly follows the global ionized fraction (see Fig. 13). However, as will be shown later the exponential sensitivity on optical depth ($\\mathcal {T}_\\rm{IGM} = e^{-\\tau _\\rm{IGM}}$) allows the same galaxy to have both large and small values, i.e. the bimodality is not due to halo mass or environment alone. In fact, there will always be sightlines with very low transmission due to filaments and other self-shielding structures common at high-z. The location and broad nature of the higher transmission peak is redshift dependent, which is an important consideration when using LAEs as probes of reionization. We emphasize that the sharp cutoffs below $\\mathcal {T}_\\rm{IGM} \\approx 1$ are not physical but are due to the global treatment of the long-range damping-wing absorption. For example, at z = 7 the red band values cannot exceed $\\mathcal {T}_\\rm{IGM} \\approx 0.95$ as every sightline on average experiences at least $5{{\\ \\rm per\\ cent}}$ absorption via cosmological integration throughout the remainder of the EoR. Finally, we note that the ultra red band is less susceptible to resonant scattering and halo proximity effects and is therefore a cleaner probe of the global state of the IGM if this can be robustly measured with deep spectroscopy for large numbers of high-z galaxies.","Citation Text":["Mason & Gronke 2020"],"Functions Text":["When also considering the central band in rare cases it is plausible to witness multiple-peaked or otherwise complex spectral line profiles (e.g. see the discussions by","This also stresses the need for systemic tracers beyond \u2009Ly\u03b1 to distinguish IGM signatures and pinpoint the origins of various spectral features."],"Functions Label":["Background","Motivation"],"Citation Start End":[[841,860]],"Functions Start End":[[650,818],[901,1046]]}
{"Identifier":"2015AandA...584A.103S__Potekhin_et_al._2013_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":["Potekhin et al. 2013"],"Functions Text":["More recently, unified EoSs for NS have been derived by the Brussels-Montreal group"],"Functions Label":["Background"],"Citation Start End":[[580,600]],"Functions Start End":[[433,516]]}
{"Identifier":"2016AandA...594A..64P__Judge_(2015)_Instance_2","Paragraph":"There is now renewed interest in the literature concerning these transitions, because some of the O\u2009iv and S\u2009iv intercombination lines, together with the Si\u2009iv resonance lines, are routinely observed with the Interface Region Imaging Spectrograph (IRIS; De Pontieu et al. 2014) at much higher spectral, spatial and temporal resolution than previously. For example, Peter et al. (2014) used the intensities of the O\u2009iv vs. Si\u2009iv lines to propose that very high densities, on the order of 1013 cm-3 or higher, are present in the so-called IRIS plasma \u201cbombs\u201d. Line ratios involving an O\u2009iv forbidden transition and a Si\u2009iv allowed transition have been used in the past to provide electron densities during solar flares and transient brightenings (e.g. Cheng et al. 1981; Hanssen 1981). However, the validity of using O\u2009iv to Si\u2009iv ratios has been hotly debated because these ratios gave very high densities compared to the more reliable ones obtained from the O\u2009iv ratios alone (see, e.g. Hayes & Shine 1987). In addition, Judge (2015) recalled several issues that should be taken into account when considering the Si\u2009iv\/O\u2009iv density diagnostic. The main ones were: (1) O\u2009iv and Si\u2009iv ions are formed at quite different temperatures in equilibrium and hence a change in the O\u2009iv\/Si\u2009iv ratio could imply a change in the temperature rather than in the plasma density (2) the chemical abundances of O and Si are not known with any great accuracy and could be varying during the observed events (3) density effects on the ion populations could increase the Si\u2009iv\/O\u2009iv relative intensities by a factor of roughly three to four. Judge (2015) has also mentioned the well known problem of the \u201canomalous ions\u201d, that is, the observed high intensities of the Li- and Na-like (as Si\u2009iv) ions (see also Del Zanna et al. 2002). Another important aspect to take into account is the effect of non-equilibrium conditions on the observed plasma diagnostics. It is well known that strong variations in the line intensities are obtained when non-equilibrium ionisation is included in the numerical calculations (see, e.g. Shen et al. 2013; Raymond & Dupree 1978; Mewe & Schrijver 1980; Bradshaw et al. 2004). In particular, Doyle et al. (2013) and Olluri et al. (2013) investigated the consequences of time-dependent ionization on the formation of the O\u2009iv and Si\u2009iv transition region lines observed by IRIS. In addition, Dud\u00edk et al. (2014) showed that non-Maxwellian electron distributions in the plasma can substantially affect the formation temperatures and intensity ratios of the IRIS Si\u2009iv and O\u2009iv lines. These authors also suggested that the observing window used by IRIS should be extended to include S\u2009iv. Recent IRIS observation sequences have indeed included the S\u2009iv line near 1406 \u00c5. The S\u2009iv line ratios have a higher limit for density sensitivity than the O\u2009iv line ratios and are thus particularly useful for diagnosing high densities which might occur in flares. Previous flare studies have in fact reported line ratios involving O ions which lay above the density sensitivity range, indicating an electron density in the excess of 1012 cm-3 (e.g. Cook et al. 1995; Polito et al. 2016). ","Citation Text":["Judge (2015)"],"Functions Text":["has also mentioned the well known problem of the \u201canomalous ions\u201d, that is, the observed high intensities of the Li- and Na-like (as Si\u2009iv) ions"],"Functions Label":["Background"],"Citation Start End":[[1621,1633]],"Functions Start End":[[1634,1778]]}
{"Identifier":"2018MNRAS.478.2576M__Marleau_et_al._2017_Instance_1","Paragraph":"The search for low-mass BHs (MBH \u2272 106 M\u2299) in dwarf galaxies is mostly based on the detection of X-ray emission (e.g. Greene & Ho 2007a; Desroches, Greene & Ho 2009; Reines et al. 2011; Dong et al. 2012; Schramm et al. 2013; Baldassare et al. 2015, 2017; Lemons et al. 2015; Secrest et al. 2015; Pardo et al. 2016; Chen et al. 2017), in some cases spatially coincident with jet radio emission (e.g. Reines et al. 2014; Nucita et al. 2017), or the use of standard virial techniques to estimate the BH mass (e.g. Barth et al. 2004; Greene & Ho 2004, 2007b; Peterson et al. 2005; Reines, Greene & Geha 2013; La Franca et al. 2015; Bentz et al. 2016; Onori et al. 2017; Liu et al. 2018; Chilingarian et al. 2018; see Mezcua 2017 for a review). Additional searches in the infrared (IR) regime have yielded a few more candidates (e.g. Satyapal et al. 2007, 2008, 2009, 2014; Sartori et al. 2015; Marleau et al. 2017). Most of these samples are however incomplete, very local (z  0.3), skewed towards high Eddington ratios, or skewed towards type-1 active galactic nucleus (AGN) in the case of optical searches (e.g. Greene & Ho 2004, 2007b; Reines et al. 2013) which can hamper the detection of BHs lighter than 105 M\u2299 if the size of the broad-line region is controlled by BH mass (e.g. Chakravorty, Elvis & Ferland 2014). Baldassare et al. (2015) found an AGN with MBH \u223c 5 \u00d7 104 M\u2299 estimated using the virial technique, Yuan et al. (2014) performed a study of four low Eddington ratio sources, and Pardo et al. (2016) searched for AGN in dwarf galaxies out to z  1. Yet, these studies include very few sources. To circumvent the biases mentioned above, in Mezcua et al. (2016) we performed an X-ray stacking analysis of \u223c50 000 dwarf galaxies selected in the COSMOS field making use of the recently completed Chandra COSMOS-Legacy survey (Civano et al. 2016). We found that a population of IMBHs with X-ray luminosities \u223c1039\u22121040\u2009erg\u2009s\u22121 does exist in dwarf galaxies out to z = 1.5 and that their detection beyond the local Universe is most likely hampered by their low luminosity and mild obscuration unless deep surveys like the Chandra COSMOS-Legacy are used.","Citation Text":["Marleau et al. 2017"],"Functions Text":["Additional searches in the infrared (IR) regime have yielded a few more candidates (e.g."],"Functions Label":["Background"],"Citation Start End":[[890,909]],"Functions Start End":[[740,828]]}
{"Identifier":"2021ApJ...923..126S__Samsing_&_Ilan_2018_Instance_1","Paragraph":"The question is, which of these proposed merger channels dominate the merger rate? Are several channels operating with a possible dependence on redshift? Or are the majority of GW sources formed through a still unknown mechanism? Several studies show that one can distinguish at least classes of channels, such as isolated binaries and dynamically induced mergers, by considering the observed distribution of merger masses (Zevin et al. 2017) or the relative spin orientation of the merging objects (Rodriguez et al. 2016c), as well as the orbital eccentricity at some reference GW frequency (G\u00fcltekin et al. 2006; Samsing et al. 2014; Samsing & Ramirez-Ruiz 2017; Samsing & Ilan 2018; Samsing et al. 2018b; Samsing 2018; Samsing et al. 2018a; Samsing & D\u2019Orazio 2018; Rodriguez et al. 2018; Zevin et al. 2019; Samsing et al. 2019a, 2020). Other \u201cindirect\u201d probes have also been suggested, such as stellar tidal disruptions (e.g., Samsing et al. 2019b; Lopez et al. 2019; Kremer et al. 2019b). In this picture, it is now largely believed that dynamically assembled mergers are likely to have mass rations near one (e.g., Rodriguez et al. 2018), random relative spin orientations (e.g., Rodriguez et al. 2016c), and a nonnegligible fraction of mergers with measurable eccentricity in LISA (Samsing & D\u2019Orazio 2018; Kremer et al. 2019c), DECIGO\/Tian-Qin (e.g., Chen & Amaro-Seoane 2017; Samsing et al. 2020), and LIGO (Samsing 2018). This is in contrast to isolated binary mergers, which likely have correlated spins (e.g., Kalogera 2000), a bimodal distribution for the effective spin parameter (Zaldarriaga et al. 2018; Hotokezaka & Piran 2017; Piran & Piran 2020), larger mass ratios, and which merge on orbits with eccentricities indistinguishable from ~0 near LISA and LIGO. This picture is rather clean when comparing mergers forming in highly dynamical systems, such as globular clusters (GCs) and GNs, to completely isolated field binary mergers; however, it becomes less clean when considering, e.g., the proposed subpopulation of field binaries that undergo secular interactions with nearby single or binary objects (e.g., Naoz et al. 2013; Naoz 2016; Toonen et al. 2016; Antonini et al. 2017; Silsbee & Tremaine 2017; Liu & Lai 2018; Rodriguez & Antonini 2018; Randall & Xianyu 2018a; Antonini et al. 2018; Liu & Lai 2019; Fragione & Loeb 2019; Fragione & Kocsis 2019; Hamers & Thompson 2019; Safarzadeh et al. 2020). In this case, secular exchanges of especially angular momentum can drive the binary to merge with random spin orientations (e.g., Liu & Lai 2017) and notable eccentricities (e.g., Randall & Xianyu 2018b; Liu et al. 2019; Fragione & Kocsis 2020), which makes it more challenging to disentangle cluster mergers from field binary mergers.","Citation Text":["Samsing & Ilan 2018"],"Functions Text":["Several studies show that one can distinguish at least classes of channels, such as isolated binaries and dynamically induced mergers,","as well as the orbital eccentricity at some reference GW frequency"],"Functions Label":["Background","Background"],"Citation Start End":[[665,684]],"Functions Start End":[[230,364],[525,591]]}
{"Identifier":"2021ApJ...916...68B__Chakrabarti_1989_Instance_1","Paragraph":"In this paper, we presented our analysis of the spectral and timing behavior of GRS 1915+105 using the TCAF paradigm. For this purpose, the \u03b8 class data of the source as obtained by the LAXPC instrument of AstroSat satellite were used. To the best of our knowledge, this is the first time that the \u03b8 class data of AstroSat are analyzed with the TCAF paradigm. In this paradigm, different spectral states as well as the QPOs resulting from resonance oscillation of the Compton cloud arise out of the interplay between two types of accretion rates, namely, the Keplerian disk rate (\n\n\n\n\n\n\n\n\nm\n\n\n\n\n\n\n\nd\n\n\n\n\n) and the sub-Keplerian halo rate (\n\n\n\n\n\n\n\n\nm\n\n\n\n\n\n\n\nh\n\n\n\n\n). In the sub-Keplerian flow, shocks are formed when the Rankine\u2013Hugoniot conditions are satisfied (Chakrabarti 1989, 1996; Chakrabarti & Das 2004). If the Keplerian disk rate increases, it increases the soft-seed photons, and the post-shock region (CENBOL, which acts as a \u201cCompton cloud\u201d) is cooled down rapidly, causing the shock to proceed toward the black hole. This has indeed been obtained in our spectral analysis. In the process, both the compressional heating timescale and the cooling timescale of the CENBOL decrease and become comparable (Chakrabarti et al. 2015), triggering shock oscillations that manifest as QPOs in the light curve. We first obtained the state of the system by estimating the photon index using the phabs*(diskbb+power-law) model. Throughout our analysis, the photon index (\u0393) was found to be above 2.4, indicating that the source was either in the soft-intermediate state (SIMS) or in the soft state (SS). In the case of both orbits, \u0393 monotonically increases with the increase in photon flux, implying the transition from SS to SIMS. In the first orbit, \u0393 made an excursion from 2.5 to 3.0 during the 400 s span of spectral analysis, indicating the transition of the source from SIMS to SS. The gradual enhancement of the disk accretion rate and the consequent decline of the shock location were also obtained from the TCAF model fitting. The shock location moves from 39rs to 17rs, while the disk accretion rate increases from 0.77 to 0.84 \n\n\n\n\n\n\n\n\nM\n\n\n\n\n\n\n\nEdd\n\n\n\n\n. A similar trend was followed in the second orbit as well (Table 2). In the two orbits (02345 and 02346) we have analyzed, within a span of 400\u2013500 s, the total flux changes by a factor of 3\u20135. The corresponding shock velocity is \u223c2200 and 900 m s\u22121, respectively. This is significantly higher than the shock velocity obtained earlier in the case of transient sources, which lie within 10\u201320 m s\u22121 (Chakrabarti et al. 2008; Nandi et al. 2012; Debnath et al. 2013). This suggests the possibility that the local modulation of the accretion rate is due to feedback from the outflow (Chakrabarti & Manickam 2000; Chakrabarti & Nandi 2000). All these results indicate that GRS 1915+105 went through repeated microflares when it was in \u03b8 class.","Citation Text":["Chakrabarti 1989"],"Functions Text":["In the sub-Keplerian flow, shocks are formed when the Rankine\u2013Hugoniot conditions are satisfied"],"Functions Label":["Background"],"Citation Start End":[[763,779]],"Functions Start End":[[666,761]]}
{"Identifier":"2015ApJ...806..152S__Ransom_et_al._2005_Instance_3","Paragraph":"One of the most astonishing characteristics of Liller 1 is the extremely large value of the collision rate parameter. Verbunt & Hut (1987) showed that Liller 1 has the second-highest value of stellar encounter rate (after Terzan 5; see also Lanzoni et al. 2010) among all star clusters in the Galaxy, thus suggesting that it represents an ideal environment where exotic objects, generated by collisions, can form. In fact, it is commonly believed that dynamical interactions in GCs facilitate the formation of close binary systems and exotic objects such as cataclysmic variables (CVs), low-mass X-ray binaries (LMXBs), millisecond pulsars (MSPs), and blue straggler stars (BSSs) (e.g., Bailyn 1992; Paresce et al. 1992; Ferraro et al. 2001, 2009a, 2012; Ransom et al. 2005; Pooley & Hut 2006). Moreover, Hui et al. (2010) found that clusters with large collisional parameters and high metallicity (see also Bellazzini et al. 1995) usually host more MSPs. Indeed, Terzan 5 hosts the largest population of MSPs among all Galactic GCs (Ransom et al. 2005).\n6\n\n\n\n6\n\nNote that Terzan 5 is suspected to not be a genuine GC, because it harbors at least three stellar populations with different iron abundances (Ferraro et al. 2009a; Origlia et al. 2011, 2013; Massari et al. 2014).\n A strong \u03b3-ray emission has been recently detected in the direction of Liller 1 by the Large Area Telescope (LAT) on board Fermi (Tam et al. 2011). This is the most intense emission detected so far from a Galactic GC, again suggesting the presence of a large number of MSPs. However, no direct radio detection of these objects has been obtained so far in this system (Ransom et al. 2005). The only exotic object identified in the cluster is the rapid burster MXB 1730-335, an LMXB observed to emit radio waves and type I and type II X-ray bursts (Hoffman et al. 1978). It seems to be located in the central region of Liller 1, but no optical\/IR counterpart of this object has been found so far (Homer et al. 2001).","Citation Text":["Ransom et al. 2005"],"Functions Text":["This is the most intense emission detected so far from a Galactic GC, again suggesting the presence of a large number of MSPs. However, no direct radio detection of these objects has been obtained so far in this system"],"Functions Label":["Motivation"],"Citation Start End":[[1645,1663]],"Functions Start End":[[1425,1643]]}
{"Identifier":"2021AandA...647A..49S__Murphy_et_al._2015_Instance_1","Paragraph":"The \u03bb Bo\u00f6tis stars are a group of chemically peculiar objects on the upper main-sequence, showing underabundances (~1\u20132 dex) of iron-peak elements and near-solar abundances of C, N, O, and S (e.g., Kamp et al. 2001; Heiter 2002; Andrievsky et al. 2002). The class was discovered by Morgan et al. (1943) and named following the bright prototype \u03bb Bo\u00f6tis, which is one extreme member of the class. This group of refractory-poor objects comprises about 2% of early B through early F stars (Gray & Corbally 1998; Paunzen 2001). However, among the pre-main-sequence Herbig Ae\/Be stars, that is to say the progenitors of A-type stars, the \u03bb Bo\u00f6tis-like fraction is about 33% (Folsom et al. 2012). The origin of the peculiarity still remains a puzzle, as we can read in the recent discussion of Murphy & Paunzen (2017). Unlike common chemical peculiarities seen in Am and Ap stars, \u03bb Bo\u00f6tis stars are not constrained to slow rotation (Abt & Morrell 1995; Murphy et al. 2015). Cowley et al. (1982) first suggested that \u03bb Bo\u00f6tis could possibly originate from the interstellar medium (ISM) with a nonsolar composition, or from the separation of grains and gas. Then, Venn & Lambert (1990) proposed that \u03bb Bo\u00f6tis stars likely occur when circumstellar gas is separated from grains and then accreted to the stars due to the similarity ofits abundance pattern with the ISM. Other proposed mechanisms include the interaction of a star with a diffuse interstellar cloud (Kamp & Paunzen 2002; Martinez-Galarza et al. 2009), where the underabundances are produced by different amounts of accreted material. Turcotte & Charbonneau (1993) estimated that once the accretion stops, the photospheric mixing and meridional circulation would erase this peculiar signature on a ~1 Myr timescale. By studying thedistribution of \u03bb Bo\u00f6tis stars on the HR diagram, Murphy & Paunzen (2017) conclude that multiple mechanisms could result in a \u03bb Bo\u00f6tis spectra, depending on the age and environment of the star.","Citation Text":["Murphy et al. 2015"],"Functions Text":["Unlike common chemical peculiarities seen in Am and Ap stars, \u03bb Bo\u00f6tis stars are not constrained to slow rotation"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[948,966]],"Functions Start End":[[813,926]]}
{"Identifier":"2018ApJ...864..154D__Liu_et_al._2016_Instance_1","Paragraph":"Feedback from massive stars plays a critical role in the star formation process and evolution of molecular clouds. In particular, expanding H ii regions may have a positive effect on star formation, i.e., they can trigger a new generation of star formation in molecular clouds either by sweeping ambient clouds into dense shells or by compressing nearby dense clouds into bound clumps\/cores (for details, see Deharveng et al. 2010). In both cases, dense material eventually fragments to form new stars. Since the PG108.3 cloud consists of three evolved H ii regions, one can thus speculate that the formation of young sources in the region might have been triggered by the expanding bubble. However, our observations largely do not favor such a process because in such a scenario one would expect the distribution of young YSOs or star-forming cores at the outskirts of the H ii regions (see, e.g., Zavagno et al. 2006; Jose et al. 2013; Panwar et al. 2014; Liu et al. 2016; Samal et al. 2018). In contrast, we find that no young Class I sources around S147 and most of Class I sources around S148 are cospatial with the Class II sources located near the center of the ionized gas. Since young clusters often harbor Class III to Class I sources as a part of the cluster formation process (e.g., Jose et al. 2017; Panwar et al. 2017, 2018), the Class I YSOs of the S148 are thus most likely part of the central cluster responsible for the ionization of S148. Given the fact that the H ii region S149 is located near the boundary of S148 and smaller in size, one can argue that S149 might have been influenced by S148. Unfortunately, from the current data we have no way to estimate precisely either the age of the H ii regions or the embedded YSOs in S148. Nonetheless, we searched for signatures of the early phases of high-mass star formation such as masers and outflows in the vicinity of S148, and our search resulted in no such observations. The small size of the S149 H ii regions could be due to the fact that it is ionized by a less massive star compared to S148, and as a result, S149 is expanding at a slow rate compared to S148, assuming that both of them formed in a similar environment. Considering the fact that S149 is optically visible, its dynamical age is close to the age of S148, and the ionized gas pressures of both regions are within a factor of two, we regard the probability that the formation of S149 is solely due to expansion of S148 to be rather low, although we cannot ignore the possibility that the expanding ionization front of S148 might have accelerated star formation in S149 after the initial clump formation and fragmentation. A way to prove such a hypothesis is to compare the molecular gas pressure of S149 with external gas pressure generated by S148 (e.g., Thompson et al. 2004; Kim et al. 2017; Liu et al. 2017) in the early stage of star formation; however, in S149, neither do we find significant cold gas nor do Azimlu & Fich (2011) find any molecular clumps to obtain some clue in this direction.","Citation Text":["Liu et al. 2016"],"Functions Text":["Since the PG108.3 cloud consists of three evolved H ii regions, one can thus speculate that the formation of young sources in the region might have been triggered by the expanding bubble. However, our observations largely do not favor such a process because in such a scenario one would expect the distribution of young YSOs or star-forming cores at the outskirts of the H ii regions (see, e.g.,","In contrast, we find that no young Class I sources around S147 and most of Class I sources around S148 are cospatial with the Class II sources located near the center of the ionized gas."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[958,973]],"Functions Start End":[[503,898],[995,1181]]}
{"Identifier":"2021MNRAS.503.2108P__Andresen_et_al._2017_Instance_2","Paragraph":"CCSNe are also of interest for GW astronomy as targets in their own right. As the sensitivity of GW detectors increases, they will begin to detect not only binary mergers but also other lower amplitude sources of GWs such as CCSNe. Accurate knowledge of the GW emission from CCSNe will be essential for detection and parameter estimation. The GW signal from rotational core bounce has already been well covered in the literature (e.g. Dimmelmeier et al. 2008; Abdikamalov et al. 2014; Fuller et al. 2015; Richers et al. 2017). In the non-rotating case, the GW emission from the post-bounce phase has been studied using self-consistent 3D simulations by many groups (Kuroda, Kotake & Takiwaki 2016; Andresen et al. 2017, 2019; Kuroda et al. 2017, 2018; Powell & M\u00fcller 2019, 2020; Radice et al. 2019; Andresen, Glas & Janka 2020; Mezzacappa et al. 2020; Pan et al. 2020). The structure of the GW emission has shown common features in different simulations from recent years. The dominant emission feature in the GW emission is due to the quadrupolar surface f\/g mode 1 of the proto-neutron star (PNS), which produces GW frequencies rising in time from a few hundred Hz up to a few kHz (M\u00fcller, Janka & Wongwathanarat 2012; Sotani et al. 2017; Kuroda et al. 2018; Morozova et al. 2018; Torres-Forn\u00e9 et al. 2018, 2019). In addition, some models (Kuroda et al. 2016, 2017; Andresen et al. 2017; Mezzacappa et al. 2020; Powell & M\u00fcller 2020) exhibit low-frequency GW emission due to the standing accretion shock instability (SASI; Blondin, Mezzacappa & DeMarino 2003; Blondin & Mezzacappa 2006; Foglizzo et al. 2007). In rapidly rotating models, very strong GW emission can also occur during the post-bounce phase due to a corotation instability (Takiwaki & Kotake 2018). The emerging understanding of the GW emission features has led to the formulation of universal relations for the GW emission (Torres-Forn\u00e9 et al. 2019) and paved the way for phenomenological modelling for CCSN signals (Astone et al. 2018). Further work is still needed, however, to extend these models to fully explore CCSN GW signals from across the progenitor parameter space. The majority of 3D simulations that include GW emission are for progenitor stars below $30\\, \\mathrm{M}_{\\odot }$. In this paper, we perform simulations of high-mass Population III (Pop-III) stars in the pulsational pair instability regime to expand the parameter space coverage of 3D simulations and to provide further insights into the massive and very massive star remnant BH population.","Citation Text":["Andresen et al. 2017"],"Functions Text":["In addition, some models","exhibit low-frequency GW emission due to the standing accretion shock instability (SASI"],"Functions Label":["Background","Background"],"Citation Start End":[[1369,1389]],"Functions Start End":[[1317,1341],[1437,1524]]}
{"Identifier":"2015MNRAS.452.1112E__Illarionov_&_Sunyaev_1975_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":["Illarionov & Sunyaev 1975"],"Functions Text":["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","see also"],"Functions Label":["Uses","Uses"],"Citation Start End":[[888,913]],"Functions Start End":[[640,842],[879,887]]}
{"Identifier":"2022ApJ...928....3A__Rappazzo_et_al._2019_Instance_1","Paragraph":"Solar vortex tubes can be spontaneously generated by turbulent convection. In simulations of quiet Sun regions, vortices are found along intergranular lanes (Shelyag et al. 2011a; Kitiashvili et al. 2012; Moll et al. 2012; Silva et al. 2020). These structures have an average lifetime of around 80 s (Silva et al. 2021) and a radius between 40 and 80 km (Shelyag et al. 2013; Silva et al. 2020). Solar kinetic vortex tubes (Silva et al. 2021) act as a sink for magnetic field, creating magnetic flux tubes that expand with height (Kitiashvili et al. 2012; Moll et al. 2012; Silva et al. 2020). The concentration of magnetic flux leads to a high magnetic field tension, which can prevent the magnetic field lines from being twisted by the rotational motion (Shelyag et al. 2011b; Moll et al. 2012; Nelson et al. 2013; Silva et al. 2021). In some cases, twisted magnetic flux tubes appear close enough to flow vortices, leading to magnetic and kinetic vortex structures closely coexisting in regions with high plasma-\u03b2 (Wedemeyer & Steiner 2014; Rappazzo et al. 2019; Silva et al. 2021). The vortical motions can still trigger perturbations along magnetic lines that could lead to wave excitation, e.g., Battaglia et al. (2021). The vorticity evolution in the magnetized solar atmosphere is mainly ruled by the magnetic field, which also influences the general shape of vortices (Shelyag et al. 2011a). Based on the analysis of swirling strength, the part of the vorticity only linked to swirling motion (Shelyag et al. 2011b; Canivete Cuissa & Steiner 2020) showed that the magnetic terms in the swirling equation evolution tend to cancel the hydrodynamic terms close to the solar surface, whereas the magnetic terms dominate alone the production of swirling motion in the chromosphere. The magnetic field also tends to play an important role in the plasma dynamics along the whole vortex tube, as the Lorentz force has a magnitude comparable to the pressure gradient (Silva et al. 2020; Kitiashvili et al. 2013). High-speed flow jets have also been linked to simulated vortex tubes, driven by high-pressure gradients close to the photosphere and by Lorentz force in the weakly magnetized upper solar photosphere (Kitiashvili et al. 2013). In general, the averaged radial profile of magnetic field, angular velocity, pressure gradient inside of the vortex tube at the lower chromosphere and photosphere levels show similar behavior (Silva et al. 2020).","Citation Text":["Rappazzo et al. 2019"],"Functions Text":["In some cases, twisted magnetic flux tubes appear close enough to flow vortices, leading to magnetic and kinetic vortex structures closely coexisting in regions with high plasma-\u03b2"],"Functions Label":["Background"],"Citation Start End":[[1044,1064]],"Functions Start End":[[837,1016]]}
{"Identifier":"2021MNRAS.503.5367B__Lu,_Kumar_&_Zhang_2020_Instance_2","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"],"Functions Text":["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"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1320,1342]],"Functions Start End":[[1164,1318]]}
{"Identifier":"2019MNRAS.485.5652D__Yamanaka_et_al._2009_Instance_1","Paragraph":"Einstein\u2019s theory has been tested successfully mainly in the regime of weak gravity through Solar system tests and laboratory experiments. But the validity of this highly accepted theory still faces stringent constraint in the regime of strong gravity, viz., the region near to a black hole, ultradense compact stars, and expanding universe. The recent discovery of peculiar highly overluminous SNeIa, e.g. SN 2003fg, SN 2006gz, SN 2007if, and SN 2009dc (Howel et al. 2006; Scalzo et al. 2010) indicates a huge Ni-mass and confirms the highly super-Chandrasekhar white dwarfs, having mass 2.1\u20132.8\u2009M\u2299, as a suitable progenitors (Howel et al. 2006; Hicken et al. 2007; Yamanaka et al. 2009; Scalzo et al. 2010; Silverman et al. 2011; Taubenberger et al. 2011). Recently, Linares, Shahbaz & Casares (2018) have discovered a highly massive pulsar of mass $2.27^{+0.17}_{-0.15}~\\mathrm{M}_{\\odot }$ in their observation of compact binaries PSR J2215+5135. Clearly, these observations are not only questioning the standard Chandrasekhar limit for the compact stellar objects but also invoking the necessity of modification of GR in the strong gravity regime. Interestingly, our study reveals that due to the effect of $f\\left(R,\\mathcal {T}\\right)$ gravity theory, the maximal mass limits rise higher than their standard values in GR for the chosen parametric values of \u03c7. Hence, the stellar systems in the framework of $f\\left(R,\\mathcal {T}\\right)$ gravity theory may also explain the observed massive stellar systems, viz., massive pulsars, super-Chandrasekhar stars, and magnetars, etc., which GR hardly can explain suitably so far. In support of the achieved result in the present investigation, it is worth mentioning that the important study by De Laurentis (2018) also reveals that the application of the Noether Symmetry Approach can explain the extreme massive stars which are supposed to be unstable in the framework of GR. However, in the limit \u03c7 = 0, one may retrieve the solutions of the standard Einstein gravity.","Citation Text":["Yamanaka et al. 2009"],"Functions Text":["The recent discovery of peculiar highly overluminous SNeIa,","indicates a huge Ni-mass and confirms the highly super-Chandrasekhar white dwarfs, having mass 2.1\u20132.8\u2009M\u2299, as a suitable progenitors"],"Functions Label":["Background","Background"],"Citation Start End":[[667,687]],"Functions Start End":[[342,401],[494,626]]}
{"Identifier":"2017ApJ...850...81P__Mu\u00f1oz-Jaramillo_et_al._2012_Instance_1","Paragraph":"Solar cycle predictions are needed to plan long-term space missions and are of high importance for space weather applications. Currently, precursor methods are the most favored models for the prediction of solar cycle strength (Conway 1998; Svalgaard et al. 2005; Kane 2008; Hathaway 2009). These precursor techniques often use geomagnetic activity levels near or before the time of solar cycle minimum (Ohl & Ohl 1979; Feynman 1982; Gonzalez & Schatten 1988; Thompson 1993; Wilson et al. 1998). Predicting the amplitude of a solar cycle can be done using solar polar magnetic fields from the previous cycle as \u201cprecursors\u201d of the next cycle (Schatten & Sofia 1987; Schatten 2005; Svalgaard et al. 2005; Wang & Sheeley 2009; Mu\u00f1oz-Jaramillo et al. 2012). The other class of precursor techniques that do not need an a priori physical understanding of the causal relations (i.e., that do not require any knowledge of the physics involved) is based on finding particular sunspot number characteristics that serve as indicators of the next cycle strength (Ramaswamy 1977; Lantos 2006; Cameron & Sch\u00fcssler 2008; Braj\u0161a et al. 2009). A number of techniques are used to predict the amplitude of a cycle during the time near and before sunspot minimum (Hathaway 2010). The two precursor types that have received the most attention are polar field precursors and geomagnetic precursors. The strength of the solar polar magnetic fields at solar minimum is a very accurate indicator of the maximum amplitude of the following solar cycle. Forecasts using the polar field method have proven to be consistently in the right range for cycles 21, 22, and 23 (Schatten & Sofia 1987; Schatten et al. 1996). The polar fields reach their maximal amplitude near the minima of the sunspot cycle. However, the maxima of the polar field curves are often rather flat, so approximate forecasts are feasible several years before the actual minimum (Hathaway 2010; Petrovay 2010). Using the rather flat and low maximum in polar field strength, Svalgaard et al. (2005) have been able to predict a relatively weak cycle 24. Such an early prediction is not always possible: early polar field predictions of cycles 22 and 23 had to be corrected later and only forecasts made shortly before the actual minimum finally converged (Hathaway 2010; Petrovay 2010).","Citation Text":["Mu\u00f1oz-Jaramillo et al. 2012"],"Functions Text":["Predicting the amplitude of a solar cycle can be done using solar polar magnetic fields from the previous cycle as \u201cprecursors\u201d of the next cycle"],"Functions Label":["Background"],"Citation Start End":[[725,752]],"Functions Start End":[[496,641]]}
{"Identifier":"2020AandA...641A.123H__Mikal-Evans_et_al._(2019)_Instance_4","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)"],"Functions Text":["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"],"Functions Label":["Similarities"],"Citation Start End":[[2047,2072]],"Functions Start End":[[1809,2046]]}
{"Identifier":"2021MNRAS.504.4626K__Kraljic_et_al._2020b_Instance_3","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"],"Functions Text":["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","using filament density as a proxy for the thickness of filaments"],"Functions Label":["Similarities","Background"],"Citation Start End":[[1504,1524]],"Functions Start End":[[1343,1502],[1526,1590]]}
{"Identifier":"2017ApJ...835..151A__Gaisser_et_al._2014_Instance_1","Paragraph":"In the southern sky, the large background of atmospheric muons reduces the efficiency to select through-going tracks induced by neutrinos below the PeV regime. A very large fraction of the aforementioned background can be rejected by imposing an active veto at the detector boundary, as, for example, used in Aartsen et al. (2013a). This reduces the detector volume to a smaller fraction of the instrumented volume sacrificing statistics for signal purity. Furthermore, the more clearly an event is identified as a starting track, the more probable it is to be an astrophysical rather than atmospheric background. Down-going atmospheric neutrino events at high energy are likely to be accompanied by muons produced in the same cosmic-ray shower that triggers the veto and reduces the atmospheric neutrino background (Sch\u00f6nert et al. 2009; Gaisser et al. 2014). In analyses using veto techniques (Aartsen et al. 2013a, 2015b), the selection is usually more efficient for cascade-like events than tracks, and high astrophysical neutrino purity demands neglecting energies below 60 TeV, where backgrounds are more abundant. In searches for point-like sources of astrophysical neutrinos, track-like events are of great importance given their good angular resolution compared to cascade-like events. Furthermore, the purity demands are lower since the signal of a point-like source is restricted to a small portion of the sky, hence reducing the background significantly. Consequently, the minimum required total charge deposited in the PMTs of the IceCube detector by an event is lowered to 1500 p.e. compared to 6000 p.e. (Aartsen et al. 2013a), resulting in a higher signal efficiency at lower energies. In addition, only down-going tracks are used, and cuts are imposed that select well-reconstructed track-like events (Aartsen et al. 2016a). For \n\n\n\n\n\n events at energies smaller than 200 TeV, the effective area of the analysis is bigger than for vertically through-going tracks (\u03b4  \u221230\u00b0; Figure 1). For energies up to 1 PeV, the effective area is smaller, but a higher purity is achieved. The angular resolution for starting tracks is shown in Figure 2 (dashed) and is \u223c1\u00b0 in the interesting energy region; the reconstruction is worse than for through-going events (solid), due to a smaller lever arm for tracks starting within the fiducial volume of the detector.","Citation Text":["Gaisser et al. 2014"],"Functions Text":["Furthermore, the more clearly an event is identified as a starting track, the more probable it is to be an astrophysical rather than atmospheric background. Down-going atmospheric neutrino events at high energy are likely to be accompanied by muons produced in the same cosmic-ray shower that triggers the veto and reduces the atmospheric neutrino background"],"Functions Label":["Background"],"Citation Start End":[[839,858]],"Functions Start End":[[457,815]]}
{"Identifier":"2020ApJ...899L...6L__Margalit_et_al._2019_Instance_2","Paragraph":"The leading FRB source model invokes magnetars as the power source to produce repeating bursts. There are two versions of this model. One version invokes rapidly spinning young magnetars that are produced in extreme stellar transients such as GRBs and SLSNe. The main motivation is that the host galaxy of FRB 121102 resembles those of LGRBs and SLSNe (Metzger et al. 2017; Nicholl et al. 2017; Wadiasingh & Timokhin 2019). The fact that the hosts of all other FRBs do not resemble that of FRB 121102 disfavors the simplest version of this proposal. A possible fix of this proposal is to introduce rapidly spinning magnetars born from binary neutron star (BNS) mergers (Margalit et al. 2019; Wang et al. 2020). In order to make this scenario work, one needs to require that rapidly spinning magnetars made from BNS mergers should be much more abundant than those made from LGRBs and SLSNe. Comparing the event rate densities of BNS mergers, LGRBs, and SLSNe (e.g., Sun et al. 2015; Abbott et al. 2017; Nicholl et al. 2017), this may be possible if a significant fraction of BNS mergers leave behind stable neutron stars (e.g., Gao et al. 2016). However, if this fraction is very low, as required if GW170817 leaves behind a black hole (Margalit et al. 2019), the fast magnetar model may fail to explain the small fraction of LGRB\/SLSN-like hosts in FRB samples. The second version of the magnetar model invokes emission (e.g., giant flares) from slowly rotating magnetars like the ones observed in the Galaxy (e.g., Popov & Postnov 2010; Katz 2014; Kulkarni et al. 2014). The births of these magnetars do not require extreme explosions such as GRBs and SLSNe (e.g., Beniamini et al. 2019). If this is the case, the host galaxy distribution may be more analogous to that of SNe II. All FRBs but FRB 121102 are consistent with this scenario (Figure 4). In order to interpret FRB 121102, the more extreme channel of forming rapid magnetars is still needed. So we conclude that the magnetar model would work only if both fast magnetars produced in extreme explosions and slow magnetars produced in regular channels (Beniamini et al. 2019) can produce FRBs. In any case, since the birth rate of these magnetars is very high (Beniamini et al. 2019), an additional factor is needed to select a small fraction of magnetars to produce FRBs (e.g., Ioka & Zhang 2020).","Citation Text":["Margalit et al. 2019"],"Functions Text":["However, if this fraction is very low, as required if GW170817 leaves behind a black hole","the fast magnetar model may fail to explain the small fraction of LGRB\/SLSN-like hosts in FRB samples."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1236,1256]],"Functions Start End":[[1145,1234],[1259,1361]]}
{"Identifier":"2018ApJ...863..162M__Liu_et_al._2013_Instance_1","Paragraph":"NLFFF extrapolation provides the reconstructed coronal magnetic field for AR 11158 from 2011 February 13 \u2212 2011 February 15 (Figures 1(d)\u2013(f)). The field lines (yellow lines) within the core of the AR have arcade-like structure with a relatively strong twist mainly near the PIL. These figures show that the magnetic field evolved during this period. Although we did not quantitatively compare the field lines with the observation, in general, the reconstructed coronal field morphologies match with the observations in Figures 1(a)\u2013(c). The general morphologies and the locations of the high-twist fields are also in agreement with many previous studies (Jing et al. 2012; Sun et al. 2012; Dalmasse et al. 2013; Inoue et al. 2013, 2014a; Liu et al. 2013; Wang et al. 2013; Aschwanden et al. 2014; Malanushenko et al. 2014; Zhao et al. 2014). Unlike Zhao et al. (2014) who could identify the twisted flux rope from the topology of the reconstructed coronal field, we could not find an obvious topological signature of a flux rope existing in our NLFFF during our analysis time window. It might be due to the fact that there was little magnetic flux with twist higher than one turn in our NLFFF and it is difficult to topologically define it as a flux rope. However, our result is consistent with other NLFFF results (Jing et al. 2012; Sun et al. 2012; Liu et al. 2013; Wang et al. 2013; Inoue et al. 2014a; Malanushenko et al. 2014). The high-twist region in our result is also in agreement with the region with high helicity flux (Dalmasse et al. 2013) and the location of the flare ribbons (Bamba et al. 2013; Liu et al. 2013), as well as the high current density region (Janvier et al. 2014). Figures 1(g)\u2013(i) show the evolution of the twist distribution map, with the magnetic twist of the field lines plotted at the footpoints of field lines according to a color scale. This shows that the high-twist (strongly right-handed twist corresponding to Tw > 0.5) areas are concentrated in only a limited part of the AR. The high-twist area grew and became even more twisted just before the X2.2 flare (Figure 1(i)). Most parts of the AR have twist values less than 0.25, but near the PIL the twist can reach more than 0.5, even up to about a full turn. This is consistent with the results of Sun et al. (2012) and Inoue et al. (2014a). A high-twist (strong negative\/left-handed twist) area also developed in the eastern part of the AR, which did not exist initially on February 13. Both of these high-twist areas produced several flares. However, here we focus on the flares that resulted from the high-twist core region near the center of the AR, where the M6.6 and X2.2 flares occurred.","Citation Text":["Liu et al. 2013"],"Functions Text":["The general morphologies and the locations of the high-twist fields are also in agreement with many previous studies"],"Functions Label":["Similarities"],"Citation Start End":[[739,754]],"Functions Start End":[[538,654]]}
{"Identifier":"2021MNRAS.502..772L__Cowley_et_al._2015_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":["Cowley et al. 2015"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[859,877]],"Functions Start End":[[667,857]]}
{"Identifier":"2016AandA...588A..74C__B\u00f6hm-Vitense_1958_Instance_1","Paragraph":"The pulsational analysis presented in this work makes use of full stellar evolution models of pre-WDs generated with the LPCODE stellar evolution code. LPCODE computes in detail the complete evolutionary stages leading to WD formation, allowing the WD and pre-WD evolution to be studied in a consistent way based on the evolutionary history of progenitors. Details of LPCODE can be found in Althaus et al. (2005, 2009, 2013) and references therein. Here, we mention only those ingredients employed which are relevant for our analysis of low-mass, He-core WD and pre-WD stars (see Althaus et al. 2013, for details). The standard Mixing Length Theory (MLT) for convection in the version ML2 is used (Tassoul et al. 1990). In this prescription, due to Bohm & Cassinelli (1971), the parameter \u03b1 (the mixing length in units of the local pressure scale height) is set equal to 1, while the coefficients a,b,c that appear in the equations for the average speed of the convective cell, the average convective flux, and the convective efficiency (see Cox 1968), have values a = 1,b = 2,c = 16. We emphasize that the results presented in this work are insensitive to the prescription of the MLT employed. In particular, we have also used the ML1 (\u03b1 = 1,a = 1 \/ 8,b = 1 \/ 2,c = 24, B\u00f6hm-Vitense 1958) and ML3 (\u03b1 = 2,a = 1,b = 2,c = 16, Tassoul et al. 1990) recipes, and we obtain the same results than for ML2. The metallicity of the progenitor stars has been assumed to be Z = 0.01. It is worth mentioning that the pulsation stability results presented in this paper do not depend on the value of Z2. Radiative opacities for arbitrary metallicity in the range from 0 to 0.1 are from the OPAL project (Iglesias & Rogers 1996). Conductive opacities are those of Cassisi et al. (2007). The equation of state during the main sequence evolution is that of OPAL for H- and He-rich compositions. Neutrino emission rates for pair, photo, and bremsstrahlung processes have been taken from Itoh et al. (1996), and for plasma processes we included the treatment of Haft et al. (1994). For the WD regime we have employed an updated version of the Magni & Mazzitelli (1979) equation of state. The nuclear network takes into account 16 elements and 34 thermonuclear reaction rates for pp-chains, CNO bi-cycle, He burning, and C ignition. Time-dependent diffusion due to gravitational settling and chemical and thermal diffusion of nuclear species has been taken into account following the multicomponent gas treatment of Burgers (1969). Abundance changes have been computed according to element diffusion, nuclear reactions, and convective mixing. This detailed treatment of abundance changes by different processes during the WD regime constitutes a key aspect in the evaluation of the importance of residual nuclear burning for the cooling of low-mass WDs. ","Citation Text":["B\u00f6hm-Vitense 1958"],"Functions Text":["In particular, we have also used the ML1 (\u03b1 = 1,a = 1 \/ 8,b = 1 \/ 2,c = 24,"],"Functions Label":["Uses"],"Citation Start End":[[1271,1288]],"Functions Start End":[[1195,1270]]}
{"Identifier":"2015AandA...574A..70F__Padova_1994_Instance_1","Paragraph":"We used the stellar population synthesis code STARLIGHT (Cid Fernandes et al. 2004) to describe the age distributions and metallicities of the stellar populations that fit the integrated light spectrum of object I. For object H it is not possible to perform the stellar population synthesis because of the low signal-to-noise ratio of the spectrum. This code is extensively discussed in Cid Fernandes et al. (2004; 2005), and it is built upon computational techniques that originally were developed for empirical population synthesis with additional ingredients from evolutionary synthesis models. This method was used by Krabbe et al. (2011) and Fa\u00fandez-Abans et al. (2012, 2013) and has been successful in describing the stellar population in interacting galaxies. Briefly, the code fits an observed spectrum with a combination of N single stellar populations (SSPs) from the models of Bruzual & Charlot (2003). These models are based on a high-resolution library of observed stellar spectra, which allows for detailed spectral evolution of the SSPs across the wavelength range of 3200\u22129500 \u00c5 \u2009with a wide range of metallicities. We used the Padova 1994 tracks, as recommended by Bruzual & Charlot (2003), with the Salpeter initial mass function (Salpeter 1955). Extinction is modeled by STARLIGHT as due to foreground dust, using the Large Magellanic Cloud average reddening law of Gordon et al. (2003) with RV = 3.1, and parametrized by the V-band extinction AV. The SSPs used in this work cover 15 ages, t = 0.001, 0.003, 0.005, 0.01, 0.025, 0.04, 0.1, 0.3, 0.6, 0.9, 1.4, 2.5, 5, 11, and 13 Gyr, and three metallicities, Z = 0.2Z\u2299, 1 Z\u2299, and 2.5 Z\u2299, adding to 45 SSP components. The fitting is carried out using a simulated annealing plus Metropolis scheme, with regions around emission lines and bad pixels excluded from the analysis. The upper panel of Fig. 4 shows the observed spectrum corrected for reddening and the model stellar population integrated light spectrum of object I. The synthesized spectrum fits the observed one closely. The results indicate that this galaxy is an old object, with 100% of the flux contribution at \u03bb5870 \u00c5 being provided by a stellar population of 13 Gyr, with Z = 2.5Z\u2299 and AV = 0.36. This spectrum is dominated principally by the bulge signal and does not represent the whole galaxy; the age estimate then is valid only for its bulge. ","Citation Text":["Padova 1994"],"Functions Text":["We used the","tracks, as recommended by Bruzual & Charlot (2003)"],"Functions Label":["Uses","Uses"],"Citation Start End":[[1144,1155]],"Functions Start End":[[1132,1143],[1156,1206]]}
{"Identifier":"2022AandA...667A.131B__Izumi_et_al._(2016)_Instance_3","Paragraph":"Molecular line ratio diagnostics are often used to investigate the physics and chemistry of the ISM in all of these environments. For example, as the gas chemistry located in the central, nuclear regions of galaxies is believed to be dominated by X-rays produced by the AGN, in so-called X-ray dominated regions (XDRs), the molecular content of the ISM surrounding such nuclei will greatly differ from that in starburst regions (Usero et al. 2004; Garc\u00eda-Burillo et al. 2010). Hence, line ratios of specific molecules have been proposed as indicators of certain energetic or physical processes, for example HCN\/HCO+ as a tracer of AGNs (Loenen et al. 2007), HCN\/HNC as a mechanical heating tracer (Hacar et al. 2020), and HCN\/CO as a density tracer (Leroy et al. 2017). In particular, the \u201csubmillimeter-HCN diagram\u201d, first proposed in Izumi et al. (2013) and later expanded upon in Izumi et al. (2016), is a very notable example of the use of molecular line ratios as a probe of AGN-galaxies; this diagram makes use of two line ratios, HCN(4\u22123)\/HCO+(4\u22123) and HCN(4\u22123)\/CS(7\u22126), where all of the molecules involved are considered tracers of dense gas. Izumi et al. (2016) observed a clear trend that AGNs, including the Seyfert composite galaxy NGC 1068, tend to show higher HCN\/HCO+ and\/or HCN\/CS than in SB galaxies as long as the observations were at high enough spatial resolutions to separate energetically discrete regions. Izumi et al. (2016) propose a scenario where it is the high temperature that is responsible for the HCN enhancement, whereby neutral-neutral reactions with high reaction barriers are enhanced (Harada et al. 2010), thus leading to the possible enhancement of HCN and the depletion of HCO+ via newly available formation and destruction paths, respectively. However, while of course higher gas temperatures are expected in AGN-dominated regions, these are not unique to these environments, as starburst regions and\/or regions where outflows dominate can also harbour high enough temperatures for such enhancement to occur. Additionally, the higher temperatures could increase HCN excitation, relative to HCO+ and CS, without necessarily changing their relative abundances (Imanishi et al. 2018a). Finally, infrared radiative pumping is also a possible explanation of the HCN intensity enhancement relative to HCO+ and CS. Infrared pumping is a result of the emission of 14 \u03bcm infrared photons due to the presence of hot dust around AGN. These photons vibrationally excite HCN to the \u03bd2\u2004=\u20041 state. Upon de-exciting from this state back to the vibrational ground state, \u03bd\u2004=\u20040, the HCN line intensities are thus pumped to higher fluxes (Imanishi et al. 2018a). However, we note that it is also not unlikely that the 12 \u03bcm infrared photons can similarly vibrationally excite HCO+, thus nullifying the extent of this effect (Imanishi et al. 2016).","Citation Text":["Izumi et al. (2016)"],"Functions Text":["propose a scenario where it is the high temperature that is responsible for the HCN enhancement, whereby neutral-neutral reactions with high reaction barriers are enhanced","thus leading to the possible enhancement of HCN and the depletion of HCO+ via newly available formation and destruction paths, respectively.","However, while of course higher gas temperatures are expected in AGN-dominated regions, these are not unique to these environments, as starburst regions and\/or regions where outflows dominate can also harbour high enough temperatures for such enhancement to occur. Additionally, the higher temperatures could increase HCN excitation, relative to HCO+ and CS, without necessarily changing their relative abundances","Finally, infrared radiative pumping is also a possible explanation of the HCN intensity enhancement relative to HCO+ and CS. Infrared pumping is a result of the emission of 14 \u03bcm infrared photons due to the presence of hot dust around AGN. These photons vibrationally excite HCN to the \u03bd2\u2004=\u20041 state. Upon de-exciting from this state back to the vibrational ground state, \u03bd\u2004=\u20040, the HCN line intensities are thus pumped to higher fluxes","However, we note that it is also not unlikely that the 12 \u03bcm infrared photons can similarly vibrationally excite HCO+, thus nullifying the extent of this effect"],"Functions Label":["Background","Background","Motivation","Motivation","Motivation"],"Citation Start End":[[1429,1448]],"Functions Start End":[[1449,1620],[1643,1783],[1784,2197],[2223,2658],[2684,2844]]}
{"Identifier":"2017ApJ...849...63R__Pascucci_et_al._2016_Instance_1","Paragraph":"The (sub)mm wavelength range is of particular interest for various reasons: at sufficiently long wavelengths, disks become optically thin, and an estimate of their dust mass can be directly obtained (via some assumptions) by simply measuring their flux (e.g., Beckwith et al. 1990). Although the bulk of the disk mass in the system is in gaseous phase, fiducial (or measured, when available) gas-to-dust ratios provide an indirect estimate of the total mass in the disk. This is a crucial parameter for planet formation theories because it determines the available reservoir for this process. Using this method, surveys of star-forming regions with (sub)mm facilities such as SMA and ALMA have determined that protoplanetary disks have typical masses of 0.1%\u20130.5% of that of their host star (e.g., Andrews & Williams 2005; Andrews et al. 2013; Pascucci et al. 2016). On the other hand, dust growth represents the initial stage of planet formation; the observed spectral index at these wavelengths can be linked to the dust opacity in the disk, Informative of its properties and grain sizes (e.g., Miyake & Nakagawa 1993; Draine 2006). In fact, the comparison of the millimeter spectral index of the interstellar medium (ISM) with that of protoplanetary disks has already revealed significant dust growth in these disks, implying the presence of mm\/cm-sized grains in many of them (e.g., D\u2019Alessio et al. 2001; Lommen et al. 2010; Ricci et al. 2010a, 2010b; Ubach et al. 2012). The combination of the mm spectral index with additional information at other wavelengths, such as the spectral index at near\/mid infrared (IR) wavelengths or silicate features may also point to links between the evolution of the inner and outer regions of the disks. As an example, Lommen et al. (2010) identified a tentative correlation between the strength of the 10 \u03bcm silicate feature and the 1\u20133 mm spectral index for a sample of T Tauri and Herbig Ae\/Be stars, suggesting a connection between the evolution of the inner and outer regions of disks, although a later study by Ricci et al. (2010b) found no signs of such a correlation for disks in the Taurus and Ophiuchus star-forming regions. Despite the obvious interest of this wavelength regime, disks have relatively weak emission at millimeter wavelengths and many of them currently lack this type of data (or, at least, sufficient observations to provide robust estimates of their spectral indices).","Citation Text":["Pascucci et al. 2016"],"Functions Text":["Using this method, surveys of star-forming regions with (sub)mm facilities such as SMA and ALMA have determined that protoplanetary disks have typical masses of 0.1%\u20130.5% of that of their host star (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[844,864]],"Functions Start End":[[593,797]]}
{"Identifier":"2015MNRAS.454.1644L__Kotze_&_Charles_2012_Instance_1","Paragraph":"The period candidates of other three ULXs may range from \u223c100 to \u223c600 d. Apart from noise and artefacts, all the candidate periods are only significant in a specific epoch. This suggests that they are not associated with any stable mechanism such as orbital motion. Instead, such long-term (> 100 d) X-ray quasi-periodic variations are likely related to superorbital periods that are thought to be due to radiation-driven warping of accretion discs (Ogilvie & Dubus 2001) or tidal interaction-induced disc precession (Whitehurst & King 1991). Alternatively, mass transfer rate-related events such as X-ray state changes and disc instability can also cause long-term modulations (Kotze & Charles 2012). In particular, there are two intermittent quasi-periodicities for both NGC 5408 X-1 and M81 X-6, suggesting that the quasi-periods are changing or evolving. They resemble some Galactic X-ray binaries that show similar behaviour (e.g. Cyg X\u20132 and SMC X\u20131; Kotze & Charles 2012) and it has been suggested that a warped disc could lead to an unstable steadily precessing disc, causing quasi-periodic behaviour (Ogilvie & Dubus 2001). We note that there are many uncertainties on the physical parameters of ULXs. To determine the origin of superorbital periods of ULXs, one has to know at least the mass ratio between the companion and the compact star (q = MC\/MX) and the binary separation. Unfortunately, it is very difficult to get these parameters for ULXs. For the three ULXs discussed here (i.e. excluding ESO 243-49 HLX-1), only M81 X-6 has better constraints on the black hole mass and the nature of the companion. The masses of the black hole and companion star are estimated (MX = 18\u2009M\u2299, MC = 23\u2009M\u2299) such that q can be derived. In this case, we can rule out a tidal interaction-induced disc precession scenario that requires q  0.25\u20130.33 (Whitehurst & King 1991). For a warped disc, the binary separation and the mass ratio suggest that M81 X-6 lies in the intermediate instability zone for radiation-driven warping in X-ray binaries (see fig. 1 of Kotze & Charles 2012). The quasi-periodic variability may represent the switching time-scale between a warped disc and a flat disc.","Citation Text":["Kotze & Charles 2012"],"Functions Text":["Alternatively, mass transfer rate-related events such as X-ray state changes and disc instability can also cause long-term modulations"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[679,699]],"Functions Start End":[[543,677]]}
{"Identifier":"2021AandA...653A..99T__Pall\u00e9_et_al._2004_Instance_1","Paragraph":"To minimize the undesired effects caused by observing different lunar locations, we conducted observations according to the following procedure. On each observing night, we first pointed the Nayuta telescope toward the crater Grimaldi (selenographic coordinate: 68.6\u00b0W, 5.2\u00b0S) in the waxing phase and the crater Neper (84.5\u00b0E, 8.8\u00b0N, east of Mare Crisium) in the waning phase, after correcting the pointing error measured using a nearby star. Both craters are near the lunar edge (distances \u22722\u2032). Then, we scanned the Moon along the RA axis until the edge of the Moon was placed near the center of the field of view (FOV). An example of the observed (and reduced) images is shown in Fig. 2. Our target locations are not on a major maria and near sites repeatedly observed in previous Earthshine photometry because they were expected to have roughly comparable albedos (Qiu et al. 2003; Pall\u00e9 et al. 2004; Monta\u00f1\u00e9s-Rodr\u00edguez et al. 2007). Half of the FOV was reserved for the sky, which allows the sky background intensities and their positional gradients to be measured. The position angle of the instrument (\u03d5inspa) was maintained at 90\u00b0 from the equatorial north, as measured counter-clockwise, so that the long side of the FOV was aligned with the RA axis. Telescope tracking was conducted in accordance with the sky motion of the Moon, which was calculated at the Jet Propulsion Laboratory (JPL) Horizons system1. Because the tracking was not perfect, we shifted the telescope east or west with a typical interval of ~ 30 min so that the lunar edge remained near the center of the FOV. Features on the Moon were hardly recognizable in the raw images because of the dim Earthshine and strong scattered light from the day side of the Moon, though we were able to visually identify the lunar edge in most cases2. Despite our efforts, the actually observed location may have varied night by night even within one phase (waxing phase or waning phase), or on an hourly timescale during a single night. Possible impacts induced by different lunar locations (namely different degrees of depolarization) are discussed in Sect. 4.1.","Citation Text":["Pall\u00e9 et al. 2004"],"Functions Text":["Our target locations are not on a major maria and near sites repeatedly observed in previous Earthshine photometry because they were expected to have roughly comparable albedos"],"Functions Label":["Motivation"],"Citation Start End":[[886,903]],"Functions Start End":[[691,867]]}
{"Identifier":"2022ApJ...940...86K__Cohen_et_al._1997_Instance_1","Paragraph":"Currently, about 20 Be stars are known to have stripped companions, most of which were confirmed with far-UV (FUV) spectroscopy, as it is in the FUV where the flux ratios are most favorable (Wang et al. 2021; Klement et al. 2022, and references therein). All of the spectroscopic FUV detections were found to be compatible with the sdO nature of the companions, while no firm case of a cooler sdB companion has been presented as of yet. There are also more than 160 confirmed and candidate Be X-ray binaries (BeXRBs; Raguzova & Popov 2005),\n12\n\n\n12\nList updated at http:\/\/xray.sai.msu.ru\/~raguzova\/BeXcat\/. which are mostly Be+NS systems that probably evolved in a similar fashion as the Be+sdO binaries but from more massive progenitor systems (Reig 2011). The Be primaries occupy a narrow range between O9 and B2 in spectral type (Reig et al. 2017). These systems are conspicuous due to the X-ray emission resulting from the (episodic) accretion of Be disk material onto the compact object, so that the sample is drawn from a much larger (partly extragalactic) volume. The WD companions to Be stars proved unexpectedly elusive (Meurs et al. 1992; Cohen et al. 1997), but several supersoft X-ray emission sources consistent with (early-)Be+WD systems undergoing a type II BeXRB outburst were recently detected in the Magellanic Clouds (Coe et al. 2020; Kennea et al. 2021). One Be star (MWC 656) was reported to have a BH companion, but this was recently shown to be questionable on the basis of new higher-quality spectra, which rather point toward another Be+sdO system (Rivinius et al. 2022). Object HD 93521 is the first candidate postmerger Be star (Gies et al. 2022), and members of a Be star runaway population were found using Hipparcos and Gaia astrometric catalogs (Berger & Gies 2001; Boubert & Evans 2018; Wang et al. 2022). On the other hand, Be stars have been found missing among B stars at high Galactic latitude (Martin2004, 2006). Meanwhile, not a single early-type Be star has been confirmed to have a close MS companion (Gies 2000; Bodensteiner et al. 2020). However, the recently studied case of the B6Ve star \u03b1 Eri, which is a highly eccentric binary with an early A-type dwarf companion on a 7 yr orbit, appears to be the first confirmed case of a Be star that does not owe its nature to mass transfer in a close binary, as the presence of a close stripped companion was ruled out (Kervella et al. 2022b). This implies that two evolutionary channels\u2014single and binary\u2014indeed exist for the formation of Be stars.","Citation Text":["Cohen et al. 1997"],"Functions Text":["The WD companions to Be stars proved unexpectedly elusive"],"Functions Label":["Background"],"Citation Start End":[[1149,1166]],"Functions Start End":[[1071,1128]]}
{"Identifier":"2022MNRAS.517.3881L__GRB_2004_Instance_1","Paragraph":"During a GF, a huge amount of the magnetic energy E > 1046 erg is subsequently released at the surface of the magnetar within the first 1 s, leading to the formation of a hot fireball similar to the case of classic GRBs (M\u00e9sz\u00e1ros & Rees 2000), which might be rich in electron\u2013positron pairs (Thompson & Duncan 1995; Fermi-LAT Collaboration 2021). The afterglow emission of a GRB is generally well explained by the synchrotron emission from electrons accelerated by the shock produced during a relativistic ejecta colliding with an external medium. The ejecta might be formed by the baryon-rich crust of the NSs undergoing a large-scale disruption, or a dense outflowing gas of electron\u2013positron pairs and radiations (Thompson 2021). Zhang et al. (2020) prefer a high fraction of electron energy for GRB 200415A, which is consistent with the situation of an electron\u2013positron pair-dominating ejecta. Chand et al. (2021) suggest that a baryon-poor outflow can explain the high-energy afterglow emission of GRB 200415A, while a baryon-rich outflow is also viable if the dissipation happens below the photosphere via internal shocks. Therefore, in this paper, we calculate two situations that the ejecta is dominated by electron\u2013positron pairs or protons, respectively. Around the magnetar, a cavity might be created by the pulsar wind and the earlier SGR activity. Holcomb et al. (2014) suggest that the Poynting flux from a pulsar can evacuate a cavity with a size of sub-parsec. Such a cavity is also needed to explain the radio afterglow of the GF from SGR 1806-20 (Gaensler et al. 2005; Gelfand et al. 2005; Taylor et al. 2005; Granot et al. 2006). Fermi-LAT Collaboration (2021) suggests a cavity environment to explain the expansion of the ejecta before its colliding with the bow-shock shell for GRB 200415A. Calculations in Zhang et al. (2020) on the afterglow of GRB 200415A also suggest a low density of ambient medium, which is consistent with the cavity environment. In this paper, we discuss two situations for the CSM density, including a density of $n = 10^{-2}\\ \\rm cm^{-3}$, and a low density of $n = 10^{-5}\\ \\rm cm^{-3}$ for the situation of a cavity environment surrounding the magnetar.","Citation Text":["Chand et al. (2021)"],"Functions Text":["suggest that a baryon-poor outflow can explain the high-energy afterglow emission of GRB 200415A, while a baryon-rich outflow is also viable if the dissipation happens below the photosphere via internal shocks. Therefore, in this paper, we calculate two situations that the ejecta is dominated by electron\u2013positron pairs or protons, respectively."],"Functions Label":["Motivation"],"Citation Start End":[[899,918]],"Functions Start End":[[919,1265]]}
{"Identifier":"2016AandA...587A..30M__Mihalas_&_Binney_1981_Instance_1","Paragraph":"For a star that moves with respect to its surrounding medium, the stellar motion adds an asymmetry to the wind velocity profile, since different parts of the wind face the ISM with different relative velocities. If the motion is supersonic, a bow shock arises at the interface where the ram pressure of the ISM and the stellar wind balance. The stand-off distance, i.e. the distance of the star to the apex of the shock front, is given by (1)\\begin{equation} \\label{R0_eq} R_0=\\sqrt{\\frac{\\dot{M} v_w}{4 \\pi \\rho_0 v_*^2}}, \\end{equation}R0=M\u0307vw4\u03c0\u03c10v\u22172,where vw is the terminal wind velocity, v\u2217 the stellar velocity with respect to the ISM, \u1e40 the mass-loss rate, and \u03c10 the density of the surrounding medium (Baranov et al. 1971). The density can be expressed in number density of hydrogen atoms (mH = 1.6727 \u00d7 10-27\u2009kg), which follows roughly (2)\\begin{equation} \\label{density} n_{\\rm H} = 2.0 \\; {\\rm e}^{-\\frac{|z|}{100\\,{\\rm pc}}}, \\end{equation}nH=2.0e\u2212|z|100\u2009pc,where z is the galactic height (Mihalas & Binney 1981) and nH is given in atoms per cm3. Wilkin (1996) demonstrated that the shape of the bow shock only depends on the stand-off distance, while Cox et al. (2012) showed that this assumption remains valid for viewing angles up to 70\u00b0. Above this value, the bow shock cone becomes broader. Therefore, we were able to use Eq. (1) to estimate the mass-loss rate from the binary system by measuring the stand-off distance. Generally, the ISM density and stellar velocity can be determined following Eq. (2) and Johnson & Soderblom (1987), respectively. While the error of the space motion is negligible, the ISM density value is only an estimate since the star could move through a dense cloud, which is not considered by Eq. (2). The respective values of these quantities for the three objects are given in Table 1. To obtain the space motion (v\u2217 ,LSR) with respect to the local standard of rest, we corrected the heliocentric motions from the solar motion vector (U,V,W)\u2299 = (8.50 \u00b1 0.29,13.38 \u00b1 0.43,6.49 \u00b1 0.26)\u2009km\u2009s-1 (Co\u015fkuno\u01e7lu et al. 2011). However, since the proper motions are quite small (a few mas\u2009yr-1), the correction for the solar motion has a large impact, especially on the PA of the motion. Interestingly, the PA of the corrected LSR motion is a worse match to the bow-shock orientation than the PA of the uncorrected motion (see Figs. 1, 3, and 5). Using (U,V,W)\u2299 = (10.00 \u00b1 0.36,5.25 \u00b1 0.62,7.17 \u00b1 0.38)\u2009km\u2009s-1 determined from the Hipparcos data by Dehnen & Binney (1998) leads to vLSR velocities which are slightly closer to the better matching heliocentric values. A similar discrepancy between bow-shock orientation and LSR motion (and less so with heliocentric motion) is found by Peri et al. (2012, 2015) for a large number of O- and B-type stars with bow shocks, as collected in the WISE E-BOSS survey. For this reason, we overplotted both the heliocentric and LSR motions on the WISE images in Figs. 1, 3, and 5. ","Citation Text":["Mihalas & Binney 1981"],"Functions Text":["The density can be expressed in number density of hydrogen atoms (mH = 1.6727 \u00d7 10-27\u2009kg), which follows roughly (2)\\begin{equation} \\label{density} n_{\\rm H} = 2.0 \\; {\\rm e}^{-\\frac{|z|}{100\\,{\\rm pc}}}, \\end{equation}nH=2.0e\u2212|z|100\u2009pc,where z is the galactic height","and nH is given in atoms per cm3."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1002,1023]],"Functions Start End":[[732,1000],[1025,1058]]}
{"Identifier":"2021MNRAS.508.1020D__Schuhmann_et_al._2019_Instance_1","Paragraph":"A few possibilities are explored for XCl: (1) The most straightforward explanation is that the Cl radical itself is present in the coma. Although it seems unlikely due to its reactivity, the CN radical was discovered in the coma of 67P (H\u00e4nni et al. 2020); therefore, the Cl radical cannot be excluded as a possible candidate. (2) Next to CH3Cl (Fayolle et al. 2017), which is too low in abundance to be a suitable neutral candidate to explain the observations, no other volatile Cl-bearing species (e.g. like Cl2) have been identified in the DFMS mass spectra. However, as all results show that EII product ion fractions decrease as the complexity of the molecules increases (Schuhmann et al. 2019) and since DFMS spectra for masses >100 have not been thoroughly analysed for the whole mission, there is still a possibility for other Cl-bearing species to be present within the DFMS mass range, albeit with low abundances. (3) One cannot exclude a priori that there would be heavier (semi)volatile Cl-bearing parents at higher m\/z outside the DFMS mass range that sublimate and are present in the coma. However, the fragmentation of such parents, e.g. of the form CxHyCl, in the DFMS ion source would create Cl-bearing fragments at lower m\/z. Alternatively, heavier semivolatile chlorine-bearing neutrals might undergo photolysis rather than sublimation, releasing intermediate chlorine-bearing neutrals at lower m\/z into the coma, which in turn fragment into smaller ions in the DFMS ion source. Both processes are improbable since apart from CH3Cl no Cl-bearing fragments at lower m\/z have been found and the effect of photodissociation should be much more pronounced during perihelion when Rosetta was farther away from the comet and this was not observed. (4) A source of Cl may lie in compounds that decompose upon sublimation and\/or ionization. As an example, it is known that ammonium perchlorate (NH4ClO4) releases HCl upon warming (Boldyrev 2006). Other perchlorates or compounds with oxidized states of Cl may provide relatively more Cl+ than HCl+ upon dissociation. According to Schauble, Rossman & Taylor (2003), molecules with oxidized Cl are enriched with 37Cl relative to non-oxidized species, which could also explain the somewhat higher coma isotopic values observed.","Citation Text":["Schuhmann et al. 2019"],"Functions Text":["However, as all results show that EII product ion fractions decrease as the complexity of the molecules increases"],"Functions Label":["Motivation"],"Citation Start End":[[677,698]],"Functions Start End":[[562,675]]}
{"Identifier":"2016ApJ...832....1R__Plekan_et_al._2011_Instance_1","Paragraph":"When dipolar species, such as CO, are allowed to deposit in the laboratory on a cold surface under vacuum to form a layer of approximately 4\u20135 monolayers (ML), the molecules may spontaneously orient throughout the layer such that their positive or negative ends protrude from the surface, creating a polarization potential (Balog et al. 2009; Cassidy et al. 2012; Field et al. 2013; Lasne et al. 2015; Rosu-Finsen et al. 2016). We have developed a semi-empirical description of this so-called spontelectric effect, which describes major features of the phenomenon. This description is based on a physical model that invokes spontaneous orientation of the dipolar molecules of which the film is composed. The analysis, described in detail in Field et al. (2013), is remarkably successful in describing important characteristics of spontelectrics, such as the deposition temperature dependence of the surface polarization charge, including the counterintuitive behavior of solid methyl formate (Plekan et al. 2011). In the case of CO, the positive oxygen end protrudes on average (Collings et al. 2014). The spontelectric effect is exemplified by a powerful spontaneous electric field in the film. It was first discovered for nitrous oxide using a direct measurement of surface potential, employing high-resolution, low-energy electron beams (Balog et al. 2009). Studies have now been extended to CO using RAIR (reflection-absorption infrared) spectroscopy (Rosu-Finsen et al. 2016). RAIR spectra were recorded of CO laid down to a thickness of 20 ML, dosed at 21, 22, and 24 K, on 50 ML of compact amorphous solid water (cASW) formed by dosing at 110 K, in order to reproduce as far as possible the nature of ices on grains in a prestellar core. The presence of spontaneous electric fields in the film of CO was demonstrated by measurements of Stark frequency shifts in RAIR spectra, the values of which are film deposition temperature dependent. These observations were used to derive the electric field in the film, yielding a temperature-averaged surface potential, \u03d5, of 6.7 \u00b1 0.5 mV per ML of CO (Rosu-Finsen et al. 2016). We note that experimental observations show that the surface potential for any spontelectric material depends on the temperature of the substrate at the moment of deposition, which is \u226426 K in the ISM. The subsequent cooling of the substrate to 10 K in the ISM has no effect on the surface potential.","Citation Text":["Plekan et al. 2011"],"Functions Text":["The analysis, described in detail in Field et al. (2013), is remarkably successful in describing important characteristics of spontelectrics, such as the deposition temperature dependence of the surface polarization charge, including the counterintuitive behavior of solid methyl formate"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[993,1011]],"Functions Start End":[[704,991]]}
{"Identifier":"2021ApJ...912...20J__Murray_et_al._1995_Instance_1","Paragraph":"Dramatic changes in the optical broad emission lines (BELs) of CLQs can also probe the physical origin of the lines themselves. Optical spectra of AGN are characterized primarily by a power-law continuum, BELs, and narrow emission lines (e.g., Vanden Berk et al. 2001). BELs are emitted from the broad-line region (BLR), which is believed to be the high velocity gas gravitationally bound to the central supermassive black hole (Peterson & Wandel 2000; Peterson et al. 2004). Results from reverberation mapping (Denney et al. 2009; De Rosa et al. 2018), and study of the quasar orientation and observed width of BEL profiles (Shen & Ho 2014; Storchi-Bergmann et al. 2017) imply that the geometry of the BLR is likely to be disk-like. Furthermore, study of BEL profiles suggests the BLR gas is structured as a smooth continuous flow (Laor et al. 2006). However, the exact origin of the BLR gas is still unclear (Elvis 2017). The disk-wind model suggests that the BLR gas comes from winds produced by the accretion disk (Emmering et al. 1992; Murray et al. 1995; Murray & Chiang 1997; Elitzur & Ho 2009; Elitzur et al. 2014; Elitzur & Netzer 2016). When the mass accretion rate drops below a certain limit (corresponding to a luminosity of \u223c\n\n\n\n\n4.7\n\n\u00d7\n\n\n10\n\n\n39\n\n\n\n\n\nM\n\n\n7\n\n\n2\n\n\/\n\n3\n\n\n\nerg\n\n\n\ns\n\n\n\u2212\n1\n\n\n\n\n, where \n\n\n\n\n\n\nM\n\n\n7\n\n\n2\n\n\/\n\n3\n\n\n\n\n is the black hole mass in \n\n\n\n\n\n\n10\n\n\n7\n\n\n\n\n\nM\n\n\n\u2299\n\n\n\n\n), winds can no longer be sustained, such that the observed BELs \u201cdisappear\u201d (Elitzur & Ho 2009). Below the critical luminosity, AGN are expected to be \u201ctrue\u201d Type 2 AGN (i.e., Type 2 AGN intrinsically absent of BLRs). Although this critical luminosity is dependent on the black hole mass, the disk-wind model predicts that more detections of BEL disappearance are expected below the 1% Eddington ratio (Elitzur & Netzer 2016), and a double-peaked BEL profile, the signature of a rotating disk, should emerge in the quasar spectrum when the accretion rate drops (Elitzur et al. 2014). Single-epoch observations show that Eddington ratios of bright quasars with prominent BELs, are always higher than 1% (Kollmeier et al. 2006; Steinhardt & Elvis 2010). Yet it is still unknown whether Eddington ratios of individual quasars varying around 1% will display the appearance\/disappearance of broad emission lines. By studying the Eddington ratios of the single changing-look AGN UGC 3223, Wang et al. (2020) find this object crosses the 1% Eddington ratio when its AGN type changes. However, it is still necessary to investigate whether this 1% Eddington ratio is associated with appearance\/disappearance of BELs with multiepoch observations of CLQs, and with a larger data set. In this work, we also study the changes in the Eddington ratio distributions from the brightest to the faintest optical spectra of CLQs, investigate the possible connection with the 1% Eddington ratio, and study the BEL profiles of those CLQs.","Citation Text":["Murray et al. 1995"],"Functions Text":["The disk-wind model suggests that the BLR gas comes from winds produced by the accretion disk"],"Functions Label":["Background"],"Citation Start End":[[1041,1059]],"Functions Start End":[[924,1017]]}
{"Identifier":"2019AandA...623A.145M__Lofthouse_et_al._2017_Instance_1","Paragraph":"Low redshift (z\u2004\u2272\u20043) LyC emitters have proven difficult to find (e.g., Leitherer et al. 1995; Bergvall et al. 2006; Siana et al. 2007, 2015; Vanzella et al. 2010; Leitet et al. 2011, 2013) due to the attenuation of LyC photons by the intergalactic medium (IGM). For decades, the only confirmed, directly detected LyC escape from a low-redshift galaxy was that in Haro 11 (Bergvall et al. 2006; Leitet et al. 2011), with \u00d6stlin et al. (2015) providing the first spatially resolved kinematic study of a confirmed LyC emitter. Only recently has significant progress been made in detecting local LyC emitters. Green peas (GPs) are the most prolific class of local LyC-leaking galaxies, showing a very high LyC detection rate. So far, 11 out of 11 GPs have been confirmed as leaking LyC by direct detection with HST\/COS (Izotov et al. 2016a,b, 2018a,b). GPs are compact galaxies which appear green in SDSS g, r, i composite images due to their extremely high equivalent width of [O\u202fIII] \u03bb\u03bb4959, 5007 and their redshifts of only z\u2004\u223c\u20040.2\u2005\u2212\u20050.4. They have subsolar metallicities, low masses, and high specific star-formation rate, and represent good analogs to high redshift SFGs (e.g., Cardamone et al. 2009; Izotov et al. 2011; Bian et al. 2016). While much closer than their high-z cousins, GPs are nevertheless at distances that still present severe difficulties for a spatially resolved analysis of their properties. Pioneering work with integral field unit (IFU) observations of four GPs (Lofthouse et al. 2017) reveals that two are rotationally supported, with seemingly undisturbed morphology, and two are dispersion-dominated, leading the authors to conclude that mergers may not be a necessary driver of their star formation properties, and by extension, for their LyC emission. At the same time, Amor\u00edn et al. (2012) find complex, multicomponent H\u03b1 line profiles, with implied high-velocity gas in all five of their observed GPs. While these GPs are not confirmed LyC emitters, and therefore their kinematical properties cannot be taken as a necessary condition for LyC escape, outflows nevertheless seem to play a role in enabling LyC leakage. For example, Chisholm et al. (2017) find that outflows can be detected in all LyC emitters they examine, albeit with velocities not statistically different from a control sample of non-leakers. Further, Martin et al. (2015) find that ULIRG galaxies with the strongest outflows manifest the lowest column densities of neutral gas. Simulations also support the scenario of outflows enabling LyC escape. Trebitsch et al. (2017) find that mechanical feedback from supernovae explosions has a vital role in disrupting dense star-forming regions and in clearing low-density escape paths for the LyC photons. While IFU observations of a large sample of GPs would help settle some of these issues, the large distances to these galaxies would likely prohibit the detailed identification and characterization of the exact escape mechanisms of LyC. Much closer galaxies, showing the same characteristics as GPs, are required for such an endeavor.","Citation Text":["Lofthouse et al. 2017"],"Functions Text":["Pioneering work with integral field unit (IFU) observations of four GPs","reveals that two are rotationally supported, with seemingly undisturbed morphology, and two are dispersion-dominated, leading the authors to conclude that mergers may not be a necessary driver of their star formation properties, and by extension, for their LyC emission."],"Functions Label":["Background","Background"],"Citation Start End":[[1487,1508]],"Functions Start End":[[1414,1485],[1510,1780]]}
{"Identifier":"2015AandA...584A...4D__Schneider_et_al._2015_Instance_1","Paragraph":"The density structure and infall kinematics discussed here are probably dominated by the most massive object clearly identified in Fig. 2b. The envelope profile defined above (see Sect. 3.5) and their corresponding values given in Table 3 have been used in Eq. (A.2) to estimate a mass of ~3500 M\u2299 in a 2.5 pc radius for the central UCH\u2009ii region. This region exhibits a steeper density gradient in its outer envelope, with qout \u2243 \u22122.5, also illustrated by the cloud structure studied with probability density function (Schneider et al. 2015; Rayner et al., in prep.). It suggests compression by external forces, as shown in numerical simulations of Hennebelle et al. (2003) and outlined in Appendix A. This compressive process could be associated with stellar winds and ionisation shocks of nearby stars such as those observed in the M\u200916 massive YSO (young stellar object) by Tremblin et al. (2014b) or with globally infalling gas driven by the dynamic formation of Mon R2 cloud through force fall. Such infalling motions have in fact been observed in CO by Loren (1977), and clumps with similar masses generally exhibit active global infall. This is the case of SDC335, which has a mass of ~5500 M\u2299 in 2.5 pc (Peretto et al. 2013) and contains a protostellar object of similar type B1 (Avison et al. 2015), as well as of DR21, Clump-14 (DR21-south, ~4900 M\u2299), and Clump-16 (DR21-north, ~3350 M\u2299) (Schneider et al. 2010a), all of which show supersonic infall (V = 0.5\u22120.7 km\u2009s-1). Moreover, the hourglass morphology of the magnetic field is thought to be a signature of global infall (Carpenter & Hodapp 2008), and Koch et al. (2014) suggest that Mon R2 is a super critical, quickly collapsing cloud. The expected infall velocity gradient, as observed in SDC13 (Peretto et al. 2014), seems to be a crucial ingredient for generating the filament crossing (Dobashi et al. 2014) characteristic of all the regions mentioned here and commonly observed in many places. ","Citation Text":["Schneider et al. 2015"],"Functions Text":["This region exhibits a steeper density gradient in its outer envelope, with qout \u2243 \u22122.5, also illustrated by the cloud structure studied with probability density function"],"Functions Label":["Similarities"],"Citation Start End":[[520,541]],"Functions Start End":[[348,518]]}
{"Identifier":"2017MNRAS.472.1052B__Bruno_&_Telloni_2015_Instance_1","Paragraph":"Typical corotating high-speed streams (HSS), i.e. those streams coming from the equatorial extension of polar coronal holes, are characterized by different regions. Within these regions, extending on daily scales, field and plasma parameters assume different average values and fluctuations have a different character. The first region is the stream interface (SI), located between fast and slow stream, right where the dynamical interaction between the two flows is stronger (Schwenn & Marsch 1990; Bruno & Carbone 2016). The SI is characterized by high level of pressure, both kinetic and magnetic, high temperature and low Alfv\u00e9nicity. This region is followed by the trailing edge (TE), which is characterized by the highest wind speed and temperature, the lowest compressibility and the highest level of Alfv\u00e9nicity. The TE is then followed by a wind that progressively slows down and becomes cooler. This region was first named rarefaction region by Hundhausen (1972). Interesting enough, these two portions of the stream are separated by a very narrow region across which Alfv\u00e9nicity changes abruptly, from high to low (Bruno & Telloni 2015). Most of the time, the rarefaction region is followed by the heliospheric current sheet crossing characterized by low Alfv\u00e9nicity, low temperature but higher field and plasma compressibility. Thus, moving from fast to slow wind, one observes a different turbulence, not only relatively to its power level, but also to its Alfv\u00e9nic character, although the spectral slope at fluid scales clearly shows a persistent Kolmogorov-like scaling. In addition, besides the above-cited differences within the fluid regime, there are also clear differences at proton kinetic scales. Bruno, Trenchi & Telloni (2014), investigating the behaviour of the spectral slope at proton scales, up to frequencies of a few Hz, beyond the high-frequency break separating fluid from kinetic scales, recorded a remarkable variability of the spectral index (Smith et al. 2006; Sahraoui et al. 2010) within the following frequency decade or so. The steepest spectra corresponded to the TEs of fast streams, while the flattest ones were found within the subsequent slow wind regions. The same authors found an empirical relationship between the power associated with the inertial range and the spectral slope observed within this narrow region, which allowed us to estimate that this slope approaches the Kolmogorov scaling within the slowest wind and reaches a limiting value of roughly 4.4 within the fast wind. In addition, Bruno et al. (2014) suggested the possible role played also by Alfv\u00e9nicity in this spectral dependence. Recent theoretical results seem to move in this direction attributing a relevant role to Alfv\u00e9nic imbalance within the inertial range (Voitenko & De Keyser 2016).","Citation Text":["Bruno & Telloni 2015"],"Functions Text":["Interesting enough, these two portions of the stream are separated by a very narrow region across which Alfv\u00e9nicity changes abruptly, from high to low"],"Functions Label":["Background"],"Citation Start End":[[1126,1146]],"Functions Start End":[[974,1124]]}
{"Identifier":"2018ApJ...869...69M__K\u00fcnzel_1960_Instance_1","Paragraph":"The AR NOAA 12673 was highly flare productive.5\n\n5\n\nhttps:\/\/www.swpc.noaa.gov\/products\/solar-region-summary\n\n It appeared in the eastern limb of the Sun on August 28 as a simple \u03b1-type AR and gradually evolved into complex \u03b2\u03b3-type on September 4. It became an even more complex \u03b2\u03b3\u03b4-type on September 5 and remained so until its disappearance over the western limb of the Sun on September 10. In total, it produced 27 M-class and 4 X-class flares between September 4\u201310. Following the occurrence of the two X-class flares on September 6, reported in this article, it went on to produce an X1.3-class flare on September 7 and an X8.2-class flare on September 10 besides several M-class major eruptive events. During the occurrence of X-class flares, the complex AR had shown \u03b4-sunspots, which are identified with a complex distribution of sunspot groups in which the umbrae of positive and negative polarities share a common penumbra (K\u00fcnzel 1960). Such complex ARs are known to produce powerful flares (see e.g., Zirin & Liggett 1987; Sammis et al. 2000; Takizawa & Kitai 2015). It is noteworthy that AR 12673 was a rather compact region (Figure 9) that displayed more spatial extension in the north\u2013south direction than the usual east\u2013west direction. The \u03b4-sunspots were concentrated at the central part of the AR where magnetic fields were very strong, and the magnetic field gradient across the PIL was extremely high (\u223c2.4 \u00d7 103 G Mm\u22121; see Figure 9(c)). Earlier studies have shown a close relationship between major flare activities and strong magnetic field, especially those with a high gradient and those that are highly sheared across the PIL (Hagyard et al. 1984; Zirin & Wang 1993; Schrivjer 2007; Barnes et al. 2016). The reported eruptive activities in AR 12673, thus, represent the capability of the AR in the rapid generation and storage of huge amount of excess magnetic energy in the corona. In this context, the evolution of normalized free magnetic energy during the X-class flares is noteworthy (Figure 12). We found that the free magnetic energy stored in the AR before the flaring activity was \u223c82% of the potential magnetic energy. After the two X-class flares, it reduced to \u223c70%. Our analysis, therefore, implies that a large amount of free magnetic energy was already stored in the AR before the flaring activities and that the large X-class flares essentially released only a small fraction of it.","Citation Text":["K\u00fcnzel 1960"],"Functions Text":["During the occurrence of X-class flares, the complex AR had shown \u03b4-sunspots, which are identified with a complex distribution of sunspot groups in which the umbrae of positive and negative polarities share a common penumbra"],"Functions Label":["Background"],"Citation Start End":[[933,944]],"Functions Start End":[[707,931]]}
{"Identifier":"2020MNRAS.493..559B__Kulow_et_al._2014_Instance_1","Paragraph":"Nearly half of the known exoplanets orbit within 0.1\u2009au from their star. At such close distances, the nature and evolution of these planets is shaped by interactions with their host star (irradiation, tidal effects, and magnetic fields). In particular, the deposition of stellar X-ray and extreme ultraviolet radiation (XUV) into an exoplanet upper atmosphere can lead to its hydrodynamic expansion and substantial escape (e.g. Lammer et al. 2003; Vidal-Madjar et al. 2003; Lecavelier des Etangs et al. 2004; Yelle 2004; Garc\u00eda Mu\u00f1oz 2007; Koskinen et al. 2010; Johnstone et al. 2015). Atmospheric loss is considered as one of the main processes behind the deficit of Neptune-mass planets at close orbital distances (the so-called hot Neptune desert, e.g. Lecavelier des Etangs 2007; Davis & Wheatley 2009; Szab\u00f3 & Kiss 2011; Lopez, Fortney & Miller 2012; Beaug\u00e9 & Nesvorn\u00fd 2013; Lopez & Fortney 2013; Owen & Wu 2013; Jin et al. 2014; Kurokawa & Nakamoto 2014; Lundkvist et al. 2016). These planets are large enough to capture much of the stellar energy, but in contrast to hot Jupiters are not massive enough to retain their escaping atmospheres (e.g. Hubbard et al. 2007; Lecavelier des Etangs 2007; Ehrenreich, Lecavelier des Etangs & Delfosse 2011). The missing hot Neptunes could have lost their entire atmosphere via evaporation, evolving into bare rocky cores at the lower radius side of the desert (e.g. Lecavelier des Etangs et al. 2004; Owen & Jackson 2012). This scenario is strengthened by the recent observations of warm Neptunes at the border of the desert, on the verge of (Kulow et al. 2014; Bourrier, Ehrenreich & Lecavelier des Etangs 2015; Ehrenreich et al. 2015; Bourrier et al. 2016; Lavie et al. 2017) or undergoing (Bourrier et al. 2018b) considerable mass-loss. Because they survive more extreme conditions than lower mass gaseous exoplanets, hot Jupiters are particularly interesting targets to study star\u2013planet interactions. Their upper atmosphere can be substantially ionized because of stellar photoionization (e.g. Schneiter et al. 2016), which could help the formation of reconnections between the stellar and planetary magnetospheres that would enhance stellar activity (e.g. Cuntz, Saar & Musielak 2000; Shkolnik, Walker & Bohlender 2003; Ip, Kopp & Hu 2004; Shkolnik et al. 2008; although see Poppenhaeger & Schmitt 2011; Scandariato et al. 2013; Llama & Shkolnik 2015 for the difficulties to detect such signatures). Atmospheric escape of neutral hydrogen and metal species has been detected via transmission spectroscopy for several Jupiter-mass planets, bringing information about their upper atmosphere and the stellar environment (HD\u2009209458b, Vidal-Madjar et al. 2003, 2004, 2008; Ehrenreich et al. 2008; Ben-Jaffel & Sona Hosseini 2010; Linsky et al. 2010; Schlawin et al. 2010; Vidal-Madjar et al. 2013; Ballester & Ben-Jaffel 2015; HD\u2009189733b, Lecavelier des Etangs et al. 2010, 2012; Bourrier et al. 2013; 55\u2009Cnc\u2009b, Ehrenreich et al. 2012; WASP-12b, Fossati et al. 2010; Haswell et al. 2012). Shocks could for example form ahead of hot Jupiters because of the interaction between the stellar wind and the planetary outflow or magnetosphere (Vidotto, Jardine & Helling 2010; Cohen et al. 2011; Llama et al. 2013; Tremblin & Chiang 2013; Matsakos, Uribe & K\u00f6nigl 2015)","Citation Text":["Kulow et al. 2014"],"Functions Text":["This scenario is strengthened by the recent observations of warm Neptunes at the border of the desert, on the verge of","considerable mass-loss."],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[1589,1606]],"Functions Start End":[[1469,1587],[1762,1785]]}
{"Identifier":"2017AandA...607A.107W__Micha\u0142owski_et_al._2015_Instance_1","Paragraph":"A likely scenario for the geometry of the entire system is presented in Fig. 15. The locations of systems A\u2013C are fairly well constrained by their metal absorption. System D is harder to place, as there is little to be learned from its negligibly low metal content. It is highly improbable that it is an outflow similar to C but with a higher velocity, or indeed that it has originated inside the galaxy at all. Instead we believe it to trace the metal-poor gas that acts as a replenishment mechanism of galaxies throughout the universe. It has indeed been suggested that GRB hosts in general are fuelled by recent metal-poor gas inflow (Micha\u0142owski et al. 2015, 2016). The \u2212320 km s-1 blueshift implies that it is not accreting onto galaxy B \u2013 if B hosted the GRB, then D must be in front of it, moving away at that velocity. The degeneracy between local velocities and cosmological redshift allows for a range of locations of system D. One possibility is that it is located between the two galaxies and is currently falling towards galaxy A, or has fallen from behind galaxy B and is being slowed by its gravity. Another possibility is that it has fallen from a large distance into the potential well of the complex and is now in the foreground, having initially passed both galaxies without being accreted, having joined in a disk-like structure (e.g. Bouch\u00e9 et al. 2013). For any of these scenarios, system D represents a rare detection of such a metal-poor inflow in a GRB galaxy, or indeed in absorption at such a small projected distance to any galaxy. Systems with positive velocities relative to the gas local to GRBs, and thus inferred to be falling in towards the star-forming regions of their hosts, have been detected previously (e.g. Prochaska et al. 2008), but these have tended to be richer in metals than the gas observed here and may represent recycled gas that has previously been expelled from the galaxy. Ly-\u03b1 emitters are thought to trace accretion in large halos around quasars (Cantalupo et al. 2014; Hennawi et al. 2015), and have been detected at impact parameters of \u226570 kpc from an LLS (Fumagalli et al. 2016). If system D is indeed bound to the host system, then it could be much closer to the galaxies than that. Finally, as discussed when defining the host complex in Sect. 3.1, it could in fact be a foreground system located ~1.5 Mpc from the interacting complex. While this would exclude it from being a direct accretion flow onto the complex, it represents at the very least a reservoir of cool metal-poor gas residing in the IGM near the dark matter halo, available for star formation at some future epoch. ","Citation Text":["Micha\u0142owski et al. 2015"],"Functions Text":["Instead we believe it to trace the metal-poor gas that acts as a replenishment mechanism of galaxies throughout the universe. It has indeed been suggested that GRB hosts in general are fuelled by recent metal-poor gas inflow"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[638,661]],"Functions Start End":[[412,636]]}
{"Identifier":"2021MNRAS.506.3511M__Foucart_et_al._2012_Instance_1","Paragraph":"In this work, we study the merger and post-merger evolution of near equal-mass BH\u2013NS binaries. Before turning to the properties of the accretion discs formed in such mergers, we first provide a very brief overview of their formation. We do this by considering the fiducial system TNT.chit.0.35. In order to illustrate the disc formation process, we begin by summarizing the dynamical formation of the disc in Fig. 1, which reports the comoving magnetic energy density b2, the rest-mass density \u03c1, the electron fraction Ye, and the local fluid temperature T. The different rows correspond to meridional (top panels) and equatorial (bottom panels) views of the accretion disc around the BH, while the different columns correspond to different times after the merger. The general dynamics of this process have been studied extensively in purely hydrodynamical simulations (Etienne et al. 2009; Kyutoku et al. 2011; Foucart et al. 2012). In order for a massive disc to form during and after merger, tidal disruption has to occur outside of the innermost stable circular orbit (ISCO) of the BH (Shibata & Taniguchi 2011; Pannarale et al. 2011b). Starting from the left-hand panel, we can see that shortly after tidal disruption, an initial accretion disc begins to form around the BH. Originating from the cold NS matter, the initial disc is very neutron rich (Ye  0.05), but already reaches temperatures $T \\lesssim 10\\, \\rm MeV$. The disc quickly grows in mass and size due to fall-back accretion from the tidal arm (middle column), begins to circularize and a steady accretion flow develops over time. As expected, this happens on the dynamical time-scales of the discs, which are proportional to the disc mass $M_{\\rm disk}^b$, so that the lightest discs circularize first. Initially, the pure neutron matter is far out of beta-equilibrium under these conditions and will rapidly re-equilibrate via beta decay of neutrons, leading to an increasing protonization especially of the low-density parts of the disc. At the same time, the magnetic-field strength is increasing throughout the disc, exceeding $10^{14}\\, \\rm G$ locally. More details on the magnetic-field evolution will be given in Section 3.3. Finally, after more than $50\\, \\rm ms$ past merger, the disc has settled into an initial quasi-equilibrium, consisting of a very neutron-rich disc, probing rest-mass densities ${\\lesssim}10^{11}\\, \\rm g\\, cm^{-3}$. A disc formed by this process will then set the initial conditions for the long-term evolution in terms of the accretion flow and mass ejection (Fern\u00e1ndez et al. 2015, 2017).","Citation Text":["Foucart et al. 2012"],"Functions Text":["The general dynamics of this process have been studied extensively in purely hydrodynamical simulations"],"Functions Label":["Background"],"Citation Start End":[[912,931]],"Functions Start End":[[765,868]]}
{"Identifier":"2015MNRAS.454.1644L__Kotze_&_Charles_2012_Instance_3","Paragraph":"The period candidates of other three ULXs may range from \u223c100 to \u223c600 d. Apart from noise and artefacts, all the candidate periods are only significant in a specific epoch. This suggests that they are not associated with any stable mechanism such as orbital motion. Instead, such long-term (> 100 d) X-ray quasi-periodic variations are likely related to superorbital periods that are thought to be due to radiation-driven warping of accretion discs (Ogilvie & Dubus 2001) or tidal interaction-induced disc precession (Whitehurst & King 1991). Alternatively, mass transfer rate-related events such as X-ray state changes and disc instability can also cause long-term modulations (Kotze & Charles 2012). In particular, there are two intermittent quasi-periodicities for both NGC 5408 X-1 and M81 X-6, suggesting that the quasi-periods are changing or evolving. They resemble some Galactic X-ray binaries that show similar behaviour (e.g. Cyg X\u20132 and SMC X\u20131; Kotze & Charles 2012) and it has been suggested that a warped disc could lead to an unstable steadily precessing disc, causing quasi-periodic behaviour (Ogilvie & Dubus 2001). We note that there are many uncertainties on the physical parameters of ULXs. To determine the origin of superorbital periods of ULXs, one has to know at least the mass ratio between the companion and the compact star (q = MC\/MX) and the binary separation. Unfortunately, it is very difficult to get these parameters for ULXs. For the three ULXs discussed here (i.e. excluding ESO 243-49 HLX-1), only M81 X-6 has better constraints on the black hole mass and the nature of the companion. The masses of the black hole and companion star are estimated (MX = 18\u2009M\u2299, MC = 23\u2009M\u2299) such that q can be derived. In this case, we can rule out a tidal interaction-induced disc precession scenario that requires q  0.25\u20130.33 (Whitehurst & King 1991). For a warped disc, the binary separation and the mass ratio suggest that M81 X-6 lies in the intermediate instability zone for radiation-driven warping in X-ray binaries (see fig. 1 of Kotze & Charles 2012). The quasi-periodic variability may represent the switching time-scale between a warped disc and a flat disc.","Citation Text":["Kotze & Charles 2012"],"Functions Text":["For a warped disc, the binary separation and the mass ratio suggest that M81 X-6 lies in the intermediate instability zone for radiation-driven warping in X-ray binaries (see fig. 1 of"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[2057,2077]],"Functions Start End":[[1872,2056]]}
{"Identifier":"2022AandA...666A.112L__Cormier_et_al._2015_Instance_1","Paragraph":"Local dwarf galaxies were the focus of large Herschel and Spitzer surveys (e.g., The Dwarf Galaxy Survey, DGS; Madden et al. 2006). Studies on both resolved and integrated-galaxy scales have highlighted some distinctively unique observational signatures of star-forming low-metallicity dwarf galaxies. A non-linear relation of the dust-to-gas mass (D\/G) with metallicity is observed, with extremely low dust masses observed for the lowest metallicity galaxies (Z \u2264 0.1 Z\u2299; Herrera-Camus et al. 2012; Fisher et al. 2014; R\u00e9my-Ruyer et al. 2015; Galliano et al. 2018, 2021; Cigan et al. 2021). Furthermore, the hard radiation fields in star-forming dwarf galaxies, along with their lower dust abundance, result in extended ionized gas regions prominent on global galaxy scales (Hunter et al. 2011; Cormier et al. 2015, 2019). The consequence is often a largely photodissociated molecular phase, existing in clumps which are difficult to observe with the usual molecular gas tracer, CO (1-0) (Cormier et al. 2014; Hunt et al. 2015; Accurso et al. 2017b), beckoning the presence of a CO-dark molecular gas phase (Grenier et al. 2005; R\u00f6llig et al. 2006; Wolfire et al. 2010; Glover & Clark 2012; Bolatto et al. 2013; Accurso et al. 2017a; Madden et al. 2020). Other emission lines, however, such as the far-infrared [C ii]\u03bbl58 \u00b5m line, are strikingly enhanced on global scales in dwarf galaxies (e.g., Cormier et al. 2015, 2019; Cigan et al. 2016; Lebouteiller et al. 2017; Jameson et al. 2018), making the [C ii]\u03bbl58 \u00b5m line a potential tool for tracing star formation activity (Malhotra et al. 2001; Papadopoulos et al. 2007; Pineda et al. 2014; De Looze et al. 2014; Olsen et al. 2015; Herrera-Camus et al. 2015, Herrera-Camus et al. 2018; Carniani et al. 2018; Matthee et al. 2019; Izumi et al. 2021; Fujimoto et al. 2021) and a tracer of the total H2 in galaxies, near and far (Poglitsch et al. 1995; Wolfire et al. 2010; Pineda et al. 2013; Nordon & Sternberg 2016; Fahrion et al. 2017; \nAccurso et al. 2017b; Zanella et al. 2018; Madden et al. 2020; Schaerer et al. 2020; Tacconi et al. 2020).","Citation Text":["Cormier et al. 2015"],"Functions Text":["Furthermore, the hard radiation fields in star-forming dwarf galaxies, along with their lower dust abundance, result in extended ionized gas regions prominent on global galaxy scales"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[796,815]],"Functions Start End":[[592,774]]}
{"Identifier":"2022AandA...663A.105P__Bonafede_et_al._2012_Instance_2","Paragraph":"Cluster radio relics are usually found in the outskirts of merging galaxy clusters. They exhibit elongated morphologies and high degrees of polarisation above 1 GHz (up to 70%, Ensslin et al. 1998; Bonafede et al. 2014; Loi et al. 2019; de Gasperin et al. 2022). The resolved spectral index in radio relics shows a gradient: it steepens towards the cluster centre and flattens towards the outskirts. Their size can reach up to \u223c2 Mpc, and high-resolution observations have revealed filamentary structures within relics themselves (Di Gennaro et al. 2018; Rajpurohit et al. 2020, 2022a,b; de Gasperin et al. 2022). The Largest Linear Sizes (LLS) and radio powers of relics are correlated, as well as the integrated spectral index and the radio power (van Weeren et al. 2009b; Bonafede et al. 2012; de Gasperin et al. 2014). Relics trace ICM shock waves with relatively low (M\u2004 \u20043) Mach numbers (Finoguenov et al. 2010; Akamatsu et al. 2013; Shimwell et al. 2015; Botteon et al. 2016). The acceleration of electrons is believed to proceed via diffusive shock acceleration (DSA) in the ICM (Ensslin et al. 1998; Roettiger et al. 1999), in which particles scatter back and forth across the shock front gaining energy at every crossing. Nevertheless, this mechanism has been shown to be rather inefficient in accelerating electrons from the thermal pool (Vazza & Br\u00fcggen 2014; Vazza et al. 2016; Botteon et al. 2020a; Br\u00fcggen & Vazza 2020; see Brunetti & Jones 2014 for a review). Recently, it has been suggested that seed electrons could originate from the tails and lobes (driven by AGN outflows) of cluster radio galaxies (Bonafede et al. 2014; van Weeren et al. 2017; Stuardi et al. 2019), which alleviates the requirements of high acceleration efficiencies at cluster shocks (e.g., Markevitch et al. 2005; Kang et al. 2012, 2017; Botteon et al. 2016; Eckert et al. 2016). In some cases, double relics have been detected on opposite sides of the cluster centre (e.g., Rottgering et al. 1997; van Weeren et al. 2010, 2012b; Bonafede et al. 2012; de Gasperin et al. 2015a). In these clusters it is possible to constrain the merger history, providing important information about the formation processes of relics.","Citation Text":["Bonafede et al. 2012"],"Functions Text":["In some cases, double relics have been detected on opposite sides of the cluster centre (e.g.,","In these clusters it is possible to constrain the merger history, providing important information about the formation processes of relics."],"Functions Label":["Background","Background"],"Citation Start End":[[2022,2042]],"Functions Start End":[[1872,1966],[2071,2209]]}
{"Identifier":"2021MNRAS.501.3781R__Nisini_et_al._2005_Instance_2","Paragraph":"While spatially extended optical jets and bipolar CO molecular outflows have been observed in numerous Class 0\/I protostars (e.g. Reipurth & Bally 2001; Bally 2016, and references therein), near-infrared high-resolution spectroscopy and spectroimaging observations in the past two decades have made it possible to study the kinematics of the outflowing gas and physical properties at the base of the jet within a few hundred au of the driving source in Class 0\/I protostars (e.g. Davis et al. 2001, 2003, 2011; Nisini et al. 2005, 2016; Caratti o Garatti et al. 2006; Takami et al. 2006; Antoniucci et al. 2008, 2011, 2017; Garcia Lopez et al. 2008, 2013). These microjets are bright in [Fe\u2009ii] forbidden and H2 rovibrational emission lines, hence showing the presence of forbidden emission-line (FEL) regions and molecular hydrogen emission-line (MHEL) regions in low-mass Class 0\/I protostars. While multiple low- and high-velocity components are observed in both MHELs and FELs, the higher velocity gas is slightly further offset from the driving source than the slower gas, and the kinematics of the H2 emission differs from [Fe\u2009ii] emission, revealing complicated kinematic structures. Evidence of H2 emission from cavity walls is also seen in some protostars, suggesting the presence of a wide-angled wind. Strong emission in the well-known accretion diagnostics of Paschen and Brackett hydrogen recombination lines is observed in protostars, with the ratio of the accretion luminosity to bolometric luminosity spanning from \u223c0.1 to \u223c1. The mass accretion and loss rates for Class 0\/I low-mass protostars span the range of 10\u22126\u201310\u22128 M\u2299 yr\u22121, and the derived jet efficiencies (ratio between mass ejection and accretion rates) range between \u223c1 per\u2009cent and 10 per\u2009cent (e.g. Davis et al. 2001, 2003, 2011; Nisini et al. 2005, 2016; Caratti o Garatti et al. 2006; Takami et al. 2006; Antoniucci et al. 2008, 2011, 2017; Garcia Lopez et al. 2008, 2013). These measurements are within the range predicted by the magnetohydrodynamic jet launching models (e.g. Frank et al. 2014).","Citation Text":["Nisini et al. 2005"],"Functions Text":["The mass accretion and loss rates for Class 0\/I low-mass protostars span the range of 10\u22126\u201310\u22128 M\u2299 yr\u22121, and the derived jet efficiencies (ratio between mass ejection and accretion rates) range between \u223c1 per\u2009cent and 10 per\u2009cent (e.g."],"Functions Label":["Background"],"Citation Start End":[[1810,1828]],"Functions Start End":[[1543,1778]]}
{"Identifier":"2022ApJ...932...98S__hand,_FRB_2018_Instance_1","Paragraph":"The band-limited nature of repeaters makes it essential to have high-resolution simultaneous observations with a large bandwidth that can help detect variations in the spectrotemporal and even polarimetric structures of the bursts at different frequencies. The study of these contemporaneous frequency-dependent variations, as well as long-term changes with time, are paramount in investigating the origins of FRBs. FRB 20121102A was one of the first repeaters to be coherently corrected for dispersion, allowing unprecedented high-time-resolution studies, revealing drifting bursts and an unusually high rotation measure (RM; Gajjar et al. 2018; Michilli et al. 2018; Hessels et al. 2019). Subsequent observations have shown a considerable decrease in the RM value with time (Hilmarsson et al. 2021). This variation is similar to what has been observed for Galactic center magnetar PSR J1745-2900 (Desvignes et al. 2018). On the other hand, FRB 20180916B has a nominal RM and has shown a gradual variation in RM value over the two year timescale (Pleunis et al. 2021). These variations may be due to changes in the viewing geometry or the result of drastic variations in the circumburst environment facilitating magnetoionic disparities. Interestingly, both of these FRBs show a flat polarization position angle (PPA) that can be indicative of emissions at higher altitudes due to relativistic shock (Beloborodov 2017; Metzger et al. 2019). Contrarily, observation of microstructures challenges this idea, as in the case of FRB 20180916B, where such structures have been seen down to 2\u20133 \u03bcs (Nimmo et al. 2020) at 1700 MHz, hinting at emission much closer to the magnetosphere. Studying the evolution of such structures with frequency will aid in constraining the emission region and help to understand the local environment. Another aspect of repeater burst morphology that remains a mystery is the \u201csad trombone\u201d effect, which describes a linear decrease in burst features with frequency and has been observed for FRB 20121102A and other repeaters; it will be interesting to undertake a comparative study of this behavior between frequencies.","Citation Text":["Pleunis et al. 2021"],"Functions Text":["On the other hand, FRB 20180916B has a nominal RM and has shown a gradual variation in RM value over the two year timescale"],"Functions Label":["Differences"],"Citation Start End":[[1048,1067]],"Functions Start End":[[923,1046]]}
{"Identifier":"2021ApJ...908...40M__Muschietti_&_Lemb\u00e8ge_2017_Instance_1","Paragraph":"These signatures are inconsistent with ultra-low frequency waves, which have circular polarization and a period similar to the upstream ion gyroperiod. The waves are also inconsistent with ion Weibel instability, which generates linearly polarized waves. Interaction of reflected ions with incoming solar wind electrons or ions can cause foot instabilities that excite waves in the whistler mode branch. Modified Two Stream Instability (MTSI) due to relative drift between reflected ions and incoming solar wind electrons (fast drift), and incoming solar wind ions and electrons (slow drift) has been frequently considered (Matsukiyo & Scholer 2003; Comi\u015fel et al. 2011; Umeda et al. 2012; Marcowith et al. 2016; Wilson et al. 2016; Muschietti & Lemb\u00e8ge 2017; Hull et al. 2020). This instability, however, if excited, creates significant ion heating throughout the foot and suppresses the reformation process (Shimada & Hoshino 2005; Matsukiyo & Scholer 2006), rather than creating episodic enhancements that we show in the foot. Furthermore, Gary et al. (1987) indicated that (fast drift) MTSI becomes dominant at low electron beta (\u03b2e  0.5), while at higher \u03b2e more resonant electrons stabilize this instability through increased electron Landau damping. Electron data for the time period we discussed in this paper show \u03b2e \u2265 1.2, and therefore fast drift mode MTSI is most likely not significant. The slow drift mode of MTSI could be a more viable candidate at high \u03b2 plasmas. Wave properties around 1.6 Hz in the middle interval (purple segment) of Figure 6, indicate that the wave is in propagation toward the ramp (\n\n\n\n\n\n = \u22120.66, \u22120.71, 0.22) with Vph\u2010sc = 34 km s\u22121 and \u03bbwave = 21.4 km  \u223c 12\u03bbe, where \u03bbe is the upstream electron inertial length. The plasma rest-frame frequency of the wave is about 8 Hz  \u223c 2.5flh. Since these characteristics are somewhat consistent with model predictions for drift mode of MTSI (Muschietti & Lemb\u00e8ge 2017), we do not rule out the possibility of some waves at certain frequencies and during some intervals being generated by the slow drift mode of MTSI.","Citation Text":["Muschietti & Lemb\u00e8ge 2017"],"Functions Text":["Modified Two Stream Instability (MTSI) due to relative drift between reflected ions and incoming solar wind electrons (fast drift), and incoming solar wind ions and electrons (slow drift) has been frequently considered"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[733,758]],"Functions Start End":[[404,622]]}
{"Identifier":"2016ApJ...821..107G__Gloeckler_&_Fisk_2015_Instance_4","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"],"Functions Text":["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","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."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[3133,3154]],"Functions Start End":[[2940,3131],[3157,3358]]}
{"Identifier":"2016ApJ...824..138Y__yanin_2011_Instance_1","Paragraph":"The above qualitative theoretical reasoning raises the question about why would Swift J1834.9\u22120846 be the only magnetar so far powering a wind nebula, given that previous searches around individual magnetars have returned no sign of extended emission attributable to wind nebulae (e.g., Vigan\u00f2 et al. 2014). With only one observed so far, it is difficult to draw any firm conclusions. Nevertheless, Swift J1834.9\u22120846 has some interesting characteristics that are not shared with the entire magnetar population. First, the environment of Swift J1834.9\u22120846 is extremely crowded, with a Fermi GeV source, an H.E.S.S. TeV source, an SNR, a GMC, and an OH maser in its vicinity (Frail et al. 2013; H.E.S.S. Collaboration et al. 2015). The relationship between all these sources is unclear. However, it is tempting to speculate that environmental effects from such a rich field could be playing a role in the production of this wind nebula (e.g., triggering of pair cascade by external gamma rays from a nearby source; Shukre & Radhakrishnan 1982; Istomin and Sob\u2019yanin 2011). Second, the Swift J1834.9\u22120846 X-ray luminosity in quiescence is \n\n\n\n\n\n erg s\u22121. Only five other magnetars (SGR 0418+5729, SGR 1745\u22122900, XTE J1810\u2212197, Swift J1822.3\u22121606, 3XMM J185246.6+003317)20\n\n20\n\nhttp:\/\/www.physics.mcgill.ca\/ pulsar\/magnetar\/main.html\n\n have luminosities \u22721032 erg s\u22121. Among these five, three have the smallest surface B fields measured (SGR 0418+5729, Swift J1822.3\u22121606, 3XMM J185246.6+003317; \n\n\n\n\n\n G), and only one source, SGR 1745\u22122900, has a rotational energy loss rate \n\n\n\n\n\n similar to Swift J1834.9\u22120846, while the rest have \n\n\n\n\n\n at least an order of magnitude lower. Hence, from an observational point of view, it seems that the combination of very weak X-ray luminosity, a magnetar-like B-field strength, and a somewhat large \n\n\n\n\n\n (properties that are only shared by the Galactic center magnetar SGR 1745\u22122900) may favor wind nebula production. Another possibility is that the Swift J1834.9\u22120846 magnetar\/nebula system is an older analog to the Kes 75 system, where the central pulsar evolves into a magnetar while preserving its originial PWN.","Citation Text":["Istomin and Sob\u2019yanin 2011"],"Functions Text":["However, it is tempting to speculate that environmental effects from such a rich field could be playing a role in the production of this wind nebula (e.g., triggering of pair cascade by external gamma rays from a nearby source"],"Functions Label":["Motivation"],"Citation Start End":[[1044,1070]],"Functions Start End":[[787,1013]]}
{"Identifier":"2016MNRAS.457.3084M__Storm_et_al._2005_Instance_1","Paragraph":"The populous blue LMC cluster NGC 1866 is already known to host an exceptionally rich sample of more than 20 Cepheids (Musella et al. 2006; Welch & Stetson 1993). One of these was identified by Musella et al. (2006) in a preliminary analysis of the proprietary BVI Very Large Telescope (VLT) data. It is unquestionable that such a unique sample of Cepheids \u2013 likely all members of the cluster and at the same distance, chemical composition and age \u2013 would constitute a milestone in our understanding of the Cepheid pulsational scenario. Indeed, it offers an unprecedented opportunity to investigate both empirical and theoretical estimates of the luminosity and colour of the pulsating structures and their relation with the observed periods. For this reason, many authors have studied the NGC 1866 Cepheids (Welch 1991; Welch & Stetson 1993; Gieren, Richtler & Hilker 1994; Walker 1995; Gieren et al. 2000; Storm et al. 2005; Testa et al. 2007) in both the optical and near-infrared bands, and have tested different methods to calibrate the PL relations in different filters. In Brocato et al. (2004), we already discussed the sample of the 23 known Cepheids in NGC 1866, concluding that unfortunately only 4\u20136 Cepheids had light curves accurate enough to allow a meaningful determination of their luminosities and colours. On the basis of such a tantalizing situation, we took advantage of assigned observing time at the ESO VLT to perform an accurate photometric investigation of the cluster field, with the aim of securing suitable data constraining the light curves of all the member Cepheids. Moreover, to get accurate information about radial velocities and chemical abundances of the stars in NGC 1866, we have performed FLAMES@VLT (Fibre Large Array Multi Element Spectrograph mounted on VLT) spectroscopic observations for 30 stars (19 belonging to the cluster and 11 to the LMC field), including three Cepheids (Mucciarelli et al. 2011; Molinaro et al. 2012). Mucciarelli et al. (2011) found that, as far as the chemical composition is concerned, the cluster stars are reasonably homogeneous. Indeed, they appear to share the same abundances within the uncertainties, and this property is independent of the evolutionary status. The average iron abundance is [Fe\/H] = \u22120.43 \u00b1 0.01 dex, with a dispersion \u03c3 = 0.04 dex. For the three spectroscopically investigated Cepheids Molinaro et al. (2012), adopting the same procedure used in Mucciarelli et al. (2011), found values fully consistent with the average iron content. Moreover, Molinaro et al. (2012) applied the CORS Baade\u2013Wesselink (BW) method (Ripepi et al. 1997) to a sample of 11 Cepheids, using radial velocities obtained both from our FLAMES investigation and from literature data, and light curves based on a part of the UBVI data used in this paper complemented with K data by Testa et al. (2007). In this way, they obtained a direct estimate of the distance modulus of NGC 1866, \u03bc0 = 18.51 \u00b1 0.03 mag (see Molinaro et al. 2012, for details).","Citation Text":["Storm et al. 2005"],"Functions Text":["It is unquestionable that such a unique sample of Cepheids \u2013 likely all members of the cluster and at the same distance, chemical composition and age \u2013 would constitute a milestone in our understanding of the Cepheid pulsational scenario. Indeed, it offers an unprecedented opportunity to investigate both empirical and theoretical estimates of the luminosity and colour of the pulsating structures and their relation with the observed periods. For this reason, many authors have studied the NGC 1866 Cepheids","in both the optical and near-infrared bands, and have tested different methods to calibrate the PL relations in different filters."],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[908,925]],"Functions Start End":[[298,807],[946,1076]]}
{"Identifier":"2020MNRAS.494.2465B__Hornik_1991_Instance_1","Paragraph":"Here we demonstrate that, over a fixed time interval, the planar three-body problem can be solved by means of a multilayered deep artificial neural network (ANN; e.g. see LeCun, Bengio & Hinto 2015). These networks are designed for high-quality pattern recognition by mirroring the function of our brains (McCulloch & Pitts 1943; Rosenblatt 1985) and have been successfully applied to a wide variety of pattern recognition problems in science and industry, even mastering the game of Go (Silver et al. 2016). The abundance of real-world applications of ANNs is largely a consequence of two properties: (i) an ANN is capable of closely approximating any continuous function that describes the relationship between an outcome and a set of covariates, known as the universal approximation theorem (Cybenko 1989; Hornik 1991); and (ii) once trained, an ANN has a predictable and a fixed computational burden. Together, these properties lead to the result that an ANN can be trained to provide accurate and practical solutions to Newton\u2019s laws of motion, resulting in major improvements in computational economy (Lee, Sode-Yome & Park 1991) relative to modern technologies. Our proof-of-principle method shows that an ANN can accurately match the results of converged solutions found using the arbitrary precision numerical integrator that, for computationally challenging scenarios, e.g. during multiple close encounters, can offer numerical solutions at a fraction of the time cost and CO2 expense. We demonstrate the importance of training an ANN on converged solutions. This enables the trained ANN to accurately predict particle locations even when a conventional \u2018double-precision\u2019 numerical integrator fails dramatically. By training an ANN that can accurately compute particle trajectories during close encounters, our work extends previous work training neural networks on an n-body-type problem (e.g. Quito, Monterola & Saloma 2001; Battaglia et al. 2016). Our findings also add to the growing body of literature that supports machine learning technologies being developed to enrich the assessment of chaotic systems (Pathak et al. 2018; Stinis 2019) and providing alternative approaches to classical numerical solvers more broadly (Hennig, Osborne & Girolami 2015).","Citation Text":["Hornik 1991"],"Functions Text":["The abundance of real-world applications of ANNs is largely a consequence of two properties: (i) an ANN is capable of closely approximating any continuous function that describes the relationship between an outcome and a set of covariates, known as the universal approximation theorem"],"Functions Label":["Background"],"Citation Start End":[[809,820]],"Functions Start End":[[509,793]]}
{"Identifier":"2016ApJ...817...12P__Sur_et_al._2007_Instance_1","Paragraph":"Large-scale magnetic fields with strength of the order of 1\u201310 \u03bcG have been observed in disk galaxies (e.g., Beck et al. 1996; Fletcher 2010; Beck 2012; Beck & Wielebinski 2013; Van Eck et al. 2015). The origin of these fields can be explained through mean-field dynamo theory (Ruzmaikin et al. 1988; Beck et al. 1996; Brandenburg & Subramanian 2005a; Kulsrud & Zweibel 2008). The conservation of magnetic helicity is one of the key constraints in these models, and also leads to the suppression of the \u03b1-effect. The operation of the mean-field dynamo automatically leads to the growth of magnetic helicity of opposite signs between the large-scale and small-scale magnetic fields (Pouquet et al. 1976; Gruzinov & Diamond 1994; Blackman & Field 2002). To avoid catastrophic suppression of the dynamo action (\u03b1-quenching), the magnetic helicity due to the small-scale magnetic field should be removed from the system (Blackman & Field 2000, 2001; Kleeorin et al. 2000). Mechanisms suggested to produce these small-scale magnetic helicity fluxes are: advection of magnetic fields by an outflow from the disk through the galactic fountain or wind (Shukurov et al. 2006; Sur et al. 2007; Chamandy et al. 2014), magnetic helicity flux from anisotropy of the turbulence produced by differential rotation (Vishniac & Cho 2001; Subramanian & Brandenburg 2004, 2006; Sur et al. 2007; Vishniac & Shapovalov 2014), and through diffusive flux (Kleeorin et al. 2000, 2002; Brandenburg et al. 2009; Mitra et al. 2010; Chamandy et al. 2014). The outflow of magnetic helicity from the disk through dynamo operation leads to the formation of a corona (Blackman & Field 2000). According to Taylor's hypothesis, an infinitely conducting corona would resistively relax to force-free field configurations under the constraint of global magnetic helicity conservation (Woltjer 1960; Taylor 1974; Finn & Antonsen 1983; Berger & Field 1984; Mangalam & Krishan 2000). In this paper, we include advective and diffusive fluxes in a simple semi-analytic model of a galactic dynamo that transfers magnetic helicity outside the disk and consequently builds up a force-free corona in course of time. We first solve the time-dependent dynamo equations by expressing them as separable in variables r and z. The radial part of the dynamo equation is solved using an eigenvector expansion constructed using the steady-state solutions of the dynamo equation. The eigenvalues of the z part of the solution are obtained by solving a fourth-order algebraic equation, which primarily depends upon the turbulence parameters and the magnetic helicity fluxes. Once the dynamo solutions are written out as parametric functions of these parameters, the evolution of the mean magnetic field is computed numerically by simultaneously solving the dynamical equations for \u03b1-quenching and the growth of large-scale coronal magnetic helicity. Since the large-scale magnetic field lines cross the boundary between the galactic disk and the corona, the magnetic helicity of the large-scale magnetic field in the disk volume is not well defined. Hence we use the concept of gauge-invariant relative helicity (Finn & Antonsen 1983; Berger & Field 1984; Berger 1985) to estimate the large-scale magnetic helicity in the disk and the corona. Here the gauge-invariant relative helicity for the cylindrical geometry is calculated using the prescription given in Low (2006, 2011). We then investigate the dependence of the saturated mean magnetic field strength and its geometry on the magnetic helicity fluxes within the disk and the corresponding evolution of the force-free field in the corona.","Citation Text":["Sur et al. 2007"],"Functions Text":["Mechanisms suggested to produce these small-scale magnetic helicity fluxes are: advection of magnetic fields by an outflow from the disk through the galactic fountain or wind"],"Functions Label":["Background"],"Citation Start End":[[1167,1182]],"Functions Start End":[[969,1143]]}
{"Identifier":"2022AandA...667A...5R__Bakx_et_al._2020_Instance_1","Paragraph":"The characterization of the dynamics of galaxies just after (4\u2004\u2272\u2004z\u2004\u2272\u20046) and within the Epoch of Reionization (z\u2004\u2273\u20046) is still in its infancy as spatially resolved observations targeting emission lines are available only for about ten of sub-mm sources (e.g., Rizzo et al. 2020; Lelli et al. 2021; Rizzo et al. 2021) and a handful of Lyman break galaxies (e.g., Jones et al. 2017; Fujimoto et al. 2021). Instead, the number of marginally resolved observations aiming at inferring the integrated properties of early galaxies has increased in the last five years (Smit et al. 2018; Le F\u00e8vre et al. 2020; Bouwens et al. 2022). In their pioneering work, Smit et al. (2018) used marginally resolved Atacama Large Millimeter\/Submillimetre Array (ALMA, Wootten & Thompson 2009) observations of the [C\u202fII]-158 \u03bcm emission line to show that two star-forming galaxies at z\u2004\u223c\u20047 have smooth velocity gradients and interpreted them as being rotationally supported disks. Similar results are obtained in studies of individual galaxies at z\u2004\u223c\u20046\u20138 (Bakx et al. 2020; Harikane et al. 2020). However, given the low angular resolution of these data, they can not rule out the possibility that one or more merging [C\u202fII]-bright satellites are mimicking the smooth gradient observed in the velocity maps (see discussion in Smit et al. 2018; Simons et al. 2019). Instead, a gradient in the velocity fields of two Lyman break galaxies at z\u2004=\u20046.1 and 7.1, combined with the identification of two compact components, has been interpreted as evidence of mergers (Jones et al. 2017; Hashimoto et al. 2019). Recently, Le F\u00e8vre et al. (2020) and Romano et al. (2021) performed the first systematic morpho-kinematic analysis of a statistically significant sample of z\u2004\u223c\u20044\u20136 main-sequence galaxies from the ALMA Large Program to INvestigate [C\u202fII] at Early times (ALPINE) survey. After combining the morphological and kinematic analysis of the [C\u202fII] observations with the rest-frame UV and optical data, Romano et al. (2021) conclude that 23 out of the 75 ALPINE galaxies are merging systems. This large fraction of mergers in the ALPINE sample could imply a significant contribution of major mergers to the mass assembly in the early Universe (Le F\u00e8vre et al. 2020; Romano et al. 2021), confirming previous results based on the count of close-pair galaxies in photometric surveys (Mantha et al. 2018; Duncan et al. 2019). However, due to the sensitivity and angular resolutions of the ALPINE observations, the kinematic characterization and, consequently, the merger fraction may be uncertain.","Citation Text":["Bakx et al. 2020"],"Functions Text":["Similar results are obtained in studies of individual galaxies at z\u2004\u223c\u20046\u20138","However, given the low angular resolution of these data, they can not rule out the possibility that one or more merging [C\u202fII]-bright satellites are mimicking the smooth gradient observed in the velocity maps"],"Functions Label":["Similarities","Compare\/Contrast"],"Citation Start End":[[1032,1048]],"Functions Start End":[[957,1030],[1073,1281]]}
{"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"],"Functions Text":["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","indicating a weak tendency that outflows associated with massive stars are less collimated than those from low-mass stars as previously thought"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1643,1657]],"Functions Start End":[[1483,1641],[1660,1803]]}
{"Identifier":"2021AandA...645A..95H__Boese_2000_Instance_1","Paragraph":"Next we chose an optimum cut-off radius for the detector FOV. The PSPC has a circular FOV with a radius 57\u2032. The PSPC entrance window has a rib support structure with an inner ring at a radius corresponding to 20\u2032 (Pfeffermann et al. 1987; Hasinger & Zamorani 2000). Both the ROSAT telescope angular resolution and its vignetting function are roughly constant within the inner 20\u2032 ring, but degrade significantly towards larger off-axis angles. The combined detector and telescope PSFs are described in detail in Boese (2000). To the first order, the PSF at each off-axis angle can be approximated by a Gaussian function with a half power radius (HPR) of 13, 22, 52, 93, 130, and 180\u2033, at off-axis angles of 0, 12, 24, 36, 48, and 57\u2032, respectively (at 1 keV). The vignetting function at 1 keV drops almost linearly to about 50% at an off-axis angle of 50\u2032. Taking into account all these effects, the HPR of the overall RASS PSF is 84\u2033 (Boese 2000). This means that the classical confusion limit (40 beams per source) is reached at a source density of about 15 sources deg\u22122, which is exceeded in the high-exposure areas of our survey. In addition, we need to optimally discriminate between extended and point-like X-ray sources, calling for an angular resolution that is as high as possible. We therefore have to reduce the detector FOV. The sharpest imaging is achieved within the inner 20\u2032 of the PSPC FOV, corresponding to the inner ring-like rib of the PSPC support structure (see Fig. 1). However, there is a trade-off between image sharpness and the number of photons required for detection and image characterization. In particular in the outer areas of our survey, where the RASS exposure times drop significantly, a 20\u2032 FOV radius does not provide sufficient exposure time. Taking into account the various competing factors in this trade-off, we made a few tests varying the FOV cut-off radius, and finally decided on an optimum FOV radius of 30\u2032. The PSPC detector coordinates have a pixel size of 0.934\u2033. We thus removed all X-ray events from the dataset, which are further than 1925 pixels from the PSPC centre pixel coordinate [4119,3929]. A similar cut had to be applied to the modified PSPC instrument map (MOIMP), which is used later for the construction of the survey exposure map.","Citation Text":["Boese (2000)"],"Functions Text":["The combined detector and telescope PSFs are described in detail in"],"Functions Label":["Background"],"Citation Start End":[[513,525]],"Functions Start End":[[445,512]]}
{"Identifier":"2018AandA...613A..76J__Kennedy_&_Kenyon_2008_Instance_2","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"],"Functions Text":["In addition,","and the snow line is located at a greater distance from the central star","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."],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1419,1440]],"Functions Start End":[[1162,1174],[1345,1417],[1443,1684]]}
{"Identifier":"2017ApJ...850...97B__Haynes_et_al._2011_Instance_1","Paragraph":"The H i mass fraction of every gas particle in the baryonic runs is calculated based on the particle\u2019s temperature and density and the cosmic UV background radiation flux while including a prescription for self-shielding of H2 and dust shielding in both H i and H2 (Christensen et al. 2012). This allows for the straightforward calculation of the total H i mass of each simulated galaxy. We create mock H i data cubes only for the 42 halos that contain \n\n\n\n\n\n. Specifically, we create mock data cubes that mimic ALFALFA observations (Haynes et al. 2011). After specifying a viewing angle (see below), our code considers the line-of-sight velocity of each gas particle. The velocity of each particle is tracked in the simulation by solving Newton\u2019s equations of motion, but any turbulent velocity of the gas is not taken into account. Velocity dispersions in dwarf galaxies can be on the order of the rotational velocity, \u223c10\u201315 km s\u22121 (e.g., Stanimirovi\u0107 et al. 2004; Tamburro et al. 2009; Oh et al. 2015). Dispersions are thought to be driven at least partially by thermal velocities or supernovae (Tamburro et al. 2009; Stilp et al. 2013a, 2013b). In our simulations, supernovae inject thermal energy, and the thermal state of the H i gas needs to be considered in the mock H i linewidth for a realistic comparison to observations. To account for the thermal velocity, the H i mass of each gas particle is assumed to be distributed along the line-of-sight in a Gaussian distribution with a standard deviation given by the thermal velocity dispersion, \n\n\n\n\n\n, where T is the temperature of the gas particle. After this thermal broadening is calculated, a mock H i data cube can be generated by specifying the spatial and velocity resolution. For all of our galaxies, we adopt a spatial resolution of 54 pixels across 2Rvir. In practice, this corresponds to a range of \u223c1 kpc resolution in our lowest-mass galaxies up to \u223c9 kpc resolution in our most massive galaxies. However, the spatial resolution plays no role in our study, since measurements of the VF are based on spatially unresolved H i data. For the velocity resolution, we match the ALFALFA specification of 11.2 km s\u22121 (two-channel boxcar-smoothed).","Citation Text":["Haynes et al. 2011"],"Functions Text":["Specifically, we create mock data cubes that mimic ALFALFA observations"],"Functions Label":["Uses"],"Citation Start End":[[534,552]],"Functions Start End":[[461,532]]}
{"Identifier":"2021MNRAS.505.2561C__Hunt_et_al._2018_Instance_1","Paragraph":"Several mechanisms have been suggested to explain the formation of moving groups. A common explanation is that these velocity structures are the remnants of open clusters, or formed by interactions with a bar (Eggen 1965; Dehnen 2000). One problem with the cluster formation idea is that stars in moving groups can have a variety of different ages and compositions, so it is unlikely that they all came from the same cluster (Eggen 1965; Kushniruk et al. 2020). Analysis of GALAH data (Quillen et al. 2018) indicates that some moving groups, such as the Hercules moving group, may be due to a resonant bar. It has also been suggested that moving groups could have been formed from perturbations due to the Magellanic Clouds via gravitational interactions (Dehnen 1998). Recent work also finds that transient spiral structure (Hunt et al. 2018) may lead to the formation of moving groups, as well as perturbations due to spiral arms in the MW (Michtchenko et al. 2018). Moving groups in Gaia data have also been identified and analysed in action space. In the (JR, Jz) plane there are at least seven overdensities that follow lines of constant slope in this plane, which correspond to known moving groups in the solar neighbourhood (Trick, Coronado & Rix 2019). It is likely that there may be multiple causal mechanisms at play in the formation of moving groups in the Milky Way. The analysis of Gaia DR2 data has revealed many facets of a Galaxy that are clearly out of equilibrium, including the so-called phase-space spiral (Antoja et al. 2018), and the Enceladus merger (Helmi et al. 2018), that have been interpreted as arising from interactions with dwarf galaxies. Analysis of Gaia DR2 data also led to the discovery of a new dwarf galaxy (Torrealba et al. 2019) that likely interacted with the Milky Way (Chakrabarti et al. 2019). However, the formation of moving groups due to dwarf galaxy interactions has not yet been studied with full N-body simulations. Motivated by these earlier works that indicate that the MW may has been perturbed by dwarf galaxies, we focus our study here in trying to understand if some of the moving groups in the Galaxy may have arisen from dwarf galaxy interactions.","Citation Text":["Hunt et al. 2018"],"Functions Text":["Recent work also finds that transient spiral structure","may lead to the formation of moving groups,"],"Functions Label":["Background","Background"],"Citation Start End":[[826,842]],"Functions Start End":[[770,824],[844,887]]}
{"Identifier":"2019AandA...632A..37B__Desidera_et_al._2015_Instance_1","Paragraph":"Of the almost 4000 exoplanets known today, more than 3700 were discovered via the radial-velocity (RV) or transit methods1. According to the database, only three of them (V830 Tau b, Donati et al. 2016; K2-33 b, David et al. 2016, and TAP 26 b, Yu et al. 2017) are younger than 100 Myr. The main reason for this is the strong stellar activity of young stars, which makes it hard to find the subtle planetary signal in the large stellar variations. This is unfortunate for two reasons: first, planet formation takes place in young systems and at least gas giants need to form before the disk has dissipated after less than a few tens of millions of years (e.g., Ercolano & Pascucci 2017). Second, planets at large orbital distances (\u227350 AU) are almost exclusively detected via direct imaging (DI), which is best applicable to young systems where the planets are still hot from their formation. Thus, in order to discover all planets in a system, one either needs to image old stars \u2013 which seems currently impossible given the already low detection rate around young stars probed by large DI surveys (e.g., Desidera et al. 2015; Lannier et al. 2016; Tamura 2016; Stone et al. 2018) \u2013 or try to minimize the impact of the stellar activity of young stars. A lot has been done to understand and characterize stellar activity (e.g., Dumusque 2018). Lindegren & Dravins (2003) further estimate the effects of stellar activity such as oscillation, granulation, meridional flow, long-term magnetic cycle, surface magnetic activity and rotation, gravitational redshift and many more on the RV measurement. Meunier & Lagrange (2019) then try to model the effect of this kind of activity signal on RV data of mature stars. With our data probing activity timescales of days to years, we are mainly probing the combined effect of stellar rotation, reconfiguration of active regions, and long-term magnetic cycles. Still, large uncertainties remain in the prediction and interpretation of any RV signal, in particular for young pre-main sequence stars. But since this is what we measure, knowledge about the typical RV variability is important, for example for developing and testing RV activity models or planning RV surveys. In this paper we therefore derive a model-free analytic relation between stellar jitter, stellar age, and lag, where lag denotes the timescale on which the jitter is measured.","Citation Text":["Desidera et al. 2015"],"Functions Text":["Thus, in order to discover all planets in a system, one either needs to image old stars \u2013 which seems currently impossible given the already low detection rate around young stars probed by large DI surveys (e.g.,","\u2013 or try to minimize the impact of the stellar activity of young stars."],"Functions Label":["Background","Background"],"Citation Start End":[[1106,1126]],"Functions Start End":[[893,1105],[1181,1252]]}
{"Identifier":"2020MNRAS.493.4868L__Stephens_et_al._2019_Instance_1","Paragraph":"Recently, polarized (sub)millimetre emission has been detected in an increasing number of discs by Atacama Large Millimeter\/submillimeter Array (ALMA) with its high sensitivity and angular resolution. However, the origin of disc polarization remains uncertain, since grains do not have to be aligned with just the magnetic field (Kataoka et al. 2017; Yang et al. 2019). They may also be aligned in the direction of the radiative anisotropy (Lazarian & Hoang 2007a; Tazaki, Lazarian & Nomura 2017) or the drift velocity of the grains relative to the ambient gas (Gold 1952; Lazarian 1995; Lazarian & Hoang 2007b). Furthermore, even spherical grains can produce polarized emission by self-scattering of large grains in an anisotropic radiation field (Kataoka et al. 2015; Yang et al. 2016, 2017; Stephens et al. 2019). The scattering interpretation of the disc polarization is favoured in several targets (e.g. Stephens et al. 2014, 2017; Kataoka et al. 2016; Bacciotti et al. 2018; Dent et al. 2019; Girart et al. 2018; Harris et al. 2018; Hull et al. 2018; Lee et al. 2018). One way to gauge the effects of scattering and identify polarization from aligned grains would be to observe at multiple wavelengths since the efficiency for scattering for grains of given sizes decreases rapidly with the wavelength in the optically thin and small-particle (or Rayleigh scattering) limit. Indeed, in the disc of Class I protostar BHB 07-11, Alves et al. (2018) detected polarization with ALMA at three wavebands (Bands 3, 6, and 7 or \u223c 3, 1.3, and 0.87\u2009mm, respectively) with consistent polarization orientations across three bands and increasing polarization fraction with wavelength, which is generally not expected for scattering-induced polarization. The rather high mean polarization fractions (\u223c7.9, 5.3, and 3.5\u2009${{\\ \\rm per\\ cent}}$ for Bands 3, 6, and 7 respectively) are also higher than those typically produced in models of scattering-induced disc polarization ($\\sim \\!1 {{\\ \\rm per\\ cent}}$). At least for this well-studied source, scattering is unlikely the main mechanism for producing the observed multiwavelength disc polarization and aligned grains are favoured.","Citation Text":["Stephens et al. 2019"],"Functions Text":["Furthermore, even spherical grains can produce polarized emission by self-scattering of large grains in an anisotropic radiation field"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[794,814]],"Functions Start End":[[613,747]]}
{"Identifier":"2022ApJ...939..117Z__Malkan_&_Moore_1986_Instance_1","Paragraph":"Blazars are a subclass of active galactic nuclei (AGNs) with relativistic jets of high-energy particles pointing near our line of sight (e.g., Urry & Padovani 1995). Their nonthermal emission is generally detected across the entire electromagnetic spectrum from radio to \u03b3-ray bands. Blazars are subclassified into flat-spectrum radio quasars (FSRQs) and BL Lac objects (BL Lacs), according to the equivalent width of the emission lines in their optical spectrum (Stickel et al. 1991; Stocke et al. 1991; Marcha et al. 1996). These two subclasses of blazars are thought to be intrinsically different, perhaps based on their accretion mode (Dermer & Giebels 2016). FSRQs have high luminosity and a thin and radiatively efficient black hole accretion disk (Malkan & Moore 1986), while BL Lacs are powered by an advection-dominated, low radiative efficiency accretion flow (Dermer & Giebels 2016; Blandford et al. 2019). The jet emission is relativistically beamed (Ghisellini 2019), with a Doppler boosting factor corresponding to a bulk Lorentz factor of several to greater than 10 (Pushkarev et al. 2009). In both cases, the broadband spectra consist of two broad humps, one peaking in the IR-to-X-ray regime and the other peaking in the \u03b3-ray regime. The low-energy peak is believed to be due to synchrotron emission, while the high-energy peak is likely due to inverse Compton scattering of low-energy photons of either the same synchrotron photons (for BL Lacs) or external photons from the disk\/BLR (for FSRQs) (e.g., Dutka et al. 2017). However, some blazars might not necessarily be detected in \u03b3-rays (e.g., Paliya et al. 2017). Indeed, a recent study showed that blazars undetected in \u03b3-rays are likely to have relatively smaller Doppler factors and more disk dominance (Paliya et al. 2017). In the case of strong Compton scattering, the beaming of \u03b3-rays could be larger than, e.g., that seen in the radio (Dermer 1995), leading to the possible nondetection (or reduced detection efficiency) of \u03b3-rays from sources not seen exactly pole-on.","Citation Text":["Malkan & Moore 1986"],"Functions Text":["FSRQs have high luminosity and a thin and radiatively efficient black hole accretion disk"],"Functions Label":["Background"],"Citation Start End":[[755,774]],"Functions Start End":[[664,753]]}
{"Identifier":"2022AandA...663A...4S__J\u00f6nsson_et_al._2020_Instance_1","Paragraph":"The pioneer of spectroscopic surveys, the Radial Velocity Experiment (RAVE), produced its first data release (DR) 15 years ago (Steinmetz et al. 2006) and its final one, DR6, last year (Steinmetz et al. 2020a,b). In the meantime, other surveys have been operated, and survey designers have developed new methodologies, learning progressively from their own and one another\u2019s experience on how to reduce biases and uncertainties in the automated determination of APs and abundances from massive datasets of stellar spectra with various resolutions and spectral coverages. Besides RAVE, other surveys have published successive DRs that are available for public use. At the time of writing, there is open access to the DR16 of the Apache Point Observatory Galactic Evolution Experiment (APOGEE; J\u00f6nsson et al. 2020), the DR3 of the Galactic ArchaeoLogy with HERMES project (GALAH; Buder et al. 2021), the DR6 of RAVE (Steinmetz et al. 2020a,b), the DR5 of the Large sky Area Multi-Object Fibre Spectroscopic Telescope (LAMOST; Luo et al. 2015, 2019), Sloan Extension for Galactic Understanding and Exploration (SEGUE; Yanny et al. 2009), and the Gaia-ESO Survey (DR3; Gilmore et al. 2012; Randich & Gilmore 2013). Additional versions of APs, based on different methods, are also provided for RAVE DR6 (Guiglion et al. 2020) and for LAMOST DR5 (Xiang et al. 2019). The next generation of optical and near-infrared spectrographs, wide-field and massively multiplexed, is in preparation and will soon provide even larger catalogues of APs and abundances, such as the WHT Enhanced Area Velocity Explorer (WEAVE; Dalton et al. 2012), the Multi-Object Optical and Near-infrared Spectrograph (MOONS) on ESO\u2019s Very Large Telescope (Taylor et al. 2018), the 4-metre Multi-Object Spectroscopic Telescope (4MOST; de Jong et al. 2019), the Prime Focus Spectrograph (PFS; Takada et al. 2014), and the Maunakea Spectroscopic Explorer (MSE; McConnachie et al. 2016). In terms of numbers of stars, the most revolutionary survey will certainly be that of the Gaia space mission (Gaia Collaboration 2016), which will deliver in its DR3 in 20221 estimates of the physical properties, including metallicities, for millions of stars obtained with various methods through an astrophysical parameter inference system (Bailer-Jones et al. 2013).","Citation Text":["J\u00f6nsson et al. 2020"],"Functions Text":["Besides RAVE, other surveys have published successive DRs that are available for public use. At the time of writing, there is open access to the DR16 of the Apache Point Observatory Galactic Evolution Experiment (APOGEE;"],"Functions Label":["Background"],"Citation Start End":[[792,811]],"Functions Start End":[[571,791]]}
{"Identifier":"2020AandA...639A..20K___2015_Instance_1","Paragraph":"where \n\n\n\n\nD\n\u03f5\n\n\nDt\n\n\n\n$ \\frac{D \\epsilon}{Dt} $\n\n\n is the advective derivative of the energy density, \n\n\n\n\n\u03ba\n\n=\n\n\u03ba\n0\n\n\nT\n\n5\n\/\n2\n\n\n\n\nb\n\n\u0302\n\n\n\n$ {\\boldsymbol{\\kappa}} = \\kappa_0 T^{5\/2} \\hat{{\\boldsymbol{b}}} $\n\n\n is the Spitzer conductivity along magnetic field lines, Qheat and Qcool are heating and radiative cooling rates, and \u03c1, v, T, P and g are the plasma density, velocity, temperature, pressure, and gravitational acceleration respectively. Such user-defined heating term usually has a form of an exponentially decreasing function along the vertical coordinate y; \n\n\n\n\nQ\nheat\n\n=\n\nc\n0\n\nexp\n\n(\n\u2212\n\ny\n\u03bb\n\n)\n\n\n\n$ Q_{\\mathrm{heat}} = c_0 \\exp (-\\frac{\\mathit{y}}{\\lambda}) $\n\n\n, where c0 is the peak heating rate and \u03bb is the heating scale height. The user defined heating is therefore highly stratified, spatially smooth and steady (e.g. M\u00fcller et al. 2003; Fang et al. 2013, 2015; Xia et al. 2017). The need for the user-defined heating in the previous coronal rain simulations arises from the fact that they typically do not include any self-consistent dissipation mechanisms. They also do not cover complete lower solar atmosphere including chromosphere, photosphere, and convection zone, therefore omitting key physical processes in the lower atmosphere, such as magneto-convection, associated magnetic buffeting, braiding, and flows. Another drawback of several coronal rain simulations is the commonly used approximation that all of the plasma cooling (i.e. the process essential for modelling the thermal instability and catastrophic cooling) occurs via optically thin radiative losses. This approximation is perfectly valid in the corona but ceases to apply for cool plasma (below temperatures of a few 100 000 K). Such assumption means that regardless whether the radiative loss function is calculated from CHIANTI (Dere et al. 2019) or using scaling law approximations (e.g. Rosner et al. 1978), there is a cut-off temperature for the radiative cooling of the plasma condensations. Once electron recombination starts during the cooling process, the energy gained (which depends on the ionisation potential and thus also on the ionisation degree of the plasma) is expected to slow down the cooling rate. The thermal evolution of the plasma condensations is therefore not modelled correctly at low temperatures in the previous coronal rain simulations.","Citation Text":["Fang et al.","2015"],"Functions Text":["The user defined heating is therefore highly stratified, spatially smooth and steady (e.g."],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[859,870],[877,881]],"Functions Start End":[[748,838]]}
{"Identifier":"2015MNRAS.450.4505H__Weiss_et_al._2013_Instance_1","Paragraph":"In the case of systems detected via transits, passage through an inclination resonance is likely to prevent a planet from being observed in transit, assuming the natal planetary system is edge-on in the first place. Such excitations may explain inner holes in otherwise densely packed systems observed by Kepler. An example is the Kepler-11 system (Lissauer et al. 2011b), which shows six planets, five of which are between 0.09 and 0.3 au, but none interior to 0.09 au. If we examine the secular architecture of this system assuming the presence of a seventh, 1\u2009M\u2295 planet located in the interior hole,1 we find avoided crossings in both eccentricity and inclination near semi-major axes \u223c0.065 au, and other secular inclination resonances at 0.0233 and 0.0763 au. Thus, it is possible that Kepler-11 contains a low-mass planet in the observed inner \u2018hole\u2019, but which is inclined relative to the outer planets. The inclination could be substantial if the current period is 8 d. The Kepler planet sample contains a large number of multiple planet systems with potential for secular resonances. KOI-94 (Weiss et al. 2013) is interesting because of the presence of a Saturn mass planet in a similar orbit to 55 Cancri. The innermost planet lies just outside a pair of secular resonances at 3.0 d (eccentricity) and 3.4 d (inclination) and so is consistent with a lack of inclination excitation. If we allow for the existence of a 1 M\u2295 planet interior to this, eccentricity resonances are found at 0.9 and 2.1 d (and a mean motion resonance at 1.9 d) and inclination resonances at 1.2 and 2.1 d. Thus, a planet that has tidally migrated to orbital periods 2 d in this system could have had its inclination pumped up and have avoided transit. Another compact multiplanet system containing a massive planet is Kepler-30 (Fabrycky et al. 2012). However, the only resonances encountered by a tidally migrating 1\u2009M\u2295 inner planet lie at 14.8 d (both inclination and eccentricity) and at 13.9 d (a mean motion resonance). These may be too far out to be strongly affected by tides. Kepler-9 (Holman et al. 2010) shows two Saturn-class planets and an interior low-mass planet, that we will assume to be 1\u2009M\u2295 again. Since this inner planet does indeed transit, we must assume little inclination excitation has taken place in this system. Indeed, the sole inclination resonance for this configuration lies at 7.5 d, so that the planet would not experience any excitation if it started the migration interior to this point. Eccentricity resonances are found at 1.0 and 2.5 d, bracketing the observed period. Kepler-51 (Steffen et al. 2013) is an example of a compact Kepler planetary system with a large inner hole but no Jovian-class planets. In this case, an inner 1\u2009M\u2295 planet experiences no secular resonances. There is a frequency commensurability at 2.9 d, but no avoided crossing because the modes are almost completely decoupled (one is driven primarily by the relativistic precession).","Citation Text":["Weiss et al. 2013"],"Functions Text":["KOI-94","is interesting because of the presence of a Saturn mass planet in a similar orbit to 55 Cancri.","The innermost planet lies just outside a pair of secular resonances at 3.0 d (eccentricity) and 3.4 d (inclination) and so is consistent with a lack of inclination excitation. If we allow for the existence of a 1 M\u2295 planet interior to this, eccentricity resonances are found at 0.9 and 2.1 d (and a mean motion resonance at 1.9 d) and inclination resonances at 1.2 and 2.1 d. Thus, a planet that has tidally migrated to orbital periods 2 d in this system could have had its inclination pumped up and have avoided transit."],"Functions Label":["Motivation","Motivation","Compare\/Contrast"],"Citation Start End":[[1101,1118]],"Functions Start End":[[1093,1099],[1120,1215],[1216,1737]]}
{"Identifier":"2021ApJ...923L..22A__Rosado_et_al._2015_Instance_3","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"],"Functions Text":["We expect to detect the GWB first, followed by detection of individual SMBHBs","that stand out above the GWB."],"Functions Label":["Background","Background"],"Citation Start End":[[1369,1387]],"Functions Start End":[[1269,1346],[1433,1462]]}
{"Identifier":"2019AandA...631A..35B__Bridges_et_al._(1996)_Instance_3","Paragraph":"The collision velocity dependence of the coefficient of restitution between particles was observed in experiments (Bridges et al. 1996; Higa et al. 1996) and is discussed in the literature (e.g., Ram\u00edrez et al. 1999; Zhang & Vu-Quoc 2002). However, the experiments by Hei\u00dfelmann et al. (2010), used in the present paper to support our assumption of a constant coefficient of restitution, do not see a variation of the coefficient of restitution between particles at low collision velocities (\u2264 1 cm s\u22121). This discrepancy in results might originate in the nature of the collisions studied in these different experiments: Bridges et al. (1996) and Higa et al. (1996) performed collisions of a particle with a flat surface, while Hei\u00dfelmann et al. (2010) observed particle-particle collisions in a free-floating environment. The latter is an experimental environment very similar to NanoRocks. In such inter-particle collisions in free-floating environments, other physical effects lead to a different behavior of the energy dissipation during collisions. In particular, the damping behavior of a large plate or surface is expected to differ from that of a same-sized particle, so that the velocity dependence of the coefficient of restitution might be an effect of the experimental setup in Bridges et al. (1996) and Higa et al. (1996). Colwell et al. (2016) and Brisset et al. (2018) studied collisions between a round cm-sized particle and a flat surface of fine grains. They also observed an increase of the coefficient of restitution with decreasing collision velocity. While the composition of the target surface was different than in Bridges et al. (1996) and Higa et al. (1996) (fine granular material vs. solid ice), the similar behavior of the coefficient of restitution supports the fact that particle-surface collisions are very different from particle-particle collisions, and coefficients of restitution are only velocity dependent for collisions with particles with very different sizes (a much larger particle can be approximated as a target surface).","Citation Text":["Bridges et al. (1996)"],"Functions Text":["In particular, the damping behavior of a large plate or surface is expected to differ from that of a same-sized particle, so that the velocity dependence of the coefficient of restitution might be an effect of the experimental setup in","and Higa et al. (1996)."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1290,1311]],"Functions Start End":[[1054,1289],[1312,1335]]}
{"Identifier":"2016ApJ...822...72C__Liu_et_al._2010_Instance_1","Paragraph":"Here, for the first time, we have identified highly dynamic non-potential activity on QS-like supergranular network scales. These events overlie mixed polarity network flux elements near the spatial resolution of HMI, and are the first non-potential structures to be unassociated with strong concentrations of bipolar magnetic flux. One event (2011b August 05) shows eruptive activity in the form of jets, which is similar to larger-scale micro-sigmoids (Raouafi et al. 2010) and even QS bright point sigmoids (Chesny et al. 2015). The existence of flaring non-potential fields in the QS-like mixed network immediately shows that supergranular-scale magnetic fields can support processes similar to sigmoid formation (Chesny et al. 2013). Strong non-potential field arcades have been observed in hot X-ray sigmoids on scales of hundreds of arseconds (Moore et al. 2001; Gibson et al. 2002; Liu et al. 2010), micro-sigmoids in soft X-ray imaging on scales of \u223c50\u2033 (Mandrini et al. 2005; Raouafi et al. 2010) and small-scale AR EUV (Zheng et al. 2012, 2013), and now at EUV temperatures in QS-like mixed network fields with lengths down to \u223c10\u2033. This range of lengths over a range of temperatures and magnetic field topologies points directly to self-similar mechanisms influencing plasma and magnetic field dynamics at a range of scales. Our findings suggest that the processes driving some large-scale eruptions (i.e., flux emergence, helicity build-up, and flux cancellation leading to non-potential field heating (Chen et al. 2014)), can also manifest in a range of configurations on sub-network size scales in QS-like magnetic field configurations. These QS flaring non-potential fields are similar to their large-scale counterparts, but not as strict in their evolution. The diversity in the observed topologies may scale with the diversity of magnetic configurations that exist in the supergranulation network. QS non-potential fields evolve in multi-polar environments, and are not restricted to strongly bipolar dominated regions as in larger-scale, higher temperature events. Despite this, one of the presented events (2011a August 05) results in a post-flare potential loop arcade, which is similar to some observed AR sigmoid fields (Moore et al. 2001).","Citation Text":["Liu et al. 2010"],"Functions Text":["Strong non-potential field arcades have been observed in hot X-ray sigmoids on scales of hundreds of arseconds","This range of lengths over a range of temperatures and magnetic field topologies points directly to self-similar mechanisms influencing plasma and magnetic field dynamics at a range of scales."],"Functions Label":["Background","Background"],"Citation Start End":[[890,905]],"Functions Start End":[[739,849],[1144,1336]]}
{"Identifier":"2016AandA...588A...2L__M\u00e4tzler_(1998)_Instance_1","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)"],"Functions Text":["Absorption coefficients for pure water ice (kH2O) at sub-mm-to-cm wavelengths are discussed extensively by","who also provides several analytic formulations to estimate them as a function of frequency and temperature along with illustrative plots."],"Functions Label":["Background","Background"],"Citation Start End":[[545,559]],"Functions Start End":[[438,544],[561,699]]}
{"Identifier":"2021AandA...655A..99D__Carigi_et_al._2005_Instance_2","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)"],"Functions Text":["This trend is in agreement with the metallicity dependent yields from","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)."],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[1909,1929]],"Functions Start End":[[1839,1908],[1931,2140]]}
{"Identifier":"2021MNRAS.500.3002B__Kremer_et_al._2020_Instance_1","Paragraph":"It is as well of wide interest and diverse implications (e.g. Abadie et al. 2010; Mandel & Farmer 2017) to consider how and under which conditions NSs and BHs would pair up in tight-enough binaries so that they can spiral in by emitting GW radiation and merge within the Hubble time. Recent numerical studies based on analytical (H\u00e9non 1975), direct N-body integration (Aarseth 2003), and Monte Carlo approach (H\u00e9non 1971; Joshi, Rasio & Portegies Zwart 2000; Hypki & Giersz 2013) show that the retention of BHs in dense stellar clusters of wide mass range, beginning from low-\/medium-mass young and open clusters (e.g. Banerjee, Baumgardt & Kroupa 2010; Ziosi et al. 2014; Mapelli 2016; Banerjee 2017, 2018a, b; Park et al. 2017; Di Carlo et al. 2019; Kumamoto, Fujii & Tanikawa 2019; Rastello et al. 2019) through globular clusters (e.g. Breen & Heggie 2013; Morscher et al. 2013; Sippel & Hurley 2013; Arca-Sedda 2016; Hurley et al. 2016; Rodriguez, Chatterjee & Rasio 2016; Wang et al. 2016; Askar et al. 2017; Chatterjee, Rodriguez & Rasio 2017a; Chatterjee et al. 2017b; Fragione & Kocsis 2018; Rodriguez et al. 2018; Antonini & Gieles 2020; Kremer et al. 2020) to galactic nuclear clusters (e.g. Antonini & Rasio 2016; Hoang et al. 2018, 2019; Antonini, Gieles & Gualandris 2019; Arca-Sedda & Capuzzo-Dolcetta 2019; Arca Sedda 2020), comprise environments where BHs can pair up through close dynamical interactions, which, furthermore, lead to general-relativistic (hereafter GR) coalescences of these BBHs. Being much more massive than the rest of the stars, the BHs, which remain gravitationally bound to a cluster after their birth, spatially segregate and remain highly concentrated in the cluster\u2019s innermost (and densest) region (e.g. Banerjee et al. 2010; Morscher et al. 2015) due to dynamical friction (Chandrasekhar 1943; Spitzer 1987) from the stellar background. This is essentially an early core collapse of the cluster leading to its post-core-collapse behaviour (H\u00e9non 1975; Spitzer 1987; Heggie & Hut 2003), i.e. energy generation in the \u2018collapsed\u2019 BH core leading to an overall expansion of the cluster with time (Breen & Heggie 2013; Antonini & Gieles 2020). Inside this core, BHs undergo close binary\u2013single and binary\u2013binary encounters giving rise to compact subsystems (triples, quadruples, or even higher order multiples) whose resonant evolution can lead to GR inspiral and merger of their innermost binaries (Leigh & Geller 2013; Samsing, MacLeod & Ramirez-Ruiz 2014; Geller & Leigh 2015; Banerjee 2018b; Samsing 2018; Zevin et al. 2019), through the binaries\u2019 eccentricity pumping. The breakup of such subsystems or simply close, flyby encounters may also lead to a sufficient boost in eccentricity of a BBH such that it merges either promptly, in between two close encounters (e.g. Kremer et al. 2019) or within a Hubble time if it gets ejected from the cluster as a result of the interaction (e.g. Rodriguez et al. 2015; Park et al. 2017; Kumamoto et al. 2019). Note that such GR mergers can also happen in hierarchical systems, containing NSs and BHs, in a galactic field that derive from field massive-stellar multiplets (e.g. Toonen, Hamers & Portegies Zwart 2016; Antonini, Toonen & Hamers 2017; Fragione & Loeb 2019; Fragione, Loeb & Rasio 2020b).","Citation Text":["Kremer et al. 2020"],"Functions Text":["Recent numerical studies","show that the retention of BHs in dense stellar clusters of wide mass range, beginning from low-\/medium-mass young and open clusters","through globular clusters","comprise environments where BHs can pair up through close dynamical interactions, which, furthermore, lead to general-relativistic (hereafter GR) coalescences of these BBHs."],"Functions Label":["Background","Background","Background","Background"],"Citation Start End":[[1148,1166]],"Functions Start End":[[284,308],[481,613],[808,833],[1341,1514]]}
{"Identifier":"2020AandA...641A..85S__Orienti_&_Dallacasa_2008_Instance_1","Paragraph":"To derive the equipartition magnetic field of J1146+4037, we predict the rest-frame 8.4 GHz (redshifted to 1.4 GHz at z\u2004=\u20045.0059) flux density from our spectral model. However, there is no source size measurement at 1.4 GHz. We make use of the full width at half maximum (FWHM) source size of 0.74\u2005\u00b1\u20050.01 mas derived by the Gaussian fit from 5 GHz VLBI mas angular resolution observations (Frey et al. 2010). We note that in our calculations, we assume a source size that is 1.8 times larger than the FWHM, following the approach of Readhead (1994) and Orienti & Dallacasa (2008). The derived equipartition magnetic field is \n\n\n\n34\n\n\u2212\n7\n\n\n+\n8\n\n\n\n$ 34^{+8}_{-7} $\n\n\n mG. This is within the range of the equipartition magnetic fields of 17 HFP radio sources (7\u201360 mG; quasars and galaxies at 0.22\u2004 \u2004z\u2004 \u20042.91; Orienti & Dallacasa 2012) and 5 HFPs at 0.084\u2004 \u2004z\u2004 \u20041.887 (18\u2013160 mG; Orienti & Dallacasa 2008). The magnetic field calculated from the turnover information listed in Table 4 is \n\n\n\n1\n.\n\n8\n\n\u2212\n2.7\n\n\n+\n2.3\n\n\n\n\n$ 1.8^{+2.3}_{-2.7} $\n\n\n G assuming an SSA origin with Eq. (3), however the uncertainty is very large. The large uncertainty is caused by the fact that we only have four data points to constrain the turnover information and we do not have source size measurements at the turnover frequency, but rather we assume the source size measured at another frequency. More data taken in other wavelength bands are needed to meaningfully constrain the turnover peak, and mas resolution observations at the peak frequency are needed to give reliable magnetic field strength measurements. This may indicate that the turnover is not caused by SSA, by comparing the large magnetic field strength measured from the spectral turnover (\n\n\n\n1\n.\n\n8\n\n\u2212\n2.7\n\n\n+\n2.3\n\n\n\n\n$ 1.8^{+2.3}_{-2.7} $\n\n\n G) with the equipartition magnetic field strength (\n\n\n\n34\n\n\u2212\n7\n\n\n+\n8\n\n\n\n$ 34^{+8}_{-7} $\n\n\n mG). As J1146+4037 is a strong blazar, the turnover may be caused by its strong variability. Another possible explanation for the spectral turnover is that high-density plasma in the nuclear region attenuates the radio emission from the central active BH. High-resolution, interstellar medium observations of the nuclear region of this target may address the latter issue.","Citation Text":["Orienti & Dallacasa (2008)"],"Functions Text":["We note that in our calculations, we assume a source size that is 1.8 times larger than the FWHM, following the approach of Readhead (1994) and"],"Functions Label":["Uses"],"Citation Start End":[[553,579]],"Functions Start End":[[409,552]]}
{"Identifier":"2022AandA...663A..77N__Krolik_et_al._1981_Instance_1","Paragraph":"In our model, we used a non-hydrodynamical approach based on assumptions of the motion of separate clouds under gravity and the action of radiation pressure acting on dust. This has considerable limitations but they are justified as the first approximation for modeling the LIL part of the BLR. As discussed in the classical paper of line-driven wind model (for HIL part of BLR) by Murray et al. (1995), the optical depth of the emitting region must be moderate (column density of order of 1023 cm\u22122) and the local density is high (for LIL part it is higher than that of HIL, many authors argue for a local density about 1012 cm\u22123, for instance, Adhikari et al. 2016; Baskin & Laor 2018; Panda et al. 2018), while the BLR is extended. There are two possibilities to support a consistent picture: it is either to assume a very narrow stream of material flowing out, with the cross-section on the order of 1012 cm, as in Murray et al. (1995), where they assume lower density so the size is actually larger 1014 cm; or to assume considerable clumpiness of the medium. We followed the second approach since there are natural thermal instabilities in the plasma, such as instability caused by X-ray irradiation (Krolik et al. 1981). In this case, the plasma spontaneously forms colder clumps (at a temperature of \u223c104 K, cooled through atomic processes) embedded in a hotter medium (at temperature \u223c107 K, set at an inverse Compton temperature value). The presence of highly or fully ionized medium is aptly supported by the observations as well as by the theory. Then the two media of density contrast of order of 103 provide the rough pressure balance. Blandford et al. (1990) discussed the typical values for ionization parameter in AGN clouds covering the range from 10\u22123 to 1, that is, up to four orders of magnitudes. The precise description of the structure of the clumpy medium is very difficult. Even in the case of a single cloud exposed to irradiation, in plane-parallel approximation requires a radiative transfer to be performed, which would then show the gradual change in the density, temperature, and ionization parameters (see e.g., Baskin & Laor 2018; Adhikari et al. 2018), with the low density first and the temperature roughly at inverse Compton temperature (depending on the shape of the incident spectrum), and then a relatively rapid decrease at the subsequent ionization fronts. A proper description of this transition, calculated under constant pressure, actually requires inclusion of the electron conduction (e.g., Begelman & McKee 1990; R\u00f3\u017ca\u0144ska & Czerny 2000). Deep within the cloud, there is a further drop in temperature and a rise in density due to a decrease in the local flux as a result of absorption. As emphasized by Baskin & Laor (2018), radiation pressure also plays a dynamical role in this process. The picture is further complicated if plane-parallel approximation is abandoned. The presence of the numerous clouds of complex shapes can be fully consistent with simple estimates of the cloud number based on line shape properties as done by Arav et al. (1998). Of course, there are also certain processes that can lead to cloud destruction, such as the action of tidal forces (M\u00fcller et al. 2022), Kelvin\u2013Helmholtz instabilities, and cloud ablation), however, the destruction rate can be strongly affected by the magnetic field (e.g., McCourt et al. 2015). The relative importance between the condensation rate and destruction rate depends on the cloud size, as it is set roughly by the field length (Field 1965). Cloud formation in AGN has been seen in the numerical simulations from Waters et al. (2021), but at distances much greater than the BLR distance which was likely related to the numerical setup and the requested spatial resolution of the computations. The issue is thus extremely complex, and simple order-of-magnitude estimates based on a single density and single temperature of the cloud and intercloud medium are not fully adequate and cannot reproduce the full ionization parameter range. However, addressing this point in detail is beyond the scope of the current paper.","Citation Text":["Krolik et al. 1981"],"Functions Text":["We followed the second approach since there are natural thermal instabilities in the plasma, such as instability caused by X-ray irradiation"],"Functions Label":["Uses"],"Citation Start End":[[1207,1225]],"Functions Start End":[[1065,1205]]}
{"Identifier":"2016MNRAS.458.2870V__Busha_et_al._2011_Instance_1","Paragraph":"Unfortunately, since the study by Onions et al. only used a single dark matter halo, albeit at exquisite numerical resolution, the comparison is limited to the relatively low-mass end of the subhalo mass function (SHMF), where the cumulative mass function N( > m), exceeds unity. In order to study the massive end of the SHMF, where N(>m)  1, one needs to average over large numbers of host haloes. The abundances of these rare but massive subhaloes has important implications for, among others, the statistics of massive satellite galaxies (e.g. Boylan-Kolchin et al. 2010; Busha et al. 2011) and the detection rate of dark matter substructure via gravitational lensing (e.g. Vegetti et al. 2010, 2012). In this paper, the second in a series, we use subhalo mass functions (SHMFs) and subhalo catalogues from a variety of numerical simulations that are publicly available, and that have been obtained using different subhalo finders, to compare SHMFs, focusing on the massive end. We confirm the findings by Onions et al., that the SHMFs are consistent at the 20 per cent level at the low-mass end. At the massive end, though, different subhalo finders yield subhalo abundances that differ by more than one order of magnitude! By comparing the simulation results with a new, semi-analytical model (Jiang & van den Bosch 2016, hereafter Paper I), we demonstrate that subhalo finders that identify subhaloes based purely on density in configuration space, such as the popular subfind\u2009and bdm, dramatically underpredict the masses, but not the maximum circular velocities, of massive subhaloes. We also show that the model predictions are in excellent agreement with the simulation results when they are analysed using more advanced subhalo finders that use phase-space and\/or time domain information in the identification of subhaloes. We discuss a number of implications of our findings, in particular with regard to the power-law slope of the subhalo mass and velocity functions.","Citation Text":["Busha et al. 2011"],"Functions Text":["The abundances of these rare but massive subhaloes has important implications for, among others, the statistics of massive satellite galaxies (e.g."],"Functions Label":["Motivation"],"Citation Start End":[[575,592]],"Functions Start End":[[399,546]]}
{"Identifier":"2020ApJ...892..103Z__Melia_&_Shevchuk_2012_Instance_1","Paragraph":"Similarly, according to Equation (2), angular diameter distances can be calculated from observed angular sizes as\n7\n\n\n\n\n\nand we also treated the length scale lm as a free parameter. In order to test the CDDR using different samples one should use a redshift matching criterion \u0394z  0.005 (Li et al. 2011; Liao et al. 2016). However, it turned out difficult to fulfill this criterion in order to compare distances derived from QSO [XUV] and QSO [CRS] directly. Therefore, proceeding in a similar manner as Li & Lin (2018), we reconstruct QSO [CRS] angular size as a function of redshift from the binned data. The angular size of intermediate-luminosity quasars was grouped into 20 redshift bins of width \u0394z = 0.1 starting from the smallest redshift of this sample. Median values of angular size plotted against the mean redshift in each bin are shown in Figure 4. The python package GaPP based on Gaussian Processes (Seikel et al. 2012) was used for the reconstruction process that depends on the mean function and the covariance function \n\n\n\n\n\n. In order to discuss the influence of the choice of these two prior functions, we studied four mean functions and three covariance functions. The prior mean functions that we discussed are the following: zero mean function, the theoretical function of angular size calculated from the angular diameter distance under the assumption of three cosmological models: flat \u039bCDM with \u03a9m = 0.27, so-called Rh = ct universe (Melia & Shevchuk 2012), and a Mirage model with \u03a9m = 0.27, w0 = \u22120.7 and w1 = \u22121.09 (Shafieloo et al. 2012). The linear size scaling factor lm = 11.42 pc was assumed as calibrated with SN Ia in Cao et al. (2017a). These functions (except the zero mean function) are shown in Figure 2. There are many possible covariance functions and we studied the three most popular ones. They comprise: the squared exponential function\n8\n\n\n\n\n\nthe \n\n\n\n\n\n\n\n9\n\n\n\n\n\nand Cauchy covariance function\n10\n\n\n\n\n\nwhere \u03c3f and \u2113 are hyperparameters that control the amplitude of deviation from the mean function and the typical length scale in x-direction, respectively. It is instructive to discuss the effects of the mean function and covariance function selection (prior assumptions) on the reconstruction. In order to show the impact of covariance function choice, we fixed the zero mean and preformed reconstruction with three different covariance functions mentioned above. The result obtained under the assumption of squared exponential covariance function is illustrated in Figure 4, where the green solid line represents the reconstructed \u03b8(z) relation and the green region around it represents a 1\u03c3 uncertainty band. The blue dashed line and black dashed\u2013dotted line represent the reconstructed \u03b8(z) relation with the Mat\u00e9rn and Cauchy covariance functions, respectively. Their uncertainty bands are not shown in order to not blur the picture since they are similar to the one displayed. One can see that differences between reconstructions performed with different choices of covariance function are insignificant. Similarly, we checked sensitivity of reconstructions with respect to the choice of the mean function fixing the covariance as a squared exponential one and using three main functions mentioned above. It turned out that the impact of mean function choice on the reconstruction was even smaller than that of the covariance function. Therefore for further calculations we assumed the zero mean function and the squared exponential covariance function to get the reconstructed \u03b8(z) function (i.e., the green line and region in Figure 4). Using this reconstructed relation we were able to have a one-to-one matching between the QSO [CRS] angular diameter distance and the QSO [XUV] luminosity distance at the same redshift.","Citation Text":["Melia & Shevchuk 2012)"],"Functions Text":["The prior mean functions that we discussed are the following: zero mean function, the theoretical function of angular size calculated from the angular diameter distance under the assumption of three cosmological models: flat \u039bCDM with \u03a9m = 0.27, so-called Rh = ct universe"],"Functions Label":["Uses"],"Citation Start End":[[1460,1482]],"Functions Start End":[[1186,1458]]}
{"Identifier":"2021AandA...649A..58L__Bemporad_et_al._(2018)_Instance_1","Paragraph":"The leading edges of the transients normally leave bright traces in the images of visible light, inspiring many methods that were developed to derive their locations and velocities, such as the icecream cone model (Fisher & Munro 1984), the graduated cylindrical shell (GCS) model (Thernisien 2011), geometric triangulation methods (Liu et al. 2010), mask-fitting methods (Feng et al. 2012), and trace-fitting methods including the point-p, fixed-\u03a6, harmonic mean, and self-similar expansion fitting methods (e.g., Sheeley et al. 1999; Howard et al. 2006; Davies et al. 2012; M\u00f6stl & Davies 2013). To derive the velocity distribution inside one transient rather than only at its leading edge, some other techniques have been proposed. Colaninno & Vourlidas (2006) applied an optical flow tool to extract the velocity vector of a coronal mass ejection (CME) in digital images. Feng et al. (2015) derived the radial velocity profiles of the whole CME from the spatial distribution of its density given by the mass continuum equation. A cross-correlation method was applied to derive continuous 2D speed maps of a CME from coronagraphic images by Bemporad et al. (2018). In their work, the radial shift pixel by pixel is determined by maximizing the cross correlation between the signal in a radial window at one frame and the signal in a radial shifted window at the previous frame, and the radial speed just equals the radial shift over the time interval between the two frames. Ying et al. (2019) improved this cross-correlation method by analyzing data in three steps: forward step (FS), backward step (BS), and average step (AS). In the FS (BS), the 2D velocity map between the current and the previous (next) frame is constructed with almost the same method as Bemporad et al. (2018). In the AS the average, velocity is obtained from the FS and BS. The velocities derived by all these methods are the component of the flow velocity vector projected onto the POS. This may underestimate the velocity especially for transients that do not propagate in the POS. Methods such as the polarizaition ratio technique (Moran & Davila 2004; DeForest et al. 2017) or the local correlation tracking (LCT) method (Mierla et al. 2009) can derive the 3D geometric information of the whole transients, but not the velocity distribution. Bemporad et al. (2018) chose the propagating direction averaged over the whole CME derived by the polarization ratio technique to correct the radial speed in the 2D maps, but the key information along the LOS is still lacking.","Citation Text":["Bemporad et al. (2018)"],"Functions Text":["A cross-correlation method was applied to derive continuous 2D speed maps of a CME from coronagraphic images by"],"Functions Label":["Uses"],"Citation Start End":[[1144,1166]],"Functions Start End":[[1032,1143]]}
{"Identifier":"2022MNRAS.516.5289M__Thompson_et_al._2015_Instance_1","Paragraph":"Given the number densities within the mass-dissociation index plane of Fig. 8, we now ask ourselves whether known dissociated clusters, such as the Bullet cluster, are expected in L210N1024NR? The Bullet Cluster has a mass of $\\sim 1.5 \\times 10^{15} \\, {\\rm M}_{\\odot }$ (e.g. Clowe et al. 2004; Brada\u010d et al. 2006; Clowe et al. 2006) and we estimated a dissociation index of SBullet \u223c 0.335 \u00b1 0.06. As seen in Fig. 8 there are no Bullet cluster analogues (structures of approximate mass and dissociation) in L210N1024NR, this is unsurprising as a simulation requires a significantly larger volume than that of L210N1024NR ((210cMpc\u2009h\u22121)3) to expect such an object (e.g. Lee & Komatsu 2010; Thompson & Nagamine 2012; Bouillot et al. 2015; Kraljic & Sarkar 2015; Thompson et al. 2015). From the distribution presented in Fig. 8, it is trivial to estimate the required cosmological volume (the effective volume, Veff) to expect structures of a given mass and dissociation index. By separating the 2D distribution on the mass-dissociation index planes into the component 1D distributions of mass and dissociation the effective volume is computed as\n(12)$$\\begin{eqnarray}\r\nV_\\text{eff}~^{-1} &amp;=&amp;\\int \\int \\,{\\rm{ d}} S \\, {\\rm{ d}} M \\phi (S, M) \\\\\r\n&amp;=&amp; \\int _{S_\\text{a}}^{S_\\text{b}} \\, {\\rm{ d}} S \\phi _S(S) \\int _{M_\\text{a}}^{M_\\text{b}} \\, {\\rm{ d}} M \\phi _M(M)~,\r\n\\end{eqnarray}$$where \u03d5S(S) is the number density function associated with S and $\\phi _\\mathit {M}(\\mathit {M})$ is the mass function presented in Fig. 7. Assuming a probable range of S = 0.335 \u00b1 0.06 and $1 \\lt M \\lt 2 \\times 10^{15} \\, {\\rm M}_{\\odot }$ we estimate a number density \u223c4.92 \u00d7 10\u221210 Mpc\u22123 or that an effective volume of \u223c2.03 Gpc3 would be required to observe a single Bullet-like cluster. This result is inline with the number density estimate of the order of \u223c10\u221210 Mpc\u22123 by Thompson et al. (2015), which improves on previous estimates (e.g. Lee & Komatsu 2010; Thompson & Nagamine 2012; Bouillot et al. 2015) due to more sophisticated halo finding methods (e.g. Behroozi, Wechsler & Wu 2013). Conversely, it was estimated by Kraljic & Sarkar (2015) (utilizing the same halo finder as Thompson et al. 2015) that given an effective volume of \u223c14.6 Gpc3, no Bullet cluster analogues are expected, however as indicated by a pairwise velocity distribution it would be expected that present binary halo\u2013halo orbits have the potential to form a Bullet-like object.","Citation Text":["Thompson et al. 2015"],"Functions Text":["As seen in Fig. 8 there are no Bullet cluster analogues (structures of approximate mass and dissociation) in L210N1024NR, this is unsurprising as a simulation requires a significantly larger volume than that of L210N1024NR ((210cMpc\u2009h\u22121)3) to expect such an object (e.g."],"Functions Label":["Uses"],"Citation Start End":[[763,783]],"Functions Start End":[[401,671]]}
{"Identifier":"2021MNRAS.504..444C__Russell_et_al._2020b_Instance_1","Paragraph":"To better visualize the jet behaviour in this phase, we show in Fig. 10 the radio light curve at the core location for the first 25 d of the outburst, taken from Fig. 2. The compact jet emission peaks on MJD 58520 and then starts to decay on MJD 58521, 1\u2009d before the system enters in the soft state and 2\u2009d after the inferred RK1 ejection date (see Section 3.3). This whole evolution is accompanied by a smooth transition between the optically thick and the optically thin regimes of synchrotron radio emission. Compact jets are usually observed to quench at the transition from the hard to soft state (e.g. Fender et al. 1999; Corbel et al. 2000). This phenomenon is observed to start at higher frequencies (where the emission is produced closer to the compact object) and terminates as the jet break evolves through the radio band (Russell et al. 2013b, 2014; Russell et al. 2020b). It is not yet clear if compact jets switch off before or during the launch of discrete ejections, and how (or if) the two events are linked (Russell et al. 2020a). We do not have a MeerKAT observation on MJD 58521, but, as the flat spectrum obtained with ATCA suggests, the flux density would have been at the \u223c50 mJy level. The following observations show instead a quick rise in flux density and a subsequent decay until MJD 58531, when the source is steadily in the soft state. With our data, we cannot conclude on the origin of the radio emission on MJD 58521, which could be produced either from the compact jet that is quenching (with no evidence of the spectral break in the radio band), or by the first self-absorbed part of the radio flare observed to peak on MJD 58523. The latter scenario implies that the jet significantly quenched in less than 2\u2009d, a shorter time-scale with respect to what observed for other sources (e.g. Russell et al. 2013b, 2020a). Therefore, we are not able to precisely order in time the ejection of RK1 and the quenching of the compact jets, and thus we cannot draw conclusions on a potential link between the two events.","Citation Text":["Russell et al. 2020b"],"Functions Text":["This phenomenon is observed to start at higher frequencies (where the emission is produced closer to the compact object) and terminates as the jet break evolves through the radio band"],"Functions Label":["Background"],"Citation Start End":[[863,883]],"Functions Start End":[[650,833]]}
{"Identifier":"2019MNRAS.488..902C__Svensson_et_al._2012_Instance_1","Paragraph":"Long-duration gamma-ray bursts (GRBs) give rise to a synchrotron afterglow, detectable at optical wavelengths if sufficiently rapid and deep follow-up observations are made. A substantial fraction, however, lack such emission even when it would be expected from extrapolation of the X-ray spectral slope (Groot et al. 1998; Fynbo et al. 2001). When the X-ray to optical spectral slope, \u03b2OX, is below the recognized threshold of 0.5, the event is classified as \u2018dark\u2019 (Jakobsson et al. 2004). This is typically evaluated at 11\u2009h post-burst to avoid contamination from early-time effects including X-ray flares and plateaus. An alternative method uses \u03b2OX  \u03b2X \u2212 0.5 to define darkness (van der Horst et al. 2009). There are two primary causes for darkness in GRBs: attenuation by dust, or rest frame ultraviolet H\u2009i absorption at high redshift (e.g. Fruchter 1999; Levan et al. 2006; Perley et al. 2009,2013; Greiner et al. 2011; Svensson et al. 2012; Zauderer et al. 2013; Chrimes et al. 2019; Higgins et al. 2019). The number of GRBs known at high-redshift (z > 5, in the epoch of reionization) is small (\u223c15, from around 500 GRBs with a known or estimated redshift, Cenko et al. 2006; Grazian et al. 2006; Jakobsson et al. 2006; Kawai et al. 2006; Ruiz-Velasco et al. 2007; Salvaterra et al. 2009; Greiner et al. 2009; Tanvir et al. 2009; Cucchiara et al. 2011; Afonso et al. 2011; Castro-Tirado et al. 2013; Laskar et al. 2014; Jeong et al. 2014b; Chornock, Fox & Berger 2014b; Tanvir et al. 2018), and each one is valuable, as they provide insight into star formation in the low mass, low luminosity galaxies which power the epoch of reionization. Because they have small projected offsets from their hosts, high-redshift GRBs with a detected afterglow uniquely allow us to place accurate, deep upper limits on the luminosities of the faintest, undetected galaxies, probing fainter galaxies than deep field studies (Berger et al. 2007; Tanvir et al. 2012; Trenti et al. 2012; McGuire et al. 2016). For those with the brightest afterglows, insight into the burst environment can be gained from absorption lines in their spectra (e.g. Kawai et al. 2006; Chornock et al. 2014a; Sparre et al. 2014; Hartoog et al. 2015).","Citation Text":["Svensson et al. 2012"],"Functions Text":["There are two primary causes for darkness in GRBs: attenuation by dust, or rest frame ultraviolet H\u2009i absorption at high redshift"],"Functions Label":["Background"],"Citation Start End":[[928,948]],"Functions Start End":[[712,841]]}
{"Identifier":"2015ApJ...807...92Y__Yi_et_al._2013_Instance_1","Paragraph":"The assumptions of the equilibrium of pressures and equality of velocities along the contact discontinuity lead to \n\n\n\n\n\n and \n\n\n\n\n\n, respectively. With the jump condition for the shocks and the equilibrium of pressures, we can obtain\n1\n\n\n\n\n\nThe Lorentz factor of the reverse shock \n\n\n\n\n\n could be approximated as\n2\n\n\n\n\n\nas long as \n\n\n\n\n\n and \n\n\n\n\n\n. Substituting Equation (2) into (1), we can obtain the following equation:\n3\n\n\n\n\n\nBecause \n\n\n\n\n\n , \n\n\n\n\n\n, and \n\n\n\n\n\n, we ignore the constant 1\/2 term in Equation (3) and thus we can obtain the solution of this equation (ignoring the negative solution, also see Panaitescu & Kumar 2004),\n4\n\n\n\n\n\nHere we obtain the relation between the Lorentz factor of the shocked fireball shell \n\n\n\n\n\n and the initial Lorentz factor \n\n\n\n\n\n (\n\n\n\n\n\n), which depends on the ratio of these two comoving densities. The number density of the ambient medium is assumed to be \n\n\n\n\n\n (Dai & Lu 1998; M\u00e9sz\u00e1ros et al. 1998; Chevalier & Li 2000; Wu et al. 2003, 2005; Yi et al. 2013); such a circumburst medium is a homogeneous ISM for k = 0, and a typical stellar wind environment for k = 2. The fireball shell is characterized by an initial kinetic energy Ek, initial Lorentz factor \n\n\n\n\n\n, and a width \u0394 in the lab frame attached to the explosion center, so the number density of the shell in the comoving frame is \n\n\n\n\n\n. The ratio of the comoving number density of the relativistic shell \n\n\n\n\n\n to the number density of the ambient medium \n\n\n\n\n\n defined in Sari & Piran (1995) is\n5\n\n\n\n\n\nwhere \n\n\n\n\n\n. The difference between the lab frame speed of the unshocked fireball shell and that of the reverse shock is (Kumar & Panaitescu 2003),\n6\n\n\n\n\n\nConsidering the thin shell case \n\n\n\n\n\n, we can calculate the radius \n\n\n\n\n\n where the reverse shock finishes crossing the fireball shell,\n7\n\n\n\n\n\nThe substitution of Equations (5) and (6) into (7) leads to\n8\n\n\n\n\n\nSo the comoving density ratio at \n\n\n\n\n\n is\n9\n\n\n\n\n\nSubstituting Equation (9) into (4), we can obtain the Lorentz factor of the reverse shock as it finishes crossing the shell:\n10\n\n\n\n\n\nTherefore, the relation between \n\n\n\n\n\n and the initial Lorentz factor \n\n\n\n\n\n is\n11\n\n\n\n\n\nand\n12\n\n\n\n\n\nFor the thin shell case, the reverse shock crossing time \n\n\n\n\n\n almost corresponds to the deceleration time Tdec, i.e., \n\n\n\n\n\n. Therefore, we can derive the initial Lorentz factor in the ISM and wind-type cases (also see Panaitescu & Kumar 2004). For k = 0 (ISM),\n13\n\n\n\n\n\nand for k = 2 (wind),\n14\n\n\n\n\n\nWith the isotropic-equivalent energy \n\n\n\n\n\n and the peak time of the afterglow onset \n\n\n\n\n\n, we can estimate the initial Lorentz factor of GRBs, where \n\n\n\n\n\n. Liang et al. (2010) discovered a tight correlation between \n\n\n\n\n\n and \n\n\n\n\n\n using 20 GRBs which show a deceleration feature in the early afterglow light curves. Other work also confirmed this correlation, but with different methods and power-law indices (Ghirlanda et al. 2012; L\u00fc et al. 2012). Using the data on \n\n\n\n\n\n and \n\n\n\n\n\n from Liang et al. (2010, 2013) and L\u00fc et al. (2012), we re-constrain the initial Lorentz factor, and also discover a tight \n\n\n\n\n\n and \n\n\n\n\n\n correlation for the ISM and wind cases. The \n\n\n\n\n\n and \n\n\n\n\n\n correlation in the wind case is even tighter than that in the ISM case, as shown in Figures 4 and 5.","Citation Text":["Yi et al. 2013"],"Functions Text":["The number density of the ambient medium is assumed to be"],"Functions Label":["Uses"],"Citation Start End":[[991,1005]],"Functions Start End":[[845,902]]}
{"Identifier":"2020AandA...637A..82D__Mason_et_al._2004_Instance_1","Paragraph":"The emission in the diffuse interstellar medium of these bands, dominated by an aromatic vibrational character, also called AIBs (aromatic infrared bands), has led to the so-called polycyclic aromatic hydrocarbon (PAH) hypothesis. Under this theory, the observed emission is related to the infrared fluorescence emission mechanism of PAH-like molecules (Leger & Puget 1984; Allamandola et al. 1985), following energetic photon absorption, although no unique PAH has been identified in the mid-infrared so far. Apart from the recent attribution of a few specific infrared emission bands to the C60 and possibly C70 fullerene molecules in some sources (Sellgren et al. 2009, 2010; Cami et al. 2010), the carriers of the AIBsremain elusive. The emission bands have been categorised in different classes (\n\n$\\mathcal{A}$\nA\n to \n\n$\\mathcal{D}$D\n) following the ascertainment of band profiles and center position variations, ensuing from a phenomenological deconvolution of the observations. (Peeters et al. 2002; van Diedenhoven et al. 2004; Matsuura et al. 2007; Sloan et al. 2007; Keller et al. 2008; Boersma et al. 2008; Pino et al. 2008; Acke et al. 2010; Carpentier et al. 2012; Gadallah et al. 2013). In the late classes of infrared emission spectra (so-called \n\n$\\mathcal{C}$C\n and \n\n$\\mathcal{D}$D\n), a mix between an aromatic and aliphatic character is observed. Class \n\n$\\mathcal{A}$A\n sources are dominated by aromatic bands, whereas classes \n\n$\\mathcal{C}$C\n and \n\n$\\mathcal{D}$D\n harboura pronounced aliphatic character. Most of the interstellar hydrocarbons are injected through the late phases of stellar evolution, post asymptotic giant branch (AGB) and protoplanetary nebula (PPN), which often display the class \n\n$\\mathcal{C}$C\n and \n\n$\\mathcal{D}$D\n spectral character, and are eventually processed later in the ISM. In absorption, bands at 3.4, 6.85, and 7.25 \u03bcm are also observed in the diffuse interstellar medium (ISM) of our Galaxy, as well as in extragalactic ISM, and can be well represented by a material with a significant amount of aliphatic character, also called HAC or a-C:H, which is a family of hydrogenated amorphous carbons (e.g. Allen & Wickramasinghe 1981; Duley & Williams 1983; Mason et al. 2004; Risaliti et al. 2006; Dartois et al. 2007; Dartois &Mu\u00f1oz-Caro 2007; Imanishi et al. 2010, and references therein). However, the AIB emission spectra are not a simple linear combination of aliphatic (such as the a-C:H observed in absorption in the diffuse ISM) and aromatic (class \n\n$\\mathcal{A}$A\n observed in emission) spectra. They show a spectral evolution of vibrational modes, including the shift of the C=C mode from 6.2 to 6.3 \u03bcm from class \n\n$\\mathcal{A}$A\n to \n\n$\\mathcal{C}$C\n (e.g. van Diedenhoven et al. 2004), and broadening and shifting back to lower wavelengths in class \n\n$\\mathcal{D}$D\n. Unlike classes \n\n$\\mathcal{A-B}$A\u2212B\n with two bands in the 7.6 to 8.2 \u03bcm range, classes \n\n$\\mathcal{C-D}$C\u2212D\n show a broad band (Szczerba et al. 2005; Matsuura et al. 2014), with the class \n\n$\\mathcal{D}$D\n one peaking at a shorter wavelength. The out of plane, predominantly aromatic, CH vibration patterns (from 11 to 13 \u03bcm) are generally difficult to reproduce with dust analogues in the laboratory. In addition to chemically pure PAH studies, a large range of dust grain analogues have been tailored using lasers, flames, VUV continuum sources, and\/or plasmas in the laboratory to tackle this identification issue (Carpentier et al. 2012; Schnaiter et al. 1999; J\u00e4ger et al. 2006; Mennella et al. 1999; Dartois et al. 2005; Furton et al. 1999; Biennier et al. 2009). Among the issues faced by laboratory synthesis of interstellar dust analogues is the ability to produce an environment that is not homogeneous for the entire batch of analogues produced (i.e. localised modifications), as many processes in the interstellar medium will only affect grain (surfaces) moieties considering, for example, hydrogen or radical accretion\/addition, cosmic ray impact, grain-grain shocks, etc. As a consequence, a large inhomogeneity at the nanometre scale between adjacent constitutive elements making the grain may be expected, introducing many local defects that will influence the nature and spectroscopic properties of such small dust grains. In this study, we address a new non-homogeneous, non-bottom-up shock approach. We used a mechanochemical synthesis method under a pressurised hydrogen atmosphere, to produce laboratory interstellar dust grain analogues to explain the infrared emission spectra of the remotely observed aliphatic and aromatic mixed interstellar dust grains observed through infrared emission bands.","Citation Text":["Mason et al. 2004"],"Functions Text":["In absorption, bands at 3.4, 6.85, and 7.25 \u03bcm are also observed in the diffuse interstellar medium (ISM) of our Galaxy, as well as in extragalactic ISM, and can be well represented by a material with a significant amount of aliphatic character, also called HAC or a-C:H, which is a family of hydrogenated amorphous carbons (e.g."],"Functions Label":["Similarities"],"Citation Start End":[[2213,2230]],"Functions Start End":[[1831,2160]]}
{"Identifier":"2015MNRAS.450.3458C__Cichowolski_et_al._2001_Instance_6","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"],"Functions Text":["there is a second OB star mentioned by","but its location is actually outside the structure"],"Functions Label":["Background","Background"],"Citation Start End":[[2695,2718]],"Functions Start End":[[2656,2694],[2719,2769]]}
{"Identifier":"2016MNRAS.457..875P__Kotov,_Churazov_&_Gilfanov_2001_Instance_1","Paragraph":"After the discovery of hard X-ray lag relative to soft X-rays in binary system (e.g. Miyamoto et al. 1988; Nowak et al. 1999), similar lags were also observed in AGN (Papadakis, Nandra & Kazanas 2001; McHardy et al. 2004). The origin of hard lag is not clearly known. One possible explanation is provided by the propagation fluctuation model in which fluctuations associated with accretion flow propagate inwards in an accretion disc and thus resulting in the emission of the soft photons from relatively outer regions earlier than the hard photons from the innermost regions (Lyubarskii 1997; Kotov, Churazov & Gilfanov 2001; Ar\u00e9valo & Uttley 2006). Recently, for example, Swift monitoring of the radio-loud NLS1 galaxy PKS 0558\u2212504 for ${\\sim } 1.5{\\rm \\ {\\rm yr}}$ has revealed that optical leads UV and UV leads soft X-rays on short time-scales of about a week (Gliozzi et al. 2013) possibly favouring the propagation model. A new type of lag has emerged from recent studies where soft photons lag to the hard photons. This is termed as the reverberation lag which is used to constrain the X-ray emitting region in AGN. Fabian et al. (2009) discovered the reverberation lag \u223c 30 s for the first time in a NLS1 galaxy 1H0707\u2212495. Since then, such lags have been observed in dozen of Seyfert galaxies (Zoghbi et al. 2010; de Marco et al. 2011; Emmanoulopoulos, McHardy & Papadakis 2011b; Zoghbi & Fabian 2011; Zoghbi, Uttley & Fabian 2011; Cackett et al. 2013; De Marco et al. 2013; Fabian et al. 2013; Kara et al. 2013). The most of above cases reveal reverberation lag \u223c100 s supporting the compact nature of X-ray-emitting region within few gravitational radii of a supermassive black hole (SMBH). In case of strong illumination, such as that implied by observation of strong blurred reflection, UV\/optical emission from AGN may be dominated by the reprocessed emission and the variations in the optical\/UV band emission lag behind the X-rays (e.g. McHardy et al. 2014). About a five year long campaign of Seyfert 1 galaxy Mrk 79 using six ground-based observatories for optical and RXTE for X-ray observations, Breedt et al. (2009) have shown zero lag between optical and X-rays on time-scale of about a day. Their study of correlated X-ray and optical emission suggests X-ray reprocessing on short time-scale of days and the changes in the optical emission on long time-scale of \u223c few years can be attributed to the variations in the accretion rate.","Citation Text":["Kotov, Churazov & Gilfanov 2001"],"Functions Text":["The origin of hard lag is not clearly known. One possible explanation is provided by the propagation fluctuation model in which fluctuations associated with accretion flow propagate inwards in an accretion disc and thus resulting in the emission of the soft photons from relatively outer regions earlier than the hard photons from the innermost regions"],"Functions Label":["Background"],"Citation Start End":[[594,625]],"Functions Start End":[[223,575]]}
{"Identifier":"2021MNRAS.503..354G__Hou_&_Han_2014_Instance_1","Paragraph":"The spatial distribution of OB stars and associations, young long-period Cepheids and open clusters, star-forming regions, H\u2009ii regions, interstellar dust, and giant molecular and neutral gas clouds in the solar vicinity that have been in existence generally \u03c4 \u2272 108 yr is known to correlate with the location of the inner Sagittarius, the closest Orion, and outer Perseus spiral arm segments. (The distances for the vast majority of these spiral tracers have been determined in the literature with trigonometric or photometric methods.) The Sun is situated at the inner edge of the Orion arm (Levine et al. 2006; Hou & Han 2014; Nakanishi & Sofue 2016; Xu et al. 2018, 2021; Lallement et al. 2019; Reid et al. 2019; Skowron et al. 2019; Cantat-Gaudin et al. 2020; Fig. 2 above).3 These three spatial features nearby to the Sun appear to form part of the global spiral structure in the Galaxy. Contrary, the objects of older population with larger random velocities, for instance, main-sequence A\u2013K stars or the oldest Cepheids and open clusters, do not currently follow the exact location of those arms (e.g. Cantat-Gaudin et al. 2020, fig. 8 therein; Griv et al. 2020, fig. 7 therein). The latter can be explained by the difference in rotation velocity between the spiral density waves and the objects. Investigating the velocity field of Xu et al.\u2019s (2018) O and early B-type stars in the framework of the Lin\u2013Shu density-wave proposal, we also found that the Sun lies within the Orion arm, at the inner edge of this spiral. The radial distance from the Sun to the centre of the Orion arm is \u22480.2 kpc in the direction of the Galactic anticentre, the centre of the Sagittarius arm is \u22481.8 kpc from the Sun in the direction of the GC, and the width of the arms is \u22480.5 kpc. The radial distance between the centres of the Orion and Sagittarius arms near the Sun is \u03bbrad \u2248 2 kpc (cf. Hou & Han 2014; Wu et al. 2014; Bovy et al. 2015). As for us, the nearest Orion spiral arm forms part of the dominant density-wave structure of the system.","Citation Text":["Hou & Han 2014"],"Functions Text":["The Sun is situated at the inner edge of the Orion arm"],"Functions Label":["Background"],"Citation Start End":[[614,628]],"Functions Start End":[[538,592]]}
{"Identifier":"2018MNRAS.480..927P__Richardson_&_Fairbairn_2014_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":["Richardson & Fairbairn 2014"],"Functions Text":["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","though this finding is still debated"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[778,805]],"Functions Start End":[[367,578],[740,776]]}
{"Identifier":"2019MNRAS.485.4343C__Rodriguez-Bernal_2012_Instance_1","Paragraph":"We obtain computationally credible samplings of the posterior probability (equation 8) by removing the burn-in steps of the random walk according to the autocorrelation time. We can then create synthetic data sets by drawing a parameter sample $\\pmb {\\theta }_k$ from the posterior and using it to draw from the likelihood to create a new data set, i.e. drawing new \u03c3Dj from the probability distribution for all galaxies in the original data set using equation (6). We then assess the validity of the model by comparing synthetic data with the observed (i.e. original) data. This comparison is done by using a discrepancy measure $\\mathcal {D}(\\sigma _\\mathrm{ D}|\\pmb {\\theta }_k)$ between data and model-derived expected values for the same data $e=\\lbrace e_j(\\pmb {\\theta }_k)\\rbrace$, where $\\boldsymbol{\\theta}_k$ is drawn from the posterior distribution and \u03c3D can be the observed errors or the model-generated synthetic errors. The discrepancy can be calculated using a statistic like \u03c72 (de la Horra 2008; de la Horra & Teresa Rodriguez-Bernal 2012), but here we will work with the Freeman\u2013Tukey discrepancy since it is weight independent (Brooks, Catchpole & Morgan 2000; Bishop, Fienberg & Holland 2007),\n\r\n\\begin{eqnarray*}\r\n\\mathcal {D}(\\sigma _\\mathrm{ D}|\\pmb {\\theta }_k)=\\sum _j^m \\left(\\sqrt{\\sigma _{\\mathrm{ D}j}\\vphantom{e_j(\\pmb {\\theta }_k)}}-\\sqrt{e_j(\\pmb {\\theta }_k)}\\right)^2.\r\n\\end{eqnarray*}\r\nFor each parameter draw k, it is possible to compare the simulated discrepancy with the observed discrepancy. If the model is representative of the data, then for many parameter draws, the simulated and observed discrepancies should be similar. We can then calculate a Bayesian \u2018p-value\u2019 as the ratio of \u2018draws when the observed discrepancies are larger than the synthetic discrepancies\u2019 to \u2018total draws\u2019. If this Bayesian p-value is too close to 0 or to 1 we can reject the model, otherwise it is generating synthetic data that are similar to the original data. This is better visualized using a discrepancy plot, where for each draw k, a synthetic discrepancy is paired with its corresponding observed discrepancy. If the discrepancy points are roughly equally distributed about the $\\mathcal {D}_\\mathrm{obs}=\\mathcal {D}_\\mathrm{sim}$ line, then we cannot reject the model. As mentioned above, we expect that galaxies with the largest number of measurements are sampling more completely the \u2018true\u2019 distribution of the distance. Therefore, we need to find the minimum number of measurements per galaxy for which the Bayesian p-value shows an agreement between on the partitioned data set and the model predictions.","Citation Text":["de la Horra & Teresa Rodriguez-Bernal 2012"],"Functions Text":["The discrepancy can be calculated using a statistic like \u03c72","but here we will work with the Freeman\u2013Tukey discrepancy since it is weight independent"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1015,1057]],"Functions Start End":[[936,995],[1060,1147]]}
{"Identifier":"2021MNRAS.507.4564P__Murgia_2003_Instance_1","Paragraph":"For the radio galaxies in our sample, we estimate the minimum jet power, Pj following W\u00f3jtowicz et al. (2020):\n(4)$$\\begin{eqnarray}\r\nP_{\\rm j} &amp;\\sim &amp; 1.5 \\times 10^{45} \\times \\left(\\frac{{\\rm LS}}{{\\rm 100 \\, pc}} \\right)^{9\/7} \\left(\\frac{\\tau _{\\rm j}}{{\\rm 100 \\, yr}} \\right)^{-1} \\\\\r\n&amp;&amp;\\times \\, \\left(\\frac{L_{\\rm 5 \\, GHz}}{10^{42} \\,{\\rm erg \\, s^{-1}}} \\right)^{4\/7} {\\rm erg\\, s^{-1}} ,\r\n\\end{eqnarray}$$where LS is the linear size, \u03c4j is the source age, and L$_{\\rm 5 \\, GHz}$ is the luminosity at 5 GHz. In equation (4), we consider the linear size and the luminosity at 5 GHz reported in Tables A1\u2013A3. We assume ages between about 100 yr, for the most compact sources, and 105 yr for sources with LS of several kpc, as derived from radiative and kinematic ages of sources (e.g. Murgia et al. 1999; Fanti & Fanti 2002; Murgia 2003; Polatidis & Conway 2003; Giroletti & Polatidis 2009). We end up with minimum jet powers for galaxies between 1040 and 1046 erg s\u22121. The higher values are obtained for sources at higher redshift, and for 3C\u2009346, which is among the marginally detected sources from our analysis. However, the majority of the galaxies have L$_{\\rm \\, 5 GHz} \\lt 10^{43}$ erg s\u22121 and estimated minimum jet power Pj  1044 erg s\u22121. As shown in Fig. 11, this requires a UV luminosity above 1045 erg s\u22121 in order to detect a cumulative signal by the stacking analysis. This is far from the expectation from Stawarz et al. (2008), in which the optimal conditions occurred for sources with jet power \u223c1046 erg s\u22121, LS  100 pc, and at redshift 0.2 (\u223c1 Gpc). For our estimated jet power, the highest expected gamma-ray luminosity for sources with LS  100 pc is 1044 erg s\u22121. The fact that young radio galaxies are faint emitters of gamma rays is also suggested by the results of the stacking analysis, which set the upper limit to their emission, as a whole population, an order of magnitude below the Fermi-LAT threshold. This indicates that only the closest sources could be detected by Fermi-LAT, while if we consider objects at higher and higher redshift, boosting effects are necessary for their detection.","Citation Text":["Murgia 2003"],"Functions Text":["We assume ages between about 100 yr, for the most compact sources, and 105 yr for sources with LS of several kpc, as derived from radiative and kinematic ages of sources (e.g."],"Functions Label":["Uses"],"Citation Start End":[[850,861]],"Functions Start End":[[634,809]]}
{"Identifier":"2020MNRAS.493...87T__Maraston_et_al._2013_Instance_1","Paragraph":"The other significant source of scatter in the size\u2013mass plane is the uncertainty in measuring the total stellar mass of galaxies from the integrated stellar mass density profile of the objects. As explained in Section 5.4, we measure our total stellar mass by integrating the stellar mass density profiles. To quantify how the uncertainty in the total stellar mass affects our results, we have assumed the following uncertainties in measuring the stellar mass: \u03b4mass = 0.24 \u00b1 0.01\u2009dex (for the entire sample), \u03b4mass = 0.19 \u00b1 0.01\u2009dex (for the E0-S0 + subsample), \u03b4mass = 0.24 \u00b1 0.01\u2009dex (for the S0\/a-Sm subsample) and \u03b4mass  = 0.25 \u00b1 0.03\u2009dex (for the Dwarfs subsample). These values were computed by an analysis of the differences between the Portsmouth stellar masses of our galaxies (Maraston et al. 2013) and those we measured using the g\u2013r colour profile (Roediger & Courteau 2015, see Appendix D for further details). To model the effect of the mass uncertainty ($\\sigma _{\\rm R_{mass}}$) on the scatter of the scaling relationship, all the observed stellar mass profiles were either scaled up or down in mass to place the galaxies on the best-fitting line through the observed stellar mass plane. This has been performed self-consistently, i.e. taking into account the change in the location of R1 due to the scaling of the profile. Once all the galaxies are located exactly on top of the best fitting stellar mass\u2013size relation (i.e. with zero scatter), we randomly scale the stellar mass density profiles up or down again, this time by a quantity compatible with a Gaussian distribution whose standard deviation is given by the above \u03b4mass values. We repeat this procedure 1000 times and on each occasion we measure the scatter of the stellar mass\u2013size plane produced by the uncertainty in measuring the stellar mass. We show an illustration of the scatter of the stellar mass\u2013size relation caused by the uncertainty in stellar mass in Fig. D2. The scatter in the stellar mass\u2013size plane generated by the uncertainty in mass is shown in Table 3. Interestingly, for R1, RH, and R23.5, i, we find that the dwarfs are the most affected by the uncertainty due to our mass determination. This is once again expected as the star formation activity of dwarf galaxies is, on average, more stochastic (Kauffmann 2014) and complicated to model than that of massive spirals and ellipticals. Therefore, a single colour is not a good proxy for the M\/L ratio of dwarfs as it is in the case for more gentle star formation histories.","Citation Text":["Maraston et al. 2013"],"Functions Text":["These values were computed by an analysis of the differences between the Portsmouth stellar masses of our galaxies"],"Functions Label":["Uses"],"Citation Start End":[[789,809]],"Functions Start End":[[673,787]]}
{"Identifier":"2019MNRAS.482.5430B__Eerten_&_MacFadyen_2012_Instance_1","Paragraph":"In light of this, the allowed structure of gamma-ray burst (GRB) jets and the efficiency at which it produces gamma-rays at large angles remains a topic of major importance, and it is useful to consider what types of jet structures are consistent with GRB observations (see also Beniamini et al. 2018b). Previous studies have considered the implications of structure models on the true energetics and rates of GRBs (Frail et al. 2001; Lipunov, Postnov & Prokhorov 2001; Rossi, Lazzati & Rees 2002; Zhang & M\u00e9sz\u00e1ros 2002; Eichler & Levinson 2004; van Eerten & MacFadyen 2012; Pescalli et al. 2015), on the shape of the afterglow light curve (Granot & Kumar 2003; Kumar & Granot 2003; Salmonson 2003) or on detectability of orphan afterglows (Lamb & Kobayashi 2017). Here, we propose a novel way to test the allowed structure of GRBs (in terms of both the energy and Lorentz factor angular distributions), by applying three independent techniques. We focus on long GRBs for which more detailed observations are available. First, we compare the predictions of these models regarding the EX\/E\u03b3 distribution (i.e. the isotropic equivalent early X-ray afterglow to prompt gamma-ray energy ratio) to the observations. We show that a variety of structure models predict large variations in this quantity, in contrast with results from GRB observations. Secondly, we reconsider the effect of the structure on the observed luminosity function and show that a large family of models can be ruled out as they lead to an overproduction of bursts with gamma-ray luminosities below the peak of the observed luminosity function. Both these considerations imply that while the energy angular profile may be steep, the Lorentz factor of GRBs must remain large at any region that produces gamma-rays efficiently. However, even such models typically lead to very peculiar light curves that can be ruled out by observations. The most likely implication is that efficient gamma-ray emission must be confined to a narrow opening angle around the jet\u2019s core, where the isotropic equivalent energy is not much lower than that of the core. This will naturally resolve all the problems mentioned above.","Citation Text":["van Eerten & MacFadyen 2012"],"Functions Text":["Previous studies have considered the implications of structure models on the true energetics and rates of GRBs"],"Functions Label":["Background"],"Citation Start End":[[546,573]],"Functions Start End":[[304,414]]}
{"Identifier":"2021MNRAS.505.5833F__Crocce_&_Scoccimarro_2006_Instance_1","Paragraph":"Besides the Patchy and the LN mocks, we also model the multipoles of the BOSS CMASS two-point correlation function using an analytic approach, which is required to run the Monte Carlo analysis (see Section 5). The 2PCF can be obtained from the Fourier transform of the matter power spectrum, P(k), for which we assume the template from Padmanabhan & White (2008):\n(10)$$\\begin{eqnarray*}\r\nP(k)=\\left[P_{\\rm {lin}}(k)-P_{\\rm {dw}}(k)\\right]e^{-k^2\\Sigma _{\\rm {nl}}^2\/2}+P_{\\rm {dw}}(k) .\r\n\\end{eqnarray*}$$In the equation above, Plin(k) is the linear matter power spectrum computed using the Boltzmann code CLASS (Lesgourgues 2011), assuming the Planck 2015 (Ade et al. 2016) fiducial cosmology. The Pdw(k) term is the de-wiggled power spectrum (Eisenstein & Hu 1998), while the \u03a3nl parameter encodes the smoothing of the BAO peak due to non-linear effects (Crocce & Scoccimarro 2006). The multipoles of the analytic 2PCF are defined as\n(11)$$\\begin{eqnarray*}\r\n\\xi _l(s) = \\frac{i^l}{2\\pi ^2}\\int _0^{\\infty } P_l(k)j_l(ks)k^2{\\rm d}k ,\r\n\\end{eqnarray*}$$from which we recover the monopole (l = 0) and the quadrupole (l = 2). In equation (11), jl(x) represents the spherical Bessel function of first kind and order l, while Pl(k) are the multipoles of the power spectrum defined as\n(12)$$\\begin{eqnarray*}\r\nP_l(k)=\\frac{2l+1}{2}\\int ^1_{-1}\\left(1+f\\mu ^2\\right)^2P(k)L_l(\\mu){\\rm d}\\mu ,\r\n\\end{eqnarray*}$$where Ll(x) is the Legendre polynomial of order l and P(k) is the template given in equation (10). By replacing equation (12) in equation (11), the analytic expressions for monopole (l = 0) and quadrupole (l = 2) are respectively (Xu et al. 2012):\n(13)$$\\begin{eqnarray*}\r\n\\xi _{\\rm {model}}^{(0)}(s) = B_0\\xi _0(\\alpha s)+a_0^{(0)}+\\frac{a_1^{(0)}}{s}+\\frac{a_2^{(0)}}{s^2} ,\r\n\\end{eqnarray*}$$(14)$$\\begin{eqnarray*}\r\n\\xi _{\\rm {model}}^{(2)}(s) = B_2\\xi _2(\\alpha s)+a_0^{(2)}+\\frac{a_1^{(2)}}{s}+\\frac{a_2^{(2)}}{s^2} ,\r\n\\end{eqnarray*}$$where \u03b1 is the shift parameter, while $(a_1^{(i)},a_2^{(i)},a_3^{(i)})$ are linear nuisance parameters.","Citation Text":["Crocce & Scoccimarro 2006"],"Functions Text":["while the \u03a3nl parameter encodes the smoothing of the BAO peak due to non-linear effects"],"Functions Label":["Uses"],"Citation Start End":[[858,883]],"Functions Start End":[[769,856]]}
{"Identifier":"2019ApJ...875...90L__Li_et_al._2018a_Instance_1","Paragraph":"When energy flows from the interior of the Sun outward into the solar atmosphere, why is the Sun\u2019s outer atmosphere, the corona, much hotter than the inner atmosphere, the underlying chromosphere and photosphere? This is the long-standing problem of the coronal heating, which is one of the eight key mysteries in modern astronomy (Kerr 2012). For about 80 yr since the discovery of the extremely hot corona around the late 1930s (Grotian 1939; Edlen 1945), people have worked hard on addressing this issue, and great advances have been made in observation and theoretical studies (Parnell & De Moortel 2012; Amari et al. 2015; Arregui 2015; Cargill et al. 2015; De Moortel & Browning 2015; Jess et al. 2015; Klimchuk 2015; Longcope & Tarr 2015; Peter 2015; Schmelz & Winebarger 2015; Velli et al. 2015; Wilmot-Smith 2015). Especially during recent decades, high-resolution observations of solar super-fine structures indicate that small spicules, minor hot jets along small-scale magnetic channels from the low atmosphere upwards to the corona, petty tornados and cyclones, and small explosive phenomena such as mini-filament eruptions and micro- and nano-flares\u2014all of these small-scale magnetic activities contribute greatly to coronal heating (De Pontieu et al. 2011; 2018; Zhang & Liu 2011; Parnell & De Moortel 2012; Klimchuk 2015; Peter 2015; Schmelz & Winebarger 2015; Henriques et al. 2016; Li et al. 2018a). Additionally, contributions of MHD waves to heating the corona have been observationally illustrated (van Ballegooijen et al. 2011; Jess et al. 2015; Kubo et al. 2016; Morton et al. 2016; Soler et al. 2017; Morgan & Hutton 2018). Meanwhile, with the progress of observational studies, two groups of theoretical models, magnetic reconnection models and magnetohydrodynamic wave models, have traditionally attempted to explain coronal heating, but so far no models can address the key mystery perfectly (van Ballegooijen et al. 2011; Arregui 2015; Cargill et al. 2015; Peter 2015; Velli et al. 2015; Wilmot-Smith 2015). Maybe we do not need to intentionally take to heart such the classical dichotomy, because waves and reconnections may interact with each other (De Moortel & Browning 2015; Velli et al. 2015). Additionally, statistical studies may look at coronal heating from a comprehensive perspective. Li et al. (2018b) found that the long-term variation of the heated corona, which is represented by coronal spectral irradiances, and that of small-scale magnetic activity are in lockstep, indicating that the corona should statistically be effectively heated by small-scale magnetic activity. Observational and theoretical model studies through heating channels and modes, and statistical studies by means of heating effect (performance of the heated corona), both suggest that coronal heating originates from small-scale magnetic activity.","Citation Text":["Li et al. 2018a"],"Functions Text":["Especially during recent decades, high-resolution observations of solar super-fine structures indicate that small spicules, minor hot jets along small-scale magnetic channels from the low atmosphere upwards to the corona, petty tornados and cyclones, and small explosive phenomena such as mini-filament eruptions and micro- and nano-flares\u2014all of these small-scale magnetic activities contribute greatly to coronal heating"],"Functions Label":["Background"],"Citation Start End":[[1400,1415]],"Functions Start End":[[824,1246]]}
{"Identifier":"2018ApJ...854..155K__Mei_et_al._2012_Instance_1","Paragraph":"From \u223c17:07 UT onward, especially during 17:12\u201317:14 UT, we detected multiple blobs in the bright, inverted-V-shaped structure below the flux rope, along with the fast rise of the filament (Figure 3). In Figure 12(b), boxes U and D encompass the upward- and downward-moving blobs, whose projected speeds are \u223c135 and 55 km s\u22121, respectively. Some blobs also appear to coalesce during their propagation. We attribute the growing linear features beneath the rising flux rope to plasma emission associated with a current sheet, analogous to the flare current sheet in CME\/eruptive flare models (e.g., Karpen et al. 2012). In this case, the multiple bright blobs are plasmoids formed by bursty reconnection in this current sheet, another phenomenon commonly found in high-Lundquist-number reconnection simulations (e.g., Daughton et al. 2006, 2014; Drake et al. 2006; Fermo et al. 2010; Uzdensky et al. 2010; Huang & Bhattacharjee 2012; Karpen et al. 2012; Mei et al. 2012; Cassak & Drake 2013; Guo et al. 2013; Wyper & Pontin 2014a, 2014b; Guidoni et al. 2016; Lynch et al. 2016). Multiple plasmoids moving bidirectionally were previously detected below flux ropes in active-region eruptive flares (Takasao et al. 2012; Kumar & Cho 2013; Kumar et al. 2015). If we assume a minimum base field strength of 50 G and an Alfv\u00e9n speed of \u223c135 km s\u22121 for an upward-moving plasmoid, we obtain an estimated minimum density of 4.5 \u00d7 1010 cm\u22123 for the flare current sheet. The curious appearance of the bright inverted-V-shaped structure diverging beneath the flux rope (see Figure 3 red and white arrows, Column 3, and the accompanying movie) underscores the 3D geometry of the flare current sheet. Here the right-hand bright line (marked by white arrows) appeared first (\u223c17:05 UT), followed by the left one (marked by a green arrow) at \u223c17:13 UT. The right-hand line disappeared by \u223c17:15 UT, while the left faded gradually through the rest of the observing period. A large downward-moving blob is visible during \u223c17:17\u201317:18 UT. Because current sheets are very thin, they become visible only if the line of sight passes through multiple folds or through regions of enhanced density. We speculate that the appearance of two plasmoid-generating regions could be a sign of patchy reconnection in a rippled current sheet, with reconnection sites appearing at different locations along the sheet.","Citation Text":["Mei et al. 2012"],"Functions Text":["In this case, the multiple bright blobs are plasmoids formed by bursty reconnection in this current sheet, another phenomenon commonly found in high-Lundquist-number reconnection simulations (e.g.,"],"Functions Label":["Uses"],"Citation Start End":[[953,968]],"Functions Start End":[[619,816]]}
{"Identifier":"2015ApJ...799..149J___2014_Instance_6","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"],"Functions Text":["Our estimate of the stellar mass fraction agrees","and with the estimates from strong lensing studies (see for example","which produced stellar mass fractions in the range 30%\u00e2\u0080\u009370% integrated inside the Einstein radius of the lenses."],"Functions Label":["Similarities","Similarities","Similarities"],"Citation Start End":[[3880,3898]],"Functions Start End":[[3402,3450],[3697,3764],[3900,4013]]}
{"Identifier":"2021AandA...656A..16C__Bruno_&_Carbone_2013_Instance_4","Paragraph":"Investigations of the turbulent nature of solar wind fluctuations have been ongoing for more than half a century (see, e.g., Bruno & Carbone 2016). Advances have been made consistently thanks to the increasingly accurate measurements of several dedicated space mission as well as to the enormous improvement of numerical calculation, new detailed models and theoretical frameworks, and the development of specific data analysis techniques. Nevertheless, the extremely complex nature of the system and the coexistence of multiple actors, scales, and dynamical regimes have led to a number of questions that remain open (Viall & Borovsky 2020). Among these, the very nature of the turbulent cascade of the solar wind flow and its relationship with the small-scale processes still need to be described in full (Tu & Marsch 1995; Bruno & Carbone 2013; Matthaeus & Velli 2011; Chen 2016). Magnetic field fluctuations have been characterized with great detail at magnetohydrodynamic and kinetic scales, for example, through spectral and high-order moments analysis (Tu & Marsch 1995; Bruno & Carbone 2013). The anisotropic nature of magnetic turbulence has also been addressed, and is still being debated, due to the limited access to three-dimensional measurements in space (see, e.g., Horbury et al. 2008, 2012; Sorriso-Valvo et al. 2010; Yordanova et al. 2015; Verdini et al. 2018; Telloni et al. 2019a; Oughton & Matthaeus 2020). Velocity fluctuations have been studied thoroughly (see, e.g., Sorriso-Valvo et al. 1999; Bruno & Carbone 2013), although the kinetic scales still remain quite unexplored for instrumental limitations, most notably in the sampling time resolution. Both the velocity and magnetic field show highly variable turbulence properties, with well developed spectra, strong intermittency (Sorriso-Valvo et al. 1999), anisotropy, and linear third-order moments scaling (Sorriso-Valvo et al. 2007; Carbone et al. 2011). The level of Alfv\u00e9nic fluctuations (mostly but not exclusively found in fast streams, see e.g., D\u2019Amicis et al. 2011; Bruno et al. 2019) are believed to be associated with the state of the turbulence. In particular, solar wind samples containing more Alfv\u00e9nic fluctuations are typically associated with less developed turbulence, as inferred from both shallower spectra and reduced intermittency (see Bruno & Carbone 2013, and references therein). This is consistent with the expectation that uncorrelated Alfv\u00e9nic fluctuations contribute to reduce the nonlinear cascade by sweeping away the interacting structures (Dobrowolny et al. 1980), as also confirmed by the observed anticorrelation between the turbulent energy cascade rate and the cross-helicity (Smith et al. 2009; Marino et al. 2011a,b).","Citation Text":["Bruno & Carbone 2013"],"Functions Text":["In particular, solar wind samples containing more Alfv\u00e9nic fluctuations are typically associated with less developed turbulence, as inferred from both shallower spectra and reduced intermittency (see","and references therein)."],"Functions Label":["Background","Background"],"Citation Start End":[[2337,2357]],"Functions Start End":[[2137,2336],[2359,2383]]}
{"Identifier":"2021MNRAS.508.2019B__K\u00f6rding,_Jester_&_Fender_2008_Instance_1","Paragraph":"In the context of AGN models, for SAD-dominated sources the ratio Lradio\/$L_{\\rm X} \\propto \\eta _{\\rm jet} \\, \\bar{a}^{2}$ (Yi & Boughn 1999) is expected to be constant and limited to a narrow range of values, as we find for our Seyferts, if the $\\bar{a}$ range is finite. Conversely, for LINERs, that are powered by RIAF discs, the relation between bolometric AGN and jet power is difficult to model because jet emission can dominate their entire SED (K\u00f6rding, Jester & Fender 2008) and [O\u2009iii]-line contamination from jet photoionizing shocks can overpredict the bolometric AGN power (Capetti, Verdoes Kleijn & Chiaberge 2005; Netzer 2009). However, in a simplistic scenario of a disc origin of the radio emission, ADAF discs would predict much shallower relations ($L_{\\rm core} \\propto L_{\\rm X}^{\\alpha }$ with \u03b1 0.6 depending on bremsstrahlung-dominated and multiple Compton scattering regimes, Yi & Boughn 1998) than the observed slopes for LINERs (1.2\u20131.3). Furthermore, the observed radio-[O\u2009iii] relation found for LINERs is even steeper (despite the large scatter) than any previous relations between the kinetic jet power and the bolometric AGN luminosities found for LINER-like RL AGN (slopes 1, e.g. Capetti & Balmaverde 2006; Merloni & Heinz 2007; Baldi et al. 2019b). Therefore, assuming that RL LINERs are the scaled-down version of RGs and normalizing the Lradio\/LX comparison with previous studies based on this assumption, the steeper slopes might be the consequence of a closer view of jet-launching mechanism even in RQ LINERs. The high sensitivity and sub-arcsec resolution (crucially intermediate between VLA and VLBI) provided by our survey allowed us to probe the parsec-scale region near the core where the jet is launched and reveal the relevant role of \u0393jet and BH spin\/mass in the jet production. These parameters could eventually induce much steeper relations than those established with shallower radio observations.","Citation Text":["K\u00f6rding, Jester & Fender 2008"],"Functions Text":["Conversely, for LINERs, that are powered by RIAF discs, the relation between bolometric AGN and jet power is difficult to model because jet emission can dominate their entire SED"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[454,483]],"Functions Start End":[[274,452]]}
{"Identifier":"2020MNRAS.498.1480S__Bharadwaj_&_Sethi_2001_Instance_1","Paragraph":"After the recombination epoch, the CMB hardly interacts with the neutral intervening medium. This restricts the CMB from probing the evolution of the structures till the end of EoR. The 21-cm radiation, which is involved in the hyperfine transition of H\u2009i, is a promising probe to study the high-redshift universe including EoR (e.g. Sunyaev & Zeldovich 1972; Hogan & Rees 1979). There are existing and the upcoming radio interferometers aiming to observe the brightness temperature fluctuations of the redshifted 21-cm signal from EoR that we coin as the \u2018EoR 21-cm signal\u2019. However, the detection of the signal is not yet possible due to the foreground contamination from galactic and extra-galactic source. The foregrounds are \u223c104\u2212105 times stronger (e.g. Ali, Bharadwaj & Chengalur 2008; Bernardi et al. 2009, 2010; Ghosh et al. 2012; Paciga et al. 2013; Beardsley et al. 2016) compared to the signal. The foregrounds, system noise, and calibration errors together keep the current observations at bay from directly detecting the EoR 21-cm signal. As a consequence, the first detection is likely to be statistical in nature. These observations plan to measure the power spectrum (PS) of the EoR 21-cm signal (e.g. Bharadwaj & Sethi 2001; Bharadwaj & Ali 2004, 2005). Several radio interferometers such as the GMRT1 (Swarup et al. 1991), LOFAR2 (van Haarlem et al. 2013), MWA3 (Tingay et al. 2013), and PAPER4 (Parsons et al. 2010) have carried out observations to measure the EoR 21-cm PS. However, only few weak upper limits on the PS amplitudes have been reported in the literature to date (e.g. GMRT: Paciga et al. 2011, Paciga et al. 2013; LOFAR: Yatawatta et al. 2013; Patil et al. 2017; Gehlot et al. 2019; Mertens et al. 2020; MWA: Dillon et al. 2014; Jacobs et al. 2016; Barry et al. 2019; Li et al. 2019; Trott et al. 2020; PAPER: Cheng et al. 2018; Kolopanis et al. 2019). A few more upcoming telescopes with improved sensitivity such as HERA5 (DeBoer et al. 2017) and SKA6 (Koopmans et al. 2014) also aim to measure the EoR 21-cm PS. Apart from PS, several other estimators such as the variance (Patil et al. 2014), bispectrum (Bharadwaj & Pandey 2005; Yoshiura et al. 2015; Shimabukuro et al. 2017; Majumdar et al. 2018), and Minkowski functional (Bag et al. 2018, 2019; Kapahtia et al. 2018; Kapahtia, Chingangbam & Appleby 2019) are being used to quantify the EoR 21-cm signal. These estimators are supposed to be rich in information about the underlying physical processes during EoR.","Citation Text":["Bharadwaj & Sethi 2001"],"Functions Text":["The foregrounds, system noise, and calibration errors together keep the current observations at bay from directly detecting the EoR 21-cm signal. As a consequence, the first detection is likely to be statistical in nature. These observations plan to measure the power spectrum (PS) of the EoR 21-cm signal (e.g."],"Functions Label":["Motivation"],"Citation Start End":[[1219,1241]],"Functions Start End":[[907,1218]]}
{"Identifier":"2015AandA...584A..75V__Essen_et_al._(2014)_Instance_3","Paragraph":"The data presented here comprise quasi-simultaneous observations during secondary eclipse of WASP-33 b around the V and Y bands. The predicted planet-star flux ratio in the V-band is 0.2 ppt, four times lower than the accuracy of our measurements. Therefore, we can neglect the planet imprint and use this band to measure the stellar pulsations, and most specifically to tune their current phases (see phase shifts in von Essen et al. 2014). Particularly, our model for the light contribution of the stellar pulsations consists of eight sinusoidal pulsation frequencies with corresponding amplitudes and phases. Hence, to reduce the number of 24 free parameters and the values they can take, we use prior knowledge about the pulsation spectrum of the star that was acquired during von Essen et al. (2014). As the frequency resolution is 1\/\u0394T (Kurtz 1983), 3.5 h of data are not sufficient to determine the pulsations frequencies. Therefore, during our fitting procedure we use the frequencies determined in von Essen et al. (2014) as starting values plus their derived errors as Gaussian priors. As pointed out in von Essen et al. (2014), we found clear evidences of pulsation phase variability with a maximum change of 2 \u00d7 10-3 c\/d. In other words, as an example after one year time a phase-constant model would appear to have the correct shape with respect to the pulsation pattern of the star, but shifted several minutes in time. To account for this, the eight phases were considered as fitting parameters. The von Essen et al. (2014) photometric follow-up started in August, 2010, and ended in October, 2012, coinciding with these LBT data. We then used the phases determined in von Essen et al. (2014) during our last observing season as starting values, and we restricted them to the limiting cases derived in Sect. 3.5 of von Essen et al. (2014), rather than allowing them to take arbitrary values. The pulsation amplitudes in \u03b4 Scuti stars are expected to be wavelength-dependent (see e.g. Daszy\u0144ska-Daszkiewicz 2008). Our follow-up campaign and these data were acquired in the blue wavelength range. Therefore the amplitudes estimated in von Essen et al. (2014), listed in Table 1, are used in this work as fixed values. This approach would be incorrect if the pulsation amplitudes would be significantly variable (see e.g., Breger et al. 2005). Nonetheless, the short time span of LBT data, and the achieved photometric precision compared to the intrinsically low values of WASP-33\u2019s amplitudes, make the detection of any amplitude variability impossible. ","Citation Text":["von Essen et al. (2014)"],"Functions Text":["Therefore, during our fitting procedure we use the frequencies determined in","as starting values plus their derived errors as Gaussian priors"],"Functions Label":["Uses","Uses"],"Citation Start End":[[1007,1030]],"Functions Start End":[[930,1006],[1031,1094]]}
{"Identifier":"2019MNRAS.490.5353M__Chen_&_Podsiadlowski_2017_Instance_1","Paragraph":"An independent approach to constrain the strength of magnetic fields in accretion flow models in binary systems is by measuring how much angular momentum is lost from the system via magnetized outflows and estimating the corresponding orbital decay in the system due to magnetic braking. The orbital decay in A0620-00 is rapid, with orbital-period derivative $\\dot{P} =-0.6 \\, {\\rm [ms \\, yr^{-1}]}$ (Gonz\u00e1lez Hern\u00e1ndez, Rebolo & Casares 2014), and it cannot be explained by the emission of gravitational waves alone. Magnetic braking of the system is a possible explanation for the measured $\\dot{P}$, but other explanations, such as resonant interactions between the binary and the possible circumbinary disc, have been proposed (Chen & Podsiadlowski 2017). Here we can estimate the magnitude of the magnetic braking of the system using a first-principles approach. The orbital angular momentum of a binary system is $J_{\\rm orb}= M_{\\rm BH}M_*\/(M_{\\rm BH}+M_*) \\sqrt{G(M_{\\rm BH}+M_*) d}$, where d is the separation between the star and the black hole. In GRMHD simulations, the loss of total angular momentum through the outer boundary is defined as $\\dot{J}(r_{\\rm out},t)=\\int _\\theta \\int _\\phi T^r_\\phi \\,\\mathrm{ d}A_{\\theta \\phi }$, where $T^r_\\phi \\equiv (\\rho + \\gamma u + b^2) u^r u_\\phi -b^r b_\\phi$ is the stress\u2013energy tensor describing the radial flux of angular momentum. The quantity u is the internal energy of the gas, \u03b3 = 4\/3 is the adiabatic index, u\u03bc is the four-velocity of the gas and b\u03bc is a four-vector that describes the magnetic field in a frame comoving with the gas. We integrate the above formula at rout = 50GM\/c2 over \u03b8 \u2208 (0, \u03c0) and \u03d5 \u2208 (0, 2\u03c0). In our best-fitting model, the ratio of angular momentum flux through the outer boundary to the orbital angular momentum is extremely low: $\\dot{J}\/J_{\\rm orb} \\approx 10^{-22}$ [s\u22121], which could account for an orbital period derivative of $\\dot{P} = 2 \\times 10^{-8} {\\rm [ms \\, yr^{-1}]}$ only. Our calculations confirm that the inner highly magnetized, rotating accretion flow (with current $\\dot{M}$) alone cannot be responsible for the rapid orbital decay observed in the system and this favours the idea of a circumbinary disc.","Citation Text":["Chen & Podsiadlowski 2017"],"Functions Text":["Magnetic braking of the system is a possible explanation for the measured $\\dot{P}$, but other explanations, such as resonant interactions between the binary and the possible circumbinary disc, have been proposed"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[732,757]],"Functions Start End":[[518,730]]}
{"Identifier":"2020MNRAS.492.1295P__Bonning_et_al._2012_Instance_2","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"],"Functions Text":["while \u2018achromatic\u2019 flux variability (no colour evolution,","have also been reported."],"Functions Label":["Background","Background"],"Citation Start End":[[1164,1183]],"Functions Start End":[[1086,1143],[1252,1276]]}
{"Identifier":"2019AandA...630A.131M__Uttley_et_al._2014_Instance_1","Paragraph":"Comptonisation Monte Carlo code (MoCA; see Tamborra et al. 2018 for a detailed description of the code) is based on a single photon approach, working in a fully special relativistic scenario. MoCA allows for various and different physical and geometrical conditions of the accretion disc and of the Comptonising corona. In this paper, the corona is assumed to have either a spherical or a slab-like geometry, and to be as extended as the disc, whose radii have been set to be Rout\u2004=\u2004500 rg and Rin\u2004=\u20046 rg, respectively. Even though arguments (e.g. variability, Uttley et al. 2014, and references therein, microlensing Chartas et al. 2009; Morgan et al. 2012 and timing Kara et al. 2016; De Marco et al. 2013) exist that favour a compact corona, we used extended coronae. In fact, as discussed by Marinucci et al. (2019), Comptonised spectra emerging from compact corona (Rout\u2004=\u2004100 rg\u2013Rin\u2004=\u20046) do not deviate significantly from those produced in more extended corona; see their Fig. 3. The adoption of even more compact coronae (Rout\u2004=\u200420 rg, Rin\u2004=\u20046) results only in the need for higher optical depths to recover the same spectral shape for a given temperature. However, in such cases, general relativity (GR) effects are not negligible (see Tamborra et al. 2018, for a detailed discussion on this topic), and the present version of MoCA does not include GR. For the slab-like geometry case, MoCA allows the user to set up the corona height above the accretion disc (set to 10\u2006rg in our simulations). We use synthetic spectra computed assuming the source BH mass and accretion rate to be the same as those of Ark 120 (e.g. Marinucci et al. 2019, and references therein), namely MBH\u2004=\u20041.5\u2005\u00d7\u2005108\u2006M\u2299 and \u1e41 = Lbol\/LEdd = 0.1. For both the slab and spherical hot electron configurations, we simulated the Comptonised spectra using a wide range of values for electron temperature and optical depth: 0.1\u2004 \u2004\u03c4\u2004 \u20047 and 20  kT  200 keV, and in Fig. 1 we show a sample of spectra obtained by MoCA. Moreover, spectra are computed from 0.01 keV up to 700 keV using 1000 logarithmic energy bins, and a Poissonian error accompanies each spectral point. The obtained spectra are averaged over the inclination angle and in Fig. 1 we show some exemplificative spectra normalised at 1 keV accounting for the two geometries considered in this work.","Citation Text":["Uttley et al. 2014"],"Functions Text":["Even though arguments (e.g. variability,","exist that favour a compact corona, we used extended coronae."],"Functions Label":["Differences","Differences"],"Citation Start End":[[561,579]],"Functions Start End":[[520,560],[709,770]]}
{"Identifier":"2021AandA...656A..44R__S\u00e9rsic_1968_Instance_1","Paragraph":"The detection of LSBGs is usually carried out by a sequence of tasks. The first step is a broad detection of LSBG candidates. Both the depth of the data and the efficiency of this detection in the low-surface-brightness regime will define the completeness of the sample, and so a high efficiency in the detection of diffuse sources in this first step is a key point in the process. Given the importance of this primary detection, it is common to use specialized software or procedures (e.g., Akhlaghi & Ichikawa 2015; Prole et al. 2018; Haigh et al. 2021). After a first detection of potential objects, a characterization of their structural and morphological properties is necessary in order to accurately define the sample, typically with certain criteria in surface brightness and radius. To obtain the structural parameters of the LSBG candidates, the detected sources are usually fitted with a S\u00e9rsic model (S\u00e9rsic 1968). Specific procedures such as GALFIT (Peng et al. 2010) or IMFIT (Erwin 2015) are used to obtain accurate morphological and structural parameters. However, the computational cost of these is very high, becoming a bottleneck in LSBG analysis pipelines. For this reason, it is advisable to minimize the number of LSBG candidates prior to their S\u00e9rsic fitting to save computational time. This is not trivial because structural parameters from automated segmentation catalogs are much less accurate than S\u00e9rsic fit modeling, making correct screening of the sample challenging. As a last step, a \u201csupervision\u201d of all the objects matching the criteria of the sample is necessary in order to discard the frequent presence of false positives. These are sources that, although meeting the imposed criteria, are clearly not LSBGs. Examples of false positives are clumped sources that have been considered as a single source, faint reflections due to the optics of the instrumentation, or artifacts present in the images. Depending on the data volume analyzed, the number of objects to be supervised can become very high. For this reason, it is increasingly common to use deep learning techniques to screen objects in order to eliminate these false-positive detections (e.g., Tuccillo et al. 2018; Burke et al. 2019; Tanoglidis et al. 2021a). However, unsupervised or automated detections can not yet reach the accuracy that human visual inspection is capable of obtaining, reaching almost 100%. Therefore, human supervision is desirable, when possible, in order to minimize the presence of false positives in samples of reasonably limited volume, such as that used here.","Citation Text":["S\u00e9rsic 1968"],"Functions Text":["To obtain the structural parameters of the LSBG candidates, the detected sources are usually fitted with a S\u00e9rsic model","Specific procedures such as GALFIT","or IMFIT","are used to obtain accurate morphological and structural parameters. However, the computational cost of these is very high, becoming a bottleneck in LSBG analysis pipelines. For this reason, it is advisable to minimize the number of LSBG candidates prior to their S\u00e9rsic fitting to save computational time. This is not trivial because structural parameters from automated segmentation catalogs are much less accurate than S\u00e9rsic fit modeling, making correct screening of the sample challenging. As a last step, a \u201csupervision\u201d of all the objects matching the criteria of the sample is necessary in order to discard the frequent presence of false positives. These are sources that, although meeting the imposed criteria, are clearly not LSBGs. Examples of false positives are clumped sources that have been considered as a single source, faint reflections due to the optics of the instrumentation, or artifacts present in the images. Depending on the data volume analyzed, the number of objects to be supervised can become very high."],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[913,924]],"Functions Start End":[[792,911],[927,961],[981,989],[1003,2035]]}
{"Identifier":"2021ApJ...919...30D__Staguhn_et_al._2014_Instance_1","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"],"Functions Text":["Other single-dish samples of SMGs have also been obtained at 1.1\u20131.3 mm using","and at 2 mm with GISMO"],"Functions Label":["Background","Background"],"Citation Start End":[[606,625]],"Functions Start End":[[324,401],[582,604]]}
{"Identifier":"2018AandA...620A..84B__Hua_et_al._1998_Instance_1","Paragraph":"The PN IC 5148 (PN G002.7-52.4) is a nebula which is to date not well investigated in detail. First listed in the Second Index Catalogue of Nebulae and Clusters of Stars (Dreyer 1910) with two independent entries as numbers IC 5148 and IC 5150 discovered by Swift (1899) and Gale (1897) independently, it was finally discovered to be the same object by Hoffmeister (1961). Morphologically it is declared as a round nebula in all catalogs and Chu et al. (1987) classified it as a multiple shell planetary nebula (MSPN) due to a small step in the H \u03b1 image with a radius ratio of only 1:1.2 between the two structures. In addition, the intensity decrease was found smaller than that for typical MSPNe. A few years later the authors searched with a larger field systematically around many nebulae for extended emission features without detecting the very low surface brightness halo we investigate here (Hua et al. 1998). While earlier spectroscopic studies took only a small fraction of one or two spectral lines to obtain radial velocity and expansion of the nebula (e.g. Meatheringham et al. 1988), to our knowledge up to now the only, more extended spectroscopic analysis were performed by Kaler et al. (1990) and by Kingsburgh & Barlow (1994) within two surveys of 75 and 80 southern PNe, respectively. Both studies used pre-CCD electronic spectral scanning devices, taking a very small aperture region of the nebula and detected only a hand full of lines. They end up coinciding in the result that the nebula has about galactic disk abundance or only slight underabundance despite its large galactic latitude (b approximately \u221252\u00b0). Further, thesurvey spectra used in the PN catalog of Acker et al. (1992) gives intensities of only four lines. The inspection of the original data file of this survey, provided now at the Hong Kong\/AAO\/Strasbourg PN data base (thereafter HASH2; Parker et al. 2016; Boji\u010di\u0107 et al. 2017), does not recover more usable lines above the noise level.","Citation Text":["Hua et al. 1998"],"Functions Text":["A few years later the authors searched with a larger field systematically around many nebulae for extended emission features without detecting the very low surface brightness halo we investigate here"],"Functions Label":["Background"],"Citation Start End":[[901,916]],"Functions Start End":[[700,899]]}
{"Identifier":"2015ApJ...811L..32H__Matthews_1994_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":["Matthews 1994"],"Functions Text":["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"],"Functions Label":["Uses"],"Citation Start End":[[941,954]],"Functions Start End":[[777,939]]}
{"Identifier":"2021MNRAS.502.4858S__Yung_et_al._2020_Instance_1","Paragraph":"One of our main long-term goals is to work towards a full forward modelling pipeline for multiwavelength galaxy surveys. Over the next decade, wide area surveys from DESI, VRO, Euclid, the Nancy Grace Roman Space Telescope, 4MOST, and other facilities will be carried out. We can use the legacy observations from surveys such as CANDELS, to build a foundation for interpreting these new surveys. What we have shown here is that the current generation of SAMs produce decent broad agreement with key properties of galaxy evolution as represented by CANDELS over the redshift range 0.5 \u2272 z \u2272 3. It has been shown elsewhere that these models produce similar results to those of numerical cosmological simulations and other SAMs (Somerville & Dav\u00e9 2015), and that they are also in agreement with higher redshift observations of galaxy populations (Yung et al. 2019a, b), the reionization history (Yung et al. 2020), and observational probes of the cold gas phase in galaxies (Popping, Somerville & Trager 2014; Popping et al. 2019). While there are certainly remaining tensions with observations, as seen here and also in, e.g. Popping et al. (2019), there is promising ongoing work to continue to improve the realism of the treatment of physical processes in SAMs (e.g. Pandya et al. 2020). In work in progress, we are using this framework to create similar mock observations for future planned surveys with the James Webb Space Telescope and the Nancy Grace Roman Space Telescope (L.Y.A., Yung et al., in preparation). SAMs coupled with light-cones extracted from large volume N-body simulations have recently been used to create a 2\u2009deg2 light-cone from 0  z  10 (Yang et al. 2020; Yung et al., in preparation). In order to create mock surveys for even larger areas \u2013 tens to hundreds of square degrees \u2013 that will be probed by the projects mentioned above, it is likely that new, even more computationally efficient techniques will need to be developed, perhaps enabled by machine-learning-based tools.","Citation Text":["Yung et al. 2020"],"Functions Text":["It has been shown elsewhere that these models produce similar results to those of numerical cosmological simulations and other SAMs","and that they are also in agreement with higher redshift observations of galaxy populations","the reionization history"],"Functions Label":["Similarities","Similarities","Similarities"],"Citation Start End":[[893,909]],"Functions Start End":[[593,724],[751,842],[867,891]]}
{"Identifier":"2019ApJ...882...65M__Kennicutt_et_al._2003_Instance_1","Paragraph":"Given this correlation between dust luminosity and SFR, it is reasonable to expect the dust mass to scale with SFR as well. In the top panel of Figure 10 we show this relation for our data with the additional parameter of the dust temperature indicated by color. This dust temperature is a weighted average of the temperatures of the two components of the dust model, namely the ISM and birth clouds. Using MAGPHYS, da Cunha et al. (2010) found a tight correlation between dust mass and SFR in a local sample derived from the Sloan Digital Sky Survey, which we show as a black line. The da Cunha et al. (2010) sample is selected to be star-forming by emission line diagnostics and lies at z \u2264 0.2. It is also important to note that the SFRs for their sample lie mostly below \u223c20 M\u2299 yr\u22121. The solid portion of the line represents the extent of their data. We also show the Spitzer Infrared Nearby Galaxy Sample (Kennicutt et al. 2003) analyzed in da Cunha et al. (2008) as purple stars for comparison. At first, our data do not appear to follow any trend, with only our highest measured dust masses falling near the da Cunha et al. (2010) relation. Given that MAGPHYS allows a range of dust temperatures when fitting for the dust mass, we investigate the effects of dust temperature in this diagram. When we examine the temperature of the dust (which MAGPHYS allows to vary from 20 to 80 K) as indicated by color in the figure, we see that sources with lower dust temperature actually lie on the published relation. If we consider bands of constant dust temperature in the diagram we see that the slope of the Mdust\u2013SFR relation for our sources actually closely matches the slope of the da Cunha et al. (2010) relation. As shown in Figures 6 and 7, we observe a slight increase in average dust temperature with redshift for our sample, so it may be tempting to interpret the variation in this diagram as a redshift dependence. We checked the relation between dust temperature and redshift for our sample and found the scatter in temperature at a given redshift to be much larger than the evolution between redshift bins. This suggests that the local calibration between dust mass and SFR can be extended to higher redshift samples through the incorporation of the additional parameter of dust temperature. Given our particular sample selection, the application of these results to a more general star-forming population would have to be done with caution.","Citation Text":["Kennicutt et al. 2003"],"Functions Text":["We also show the Spitzer Infrared Nearby Galaxy Sample","analyzed in da Cunha et al. (2008) as purple stars for comparison."],"Functions Label":["Uses","Uses"],"Citation Start End":[[911,932]],"Functions Start End":[[855,909],[934,1000]]}
{"Identifier":"2015ApJ...806..199B__Balser_et_al._2011_Instance_1","Paragraph":"The most prominent metallicity structure in the Milky Way disk is the decrease of metallicity with increasing Galactic radius\u2014the radial metallicity gradient. All of the major tracers reveal radial gradients, typically with slopes between \u22120.03 and \u22120.09 \n\n\n\n\n\n. The radial metallicity gradient can be explained by an inside-out galaxy formation where the disk grows via a radially dependent gas infall rate and star formation rate (e.g., Matteucci & Fran\u00e7oi 1989). But why is there such a wide range of radial gradient slopes measured? There are several possibilities. (1) Measurement uncertainty. Measurement errors will cause some variations, but it is unlikely to produce a factor of 3 difference in slope. Homogeneity in observing procedures and data analysis may be more important given the variations in abundance that can be derived for the same source (for further discussion see Rudolph et al. 2006; Henry et al. 2010; Balser et al. 2011). (2) Dynamical evolution. Radial gradients calculated with stellar tracers may be affected by radial migration, where stars are scattered into different orbits (Sellwood & Binney 2002). Radial migration should flatten the radial metallicity gradient, but there is some evidence that this may not be a large factor for stars in the Milky Way (Di Matteo et al. 2013; Kubryk et al. 2013; Bovy et al. 2014). (3) Temporal evolution. Many tracers have a wide range of age and therefore are probing the Milky Way disk at different times. For example, the radial gradient has been observed to flatten with time when using Open clusters (e.g., Friel et al. 2002). But accurate ages are crucial to separate out these temporal effects. PN studies indicate a flattening of the radial gradient with time (Maciel et al. 2003), a steepening with time (Stanghellini & Haywood 2010), or no temporal variation (Henry et al. 2010). These vastly different conclusions reveal the difficulty in deriving accurate PN ages and distances. (4) Sample Evolution. Hayden et al. (2013) measure a flattening of the radial gradient away from the Galactic mid-plane. Therefore, including sources out of the Galactic plane, which may be from an older population, may alter the derived radial gradient slope (see Minchev et al. 2014). Azimuthal abundance variations have been reported (e.g., Pedicelli et al. 2009; Balser et al. 2011, Section 3.3). If real, these will complicate any analysis of radial gradients. Here we detect metallicity gradient slopes over different azimuth ranges that span the values observed in the literature. The usual assumption is that the disk is well mixed at a given radius but this may not be true.","Citation Text":["Balser et al. 2011"],"Functions Text":["But why is there such a wide range of radial gradient slopes measured? There are several possibilities. (1) Measurement uncertainty. Measurement errors will cause some variations, but it is unlikely to produce a factor of 3 difference in slope. Homogeneity in observing procedures and data analysis may be more important given the variations in abundance that can be derived for the same source (for further discussion see"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[929,947]],"Functions Start End":[[466,888]]}
{"Identifier":"2019ApJ...871..243Y__Saito_et_al._1999_Instance_1","Paragraph":"There are two possibilities resulting in the different magnetic field strengths inferred from the polarimetric and molecular-line observations: (1) the rotational-to-gravitational energy \u03b2rot is overestimated, and (2) there are additional contributions in the polarized intensity from other mechanisms, such as dust scattering. In our MHD simulations, \u03b2rot is adopted to be 0.4% based on the observational estimates of the core mass of \u223c1 M and the angular speed of the core rotation of 4 \u00d7 10\u221214 s\u22121. The angular speed was estimated based on the global velocity gradient along the major axis of the dense core observed with single-dish telescopes (Saito et al. 1999; Yen et al. 2011; Kurono et al. 2013). Numerical simulations of dense cores including synthetic observations show that the specific angular momentum derived from the synthetic images of the dense cores can be a factor of 8\u201310 higher than their actual specific angular momentum computed by the sum of the angular momenta contributed by the individual gas parcels in the dense cores (Dib et al. 2010). In addition, if there are filamentary structures in the dense core in B335, which could not be resolved with the single-dish observations, infalling motions along the filamentary structures could also contribute to the observed velocity gradient, leading to an overestimated angular speed of the core rotation (Tobin et al. 2012). We have also performed our simulations with a lower \u03b2rot, and we find that the rotational velocity on a 100 au scale in the simulations decreases with decreasing \u03b2rot. Thus, the discrepancy in the magnetic field strengths inferred from the field structures and the gas kinematics can be reconciled, if the core rotation in B335 is overestimated by a factor of a few in the observations, and these results would suggest a weak magnetic field of initial \u03bb of 9.6 in B335. Further observations combining single dishes and interferometers to have a high spatial dynamical range and to map the velocity structures of the entire dense core in B335 at a high angular resolution are needed to study coherent velocity features and provide a better estimate of the core rotation.","Citation Text":["Saito et al. 1999"],"Functions Text":["The angular speed was estimated based on the global velocity gradient along the major axis of the dense core observed with single-dish telescopes"],"Functions Label":["Uses"],"Citation Start End":[[649,666]],"Functions Start End":[[502,647]]}
{"Identifier":"2015AandA...575A.111D__Pinsonneault_et_al._(2001)_Instance_1","Paragraph":"The difference in iron abundance between the two XO-2 stellar components poses an interesting question. They belong to a visual binary and, as normally assumed for such systems, they should share the same origin and initial bulk metallicity. A relevant characteristic of this system is that both of the stars host planets. Thus, for components of wide binaries where at least one star has a planet, a reasonable hypothesis to explain any measured and significant difference in their present-day elemental abundances is that the planet formation process had played a relevant role. The higher iron abundance of XO-2N when compared to XO-2S might be due to past ingestion of dust-rich or rocky material that came from the inner part of the proto-planetary disk and that was pushed into the host star by the hot Jupiter XO-2N as it migrated inward to its current orbit. A second mechanism, which acts on a different time scale and after the pre-main-sequence stage, is presented by Fabrycky & Tremaine (2007), who discuss the pollution of the stellar photosphere with metals produced by mass loss of inward migrating hot Jupiters. In addition, Kaib et al. (2013) showed with their simulations that very distant binary companions may severely affect planetary evolution and influence the orbits of any planet around the other component of the system, perhaps favouring the ingestion of material by the host star. Following the results shown in Fig. 2 of Pinsonneault et al. (2001), we roughly estimated that XO-2N could have ingested an amount of iron slightly higher than 5 M\u2295 to increase its photospheric iron content by ~0.05 dex, given its Teff ~ 5300 K. The difference in the iron abundance of the two XO-2 stellar components is similar to that found by Ram\u00edrez et al. (2011) for the solar twins 16 Cyg A and 16 Cyg B (\u0394 [Fe\/H] = 0.042  \u00b1  0.016 dex). These stars have masses close to those of the XO-2 companions and, together with a third companion 16 Cyg C, are members of a hierarchical triple system, where the A and C components form a close binary with a projected separation ~70 AU. The component B is known to host a giant planet in a long-period and highly eccentric orbit (P ~ 800 days and e ~ 0.69)5, but unlike the case of XO-2 it is the star with a lower iron content. To explain this deficiency, Ram\u00edrez et al. (2011) suggest that an early depletion of metals happened during the formation phase of the 16 Cyg Bb planet. An alternative and suggestive possibility is that 16 Cyg A may have ingested a massive planet that enriched the star with iron, while the large orbit of 16 Cyg Bb (with semi-major axis of 1.68 AU) very likely prevented any mass loss from the planet. We note that, although \u0394[Fe\/H] is similar to that found for XO-2, the amount of iron involved should be different. The XO-2 stars are cooler than 16 Cygni A and B and therefore have more massive convective envelopes, implying that more iron should be [is?] necessary to pollute the XO-2N photosphere and produce almost the same \u0394[Fe\/H]. Laws & Gonzalez (2001) also found the primary of the 16-Cyg system enhanced in Fe relative to the secondary. Similar studies conducted in binary stars hosting planets (see e.g. Gratton et al. 2001; Desidera et al. 2004, 2006; Schuler et al. 2011a; Teske et al. 2013; Liu et al. 2014; Mack et al. 2014) did not find relevant differences in elemental abundance among the components. All these results imply that the presence of giant planets does not necessarily imply differences in the chemical composition of the host star. ","Citation Text":["Pinsonneault et al. (2001)"],"Functions Text":["Following the results shown in Fig. 2 of","we roughly estimated that XO-2N could have ingested an amount of iron slightly higher than 5 M\u2295 to increase its photospheric iron content by ~0.05 dex, given its Teff ~ 5300 K."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1450,1476]],"Functions Start End":[[1409,1449],[1478,1654]]}
{"Identifier":"2017AandA...607A..85L__Zeipel_1924_Instance_1","Paragraph":"The LCs analyses were made using the phoebe v.0.29d software (Pr\u0161a & Zwitter 2005) that is based on the 2003 version of the Wilson-Devinney (W-D) code (Wilson & Devinney 1971; Wilson 1979, 1990). In the absence of spectroscopic mass ratios, the \u201cq-search\u201d method (for details see Liakos & Niarchos 2012) was applied in modes 2 (detached system), 4 (semi-detached system with the primary component filling its Roche lobe) and 5 (conventional semi-detached binary) to find feasible (\u201cphotometric\u201d) estimates of the mass ratio. The step of q change during the search was 0.05\u20130.1 starting from q = 0.05\u22120.1. The effective temperatures of the primaries (T1) were given the values derived from the spectral classification (see Sect. 2) and were kept fixed during the analysis, while the temperatures of the secondaries T2 were adjusted. The Albedos, A1 and A2, and gravity darkening coefficients, g1 and g2, were set to generally adopted values for the given spectral types of the components (Rucinski 1969; von Zeipel 1924; Lucy 1967). The (linear) limb darkening coefficients, x1 and x2, were taken from the tables of van Hamme (1993). The dimensionless potentials \u03a91 and \u03a92, the fractional luminosity of the primary component L1, and the orbital inclination of the system i were set as adjustable parameters. At this point it should to be noted that since the Kepler\u2019s photometer has a spectral response range between approximately 410\u2013910 nm with a peak at ~ 588 nm, the R filter (Bessell photometric system-range between 550\u2013870 nm and with a transmittance peak at 597 nm) was selected as the best representative for the filter depended parameters (i.e. x and L). Moreover, there is evidence of maxima brightness changes in all of the systems, therefore parameters of photospheric spots on the surface of the secondary were also adjusted. The selection of the magnetically active component was based on the effective temperatures of the members of the systems. In all cases the secondaries are clearly cooler than the primaries (i.e. large minima difference), therefore, they host a convective envelope that better suits a magnetically active star. In addition, the hotter stars are candidates for \u03b4 Sct type pulsations and it is rather rare to present magnetic activity also. For all EBs, except for KIC 111, the third light parameter (l3) was also adjusted because the systems are candidates for triplicity (see Sect. 1). However, during the iterations it resulted in unrealistic values, therefore, it was omitted in the final analysis. Finally, all systems were found to have the minimum \u2211 res2 in mode 5. KIC 066 and KIC 111 have a minimum at q = 0.3, while KIC 105 and KIC 106 at q = 0.15. In Fig. 7 the respective q-search plots are shown. ","Citation Text":["von Zeipel 1924"],"Functions Text":["The Albedos, A1 and A2, and gravity darkening coefficients, g1 and g2, were set to generally adopted values for the given spectral types of the components"],"Functions Label":["Uses"],"Citation Start End":[[1003,1018]],"Functions Start End":[[832,986]]}
{"Identifier":"2015ApJ...800...72T__Kroupa_1995_Instance_1","Paragraph":"The binarity or binary fraction, f, is defined as the ratio of the number of binary or higher-order systems, Nbin, to the total number of systems, Nsys. Here, the term system includes multiple systems and singles (their number being noted as Nsng) as well. Then\n6For the star-like population, we apply the binary DPS method developed by Marks et\u00c2 al. (2011) and Marks & Kroupa (2011), hereafter referred to as dynamical or DPS pairing. In DPS the binary stars are formed in a population of embedded clusters, within which they are dynamically processed to yield the Galactic disk stellar single-plus-binary population. An attractive feature of this theory is its underlying assumption of the universality of binary properties of late-type stars being consistent with observational data (Marks & Kroupa 2012; Leigh et\u00c2 al. 2014). For initial binaries with intermediate to large separations, the DPS pairing method applies random pairing11M. Marks et\u00c2 al. (2014, in preparation) show that random pairing with subsequent dynamical processing does indeed reproduce at least the observed low-mass stellar population (see also Kroupa 1995).\nbelow a primary mass of 5\u00e2\u0080\u0089M and ordered pairing (such that q \u00e2\u00a9\u00be 0.9) above. Here, q = mcomp\/mprim \u00e2\u00a9\u00bd 1, where mcomp is the companion mass and mprim is the mass of the primary star. This initial binary population is then altered by dynamical evolution. Close binaries with orbital periods below about 10\u00c2 days undergo eigenevolution (Kroupa 1995) and tend to equalize the companion masses. Note that this eigenevolution term alters the very-low-mass end of the star-like IMF. For the purpose of this work, however, these effects only play a negligible role. Here, the initial or primordial binary fraction is 100%; that is, it is assumed that all stars form in binaries. The final (after dynamical processing in the embedded cluster) overall binary fraction is about 40% (i.e., f = 0.4) but varies as a function of the primary-star mass. For M dwarfs, in particular, it is as low as 25%, whereas G dwarfs show about 56% binarity. The binary fraction approaches 90% for O stars. For the BD-like population, we chose an overall binary fraction of 20% (i.e., f = 0.20), in accordance with TK07 and TK08. About half of the members of observed average stellar populations are binaries, most of them remaining unresolved in typical star-cluster surveys. However, very young and likely dynamically unevolved populations like the Taurus\u00e2\u0080\u0093Auriga association exhibit almost 100% binarity (Kroupa et\u00c2 al. 2013; Duch\u00c3\u00aane & Kraus 2013; Reipurth et\u00c2 al. 2014). The number of systems must not be confused with the number of individual bodies: Nbod = Nsng + 2Nbin. Because higher-order multiples are relatively rare (Goodwin & Kroupa 2005), they are summarized within the binary population in this work, so the total number of bodies is\n7The CMRD describes the relative number of binaries as a function of the companion-to-primary mass ratio. Observations reveal a continuous decline of f as a function of the primary-object mass, which has been interpreted as a continuous transition from the stellar to the substellar regime (Joergens 2008; Kraus & Hillenbrand 2012; but see Thies & Kroupa 2008). There is also a shift toward more equal-mass binaries (q = 1) for VLMSs and BDs (Dieterich et\u00c2 al. 2012). These properties of the stellar population are well reproduced by DPS such that the origin and properties of binary populations are well understood.","Citation Text":["Kroupa 1995","Kroupa 1995"],"Functions Text":["Marks et\u00c2 al. (2014, in preparation) show that random pairing with subsequent dynamical processing does indeed reproduce at least the observed low-mass stellar population (see also","This initial binary population is then altered by dynamical evolution. Close binaries with orbital periods below about 10\u00c2 days undergo eigenevolution","and tend to equalize the companion masses. Note that this eigenevolution term alters the very-low-mass end of the star-like IMF. For the purpose of this work, however, these effects only play a negligible role."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[1121,1132],[1472,1483]],"Functions Start End":[[940,1120],[1320,1470],[1485,1695]]}
{"Identifier":"2016ApJ...822...72C__Mandrini_et_al._2005_Instance_1","Paragraph":"Here, for the first time, we have identified highly dynamic non-potential activity on QS-like supergranular network scales. These events overlie mixed polarity network flux elements near the spatial resolution of HMI, and are the first non-potential structures to be unassociated with strong concentrations of bipolar magnetic flux. One event (2011b August 05) shows eruptive activity in the form of jets, which is similar to larger-scale micro-sigmoids (Raouafi et al. 2010) and even QS bright point sigmoids (Chesny et al. 2015). The existence of flaring non-potential fields in the QS-like mixed network immediately shows that supergranular-scale magnetic fields can support processes similar to sigmoid formation (Chesny et al. 2013). Strong non-potential field arcades have been observed in hot X-ray sigmoids on scales of hundreds of arseconds (Moore et al. 2001; Gibson et al. 2002; Liu et al. 2010), micro-sigmoids in soft X-ray imaging on scales of \u223c50\u2033 (Mandrini et al. 2005; Raouafi et al. 2010) and small-scale AR EUV (Zheng et al. 2012, 2013), and now at EUV temperatures in QS-like mixed network fields with lengths down to \u223c10\u2033. This range of lengths over a range of temperatures and magnetic field topologies points directly to self-similar mechanisms influencing plasma and magnetic field dynamics at a range of scales. Our findings suggest that the processes driving some large-scale eruptions (i.e., flux emergence, helicity build-up, and flux cancellation leading to non-potential field heating (Chen et al. 2014)), can also manifest in a range of configurations on sub-network size scales in QS-like magnetic field configurations. These QS flaring non-potential fields are similar to their large-scale counterparts, but not as strict in their evolution. The diversity in the observed topologies may scale with the diversity of magnetic configurations that exist in the supergranulation network. QS non-potential fields evolve in multi-polar environments, and are not restricted to strongly bipolar dominated regions as in larger-scale, higher temperature events. Despite this, one of the presented events (2011a August 05) results in a post-flare potential loop arcade, which is similar to some observed AR sigmoid fields (Moore et al. 2001).","Citation Text":["Mandrini et al. 2005"],"Functions Text":["Strong non-potential field arcades have been observed in","micro-sigmoids in soft X-ray imaging on scales of \u223c50\u2033","and now at EUV temperatures in QS-like mixed network fields with lengths down to \u223c10\u2033. This range of lengths over a range of temperatures and magnetic field topologies points directly to self-similar mechanisms influencing plasma and magnetic field dynamics at a range of scales."],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[964,984]],"Functions Start End":[[739,795],[908,962],[1057,1336]]}
{"Identifier":"2019MNRAS.482.3757B__Rankin_1990_Instance_1","Paragraph":"Subpulse drifting has been a topic of intensive research with the phenomenon expected to be present in 40\u201350 per\u2009cent of the pulsar population. There are around 120 pulsars known at present to exhibit some form of periodic modulations in their single-pulse sequence (Weltevrede et al. 2006, 2007; Basu et al. 2016). The drifting effects are very diverse and associations are seen with other phenomena like mode changing and nulling in some cases (Wright & Fowler 1981; Deich et al. 1986; Hankins & Wolszczan 1987; Vivekanand & Joshi 1997; Redman et al. 2005; Basu & Mitra 2018b). However, to gain a deeper understanding of the phenomenon at large, more comprehensive studies involving classification of the drifting population and identification of underlying traits are essential. The first such attempt was undertaken by Rankin (1986) who estimated the drifting behaviour within the empirical core\u2013cone model of the emission beam. The pulsar profile, which is formed after averaging several thousand single pulses, has a highly stable structure and is made up of one or more (typically 5) components. The components can be classified into two distinct categories: the central core component and the adjacent conal components. A number of detailed studies suggest that the radio emission beam consists of a central core emission surrounded by conal emission arranged in nested rings (Rankin 1990, 1993; Mitra & Deshpande 1999). The differences in the observed profiles are a result of the different line of sight (LOS) traverses across the emission beam. Rankin (1986) suggested subpulse drifting to be primarily a conal phenomenon, and the drifting to arise due to the circulation of the conal emission around the magnetic axis. This would result in an association between the drifting properties (particularly the phase variation of the drifting feature across the pulse) and the profile classification. The most prominent drifting with clear drift bands and large phase variations was associated with the conal single (Sd) and barely resolved conal double (D) profile classes. This corresponded to the LOS traversing the emission beam towards the outer edge. Progressively more interior LOS traverses of the emission beam result in well-resolved conal double (D), conal Triple (cT), and conal Quadruple (cQ) profile shapes. In the above classification scheme, such profile classes were expected to show primarily longitude stationary drift with very little phase variations. The core-dominated profiles were associated with central LOS traverses of the emission beam. The principal profile classes were categorized as core single (St), Triple (T) with a central core component and pair of conal outriders, and Multiple (M) with central core and two pairs of conal components. Subpulse drifting was expected to be phase stationary and only seen in the conal components of the T and M class profiles. However, the core components sometimes show longer periodic structures that were not classified as subpulse drifting.","Citation Text":["Rankin 1990"],"Functions Text":["A number of detailed studies suggest that the radio emission beam consists of a central core emission surrounded by conal emission arranged in nested rings"],"Functions Label":["Background"],"Citation Start End":[[1385,1396]],"Functions Start End":[[1228,1383]]}
{"Identifier":"2020ApJ...905..111Z__Jiri\u010dka_et_al._2001_Instance_1","Paragraph":"Surveys of radio bursts in decimetric wavelengths is presented in papers by Isliker & Benz (1994) and Jiri\u010dka et al. (2001), within 1\u20133 GHz and 0.8\u20132.0 GHz frequency ranges, respectively. Some of these bursts are still not well understood. This is a case of the slowly positively drifting bursts (SPDBs). They appear in groups or as single bursts, with a duration of an individual burst from 1 to several seconds and their frequency drift is lower than about 100 MHz s\u22121 (Jiri\u010dka et al. 2001). The SPDBs seem to be similar to the reverse type III bursts (Aschwanden 2002) but their frequency drift is much smaller. The majority of observed SPDBs are connected to solar flares (Jiri\u010dka et al. 2001), and they appear many times at the very beginning of the flares (Benz & Simnett 1986; Kotr\u010d et al. 1999; Kaltman et al. 2000; Karlick\u00fd et al. 2018). Kaltman et al. (2000) reported on several SPDBs observed during three solar flares in the 0.8\u20132 GHz frequency range. They found frequency drifts of the observed SPDBs to be within the 20\u2013180 MHz s\u22121 range. Kotr\u010d et al. (1999) studied one of those flares. By combining the radio and spectral plus imaging H\u03b1 observations, they explained the observed SPDBs as radio emission generated by downwards propagating shock waves. Based on numerical simulations of the formation of thermal fronts in solar flares, Karlick\u00fd (2015) proposed that SPDBs observed in the 1\u20132 GHz range could be a signature of a thermal front. Furthermore, Karlick\u00fd et al. (2018) reported the observation of an SPDB (1.3\u20132.0 GHz) observed during the impulsive phase of an eruptive flare. They found time coincidence between the SPDB occurrence, an appearance of an ultraviolet (UV)\/EUV multithermal plasma blob moving down along the dark H\u03b1 loop at approximately 280 km s\u22121, and the observed change of H\u03b1 profile at the footpoint of that dark loop. Combining these observations they concluded that observed SPDB was likely generated by the thermal front formed in front of the falling EUV blob.","Citation Text":["Jiri\u010dka et al. (2001)"],"Functions Text":["Surveys of radio bursts in decimetric wavelengths is presented in papers by Isliker & Benz (1994) and","within 1\u20133 GHz and 0.8\u20132.0 GHz frequency ranges, respectively."],"Functions Label":["Background","Background"],"Citation Start End":[[102,123]],"Functions Start End":[[0,101],[125,187]]}
{"Identifier":"2022MNRAS.515.4430M__Iucci_et_al._1979b_Instance_1","Paragraph":"Forbush decreases (FD) are the results of influence of solar wind (SW) large-scale disturbances on the background cosmic ray (CR) flux. They often demonstrate relatively fast CR intensity decrease, accompanied by large and extremely variable CR anisotropy, which is followed by a slower recovery (Forbush 1937; Lockwood 1971; Iucci et al. 1979a; Belov 2009). Most FDs have a sporadic character and are induced by interplanetary disturbances such as Interplanetary Coronal Mass Ejections (ICMEs) caused by coronal mass ejections (CMEs) (Cane 2000; Gopalswamy 2010a; Richardson & Cane 2011a). FDs caused by high-speed streams (HSS) from coronal holes (CHs) have a recurrent character (Iucci et al. 1979b; Richardson 2004; Singh & Badruddin 2007b). The influence of different types of SW disturbances on galactic CR has been well-documented in scientific literature (e.g. Dumbovic et al. 2012; Chertok et al. 2013; Kumar & Baddruddin 2014a,b; Melkumyan et al., 2018, 2019). It is known that most of CMEs, and almost all energetic CMEs, are originated from the active regions being accompanied by solar flares which intensity correlated with the CME kinetic energy (e.g. Yashiro & Gopalswamy 2009). Comparison between the impact on the heliosphere of active-region (AR) and non-AR CMEs is presented in Gopalswamy et al. (2010b). The authors reveal that: (i) active regions produce almost all CMEs with above-average energy, for example geoeffective CMEs, CMEs associated with solar energetic particles, shock-driving CMEs; (ii) the quiescent filament regions produce the other type of CMEs which are not related to the sunspots and occur mostly at high latitudes during the maximum phase of solar activity. One of the factors responsible for the variability of FDs is different origin of related interplanetary transients. In Marii et al. (2020, 2021) the influences of different parts of ICMEs (turbulent sheath and magnetic obstacle) on the associated FDs are examined according to the CME origin. The events are separated into three subsets: AR CMEs, disappearing filament CMEs, and stealthy CMEs. CR variations caused by turbulent sheath and magnetic obstacle parts of the corresponding ICME are compared for the three subsets and two different linear relations are found: for AR ICMEs on one side and for filament\/stealthy ICMEs on the other side.","Citation Text":["Iucci et al. 1979b"],"Functions Text":["FDs caused by high-speed streams (HSS) from coronal holes (CHs) have a recurrent character"],"Functions Label":["Background"],"Citation Start End":[[683,701]],"Functions Start End":[[591,681]]}
{"Identifier":"2020MNRAS.492.3420B__Hildebrandt_et_al._2018_Instance_1","Paragraph":"Given the potential huge information content of these observables advocated by some recent works (Patton et al. 2017), the predictions from first principles developed in this paper could be successfully applied to forthcoming data along with the standard power spectrum based analysis and could bring additional information beyond \u039bCDM parameters like massive neutrinos (Liu & Madhavacheril 2019) or dark energy (Codis et al. 2016a). Let us stress that implementing nulling in weak-lensing analysis is central in order to avoid extracting biased information from the small scales that lack a full theoretical understanding (including due to the effect of baryon physics that needs to be modelled in weak-lensing surveys Hildebrandt et al. 2018; Yoon et al. 2019). Not only this general nulling technique should be used for one-point statistics but could also be applied to standard power spectrum analysis (and more generally to the full two-point PDF) in order to disentangle the effects of the different physical scales. However, more realistic effects have to be accounted for before the here mentioned formalism could be directly applied to real data. In particular, we have not investigated the precise impact of the galaxy redshift distribution ns(z) (for which one needs to go from a set of discrete source planes to a source distribution), photometric redshift errors or shape noise, which are left for future works. Promising extensions include an application of the formalism to (i) compensated filters such as for aperture mass which require the joint modelling of the field at two different scales, (ii) two-point statistics in order to model cosmic variance, and (iii) the joint analysis of multiple redshift bins. All of these ideas are within reach of LDT as was shown in the case of the three-dimensional matter density in Bernardeau, Codis & Pichon (2015) and Codis et al. (2016b) for, respectively, the multiscale and two-point statistics.","Citation Text":["Hildebrandt et al. 2018"],"Functions Text":["Let us stress that implementing nulling in weak-lensing analysis is central in order to avoid extracting biased information from the small scales that lack a full theoretical understanding (including due to the effect of baryon physics that needs to be modelled in weak-lensing surveys"],"Functions Label":["Future Work"],"Citation Start End":[[720,743]],"Functions Start End":[[434,719]]}
{"Identifier":"2019MNRAS.484.1645O__Hummel_et_al._1991_Instance_1","Paragraph":"We justify our use of the Two-Zone approximation as follows: we anticipate that the magnetic structure within the outflow would be perpendicular to the plane of the host galaxy and thus perpendicular to the magnetic field orientation within the Zone B ISM region. Indeed, such perpendicular magnetic structure in outflows is seen in simulation work where, e.g. the action of a CR-driven dynamo yields a perpendicular magnetic field configuration compared to the host galactic plane (Kulpa-Dybe\u0142 et al. 2011), or by the advection of the magnetic fields by the flows themselves (Bertone et al. 2005), by magnetic amplification via the CR streaming instability (Uhlig et al. 2012) along the outflow. This magnetic structure would also be consistent with polarised radio synchrotron emission above and below the planes of galaxies known to host outflows in the nearby Universe, with the polarisation direction aligned with the orientation of the outflow cone (see, e.g. Hummel et al. 1988; Sukumar & Allen 1990, 1991; Hummel et al. 1991; Brandenburg et al. 1993; Chy\u017cy et al. 2006; Soida et al. 2011; Mora & Krause 2013). We argue that the principle mechanism for CRs to permeate the Zone A\/Zone B interface would be via diffusion. With magnetic field lines aligned in a direction parallel to the inter-zone boundary, diffusion across the interface would be severely hampered \u2013 the cross-boundary diffusion coefficient would effectively be perpendicular the the local magnetic field lines, and so would be around two orders of magnitude smaller than that along the field directions (e.g. Shalchi et al. 2004, 2006; Hussein & Shalchi 2014), and substantially less than the effective ISM diffusion coefficient. The detailed substructure of the magnetic fields in these interfacing regions is not yet fully understood (Veilleux et al. 2005), but we argue that our prescription is consistent with existing work on relevant scales and that adopting an alternative model for CR transport across this boundary at this point would not imply an interpretation that is any more physical than that adopted here. We acknowledge that, in future studies, it will be critical to assess the magnetic fields in these interfacing regions across a range of length-scales to properly determine the permeability of the Zone A\/Zone B interface to diffusing CRs.","Citation Text":["Hummel et al. 1991"],"Functions Text":["This magnetic structure would also be consistent with polarised radio synchrotron emission above and below the planes of galaxies known to host outflows in the nearby Universe, with the polarisation direction aligned with the orientation of the outflow cone (see, e.g."],"Functions Label":["Similarities"],"Citation Start End":[[1014,1032]],"Functions Start End":[[697,965]]}
{"Identifier":"2017MNRAS.469.3322G__Narayan_&_Yi_1994_Instance_1","Paragraph":"The induction equation of the magnetic field is\n(10)\r\n\\begin{equation}\r\n\\frac{\\partial \\boldsymbol {B}}{\\partial t}={\\nabla }\\times \\left(\\boldsymbol {v}\\times \\boldsymbol {B}-\\frac{4\\pi }{c}\\eta \\boldsymbol {J}\\right),\r\n\\end{equation}\r\nwhere $\\boldsymbol {J}=\\frac{c}{4\\pi }{\\nabla }\\times \\boldsymbol {B}$ is the current density. The induction equation is the field escaping\/creating rate due to magnetic instability or the dynamo effect. For the steady-state accretion flow, we neglect the dynamo effect. In the energy equation, we assumed a balance between the heating due to viscosity and cooling due to advection, convection and radiation as\n(11)\r\n\\begin{equation}\r\nQ^{-}_{{\\rm adv}}+Q^{-}_{{\\rm rad}}+Q^{-}_{{\\rm con}}=Q^{+}_{{\\rm diss}},\r\n\\end{equation}\r\nwhere $Q^{-}_{{\\rm adv}}=\\scriptstyle\\rho (\\frac{{\\rm d} e}{{\\rm d}t}-\\frac{p}{\\rho ^{2}}\\frac{{\\rm d}\\rho }{{\\rm d}t})$ is advection cooling in ADAFs. e is the gas internal energy $(e=\\frac{c^{2}_{{\\rm s}}}{\\gamma -1})$ and \u03b3 is the specific heat ratio. Also, we consider $Q^{+}_{{\\rm diss}}-Q^{-}_{{\\rm rad}}=fQ^{+}_{{\\rm diss}}$ in ADAFs (where f, the advection parameter, is defined by Narayan & Yi 1994). Also, the viscous heating is defined as $Q_{{\\rm diss}}= f(\\nu +\\frac{1}{3}\\nu _{{\\rm con}})\\Sigma r^{2} (\\frac{{\\rm d} \\Omega }{{\\rm d}r})^{2}$. The outward energy flow by convection is $Q^{-}_{{\\rm con}}=- {\\nabla }\\cdot \\boldsymbol {F}_{c}$, where $F_{{\\rm con}}=-\\nu _{{\\rm con}} \\frac{1}{\\gamma -1}\\frac{{\\rm d} c^{2}_{{\\rm s}}}{{\\rm d}r}-\\frac{c^{2}_{{\\rm s}}}{\\rho }\\frac{{\\rm d} \\rho }{{\\rm d}r}$. So, the energy equation is written by considering the above definitions and the specific internal energy of inflow \u03b5 and outflow \u03b5w on the surface of the disc as follows:\n(12)\r\n\\begin{eqnarray}\r\n&amp;&amp;{\\frac{\\Sigma v_{r}}{\\gamma -1}\\frac{{\\rm d} c^{2}_{{\\rm s}}}{{\\rm d}r}-2Hc^{2}_{{\\rm s}}v_{r}\\frac{{\\rm d}\\rho }{{\\rm d}r}+\\frac{1}{2\\pi r}\\frac{{\\rm d}\\dot{M}(r)}{{\\rm d}r}(\\epsilon _{w}-\\epsilon )}\\nonumber \\\\\r\n&amp;&amp;-\\frac{1}{r}\\frac{{\\rm d}}{{\\rm d}r}\\left(r\\nu _{{\\rm con}}\\frac{\\Sigma }{\\gamma -1}\\frac{{\\rm d} c^{2}_{{\\rm s}}}{{\\rm d}r}-2Hr\\nu _{m}c^{2}_{{\\rm s}}\\frac{{\\rm d}\\rho }{{\\rm d}r}\\right)\\nonumber \\\\\r\n&amp;&amp; -\\frac{\\nu \\rho }{\\gamma -1}\\frac{{\\rm d} c^{2}_{{\\rm s}}}{{\\rm d}r}+\\nu _{{\\rm con}}c^{2}_{{\\rm s}}\\frac{{\\rm d}\\rho }{{\\rm d}r} = f\\Sigma \\left(\\alpha -\\frac{1}{3}\\alpha _{\\rm c}\\right) c_{{\\rm s}}H r^{2}\\left(\\frac{{\\rm d} \\Omega }{{\\rm d}r}\\right)^{2}. \\nonumber\\\\\r\n\\end{eqnarray}\r\n","Citation Text":["Narayan & Yi 1994"],"Functions Text":["Also, we consider $Q^{+}_{{\\rm diss}}-Q^{-}_{{\\rm rad}}=fQ^{+}_{{\\rm diss}}$ in ADAFs (where f, the advection parameter, is defined by"],"Functions Label":["Uses"],"Citation Start End":[[1153,1170]],"Functions Start End":[[1018,1152]]}
{"Identifier":"2020AandA...639A..46B__\u0160tver\u00e1k_et_al._(2009)_Instance_1","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"],"Functions Text":["The linear relationship that we observe between breakpoint energy and core temperature is in line with previous measurements (e.g.","for both the halo and strahl."],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[152,171]],"Functions Start End":[[0,130],[174,203]]}
{"Identifier":"2018MNRAS.479.4509R__Kingma_&_Ba_2014_Instance_1","Paragraph":"After each step of calculations, the network should optimize the model based on its current and previous states to improve the subsequent mapping. Our model utilizes a computationally memory efficient optimization due to its dependence to only the first-order gradients, namely the \u2018adaptive moment estimation\u2019 (or Adam). For more details, we refer the readers to Kingma & Ba (2014). Adam optimization, compared to other gradient-based optimization, is very suitable for noisy and sparse gradients, and for simulated data that show very large scatter with respect to a given quantity of parameter (Kingma & Ba 2014). With this optimizer, we have to decide few parameters in advance. The learning step \u03b1 and the parameters controlling the moving averages of the first- and second-order moments, namely \u03b21 and \u03b22 (both \u2208[0,1)), respectively. For this purpose, we chose to minimize the MSE between the target and the prediction from the model: in what follows, we will alternatively call the MSE the \u2018objective function\u2019 f($\\bf x$): with ${\\bf x}$ the parameters of the model to be updated, such as weights and biases. At a given time t \u2264 T, where T is the maximal learning time-step, we can update the parameters of the model as shown in the following:\n(11)\r\n\\begin{eqnarray*}\r\ng_t &amp;=\\nabla _\\mathrm{ \\text{$x$}} f(\\mathrm{\\text{$x$}}_{t-1}),\r\n\\end{eqnarray*}\r\n(12)\r\n\\begin{eqnarray*}\r\n\\mu _{1,t} &amp;=\\beta _1 \\times \\mu _{1,t-1} + (1-\\beta _1)\\times g_t,\r\n\\end{eqnarray*}\r\n(13)\r\n\\begin{eqnarray*}\r\n\\bar{\\mu }_{1,t} &amp;=\\mu _{1,t}\/(1-\\beta _1^t),\r\n\\end{eqnarray*}\r\n(14)\r\n\\begin{eqnarray*}\r\n\\mu _{2,t} &amp;=\\beta _2 \\times \\mu _{2,t-1} + (1-\\beta _2)\\times g_t^2,\r\n\\end{eqnarray*}\r\n(15)\r\n\\begin{eqnarray*}\r\n\\bar{\\mu }_{2,t} &amp;=\\mu _{2,t}\/(1-\\beta _2^t),\r\n\\end{eqnarray*}\r\n(16)\r\n\\begin{eqnarray*}\r\n\\mathrm{\\text{$x$}}_t &amp;=\\mathrm{\\text{$x$}}_{t-1} - \\alpha _t \\times \\bar{\\mu }_{1,t}\/ (\\sqrt{\\bar{\\mu }_{2,t}} + \\epsilon),\r\n\\end{eqnarray*}\r\nwhere $\\alpha _t=\\alpha \\sqrt{1-\\beta _2^t}\/(1-\\beta _1^t)$ is the time-step at t. Equation (11) shows the gradients of the objective function at t with respect to the model parameters. Equations (12) and (14) update the estimations of the first and second moments. Our moments are biased towards the initial values; thus, we require equations (13) and (15) to account for the corrections. Finally, we update the model parameters with equation (16).","Citation Text":["Kingma & Ba (2014)"],"Functions Text":["For more details, we refer the readers to"],"Functions Label":["Background"],"Citation Start End":[[364,382]],"Functions Start End":[[322,363]]}
{"Identifier":"2021MNRAS.504..146V__Vink_&_Gr\u00e4fener_2012_Instance_1","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"],"Functions Text":["While we know that very massive stars (VMS) above 100\u2009\u2009M\u2299 exist","this mass is significantly diminished via stellar winds already during core hydrogen (H) burning"],"Functions Label":["Background","Background"],"Citation Start End":[[955,975]],"Functions Start End":[[751,814],[857,953]]}
{"Identifier":"2019AandA...622A.146M__Arribas_et_al._(2014)_Instance_1","Paragraph":"Previous works (e.g. Holt et al. 2011; Arribas et al. 2014; Villar Mart\u00edn et al. 2014, 2015) have found very high reddening and densities associated with ionised outflows in local objects (e.g. H\u03b1\/H\u03b2\u2004\u223c\u20044.91 and ne\u2004\u2273\u20041000 cm\u22123, Villar Mart\u00edn et al. 2014). Concerning the reddening, although we find that the outflowing gas is generally less affected by dust extinction than the disc, the median value of the total distribution is significantly affected by dust (H\u03b1\/H\u03b2\u2004\u223c\u20044.16), with tails up to H\u03b1\/H\u03b2\u2004\u2273\u20046. Similarly, the outflow density of MAGNUM galaxies is higher than the values in the disc gas, but appears to be far lower than the values found by these authors. This could stem from the fact that the galaxies studied by Holt et al. (2011), Arribas et al. (2014) are local luminous or ultra-luminous infrared galaxies (U\/LIRGs), and those of Villar Mart\u00edn et al. (2014, 2015) are highly obscured Seyfert 2, thus sampling sources that are more gas and dust rich compared to our sample. However, our values are also lower than the outflow densities found in Perna et al. (2017) (ne\u2004\u223c\u20041200 cm\u22123), who targeted optically selected AGNs from the SDSS, and F\u00f6rster Schreiber et al. (2018b), who presented a census of ionised gas outflows in high-z AGN with the KMOS3D survey (ne\u2004\u223c\u20041000 cm\u22123). A possible explanation could be related to the high quality of our MUSE data, which also allows us to detect the faint [S\u202fII] emission associated with lower density regions. If we calculate the median densities of the disc and outflow components, weighting for the [S\u202fII] line flux, we obtain higher values (ne\u2004\u223c\u2004170 cm\u22123 and ne\u2004\u223c\u2004815 cm\u22123, for disc and outflow, respectively). This shows that previous outflow density values from the literature could be biased towards higher ne because they are based only on the most luminous outflowing regions, characterised by a higher S\/N. This could also mean that outflows at high-z could be far more extended than the values we observe.","Citation Text":["Arribas et al. 2014"],"Functions Text":["Previous works","have found very high reddening and densities associated with ionised outflows in local objects"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[39,58]],"Functions Start End":[[0,14],[93,187]]}
{"Identifier":"2019AandA...628L...2C__Cukanovaite_et_al._(2018)_Instance_1","Paragraph":"As a final note, we point out that correcting for the effect of neutrinos did not change the finding that KIC 08626021 has a very thin helium-pure layer on top of its envelope. Such a thin layer might be explained by the presence of any mechanism competing against the separation of He, C, and O in the envelopes of DO and hot DB white dwarfs. For instance, Fontaine & Brassard (2005; see also Brassard et al. 2007) postulated the existence of a weak stellar wind of the order of \u223c2\u2005\u00d7\u200510\u221213M\u2299\u2006yr\u22121 in order to account for the presence of traces of C in Teff\u2004\u223c\u200420\u2006000\u201330 000 K DB white dwarfs. Such a wind, similar to the solar wind, would be driven by deposition of acoustic energy generated by strong convective motions which characterize the outer layers of these stars. Perhaps a more likely mechanism is convective overshooting, which was first studied by Freytag et al. (1996) in a white dwarf context and more recently by Tremblay et al. (2015) and Cukanovaite et al. (2018) through detailed 3D simulations. Brassard & Fontaine (2015) explicitly used the overshooting prescription of Freytag et al. to demonstrate that element separation is indeed considerably slowed down in cooling hot white dwarfs. They found that only a \u201cthin\u201d layer of pure He (logq1\u2004\u223c\u2004\u22128) has accumulated at the surface of their evolving model by the time the latter has cooled down to Teff \u223c 31 000 K. This is in line with our seismic inference for KIC 08626021. In addition, the new solution for KIC 08626021 does not change the fact that no C\/O\/He triple transition at the core boundary is found in the model that best reproduces the pulsation frequencies of this star. If puzzling from an evolutionary perspective, our seismic probing seemingly requires that mode trapping generated by two chemical transitions around logq\u2004\u223c\u2004\u22122.4 and logq\u2004\u223c\u2004\u22123.5 must be present. This can be obtained with our current models only if a carbon buffer exists between the envelope and the core. Alternately, enforcing a triple transition and searching for a new optimal solution within this additional constraint strongly degrades the quality of the achieved best frequency match. Improving the capacity of such a configuration to reproduce the observed pulsation modes would likely require the presence of an additional structure (of unclear origin) in the envelope to provide the needed mode trapping. We therefore leave that issue open at this stage.","Citation Text":["Cukanovaite et al. (2018)"],"Functions Text":["Perhaps a more likely mechanism is convective overshooting, which was first studied by Freytag et al. (1996) in a white dwarf context and more recently by Tremblay et al. (2015) and","through detailed 3D simulations."],"Functions Label":["Background","Background"],"Citation Start End":[[955,980]],"Functions Start End":[[773,954],[981,1013]]}
{"Identifier":"2018AandA...618A.145O__Brouillet_et_al._2013_Instance_1","Paragraph":"Such a chemical differentiation has been reported in other multiple systems, like the low-mass protostars IRAS16293-2422 (2004a; 2016; 2011) and IRAS4A (Santangelo et al. 2015; L\u00f3pez-Sepulcre et al. 2017), and, recently, towards the intermediate-mass protostars NGC 2264 CMM3 (Watanabe et al. 2017). With four examples at hand, we speculate that it could be a general feature of multiple protostellar systems and not a \u201cpathological anomaly\u201d. There is no systematic trend between the millimetre thermal dust and molecular line emissions. Towards IRAS16293-2422 and IRAS4A, a rich content in COMs is observed towards the source with the less massive continuum source. Towards NGC 2264 CMM3, it is the most massive continuum component that displays a rich molecular content. The case of Cep E-mm appears similar to the latter one. High-mass star forming regions (HMSFRs) also present a rich chemical diversity. One of the best known examples is provided by Orion-KL. This source harbours: (i) a dichotomy between the spatial distribution of complex O-bearing and complex N-bearing species, with the latter species probing the hotter gas (see, e.g. Gu\u00e9lin et al. 2008; Favre et al. 2011; Friedel & Widicus Weaver 2012; Peng et al. 2013; Brouillet et al. 2013; Crockett et al. 2014, 2015) but also (ii) differences between supposed chemically related species (see, e.g. Favre et al. 2017; Pagani et al. 2017). Other examples are provided by W3(OH) (Qin et al. 2015; Nishimura et al. 2017) and SgrB2 (Belloche et al. 2008, 2013). In an observational study of four HMSFRs (Orion KL, G29.96, IRAS 23151+5912, and IRAS 05358+3543), Beuther et al. (2009) showed that the properties of CH3OH can be easily accounted for by the physical conditions (temperature) in the cores, whereas the N-bearing species appear to be more selective as they are detected only towards the sources at the (evolved) hot core stage. Recently, in an ALMA study of the filamentary HMSFR G35.20, S\u00e1nchez-Monge et al. (2014) found that only three out of the six continuum cores of the filament display COM emission typical of hot cores. Several hypotheses have been proposed to account for the observed chemical differentiation. L\u00f3pez-Sepulcre et al. (2017) proposed that the COM-rich protostar is either more massive and\/or subject to a higher accretion rate, resulting in a lower envelope mass. Watanabe et al. (2017) suggest that the less massive protostar is related to a younger evolutionary stage in which the hot corino (hot core) is not yet developed, meaning that its dimensions are still very small. In Cep E-mm, the presence of high-velocity SiO jets provides evidence of active mass ejection around both protostars. The short dynamical timescales (500\u20131000 yr) also indicate that these ejections began recently, meaning that both sources are still in an early evolutionary stage. Incidentally, Lykke et al. (2015) found an apparent correlation between the source luminosities and the relative abundance of complex organic molecules in a sample of sources including high-mass protostars. The authors have suggested that this could be the result of the timescale and the temperature experienced by a source during its evolution. The sample of sources with evidence for chemical differentiation should be increased in order to confirm this observational trend.","Citation Text":["Brouillet et al. 2013"],"Functions Text":["High-mass star forming regions (HMSFRs) also present a rich chemical diversity. One of the best known examples is provided by Orion-KL. This source harbours: (i) a dichotomy between the spatial distribution of complex O-bearing and complex N-bearing species, with the latter species probing the hotter gas (see,"],"Functions Label":["Background"],"Citation Start End":[[1234,1255]],"Functions Start End":[[829,1140]]}
{"Identifier":"2015AandA...584A..75V__Essen_et_al._(2014)_Instance_5","Paragraph":"The data presented here comprise quasi-simultaneous observations during secondary eclipse of WASP-33 b around the V and Y bands. The predicted planet-star flux ratio in the V-band is 0.2 ppt, four times lower than the accuracy of our measurements. Therefore, we can neglect the planet imprint and use this band to measure the stellar pulsations, and most specifically to tune their current phases (see phase shifts in von Essen et al. 2014). Particularly, our model for the light contribution of the stellar pulsations consists of eight sinusoidal pulsation frequencies with corresponding amplitudes and phases. Hence, to reduce the number of 24 free parameters and the values they can take, we use prior knowledge about the pulsation spectrum of the star that was acquired during von Essen et al. (2014). As the frequency resolution is 1\/\u0394T (Kurtz 1983), 3.5 h of data are not sufficient to determine the pulsations frequencies. Therefore, during our fitting procedure we use the frequencies determined in von Essen et al. (2014) as starting values plus their derived errors as Gaussian priors. As pointed out in von Essen et al. (2014), we found clear evidences of pulsation phase variability with a maximum change of 2 \u00d7 10-3 c\/d. In other words, as an example after one year time a phase-constant model would appear to have the correct shape with respect to the pulsation pattern of the star, but shifted several minutes in time. To account for this, the eight phases were considered as fitting parameters. The von Essen et al. (2014) photometric follow-up started in August, 2010, and ended in October, 2012, coinciding with these LBT data. We then used the phases determined in von Essen et al. (2014) during our last observing season as starting values, and we restricted them to the limiting cases derived in Sect. 3.5 of von Essen et al. (2014), rather than allowing them to take arbitrary values. The pulsation amplitudes in \u03b4 Scuti stars are expected to be wavelength-dependent (see e.g. Daszy\u0144ska-Daszkiewicz 2008). Our follow-up campaign and these data were acquired in the blue wavelength range. Therefore the amplitudes estimated in von Essen et al. (2014), listed in Table 1, are used in this work as fixed values. This approach would be incorrect if the pulsation amplitudes would be significantly variable (see e.g., Breger et al. 2005). Nonetheless, the short time span of LBT data, and the achieved photometric precision compared to the intrinsically low values of WASP-33\u2019s amplitudes, make the detection of any amplitude variability impossible. ","Citation Text":["von Essen et al. (2014)"],"Functions Text":["The","photometric follow-up started in August, 2010, and ended in October, 2012, coinciding with these LBT data."],"Functions Label":["Background","Background"],"Citation Start End":[[1515,1538]],"Functions Start End":[[1511,1514],[1539,1645]]}
{"Identifier":"2016AandA...588A..25M__Vassiliadis_&_Wood_(1994)_Instance_1","Paragraph":"The faster evolutionary timescales and higher luminosities of our H-burning sequences should have an important impact on the study of the formation of PNe (Sch\u00f6nberner et al. 2014; Toal\u00e1 & Arthur 2014). The low-mass models of Sch\u00f6nberner (1983), which are still in use to complement the sequences of Bl\u00f6cker (1995a), show crossing timescales of ~340 kyr (0.546 M\u2299) and ~20 kyr (0.565 M\u2299). Our H-burning sequences of similar mass and metallicity (Z0 = 0.02) show crossing timescales about ~3.5 to \u227315 times faster and even faster in the case of M15 models. The discrepancy is even larger in the case of the low-mass models of Vassiliadis & Wood (1994). In order to be able to produce PNe, the central stars need to evolve in less than a few tens of thousand years. If that is not the case, the circumstellar material dissipates before the star becomes hot enough to ionize it. This fact, together with the very long timescales of \u03c4cross \u2273 100 kyr of the low-mass models of older grids, leads to the conventional wisdom that low-mass post-AGB stars (\u22720.55 M\u2299) cannot form PNe; see, e.g., Jacoby et al. (2013), Bond (2015). In this context, the new models might help to explain the existence of single CSPNe with masses lower than ~0.55 M\u2299 (Althaus et al. 2010; Werner & Rauch, in prep.). The new models should also have an impact on the question of whether single stellar evolution can form PNe in globular clusters (Jacoby et al. 2013; Bond 2015). Our H-burning post-AGB sequences with ages similar to that of globular clusters (9 to 12 Gyr) have values of \u03c4cross ~ 25\u201370 kyr. Timescales drop to \u03c4cross ~ 5\u201310 kyr for post-AGB sequences with slightly younger progenitors (5 to 7 Gyr); see Tables 2 and 3. Timescales are even shorter in the case of the models of Weiss & Ferguson (2009) and M15 (see Appendix B). The much shorter timescales of the new H-burning post-AGB sequences call into question the idea that single stellar evolution cannot produce PNe in globular clusters. ","Citation Text":["Vassiliadis & Wood (1994)"],"Functions Text":["The discrepancy is even larger in the case of the low-mass models of"],"Functions Label":["Differences"],"Citation Start End":[[625,650]],"Functions Start End":[[556,624]]}
{"Identifier":"2018MNRAS.481.5630S__Tody_1993_Instance_1","Paragraph":"We also introduce here NB2071 data, taken as part of the MDCS with the Multi-Object Infrared Camera and Spectrograph (MOIRCS; Ichikawa et al. 2006; Suzuki et al. 2008) on the Subaru Telescope (the same instrument that was used in the past MAHALO-Subaru survey; Koyama et al. 2013a). The observations were executed between April 30 and May 6, 2015, under photometric conditions with seeing FWHM \u223c0.6 arcsec. The integration time is 125 min which was split into 180 s individual exposures. After combining with the existing NB2071 data (186 min integration), we reconstructed all the data using the reduction pipeline mcsred2 (Tanaka et al. 2011), which is written as iraf3 scripts (Tody 1993). As described in Shimakawa et al. (2018), we executed flat-fielding, masking objects from the combined data in the first run (thus the whole reduction process was conducted twice to remake secure object masks), sky subtraction (by median sky and then the polynomially fitted plane for residual sky subtraction), distortion correction, cross-matching, and image mosaicing with this pipeline. The reconstructed NB2071 image reaches 23.95 mag in 3\u03c3 limiting magnitude using a 1.4 arcsec diameter aperture, and its seeing FWHM is 0.63 arcsec. The image depth becomes deeper by 0.5 mag than the previous data (Koyama et al. 2013a). The world coordinate system (WCS, Calabretta & Greisen 2002; Greisen & Calabretta 2002) of the narrow-band image is carefully matched by the iraf scripts (ccmap and ccsetwcs) to that of the F814W image, based on 67 point sources. F814W has one of the best spatial resolutions amongst our data set. The standard deviation of point source separations between the NB2071 and F814 images suggests that the relative WCS uncertainty would be around 0.04 arcsec in the survey area. One should note, however, that the absolute astrometry would have 0.3 arcsec errors in right ascension and declination based on comparison with the Guide Star Catalogue 2 (Lasker et al. 2008).","Citation Text":["Tody 1993"],"Functions Text":["After combining with the existing NB2071 data (186 min integration), we reconstructed all the data using the reduction pipeline mcsred2","which is written as iraf3 scripts"],"Functions Label":["Uses","Uses"],"Citation Start End":[[681,690]],"Functions Start End":[[488,623],[646,679]]}
{"Identifier":"2019AandA...626A..34C__Ohba_et_al._2005_Instance_1","Paragraph":"The uncertainty for most frequency determinations in the spectrum of NEFA is of a few kHz in view of the high S\/N of the data (see Fig. 20). However, we have assigned an uncertainty of 10 kHz for all lines showing a single component and intensity larger than 20 mK (S\/N\u2004\u2265\u200420). Unlike linear molecules with a nitrogen atom, for which hyperfine splittings rapidly collapse, symmetric or asymmetric tops can show significant hyperfine splitting even for high-J and high-Ka transitions. As an example, the bottom panel of Fig. 10 shows the hyperfine structure of K\u2004=\u20043 component of the J\u2004=\u20046\u2004\u2192\u20045 rotational transition of CH3CN. The K\u2004=\u20040 and 1 components of the same transition do not show any measurable splitting, while the K\u2004=\u20042 exhibits an emerging shoulder (clearly seen in that figure) due to its hyperfine structure. Using the diagonal elements of the 14N quadrupole coupling tensor in NEFA (Ohba et al. 2005), we have computed the expected hyperfine splitting of all lines in the W band. Most of the lines we have identified have the hyperfine splitting concentrated around \u00b110 kHz of the central frequency. However, several lines present three hyperfine components with similar intensities. Two of these components are at practically the same frequency, while the third component is always on the opposite side producing a global splitting of 60\u2013100 kHz. Hence, some of the observed lines could have a two-component structure, with intensity ratio 1:2, with a strong feature accompanied by less intense shoulder (a factor of two lower), either left or right of the central frequency. Many other lines also show a complex profile due to blending with lines coming from the vibrationally excited states and the other conformers of NEFA (see, for example, the bottom right panel of Fig. 20). We have fitted in most cases one single Gaussian profile to the observed lines. However, when the line profile exhibits two components we fitted two Gaussians. Depending on the degree of overlap between these components, the assigned uncertainty varies between 20 and 80 kHz. When two components were obvious in the line profile, the feature assigned to NEFA was based on the expected intensity of the transition (determined with the MADEX code; Cernicharo 2012), and by comparison with the intensity of other nearby lines of NEFA.","Citation Text":["Ohba et al. 2005"],"Functions Text":["Using the diagonal elements of the 14N quadrupole coupling tensor in NEFA","we have computed the expected hyperfine splitting of all lines in the W band."],"Functions Label":["Uses","Uses"],"Citation Start End":[[895,911]],"Functions Start End":[[820,893],[914,991]]}
{"Identifier":"2019MNRAS.486.1781R__Gu_et_al._2006_Instance_1","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":["Gu et al. 2006"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[906,920]],"Functions Start End":[[734,904]]}
{"Identifier":"2016ApJ...825...10T___2009_Instance_1","Paragraph":"Although this penalty is small, we can still provide a quantitative estimate based on a few assumptions. The observations of M15 that measured the parallax to VLA J2130+12 (K14) employed an interferometric technique where the position of a weaker potential in-beam calibrator source (M15 S1 and VLA J2130+12) can be measured against a brighter primary calibration source, with the potential for using the in-beam calibrator to transfer more accurate calibration solutions to other in-beam targets. Based on a literature review for papers where VLBI parallax measurements were determined using an in-beam calibrator, we find 24 similar measurements, the majority of which have been used to measure parallaxes to pulsars (Fomalont et al. 1999; Brisken et al. 2003; Chatterjee et al. 2005, 2009; Ng et al. 2007; Middelberg et al. 2011; Deller et al. 2012, 2013; Ransom et al. 2014; Liu et al. 2016). As these measurements represent additional opportunities to detect the parallax of a VLA J2130+12-like object, they could be considered as potential trials. However, all of the reported in-beam calibrators were significantly brighter (4\u201386 mJy) than VLA J2130+12 (\u223c0.1\u20130.5 mJy). Given our conservative assumption of a uniform volume density of VLA J2130+12-like objects in the Galaxy, we should down-weight the number of trials from brighter sources by (f\u03bd,in-beam\/f\u03bd,VLA J2130+12)\u22121.5. In that case, the number of trials penalty is only an additional \u223c0.2 trials. In addition, we note that the PSR\u03c0 parallax project (a large VLBA program) has reported15\n\n15\n\nhttps:\/\/safe.nrao.edu\/vlba\/psrpi\/\n\n 111 additional in-beam calibrator sources that they used to measure parallaxes. Although they do not provide the flux densities of individual sources, they note that their median in-beam calibrator source is 9.2 mJy. We have measured the flux density function of secure 1\u201320 mJy FIRST sources (Helfand et al. 2015; \n\n\n\n\n\n) to estimate the expected distribution of the flux densities of in-beam calibrators. We found that the minimum flux density is likely \u223c3.2 mJy, and using the same down-weighting we estimate an additional penalty of \u223c0.9 trials.","Citation Text":["Chatterjee et al.","2009"],"Functions Text":["Based on a literature review for papers where VLBI parallax measurements were determined using an in-beam calibrator, we find 24 similar measurements, the majority of which have been used to measure parallaxes to pulsars"],"Functions Label":["Uses"],"Citation Start End":[[763,780],[787,791]],"Functions Start End":[[498,718]]}
{"Identifier":"2020MNRAS.497.4293G__Ferguson_et_al._2004_Instance_1","Paragraph":"Recently Izotov et al. (2016a,b, 2018a,b) have discovered significant emission of Lyman continuum (LyC) ionizing radiation leaking with the escape fractions of 2\u201376 per cent in a sample of 11 low-z compact active star-forming galaxies (SFGs) observed with the Hubble Space Telescope (HST) in conjunction with the Cosmic Origins Spectrograph (COS). These galaxies hereafter referred to as LyC leakers, possess many properties similar to those of high-redshift galaxies both at z \u223c 2\u20133 and z$\\gtrsim$ 6 such as compact morphology with similar galaxy radii (e.g. Bouwens et al. 2004; Ferguson et al. 2004; Oesch et al. 2010; Ono et al. 2012; Shibuya, Ouchi & Harikane 2015; Curtis-Lake et al. 2016; Paulino-Afonso, Sobral & Ribeiro 2018), strong emission lines with high equivalent widths (EWs; e.g. Schaerer & de Barros 2009; Smit et al. 2014, 2015; Roberts-Borsani et al. 2016; Bowler et al. 2017; Castellano et al. 2017; Fletcher et al. 2019; Bian & Fan 2020; Endsley et al. 2020), similar low stellar masses, low metallicities, and high specific star formation rates (SFRs; e.g. Jaskot & Oey 2013; Nakajima et al. 2013; Stark et al. 2013a; de Barros, Schaerer & Stark 2014; Duncan et al. 2014; Gonz\u00e1lez et al. 2014; Nakajima & Ouchi 2014; Becker, Bolton & Lidz 2015; Grazian et al. 2015; Salmon et al. 2015; Huang et al. 2016; Stark 2016; Santini et al. 2017; Stark et al. 2017; Dors et al. 2018), small dust content (e.g. Ouchi et al. 2013; Ota et al. 2014; Maiolino et al. 2015; Schaerer et al. 2015; Watson et al. 2015), and are considered as the main sources of reionization of the Universe after the cosmic \u2018Dark Ages\u2019. This makes low-z LyC leakers the best local analogues of reionization galaxies (see e.g. Schaerer et al. 2016; Stark 2016; Ma et al. 2020). Given their proximity, these galaxies represent excellent laboratories for a detailed study of their physical conditions, and the main mechanisms responsible for LyC leakage. Ground-based spectroscopic observations in the visible and near-infrared ranges are necessary for that.","Citation Text":["Ferguson et al. 2004"],"Functions Text":["These galaxies hereafter referred to as LyC leakers, possess many properties similar to those of high-redshift galaxies both at z \u223c 2\u20133 and z$\\gtrsim$ 6 such as compact morphology with similar galaxy radii (e.g."],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[581,601]],"Functions Start End":[[348,559]]}
{"Identifier":"2019AandA...629A..95S___2017a_Instance_1","Paragraph":"A rough idea of the location of the inert Oort cloud can be obtained from previous works. Gladman et al. (2002) showed that the scattering effect by Neptune is significant over long timescales only for perihelion distances below about 45 astronomical units (au). The precise limit actually increases with the semi-major axis (Gallardo et al. 2012), because energy kicks result in larger variations of semi-major axis if the semi-major axis is large. In fact, some observed objects with perihelion beyond 45 au are known to experience scattering (Bannister et al. 2017). In any case, we look here for a rough limit only. The scattering process mostly affects the semi-major axis of small bodies, which diffuses chaotically, while the perihelion distance does not vary much. Later on, Gomes et al. (2005), Gallardo et al. (2012), and Saillenfest et al. (2016, 2017a), showed that the Lidov\u2013Kozai mechanism raised by the giantplanets inside a mean-motion resonance with Neptune is able to raise the perihelion of small bodies beyond 60 au in a few thousands of million years. Contrary to scattering effects, this mechanism induces a variation of perihelion distance and inclination, while the semi-major axis remains at the resonance location. This mechanism, however, is only efficient for semi-major axes smaller than about 500 au. From these studies, one can deduce that the action of the planets is limited to orbits with perihelion distances smaller than about 80 au, and that for perihelion beyond 45 au, the semi-major axis should be smaller than 500 au for the planets to possibly have a substantial effect through mean-motion resonances. As regards the effects of the galactic tides, Fouchard et al. (2017) showed that an object with perihelion in the Jupiter-Saturn region, that is, below 15 au from the Sun, should have a semi-major axis larger than 1600 au for the tides to be able to raise its perihelion beyond 45 au in less than the age of the solar system. In other words, the tides can move its perihelion out of reach of any significant planetary scattering.","Citation Text":["Saillenfest et al.","2017a"],"Functions Text":["Later on, Gomes et al. (2005), Gallardo et al. (2012), and","showed that the Lidov\u2013Kozai mechanism raised by the giantplanets inside a mean-motion resonance with Neptune is able to raise the perihelion of small bodies beyond 60 au in a few thousands of million years."],"Functions Label":["Background","Background"],"Citation Start End":[[832,850],[858,863]],"Functions Start End":[[773,831],[866,1072]]}
{"Identifier":"2021ApJ...908...40M__Muschietti_&_Lemb\u00e8ge_2017_Instance_2","Paragraph":"These signatures are inconsistent with ultra-low frequency waves, which have circular polarization and a period similar to the upstream ion gyroperiod. The waves are also inconsistent with ion Weibel instability, which generates linearly polarized waves. Interaction of reflected ions with incoming solar wind electrons or ions can cause foot instabilities that excite waves in the whistler mode branch. Modified Two Stream Instability (MTSI) due to relative drift between reflected ions and incoming solar wind electrons (fast drift), and incoming solar wind ions and electrons (slow drift) has been frequently considered (Matsukiyo & Scholer 2003; Comi\u015fel et al. 2011; Umeda et al. 2012; Marcowith et al. 2016; Wilson et al. 2016; Muschietti & Lemb\u00e8ge 2017; Hull et al. 2020). This instability, however, if excited, creates significant ion heating throughout the foot and suppresses the reformation process (Shimada & Hoshino 2005; Matsukiyo & Scholer 2006), rather than creating episodic enhancements that we show in the foot. Furthermore, Gary et al. (1987) indicated that (fast drift) MTSI becomes dominant at low electron beta (\u03b2e  0.5), while at higher \u03b2e more resonant electrons stabilize this instability through increased electron Landau damping. Electron data for the time period we discussed in this paper show \u03b2e \u2265 1.2, and therefore fast drift mode MTSI is most likely not significant. The slow drift mode of MTSI could be a more viable candidate at high \u03b2 plasmas. Wave properties around 1.6 Hz in the middle interval (purple segment) of Figure 6, indicate that the wave is in propagation toward the ramp (\n\n\n\n\n\n = \u22120.66, \u22120.71, 0.22) with Vph\u2010sc = 34 km s\u22121 and \u03bbwave = 21.4 km  \u223c 12\u03bbe, where \u03bbe is the upstream electron inertial length. The plasma rest-frame frequency of the wave is about 8 Hz  \u223c 2.5flh. Since these characteristics are somewhat consistent with model predictions for drift mode of MTSI (Muschietti & Lemb\u00e8ge 2017), we do not rule out the possibility of some waves at certain frequencies and during some intervals being generated by the slow drift mode of MTSI.","Citation Text":["Muschietti & Lemb\u00e8ge 2017"],"Functions Text":["Since these characteristics are somewhat consistent with model predictions for drift mode of MTSI",", we do not rule out the possibility of some waves at certain frequencies and during some intervals being generated by the slow drift mode of MTSI."],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[1922,1947]],"Functions Start End":[[1823,1920],[1948,2095]]}
{"Identifier":"2015AandA...577A.117S__Kaluzny_et_al._(2013a)_Instance_1","Paragraph":"A lower CB frequency in GCs than in the field has already been noticed by Rucinski (2000). It results most likely from a very low binary frequency of ~0.1 in GCs (but with a significant scatter among different clusters), compared to 0.5 in the field (Milone et al. 2012), together with a low percentage of binaries in the contact phase at a given age, as shown above. Any statistical comparison of models with observations is therefore extremely uncertain because we have so few observational data. Nevertheless, we can compare the predicted to observed fraction of BSs among all CBs in a GC. We used data for two clusters observed by Kaluzny and his collaborators with the highest number of CBs (except for \u03c9 Centauri). In M4, Kaluzny et al. (2013a) detected nine CBs with one BS among them, and in 47 Tuc, Kaluzny et al. (2013b) identified 15 CBs or NCBs, of which six are BS. The resulting fractions are 0.1 and 0.4, respectively. We did not take into account \u03c9 Cen with the richest population of CBs (Kaluzny et al. 2004) because it is highly atypical and the accurate number of member CBs is not well known, although the approximate data indicate a similarly high ratio as in case of 47 Tuc. The predicted fractions range from 0.08 to 0.2 for the three considered cluster ages. As we see, they are close to the lower observed value but are at odds with the higher value. This may indicate some deficiencies of the CCBM, for instance, too short initial cut-off period, a too high AML rate at short periods, or a too low mass-transfer rate. With the lower AML rate and\/or higher mass transfer rate, a binary stays longer in contact and can reach a lower mass ratio before merging. Many of the short-period CBs have a rather high mass ratio of about 0.7\u22120.8 at the time of merging, whereas field CBs with P \u2272 0.3 d center around a value of 0.5 (Rucinski 2010; St\u0229pie\u0144 & Gazeas 2012). The lower mass ratio means more mass transferred to the gainer, hence its higher mass and higher position on the HRD, that is, a higher probability for entering the BS region. There may also still be another explanation of the discrepancy: strong fluctuations of this fraction among different clusters suggest a nonuniform BSs formation rate in some clusters, with individual bursts occurring in the recent past (see below). Regardless of the reason, it is apparent that profound differences occur among GCs, which makes a comparison of theoretical predictions with individual clusters uncertain. ","Citation Text":["Kaluzny et al. (2013a)"],"Functions Text":["We used data for two clusters observed by Kaluzny and his collaborators with the highest number of CBs (except for \u03c9 Centauri).","In M4,","detected nine CBs with one BS among them"],"Functions Label":["Uses","Background","Background"],"Citation Start End":[[728,750]],"Functions Start End":[[593,720],[721,727],[751,791]]}
{"Identifier":"2019AandA...623A.140G__Pohl_et_al._2017_Instance_2","Paragraph":"HD 169142 is a very young Herbig Ae-Be star with a mass of 1.65\u20132 M\u2299 and an age of 5\u201311 Myr (Blondel & Djie 2006; Manoj et al. 2007) that is surrounded by a gas-rich disk (i = 13\u00b0; Raman et al. 2006; PA = 5\u00b0; Fedele et al. 2017) that is seen almost face-on. The parallax is 8.77 \u00b1 0.06 mas (Gaia DR2 2018). Disk structures dominate the inner regions around HD 169142 (see, e.g., Ligi et al. 2018). Figure 1 shows the view obtained from polarimetric observations: the left panel shows the Q\u03a6 image in the J band obtained by Pohl et al. (2017) using SPHERE on a linear scale, and the two rings are clearly visible. The right panel shows a pseudo-ADI image of the inner regions obtained by differentiating the Q\u03a6 image (see Ligi et al. 2018, for more details). Biller et al. (2014) and Reggiani et al. (2014) discussed the possible presence of a point source candidate at small separation (0.2 arcsec from the star). However, the analysis by Ligi et al. (2018) based on SPHERE data does not support or refute these claims; in particular, they suggested that the candidate identified by Biller et al. (2014) might be a disk feature rather than a planet. Polarimetricimages with the adaptive optics system NACO at the Very Large Telescope (VLT; Quanz et al. 2013b), SPHERE (Pohl et al. 2017; Bertrang et al. 2018) and GPI (Monnier et al. 2017) show a gap at around 36 au, with an outer ring at a separation >40 au from the star. This agrees very well with the position of the rings obtained from ALMA data (Fedele et al. 2017); similar results were obtained from VLA data (Osorio et al. 2014; Mac\u00edas et al. 2017). We summarize this information about the disk structure in Table 1 and call the ring at 0.17\u20130.28 arcsec from the star Ring 1 and the ring at 0.48\u20130.64 arcsec Ring 2. We remark that in addition to these two rings, both the spectral energy distribution (Wagner et al. 2015) and interferometric observations (Lazareff et al. 2017; Chen et al. 2018) show an inner disk at a separation smaller than 3 au. This inner disk isunresolved from the star in high-contrast images and consistent with ongoing accretion from it onto the young central star.While the cavities between the rings seem devoid of small dust, some gas is present there (Osorio et al. 2014; Mac\u00edas et al. 2017; Fedele et al. 2017). Fedele et al. (2017) and Bertrang et al. (2018) have suggested the possibility that the gap between Rings 1 and 2 is caused by a planet with a mass slightly higher than that of Jupiter. However, this planet has not yet been observed, possibly because it is at the limit of or beyond current capabilities of high-contrast imagers. On the other hand, Bertrang et al. (2018) found a radial gap in Ring 1 at PA ~ 50\u00b0 that might correspond to a similar radial gap found by Quanz et al. (2013b) at PA ~ 80\u00b0. The authors noted that if this correspondence were real, then this gap might be caused by a planet at about 0.14 arcsec from the star. So far, this planet has not been unambiguously detected either.","Citation Text":["Pohl et al. 2017"],"Functions Text":["Polarimetricimages with","SPHERE","show a gap at around 36 au, with an outer ring at a separation >40 au from the star."],"Functions Label":["Uses","Uses","Uses"],"Citation Start End":[[1269,1285]],"Functions Start End":[[1150,1173],[1261,1267],[1339,1423]]}
{"Identifier":"2022MNRAS.516.2597S__Padmanabhan_2002_Instance_1","Paragraph":"Observational data from distant Type Ia supernovae (SNIa) revealed a hidden fact that the present Universe is expanding in an accelerated phase (Riess et al. 1998; Perlmutter et al. 1999; Kowalski et al. 2008). This important fact has been confirmed by other cosmological observations such as the cosmic microwave background (CMB) (Komatsu et al. 2009; Jarosik et al. 2011; Ade et al. 2016), large-scale structure (LSS), and baryonic acoustic oscillation (BAO) (Tegmark et al. 2004; Cole et al. 2005; Eisenstein et al. 2005; Percival et al. 2010; Blake et al. 2011; Reid et al. 2012), high-redshift galaxies (Alcaniz 2004), high-redshift galaxy clusters (Wang & Steinhardt 1998; Allen et al. 2004), and weak gravitational lensing (Benjamin et al. 2007; Amendola, Kunz & Sapone 2008; Fu et al. 2008). The accelerated phase of the expansion can be interpreted by modifying the standard theory of gravity on cosmological scales or by adding an exotic cosmic fluid with negative pressure, the so-called dark energy (DE) (Riess et al. 1998; Perlmutter et al. 1999; Kowalski et al. 2008). Historically, Einstein\u2019s cosmological constant \u039b with a constant equation of state (EoS) parameter equal to \u22121 is the first and simplest DE model to interpret the current accelerated phase of the expansion. The standard model of cosmology, the \u039bCDM model, which accounts for about 70\u2009per cent of the total energy density of the Universe from \u039b and about 30\u2009per cent from cold dark matter (CDM) is consistent with the most of the cosmological data. From a theoretical point of view, however, this model has two fundamental problems, namely the problem of fine-tuning and the problem of cosmic coincidence (Weinberg 1989; Sahni & Starobinsky 2000; Carroll 2001; Padmanabhan 2003; Copeland, Sami & Tsujikawa 2006). In addition, the standard model of cosmology suffers from the big tension between the local measurement values of Hubble constant H0 with that of the Planck CMB estimation. Moreover, there are big tensions between the Planck CMB data with weak lensing measurements and redshift surveys, concerning the value of non-relativistic matter density \u03a9m and the amplitude of the growth of perturbations \u03c38. The above observational tensions can make the standard model as an approximation of a general gravitational scenario yet to be found. For a recent review about the cosmological tensions of the standard model, we refer the reader to see (Abdalla et al. 2022). In this concern, a large family of dynamical DE models with time-varying EoS parameters has been proposed in the literature (for some earlier attemts, see Armendariz-Picon, Mukhanov & Steinhardt 2000; Caldwell 2002; Padmanabhan 2002; Elizalde, Nojiri & Odintsov 2004). Parallel to the solution of DE, the positive cosmic acceleration can be seen as the expression of a new theory of gravity on large cosmological scales. Indeed, modifying the standard Einstein\u2013Hilbert action in the context of the Friedmann\u2013Robertson\u2013Walker (FRW) metric leads to the modified Friedmann equations, which can be used to justify the current accelerated expansion of the Universe without resorting to DE fluid. One of the most popular modified gravity theories is the f(R) scenario, in which the Lagrangian of the modified Einstein\u2013Hilbert action is extended to the function of the Ricci scalar R (Capozziello, Stabile & Troisi 2007; Capozziello & Francaviglia 2008; Sotiriou & Faraoni 2010; Nojiri & Odintsov 2011). Besides the f(R) theories of gravity, the so-called f(T) theory of gravity is the other solution to solve the puzzle of cosmic acceleration. This theory is defined on the basis of the old definition of the teleparallel equivalent of general relativity (TEGR), first introduced by Einstein (Einstein 1928) and extended by (Hayashi & Shirafuji 1979; Maluf 1994). A comprehensive study on the theory of teleparallel gravity (TG) can be found in the recent review by (Bahamonde et al. 2021). In general, an extended theory of gravity involves curvature, torsion, and non-metricity components. If only torsion is non-vanishing, one can obtain the torsional teleparallel geometry which is the basic geometry for the most f(T) theories in literature. TG theory can be used to formulate a TEGR formalism, which is dynamically equivalent to GR but may have different behaviors for other scenarios, such as quantum gravity. The Horndeski gravity can also be formulated using the teleparallel geometry as a possible revival modes for regular Hordenski gravity models (see Bahamonde et al. 2021, for more detils). TG as a theory built on the tangent sapce must be invariant under general coordinate transformation and local lorentz transformation. In the first formulation of teleparallel theories, it was assumed that the spin connection was always zero and then torsion tensor depends on the tetrads. This torsion tensor is a particular case which is computed in the so-called Weitzenb\u00f6ck gauge (see section 2.2.3 of Bahamonde et al. 2021). By taking a local Lorentz transformation only in the tetrads, one can conclude that the torsion tensor is non-covariant quantity under the local Lorentz transformation. In this context, the action of TEGR has a total divergence term which can be removed as being a boundary term. Hence, the TEGR action is invariant under local Lorentz transformation up to a boundary term. In modified teleparallel theories of gravity like f(T) gravity, we have no boundary term anymore, meaning that f(T) gravity models break the local Lorentz invariance. Notice that the problem of breaking the local Lorentz invariance is related to the particular Weitzenb\u00f6ck gauge (spin connection zero) (Kr\u0161\u0161\u00e1k & Saridakis 2016). If we take the simultaneous transformations in the tetrads and the spin connection, the teleparallel theories are then fully invariant (diffeomorphisms and local Lorentz). Thus, any action and consequently field equations constructed based on the torsion tensor will be fully invariant. In the context of cosmology, there is a large body of works that has examined the cosmological properties of various f(T) models. In this framework, the dynamics and various aspects of the Universe with homogeneous and isotropic background are studied in recent works (Bengochea & Ferraro 2009; Linder 2010; Wu & Yu 2011, 2010b; Bamba et al. 2011, 2012; Dent, Dutta & Saridakis 2011; Geng et al. 2011; Myrzakulov 2011; Zhang et al. 2011). In general, cosmological consequences for the various formulations of TG at both background and perturbation levels have been discussed in (Bahamonde et al. 2021). The f(T) models have been studied and constrained using the various cosmological data (Wu & Yu 2010a; Capozziello, Luongo & Saridakis 2015; Iorio, Radicella & Ruggiero 2015; Nunes, Pan & Saridakis 2016; Saez-Gomez et al. 2016). As a more recent study, we refer the work of Briffa et al. 2022, where the various versions of f(T) model have been studied observationally. Using the combinations of cosmological data including the cosmic chronometers, the SNIa observations from Pantheon catalogue, BAO observations, and different model-independent values for Hubble constant H0, Briffa et al. (2022) studied a detailed analysis of the impact of the H0 priors on the late time properties of f(T) cosmologies. They found a higher value of H0 compared to equivalent analysis without considering H0 priors. In addition, they showed that the f(T) model produces the higher value of H0 and slightly lower value of matter density \u03a9m0 as compared to standard \u039bCDM model. In general, studies on the cosmological tensions and their possible alleviation in TG formalism have been addressed in section 10 of review article (Bahamonde et al. 2021). For a review of cosmological tensions H0 and \u03c38, we refer the reader to see the current studies in (Verde, Treu & Riess 2019; Di Valentino et al. 2021a, b, c; Perivolaropoulos & Skara 2021).","Citation Text":["Padmanabhan 2002"],"Functions Text":["In this concern, a large family of dynamical DE models with time-varying EoS parameters has been proposed in the literature (for some earlier attemts, see"],"Functions Label":["Background"],"Citation Start End":[[2669,2685]],"Functions Start End":[[2453,2607]]}
{"Identifier":"2022ApJ...928..120G__Calzetti_et_al._2007_Instance_1","Paragraph":"As part of this study, we investigate the 8 \u03bcm emission within our sample of galaxies using archival data from IRAC on the Spitzer Space Telescope (SST). The rest-frame 8 \u03bcm emission from galaxies has historically been used as a monochromatic SFR indicator. This is justified by the fact that polycyclic aromatic hydrocarbons (PAHs) and small dust grains can be heated by the photodissociation regions (PDRs) that surround actively star-forming regions and the subsequent emission from PAHs is brightest at about 8 \u03bcm (Draine & Li 2007; Smith et al. 2007; Kennicutt et al. 2009; Elbaz et al. 2011). However, there is an important caveat; PAHs are destroyed by ionizing radiation from newly formed stars (Helou et al. 2004; Povich et al. 2007; Bendo et al. 2008; Rela\u00f1o & Kennicutt 2009), leading to a deficit in the 8 \u03bcm luminosity in galaxies with low metallicity, whereas PAHs are less well shielded by metals (Engelbracht et al. 2005; Calzetti et al. 2007; Smith et al. 2007; Cook et al. 2014; Shivaei et al. 2017). It has also been argued that metals can act as catalysts for the formation and growth of PAHs, leading to smaller average sizes in low-metallicity environments (Sandstrom et al.2012). More recently, Lin et al. (2020) find an anticorrelation between the dust-only 8 \u03bcm luminosity and the age of young stellar clusters, suggesting the 8 \u03bcm luminosity decreases with increasing age of the stellar population. The existence of a strong interstellar radiation field is also found to suppress the emission from PAHs, independent of metallicity (Madden et al. 2006; Gordon et al. 2008; Lebouteiller et al. 2011; Shivaei et al. 2017; Binder & Povich 2018). Adding to the complexity, spatially resolved studies have shown that a significant amount of 8 \u03bcm emission is associated with the cold, diffuse ISM, which suggests an important heating source other than recent (100 Myr) star formation (Bendo et al. 2008; Calapa et al. 2014; Lu et al. 2014). These results suggest that the luminosity at 8 \u03bcm is not a straightforward indicator of the SFR. In this work, we further investigate these issues by exploring how metallicity gradients within our sample affect the observed \u201cred\u201d side versus \u201cblue\u201d side IR color\u2013color correlations\u2014where we derive the \u201cblue\u201d side color from the flux density ratios at 8 and 24 \u03bcm.","Citation Text":["Calzetti et al. 2007"],"Functions Text":["However, there is an important caveat; PAHs are destroyed by ionizing radiation from newly formed stars","leading to a deficit in the 8 \u03bcm luminosity in galaxies with low metallicity, whereas PAHs are less well shielded by metals"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[938,958]],"Functions Start End":[[599,702],[788,911]]}
{"Identifier":"2019AandA...631A...5Z__Mink_et_al._(2013)_Instance_1","Paragraph":"Post-main-sequence mergers. Mergers of evolved stars with their MS companions originate from wider systems than main-sequence mergers (log P\u2004\u2273\u20040.7). We thus roughly estimate \n\n\n\n\nf\n\nP\n,\ni\n\n\n\u223c\n\nIPF\n\n0.7\n\n\n3.0\n\n\n\u223c\n0.66\n\n\n$ f_{P , \\mathrm{i}} \\sim \\mathrm{IPF}_{0.7}^{3.0} \\sim 0.66 $\n\n\n, using log10P\u2004=\u20043 as the maximum initial period for binary interaction (e.g., Claeys et al. 2011; Yoon et al. 2017). For merging to occur, a CEE phase needs to be initiated. Extreme mass ratio systems are prone to unstable mass transfer and thus to CEE, and although the exact value of the boundary is not well-constrained, we assume that this occurs for evolved donors in systems with q\u2004\u2272\u20040.4, following previous works such as Wellstein et al. (2001), Hurley et al. (2002) and de Mink et al. (2013). In principle, CEE can alternatively lead to the ejection of the envelope, but in this simple estimate we assume that all these systems eventually merge, leaving some hydrogen-rich layers on the surface of the formed star. This is consistent in most of the cases with the findings from our computational simulations for our standard assumption of using the entire orbital energy change to eject the envelope (\u03b1CEE\u2004=\u20041, as we will introduce and discuss in Sect. 4.1). Thus, \n\n\n\n\nf\n\nq\n,\ni\n\n\n\u223c\n\nIQF\n\n0.1\n\n\n0.4\n\n\n\u223c\n0.33\n\n\n$ f_{q , \\mathrm{i}} \\sim \\mathrm{IQF}_{0.1}^{0.4} \\sim 0.33 $\n\n\n. The projection of the initial parameter space for this channel at the log P\u2005\u2212\u2005q plane is depicted in magenta on the left panel of Fig. 2. The donor star in such systems originate from a roughly similar part in M1 space as main-sequence mergers (\n\n\n\n\nf\n\n\nM\n1\n\n,i\n\n~\n\nIMF\n7\n\n25\n\n~1.25\n\n$ f_{M_1 , {\\rm i}} \\sim {\\rm IMF}_7^{25} \\sim 1.25 $\n\n\n\n). In contrast, binaries of less extreme mass ratio are able to either follow a phase of stable mass transfer onto the secondary star or are assumed to always survive a CEE by ejecting the envelope, avoiding a merger. Such donors are expected to eventually produce a hydrogen-poor, stripped-envelope SN, not contributing to the SN II population that we focus on in this study. We find XpostMS\u2005+\u2005MS\u2006mergers\u2004\u223c\u2004C\u2005\u00d7\u20050.136\u2004\u223c\u200415%.","Citation Text":["de Mink et al. (2013)"],"Functions Text":["Extreme mass ratio systems are prone to unstable mass transfer and thus to CEE, and although the exact value of the boundary is not well-constrained, we assume that this occurs for evolved donors in systems with q\u2004\u2272\u20040.4, following previous works such as Wellstein et al. (2001), Hurley et al. (2002) and"],"Functions Label":["Background"],"Citation Start End":[[763,784]],"Functions Start End":[[459,762]]}
{"Identifier":"2022AandA...661A..10B__Ghirardini_et_al._2021a_Instance_3","Paragraph":"It is also possible that these clusters have a smaller extent and can just be missed by our extent selection as our detection algorithm sets the extent to zero if it is smaller than 6 (Brunner et al. 2022). Following the method presented in Ghirardini et al. (2021a), we estimated several dynamical properties of the clusters in the point source sample and compared them with the extent-selected sample presented in Ghirardini et al. (2021a). In Fig. 5 we compare the distributions of the core radii (Rcore) constrained by the V06 model and the concentration parameter (cSB) between these two samples. The concentration parameter is defined as the ratio of the surface brightness within 0.1 R500 to the surface brightness within R500 (Ghirardini et al. 2021a; Santos et al. 2008; Maughan et al. 2012). Intuitively, the expectation is that the smaller the core radius, the more compact the cluster. The left panel of Fig. 5 clearly shows that the clusters in the point source sample have relatively smaller core radii, hence the emission is more concentrated in a smaller area. Consistently, the concentration of the point source sample shows a clear excess in higher values than the extent-selected sample, indicating that a significantly larger fraction of cool-core clusters and clusters host a central AGN. We performed the same experiment by applying cuts in flux 1.5 \u00d7 10\u221214 ergs s\u22121 cm\u22122 and in detection likelihood to test whether the clusters are missed by the extent selection because they are fainter and\/or more compact than the extent-selected clusters. The distribution of number density, core radius, and concentration parameter remains the same, indicating that the population of clusters in the point source catalog is more compact than the extent-selected sample. The extent-selected sample does not show a clear bias toward cool-core clusters or clusters with a central AGN, but contains the fraction of cool-cores is similar to that of SZ surveys (Ghirardini et al. 2021a). In this sample, we observe the opposite trend.","Citation Text":["Ghirardini et al. 2021a"],"Functions Text":["In Fig. 5 we compare the distributions of the core radii (Rcore) constrained by the V06 model and the concentration parameter (cSB) between these two samples.","The concentration parameter is defined as the ratio of the surface brightness within 0.1 R500 to the surface brightness within R500","Intuitively, the expectation is that the smaller the core radius, the more compact the cluster.","The left panel of Fig. 5 clearly shows that the clusters in the point source sample have relatively smaller core radii, hence the emission is more concentrated in a smaller area. Consistently, the concentration of the point source sample shows a clear excess in higher values than the extent-selected sample, indicating that a significantly larger fraction of cool-core clusters and clusters host a central AGN."],"Functions Label":["Compare\/Contrast","Background","Background","Compare\/Contrast"],"Citation Start End":[[735,758]],"Functions Start End":[[443,601],[602,733],[802,897],[898,1309]]}
{"Identifier":"2018MNRAS.477.2220T__Miettinen_2014_Instance_1","Paragraph":"We further restricted our sample to well-defined N2H+ (1\u20130) spectra that we used to estimate the gas velocity dispersion. The N2H+ (1\u20130) emission of each clump was evaluated from the MALT90 data cubes by averaging the spectrum across all the pixels within one MALT90 beam, \u2243 38 arcsec. We assumed that all the N2H+ emission comes from the clumps, and we estimate the filling factor from the comparison of the radius of each Hi-GAL clump with the radius of a region equal to the MALT90 beam (Fig. 1). There is a strong correlation between these two quantities, and the size of the Hi-GAL clumps is systematically smaller than the radius estimated from the MALT90 beam for a factor of 0.64, on average. We assumed an average filling factor of 0.64 for the entire sample. The MALT90 data cubes are given in antenna temperature $T_{\\rm A}^{*}$ and they have been converted to the main beam temperature $T_{\\rm MB}=T_{\\rm A}^{*}\/\\eta _{\\rm MB}$, assuming a mean beam efficiency \u03b7MB = 0.49 (Miettinen 2014). The properties of each N2H+ (1\u20130) averaged spectrum have been extracted in idl using a hyperfine fitting routine and the mpfitfun algorithm (Markwardt 2009), after smoothing the data to a spectral resolution of 0.3\u2009km\u2009s\u22121 to enhance the signal-to-noise (S\/N) ratio. We excluded all clumps with a S\/N ratio below 5, where the rms in each smoothed data cube has been measured in a 100\u2009km\u2009s\u22121 wide spectral window near the N2H+ emission. We further excluded clumps for which the fit converged but the spectrum was affected by spikes and\/or by multiple components along the line of sight. Using these criteria, we obtained 308 clumps. We completed our selection by excluding all clumps without a clear distance assignation, in particular without a well-defined resolution of the near\u2013far distance ambiguity. First, we have refined the kinematic distances in the Elia et al. (2017) catalogue (and the quantities that depend on them) with the newest set of distances developed for the Hi-GAL survey under the VIALACTEA project (Mege et al., in preparation). The method used by Elia et al. (2017) was the same as in Russeil et al. (2011): the brightest emission lines in the 12CO or 13CO spectra along the line of sight of each source are used to estimate the velocities of the local standard of rest and converted into heliocentric distances using the Brand & Blitz (1993) rotation curve. The distances in Mege et al. (in preparation) have been determined using a similar approach, but including all the recent surveys of the Galactic plane to trace structures along the line of sight, and using the more recent Reid et al. (2009) rotation curve. Then, in order to identify only clumps with a well-defined distance estimation, we have compared the distances assigned to our 308 clumps with the distances of the MALT90 sample estimated in Whitaker et al. (2017) and of the ATLASGAL sources published in Urquhart et al. (2018). We excluded from the sample all sources with a difference in the distance estimation larger than 20 per cent among the three surveys.","Citation Text":["Miettinen 2014"],"Functions Text":["assuming a mean beam efficiency \u03b7MB = 0.49"],"Functions Label":["Uses"],"Citation Start End":[[985,999]],"Functions Start End":[[941,983]]}
{"Identifier":"2022ApJ...925..123N__Tielens_&_Charnley_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":["Tielens & Charnley 1997"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[981,1004]],"Functions Start End":[[794,979]]}
{"Identifier":"2022MNRAS.515.5135H__Zenteno-Quinteros,_Vi\u00f1as_&_Moya_2021_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":["Zenteno-Quinteros, Vi\u00f1as & Moya 2021"],"Functions Text":["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.","Observations have struggled to confirm these theories."],"Functions Label":["Background","Differences"],"Citation Start End":[[637,673]],"Functions Start End":[[365,534],[729,783]]}
{"Identifier":"2015ApJ...800...38G___2003_Instance_1","Paragraph":"The currently accepted cold dark matter dominated model with the cosmological constant (\u00ce\u009bCDM) predicts that structures in our Universe assemble hierarchically, with more massive systems forming later through accretion and mergers of smaller, self-bound dark-matter halos (e.g., Tormen 1997; Moore et\u00c2 al. 1999; Klypin et\u00c2 al. 1999; Springel et\u00c2 al. 2001). In N-body cosmological simulations, dark matter halos of all masses converge to a roughly \u00e2\u0080\u009cuniversal\u00e2\u0080\u009d and cuspy density profile that steepens with radius, the so-called Navarro\u00e2\u0080\u0093Frenk\u00e2\u0080\u0093White profile (NFW profile; Navarro et\u00c2 al. 1996, 1997). Moreover, the degree of central concentration of a halo depends on its formation epoch and hence on its total mass (e.g., Wechsler et\u00c2 al. 2002; Zhao et\u00c2 al. 2003). Within this scenario, early virialized objects are compact when they get accreted into a larger halo. Such objects are usually referred to as subhalos or substructures of their host and, as they orbit within the host potential well, they are strongly affected by tidal forces and dynamical friction, causing mass, angular momentum, and energy loss (e.g., Ghigna et\u00c2 al. 1998; Tormen et\u00c2 al. 1998; De Lucia et\u00c2 al. 2004; Gao et\u00c2 al. 2004). In the \u00ce\u009bCDM framework, more massive halos are predicted to have a larger fraction of mass in subhalos than lower mass halos because in the former there has been less time for tidal destruction to take place (e.g., Gao et\u00c2 al. 2004; Contini et\u00c2 al. 2012). On galaxy cluster scales, observational tests of these predictions have been attempted in some previous works (e.g., Natarajan et\u00c2 al. 2007, 2009), but highly accurate analyses are becoming possible only now, thanks to the substantially improved quality of the available photometric and spectroscopic data. From an observational point of view, more investigations are still required to fully answer key questions on the formation and evolution of subhalos. How much mass of subhalos is stripped as they fall into the host potential? How many subhalos survive as bound objects? What are the spatial and mass distributions of the subhalos?","Citation Text":["Zhao et\u00c2 al. 2003"],"Functions Text":["Moreover, the degree of central concentration of a halo depends on its formation epoch and hence on its total mass (e.g.,"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[750,767]],"Functions Start End":[[605,726]]}
{"Identifier":"2015MNRAS.453.2126M__Hubrig_et_al._2013_Instance_1","Paragraph":"There are additional lines of evidence suggesting that accretion in HD 100546 could be magnetospheric, and not through a BL. Equation 3 in Johns-Krull, Valenti & Koresko (1999) provides a lower limit to the magnetic field necessary to drive MA, in terms of the stellar parameters and accretion rate. Assuming a mass accretion rate of \u223c10\u22127 M\u2299 yr\u22121, the minimum stellar mass and maximum stellar radius allowed by the uncertainties provided by Fairlamb et al. (2015), and a minimum rotation period of 0.26 d (from a maximum projected rotational velocity vsin\u2009i = 65 km s\u22121 and a minimum inclination i = 22\u00b0, from Guimar\u00e3es et al. 2006, and TAT11), HD 100546 could require a magnetic field of only several tens of Gauss to drive accretion magnetospherically.1 This is consistent with the magnetic field measured by Hubrig et al. (2009), 89 \u00b1 26 G, although non detections have also been reported (Donati et al. 1997; Hubrig et al. 2013). In addition, if accretion in HD 100546 is actually magnetospheric the Keplerian gaseous disc should be truncated at a stelleocentric radius that increases with the stellar magnetic field and radius, and decreases with the accretion rate and stellar mass (Elsner & Lamb 1977). Using equation 6 in Tambovtseva et al. (2014), the disc truncation radius of HD 100546 should be \u2272 0.01\u2009au (Fig. 3). This value is consistent with the Keplerian radius inferred from the width of the Br \u03b3 line profile (\u223c \u00b1 200 km s\u22121; Section 3); once this value is de-projected using the inclination in TAT11, the corresponding Keplerian distance is \u2272 0.02 au. It is noted that, in contrast to optically thick lines like H \u03b1, Br \u03b3 is mainly broadened by the Doppler effect (Tambovtseva et al. 2014). If the Keplerian disc is truncated by the stellar magnetic field, the gas would then fall ballistically on to the stellar surface. This could eventually be traced from the presence of redshifted absorptions at free-fall velocities in the profiles of several lines. In fact, previous spectroscopic analysis involving optical\/UV lines suggest MA\/ejection processes in HD 100546 (Vieira et al. 1999; Deleuil et al. 2004), with redshifted absorptions at velocities comparable to free-fall (Guimar\u00e3es et al. 2006). However, those signatures are not observed in the Br \u03b3 profile of HD 100546. The unresolved component contributing \u223c35 per cent to the visibility (Section 4.1) could be related to magnetospheric infall, but the small spatial scales involved cannot be probed from our observations.","Citation Text":["Hubrig et al. 2013"],"Functions Text":["This is consistent with the magnetic field measured by Hubrig et al. (2009), 89 \u00b1 26 G, although non detections have also been reported"],"Functions Label":["Similarities"],"Citation Start End":[[914,932]],"Functions Start End":[[757,892]]}
{"Identifier":"2022AandA...664A...2T__Virtanen_et_al._2017_Instance_1","Paragraph":"The earliest observations of magnetic fields in astrophysics were sunspot field strength measurements at the Mount Wilson Observatory (MWO) in California, USA, in 1908 (Hale 1908). Observations employed the Zeeman effect and were based on measuring the separation (splitting) between the two components of a spectral line, first Fe\u202fI 6173 \u00c5 and later Fe\u202fII 5250 \u00c5. Daily observations of sunspot magnetic fields have been conducted since 1917 (Hale et al. 1919). In the early 1950s, the invention of an electronic magnetograph (Babcock 1953) allowed the measurement of regions with weaker magnetic fields than in sunspots, such as plages. Regular full-disk magnetograms have been observed since early 1960, first at MWO (Howard 1974) and then at the National Solar Observatory (NSO) at Kitt Peak starting in 1973 (Livingston et al. 1976). In our long-term project aimed at reconstructing the past magnetic activity on the Sun (Pevtsov et al. 2016, 2019; Virtanen et al. 2017, 2018, 2019a,b), we also use indirect measurements of magnetic fields obtained, in particular, from chromospheric spectroheliograms. Recently, Chatzistergos et al. (2021) and Shin et al. (2020) have used chromospheric observations to reconstruct the solar magnetic activity in the past. Reconstructing past magnetic fields is important for understanding the long-term behavior of the Sun since it is the main factor affecting space climate. Ca II K spectroheliograms are essential for this task, since the first ones were taken already in the 1890s in Europe (Paris and Meudon Observatories, France; Deslandres 1909; Malherbe & Dalmasse 2019) and in the United States (Kenwood Observatory; Hale 1893). Continuous observation campaigns in the Ca II K line started in the early 20th century in Kodaikanal (India), Mount Wilson Observatory (USA), the National Observatory of Japan (Japan), Paris-Meudon Observatory (France), the Arcetri Astrophysical Observatory (Italy), and the Astronomical Observatory of Coimbra University (Portugal) (for a historical overview see, e.g., Bertello et al. 2016; Chatzistergos et al. 2020a). Chromospheric spectroheliograms, together with flux transport simulations, have also been successfully used to study the evolution of large-scale solar magnetic fields, particularly the polar fields, which are only partially observable (Virtanen et al. 2019a).","Citation Text":["Virtanen et al. 2017"],"Functions Text":["In our long-term project aimed at reconstructing the past magnetic activity on the Sun","we also use indirect measurements of magnetic fields obtained, in particular, from chromospheric spectroheliograms."],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[953,973]],"Functions Start End":[[838,924],[991,1106]]}
{"Identifier":"2016ApJ...832..195N__Jin_et_al._2012_Instance_3","Paragraph":"We ignore the density stratification effect in Case I, II, and IIa, because the width of the horizontal current sheet in our simulations is much shorter than the length. The simulation domain extends from x = 0 to x = L0 in the x-direction, and from \n\n\n\n\ny\n=\n\u2212\n0.5\n\n\nL\n\n\n0\n\n\n\n\n to \n\n\n\n\ny\n=\n0.5\n\n\nL\n\n\n0\n\n\n\n\n in the y-direction, in the three cases, with \n\n\n\n\n\n\nL\n\n\n0\n\n\n=\n\n\n10\n\n\n6\n\n\n\n\n m. Outflow boundary conditions are used in the x-direction and inflow boundary conditions in the y-direction. For the inflow boundary conditions, the fluid is allowed to flow into the domain but not to flow out; the gradient of the plasma density vanishes; the total energy is set such that the gradient in the thermal energy density vanishes; a vanishing gradient of parallel components plus divergence-free extrapolation of the magnetic field. For the outflow boundary conditions, the fluid is allowed to flow out of the domain but not to flow in, and the other variables are set by using the same method as the inflow boundary conditions. The horizontal force-free Harris current sheet is used as the initial equilibrium configuration of magnetic fields in Case I,\n13\n\n\n\n\n\n\nB\n\n\nx\n0\n\n\n=\n\u2212\n\n\nb\n\n\n0\n\n\ntanh\n[\ny\n\n\/\n\n(\n0.05\n\n\nL\n\n\n0\n\n\n)\n]\n\n\n\n\n14\n\n\n\n\n\n\nB\n\n\ny\n0\n\n\n=\n0\n\n\n\n\n15\n\n\n\n\n\n\nB\n\n\nz\n0\n\n\n=\n\n\nb\n\n\n0\n\n\n\n\/\n\ncosh\n[\ny\n\n\/\n\n(\n0.05\n\n\nL\n\n\n0\n\n\n)\n]\n.\n\n\nThe magnetic fields in the low solar atmosphere could be very strong (Jin et al. 2009, 2012; Khomenko et al. 2014; Peter et al. 2014; Vissers et al. 2015) and the magnetic field can exceed 0.15 T in both the intranetwork and the network quiet region (e.g., Orozco Su\u00e1rez et al. 2007; Mart\u00ednez Gonz\u00e1lez et al. 2008; Jin et al. 2009, 2012). In the work by Jin et al. 2012, the maximum of the field strength was found to be 0.15 T. The magnetic field could be even stronger in the active region near the sunspot. Therefore, we set b0 = 0.05 T in Case I and Case II, and b0 = 0.15 T in Case IIa. Due to the force-freeness and neglect of gravity, the initial equilibrium thermal pressure is uniform. The initial temperature and plasma density are set as T0 = 4200 K and \u03c10 = 1.66057 \u00d7 10\u22126 kg m\u22123 in Case I, and T0 = 4800 K and \u03c10 = 3.32114 \u00d7 10\u22125 kg m\u22123 in Case II and Case IIa. Therefore, the initial plasma \u03b2 is calculated as \u03b2 \u2243 0.0583 in Case I, \u03b2 \u2243 1.332 in Case II, and \u03b2 \u2243 0.148 in Case IIa. The initial ionization degree is assumed as Yi = 10\u22123 in Case I, and Yi = 1. 2 \u00d7 10\u22124 in Case II and IIa. The magnetic diffusion in this work matches the form computed from the solar atmosphere model in Khomenko & Collados (2012), and we set \n\n\n\n\n\u03b7\n=\n\n[\n\n5\n\u00d7\n\n\n10\n\n\n4\n\n\n\n\n(\n4200\n\n\/\n\nT\n)\n\n\n1.5\n\n\n+\n1.76\n\u00d7\n\n\n10\n\n\n\u2212\n3\n\n\n\n\nT\n\n\n0.5\n\n\n\n\nY\n\n\ni\n\n\n\u2212\n1\n\n\n\n]\n\n\n\n m2 s\u22121 in Case I, and \n\n\n\n\n\u03b7\n=\n[\n5\n\u00d7\n\n\n10\n\n\n4\n\n\n\n\n(\n4800\n\n\/\n\nT\n)\n\n\n1.5\n\n\n\n+\n1.76\n\n\u00d7\n\n\n\n10\n\n\n\u2212\n3\n\n\n\n\nT\n\n\n0.5\n\n\n\n\nY\n\n\ni\n\n\n\u2212\n1\n\n\n]\n\n\n m2 s\u22121 in Case II and IIa. The first part \u223c T\u22121.5 is contributed by collisions between ions and electrons, the second part \n\n\n\n\n\u223c\n\n\nT\n\n\n0.5\n\n\n\n\nY\n\n\ni\n\n\n\u2212\n1\n\n\n\n\n is contributed by collisions between electrons and neutral particles. Small perturbations for both magnetic fields and velocities at t = 0 make the current sheet to evolve and secondary instabilities start to appear later in the three cases. The forms of perturbations are listed below:\n16\n\n\n\n\n\n\nb\n\n\nx\n1\n\n\n=\n\u2212\npert\n\u00b7\n\n\nb\n\n\n0\n\n\n\u00b7\nsin\n\n\n2\n\u03c0\n\n\n\ny\n+\n0.5\n\n\n\nL\n\n\n0\n\n\n\n\n\n\nL\n\n\n0\n\n\n\n\n\n\n\n\u00b7\ncos\n\n\n2\n\u03c0\n\n\n\nx\n+\n0.5\n\n\nL\n\n\n0\n\n\n\n\n\n\nL\n\n\n0\n\n\n\n\n\n\n\n\n\n\n\n17\n\n\n\n\n\n\nb\n\n\ny\n1\n\n\n=\npert\n\u00b7\n\n\nb\n\n\n0\n\n\n\u00b7\ncos\n\n\n2\n\u03c0\n\n\n\ny\n+\n0.5\n\n\nL\n\n\n0\n\n\n\n\n\n\nL\n\n\n0\n\n\n\n\n\n\n\n\u00b7\nsin\n\n\n2\n\u03c0\n\n\n\nx\n+\n0.5\n\n\nL\n\n\n0\n\n\n\n\n\n\nL\n\n\n0\n\n\n\n\n\n\n\n\n\n\n\n18\n\n\n\n\n\n\nv\n\n\ny\n1\n\n\n=\n\u2212\npert\n\u00b7\n\n\nv\n\n\nA\n0\n\n\n\u00b7\nsin\n\n\n\u03c0\n\n\n\ny\n\n\n\n\nL\n\n\n0\n\n\n\n\n\n\n\n\u00b7\n\n\n\n\n\nrandom\n\n\nn\n\n\n\n\nMax\n\n(\n\n\u2223\n\n\nrandom\n\n\nn\n\n\n\u2223\n\n)\n\n\n\n\n,\n\n\nwhere pert = 0.08, vA0 is the initial Alfv\u00e9n velocity, randomn is the random noise function in our code, and \n\n\n\n\nMax\n(\n\u2223\n\n\nrandom\n\n\nn\n\n\n\u2223\n)\n\n\n is the maximum of the absolute value of the random noise function. This random noise function makes the initial perturbations for the velocity in the y-direction to be asymmetric, and such an asymmetry makes the current sheet gradually become more tilted, especially after secondary islands appear. The reconnection process is not really symmetrical in nature (Murphy et al. 2012), this is one of the reasons that we use such a noise function. Another reason is that the asymmetric noise function makes the secondary instabilities develop faster. Figure 1(a) shows the distributions of the current density and magnetic fields at t = 0 in case I.","Citation Text":["Jin et al. 2012"],"Functions Text":["In the work by",", the maximum of the field strength was found to be 0.15 T. The magnetic field could be even stronger in the active region near the sunspot. Therefore, we set b0 = 0.05 T in Case I and Case II, and b0 = 0.15 T in Case IIa."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1676,1691]],"Functions Start End":[[1661,1675],[1691,1913]]}
{"Identifier":"2020AandA...635A..47H__Heckman_et_al._2000_Instance_1","Paragraph":"Finally, the multiphase nature of galactic outflows implies that measurements of the outflow properties based on a single gas phase can lead to misleading conclusions (for a discussion, see e.g., Cicone et al. 2018b). Historically, systematic studies of galactic outflows in nearby and high-z galaxies have focused on the ionized gas \u2013 for example, as observed as broad wing emission in the spectra of the H\u03b1, [O\u202fIII] or Pa\u03b1 lines \u2013 (e.g., Heckman et al. 1990; Rupke & Veilleux 2013a; Woo et al. 2016; Harrison et al. 2016; F\u00f6rster Schreiber et al. 2019; Ramos Almeida et al. 2019) and the atomic phase \u2013 based on the Na D or Mg II lines in absorption (e.g., Heckman et al. 2000; Rupke et al. 2002, 2005; Weiner et al. 2009; Roberts-Borsani & Saintonge 2019). The molecular component of outflows, on the other hand, has been much more difficult to study. Great progress was made with the Herschel Space Observatory using the OH 119 \u03bcm line in absorption to study molecular outflows in Seyfert and luminous infrared galaxies (Fischer et al. 2010; Sturm et al. 2011; Veilleux et al. 2013; Bolatto et al. 2013; Spoon et al. 2013; George et al. 2014; Stone et al. 2016; Gonz\u00e1lez-Alfonso et al. 2017; Zhang et al. 2018). More recently, the advent of powerful millimeter-wave interferometers such as the Atacama Large Millimeter\/submillimeter Array (ALMA) and the NOrthern Extended Millimeter Array (NOEMA) are rapidly increasing the number of molecular outflows detected based on observations of the CO line (e.g., Combes et al. 2013; Sakamoto et al. 2014; Garc\u00eda-Burillo et al. 2014; Leroy et al. 2015; Feruglio et al. 2015; Morganti et al. 2015; Dasyra et al. 2016; Pereira-Santaella et al. 2016, 2018; Veilleux et al. 2017; Fluetsch et al. 2019; Lutz et al. 2020). At high-z, so far only a handful of large-scale, molecular outflows have been studied in QSOs (e.g., Cicone et al. 2015; Vayner et al. 2017; Feruglio et al. 2017; Carniani et al. 2017; Fan et al. 2018; Brusa et al. 2018), sub-millimeter galaxies (e.g., Spilker et al. 2018), and main-sequence, star-forming galaxies (e.g., Herrera-Camus et al. 2019).","Citation Text":["Heckman et al. 2000"],"Functions Text":["Historically, systematic studies of galactic outflows in nearby and high-z galaxies have focused on","and the atomic phase \u2013 based on the Na D or Mg II lines in absorption (e.g.,"],"Functions Label":["Background","Background"],"Citation Start End":[[659,678]],"Functions Start End":[[218,317],[582,658]]}
{"Identifier":"2022ApJ...940...72R__Camilo_et_al._2006_Instance_2","Paragraph":"Several studies have discussed the radio luminosity of GLEAM-X J1627 during its radio outburst in comparison with the limits of its rotational energy (Erkut 2022; Hurley-Walker et al. 2022). In particular, assuming isotropic emission, the radio luminosity of the brightest single peaks (L\nradio \u223c 1030\u20131031 erg s\u22121; Hurley-Walker et al. 2022) exceeds the limits on the rotational power of the source by a few orders of magnitude. Figure 6 shows those peak radio luminosities and the rotational power of GLEAM-X J1627 in comparison with other pulsars, rotating radio transients (RRATs) and radio-loud magnetars. For the radio-loud magnetars, given their large variability, we have chosen the brightest radio pulses reported in the literature (data collected from Camilo et al. 2006, 2007; Weltevrede et al. 2011; Deller et al. 2012; Lynch et al. 2015; Majid et al. 2017; Pearlman et al. 2018; Lower et al. 2020, and Esposito et al. 2021). It is well known that assuming isotropic radio emission is not realistic, and a beaming factor necessarily has to be present (see, e.g., Erkut 2022). However, the relation between the duty cycle and the spin period of canonical pulsars has a large spread (Manchester et al. 2005). Moreover, it is observed that this relationship does not apply to radio-loud magnetars, which in general show larger duty cycles than what one would expect from the extrapolation of this tentative relation for radio pulsars to magnetars (see, e.g., Camilo et al. 2006, 2007). To avoid the uncertainty of beaming models, which for magnetars are mostly unknown even theoretically, we plotted the isotropic radio luminosity for all the different pulsar classes in Figure 6. From this plot, at variance with canonical radio pulsars, we see how the brightest single peaks for radio-loud magnetars might exceed their rotational powers, in line with what is possibly observed for GLEAM-X J1627. While not resolving uncertainties related to the exact mechanism of radio emission or the beaming factor, Figure 6 shows that, under the assumption of isotropic emission, even for magnetars the brightest single peaks exceed their rotational energy budget. Considering all the uncertainties in the assumptions used to derive the radio luminosities plotted in Figure 6, GLEAM-X J1627\u2019s radio luminosity excess over its rotational power cannot be used as an argument for or against its neutron star nature.","Citation Text":["Camilo et al. 2006"],"Functions Text":["Moreover, it is observed that this relationship does not apply to radio-loud magnetars, which in general show larger duty cycles than what one would expect from the extrapolation of this tentative relation for radio pulsars to magnetars (see, e.g.,"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1468,1486]],"Functions Start End":[[1219,1467]]}
{"Identifier":"2019AandA...631A..88Y__Bohren_&_Huffman_(1998)_Instance_1","Paragraph":"Starting from the four aforementioned materials, we consider several composition mixtures and grain structures. For the sake of comparison, we first consider compact grains of purely a-Sil, a-C, or a-C:H. Subsequently, according to K\u00f6hler et al. (2015), we consider compact grains made of two thirds a-Sil and one third a-C (Mix 1) or one third a-C:H (Mix 2), in terms of volume fractions. These allow reproduction of the mass fractions derived by Jones et al. (2013) for the diffuse ISM. The effect of porosity is tested for the Mix 1 mixture, with a porosity degree of 50% (Mix 1:50). We also evaluate theimpact of the presence of a water ice mantle on compact Mix 1 grains (Mix 1:ice). We further consider two material compositions defined in Pollack et al. (1994) based on depletion measurements: (i) 21% a-Sil and 79% a-C (Mix 3); and (ii) 8% a-Sil, 30% a-C, and 62% water ice (Mix 3:ice). The various grain compositions are summarised in Table 1. For each grain composition, we derive the absorption and scattering efficiencies Qabs and Qsca, respectively, and the asymmetry factor of the phase function g = \u27e8cos\u03b8\u27e9. To allow fast calculations, we make the major assumption that the grains are spherical and compute their optical properties using the Mie theory (Mie 1908; Bohren & Huffman 1983) with the Fortran 90 version of the BHMIE routine given in Bohren & Huffman (1998). For grains consisting of two or three materials, we first derive effective optical constants following the Maxwell Garnett mixing rule (Maxwell Garnett 1904; Bohren & Huffman 1998). Indeed, we assume that in Mix 1 grains, for example, carbon appears as proper inclusions in the silicate matrix rather than assuming a completely random inhomogeneous medium. Mishchenko et al. (2016a,b) performed exhaustive studies of the applicability of the Maxwell Garnett mixing rule to heterogeneous particles. These latter authors showed that this rule can provide accurate estimates of the scattering matrix and absorption cross-section of heterogeneous grains at short wavelengths (typically up to the visible for a 0.1 \u03bcm grain and to the mid-infrared(MIR) for a 10 \u03bcm grain) if twocriteria are met: both the size parameter of the inclusions and the refractive index contrast between the host material and the inclusions have to be small. Moreover, Mishchenko et al. (2016a) demonstrated that the extinction and asymmetry-parameter errors of the Maxwell Garnett mixing rule are significantly smaller than the scattering-matrix errors, remaining small enough for most typical applications and in particular the kind of applications we perform here. It is however well known that this kind of mixing rule systematically underestimates the absorption efficiency in the FIR to millimetre wavelength range, the implications of which are discussed in Sect. 3.2. We perform our computations with the emc routine of V. Ossenkopf3. For Mix 1 and Mix 2, we assume a matrix of a-Sil with inclusions of a-C or a-C:H, and for Mix 3 a matrix of a-C with inclusions of a-Sil. For grains surrounded by an ice mantle, the optical properties are derived with the core-mantle Mie theory using the BHCOAT routine given in Bohren & Huffman (1998).","Citation Text":["Bohren & Huffman (1998)"],"Functions Text":["To allow fast calculations, we make the major assumption that the grains are spherical and compute their optical properties using the Mie theory","with the Fortran 90 version of the BHMIE routine given in"],"Functions Label":["Uses","Uses"],"Citation Start End":[[1359,1382]],"Functions Start End":[[1122,1266],[1301,1358]]}
{"Identifier":"2022AandA...667A..35R__Nakajima_&_Ouchi_(2014)_Instance_1","Paragraph":"In the optical range, several proxies can serve as indicators of LyC leakage. In particular, line ratios involving ions with different ionization potentials, produced at different depths, have been proposed as indicators to discriminate between radiation-bounded and density-bounded H\u202fII regions. Radiation-bounded regions correspond to ionized spheres delimited by their Str\u00f6mgren radii set by the equilibrium between production of photons by stars and ionization of the surrounding gas. Density-bounded regions are instead delimited by the lack of matter, which sets their outer radius before the Str\u00f6mgren radius. Hence, they allow part of the LyC-photons produced by stars to escape from H\u202fII regions. The oxygen line ratio [O\u202fIII]\u03bb 5007 \u00c5\/[O\u202fII] \u03bb\u03bb3726, 3728 \u00c5 (O32) proposed by Jaskot & Oey (2013) and Nakajima & Ouchi (2014) was successfully used to select LyC-leaking candidates but no strong correlation was found with the measured values of escape fraction (see Izotov et al. 2018b; Naidu et al. 2018; Bassett et al. 2019; Nakajima et al. 2020, and discussions therein). Based on a similar idea, the lack of emission from ions with low ionization potentials like [S\u202fII] \u03bb\u03bb6716,6731 \u00c5 has also been proposed to target leaking candidates (Wang et al. 2019, 2021; Katz et al. 2020). This lack of emission of some low ionization species was first interpreted as the signature of a density-bounded galaxy where the outer part of H\u202fII regions were completely stripped out. However, using simple photoionization models, Stasi\u0144ska et al. (2015) have shown that on average galaxies with high O32 cannot have massive escapes of ionizing photons, since low ionization lines like [O\u202fI]6300 \u00c5 are often also detected in these galaxies, implying the presence of radiation-bounded regions. Subsequently Plat et al. (2019) and Ramambason et al. (2020) noted that several strong LyC-emitters show surprisingly strong [O\u202fI]6300 \u00c5 emission, and proposed several explanations. While Stasi\u0144ska et al. (2015) and Plat et al. (2019) suggested that such emission could be powered by the presence of AGN or radiative-shocks, we proposed in Ramambason et al. (2020) a 2-component model combining both density- and ionization-bounded regions.","Citation Text":["Nakajima & Ouchi (2014)"],"Functions Text":["The oxygen line ratio [O\u202fIII]\u03bb 5007 \u00c5\/[O\u202fII] \u03bb\u03bb3726, 3728 \u00c5 (O32) proposed by Jaskot & Oey (2013) and","was successfully used to select LyC-leaking candidates"],"Functions Label":["Uses","Uses"],"Citation Start End":[[808,831]],"Functions Start End":[[706,807],[832,886]]}
{"Identifier":"2020MNRAS.492.4975M__Iben_1967_Instance_1","Paragraph":"We emphasize that 7Be\/H remains larger by at least one order of magnitude than predicted by nova models (Starrfield et al. 1978; Hernanz et al. 1996; Jos\u00e9 & Hernanz 1998). We note, however, that the final amount of 7Be is sensitive to the amount of 3He in the donor star as a higher abundance of 3He is expected to result in a higher 7Be abundance. Boffin et al. (1993) and Hernanz et al. (1996) found a logarithmic dependence of the 7Be output to the initial 3He abundance. The non-linearity of 7Be yields results from 3He(3He,\u20092p)4He and its importance increases as the square of the initial 3He abundance. Therefore, it produces a leaking of the available 3He for the 3He(\u03b1,\u2009\u03b3) 7Be, whose rate increases only linearly for the initial 3He abundance. For 3He enhancements up to 100 solar, Boffin et al. (1993) derive X(7Be)\/X(7Be0) = 1 + 1.5log\u2009X(3He)\/X(3He\u2299), where 7Be0 is the 7Be final mass fraction obtained with a solar initial 3He mass fraction. From the theoretical point of view, it is believed that low-mass main-sequence stars synthesize 3He through the p\u2013p chains with peak abundances of few 10\u22123 by number (Iben 1967). As the star ascends, the red giant branch convection dredges up 3He-enriched material to the surface, which is later expelled into the interstellar medium by wind or during the planetary phase. 3He is a particularly difficult element to measure. It can be measured in H\u2009ii regions by using measurements at a frequency of 8.665 GHz (i.e. 3.46 cm), which is emitted naturally by ionized 3He (Bania, Rood & Balser 2010; Balser & Bania 2018) or in the stellar atmospheres of hot stars (Geier et al. 2012). Surprisingly, interstellar medium observations indicate that there is far less of this element in the Galaxy than the current models predict. In order not to overproduce 3He in the course of chemical evolution, it has become customary to assume that some unknown 3He-destruction mechanism is at work in low-mass giants (Dearborn, Steigman & Tosi 1996; Galli et al. 1997; Chiappini, Renda & Matteucci 2002; Romano & Matteucci 2003). For instance, Charbonnel & Zahn (2007) suggested a thermohaline mixing during the red giant branch phase of low-mass stars. However, in a few planetary nebulae, 3He\/H is found to be high at the level of 10\u22123, consistent with predictions from standard stellar models (Rood, Bania & Wilson 1992; Balser & Bania 2018).","Citation Text":["Iben 1967"],"Functions Text":["From the theoretical point of view, it is believed that low-mass main-sequence stars synthesize 3He through the p\u2013p chains with peak abundances of few 10\u22123 by number"],"Functions Label":["Background"],"Citation Start End":[[1120,1129]],"Functions Start End":[[953,1118]]}
{"Identifier":"2022AandA...663A.110M__Tschudi_&_Schmid_2021_Instance_1","Paragraph":"The presented model calculations provide two-dimensional images for the scattered intensity I(x, y), azimuthal polarization Q\u03c6(x, y), and other radiation parameters. It is useful for the comparison with observations to deduce disk-integrated radiation parameters from these model images that are scaled to the stellar intensity, as in ${{\\bar I\\left( i \\right)} \\mathord{\\left\/ {\\vphantom {{\\bar I\\left( i \\right)} {{I_ \\star }}}} \\right. \\kern-\\nulldelimiterspace} {{I_ \\star }}}{{,{{\\bar Q}_\\varphi }\\left( i \\right)} \\mathord{\\left\/ {\\vphantom {{,{{\\bar Q}_\\varphi }\\left( i \\right)} {{I_ \\star }}}} \\right. \\kern-\\nulldelimiterspace} {{I_ \\star }}}$I\u00af(i)\/I\u22c6,Q\u00af\u03c6(i)\/I\u22c6, the disk-averaged fractional polarization \u2329p\u03c6\u232a, or quadrant polarization values Qxxx(i)\/I\u22c6 or Uxxx(i)\/I\u22c6 which can also be deduced from observations without introducing significant ambiguities by the diversity of disk morphologies. We note that the measured values must be corrected for instrumental effects, in particular for the signal convolution with the instrumental PSF, which can introduce significant cancellation of the disk signal (Schmid et al. 2006; Avenhaus et al. 2014, Avenhaus et al. 2017; Heikamp & Keller 2019; Tschudi & Schmid 2021). Unfortunately, available observational data with high accuracy and quantified uncertainties are still very limited and often only some of the radiation parameters can be determined for a given disk. Therefore, the derivation of disk parameters from the comparison with model results can be rather ambiguous. The model results presented in this work allow us to carry out a detailed investigation of the important parameter ambiguities involved in the interpretation of observations of transition disks and provide diagnostic relations that can be used to constrain key model parameters for different types of observational data. For example, Fig. B.1 shows that, for a given inclination $i,{{\\bar I} \\mathord{\\left\/ {\\vphantom {{\\bar I} {{I_ \\star }}}} \\right. \\kern-\\nulldelimiterspace} {{I_ \\star }}}{{,{{\\bar Q}_\\varphi }} \\mathord{\\left\/ {\\vphantom {{,{{\\bar Q}_\\varphi }} {{I_ \\star }}}} \\right. \\kern-\\nulldelimiterspace} {{I_ \\star }}}$i,I\u00af\/I\u22c6,Q\u00af\u03c6\/I\u22c6, and ${{\\bar Q} \\mathord{\\left\/ {\\vphantom {{\\bar Q} {{I_ \\star }}}} \\right. \\kern-\\nulldelimiterspace} {{I_ \\star }}}$Q\u00af\/I\u22c6 all strongly correlate with the angular wall height \u03b1 and single scattering albedo \u03c9 , and anti-correlate with wall slope \u03c7 and scattering asymmetry parameter \u0261. In addition, the polarization parameters ${{{{\\bar Q}_\\varphi }} \\mathord{\\left\/ {\\vphantom {{{{\\bar Q}_\\varphi }} {{I_ \\star }}}} \\right. \\kern-\\nulldelimiterspace} {{I_ \\star }}}$Q\u00af\u03c6\/I\u22c6 and ${{\\bar Q} \\mathord{\\left\/ {\\vphantom {{\\bar Q} {{I_ \\star }}}} \\right. \\kern-\\nulldelimiterspace} {{I_ \\star }}}$Q\u00af\/I\u22c6 also correlated with pmax. If only the scattered intensity ${{\\bar I} \\mathord{\\left\/ {\\vphantom {{\\bar I} {{I_ \\star }}}} \\right. \\kern-\\nulldelimiterspace} {{I_ \\star }}}$I\u00af\/I\u22c6 or only the polarized intensity is measured then it is almost impossible to constrain individual disk parameters without additional information. Therefore, it is important to select and obtain more observational information and better diagnostic parameters with more diagnostic power.","Citation Text":["Tschudi & Schmid 2021"],"Functions Text":["We note that the measured values must be corrected for instrumental effects, in particular for the signal convolution with the instrumental PSF, which can introduce significant cancellation of the disk signal"],"Functions Label":["Uses"],"Citation Start End":[[1202,1223]],"Functions Start End":[[905,1113]]}
{"Identifier":"2015MNRAS.453.3414A__Filippenko_&_Chornock_2001_Instance_1","Paragraph":"Filippenko & Chornock (2001) first presented the dynamical estimate of mass of the source to be around 7.4 \u00b1 1.1\u2009M\u2299. Recently, Radhika & Nandi (2014) claimed that the mass of XTE J1859+226 is perhaps in between 6.58 and 8.84\u2009M\u2299 which is similar to the prediction of Shaposhnikov & Titarchuk (2009), although the lower mass limit is estimated as 5.4\u2009M\u2299 by Corral-Santana et al. (2011). However, we consider the typical mass of the source as 7\u2009M\u2299. The distance of this source is around d \u223c 11\u2009kpc (Filippenko & Chornock 2001). Steiner et al. (2013) measured the spin as ak \u223c 0.4; however, Motta et al. (2014b) recently reported that the spin of the source is ak \u223c 0.34. Since the spin predictions are quite close, we use ak \u223c 0.4 for this analysis. We estimate the fluxes Fx (see Table 1) of LHS and HIMS of the 1999 outburst of the source (Radhika & Nandi 2014). The corresponding disc luminosities are calculated as $L_{\\rm disc}^{{\\rm LHS}}=8.26 \\times 10^{37}\\ {\\rm erg\\ s^{-1}}$ and $L_{\\rm disc}^{{\\rm HIMS}}=1.85 \\times 10^{38}\\ {\\rm erg\\ s^{-1}}$, respectively. Now, it is reasonable to assume the accretion efficiency for rotating BH as \u03b7 = 0.3 which corresponds to the accretion rate of the inflowing matter as ${\\dot{M}}_{{\\rm acc}}^{{\\rm LHS}} = 0.304 {\\dot{M}}_{{\\rm Edd}}$ in LHS and ${\\dot{M}}_{{\\rm acc}}^{{\\rm HIMS}} = 0.680 {\\dot{M}}_{{\\rm Edd}}$ in HIMS. For LHS, we use $R_{\\dot{m}}=9.83$\u2009per\u2009cent following our theoretical estimate where xs = 64.6rg for ak = 0.4, ${\\mathcal {E}}=0.001\\,98$ and \u03bb = 3.18. Incorporating these inputs in equation (15), we obtain the jet kinetic power as $L^{{\\rm LHS}}_{{\\rm jet}} = 2.52\\times 10^{37}\\ {\\rm erg\\ s^{-1}}$. The maximum mass outflow rate for HIMS corresponding to ak = 0.4 is obtained from Fig. 9 as $R^{\\rm max}_{\\dot{m}}=17.5$\u2009per\u2009cent for ${\\mathcal {E}}=0.005\\,47$ and \u03bb = 3.1, where the shock transition occurs at 21.9rg. Using these values in equation (15), we obtain the maximum jet kinetic power as $L^{{\\rm HIMS}}_{{\\rm jet}} = 1.08\\times 10^{38}\\ {\\rm erg\\ s^{-1}}$ which we regard to be associated with the HIMS of this source.","Citation Text":["Filippenko & Chornock (2001)"],"Functions Text":["first presented the dynamical estimate of mass of the source to be around 7.4 \u00b1 1.1\u2009M\u2299."],"Functions Label":["Uses"],"Citation Start End":[[0,28]],"Functions Start End":[[29,116]]}
{"Identifier":"2022MNRAS.514.5570N__in_1928_Instance_1","Paragraph":"In this study, we examine the stellar populations of NGC 5053 using the UVIT, on-board the AstroSat. NGC 5053 is a galactic GC that lies in the northern constellation of Coma Berenices, having Galactic coordinates l = 335${_{.}^{\\circ}}$70, b = 78${_{.}^{\\circ}}$95. This cluster is located at a distance of 17.54 \u00b1 0.23 kpc (Baumgardt & Vasiliev 2021) and the metallicity is estimated to be [Fe\/H] = \u22122.27 (Harris 2010). It is, therefore, placed amongst the metal-poor galactic GCs. The tidal radius of this cluster has been estimated as rt = 15.2 \u00b1 3.3\u2009arcmin by de Boer et al. (2019). When it was discovered in 1784, it was not classified as a GC because of its appearance and structure; there was no densely packed nucleus, and the central region was resolvable (Herschel 1786). These observations placed this GC in a grey-area between globular and open clusters. It was only in 1928 that Baade first classified it as a GC owing to its high latitude, richness in faint stars and presence of variable stars (Baade 1928). The variable stars in this cluster include RR Lyrae variables and SX Phoenicis (SX Phe) stars (Sawyer 1946; Nemec, Mateo & Schombert 1995a; Nemec et al. 1995b; Nemec 2004). The cluster is known to have a HB that is extended predominantly towards the blue flank of the RR Lyrae instability strip, which is an archetypal feature in many metal-poor GCs (Sarajedini & Milone 1995). A rich population of BSSs has been identified in this cluster (Sarajedini & Milone 1995). The presence of relatively large number of BSSs in such low density GCs has been examined to propose alternate formation mechanisms of BSSs (Leonard & Fahlman 1991). The cluster has been examined in the UV regime by Schiavon et al. (2012) using photometric data from GALEX, and they have published a catalogue of the UV bright stars present in this cluster. The location of this cluster in the sky and its velocity have given rise to a debate on whether it belongs to the Milky Way or Sgr dSph (e.g. Law & Majewski 2010; Boberg, Friel & Vesperini 2015; Sbordone et al. 2015; Tang et al. 2018).","Citation Text":["Baade 1928"],"Functions Text":["It was only in 1928 that Baade first classified it as a GC owing to its high latitude, richness in faint stars and presence of variable stars"],"Functions Label":["Background"],"Citation Start End":[[1011,1021]],"Functions Start End":[[868,1009]]}
{"Identifier":"2021MNRAS.506.1258W__Borsa_et_al._2021_Instance_1","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"],"Functions Text":["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"],"Functions Label":["Motivation"],"Citation Start End":[[580,597]],"Functions Start End":[[324,506]]}
{"Identifier":"2019MNRAS.482.5651M__Schweizer_&_Middleditch_1980_Instance_2","Paragraph":"Therefore, the kinetics characteristics of the star could be the only piece to judge whether or not the SM star is the surviving companion of SN 1006. If its space velocity is significantly different from the other stars in the remnant of SN 1006, the probability to be the surviving companion would become high. Otherwise, the probability becomes low. We check the proper motion of the stars within 5\u2009arcmin of the remnant centre from Gaia DR2, as shown in Fig. 19. From the figure, it seems that there is not difference between the SM star and other stars in the remnant of SN 1006 in the aspect of proper motion, i.e. the proper motion of the SM star only slightly deviates from the median value of the proper motions of the stars at the direction of the SNR centre of SN 1006, and such a proper motion disfavours the SM star as the surviving companion of SN 1006 (Schweizer & Middleditch 1980; Burleigh et al. 2000). So, a 3D space velocity is helpful to judge the nature of the SM star. However, unfortunately, some data of the SM star in Gaia DR2 are so uncertain that we cannot use them to constrain its 3D space velocity, otherwise we could obtain a complete wrong conclusion.5 For example, the parallax of the SM star is \u03d6 = 0.0736 \u00b1 0.1244, and then \u03c3\u03d6\/\u03d6 = 1.69 which is much larger than the threshold value of 0.2 for distance estimation from GAIA DR2 data (Astraatmadja & Bailer-Jones 2016; Katz et al. 2018). The distance of the SM star from this parallax is much larger than all the previous measurements from spectrum by at least a factor of 2 (see summary in Burleigh et al. 2000). Considering that some other astrometric measurements of the SM star are also very uncertain, we applied the measurements in the previous literatures as the distance of the SM star. Based on a radial velocity of $-13\\pm 17\\, {\\rm km^{\\rm -1}}$ and a distance of 2.07 \u00b1 0.18 kpc (Schweizer & Middleditch 1980; Winkler et al. 2003; Kerzendorf et al. 2018), we can calculate the UVW velocities of the SM star, i. e. $U=-5.2\\pm 14\\, {\\rm km^{\\rm -1}}$, $V=197\\pm 10\\, {\\rm km^{\\rm -1}}$, and $W=3.1\\pm 5\\, {\\rm km^{\\rm -1}}$. The V value of the SM star is smaller than that of a normal disc star. We then transform these velocities into the Galactic rotational velocity at a Galactocentric distance of \u223c6.67 kpc, i.e. $V_{\\rm c}=196\\pm 12\\, {\\rm km^{\\rm -1}}$, which is smaller than the Galactic rotational velocity of the disc stars at the Galactocentric distance by $50\\pm 19\\, {\\rm km^{\\rm -1}}$ (Huang et al. 2016). This velocity difference is marginally consistent with the predicted orbital velocity here (see Fig. 7). In addition, the smaller rotational velocity of the SM star may explain its small proper motion shown in Fig. 19. So, the SM star is still possible to be the surviving companion of SN 1006.","Citation Text":["Schweizer & Middleditch 1980"],"Functions Text":["Based on a radial velocity of $-13\\pm 17\\, {\\rm km^{\\rm -1}}$ and a distance of 2.07 \u00b1 0.18 kpc",", we can calculate the UVW velocities of the SM star, i. e. $U=-5.2\\pm 14\\, {\\rm km^{\\rm -1}}$, $V=197\\pm 10\\, {\\rm km^{\\rm -1}}$, and $W=3.1\\pm 5\\, {\\rm km^{\\rm -1}}$."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1876,1904]],"Functions Start End":[[1779,1874],[1950,2118]]}
{"Identifier":"2016MNRAS.463..382U__Zdziarski_et_al._2000_Instance_1","Paragraph":"Active galactic nuclei (AGN) are thought to be powered by an accretion disc around a supermassive black hole, mostly emitting in the optical\/ultraviolet (UV) band. According to the standard paradigm, the X-ray emission is due to thermal Comptonization of the soft disc photons in a hot region, the so-called corona (Haardt & Maraschi 1991; Haardt, Maraschi & Ghisellini 1994, 1997). This process explains the power-law shape of the observed X-ray spectrum of AGN. A feature of thermal Comptonization is a high-energy cut-off, which has been observed around \u223c100 keV in several sources thanks to past observations with Compton Gamma-Ray Observatory (CGRO)\/Oriented Scintillation Spectrometer Experiment (OSSE; Zdziarski, Poutanen & Johnson 2000), BeppoSAX (Perola et al. 2002) and more recently Swift\/Burst Alert Telescope (BAT; Baumgartner et al. 2013) and INTEGRAL (Malizia et al. 2014; Lubi\u0144ski et al. 2016). From the application of Comptonization models, such high-energy data allow to constrain the plasma temperature, which is commonly found to range from 50 to 100 keV (e.g. Zdziarski et al. 2000; Lubi\u0144ski et al. 2016). Furthermore, the cut-off energy is now well constrained in an increasing number of sources thanks to the unprecedented sensitivity of NuSTAR up to \u223c80 keV (e.g. Ballantyne et al. 2014; Brenneman et al. 2014; Marinucci et al. 2014; Balokovi\u0107 et al. 2015; Matt et al. 2015; Ursini et al. 2015). The primary X-ray emission can be modified by different processes, such as absorption from neutral or ionized gas (the so-called warm absorber), and Compton reflection from the disc (e.g. George & Fabian 1991; Matt, Perola & Piro 1991) or from more distant material, like the molecular torus at pc scales (e.g. Matt, Guainazzi & Maiolino 2003). A smooth rise below 1\u20132 keV above the extrapolated high-energy power law is commonly observed in the spectra of AGN (see e.g. Bianchi et al. 2009). The origin of this so-called soft excess is uncertain (see e.g. Done et al. 2012). Ionized reflection is able to explain the soft excess in some sources (e.g. Crummy et al. 2006; Ponti et al. 2006; Walton et al. 2013), while Comptonization in a \u2018warm\u2019 region is favoured in other cases (e.g. Mehdipour et al. 2011; Boissay et al. 2014).","Citation Text":["Zdziarski et al. 2000"],"Functions Text":["From the application of Comptonization models, such high-energy data allow to constrain the plasma temperature, which is commonly found to range from 50 to 100 keV"],"Functions Label":["Background"],"Citation Start End":[[1081,1102]],"Functions Start End":[[911,1074]]}
{"Identifier":"2021AandA...650A.205V__Hippke_&_Heller_2019_Instance_1","Paragraph":"To search for transit events, we will make use of our custom pipeline SHERLOCK (Pozuelos et al. 2020)5. This pipeline provides the user with easy access to Kepler, K2, and TESS data for both SC and LC. The pipeline searches for and downloads the pre-search data conditioning simple aperture (PDC-SAP) flux data from the NASA Mikulski Archive for Space Telescope (MAST). Then, it uses a multi-detrend approach in the WOTAN package (Hippke et al. 2019), whereby the nominal PDC-SAP flux light curve is detrended several times using a biweight filter or a Gaussian process, by varying the window size or the kernel size, respectively. This multi-detrend approach is motivated by the associated risk of removing transit signals, in particular, short and shallow signals. Each of the new detrended light curves, jointly with the nominal PDC-SAP flux, is then processed through the transit least squares (TLS) package (Hippke & Heller 2019) in the search for transits. In contrast to the classical box least-squares (BLS) algorithm (Kov\u00e1cs et al. 2002), the TLS algorithm uses an analytical transit model that takes the stellar parameters into account. Then, it phase folds the light curves over a range of trial periods (P), transit epochs (T0), and transit durations (d). It then computes the \u03c72 between the model and the observed values, searching for the minimum \u03c72 value in the 3D parameter space (P, T0, and d). The TLS algorithm has been found to be more reliable than the classical BLS in finding any type of transiting planet, and it is particularly well suited for the detection of small planets in long time series, such as those coming from Kepler, K2, and TESS. The TLS algorithm also allows the user to easily fine-tune the parameters to optimize the search in each case, which is particularly interesting for shallow transits. In addition, SHERLOCK incorporates a vetting module that combines the TPFplotter (Aller et al. 2020), LATTE (Eisner et al. 2020), and TRICERATOPS (Giacalone et al. 2021) packages, which allows the user to explore any contamination source in the photometric aperture used, momentum dumps, background flux variations, x\u2013y centroid positions, aperture size dependences, flux in-and-out transits, each individual pixel of the target pixel file, and to estimate the probabilities for different astrophysical scenarios such as transiting planet, eclipsing binary, and eclipsing binary with twice the orbital period. Collectively, these analyses help the user estimate the reliability of a given detection.","Citation Text":["Hippke & Heller 2019"],"Functions Text":["Each of the new detrended light curves, jointly with the nominal PDC-SAP flux, is then processed through the transit least squares (TLS) package","in the search for transits.","In contrast to the classical box least-squares (BLS) algorithm","the TLS algorithm uses an analytical transit model that takes the stellar parameters into account."],"Functions Label":["Uses","Uses","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[913,933]],"Functions Start End":[[767,911],[935,962],[963,1025],[1048,1146]]}
{"Identifier":"2021AandA...647A..49S__P\u00e9rez_et_al._2019_Instance_1","Paragraph":"It has been claimed that, for solar-type stars, CA models could face some difficulties (in principle) explaining the presence of giant planets around metal-poor stars, or massive planets at long radial distances (see e.g., Helled et al. 2014). In particular, HR 8799 is a metal-poor A5 star included in our sample, hosting three giant planets orbiting beyond 10 AU (Marois et al. 2008). Then, some authors proposed that the planets orbiting HR 8799 are likely formed by GI, given that GI models can take place at large radii and in low metal environments (Marois et al. 2008; Dodson-Robinson et al. 2009; Meru & Bate 2010). Other works cast some doubts about the planet formation around HR 8799, showing that CA models are also possible (Currie et al. 2011) or even a combination of GI and CA (Marois et al. 2010). For the case of \u03b2 Pictoris and HD 169142, also having giant planets at long distances, CA models seem to be favored (GRAVITY Collaboration 2020; Nowak et al. 2020; P\u00e9rez et al. 2019). However, an important word of caution about these works is in order. Some of the stars mentioned (HR 8799 and HD 169142) display (superficial) metal-poor abundances, showing in fact a \u03bb Bo\u00f6tis pattern (see for example Fig. 8). As previously mentioned, the most accepted idea about the origin of this peculiar signature, suppose a solar-like composition for the original molecular cloud where the stars born, and then some kind of selective accretion to obtain a \u03bb Bo\u00f6tis pattern. In this way, it would be not entirely appropriate to assume a metal-poor natal environment for stars like HR 8799, as assumed by some works to support a GI planet formation (e.g., Meru & Bate 2010). This fact was early noted by Paunzen et al. (2014), in their comparison of \u03bb Bo\u00f6tis stars and Population II type stars. Numerical simulations of planet formation around \u03bb Bo\u00f6tis stars should assume a solar-like composition (rather than a metal-poor natal environment), and this could have important consequences for the subsequent results.","Citation Text":["P\u00e9rez et al. 2019"],"Functions Text":["For the case of \u03b2 Pictoris and HD 169142, also having giant planets at long distances, CA models seem to be favored"],"Functions Label":["Background"],"Citation Start End":[[979,996]],"Functions Start End":[[815,930]]}
{"Identifier":"2018MNRAS.473.4130M__Martin_&_Drissen_2016_Instance_1","Paragraph":"The flux calibration was performed from two calibration sources: (1) the spectrum of the spectrophotometric standard star GD71 (Bohlin 2003), obtained in 2016 January, which is used to eliminate as much as possible any strong wavelength dependence; (2) the median combination of a set of 10 images of the standard star HZ\u20094 (Bohlin, Dickinson & Calzetti 2001) obtained right after the end of the cube observation with photometric conditions similar to the observation conditions. The exact value of the interferometer\u2019s modulation efficiency (ME), which acts essentially as an additional throughput loss, is the major source of uncertainty on the absolute flux calibration (Martin & Drissen 2016). Interferometric images of the laser source have been obtained before and after the observation of the target in order to measure the variation of ME at the calibration laser wavelength (543.5\u2009nm) with respect to its nominal value (85\u2009per\u2009cent). We have measured a loss of 11.7\u2009per\u2009cent. The initial flux calibration of M\u200931 in the SN3 filter has been corrected for this loss (thus multiplied by a factor 1.13). But given the possible uncertainty on this estimate, we have double-checked it using Hubble Space Telescope (HST) narrow-band images of the target. The advantage of HST\u2019s narrow-band filters is that they can easily be simulated by integrating the spectral cube over the filter\u2019s well-known transmission curve. As SITELLE\u2019s cube flux is expressed in erg\u2009cm\u22122\u2009s\u22121\u2009\u00c5\u22121 and given the filter transmission curve F(\u03c3), the integrated flux $\\tilde{\\phi }_{\\rm F}$, expressed in erg\u2009cm\u22122\u2009s\u22121\u2009\u00c5\u22121, is\n(15)\r\n\\begin{equation}\r\n\\tilde{\\phi }_{\\rm F} = \\frac{\\int _{-\\infty }^{+\\infty }\\phi (\\sigma )F(\\sigma ){\\rm d}\\sigma }{\\int _{-\\infty }^{+\\infty }F(\\sigma ){\\rm d}\\sigma }\\, .\r\n\\end{equation}\r\nIf one converts the flux in terms of surface brightness, expressed in erg\u2009cm\u22122\u2009s\u22121\u2009\u00c5\u22121 per HST pixel surface (SHSTCamera), i.e.\n(16)\r\n\\begin{equation}\r\n\\tilde{B}_{\\rm F} = \\frac{\\tilde{\\phi }_F }{S_{\\rm SITELLE}}S_{\\rm HSTCamera},\r\n\\end{equation}\r\nwith SSITELLE = 0.322\u2009arcsec2, the image $\\tilde{B}_{\\rm F}$, once properly aligned, can be directly compared with the HST frame. We have made this comparison for three different regions of the FOV (called WFC3, ACS1 and ACS2, shown in Fig. 9) and four different filters (F656N, F658N, F665N and F660N). Histograms of the flux ratio between the integrated frames and the HST frames are presented in Fig. 10 and the first moments of their distributions are listed in Table 3.","Citation Text":["Martin & Drissen 2016"],"Functions Text":["The exact value of the interferometer\u2019s modulation efficiency (ME), which acts essentially as an additional throughput loss, is the major source of uncertainty on the absolute flux calibration"],"Functions Label":["Background"],"Citation Start End":[[674,695]],"Functions Start End":[[480,672]]}
{"Identifier":"2015ApJ...799..149J___2014_Instance_2","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 al (2014)"],"Functions Text":["We also show as a gray band the best fit for the stellar fraction in the form of stars determined by","in a study of a large sample of lens galaxies using strong lensing and photometry,"],"Functions Label":["Uses","Uses"],"Citation Start End":[[937,955]],"Functions Start End":[[836,936],[956,1038]]}
{"Identifier":"2018MNRAS.475.1160H__Tumlinson_et_al._2013_Instance_1","Paragraph":"Galaxies are surrounded by vast gaseous haloes which extend well beyond the hosts\u2019 stellar components: Early observations of quasar sight lines attributed the presence of absorption at multiple intermittent redshifts to gaseous haloes of intervening galaxies (e.g. Bergeron 1986; Bergeron & Boiss\u00e9 1991; Lanzetta et al. 1995; Tripp, Savage & Jenkins 2000; Chen, Lanzetta & Webb 2001). In the past decade, owing to the rise of large spectroscopic surveys of galaxies with well-determined physical properties (e.g. SDSS), all sky UV surveys (e.g. GALEX), and improved sensitivity of UV spectrographs (e.g. COS), studies of the gaseous haloes of galaxies could systematically connect gas absorption properties to galaxy properties in statistically meaningful samples (e.g. Cooksey et al. 2010; Prochaska et al. 2011; Tumlinson et al. 2013; Liang & Chen 2014; Lehner, Howk & Wakker 2015). The aforementioned gaseous haloes are commonly referred to as the circum-galactic medium (CGM) and are ubiquitous in galaxies regardless of mass or star formation activity: even sub-L* galaxies (Bordoloi et al. 2014), and passive galaxies host a CGM (Thom et al. 2012). The current model of the CGM suggests the presence of a clumpy multiphase medium which extends beyond the virial radius of the host galaxy, with a declining radial density profile, containing a substantial amount of gas and metals (e.g. Werk et al. 2013, 2014; Liang & Chen 2014; Lehner et al. 2014, 2015; Prochaska et al. 2017). Observational studies targeting the CGM of L* galaxies showed that the CGM gas content is comparable to the mass of the interstellar medium (ISM; e.g. Chen et al. 2010; Tumlinson et al. 2011; Werk et al. 2014; Prochaska et al. 2017) and correlates positively with ISM properties (Borthakur et al. 2015). Additionally, CGM observations infer a significant amount of metals (e.g. Werk et al. 2013; Peeples et al. 2014) where CGM metallicities can extend to supersolar metallicities (Prochaska et al. 2017). The clumpy multiphase CGM consists of a warm gas T \u223c 104 \u2212 5\u2009K (clumpy in nature) embedded within a hot diffuse T \u223c 106\u2009K medium (e.g. Heitsch & Putman 2009; Armillotta et al. 2017; Bordoloi et al. 2017). The multiphase structure of the CGM is corroborated by the variety of observed ionic species which survive at a vast range of temperatures: While the warm gas hosts the low ionization species (e.g. H\u2009i, Si\u2009ii, Si\u2009iii, C\u2009ii, C\u2009iv), the hot medium is home for the most highly ionized species (e.g. O\u2009vi, O\u2009vii). Additionally, the spectral line profiles of absorbers in the CGM can be reproduced by invoking a patchy medium (e.g. Stern et al. 2016; Werk et al. 2016), i.e. multiple high density gas clouds contribute to the optical depth along the line of sight thus leaving their kinematic imprint on the absorption line profile. For a review of the CGM, see Putman, Peek & Joung (2012) and Tumlinson, Peeples & Werk (2017).","Citation Text":["Tumlinson et al. 2013"],"Functions Text":["In the past decade, owing to the rise of large spectroscopic surveys of galaxies with well-determined physical properties (e.g. SDSS), all sky UV surveys (e.g. GALEX), and improved sensitivity of UV spectrographs (e.g. COS), studies of the gaseous haloes of galaxies could systematically connect gas absorption properties to galaxy properties in statistically meaningful samples (e.g."],"Functions Label":["Background"],"Citation Start End":[[814,835]],"Functions Start End":[[385,769]]}
{"Identifier":"2018ApJ...855...26A__Hsu_et_al._2016_Instance_1","Paragraph":"Ever since the initial measurements of the cosmic infrared background (CIB; for a review, see Hauser & Dwek 2001) revealed that the amount of energy radiated in the far-infrared (IR) and submillimeter spectral windows is comparable to that measured at ultraviolet (UV) and optical wavelengths, it has been widely recognized that one of the keys to a comprehensive understanding of the star formation history of the universe is the study of the multiwavelength properties of dusty star-forming galaxies (DSFGs), whose integrated radiation produces the CIB. These systems host intense star-forming activity obscured by large columns of dust, which re-emit the UV radiation of young hot stars at longer wavelengths, so that the peak of their rest-frame spectral energy distribution (SED) falls in the far-IR. In the local universe, DSFGs are typically identified as luminous or ultraluminous infrared galaxies (LIRGs\/ULIRGs; Sanders & Mirabel 1996), whereas more distant DSFGs\u2019 emission can be redshifted into the submillimeter domain, allowing many to manifest as submillimeter galaxies (SMGs; Blain et al. 2002; Casey et al. 2014). SMGs were first detected with the Submillimeter Common-User Bolometer Array (SCUBA; Holland et al. 1999) on the James Clerk Maxwell Telescope both in blank-field surveys (e.g., Hughes et al. 1998; Barger et al. 1999; Scott et al. 2002; Serjeant et al. 2003; Webb et al. 2003; Coppin et al. 2006) and behind galaxy clusters (e.g., Smail et al. 1997; Chapman et al. 2002; Cowie et al. 2002; Knudsen et al. 2008). With the subsequent advent of comparable single-dish telescopes and larger format instruments covering the 870 \u03bcm atmospheric window like the Large APEX Bolometer Camera (LABOCA; Siringo et al. 2009) on the 12 m Atacama Pathfinder Experiment telescope (APEX; G\u00fcsten et al. 2006), and more recently SCUBA-2 (Holland et al. 2013), the number of known SMGs is now of the order of a few thousand (e.g., Wei\u00df et al. 2009; Johansson et al. 2011; Chen et al. 2013; Hsu et al. 2016; Geach et al. 2017), and intensive observational efforts have been devoted to understanding their physical properties. In this quest, one of the main challenges has been the coarse (15\u2033\u201320\u2033) resolution of single-dish observations, which hinders the identification of counterparts at different wavelengths and can result in the blending of multiple, fainter SMGs into a single brighter object. However, the persistence of the local radio\u2013FIR correlation to higher redshifts (e.g., Condon 1992) allows accurate SMG positions from deep radio imaging obtained with the Very Large Array (VLA) at 1.4 GHz and the Australia Telescope Compact Array (ATCA) at 2.1 GHz (e.g., Ivison et al. 1998, 2000, 2002; Smail et al. 2000; Chapman et al. 2002) to be determined, thus enabling the identification of optical and near-IR counterparts, determination of photometric and spectroscopic redshifts (e.g., Chapman et al. 2005), modeling of SEDs, analysis of individual morphologies, and characterization of dust and stellar components (for full reviews of these results, see Blain et al. 2002 and Casey et al. 2014). The general picture derived from such studies is that SMGs are massive, gas-rich galaxies with high IR luminosities (\n\n\n\n\n\n) and complex optical\/near-IR morphologies, in which respects they resemble the local ULIRG population. However, SMGs have a median redshift z \u223c 2.5 (Chapman et al. 2005) and a significantly higher number density than ULIRGs. Complementary observations with centimeter and (sub)millimeter telescopes have been used as well to study the cool, molecular gas of SMGs (e.g., Carilli & Walter 2013), an effort that has been transformed in recent years thanks to the exceptional spatial resolution and sensitivity provided by the Jansky VLA and the Atacama Large Millimeter\/submillimeter Array (ALMA). Moreover, high-resolution continuum imaging at 870 \u03bcm with ALMA has made it possible to resolve the structure of the dust emission from SMGs and identify their counterparts in an unbiased way, revealing that a large fraction of bright single-dish detections actually \u201cbreak up\u201d into multiple, fainter (\n\n\n\n\n\n mJy) SMGs blended together at the coarse resolution of the maps in which they were detected (Hodge et al. 2013; Karim et al. 2013; Simpson et al. 2015a, 2015b).","Citation Text":["Hsu et al. 2016"],"Functions Text":["the number of known SMGs is now of the order of a few thousand (e.g.,","and intensive observational efforts have been devoted to understanding their physical properties."],"Functions Label":["Background","Background"],"Citation Start End":[[2000,2015]],"Functions Start End":[[1871,1940],[2037,2134]]}
{"Identifier":"2022MNRAS.516.3381J__Lindblom_&_Owen_2002_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":["Lindblom & Owen 2002"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[1136,1156]],"Functions Start End":[[924,1091]]}
{"Identifier":"2020AandA...633A.163C__Aalto_et_al._2015_Instance_2","Paragraph":"By using the RADEX2 dense cloud models developed by Aalto et al. (2015) to reproduce the HCN(3\u20132)\/(1\u20130) line luminosity ratios in the outflow of Mrk 231, we can attempt to find a combination of XHCN, XCN, Tkin, and nH2 solutions that can also fit the CN\/HCN and CN spin doublet line ratios (Table 2). We assume that the HCN and CN line emissions arise from the same dense cloud population, while the low-J CO line emission is due to a different, more diffuse phase of the outflow. We recall that in these models (see also Aalto et al. 2015), the dense clouds can be either self-gravitating virialised clouds, which implies that their internal velocity dispersion (\u0394vsg) is locked to their mass (Mvir) and size (R) through \u0394vsg\u2004=\u2004(GMvir\/G)1\/2, or unbound clouds, for which \u0394v\u2004\u226b\u2004\u0394vsg. We explored CN and HCN abundances in the range between 10\u22128 and 10\u22126. We find that depending on whether the clouds are self-gravitating or unbound, the models produce very different values for the absolute CN and HCN abundances, hence XCN and XHCN remain quantitatively unconstrained for the outflow with current data. However, all possible solutions that fit the observed line ratios consistently require XCN\u2004> \u2004XHCN, with a CN abundance that is at least a factor of three higher than the HCN abundance. Gas densities for this outflow phase (traced by the CN and HCN emissions) are nH2\u2004\u223c\u2004105\u2005\u2212\u2005106 cm\u22123, with temperatures not much higher than Tkin\u2004\u223c\u200420 K. Because CN is a well-known PDR tracer (see also Sect. 1), these results strongly suggest that the whole dense cloud population in outflow is affected by UV radiation. We should mention that high CN abundances may also be due to cosmic rays (e.g. see work done on the Galactic centre by Harada et al. 2015), which are known to permeate the outflow of Mrk 231, as inferred by Gonz\u00e1lez-Alfonso et al. (2018) based on the high OH+ abundance. However, it is not clear whether a cosmic-ray chemistry would also explain XCN\u2004> \u2004XHCN.","Citation Text":["Aalto et al. 2015"],"Functions Text":["We recall that in these models (see also","the dense clouds can be either self-gravitating virialised clouds, which implies that their internal velocity dispersion (\u0394vsg) is locked to their mass (Mvir) and size (R) through \u0394vsg\u2004=\u2004(GMvir\/G)1\/2, or unbound clouds, for which \u0394v\u2004\u226b\u2004\u0394vsg."],"Functions Label":["Uses","Uses"],"Citation Start End":[[522,539]],"Functions Start End":[[481,521],[542,782]]}
{"Identifier":"2019MNRAS.486..360K__N\u00e4ttil\u00e4_&_Pihajoki_2018_Instance_1","Paragraph":"We have presented the principles and framework for calculating the radio signal from a PSR in an EMRB. We restrict our study to the extreme mass ratio of EMRB systems and so do not consider PSRs in stellar-mass BH binaries with finite mass ratios (e.g. Blanchet 2014; Liu et al. 2014). We account for both relativistic and astrophysical effects and the convolution between the two. This includes gravitational and relativistic time dilation and energy shift, gravitational light bending, complex orbital dynamics induced by spin couplings, temporal variation and distortion of the pulse profile due to spin axis precession and relativistic aberration, second-order pulses due to gravitational bending, and dispersions (temporal and spatial) induced by the material along the line of sight. We have demonstrated that within our framework we are able to determine the time\u2013frequency behaviour accounting for all these effects. The framework also applies for any orbital configuration, e.g. we are not restricted to orbital motion in the equatorial plane or beaming confined to the orbital plane. The methods used are entirely covariant and general relativistic, rather than working under any post-Newtonian approximation and so are inherently more accurate. Indeed, the post-Newtonian method is an explicitly weak-field method, and the validity of its application to strong-field dynamical regimes is unclear (Will 2011). Whilst working explicitly in the Kerr metric means that we are unable to independently probe either alternative gravitational theories or extensions to Kerr (e.g. Kerr space\u2013time with an arbitrary mass quadrupole; see Bini et al. 2009), our framework provides the basis for a theoretical timing model that can then be compared with observations for tests of strong field GR. We approximate the PSR body as a perfect sphere. However, due to the spin of the PSR the true shape is more oblate. This will ultimately influence the pitch angle of the ray with the neutron star surface. This effect is considered to be minor, but the method could easily be extended to account for this oblateness (see N\u00e4ttil\u00e4 & Pihajoki 2018). We neglect the effects of hydrodynamic drag due to the plasma that surrounds that BH since at compact radii (\u2272 104rg) the gravitational and relativistic effects dominate (Psaltis 2012). We also do not take account of any potential Newtonian perturbations on the motion of the pulsar (e.g. Merritt et al. 2011) due to the presence of other masses (e.g. stars, other compact objects etc.) since these factors are likely negligible for the orbital periods considered in this work (\u22720.3\u2009yr; Liu et al. 2012). Indeed, the potential for external perturbations to hamper tests of strong-field GR necessitates that an ideal PSR\u2013EMRB systems should have orbital periods on the order of 0.1\u2009yr (or better), or else observations should be taken close to periapsis (see discussion in Psaltis, Wex & Kramer 2016). These are precisely the regions where the space\u2013time curvature and orbital acceleration is greatest, further stressing the importance of a strong-field timing model. We also neglect any influence of gravitational radiation on the orbit or the ray trajectory. The neglect of gravitational radiation is justified since in the extreme mass ratio limit, the time-scale for orbital decay due to gravitational wave emission is (Misner, Thorne & Wheeler 1973)\n(45)\r\n\\begin{eqnarray*}\r\n\\tau _{\\mathrm{ GW}} \\sim \\frac{5 r^4}{96mM(m+M)} f(e)^{-1} \\ ,\r\n\\end{eqnarray*}\r\nwhere M is the mass of the BH, m the pulsar mass, and r the orbital separation. The eccentricity function is\n(46)\r\n\\begin{eqnarray*}\r\nf(e) = (1-e^2)^{-7\/2} \\left(1 + \\frac{73}{24} e^2 + \\frac{37}{96} e^4\\right) \\ .\r\n\\end{eqnarray*}\r\nIf we take the PSR orbital period P to be Keplerian, then for a pulsar with mass $1.4 \\, \\mathrm{M}_{\\odot }$ on an eccentric (e = 0.8), P = 0.1\u2009yr orbit around a BH with mass $4.3 \\times 10^6 \\, \\mathrm{M}_{\\odot }$(47)\r\n\\begin{eqnarray*}\r\n\\frac{\\tau _{\\mathrm{ GW}}}{P} \\sim 10^9 \\gt \\gt 1 \\ ,\r\n\\end{eqnarray*}\r\nand so the effects of gravitational radiation can be neglected. Even for smaller radii and more eccentric orbits the space\u2013time is well approximated as stationary (e.g. \u03c4GW\/P\u2009 \u223c 105 for e = 0.9, r = 100\u2009M). Whilst the effects of gravitational radiation are then not important for a single orbit, for observations over longer periods of time the effect of gravitational emission on the orbit and hence the timing solution will need to be considered. The PSR may also emit a gravitational wave burst during passage through periastron (Berry & Gair 2013a,b). The influence of this gravitational radiation on both the PSR trajectory and the photon ToA is highly non-trivial and not considered here.","Citation Text":["N\u00e4ttil\u00e4 & Pihajoki 2018"],"Functions Text":["We approximate the PSR body as a perfect sphere. However, due to the spin of the PSR the true shape is more oblate. This will ultimately influence the pitch angle of the ray with the neutron star surface. This effect is considered to be minor, but the method could easily be extended to account for this oblateness (see"],"Functions Label":["Future Work"],"Citation Start End":[[2115,2138]],"Functions Start End":[[1795,2114]]}
{"Identifier":"2018ApJ...866L...1S__Pecharrom\u00e1n_et_al._1999_Instance_6","Paragraph":"It was found that the complex dielectric function from Pecharrom\u00e1n et al. (1999) for the sample obtained by heating bayerite at 1273 K, assuming a spheroid with depolarization parameters of (0.35, 0.003), produced an opacity with 11, 20, 28, and 32 \u03bcm features, so this component was included in the models. However, with only this component, the observed 20 \u03bcm features in the residual spectra were found to be wider than those in the models. By adding the opacity of the sample obtained by heating boehmite at 1173 K, the width of the 20 \u03bcm feature could be matched. This was done using the complex dielectric function for the sample obtained by heating boehmite at 1173 K from Pecharrom\u00e1n et al. (1999), assuming a spheroid with depolarization parameters of (0.35, 0.035). The complex dielectric functions of the samples obtained by heating bayerite and boehmite to various temperatures (Pecharrom\u00e1n et al. 1999) were derived by modeling the reflectance spectra of pellets obtained by pressing powders of these materials under great pressure. This method required Pecharrom\u00e1n et al. (1999) to assume an effective medium theory, such that a pellet is a mixture of one of their samples with a matrix of air. Pecharrom\u00e1n et al. (1999) noted that heating bayerite at 500\u00b0C eliminates the XRD pattern of bayerite, and they note that at 700\u00b0C, the infrared reflectance spectrum of the boehmite sample no longer shows OH\u2212 stretching bands. This must mean that the samples obtained from heating bayerite at 1273 K and from heating boehmite at 1173 K are no longer bayerite or boehmite, respectively. XRD performed by Pecharrom\u00e1n et al. (1999) of the sample of bayerite prepared at 1273 K suggests only \u03b8-alumina was present, and their infrared and NMR spectroscopy confirms this. XRD of their sample obtained from heating boehmite to 1173 K (Pecharrom\u00e1n et al. 1999) suggests \u03b4-alumina to be present, though some amounts of \u03b8-alumina and \u03b1-alumina are present, as they deduce from XRD and infrared and NMR spectroscopy.","Citation Text":["Pecharrom\u00e1n et al. (1999)"],"Functions Text":["XRD performed by","of the sample of bayerite prepared at 1273 K suggests only \u03b8-alumina was present, and their infrared and NMR spectroscopy confirms this."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1612,1637]],"Functions Start End":[[1595,1611],[1638,1774]]}
{"Identifier":"2022AandA...667A..97A__Xiong_(1986)_Instance_1","Paragraph":"The temperature gradient of the 3-equation model is comparable to results of different TCM approaches (Xiong & Deng 2001; Li & Yang 2007) for the base of the solar convective envelope. Both Zhang & Li (2012; their Figs. 6 and 7) and Xiong & Deng (2001; their Fig. 8) find a temperature gradient that transitions gradually from the adiabatic to the radiative value. They also find a Deardorff layer with a degree of sub-adiabaticity at the formal Schwarzschild boundary comparable to our findings. From the convective flux as presented in Xiong (1986) one also would expect a similar temperature gradient in the overshooting zone. Furthermore, the shape of the model temperature gradient is also in qualitative agreement with the discussion in Viallet et al. (2015). They argue that under the physical conditions in convective cores, in regions of overshooting efficient chemical mixing and a gradually transitioning temperature gradient are expected. In the 3-equation non-local model, the extent of the nearly adiabatic overshooting zone is controlled by the shape of the negative convective flux in the overshooting zone. For smaller (more negative) values of the convective flux (i.e. more efficient buoyancy braking) the temperature gradient is expected to be closer to the adiabatic value, while for larger (less negative) values it will be closer to the radiative temperature gradient. In Eq. (1) the negative convective flux and the dissipation term act as sink terms in the overshooting zone. Hence, the behaviour of the dissipation term will impact also on the convective flux and in turn on the value of the temperature gradient in the overshooting zone. In computations with the 1-equation non-local version of the theory, the negative convective flux is the dominant sink term for the TKE and the actual dissipation term is negligible (Fig. 8 of Paper I). This leads to more negative values of the convective flux and thus to a mostly adiabatic temperature gradient in the overshooting zone. We discuss this in more detail in Sect. 5.1 (we also refer to Fig. 10).","Citation Text":["Xiong (1986)"],"Functions Text":["From the convective flux as presented in","one also would expect a similar temperature gradient in the overshooting zone."],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[538,550]],"Functions Start End":[[497,537],[551,629]]}
{"Identifier":"2022MNRAS.514.2010M__Feng_&_Holder_2018_Instance_2","Paragraph":"In the last few years, several experiments have reported upper limits on the power spectrum of 21-cm fluctuations during reionization (Parsons et al. 2014; Patil et al. 2017; Barry et al. 2019; Mertens et al. 2020; The HERA Collaboration 2021b) and the earlier cosmic-dawn era (Eastwood et al. 2019; Gehlot et al. 2019, 2020; Garsden et al. 2021; Yoshiura et al. 2021). Scenarios in which the bulk IGM is still colder than the cosmic microwave background (CMB) during reionization give rise to the strongest fluctuations and so will be the first models to be tested as upper limits continue to improve (e.g. Parsons et al. 2014; Pober et al. 2015; Greig, Mesinger & Pober 2016). Similarly, stronger-than-expected 21-cm signals can arise if the cosmic radio background has contributions other than the CMB (Feng & Holder 2018), e.g. synchrotron emission from accreting black holes (Ewall-Wice et al. 2018), star-forming galaxies (Mirocha & Furlanetto 2019), or from decaying particles (Fraser et al. 2018; Pospelov et al. 2018). Indeed, constraints from MWA, HERA, and LoFAR disfavour models with negligible X-ray heating at z \u223c 8\u20139 or very strong radio backgrounds (Ghara et al. 2020, 2021; Mondal et al. 2020; Greig et al. 2021a, b; The HERA Collaboration 2021a). Of course, the recent report of an absorption signal in the sky-averaged spectrum at z \u223c 17 from EDGES (Bowman et al. 2018) requires an even colder IGM (Barkana 2018; Boddy et al. 2018; Fialkov, Barkana & Cohen 2018; Kovetz et al. 2018; Mu\u00f1oz & Loeb 2018) or a brighter background (Ewall-Wice et al. 2018; Feng & Holder 2018; Fialkov & Barkana 2019; Mirocha & Furlanetto 2019) than models in \u039bCDM cosmologies generally predict. However, the most stringent power spectrum upper limits from The HERA Collaboration (2021b) are derived at sufficiently low redshifts relative to EDGES (z \u2272 10 versus z \u2243 18) that they cannot yet directly address the EDGES controversy (Hills et al. 2018; Bradley et al. 2019; Singh & Subrahmanyan 2019; Sims & Pober 2020; Tauscher, Rapetti & Burns 2020; Singh et al. 2021).","Citation Text":["Feng & Holder 2018"],"Functions Text":["Of course, the recent report of an absorption signal in the sky-averaged spectrum at z \u223c 17 from EDGES (Bowman et al. 2018)","or a brighter background","than models in \u039bCDM cosmologies generally predict."],"Functions Label":["Motivation","Motivation","Motivation"],"Citation Start End":[[1571,1589]],"Functions Start End":[[1265,1388],[1521,1545],[1642,1692]]}
{"Identifier":"2022AandA...659A...5Y__Mills_et_al._2018_Instance_1","Paragraph":"Since its discovery more than five decades ago (Cheung et al. 1968), ammonia (NH3) has been a most valuable molecule for investigating the physical properties of molecular clouds (e.g., Ho & Townes 1983). While thermally excited transitions in the centimeter-wavelength inversion transitions of ammonia are regarded as a reliable thermometer of molecular clouds (e.g., Walmsley & Ungerechts 1983; Danby et al. 1988), ammonia masers have attracted attention since the first detection of maser action in the (J, K) = (3,3) metastable (J = K) line toward the massive star-forming region W33 (Wilson et al. 1982). Subsequent observations have led to the detection of new metastable ammonia masers, including 15NH3 (3,3) (Mauersberger et al. 1986), NH3 (1,1) (Gaume et al. 1996), NH3 (2,2) (Mills et al. 2018), NH3 (5,5) (Cesaroni et al. 1992), NH3 (6,6) (Beuther et al. 2007), NH3 (7,7), NH3 (9,9), and NH3 (12,12) (Henkel et al. 2013). These have led to the discovery of metastable maser lines in 22 different regions (Mauersberger et al. 1986, 1987; Wilson & Henkel 1988; Wilson et al. 1990; Pratap et al. 1991; Cesaroni et al. 1992; Wilson & Schilke 1993; Mangum & Wootten 1994; Kraemer & Jackson 1995; Zhang & Ho 1995; Zhang et al. 1999; Walsh et al. 2007; Hunter et al. 2008; Galv\u00e1n-Madrid et al. 2009; Brogan et al. 2011; Urquhart et al. 2011; Walsh et al. 2011; Wang et al. 2012; Henkel et al. 2013; Hoffman & Joyce 2014; McEwen et al. 2016; Mills et al. 2018; Hogge et al. 2019; Mei et al. 2020; Towner et al. 2021). Compared with the metastable ammonia masers, detected non-metastable (J > K) ammonia maser transitions are more numerous. The first highly excited non-metastable ammonia maser was detected by Madden et al. (1986) in the (J, K) = (9,6) and (6,3) lines. Thereafter, many other NH3 non-metastable inversion transition lines have been identified as masers, including the (5,3), (5,4), (6,1), (6,2), (6,4), (6,5), (7,3), (7,4), (7,5) (7,6), (8,3), (8,4), (8,5), (8,6), (9,3), (9,4), (9,5), (9,7), (9,8), (10,7), (10,8), (10,9), and (11,9) transitions (e.g., Mauersberger et al. 1987, 1988; Walsh et al. 2007; Henkel et al. 2013; Mei et al. 2020). Except for the NH3 (3,3) masersproposed to be associated with four supernova remnants (McEwen et al. 2016), almost all the other ammonia masers are detected in high-mass star-forming regions (HMSFRs). However, while many HMSFRs host water (H2O), hydroxyl (OH), or methanol (CH3OH) masers, ammonia masers are quite rare in these sources, and the role that the environment of a young high-mass star plays in their excitation remains unclear. Therefore, dedicated searches for ammonia masers in HMSFRs are indispensable in regard to their overall incidence and association with different environments, which can provide additional constraints on the pumping mechanism of ammonia masers.","Citation Text":["Mills et al. 2018"],"Functions Text":["Subsequent observations have led to the detection of new metastable ammonia masers, including","NH3 (2,2)"],"Functions Label":["Background","Background"],"Citation Start End":[[786,803]],"Functions Start End":[[610,703],[775,784]]}
{"Identifier":"2019AandA...622A.106M__Herranz_et_al._2009_Instance_1","Paragraph":"The standard single-frequency detection methods for point sources in the CMB and far IR are based on wavelet techniques (Vielva et al. 2003; Barnard et al. 2004; Gonz\u00e1lez-Nuevo et al. 2006) or on the matched filter (or MF hereafter, Tegmark & de Oliveira-Costa 1998; Herranz et al. 2002; Barreiro et al. 2003; L\u00f3pez-Caniego et al. 2006, see also Herranz & Vielva 2010 for a review.). Wavelets are well suited for the detection of compact sources due to their good position-scale determination properties, whereas the MF is the optimal linear detector-estimator because it provides the maximum signal-to-noise (S\/N) amplification for a source with a known shape (usually the point-spread function, or PSF hereafter, of the telescope) embedded in statistically homogeneous and spatially correlated noise. By default, these techniques are applicable only to single-frequency sky images: even for multiwavelength observatories such as the Herschel Space Observatory (Pilbratt et al. 2010) or Planck (Tauber et al. 2010), the standard detection pipelines have produced individual source catalogs for each frequency band (see e.g., Planck Collaboration VII 2011; Planck Collaboration XXVIII 2014; Planck Collaboration XXVI 2016; Maddox et al. 2018). The next logical step is to boost the signal of faint sources by combining the different bands into a single detection, that is, \u201cmultifrequency detection\u201d. Most of the blind component separation algorithms that are used for diffuse components in microwave and far IR astronomy can not deal with the high diversity of spectral behaviors associated to the different populations of extragalactic compact sources (see for example Leach et al. 2008). However, over the last few years a number of multifrequency compact source detection techniques have been proposed in the literature (Herranz & Sanz 2008; Herranz et al. 2009; Lanz et al. 2010, 2013; Planck Collaboration Int. LIV 2018). A review on the topic can be found in Herranz et al. (2012). In particular, if the spatial profile and the spectral energy distribution (SED) of the sources are known, and if the cross-power spectrum is known, or can be estimated from the data, the optimal linear detection method is the matched multifilter (or MMF hereafter, Herranz et al. 2002). Lanz et al. (2010) also showed that the MMF can be generalized for the case where the SED of the sources is not known. This generalization outperforms the single-frequency MF in terms of S\/N and can be used to infer the spectral index of synchrotron-dominated radio sources, as shown in Lanz et al. (2013). However, in this paper we will incorporate a specific SED to the MMF in order to derive a photometric redshift estimation of dusty galaxies and high-redshift star forming galaxies detected in the IR part of the spectrum1. We will do so by applying the multifrequency MMF filter to the first and second data releases of the Herschel Astrophysical Terahertz Large Area Survey (the Herschel-ATLAS or H-ATLAS, Eales et al. 2010), the largest single key project carried out in open time with the Herschel Space Observatory. We restrict our multifrequency analysis to the three wavelength bands covered by the SPIRE instrument aboard Herschel (Griffin et al. 2010), centered around 250, 350 and 500 \u03bcm. As discussed in Hopwood et al. (2010), Lapi et al. (2011), Gonz\u00e1lez-Nuevo et al. (2012), Pearson et al. (2013) and Donevski et al. (2018), the SPIRE bands are ideal for capturing the peak in the SED corresponding to dust emission of star-forming galaxies at z\u2004\u223c\u20042, that is redshifted from its rest-frame wavelength around 70\u2013100 \u03bcm to the SPIRE wavelengths: This is the redshift range where galaxies have formed most of their stars. At higher redshifts, dusty star-forming galaxies (DSFGs) occupy the most massive halos and are among the most luminous objects found at z\u2004\u2273\u20044 (Micha\u0142owski et al. 2014; Oteo et al. 2016; Ikarashi et al. 2017). These high-redshift DSFGs have markedly red colors as seen by SPIRE, with rising flux densities from 250 to 500 \u03bcm (the so-called \u201c500 \u03bcm-risers\u201d), and have received a great deal of attention in the recent years (see for example Ivison et al. 2016; Negrello et al. 2017; Strandet et al. 2017). The DSFGs, and particularly the 500 \u03bcm risers uncovered by Herschel, are providing much insight into the early star forming history of the universe. However, sensitivity and limited angular resolution severely constrain the power of this type of objects as astrophysical probes. The sensitivity of SPIRE allows for the direct detection of only the brightest, and thus rarest objects, at the bright end of the luminosity function. By means of our multifrequency MMF technique, we intend to enhance the detectability and statistical significance of very faint red objects in the H-ATLAS source catalog and so expand the list of reliable 500 \u03bcm-riser candidates.","Citation Text":["Herranz et al. 2009"],"Functions Text":["However, over the last few years a number of multifrequency compact source detection techniques have been proposed in the literature"],"Functions Label":["Motivation"],"Citation Start End":[[1846,1865]],"Functions Start End":[[1691,1823]]}
{"Identifier":"2019AandA...629A..93Y__Hopkins_2018_Instance_1","Paragraph":"The evolution of galaxies sensitively depends on the stellar initial mass function (IMF). The understanding of the IMF has changed rapidly in the past decade. Despite a direct conflict (Kroupa et al. 2013) with the theoretical expectation that the IMF should vary as the star-forming environment alters (e.g. ambient gas temperature, metallicity, density, and pressure dependence, as argued by Adams & Fatuzzo 1996; Larson 1998; Elmegreen 2004; Dib et al. 2007; Papadopoulos 2010) the observed IMFs of star clusters in the local Universe are consistent with no variation (Kroupa 2001, 2002; Bastian et al. 2010; Offner et al. 2014; Hopkins 2018). Thus, the universal and invariant canonical IMF assumption was widely applied. The more recent observations, which are able to probe physical regimes further from the solar and Galactic neighbourhood, have been consistently suggesting a variation of the galaxy-wide IMF (gwIMF1). For the distribution of low-mass stars, a bottom-heavy IMF (excess in the number of low-mass stars) in the inner regions of massive metal-rich elliptical galaxies is indicated by the galaxy mass-to-light ratio (Li et al. 2017), spectral analysis of stellar-mass sensitive features (Vazdekis et al. 1997, 2003; Cenarro et al. 2003; van Dokkum & Conroy 2010; Conroy & van Dokkum 2012; Ferreras et al. 2013; Mart\u00edn-Navarro et al. 2015; La Barbera et al. 2017; Parikh et al. 2018) and lensing studies (Auger et al. 2010; Oldham & Auger 2018)2. For the distribution of massive stars, independent evidence strongly indicates a systematically varying IMF (galaxy photometry: Hoversten & Glazebrook 2008; Meurer et al. 2009; Lee et al. 2009; Gunawardhana et al. 2011, metal abundance of galaxy clusters: Renzini & Andreon 2014; Urban et al. 2017, isotope abundance: Romano et al. 2017; Zhang et al. 2018), being top-heavy (more massive stars than predicted when assuming the canonical IMF) when the SFR is high and\/or when the metallicity is low, as summarized by Kroupa et al. (2013), Yan et al. (2017), and Je\u0159\u00e1bkov\u00e1 et al. (2018). We note that the IMF can be bottom-heavy and top-heavy at the same time (see Je\u0159\u00e1bkov\u00e1 et al. 2018).","Citation Text":["Hopkins 2018"],"Functions Text":["Despite a direct conflict","with the theoretical expectation that the IMF should vary as the star-forming environment alters","the observed IMFs of star clusters in the local Universe are consistent with no variation"],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[632,644]],"Functions Start End":[[159,184],[206,302],[481,570]]}
{"Identifier":"2017MNRAS.470.1442C__Hurley_et_al._2002_Instance_1","Paragraph":"We then allow the synthetic single or binary system to evolve until present time, adopting for our reference model a thin disc age of 10 Gyr (Cojocaru et al. 2014) and a thick disc age of 12 Gyr. This is motivated by the findings of Feltzing & Bensby (2009), who presented a sample of very likely thick disc candidates with ages, on average, well above 10 Gyr and of Ak et al. (2013), who found that thick disc cataclysmic variables have ages up to 13 Gyr. If the synthetic star is single and has time to become a white dwarf, it evolves following the cooling tracks detailed in the following section. If that is the case, the mass of the white dwarf is obtained from the initial-to-final mass relation (IFMR) according to the prescription from Hurley, Tout & Pols (2002). If the object is member of a binary system and the primary star has time to become a white dwarf, then the pair can evolve through two different scenarios. In the first scenario, the binary evolves without mass transfer interactions as a detached system and the primary star evolves into a white dwarf that subsequently cools down following the cooling sequences described in the next section. In this case, the mass of the white dwarf is also calculated from the IMFR of Hurley et al. (2002). The second scenario involves mass transfer episodes and the evolution of the binary is obtained following the prescriptions of the bse package (Hurley et al. 2002), following the parameter assumptions detailed in Camacho et al. (2014). If the system evolves though the common envelope phase, we use the \u03b1-formalism as described in Tout et al. (1997), with \u03b1CE being the efficiency in converting orbital energy into kinetic energy to eject the envelope (assumed to be 0.3 in our reference model). This implementation also takes into account the \u03b1int parameter (assumed to be 0.0 in our reference model), first presented in Han, Podsiadlowski & Eggleton (1995), describing the fraction of the internal energy (thermal, radiation and recombination energy) used to eject the envelope. As described in Camacho et al. (2014), the \u03b1int parameter is used to include the effects of the internal energy in the binding energy parameter \u03bb, which is thus not taken as a constant, but computed using a specific algorithm (Claeys et al. 2014) in bse. In the current version of the code, provided that a positive value is used, the parameter \u03b1int represents the fraction of recombination energy that contributes to eject the envelope. It is important to note that the thermal energy of the envelope is always taken into account (using the virial theorem) even if \u03b1int is set to zero. For a more detailed discussion on how this is implemented in the latest version of BSE and important comments on the correct use of BSE and the notations used in the code itself, we direct the reader to Zorotovic, Schreiber & Parsons (2014a), mentioning that the notations \u03b1int or \u03b1rec are, in our case, equivalent.","Citation Text":["Hurley et al. (2002)"],"Functions Text":["In this case, the mass of the white dwarf is also calculated from the IMFR of"],"Functions Label":["Uses"],"Citation Start End":[[1245,1265]],"Functions Start End":[[1167,1244]]}
{"Identifier":"2021ApJ...920...89C__Arag\u00f3n-Calvo_et_al._2007_Instance_1","Paragraph":"As mentioned in Section 1, once the sample of matching halos between the twin N-body simulations is established, we can study the halo\u2212LSS correlations. We employ the Hessian matrix method used in many previous articles (e.g., Hahn et al. 2007a, 2007b; Zhang et al. 2009; Kang & Wang 2015) to define the matrix as:\n4\n\n\n\n\n\n\nH\n\n\nij\n\n\n=\n\n\n\n\n\n\u2202\n\n\n2\n\n\n\n\n\u03c1\n\n\ns\n\n\n\n\nx\n\n\n\n\n\u2202\n\n\nx\n\n\nj\n\n\n\u2202\n\n\nx\n\n\ni\n\n\n\n\n\n.\n\n\nThis method is based on the smoothed density field \n\n\n\n\n\n\n\u03c1\n\n\ns\n\n\n\n\nx\n\n\n\n\n at the halo position (based on a more accurate and improved algorithm by Wang et al. 2020), which can be given by the Cloud-in-Cell (CIC) technique (MacNeice 1995). The smoothing length Rs is the only parameter of the CIC, which can be regarded as the typical scale of a halo LSS environment identified by the Hessian matrix method. Many previous works (e.g., Arag\u00f3n-Calvo et al. 2007; Hahn et al. 2007a; Zhang et al. 2009; Codis et al. 2012; Hoffman et al. 2012; Libeskind et al. 2013; Trowland et al. 2013, but see also Libeskind et al. 2014) used a constant smoothing length. To determine which Rs value should be chosen, we test some LSS properties (e.g., the environment and eigenvectors) of matching halos for three fixed Rs: 2.5,  5, and  10 h\u22121 Mpc. We find that the halo\u2212LSS correlation, as well as other main conclusions we make below, become stronger as Rs decreases, which is a reasonable result according to our previous discussions. Consequently Rs = 2.5 h\u22121 Mpc is chosen hereafter. The three eigenvalues of the Hessian matrix are marked as \u03bb1, \u03bb2, and \u03bb3, with corresponding eigenvectors e1, e2, and e3. Eigenvalues \u03bbi can be used to define the LSS environment of dark matter halos according to the number of positive eigenvalues (Zel\u2019Dovich 1970; Hahn et al. 2007a, 2007b; Libeskind et al. 2018; Zhang & Yang 2019); i.e.,\n\n1.\nvoid: 0  \u03bb1 \u03bb2 \u03bb3\n\n\n2.\nwall: \u03bb1  0  \u03bb2 \u03bb3\n\n\n3.\nfilament: \u03bb1 \u03bb2  0  \u03bb3\n\n\n4.\ncluster: \u03bb1 \u03bb2 \u03bb3  0\n\nand the eigenvectors ei stand for the three compressed directions of the smoothed density field. The e3 vector indicates the least compressed direction, which is a robust and universal definition of the LSS. In this work, we will focus on the alignment of halo spin and shape with the e3 vector; i.e., \n\n\n\n\ncos\n\n\n\u03b8\n\n\n3\n\n\n=\na\n\u00b7\n\n\ne\n\n\n3\n\n\n\n\n, where a is the halo spin or major-axis vector.","Citation Text":["Arag\u00f3n-Calvo et al. 2007"],"Functions Text":["Many previous works","used a constant smoothing length."],"Functions Label":["Uses","Uses"],"Citation Start End":[[831,855]],"Functions Start End":[[804,823],[1016,1049]]}
{"Identifier":"2020AandA...635A..81P__simulations,_Georgobiani_et_al._(2003)_Instance_2","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)"],"Functions Text":["Using 3D simulations,","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."],"Functions Label":["Background","Background"],"Citation Start End":[[1688,1713]],"Functions Start End":[[1666,1687],[1714,2057]]}
{"Identifier":"2016MNRAS.458.1504S__Wang_et_al._2012_Instance_1","Paragraph":"The energy levels calculated by Fritzsche, Fischer & Fricke (1998) with grasp92 code are more closer to the NIST values as compared with other calculations within 3.6 per cent for the levels in the ground configurations. Our calculated MCDF energies agree well with NIST data within 6.8 per cent and the maximum error appears for ground state configuration (3s23p3) levels. The average percentage difference between our grasp1 energies and the measured values is 1.9 per cent. For the lowest 41 energy levels, we see that our calculated energies and previous fac energies (Landi & Bhatia 2010; Wang et al. 2012) belonging to the n = 3 complex in Ni\u2009xiv agree well with each other and both set of calculation are of the same accuracy. However, our calculated ground state configuration levels show better agreement with measured energies in comparison to those of fac calculations (7.2 per cent). When compared to MCDF calculations performed by Aggarrwal, Keenan & Msezane (2003) with grasp Dyall et al. (1989) code using a model of 14 configurations, our grasp1 levels in the excited configurations lie closer to the NIST ones. In the case of four energy levels (2D5\/2,3\/2,4P5\/2,3\/2), our calculated energy levels are more closer to NIST than those of Aggarrwal et al. (2003). The average percentage difference of these grasp energies with NIST values is 2.2 per cent, and the maximum error of 6.3 per cent appears for the first excited state (2D3\/2) of ground configuration. For more accurate target states we adopt a model of 13 electronic configurations which yields target-states energies as closer to NIST values as possible. For instance, our calculated energy levels 2D3\/2,5\/2 associated with the configuration 3s3p4 lie more closer to the observed values as compared to existing calculations. The remaining set of energy levels are almost of the same accuracy as those of previous data when compared to experimental energies. From this comparison it is evident that present grasp1 calculated N-electron target wavefunctions demonstrate good agreement with experimental values, which is necessary for more accurate PI cross-sections of Ni\u2009xiii.","Citation Text":["Wang et al. 2012"],"Functions Text":["For the lowest 41 energy levels, we see that our calculated energies and previous fac energies","belonging to the n = 3 complex in Ni\u2009xiv agree well with each other and both set of calculation are of the same accuracy.","However, our calculated ground state configuration levels show better agreement with measured energies in comparison to those of fac calculations (7.2 per cent)."],"Functions Label":["Similarities","Similarities","Differences"],"Citation Start End":[[594,610]],"Functions Start End":[[477,571],[612,733],[734,895]]}
{"Identifier":"2018AandA...618A..24V__Reeth_et_al._2015_Instance_1","Paragraph":"Gamma Doradus (\u03b3 Dor) stars, with 1.4 M\u2299 \u2272 M* \u2272 2.0 M\u2299, and slowly-pulsating B-type (SPB) stars, with 2.5 M\u2299 \u2272 M* \u2272 8 M\u2299, exhibit high-order gravity-mode (g-mode) pulsations, gravito-inertial pulsations (Van Reeth et al. 2016), and\/or purely inertial pulsations, such as r-modes (Saio et al. 2018). The restoring forces for gravity-modes and purely inertial pulsation modes are buoyancy and the Coriolis force, respectively. In the case of gravito-inertial pulsation modes, both forces contribute. As predicted by asymptotic theory, the pulsation periods for \u03b3 Dor and SPB stars were observed to form period spacing patterns (e.g. Chapellier et al. 2012; Chapellier & Mathias 2013; Kurtz et al. 2014; Bedding et al. 2015; Saio et al. 2015; Van Reeth et al. 2015; Ouazzani et al. 2017). The pulsation periods are equidistant in the asymptotic regime (with radial order n \u226b spherical degree l) for a non-rotating chemically homogeneous star (Tassoul 1980). Chemical gradients in the deep stellar interior cause pulsation mode trapping, which introduces non-uniform variations in the spacings (Miglio et al. 2008). On the other hand, the stellar rotation leads to shifts in the observed pulsation mode frequencies (Bouabid et al. 2013; Salmon et al. 2014; Van Reeth et al. 2015; Moravveji et al. 2016). For slowly rotating stars, the observed pulsation modes are split into frequency multiplets that depend on the mode identification. For moderate to fast rotators, that is, with rotation in the order of or more than 20% of the critical rotation rate, the observed period spacing patterns have a clear slope (e.g. Van Reeth et al. 2016; Ouazzani et al. 2017). Prograde (azimuthal order m >  0) and zonal modes have a downward slope, that is, the spacing between consecutive pulsation periods decreases with increasing pulsation period, i.e. radial order n. The period spacing patterns of retrograde modes (with m   0) mostly have an upward slope. For stars with detected period spacing patterns, this has been exploited to derive the near-core stellar rotation (e.g. Kurtz et al. 2014; Saio et al. 2015, 2018; Triana et al. 2015; Murphy et al. 2016; Schmid & Aerts 2016; Van Reeth et al. 2016; Ouazzani et al. 2017).","Citation Text":["Van Reeth et al. 2015","Van Reeth et al. 2015"],"Functions Text":["As predicted by asymptotic theory, the pulsation periods for \u03b3 Dor and SPB stars were observed to form period spacing patterns (e.g.","On the other hand, the stellar rotation leads to shifts in the observed pulsation mode frequencies"],"Functions Label":["Background","Motivation"],"Citation Start End":[[740,761],[1253,1274]],"Functions Start End":[[498,630],[1112,1210]]}
{"Identifier":"2017MNRAS.469.2720G__Maoz_et_al._2005_Instance_1","Paragraph":"With all this in mind, our last question is: What powers soft X-rays and [O\u2009III] in LINERs? This has no clear answer, but both are not tracing the same mechanism, since none of them match in morphologies. This is a clear difference between type-2 Seyferts and LINERs. In favour of the soft X-ray emission being originated by AGN photoionization, the RGS spectra studied by Gonzalez-Martin et al. (2010b) showed that in at least 30 per cent of their sample, a contribution of photoionization by the AGNs is required due to the presence of radiative recombination continua (RRC) from CV emission line. However, this does not guarantee a dominance of this emission mechanism. Moreover, cone-like morphologies at soft X-rays in some objects in this study point out again to the photoionization by the AGNs being responsible for the soft X-ray emission. Nevertheless, this assumes that we are seing LINERs with an LOS perpendicular to the accretion disc, which might not be the case. Indeed, the ultraviolet and X-ray variability detected for many of these LINERs (Maoz et al. 2005; Hern\u00e1ndez-Garc\u00eda et al. 2015) is in favour of a direct view of the AGNs (i.e. perpendicular to the disc under the UM). This is consistent with the fact that most of the [O\u2009III] morphologies found for LINERs are spheroids, if we assume that the [O\u2009III] traces the NLR. In addition, the fact that we detected a clear correspondence between soft X-ray and [O\u2009III] morphologies only in objects with log\u2009(LHX)>40, and also that all the objects where soft X-rays and [O\u2009III] match their morphologies seem to better follow the previously found relation between the size of the region and the hard X-ray luminosity (see Fig. 2 and Section 5.2), may argue in favour of the scenario in which the AGNs do not have enough thrust to ionize in the low-luminosity regime (Elitzur & Shlosman 2006; Elitzur & Ho 2009), ruling out photoionization by the AGNs at both soft X-ray and [O\u2009III] emissions. In this case, the most reasonable explanation for the [O\u2009III] is the host galaxy emission, which, anyhow, could also be on top of the AGNs, preventing its detection and erasing the connection (Gonz\u00e1lez-Mart\u00edn et al. 2014). The host galaxy can contribute either as star formation or shocks to the total [O\u2009III] emission. Regarding the soft X-ray origin, Mingo et al. (2014) confirmed that jets are the main responsible for soft X-ray emission from their sources. In our sample, jets are identified in NGC 1052 (Kadler et al. 2004), where the jet position angle would be consistent with the extended soft X-ray emission shown here.","Citation Text":["Maoz et al. 2005"],"Functions Text":["Indeed, the ultraviolet and X-ray variability detected for many of these LINERs","is in favour of a direct view of the AGNs (i.e. perpendicular to the disc under the UM).","This is consistent with the fact that most of the [O\u2009III] morphologies found for LINERs are spheroids, if we assume that the [O\u2009III] traces the NLR."],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Similarities"],"Citation Start End":[[1060,1076]],"Functions Start End":[[979,1058],[1108,1196],[1197,1345]]}
{"Identifier":"2018AandA...615A..57M__Remillard_&_McClintock_2006_Instance_1","Paragraph":"A huge amount of data at all wavelengths has been collected in the past 20 years on black hole X-ray binaries, hereafter XrBs (for a global review see Dunn et al. 2010). These objects spend most of their time in quiescence at very low accretion rates, but occasionally, they produce outbursts that last from a few months to a year. Their flux then rises by several orders of magnitude across the whole electromagnetic spectrum (see, e.g., Corbel et al. 2004, Fender et al. 2006; Remillard & McClintock 2006; Done et al. 2007, for recent reviews). During an outburst, XrBs show very different spectral and temporal states that can be easily distinguished in a hardness-intensity diagram (HID) where the X-ray luminosity is plotted versus the hardness ratio of the X-ray spectrum (see, e.g., Homan et al. 2001; Fender et al. 2004). The evolutionary track produces a typical q-shaped figure that reveals a hysteresis: outbursting XrBs have two distinct spectra with the same X-ray luminosity above 1\u20132% Eddington luminosity. At the beginning of the outburst, the system is in the so-called hard state: the spectrum has a hard power-law shape up to a few tens to hundreds of keV, requiring a very hot, optically thin plasma (referred to as the \u201ccorona\u201d). Then, when the system reaches high luminosities (up to a few tens of the Eddington luminosity), it transits within a few days through a bright intermediate state into the so-called soft state (referred to as the \u201ccold disk\u201d). In this state, the spectrum is dominated by strong and soft X-ray emission, which is commonly interpreted as thermal emission from an optically thick geometrically thin accretion flow. In the latter state, the luminosity starts to decrease and the system returns to the hard state, transiting through a faint intermediate state. The luminosities at which a system transits from hard to soft states are several times higher than the luminosity of the reverse transition (see Appendix in Dunn et al. 2010).","Citation Text":["Remillard & McClintock 2006"],"Functions Text":["These objects spend most of their time in quiescence at very low accretion rates, but occasionally, they produce outbursts that last from a few months to a year. Their flux then rises by several orders of magnitude across the whole electromagnetic spectrum (see, e.g.,"],"Functions Label":["Background"],"Citation Start End":[[479,506]],"Functions Start End":[[170,438]]}
{"Identifier":"2021MNRAS.501.2897G__Pribulla_&_Rucinski_2006_Instance_2","Paragraph":"EE Cet (ADS 2163 B) is the southern (slightly fainter) component of the visual binary WDS 02499+0856 (Mason et al. 2001). It was discovered by the HIPPARCOS mission (Perryman et al. 1997), by noticing the variability of the combined light of both visual components. Lampens et al. (2001) performed photometric measurements of the visual pair and gave (but only for one epoch) the following values V(A) = 9.47\u2009mag and V(B) = 9.83\u2009mag. Pribulla & Rucinski (2006) lists the orbital parameters (orientation and separation) of WDS 02499+0856 and gave \u03b8 = 194\u00b0, \u03c1 = 5.66\u2009arcsec and magnitude difference \u0394V = 0.07\u2009mag (the magnitude difference can be as large as \u0394V = 0.36\u2009mag, due to photometric variability of the eclipsing binary). WDS 02499+0856 turned out to be a quadruple system, when the northern component was found to be a double-lined (SB2) binary from the DDO spectroscopic observations (Pribulla & Rucinski 2006). D\u2019Angelo et al. (2006) re-confirmed the multiplicity of the system, and listed it among the contact binaries with additional components. Radial velocity observations from Rucinski et al. (2002) resulted in a well-defined circular orbit of the contact binary, with K1 = 84.05\u2009km\u2009s\u22121, K2 = 266.92\u2009km\u2009s\u22121 (q = 0.315), and an F8V spectral type. Karami & Mohebi (2007) using their own velocity curve analysis method, arrived at almost identical results for the mass ratio. Djura\u0161evi\u0107 et al. (2006) presented the first model, resulting in orbital inclination of i = 78.5\u00b0 and a fill-out factor of f = 32.69 per\u2009cent, T2 = 6314\u2009K, and T1 = 6095\u2009K, when spots were added. Their no-spot model resulted in very close value for the fill-out factor but slightly different geometrical and orbital parameters. The physical parameters derived in this study were: M1 = 1.37\u2009M\u2299, M2 = 0.43\u2009M\u2299, and mean radii R1 = 1.35\u2009R\u2299, R2 = 0.82\u2009R\u2299. It is worth noting here that the light curves analysed by these authors included the visual component in the photometric aperture with a contamination of about 54 per\u2009cent.","Citation Text":["Pribulla & Rucinski 2006"],"Functions Text":["WDS 02499+0856 turned out to be a quadruple system, when the northern component was found to be a double-lined (SB2) binary from the DDO spectroscopic observations"],"Functions Label":["Background"],"Citation Start End":[[893,917]],"Functions Start End":[[728,891]]}
{"Identifier":"2016ApJ...833..216G__Gopalswamy_et_al._2014a_Instance_1","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"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[235,252],[260,265]],"Functions Start End":[[63,233]]}
{"Identifier":"2016ApJ...833..216G___2010_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.","2010"],"Functions Text":["SEP events with gigaelectronvolt components are accompanied by type II radio bursts from meter (m) wavelengths to kilometer (km) wavelengths"],"Functions Label":["Background"],"Citation Start End":[[2149,2166],[2174,2178]],"Functions Start End":[[2007,2147]]}
{"Identifier":"2022MNRAS.516.2500C__Lin_et_al._2009_Instance_2","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"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[2166,2181]],"Functions Start End":[[1861,2164]]}
{"Identifier":"2022MNRAS.515.1276C__Reg\u00e1ly_et_al._2012_Instance_1","Paragraph":"Vortex formation in protoplanetary discs has been shown to proceed through various mechanisms. Klahr & Bodenheimer (2003) demonstrated that in discs with a radial entropy gradient, azimuthal perturbations are baroclinically unstable and can lead to the formation of vortices. The most commonly invoked mechanism to form a large-scale vortex in a protoplanetary disc is the Rossby wave instability (RWI). The RWI can be triggered at steep radial density gradients (Lovelace et al. 1999; Li et al. 2000, 2001); such a gradient may be generated at the viscosity transition between regions in which the magnetorotational instability can operate (e.g. Varni\u00e8re & Tagger 2006; Lyra et al. 2009b; Lyra & Mac Low 2012; Reg\u00e1ly et al. 2012). Steep radial density gradients are also found at the outer edge of a gap carved in the disc by a giant planet. Indeed, hydrodynamics simulations of protoplanetary discs containing a giant planet have successfully produced large-scale vortices (e.g. Li et al. 2005; Masset et al. 2006; de Val-Borro et al. 2007; Meheut et al. 2010; Lin & Papaloizou 2011; Lin 2012), and Reg\u00e1ly, Juh\u00e1sz & Neh\u00e9z (2017) presented a comparison of vortices formed at a gap edge to those formed at a viscosity transition. However, Hammer, Kratter & Lin (2017) pointed out that the short growth time-scales of the giant planet prescribed in these simulations result in unphysically large perturbations that are capable of setting up the RWI-unstable density gradient at the gap\u2019s outer edge. When the giant planet is grown over more realistic time-scales, vortex production is significantly suppressed because viscosity has time to smooth out any steep gradients and the planetary torque itself can reshape the gap edge. Under the most favourable conditions for vortex formation, vortex lifetimes were found to be limited. One may then argue that if these asymmetries are indeed vortices formed through the RWI triggered by giant planets, we must be seeing them just after the planet has formed. However, if a RWI-generated vortex typically survives for \u223c103 orbits and discs are \u223c105 orbits old, observing the vortex so soon after its formation is unlikely. It is therefore uncertain whether the non-axisymmetric substructures seen in observations can be explained by this mechanism.","Citation Text":["Reg\u00e1ly et al. 2012"],"Functions Text":["he RWI can be triggered at steep radial density gradients","such a gradient may be generated at the viscosity transition between regions in which the magnetorotational instability can operate (e.g."],"Functions Label":["Background","Background"],"Citation Start End":[[711,729]],"Functions Start End":[[405,462],[509,646]]}
{"Identifier":"2021AandA...646L...4I__Zheng_et_al._2014_Instance_1","Paragraph":"Our initial sample comprised all 40 stars of the MMT-HVS survey that were revised by Kreuzer et al. (2020), together with 30 runaway stars from the collection of Silva & Napiwotzki (2011) for which we were able to obtain spectra. This covers the majority of objects with high ejection velocities in that compilation. This group was complemented by the prototype hyper-runaway star HD 271791 (Heber et al. 2008), the potential hyper-runaway stars SDSS J013655.91+242546.0 (J0136+2425 for short, Tillich et al. 2009) and HIP 60350 (Irrgang et al. 2010), the extreme disk-runaway star PG 1610+062 (Irrgang et al. 2019), the four candidate HVSs from the LAMOST survey (Zheng et al. 2014; Huang et al. 2017; Li et al. 2018), and the probable Hills star S5-HVS 1 (Koposov et al. 2020). Based on proper motions from Gaia EDR3, we then carried out spectroscopic and kinematic analyses (see Sects. 3 and 4 for details) of all members of this initial sample to filter out those with both a high ejection velocity and a relatively well-constrained origin within the Galactic disk. In order to account for the fact that massive stars are typically ejected at lower velocity, we chose a mass-dependent threshold for the deduced 1\u03c3 upper limit of the ejection velocity, that is, 400 km s\u22121 or 320 km s\u22121 for stars with masses below or above 5\u2006M\u2299, respectively. The first cut applies to almost all stars from the MMT-HVS survey, while most of the others fall into the second category. The chosen thresholds roughly represent the limits for the classical ejection scenarios (see, e.g., Tauris 2015; Irrgang et al. 2018a, and references therein). A disk origin was granted when the 1\u03c3 lower limit of the inferred galactocentric plane-crossing radius was below 25 kpc (motivated by Xu et al. 2015), while visual inspection was used to judge whether the origin was sufficiently well constrained. These criteria left us with 14 stars from the MMT-HVS survey, 7 stars from the Silva & Napiwotzki (2011) sample, all 4 stars from the LAMOST survey, and the 5 individual targets mentioned above, yielding a final sample of 30 program stars, the names of which are listed in Table 1.","Citation Text":["Zheng et al. 2014"],"Functions Text":["This group was complemented by","the four candidate HVSs from the LAMOST survey"],"Functions Label":["Uses","Uses"],"Citation Start End":[[665,682]],"Functions Start End":[[317,347],[617,663]]}
{"Identifier":"2020ApJ...899..147F__Venot_et_al._2015_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":["Venot et al. 2015"],"Functions Text":["Numerous studies have been performed using chemical models"],"Functions Label":["Background"],"Citation Start End":[[797,814]],"Functions Start End":[[700,758]]}
{"Identifier":"2019AandA...632A.104G__Hirabayashi_et_al._2016_Instance_3","Paragraph":"Finally, our observations are consistent with the bilobate shape of the nucleus of comet 8P\/Tuttle. As noted in Sect. 1, this shape is likely common among comets because it was found for four out of the six comets for which we have spatially resolved images. This is also the case of the trans-Neptunian object 2014 MU69 (Ultima Thule) observed by the New Horizon spacecraft (Stern et al. 2019). This binary configuration has some implications for the formation and evolution of 8P\/Tuttle. A contact binary could result from (i) the accretion at low velocity of two primordial objects (Massironi et al. 2015; Davidsson et al. 2016), (ii) the disruption of a monolithic object due to excessive spin-up resulting from non-gravitational forces or YORP5 effect followed by a reaccretion (Boehnhardt 2004; \u0106uk 2007; Hirabayashi et al. 2016), or (iii) the catastrophic disruption of a monolithic object by a collision followed by a re-accretion (Jutzi & Benz 2017; Schwartz et al. 2018). On the one hand, with a low thermal inertia compared with NEAs, the YORP effect is low for comets, in particular for NIC, which have an elongated orbit and spend most of their time far from the Sun, and it may not be sufficient to increase the spin rate of the nucleus to the point where centrifugal exceed gravitational forces. On the other hand, comet 8P\/Tuttle has been on a very stable orbit for centuries, and it is likely an evolved comet, as suggested by its low activity, so that it could have been much more active in the past. For cometary nuclei, the primary cause for spin-up is torques caused by outgassing, therefore it is possible that 8P\/Tuttle formed as a monolithic body and became a contact binary after its injection into the inner Solar System as a result of excessive spin-up resulting from non-gravitational forces. This scenario has been proposed for comet 67P\/Churyumov-Gerasimenko by Hirabayashi et al. (2016). Alternatively, if the binary nature of comet 8P\/Tuttle is the result of a primordial accretion or a catastrophic collision in the early Solar Sytem, it could have persisted until now. Similar examples are offered by some binary asteroids that can be stable over the age of the Solar System (Chauvineau et al. 1991), or as proposed by Davidsson et al. (2016) for comet 67P\/Churyumov-Gerasimenko. For comet 8P\/Tuttle, it is however not possible to distinguish the solution of a binary nucleus that formed in the first billion years of our Solar System (e.g., Matonti et al. 2019) from a more recent origin following its injection into the inner Solar System (e.g., Hirabayashi et al. 2016).","Citation Text":["Hirabayashi et al. 2016"],"Functions Text":["For comet 8P\/Tuttle, it is however not possible to distinguish the solution of a binary nucleus that formed in the first billion years of our Solar System","from a more recent origin following its injection into the inner Solar System (e.g.,"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[2582,2605]],"Functions Start End":[[2314,2468],[2497,2581]]}
{"Identifier":"2020AandA...643A.178T__Friesen_et_al._(2016)_Instance_1","Paragraph":"Most of the clumps analyzed in this study lie in the optimal range of precise Tkin determination when using NH3 (1,1) and (2,2) lines (see Sects. 1 and 3.3). Figure 8 indicates that there is a large number of cold clumps with Tkin  20 K. The gas and dust are expected to be coupled at densities above about 104.5 or 105 cm\u22123 (Goldsmith 2001; Young et al. 2004). The temperatures derived from dust and gas are often in agreement in the active and dense clumps of Galactic disk clouds (Dunham et al. 2010; Giannetti et al. 2013; Battersby et al. 2014; Merello et al. 2019). This is also the case for Serpens South. Low gas temperatures are associated with Serpens South ranging from 8.9 to 16.8 K with an average of 12.3 \u00b1 1.7 K, which is consistent with the mean value of 11 \u00b1 1 K found by Friesen et al. (2016). The gas and dust temperatures (mean and standard deviations Tgas,avg ~ 12.3 \u00b1 1.7 K versus Tdust,avg ~ 13.4 \u00b1 0.9 K) scatter in Serpens South, but agree reasonably well as can be most directly seen in the right panel of Fig. 8 (blue points). However, in the HMSF region W 40, we find that the measured gas kinetic temperatures are higher than the dust temperatures (mean and standard deviations Tgas,avg ~ 25.1 \u00b1 4.9 K versus Tdust,avg ~ 19.1 \u00b1 2.2 K), which indicates that the gas and dust are not well-coupled in W 40 and that the dust can cool more efficiently than the gas. This is consistent with the relatively weak NH3 lines associated with the core region of W 40, indicating the presence of only small amounts of dense gas. This illustrates that the interplay between gas and dust cooling and heating is not uniform in the area covered by our observations. Such a difference between Tgas and Tdust is also seen in other regions and appears to be an often encountered property of massive-star formation regions. Battersby et al. (2014) and Koumpia et al. (2015) compare gas and dust temperatures in the massive-star-forming infrared dark cloud G32.02+0.05 and the high-mass-star-forming PDR S140, respectively, and find similar discrepancies between gas and dust temperatures. This likely indicates a lack of coupling between the gas and dust (Battersby et al. 2014) or could be due to the clouds beingclumpy (Koumpia et al. 2015). These may be potential mechanisms relevant to W 40, where the gas temperature is higher than the dust temperature.","Citation Text":["Friesen et al. (2016)"],"Functions Text":["Low gas temperatures are associated with Serpens South ranging from 8.9 to 16.8 K with an average of 12.3 \u00b1 1.7 K, which is consistent with the mean value of 11 \u00b1 1 K found by"],"Functions Label":["Similarities"],"Citation Start End":[[789,810]],"Functions Start End":[[613,788]]}
{"Identifier":"2021MNRAS.500.3083C__Lupi_&_Bovino_2020_Instance_1","Paragraph":"Previous theorethical studies have outlined that the [C\u2009ii] emission originates from the cold (with temperatures of a few 100 K) neutral medium and from photo-dissociation regions (PDR, Vallini et al. 2013). This seems to suggest that its presence closely traces star formation sites, resulting in a linear relation, as found by De Looze et al. (2014) and Herrera-Camus et al. (2015). While at low-redshift and close to solar metallicity such a relation is well established, as shown by several observations (De Looze et al. 2014; Herrera-Camus et al. 2018) and also numerical simulations (see, e.g. Lupi & Bovino 2020), significant deviations can arise in different ISM conditions, like at lower metallicity or in presence of a strong ionizations field, that are more typically found in the high-redshift Universe. To address the impact of such conditions, several studies have analysed the [C\u2009ii] emission from typical high-redshift galaxies, by post-processing hydrodynamic zoom-in cosmological simulations with cloudy (Ferland et al. 2017; see, e.g. Olsen et al. 2017; Pallottini et al. 2017, 2019; Katz et al. 2019), or via ad-hoc methods, as in Arata et al. (2020), or also via on-the-fly non-equilibrium chemistry (Lupi et al. 2020). The main conclusion in all these studies is that a [C\u2009ii] deficit exists at high redshift, most likely due to the starbursting nature of these galaxies rather than their metallicity, since most of these systems are close to solar (see, e.g. Vallini et al. 2015; Lupi & Bovino 2020). Other studies have evidenced a weak dependence of [C\u2009ii] on metallicity. Harikane et al. (2020) showed that the L[C\u2009ii]\/SFR ratio does not show a strong dependence on metallicity, which to first approximation was interpreted as the result of the proportionality between C abundance and metallicity Z and an inverse proportionality between PDR column density and Z in a dust-dominated shielding regime (Kaufman, Wolfir & Hollenbach 2006). In this framework, if PDR give a large conribution to the [C\u2009ii] emissivity, the [C\u2009ii] luminosity is not expected to strongly depend on Z (see also Ferrara et al. 2019; Pallottini et al. 2019).","Citation Text":["Lupi & Bovino 2020"],"Functions Text":["This seems to suggest that its presence closely traces star formation sites, resulting in a linear relation, as found by De Looze et al. (2014) and Herrera-Camus et al. (2015). While at low-redshift and close to solar metallicity such a relation is well established, as shown by several observations","and also numerical simulations (see, e.g.","significant deviations can arise in different ISM conditions, like at lower metallicity or in presence of a strong ionizations field, that are more typically found in the high-redshift Universe."],"Functions Label":["Similarities","Similarities","Compare\/Contrast"],"Citation Start End":[[600,618]],"Functions Start End":[[208,507],[558,599],[621,815]]}
{"Identifier":"2018ApJ...856...65W__Betoule_et_al._2014_Instance_1","Paragraph":"Improvements in cosmological measurement in recent years have been said to hail an era of \u201cprecision cosmology,\u201d with observations of the cosmic microwave background (CMB) temperature anisotropies (Hinshaw et al. 2013; Planck Collaboration et al. 2016a, 2016b), baryon acoustic oscillation (BAO) wiggles in the galaxy power spectrum (Beutler et al. 2011; Anderson et al. 2014; Ross et al. 2015), luminosity distance\u2013redshift relation of Type Ia supernovae (SNIa; Riess et al. 2004, 2007; Kowalski et al. 2008; Betoule et al. 2014), local distance ladder (Riess et al. 2016), galaxy clustering and weak lensing (DES Collaboration et al. 2017), and direct detection of gravitational waves (Abbott et al. 2017), providing constraints on cosmological model parameters at percent, or subpercent, level precision. Since the discovery of the accelerated expansion of the universe, these observations have cemented the emergence of the flat \u039bCDM model as the standard model of cosmology, in which global spatial curvature is zero, and the energy budget of the universe is dominated by \u201cdark energy\u201d in the form of a cosmological constant, \u039b. However, beyond the \u039bCDM paradigm, there is a large number of dark energy models aimed at explaining the accelerated expansion of the universe (see reviews by Li et al. 2011 and Joyce et al. 2015, and references therein), and so understanding the nature of dark energy remains one of the central pursuits in modern cosmology. To this end, it has become common observational practice to constrain the dark energy equation of state, w(z), and check for deviations from the \u039bCDM value of w = const. = \u22121. While observational probes do not indicate any significant departure from \u039bCDM (Huterer & Shafer 2017), there is still room to tighten constraints and thereby rule out competing alternatives for dark energy. In particular, by tuning the parameters of alternative theories of dark energy, one can recover the behavior of the \u039bCDM model at both the background expansion and perturbation levels (Li et al. 2011; Joyce et al. 2015).","Citation Text":["Betoule et al. 2014"],"Functions Text":["Improvements in cosmological measurement in recent years have been said to hail an era of \u201cprecision cosmology,\u201d with observations of","luminosity distance\u2013redshift relation of Type Ia supernovae (SNIa;","providing constraints on cosmological model parameters at percent, or subpercent, level precision."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[510,529]],"Functions Start End":[[0,133],[396,462],[709,807]]}
{"Identifier":"2020ApJ...898...25T__Pavlovskii_et_al._2017_Instance_1","Paragraph":"Recent detections of gravitational waves (GWs) have shown evidence for a high rate of black hole (BH)\u2013BH and neutron star (NS)\u2013NS mergers in the universe (Abbott et al. 2016a, 2016b, 2017a, 2017b, 2017c, 2019a; Zackay et al. 2019, 2020; Venumadhav et al. 2020). However, the proposed astrophysical pathways to mergers remain highly debated. Possible compact-object merger pathways include isolated binary evolution (Dominik et al. 2012; Kinugawa et al. 2014; Belczynski et al. 2016, 2017; Breivik et al. 2016; Giacobbo et al. 2018; Bavera et al. 2019; Spera et al. 2019) accompanied by mass transfer (Inayoshi et al. 2017a; Pavlovskii et al. 2017; van den Heuvel et al. 2017), common-envelope ejection (e.g., Paczynski 1976; Ivanova et al. 2013), envelope expansion (Tagawa et al. 2018), or chemically homogeneous evolution in a tidally distorted binary (de Mink & Mandel 2016; Mandel & de Mink 2016; Marchant et al. 2016), evolution of triple or quadruple systems (Antonini et al. 2017; Liu & Lai 2017, 2018, 2019; Silsbee & Tremaine 2017; Arca-Sedda et al. 2018; Hoang et al. 2018b; Randall & Xianyu 2018; Fragione & Kocsis 2019; Fragione et al. 2019; Michaely & Perets 2019), gravitational capture (O\u2019Leary et al. 2009; Kocsis & Levin 2012; Gond\u00e1n et al. 2018b; Rodriguez et al. 2018; Rasskazov & Kocsis 2019; Zevin et al. 2019; Samsing et al. 2020), dynamical evolution in open clusters (Banerjee 2017, 2018a, 2018b; Bouffanais et al. 2019; Kumamoto et al. 2019; Rastello et al. 2019) and dense star clusters (e.g., Portegies Zwart & McMillan 2000; O\u2019Leary et al. 2006, 2016; Samsing et al. 2014; Ziosi et al. 2014; Mapelli 2016; Rodriguez et al. 2016a, 2016b; Askar et al. 2017; Fujii et al. 2017; di Carlo et al. 2019; Zevin et al. 2019; Zhang et al. 2019), and dynamical interaction in gas-rich nuclear regions (McKernan et al. 2012, 2014, 2018; Bellovary et al. 2016; Bartos et al. 2017; Stone et al. 2017; Leigh et al. 2018; Tagawa & Umemura 2018; Yi et al. 2018; Secunda et al. 2019; Yang et al. 2019a, 2019b; Gayathri et al. 2020; McKernan et al. 2020; Tagawa et al. 2020).","Citation Text":["Pavlovskii et al. 2017"],"Functions Text":["Possible compact-object merger pathways include isolated binary evolution","accompanied by mass transfer"],"Functions Label":["Background","Background"],"Citation Start End":[[624,646]],"Functions Start End":[[341,414],[571,599]]}
{"Identifier":"2020AandA...640L..11B__Segretain_1996_Instance_3","Paragraph":"Another possibly important cooling delay may arise from the phase separation of 22Ne during crystallization (Isern et al. 1991; Althaus et al. 2010). Our current best understanding is that at the small 22Ne concentrations typical of C\/O white dwarfs (\u223c1% by number), the presence of 22Ne should not affect the phase diagram, except near the azeotropic point of the C\/O\/Ne phase diagram. Thus, the crystallization of the C\/O core initially proceeds as in the case without 22Ne with no redistribution of neon ions between the solid and liquid phases. After a significant fraction of the core has crystallized, the temperature approaches the azeotropic point and the existing calculations indicate that the liquid phase is enriched in 22Ne relative to the solid (Segretain 1996; Garc\u00eda-Berro et al. 2008). The 22Ne-poor solid is lighter than the surrounding liquid and floats upward where it eventually melts. This gradually displaces the 22Ne-rich liquid downward toward the solid\u2013liquid interface until the azeotropic composition is reached, thereby releasing a considerable amount of gravitational energy. Given our very limited knowledge of the ternary C\/O\/Ne phase diagram (Segretain 1996; Hughto et al. 2012), this effect cannot be quantitatively implemented in our evolution models. However, we note that our current understanding of 22Ne phase separation is remarkably consistent with the missing cooling delay. In Fig. 2 we show the luminosity function obtained by adding an artificial 0.6 Gyr delay when 60% of the core is crystallized. These parameters are entirely consistent with those found in preliminary studies (Segretain 1996; Garc\u00eda-Berro et al. 2008) and yield an excellent fit to the crystallization pile-up3. Based on the current (albeit limited) knowledge of the C\/O\/Ne phase diagram, we propose that the phase separation of 22Ne in the advanced stage of crystallization significantly contributes to the pile-up in the luminosity function of 0.9\u22121.1\u2006M\u2299 white dwarfs (Fig. 2).","Citation Text":["Segretain 1996"],"Functions Text":["These parameters are entirely consistent with those found in preliminary studies"],"Functions Label":["Similarities"],"Citation Start End":[[1626,1640]],"Functions Start End":[[1544,1624]]}
{"Identifier":"2021MNRAS.503.1319G__Kochanek_1992_Instance_1","Paragraph":"Independent of the traditional redshift surveys, in this work, we constrain the VDF of ETGs using the statistics of strong gravitational lensing systems (Turner et al. 1984; Biesiada 2006; Cao et al. 2012b,c). Assuming the concordance cosmological model (\u039bCDM), several efforts have been made to include the distribution of lensed image separations in the study of the galaxy mass profiles and the evolution history of galaxies. Based on the CLASS and PANELS lens sample, the first attempt to constrain the redshift evolution of galaxies (since redshift z \u223c 1) was presented in Chae & Mao (2003). This study was then extended to the study of the shape of the VDF and the characteristic velocity dispersion (Chae 2005). However, the sample size of the data available at that time did not allowed for a firm determination of the galaxy VDF. Here, we present a new approach to derive the VDF based on the lens redshift distribution (Kochanek 1992) and to constrain its evolution out to z \u223c 1, given its strong dependency on the dynamical properties of galaxies (i.e. stellar velocity dispersion) and the number density of gravitational lenses (i.e. galaxy evolution). Compared with the previous works focusing on image separation distributions, constraining a VDF through lens redshift test is unique and promising, since it does not require the knowledge of the total lensing probability and the magnification bias in the sample (Ofek, Rix & Maoz 2003). The advantages of the lens redshift test have been extensively discussed in Matsumoto & Furamase (2008), Koopmans et al. (2009), Oguri et al. (2012), and Cao et al. (2012a). Therefore, it will be rewarding to investigate the VDF and evolution of the lensing galaxies by adopting the cosmological parameters determined by Planck and using a new lens sample better representing the distribution of the galaxy properties. In this work, we focus on a newly compiled sample of 157 galaxy-scale strong lensing systems, which are all early-type lenses (E or S0 morphologies) without significant substructures or close companion galaxies (Chen, Li & Shu 2019). Throughout the paper, we assume the concordance cosmology by adopting the cosmological parameters determined by the Planck 2016 data (Ade et al. 2016).","Citation Text":["Kochanek 1992"],"Functions Text":["Here, we present a new approach to derive the VDF based on the lens redshift distribution","and to constrain its evolution out to z \u223c 1, given its strong dependency on the dynamical properties of galaxies (i.e. stellar velocity dispersion) and the number density of gravitational lenses (i.e. galaxy evolution)."],"Functions Label":["Uses","Uses"],"Citation Start End":[[930,943]],"Functions Start End":[[839,928],[945,1164]]}
{"Identifier":"2022AandA...663A..44K__Blundell_et_al._1999_Instance_1","Paragraph":"The remarkably bright nature of radio-jetted active galactic nuclei (AGNs) allows us to observe them at extreme redshifts and thus use their properties as an observational tracer of cosmological principles (Wang et al. 2021). Blazars, in particular, benefit from increased apparent luminosities due to relativistic boosting effects (Cohen et al. 2007b), and can thus be detected over a wide range of the electromagnetic spectrum, although at the cost of being observationally compact objects with small angles between their jets and the line of sight. Conversely, AGNs with radio jets that are at a larger angle to the line of sight are harder to detect with increasing distance. Because they can be observed over such a wide range of redshifts, they provide a unique insight into cosmology, galaxy evolution, and the evolution of AGNs (Dunlop & Peacock 1990; Georgakakis et al. 2017). In addition to probing the universe by observing targets at different redshifts, one should also account for differences in their evolutionary stage, the environment they are embedded in at that time, and their interactions with that environment. Many radio surveys have already been performed to investigate radio-loud AGN populations (e.g. Becker et al. 1995; Condon et al. 1998; Cohen et al. 2007a; Intema et al. 2017) and most find self-consistent relative number ratios of radio galaxies and blazars up to a redshift of ~3 (Volonteri et al. 2011). Beyond this redshift, matters are complicated by uncertainties about the density evolution and the build-up of high black hole masses in the early universe (Blundell et al. 1999; Shankar et al. 2008). However, there seems to be a consensus that there is a relative lack of higher redshift radio galaxies even when accounting for evolutionary effects and detection limits (e.g. Wu et al. 2017; Hodges-Kluck et al. 2021, and references therein). The reason for the deficit is still not fully understood. The currently favoured explanation was suggested by Ghisellini et al. (2015): interaction of the extended radio emission with the cosmic microwave background (CMB) could efficiently quench the brightness of the extended radio lobes. Morabito & Harwood (2018) find evidence to support this model based on a comparison of simulations and observational data. In this scenario the CMB energy density dominates over the magnetic energy density at very high redshifts, so that the jet electrons interact with the CMB photons by inverse Compton (IC) scattering to cool, while the synchrotron radiation is suppressed. Although quenched, the steep-spectrum isotropic radiation of the extended structures could nevertheless be detected by telescopes operating in the long wavelength radio regime, which can test and guide the theoretical models. Ghisellini et al. (2015) proposed a number of suitable blazars to spotlight this issue, and published expected radio fluxes for different model parameters. The International LOFAR Telescope (ILT) (van Haarlem et al. 2013) offers the resolution, sensitivity, and observing wavelengths necessary to detect the extended emission from these blazars and to address the questions associated with its possible suppression. For this study, we processed and analysed one ILT observation of GB 1508+5714 at z = 4.30 (Hook et al. 1995), which is one of the most distant quasars with a detected X-ray jet (Yuan et al. 2003; Siemiginowska et al. 2003). Throughout the paper we use the following cosmological parameters: H0 = 71 km s\u22121 Mpc\u22121, \u03a9m = 0.27, and \u03a9\u039b = 0.73, hence a luminosity distance of 39.8 Gpc and a conversion scale where 1\u2033 is about 6.9 kpc for the given source.","Citation Text":["Blundell et al. 1999"],"Functions Text":["Beyond this redshift, matters are complicated by uncertainties about the density evolution and the build-up of high black hole masses in the early universe"],"Functions Label":["Background"],"Citation Start End":[[1596,1616]],"Functions Start End":[[1439,1594]]}
{"Identifier":"2016AandA...585A..48G__Mostardi_et_al._(2013)_Instance_2","Paragraph":"On the same SSA22 field, Nestor et al. (2011) and Nestor et al. (2013) show nine LBGs and 20 Lyman-\u03b1 emitters (LAEs) with LyC detection out of a sample of 41 LBGs and 91 LAEs (all spectroscopically confirmed). They started from a different narrowband image centred at ~3640 \u00c5, which is deeper than that used by Iwata et al. (2009), at ~3590 \u00c5. A careful analysis of their LBG detections, however, shows that even in this case the LyC emission for many z ~ 3 sources is offset by 0.4\u20131.0 arcsec. The observed ratio between the 900 and 1500 \u00c5 rest-frame emission is difficult to reconcile with that expected by standard stellar populations (Vanzella et al. 2012a). The HST images in I814W filter available for a few of them show the presence of clearly separated galaxies, sometimes fainter at 1500 \u00c5 than in the ionizing continuum (i.e. their C16 object), with a resulting escape fraction well exceeding 1000% if estimated in the LyC position. For the majority of them, no HST imaging is available but even ground-based images often show the presence of slightly offset emission in LyC, w.r.t. the non-ionizing continuum. Similar conclusions can be reached for the Mostardi et al. (2013) sample where they adopt the same analysis as in Nestor et al. (2013). At z ~ 2.8, they found four LyC emitters out of 49 LBG galaxies and seven LyC emitters out of 91 LAEs. In this case the lack of high spatial resolution data from HST for the majority of the sample prevents any detailed analysis about possible contamination by interlopers\/foregrounds. These conclusions have been strengthened by recent observations by Siana et al. (2015), who found no convincing detection in their deep HST imaging with WFC3-UVIS of five LyC emitters extracted from the sample of Nestor et al. (2011), or by Mostardi et al. (2015), who only found one robust LyC emitter after a reanalysis of a sample of 16 galaxies by Mostardi et al. (2013). More interestingly, the only robust candidate LyC emitter by Mostardi et al. (2015), the galaxy MD5b, has an observed ratio FUV\/FLyC = 4.0 \u00b1 2.0, equivalent to a relative escape fraction of 75%, when assuming complete transmission of the IGM. Instead, if a mean value of \u27e8 exp(\u2212\u03c4IGM) \u27e9  = 0.4 at z ~ 3.1 is adopted, following Inoue et al. (2014), the relative escape fraction of MD5b turns out to be 188%. Imposing the constraint of a physical value for the relative escape fraction of \\hbox{$f^{\\rm rel}_{\\rm esc}<100\\%$}fescrel100%, the LoS of MD5b must be very transparent, exp(\u2212\u03c4IGM) > 0.75, which corresponds to a probability 10-4, according to Inoue et al. (2014). This could be an indication that this galaxy is also a low-z contaminant, similar to the other cases studied by Mostardi et al. (2015). ","Citation Text":["Mostardi et al. (2013)"],"Functions Text":["These conclusions have been strengthened by recent observations by Siana et al. (2015), who found no convincing detection in their deep HST imaging with WFC3-UVIS of five LyC emitters extracted from the sample of Nestor et al. (2011), or by Mostardi et al. (2015), who only found one robust LyC emitter after a reanalysis of a sample of 16 galaxies by"],"Functions Label":["Similarities"],"Citation Start End":[[1894,1916]],"Functions Start End":[[1542,1893]]}
{"Identifier":"2019MNRAS.488.4638L__Drabek-Maunder_et_al._2016_Instance_2","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)"],"Functions Text":["However,","reported that there is no correlation between the turbulent and outflow energies."],"Functions Label":["Differences","Differences"],"Citation Start End":[[1780,1808]],"Functions Start End":[[1771,1779],[1832,1913]]}
{"Identifier":"2021MNRAS.500.3083C__Meneghetti_et_al._2017_Instance_1","Paragraph":"A natural help in observational studies of particularly faint objects is provided by strong gravitational lensing. In the last few years, much progress in this field has been driven by deep observations of massive galaxy clusters, carried out in the context of large Hubble Space Telescope programmes, in particular the Hubble Frontier Fields (HFF) survey (Lotz et al. 2017). Robust lens models of galaxy clusters are built thanks to the identification of large numbers of multiply lensed sources, which can span a large redshift range (Meneghetti et al. 2017; Bergamini et al. 2019). Thanks to strong-lensing effects produced by massive clusters of galaxies located along the line of sight, sources can be magnified by large factors (ranging from \u03bc of the order of a few up to \u223c100), allowing faint, compact objects to be studied with very high spatial resolution and signal-to-noise ratios (SNR) (and suprequent recurrencies) and in some cases, to probe their structural parameters down to scales of a few \u223c10 pc. In this context, the determination of the redshift of the images is a key problem. In many cases, photometric redshift estimates are accessible, but a safe determination of the redshift generally has to rely upon a spectroscopic detection, achievable only with deep observations. In this regard, in recent times considerable, further progress was possible thanks to the significant use of powerful instruments such as the Multi Unit Spectroscopic Explorer (MUSE, Bacon et al. 2015) mounted on the VLT, which has enabled the spectroscopic confirmation of hundreds of multiple images at high redshift (z > 3, e.g. Caminha et al. 2017; Lagattuta et al. 2019). This has enhanced the production of highly accurate lens models, significantly mitigating systematic uncertainties in the computation of magnification maps in lensed fields. This recent progress has allowed one to determine absolute physical quantities such as luminosities, sizes, stellar mass values and SFRs of new objects, which before were impossible to study in non-lensed fields.","Citation Text":["Meneghetti et al. 2017"],"Functions Text":["Robust lens models of galaxy clusters are built thanks to the identification of large numbers of multiply lensed sources, which can span a large redshift range"],"Functions Label":["Background"],"Citation Start End":[[537,559]],"Functions Start End":[[376,535]]}
{"Identifier":"2020ApJ...904...91V__Zanon_et_al._2018_Instance_1","Paragraph":"When in a rotation-powered state, many spiders exhibit a nonthermal (synchrotron radiation; SR) orbitally modulated emission component in the X-ray band, attributed to particle acceleration in an intrabinary pulsar wind termination shock (IBS) and Doppler boosting of a bulk flow along the shock tangent (e.g., Wadiasingh et al. 2015, 2017; Romani & Sanchez 2016). This type of geometry\/radation picture was identified over two decades ago for the BW system involving the MSP B1957+20 (Arons & Tavani 1993) and later for the gamma-ray-emitting PSR B1259-63\/Be star binary system (Tavani & Arons 1997). The existence of a shock is also indicated by orbital-phase- and frequency-dependent radio eclipses of the MSP (since plasma structures at the shock attenuates the radio), where in RBs the eclipse fraction can be >50% (e.g., Archibald et al. 2009, 2013; Broderick et al. 2016; Miraval Zanon et al. 2018) of the orbital phase (while the MSP remains largely uneclipsed at inferior conjunction of the pulsar). The hard power laws inferred in X-rays imply emission due to an energetic electron population. The X-ray spectra may furthermore extend to at least \u227350 keV with no suggestion of spectral cutoffs (e.g., Tendulkar et al. 2014; Kong et al. 2017b; Al Noori et al. 2018), constraining the shock magnetic field to Bsh \u2273 1 G (Wadiasingh et al. 2017). For a given pulsar spin-down power \n\n\n\n\n\n erg s\u22121, the magnetic field Bsh at the shock can be bounded using the Poynting flux \n\n\n\n\n\n for a magnetic field Bw \u221d R\u22121 in the striped wind outside the light cylinder. The flux can be integrated over a spherical surface of area \n\n\n\n\n\n at the shock radius Rsh to yield an isotropic electromagnetic luminosity \n\n\n\n\n\n. One therefore arrives at the constraint \n\n\n\n\n\n, which is detailed in Equation (1), being generally less than 100 G. This may be recast as a condition that the ratio of electromagnetic energy density to particle pressure, \u03c3, is less than unity, i.e., that \n\n\n\n\n\n is dominated by the plasma wind contribution. As will be apparent in due course (see Section 3), if the particle acceleration is as fast and efficient (attaining the synchrotron burn-off limit) as in pulsar wind nebulae, MeV synchrotron components may be observable by future medium-energy, gamma-ray Compton\/pair telescopes, such as e-ASTROGAM (De Angelis et al. 2017) and AMEGO.8\n\n8\nAMEGO: https:\/\/asd.gsfc.nasa.gov\/amego\/index.html (McEnery et al. 2019).\n\n","Citation Text":["Miraval Zanon et al. 2018"],"Functions Text":["The existence of a shock is also indicated by orbital-phase- and frequency-dependent radio eclipses of the MSP (since plasma structures at the shock attenuates the radio), where in RBs the eclipse fraction can be >50% (e.g.","of the orbital phase (while the MSP remains largely uneclipsed at inferior conjunction of the pulsar)."],"Functions Label":["Background","Background"],"Citation Start End":[[879,904]],"Functions Start End":[[602,825],[906,1008]]}
{"Identifier":"2015MNRAS.454.3134M__Wilkinson_&_Uttley_2009_Instance_1","Paragraph":"In order to study the behaviour of the soft residuals and attempt to distinguish between possible origins, we require a source that evolves across a dynamic range in spectral hardness; as can be seen from fig. 8 and tables 2 and 3 of M15, the ideal source for such a study is NGC 1313 X-1. However, both the inferred (deabsorbed) hardness ratio and our ability to reliably characterize the residuals is sensitive to the modelling of the continuum. M15 show, via use of covariance spectra (see Wilkinson & Uttley 2009; Uttley et al. 2014 for a review of the technique), that the ULX continuum emission originates from at least two components that cross at \u223c1 keV. Mismodelling of the continuum will therefore have an adverse effect on detecting atomic features imprinted between 0.7 and 2 keV, whilst an inability to separate out the components in the time-averaged spectrum can lead to misleading hardness ratios. To better address this problem we model each of the observations (using both PN and MOS data) using the continuum described above (and including a normalization offset to account for differing responses, typically within \u223c10 per cent of unity), finding that in five out of the available 11 observations we cannot distinguish between a solution where the soft component is intrinsically strong or weak. Middleton et al. (2015) show, through use of the covariance spectrum, that the variability is entirely contained within the hard component but that a soft excess must still be present (see also Middleton et al. 2011) whilst the highest S\/N data clearly show that the soft component dominates over the hard component below \u223c1 keV. Whilst we can force the spectra for the five observations into what we believe to be a more appropriate deconvolution (namely by fixing the normalization of the soft component \u2013 see M15), accurate modelling of the residuals requires the normalizations to be free to vary and so we exclude these observations from our subsequent analysis.","Citation Text":["Wilkinson & Uttley 2009"],"Functions Text":["M15 show, via use of covariance spectra (see","for a review of the technique), that the ULX continuum emission originates from at least two components that cross at \u223c1 keV."],"Functions Label":["Background","Background"],"Citation Start End":[[493,516]],"Functions Start End":[[448,492],[537,662]]}
{"Identifier":"2016ApJ...826..134D__Wiedenbeck_et_al._2013_Instance_1","Paragraph":"Solar energetic particle (SEP) events as can be observed in the Earth\u2019s orbit are determined by a combination of the underlying acceleration, injection and transport conditions, and therefore carry fundamental information on these processes. In the classical picture introduced by Reames (1999), SEP events are separated into two groups according to their characteristics: impulsive and gradual. Impulsive events, which are electron and 3He-rich, are believed to be flare-associated and therefore show a narrow angular spread in interplanetary space. These events have an impulsive increase and a shorter decay. Gradual events, on the other hand, show much broader angular SEP spreads because of their association to extended shock fronts. The gradual increase is caused by the continuous acceleration by the coronal mass ejection (CME)-driven shock front. Gradual events are usually observed in protons and heavy ions, which can be efficiently accelerated by shocks. Recent observations during the STEREO era, however, question this simple picture, especially since the nature of events showing wide longitudinal SEP spreads appears to be more complicated. It was found that 3He-rich events extend over more than 130\u00b0 in longitude (Wiedenbeck et al. 2013), and that near-relativistic electron events can spread up to almost all around the Sun. Dresing et al. (2012) suggested that a strong perpendicular transport may cause the wide SEP distribution. On the other hand, it is also possible that the injection of SEPs at the Sun can be much wider than expected from a flare-like point source. The driver of such extended injection regions may be of a different nature. Diverging magnetic field lines below the source surface (e.g., Klein et al. 2008) as well as the presence of a coronal shock may play a role. The importance of EIT-waves intersecting the magnetic footpoints of far separated spacecraft are also under discussion (Rouillard et al. 2012; Park et al. 2013; Lario et al. 2014). Statistical analyses of multi-spacecraft SEP events (e.g., Lario et al. 2013) show that event characteristics, such as peak intensities, onset delays, rise times, and anisotropies, show strongly varying longitudinal dependencies from event to event. The lack of a clear correlation to a single mechanism, such as the CME speed or width as a representative of the shock strength or extent (Richardson et al. 2014), as well as varying event characteristics suggests that there is not a single process providing these wide SEP distributions (Dresing et al. 2014). Therefore, when describing an SEP event, a possible mixture of different processes as well as a varying strength of these mechanisms must be taken into account.","Citation Text":["Wiedenbeck et al. 2013"],"Functions Text":["Recent observations during the STEREO era, however, question this simple picture, especially since the nature of events showing wide longitudinal SEP spreads appears to be more complicated. It was found that 3He-rich events extend over more than 130\u00b0 in longitude","and that near-relativistic electron events can spread up to almost all around the Sun."],"Functions Label":["Background","Background"],"Citation Start End":[[1233,1255]],"Functions Start End":[[968,1231],[1258,1344]]}
{"Identifier":"2022MNRAS.516.5618P__Helton_et_al._2021_Instance_1","Paragraph":"At present, direct detection of the CGM in emission poses an observational challenge due to its diffuse nature (with hydrogen densities of the order of nH 0.1 cm\u22123). Cosmological hydrodynamical simulations concur that the emission signal is faint by current observational standards (Augustin et al. 2019; P\u00e9roux et al. 2019; Corlies et al. 2020; Wijers, Schaye & Oppenheimer 2020; Byrohl et al. 2021; Nelson et al. 2021; Wijers & Schaye 2021). For these reasons, detections in emission at high-redshifts are currently limited to deep fields (Wisotzki et al. 2016; Leclercq et al. 2017; Wisotzki et al. 2018; Leclercq et al. 2020, 2022) or regions around bright quasars (Cantalupo et al. 2005; Arrigoni Battaia et al. 2015; Farina et al. 2019; Lusso et al. 2019; Mackenzie et al. 2021) while detections at z1 are now becoming available (Epinat et al. 2018; Johnson et al. 2018; Chen et al. 2019; Rupke et al. 2019; Burchett et al. 2021; Helton et al. 2021; Zabl et al. 2021). Given this limitation, absorption lines detected against bright background quasars at UV and optical wavelengths provide the most compelling way to study the distribution, kinematics and chemical properties of CGM atomic gas to date. In these quasar absorbers, the minimum column density (which is tightly correlated to the volumic gas density, see Rahmati et al. 2013) that can be detected is set by the apparent brightness of the background sources and thus the detection efficiency is independent of the redshift of the foreground absorber host galaxy. In addition, absorption line-based metallicity measurements are independent of excitation conditions (Kewley, Nicholls & Sutherland 2019; Maiolino & Mannucci 2019). In fact, unlike emission lines metallicity estimates, they are largely insensitive to density or temperature and high column density systems tracing neutral gas require no assumption on a local source of excitation (Vladilo et al. 2001; Dessauges-Zavadsky et al. 2003). Importantly, multiple state-of-the-art cosmological hydrodynamical simulations and early observational results indicate that the chemical properties of the CGM gas probed in absorption show an inhomogeneous metal distribution around galaxies with indication of a trend with galaxy orientation (P\u00e9roux & Howk 2020; Wendt et al. 2021).","Citation Text":["Helton et al. 2021"],"Functions Text":["For these reasons, detections in emission at high-redshifts are currently limited to deep fields","while detections at z1 are now becoming available"],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[936,954]],"Functions Start End":[[444,540],[785,834]]}
{"Identifier":"2015MNRAS.450.3840C__Dunn_et_al._2010_Instance_1","Paragraph":"In order to understand the uniqueness of the recurrence of the 4U 1630\u2013472 outbursts and to correlate this with the lack of high-energy emission, we investigated the ASM and BAT light curves of some of the brightest and active transient black hole X-ray binaries (BHXRBs) during the seven years (2005\u20132011) in which both instruments were operational. Many studies of the outburst evolution of these sources have been reported in the literature (e.g. Yu & Yan 2009; Capitanio et al. 2010; Dunn et al. 2010, and references therein); the various outbursts observed do not have any detectable recurrence period and, in addition, show a complex behaviour in both soft and hard energy ranges and very different luminosities. The only exception is H 1743\u2013322; this BHC has shown only for some years outbursts equally spaced in time, as reported by Capitanio et al. (2010). In analogy with 4U 1630\u2013472, the outburst of H 1743\u2013322 that occurred in 2007 has a 2\u201312 keV behaviour mostly identical to some of the subsequent periodical outbursts but showing a fainter hard X-ray emission; the H 1743\u2013322 ASM 2007 peak flux (1\u201312 keV) is \u223c250 mCrab and the BAT 2007 peak flux (15\u201350 keV) is \u223c60 mCrab. The subsequent H 1743\u2013322 outburst (2008\u20132010) presents a similar 1\u201312 keV flux but with a 15\u201350 keV flux of \u223c200 mCrab. As Fig. 10 shows, the difference in the ratio between the BAT HS\u2013HSS transition luminosity and the ASM peak flux of H 1743\u2013322 is not enough to bring the outburst totally out of the correlation reported by Yu & Yan (2009) as in the case of the 2006 and 2008 outbursts of 4U 1630\u2013472. However, as reported by Corbel & Tzioumis (2008), the HS\u2013HSS transition occurred before the peak of the hard X-ray emission and thus it is only an upper limit. For the 2007 and 2008\u20132010 outbursts of H 1743\u2013322, the relation between the hard X-ray luminosity peak and the outburst waiting time is also not respected. In contrast, during the subsequent outbursts, a similar waiting time corresponds to a similar peak luminosity in the 15\u201350 keV energy range.","Citation Text":["Dunn et al. 2010"],"Functions Text":["Many studies of the outburst evolution of these sources have been reported in the literature (e.g.","the various outbursts observed do not have any detectable recurrence period and, in addition, show a complex behaviour in both soft and hard energy ranges and very different luminosities."],"Functions Label":["Background","Background"],"Citation Start End":[[488,504]],"Functions Start End":[[351,449],[531,718]]}
{"Identifier":"2021ApJ...923..233G__Carlton_et_al._2011_Instance_1","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"],"Functions Text":["The angular expansion rate of 0.64 arcsec yr\u22121 obtained from comparing X-ray images from 2007 and 2009","gives an upper limit for the age of about 160 yr, less if (as is almost certainly the case) deceleration has occurred;"],"Functions Label":["Uses","Uses"],"Citation Start End":[[866,885]],"Functions Start End":[[762,864],[887,1005]]}
{"Identifier":"2020AandA...644A..97C__Leroy_et_al._2013_Instance_1","Paragraph":"Major nearby galaxy cold gas mapping surveys (Regan et al. 2001; Wilson et al. 2009; Rahman et al. 2011; Leroy et al. 2009; Donovan Meyer et al. 2013; Bolatto et al. 2017; Sorai et al. 2019; Sun et al. 2018) have focused on observations of the molecular gas (through CO lines). Despite a few notable exceptions (e.g. Alatalo et al. 2013; Saintonge et al. 2017), these surveys observed mainly spiral or infrared-bright galaxies (i.e. galaxies with significant star formation) and have furthered our understanding of how star formation happens, rather than how it stops. This boils down to quantifying the relation between molecular gas and star formation rate (SFR), which appears nearly linear in nearby discs (Kennicutt 1998; Bigiel et al. 2008; Leroy et al. 2013; Lin et al. 2019). This relationship is often parametrised via the ratio between the SFR and the molecular gas mass (Mmol), which is called the molecular star formation efficiency (SFE = SFR\/Mmol = 1\u2215\u03c4dep), where the inverse of the SFE is the depletion time, \u03c4dep. The depletion time indicates how much time is necessary to convert all the available molecular gas into stars at the current star formation rate. On kpc scales and in the discs of nearby star-forming galaxies, \u03c4dep is approximately constant around 1\u20132 Gyr (Bigiel et al. 2011; Rahman et al. 2012; Leroy et al. 2013; Utomo et al. 2017), and it appears to weakly correlate with many galactic properties such as stellar mass surface density or environmental hydrostatic pressure (Leroy et al. 2008; Rahman et al. 2012). Nevertheless, small but important deviations for a constant SFE have been noticed, which can be the first hints of star formation quenching. In some galaxies, the depletion time in the centres appear shorter (Leroy et al. 2013; Utomo et al. 2017) or longer (Utomo et al. 2017) with respectto their discs. These differences may correlate with the presence of a bar or with galaxy mergers (Utomo et al. 2017; see also Muraoka et al. 2019) and do not seem to be related to unaccounted variation in the CO-to-H2 conversion factor (Leroy et al. 2013; Utomo et al. 2017). Spiral arm streaming motions have also been observed to lengthen depletion times (Meidt et al. 2013; Leroy et al. 2015).","Citation Text":["Leroy et al. 2013"],"Functions Text":["This boils down to quantifying the relation between molecular gas and star formation rate (SFR), which appears nearly linear in nearby discs"],"Functions Label":["Background"],"Citation Start End":[[747,764]],"Functions Start End":[[569,709]]}
{"Identifier":"2015MNRAS.451.2610H__Hinshaw_et_al._2013_Instance_1","Paragraph":"The ability to test the hypothesis that a set of data samples is drawn from some general n-dimensional probability distribution $f(\\boldsymbol {x})$ has an interesting application in the validation of Bayesian inference analyses (indeed this application provided our original motivation for seeking to extend the K\u2013S test to multiple dimensions). Bayesian methods are now pervasive across all branches of science and engineering, from cognitive neuroscience (Doya et al. 2007) and machine learning (Bishop 2006) to spam filtering (Sahami et al. 1998) and geographic profiling (Collins, Gao & Carin 1999; Le Comber & Stevenson 2012). In precision cosmology, Bayesian inference is the main tool for setting constraints on cosmological parameters (Hinshaw et al. 2013; Planck Collaboration XVI 2014), but very few attempts have been made to assess whether the derived posterior probability distributions are a truthful representation of the actual parameter constraints one can infer from the data in the context of a given physical model. This lack of validation has been highlighted by several authors, with the strong dependence of the inference results on the priors being of particular concern (Efstathiou 2008; Linder & Miquel 2008). There have been attempts to address this issue, ranging from the approach of Cook, Gelman & Rubin (2006), which was designed with software validation solely in mind, to a method based on the inverse probability integral transform (Smirnov transform) applied to posterior distributions that extends to spaces of higher dimensionality via marginalization (Dorn, Oppermann & En\u00dflin 2013). Also, validation of the Bayesian source-finding algorithm of Carvalho et al. (2012) was performed in Planck Collaboration XXIX (2014), but only point estimates deduced from the posterior distributions were actually verified. Our method for addressing this problem is based on our applying our multidimensional extension of the K\u2013S test to sets of Monte Carlo simulations of the data and the posterior distributions derived therefrom. In particular, it can take advantage of the fact that one may typically generate simulations that are of greater sophistication and realism than may be modelled in the inference process, and thus allows for a more thorough validation of the inference than has been possible with the methods developed previously. In particular, our validation procedure enables us to test all the assumptions made (explicitly or implicitly) in the inference process, such as the statistical description of the data, model assumptions and approximations, as well as the software implementation of the analysis. Moreover, we consider the full posterior distribution, regardless of its dimensionality and form, without the need to resort to marginalization, and thereby keeping intact its n-dimensional character.","Citation Text":["Hinshaw et al. 2013"],"Functions Text":["In precision cosmology, Bayesian inference is the main tool for setting constraints on cosmological parameters"],"Functions Label":["Background"],"Citation Start End":[[745,764]],"Functions Start End":[[633,743]]}
{"Identifier":"2021ApJ...913..115A__Linden_et_al._2012_Instance_2","Paragraph":"The origin of the GC VHE emission remains undetermined, due in part to source confusion and the limitations of current instruments. The source of VER J1745\u2013290 may be Sgr A* (Atoyan & Dermer 2004; Aharonian & Neronov 2005; Ballantyne et al. 2011; Chernyakova et al. 2011; Fatuzzo & Melia 2012; Kusunose & Takahara 2012; Fujita et al. 2017; Rodr\u00edguez-Ram\u00edrez et al. 2019) or PWN G359.95-0.04 (Wang et al. 2006; Hinton & Aharonian 2007), with which it is spatially coincident (Acero et al. 2010). Other possible origins include the annihilation of dark matter particles (Bergstr\u00f6m et al. 2005a, 2005b; Horns 2005; Profumo 2005; Aharonian et al. 2006c; Belikov et al. 2012; Cembranos et al. 2012, 2013; Gammaldi et al. 2016) or a population of millisecond pulsars (Bednarek & Sobczak 2013; Bartels et al. 2016; Gu\u00e9pin et al. 2018). The mechanism of gamma-ray emission may be predominantly due to hadronic processes, where relativistic protons interact with gas and subsequently produce gamma rays through neutral pion decay (Aharonian & Neronov 2005; Ballantyne et al. 2011; Chernyakova et al. 2011; Fatuzzo & Melia 2012; Linden et al. 2012; Gu\u00e9pin et al. 2018), leptonic processes where gamma rays are produced when electrons and positrons undergo inverse Compton scattering off a radiation field (Atoyan & Dermer 2004; Hinton & Aharonian 2007; Kusunose & Takahara 2012; Lacroix et al. 2016), or a combination of processes (hybrid scenario), where leptons produce high-energy, but not VHE, gamma rays (Guo et al. 2013). Both the correlation of VHE emission with the CMZ and the lack of a cutoff in the diffuse spectrum support a hadronic scenario, capable of explaining both VER J1745\u2013290 and the diffuse emission (Aharonian et al. 2006b; Linden et al. 2012; Abramowski et al. 2016). Measurement of the diffuse spectrum by H.E.S.S. up to energies of tens of TeV with no evidence of a cutoff has also been interpreted as evidence for the existence of PeV protons within the central 10 parsecs of the GC, accelerated by Sgr A* (Abramowski et al. 2016). While cosmic rays are known to extend up to PeV energies (e.g., H\u00f6randel 2003), few, if any, accelerators of PeV cosmic rays, or \u201cPeVatrons,\u201d have been clearly established (e.g., Abramowski et al. 2016; Abeysekara et al. 2020). Discovering the nature of PeVatrons in our Galaxy is thus a particularly important step in understanding the origins of cosmic rays.","Citation Text":["Linden et al. 2012"],"Functions Text":["Both the correlation of VHE emission with the CMZ and the lack of a cutoff in the diffuse spectrum support a hadronic scenario, capable of explaining both VER J1745\u2013290 and the diffuse emission"],"Functions Label":["Background"],"Citation Start End":[[1737,1755]],"Functions Start End":[[1518,1711]]}
{"Identifier":"2018MNRAS.479.3923Z__Welsh_&_Lallement_2012_Instance_1","Paragraph":"Simulations are performed in Fig. 6 to compare the spectrum profiles of the S\u2009ii triplets for a synthetic absorption spectrum of a DLA with and without the magnetic alignment included. Fig. 6 was designed up to illustrate the effect in a realistic instrumental set-up, with a given typical pixel size and S\/N for an optical spectrum taken by an 8m-class telescope. The graphical use of steps instead of curves, which is common in absorption spectroscopy, is intended to visualize the pixel-by-pixel noise variations in the data. The three transitions of singly ionized Sulfur (Sii; upper ionization potential 23.3 eV) at \u03bb\u03bb1250.58, 1253.81, 1259.52\u00c5 represent important tracers for neutral and weakly ionized gas in the local interstellar and intergalactic medium and in distant galaxies (e.g. Richter et al. 2001; Welsh & Lallement 2012; Fox, Richter & Fechner 2014a; Kisielius et al. 2014). Being an \u03b1 element, singly ionized Sulfur only has a weak depletion into dust grains (e.g. Savage & Sembach 1996). Thus the interstellar Sulfur abundance is often used as a proxy for the \u03b1-abundance in the gas. In addition, Sulfur has a relatively low-cosmic abundance (Asplund et al. 2009) and under typical interstellar conditions (in particular in low-metallicity environments) these lines are not saturated. The three lines are observed in the same wavelength region with identical S\/N. The important parameters for simulations are presented in the caption, such as the assumption of S\u2009ii column density, etc. The synthetic spectra were generated using the fitlyman routine (Fontana & Ballester 1995) implemented in the eso-midas software package. Atomic data were taken from Morton (2003). To show the effect clearly, we zoom in the spectrum to the radial velocity range [ \u2212 10km \u00b7 s\u22121, 10km \u00b7 s\u22121]. The enhancement and reduction of the spectral line profile due to the magnetic realignment change among the triplets. Given the fact that meanwhile optical QSO spectra reach up to a new standard of S\/N of a few hundred (e.g. D\u2019Odorico et al. 2016), the predicted effect is already VISIBLE, if the component structure of the DLA allows a detailed investigation.","Citation Text":["Welsh & Lallement 2012"],"Functions Text":["The three transitions of singly ionized Sulfur (Sii; upper ionization potential 23.3 eV) at \u03bb\u03bb1250.58, 1253.81, 1259.52\u00c5 represent important tracers for neutral and weakly ionized gas in the local interstellar and intergalactic medium and in distant galaxies (e.g."],"Functions Label":["Uses"],"Citation Start End":[[815,837]],"Functions Start End":[[529,793]]}
{"Identifier":"2018AandA...616A..99K__Narang_et_al._2016_Instance_4","Paragraph":"The high-resolution imaging observations of TR from IRIS reveal the ubiquitous presence of network jets. We have used three different IRIS observations of the quiet sun, which are located near the disk center. On the basis of careful inspection, 51 network jets are identified from three QS observations and used for further analysis. These 51 network jets are very well resolved and are not affected by the dynamics of other jets. The study is focused on the rotating motion of network jets along with the estimation of their other properties (speed, height, and lifetime). The mean speed, as predicted by statistical distributions of the speed, is 140.16 km s\u22121 with a standard deviation of 39.41 km s\u22121. The mean speed of network jets is very similar, as reported in previous works (e.g., Tian et al. 2014; Narang et al. 2016). However, in case of their lifetimes, we found a value that is almost double (105.49 s) that of the previously reported mean lifetime of the network jets (49.6 s; Tian et al. 2014). As mentioned above, we took only those network jets that are very well resolved in space and in the time; these criteria exclude short lifetime network jets. Therefore, our statistical distribution of the lifetime predicts a higher mean lifetime. The mean length of the network jets is 3.16 Mm with a standard deviation of 1.18 Mm. In the case of CH network jets, Tian et al. (2014) have reported that most of the network jets have lengths from 4.0 to 10.0 Mm. However, the mean length for QS network jets is smaller (3.53 Mm; Narang et al. 2016). So, the mean length for QS network jets from the present work is in good agreement with Narang et al. (2016). In addition, the apparent speed and length of these network jets are positively correlated, which is very similar to what has already been reported in previous works (Narang et al. 2016). Finally, we can say that these networks jets are very dynamic features of the solar TR, as revealed by their estimated properties.","Citation Text":["Narang et al. 2016"],"Functions Text":["In addition, the apparent speed and length of these network jets are positively correlated, which is very similar to what has already been reported in previous works"],"Functions Label":["Similarities"],"Citation Start End":[[1837,1855]],"Functions Start End":[[1670,1835]]}
{"Identifier":"2022AandA...666A..80N__Piatti_2022_Instance_1","Paragraph":"Many studies indicate a much more complicated star formation history (SFH; e.g., Pagel & Tautvai\u0161vien\u0117 1998; Carrera et al. 2008; Harris & Zaritsky 2009; Rubele et al. 2012; Meschin et al. 2014; Palma et al. 2015; Perren et al. 2017) or AMR for star clusters (e.g., Olszewski et al. 1991; Hill et al. 2000; Dirsch et al. 2000; Rich et al. 2001; Leonardi & Rose 2003; Kerber et al. 2007) in the LMC than in the SMC. Early studies of the cluster AMR in the LMC (e.g., the spectroscopic study of Olszewski et al. 1991) revealed a mysterious gap in the star cluster formation between 10 and 3 Gyr ago (with the sole exception of cluster ESO121-3 and very recently also KMHK1592, as reported by Piatti 2022), which follows the initial formation of old, metal-poor globular clusters, and precedes more recent active formation of intermediate-age clusters. This so-called \u201cage gap\u201d, which is simultaneously a metallicity gap, was later reported also by other authors (e.g., Harris & Zaritsky 2009; Sharma et al. 2010; Kerber et al. 2007, and references therein), even though it has not been fully confirmed when only the field stars are considered. For example, Piatti & Geisler (2013) reported that no clear age gap in the field star formation is observed in their study of fields in the LMC main body, based on Washington photometry. Furthermore, based on their spectroscopic analysis of four fields to the north of the LMC bar, Carrera et al. (2008) maintain that the disk and cluster AMR are similar. However, unlike clusters, there is no age gap in the field population. On the contrary, Harris & Zaritsky (2009) claim that such a gap is evident in the field population of the bar, mostly omitted in previous studies. Also Meschin et al. (2014), using optical photometry, identified two main star-forming epochs with a period characterized by a lower activity in between; however, in their work that interval of time was shorter and lasted from \u223c8 up to \u223c4 Gyr ago. All the aforementioned authors agree that star formation in the LMC is continuing to this day.","Citation Text":["Piatti 2022"],"Functions Text":["with the sole exception of cluster ESO121-3 and very recently also KMHK1592, as reported by"],"Functions Label":["Differences"],"Citation Start End":[[690,701]],"Functions Start End":[[598,689]]}
{"Identifier":"2016AandA...591A..13V__Clarke_2004_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":["Clarke 2004"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[615,626]],"Functions Start End":[[284,526]]}
{"Identifier":"2019AandA...622A..62A__Aviles_et_al._(2018)_Instance_1","Paragraph":"On the other hand, PT has experienced many developments in recent years (Matsubara 2008a; Baumann et al. 2012; Carlson et al. 2013), in part because it can be useful to analytically understand different effects in the power spectrum and correlation function for the dark matter clustering. These effects can be confirmed or refuted, and further explored with simulations to ultimately understand the outcomes of present and future galaxy surveys, such as eBOSS (Zhao 2016), DESI (Aghamousa et al. 2016), EUCLID (Amendola et al. 2013), and LSST (LSST Dark Energy Science Collaboration 2012). Two approaches have been used to study PT: the Eulerian standard PT (SPT) and Lagrangian PT (LPT), which both have advantages and drawbacks, but they are complementary in the end (Tassev 2014). The nonlinear PT for MG was developed initially in (Koyama et al. 2009), and has been further studied in several other works (Taruya et al. 2014a,b; Brax & Valageas 2013; Bellini & Zumalacarregui 2015; Taruya 2016; Bose & Koyama 2016, 2017; Barrow & Mota 2003; Akrami et al. 2013; Fasiello & Vlah 2017; Aviles & Cervantes-Cota 2017; Hirano et al. 2018; Bose et al. 2018; Bose & Taruya 2018; Aviles et al. 2018). The LPT for dark matter fluctuations in MG was developed in Aviles & Cervantes-Cota (2017), and further studies for biased tracers in Aviles et al. (2018). The PT for MG has the advantage that it allows us to understand the role that these physical parameters play in the screening features of dark matter statistics. We here study some of these effects through screening mechanisms by examining them at second- and third-order perturbation levels using PT for some MG models. To this end, we build on the LPT formalism developed in Aviles & Cervantes-Cota (2017), which was initially posited for MG theories in the Jordan frame, in order to apply it to theories in the Einstein frame. Because of the direct coupling of the scalar field and the dark matter in the Klein\u2013Gordon equation, the equations that govern the screening can differ substantially from those in Jordan-frame MG theories. In general, screening effects depend on the type of nonlinearities that are introduced in the Lagrangian density. We present a detailed analysis of screening features and identify the theoretical roots of their origin. Our results show that screenings possess peculiar features that depend on the scalar field effective mass and couplings, and that may in particular cases cause anti-screening effects in the power spectrum, such as in the symmetron model. We perform this analysis by separating the growth functions into screening and non-screened parts. We note, however, that we do not compare the perturbative approach with a fully nonlinear simulation. We refer to (Koyama et al. 2009), for instance, for investigations like this at the level of the power spectrum.","Citation Text":["Aviles et al. 2018"],"Functions Text":["The nonlinear PT for MG was developed initially in","and has been further studied in several other works"],"Functions Label":["Uses","Extends"],"Citation Start End":[[1176,1194]],"Functions Start End":[[785,835],[858,909]]}
{"Identifier":"2015MNRAS.450...45G__Klypin_et_al._1999_Instance_1","Paragraph":"The \u039b cold dark matter (CDM) model has had tremendous success in providing a cosmogony that links the state of the very high redshift Universe as inferred from cosmic microwave background observations and big bang nucleosynthesis considerations, with the observed large-scale distribution of galaxies and its evolution at later times. In this paradigm, the observed structures grew due to gravitational amplification of initially small perturbations, perhaps seeded by inflation (Guth 1981), with most of the gravitating mass in the form of \u2018dark matter\u2019, an as of yet unknown type of matter that apart from gravitationally, interacts only very weakly if at all with baryonic matter and radiation, see e.g. Frenk & White (2012) for a recent review. Observations on smaller, sub galactic scales, have proven more problematic for \u039bCDM, with suggestions that the low abundance (e.g. Klypin et al. 1999; Moore et al. 1999) and shallow density profiles (e.g Gilmore et al. 2007; de Vega & Sanchez 2010) inferred for haloes hosting Milky Way satellites are inconsistent with the numerous substructures with cuspy density profiles that form in Milky Way-like CDM haloes. Indeed the substructures in the haloes of the Aquarius project (Springel et al. 2008) and the GHALO project (Stadel et al. 2009) may be difficult to reconcile with those inferred to host Milky Way satellites. Even if many of the smaller DM substructures may not have a sufficiently deep potential well to form stars after reionization (e.g Efstathiou 1992; Benson et al. 2002) and so may well be dark, some of the more massive ones are probably too big to be affected (Okamoto, Gao & Theuns 2008) and would hence appear to be \u2018too big to fail\u2019 (Boylan-Kolchin, Bullock & Kaplinghat 2011, 2012). However, the jury on this is still out: the Milky Way's halo may simply be of lower mass than those of the Aquarius haloes and hence have fewer massive satellites (Wang et al. 2012). The status of the density profiles of the satellite's haloes \u2013 cored versus cuspy \u2013 is similarly unresolved. Strigari, Frenk & White (2010) claim that the stellar dynamics observations of the satellites do not imply cores at all and hence may be consistent with CDM cusps. But even if the satellites had cored profile, this might result from the action of baryonic feedback processes (Governato et al. 2012; Pontzen & Governato 2012), and hence still be consistent with CDM.","Citation Text":["Klypin et al. 1999"],"Functions Text":["Observations on smaller, sub galactic scales, have proven more problematic for \u039bCDM, with suggestions that the low abundance (e.g.","inferred for haloes hosting Milky Way satellites are inconsistent with the numerous substructures with cuspy density profiles that form in Milky Way-like CDM haloes."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[880,898]],"Functions Start End":[[749,879],[998,1163]]}
{"Identifier":"2015AandA...584A.103S__Lattimer_&_Swesty_1991_Instance_1","Paragraph":"The energy in the inner crust is largely influenced by the properties of the neutron gas and, therefore, the EoS of neutron matter of the different calculations plays an essential role. The NV calculation (Negele & Vautherin 1973) is based on a local energy density functional that closely reproduces the Siemens-Pandharipande EoS of neutron matter (Siemens & Pandharipande 1971) in the low-density regime. The Moskow calculation (Baldo et al. 2007) employs a semi-microscopic energy density functional obtained by combining the phenomenological functional of Fayans et al. (2000) inside the nuclear cluster with a microscopic part calculated in the Brueckner theory with the Argonne v18 potential (Wiringa et al. 1995) to describe the neutron environment in the low-density regime (Baldo et al. 2004). The BBP calculation (Baym et al. 1971a,b) gives the EoS based on the Brueckner calculations for pure neutron matter of Siemens (Siemens & Pandharipande 1971). The LS-Ska (Lattimer & Swesty 1991; Lattimer 2015) and DH-SLy4 (Douchin & Haensel 2001) EoSs were constructed using the Skyrme effective nuclear forces Ska and SLy4, respectively. The SLy4 Skyrme force (Chabanat et al. 1998) was parametrized, among other constraints, to be consistent with the microscopic variational calculation of neutron matter of Wiringa et al. (1988) above the nuclear saturation density. The Shen-TM1 EoS (Shen et al. 1998b,a; Sumiyoshi 2015) was computed using the relativistic mean field parameter set TM1 for the nuclear interaction. The calculations of LS (Lattimer & Swesty 1991; Lattimer 2015) and Shen et al. (Shen et al. 1998b,a; Sumiyoshi 2015) are the two EoS tables in more widespread use for astrophysical simulations. The BSk21 EoS (Pearson et al. 2012; Fantina et al. 2013; Potekhin et al. 2013; Goriely et al. 2010) is based on a Skyrme force with the parameters accurately fitted to the known nuclear masses and constrained, among various physical conditions, to the neutron matter EoS derived within modern many-body approaches which include the contribution of three-body forces. ","Citation Text":["Lattimer & Swesty 1991"],"Functions Text":["The LS-Ska","EoSs were constructed using the Skyrme effective nuclear forces Ska and SLy4, respectively."],"Functions Label":["Background","Background"],"Citation Start End":[[974,996]],"Functions Start End":[[962,972],[1050,1141]]}
{"Identifier":"2018MNRAS.479.3923Z__Yan_&_Lazarian_2006_Instance_1","Paragraph":"As illustrated in this paper, the variation of spectral line intensity induced by GSA varies among different spectral lines. Thus, such influence could be precisely analysed if multispectral lines for the same element is achievable. In addition, the alignment on the ground state is transferred to the levels of atoms on the excited states through absorption process, as illustrated in Section 3 and 4. The intensity of the ultraviolet-pumped fluorescence lines, which are derived from successive decays to different levels of atoms and applied to the modelling of reflection nebulae (Sellgren 1984, 1986), are dependent on the initial upper levels, and thus is influenced by GSA through the scattering process. The influence of collision is neglected in this paper, which applies to most diffuse ISM and IGM. Collisionsreduce the alignment efficiency (see Hawkins 1955). The collision effect can become important in the case of higher density medium where the collision rate $\\tau _c^{-1}$ (either inelastic collision rate or Van der Waals collision rate) dominates over optical pumping rate $B_{lu}\\bar{J}^0_0$ (see Yan & Lazarian 2006 for details). The focus of our paper is on the GSA effect that is a saturated state. When the magnetic precession rate is comparable to the optical pumping from the ground state ($2\\pi \\nu _L\\sim B_{lu}\\bar{J}^0_0$), the ground-level Hanle effect is applicable (see Landolfi & Landi Degl\u2019Innocenti 1986), which means the magnetic field influence on the spectrum is not limited to the change of direction but also the magnetic field strength. As demonstrated in Yan & Lazarian (2008), the effect becomes saturated when $2\\pi \\nu _L\\rightarrow 10B_{lu}\\bar{J}^0_0$. Therefore, we set rHanlewhere $2\\pi \\nu _L\\simeq 10B_{lu}\\bar{J}^0_0$ as the boundary between the ground-level Hanle regime and the GSA regime (see Fig. 12a):\n(10)\r\n\\begin{eqnarray*}\r\nr_{Hanle}=r_\\ast \\sqrt{\\frac{A_{ul}[J_u]}{1.76(B\/{\\mu }G)(\\exp (h\\nu \/(k_BT))-1)[J_l]}}\r\n\\end{eqnarray*}\r\nWe calculate the boundary rHanle for C\u2009ii \u03bb1037\u2009\u00c5 in the presence of stars with different effective temperature Teff and radius r* for the magnetic field with the strength range from \u03bcG to mG in Fig. 12(b). The rHanle increases as the magnetic field becomes weaker and as the effective temperature Teff and radius r* increase. Nevertheless, even in the most optimistic scenario with Teff = 4 \u00d7 104K, r* = 10r\u2299, B = 1\u03bcG (though rarely applies), the boundary rHanle \u2243 180Au = 8.5 \u00d7 10\u22124pc, which is a thousand times smaller than the normal H\u2009ii Region which is in pc scale. All the analysis with GSA and their observational implication in this paper is applicable to most of the ISM, except when performing very high resolution spectral analysis on regions very close to the very bright O-type star.","Citation Text":["Yan & Lazarian 2006"],"Functions Text":["The collision effect can become important in the case of higher density medium where the collision rate $\\tau _c^{-1}$ (either inelastic collision rate or Van der Waals collision rate) dominates over optical pumping rate $B_{lu}\\bar{J}^0_0$ (see"],"Functions Label":["Background"],"Citation Start End":[[1118,1137]],"Functions Start End":[[872,1117]]}
{"Identifier":"2019MNRAS.490.5478W__Facchini_et_al._2016_Instance_1","Paragraph":"The focus of this work is to address the proplyd lifetime problem from a modelling perspective using recent developments in both the theory of photoevaporating discs and observations of the star and disc population in the ONC. The most similar study to this work is that of Scally & Clarke (2001), who produced an N-body model of the ONC coupled with theoretical photoevaporative mass-loss rates. The aim was to test the idea put forward by St\u00f6rzer & Hollenbach (1999) that radial orbits of stars could result in shorter periods of exposure to strong FUV flux close to \u03b81C, thus increasing the PPD dispersal time-scale. The models demonstrated that such dynamical orbits alone are insufficient to produce significantly extended disc lifetimes. However, since that study a number of developments mean that the problem is due to be revisited. First, thermodynamic calculations of conditions in photodissociation regions (PDRs) have been coupled with self-consistent equations for the thermally driven disc wind (Adams et al. 2004; Facchini et al. 2016; Haworth et al. 2018; Haworth & Clarke 2019). In particular, the recent Fried grid (Haworth et al. 2018) can be interpolated across FUV flux, disc outer radius, and disc mass to calculate an instantaneous mass-loss rate for a given PPD. Secondly, the recent sub-mm survey of disc masses and radii by Eisner et al. (2018, hereafter E18) offers further observational constraints on a successful model of PPD evolution, and the host mass-dependent disc properties permit a test of theoretical predictions. Thirdly, Beccari et al. (2017, see also Jerabkova et al. 2019) recently found evidence of three stellar populations with distinct ages in the ONC. This has multiple consequences for models of a PPD population, most obviously that a subset of discs has evolved for a shorter period than the oldest stars in the region. Additionally, since stars are formed from gas, during the evolution of the cluster prior to the most recent formation event there may have been considerable intracluster extinction of UV photons. Understanding how such a scenario influences the disc population would also represent a test of the hypothesis that discrete epochs of star formation occurred in the ONC.","Citation Text":["Facchini et al. 2016"],"Functions Text":["First, thermodynamic calculations of conditions in photodissociation regions (PDRs) have been coupled with self-consistent equations for the thermally driven disc wind"],"Functions Label":["Uses"],"Citation Start End":[[1029,1049]],"Functions Start End":[[841,1008]]}
{"Identifier":"2020ApJ...901..130Z___2020_Instance_1","Paragraph":"Fast radio bursts (FRBs), a promising new and mysterious astrophysical phenomenon, are a class of bright transients with millisecond durations detected at \u223cGHz (Lorimer et al. 2007; Keane et al. 2011; Thornton et al. 2013; Burke-Spolaor & Bannister 2014; Spitler et al. 2014; Masui & Sigurdson 2015; Petroff et al. 2015, 2016; Ravi et al. 2015, 2016; Champion et al. 2016; Keane et al. 2016; Caleb et al. 2017). Although the physical origin of these intense short pulses has not been figured out so far, most of them have been detected with relatively large dispersion measures (DM) (greater than the maximum produced by the Milky Way). This suggests that they may occur at a cosmological distance, probably at distances of the order of gigaparsecs (Ioka 2003; Inoue 2004; Deng & Zhang 2014; Zheng et al. 2014; Zhang 2018). So far, the cosmological origin of these kinds of mysterious flashes has been confirmed by successfully localizing several bursts (including two repeaters and several apparently nonrepeating ones; Spitler et al. 2016; Scholz et al. 2016; Chatterjee et al. 2017; Marcote et al. 2017; Tendulkar et al. 2017; Bannister et al. 2019; Prochaska et al. 2019; Ravi et al. 2019; Wang & Zhang 2019; Zhang & Wang 2019; Marcote et al. 2020). The confirmation of their cosmological origin allows them to be widely proposed as promising tools for studying the universe and fundamental physics, e.g., the distribution of baryons in the intergalactic medium (IGM) or diffuse gas (Deng & Zhang 2014; Mcquinn 2014; Mu\u00f1oz & Loeb 2018; Li et al. 2019, 2020; Walters et al. 2019; Wei et al. 2019; Macquart et al. 2020), dark energy (Gao et al. 2014; Zhou et al. 2014; Walters et al. 2018; Wei et al. 2018; Zhao et al. 2020), cosmic ionization history (Deng & Zhang 2014; Zheng et al. 2014), the large-scale structure of the universe (Masui & Sigurdson 2015), the Einstein\u2019s equivalence principle (Wei et al. 2015; Nusser 2016; Tingay & Kaplan 2016), the rest mass of the photon (Wu et al. 2016; Shao & Zhang 2017; Xing et al. 2019), the cosmic proper distance measurements (Yu & Wang 2017), and constraints on the magnetic fields in the IGM (Akahori et al. 2016), as well as measuring the cosmic expansion rate (Wu et al. 2020). Because the event rate of this phenomenon inferred from observations is very high (\u223c103\u2013104 per day all sky; Thornton et al. 2013; Champion et al. 2016) and due to their extragalactic origin, these bursts are likely to be gravitationally lensed by intervening objects with a different magnitude of mass. Moreover, high time resolutions or short durations of these bursts (\u223c(1\u201310) ms) (Lorimer et al. 2007; Keane et al. 2011; Thornton et al. 2013; Spitler et al. 2014, 2016; Petroff et al. 2015, 2016; Ravi et al. 2015; Champion et al. 2016) guarantee that the difference between arrival times of two images can be resolved even if they are lensed by an object as small as 10  M\u2299. Therefore, different scale lensing effects of FRBs have been proposed as a probe of compact dark matter (Mu\u00f1oz et al. 2016; Wang & Wang 2018; Liao et al. 2020; Ranjan 2020), motion of the FRB source (Dai & Lu 2017), and precision cosmology (Li et al. 2018; Liu et al. 2019; Wucknitz et al. 2020). Recently, taking into consideration the theoretical prediction that a portion of FRBs might be associated with GWs, Wei et al. (2018) proposed the joint measurements of DM from FRB observations and DL from GW detections in the same FRB\/GW association system (i.e., the combination DM \u00b7 DL as a function of redshift) as a complementary cosmic probe. The most striking merit of this probe is the independence of the Hubble constant H0, which is a fundamental cosmological parameter now under intense debate because of the well-known Hubble constant tension (Planck Collaboration et al. 2018; Riess et al. 2019). In this paper, we propose that the combination DM \u00b7 DL as a function of redshift could be further extended to the independent measurements of DMs of localized FRBs and DL of other distance indicators (e.g., SNe Ia) at a similar redshift according to the fundamental assumption that luminosity distance monotonously increases with increasing redshift. Following Masui & Sigurdson (2015), who call brief broadband radio impulses like FRBs \u201cstandard pings,\u201d the extended combination DM \u00b7 DL as a function of redshift can be easily constructed from independent measurements of standard pings and standard candles at similar redshifts (Wei et al. 2018). Moreover, in the CPL framework, we investigate the constraining power on the equation of state of dark energy from this extended combination.","Citation Text":["Li et al.","2020"],"Functions Text":["The confirmation of their cosmological origin allows them to be widely proposed as promising tools for studying the universe and fundamental physics, e.g., the distribution of baryons in the intergalactic medium (IGM) or diffuse gas"],"Functions Label":["Motivation"],"Citation Start End":[[1540,1549],[1556,1560]],"Functions Start End":[[1254,1486]]}
{"Identifier":"2015ApJ...805...88Y__Metzger_et_al._2011_Instance_1","Paragraph":"GRB 100814A is another special event with an early-time shallow decay phase and late-time significant rebrightening in its optical afterglow light curve (Nardini et al. 2014). The power-law (\n\n\n\n\n\n) temporal index of the early shallow decay is \n\n\n\n\n\n, which is inconsistent with expectations from the external shock model. It is argued that the shallow decay phases come from energy injection. Usually, the injection luminosity is assumed as \n\n\n\n\n\n (Nousek et al. 2006; Zhang et al. 2006; Yu & Huang 2013), which may naturally come from the magnetic dipole radiation of a newborn millisecond magnetar (Dai & Lu 1998; Zhang & M\u00e9sz\u00e1ros 2001; Dall\u2019Osso et al. 2011). As a result, magnetars have been suggested as the central engines for some GRBs, including both long and short events (Zhang & M\u00e9sz\u00e1ros 2001; Troja et al. 2007; Metzger et al. 2011; Bernardini et al. 2012; Rowlinson et al. 2013). In both Dai & Lu (1998) and Dall\u2019Osso et al. (2011), works considering a strongly magnetized neutron star as the central engine of a GRB, the energy injection power is more realistically derived as \n\n\n\n\n\n, where T is the spin-down timescale and \n\n\n\n\n\n is the initial luminosity. In particular, while considering the exact form for the energy injection power of a spinning down magnetar due to magnetic dipole radiation, Dall\u2019Osso et al. (2011) found that the luminosity of the X-ray afterglow naturally has a shallow decay phase with a temporal power-law index of \n\n\n\n\n\n. Recently, a nearly constant dipole radiation luminosity (\n\n\n\n\n\n) during the spin-down timescale was favored by observations from GRBs, such as 050801 (de Pasquale et al. 2007), 060729 (Grupe et al. 2007), and 080913 (Greiner et al. 2009). However, some observations of GRB afterglows with rebrightenings or bumps (i.e., \n\n\n\n\n\n) require that the injection luminosity increases with time (i.e., \n\n\n\n\n\n). Additionally, there is a plateau phase in the X-ray band of GRB 100814A between 103 and 105 s (Nardini et al. 2014) indicating continuous energy injection from the central engine during this prolonged period.","Citation Text":["Metzger et al. 2011"],"Functions Text":["As a result, magnetars have been suggested as the central engines for some GRBs, including both long and short events"],"Functions Label":["Background"],"Citation Start End":[[825,844]],"Functions Start End":[[664,781]]}
{"Identifier":"2017AandA...607A.120P__Warnecke_et_al._(2016)_Instance_1","Paragraph":"Outside the TC, the MHD simulation behaves similarly to its HD counterpart, with baroclinicity being the dominant contribution. In terms of balance between left- and right-hand sides of (13), panels A and D show a very good agreement in most of the CZ. Differences are only evident inside the TC in the top layers and slightly near the poles. We attribute these differences to the upper and polar boundary conditions for the magnetic field. The difference observed between panels 12A and D in the top boundary might be due to the radial magnetic field we enforce at the surface. This means that when we have a strong poloidal field located near the surface, it will be forced over a couple grid points by the boundary condition into the radial direction. This is exactly the case of the poloidal field configuration during cycle maximum (see Fig. 3E). This problem might be alleviated in future simulations by introducing more realistic top boundary conditions like that used in Warnecke et al. (2016). The differences between the left- and right-hand sides during cycle minimum (not shown here) are much smaller. It is the radial derivative of the poloidal field present in the second magnetic term of Eq. (13) that is responsible for this \u201cartificial\u201d contribution. There are two other possible sources of error that can explain the minute differences we find in this balance calculation (and in the previous section as well). The first is numerical diffusivity, which we cannot measure directly. The other issue is related with the different numerical methods used to compute derivatives and other composite quantities in the main code during the simulation and a posteriori. During the simulation run, EULAG numerics computes central cell values and fluxes across the cell borders, while the type of analysis that we perform a posteriori assumes values computed in the cell corners using centered finite differences. Differentiation across the poles can also introduce some artifacts. Nevertheless, the very good match obtained in the HD case indicates that these two sources of error are in fact very small, and that the main issue here seems related to the magnetic field boundary conditions. Despite these possible sources of uncertainty, we consider that there is a general good agreement between left- and right-hand sides for most of the CZ. ","Citation Text":["Warnecke et al. (2016)"],"Functions Text":["This problem might be alleviated in future simulations by introducing more realistic top boundary conditions like that used in"],"Functions Label":["Future Work"],"Citation Start End":[[979,1001]],"Functions Start End":[[852,978]]}
{"Identifier":"2019ApJ...871L..22W__Alexandrova_2008_Instance_1","Paragraph":"In analogy to the hydrodynamic case, the nonlinear coherent vortex structure also plays an important role in plasma dynamics and transport processes (Hasegawa & Mima 1978; Shukla et al. 1985; Petviashvili & Pokhotelov 1992; Horton & Hasegawa 1994). These vortices tend to have a long lifetime and are widely observed in space, laboratory, and numerical simulation of plasma (Chmyrev et al. 1988; Burlaga 1990; Volwerk et al. 1996; Stasiewicz et al. 2000; Sundkvist et al. 2005; Alexandrova et al. 2006; Alexandrova 2008; Alexandrova & Saur 2008; Vianello et al. 2010; Servidio et al. 2015). An essential subset of these plasma vortices is known as Alfv\u00e9n vortices, which can be viewed as the cylindrical analog of the nonlinear Alfv\u00e9n wave (Petviashvili & Pokhotelov 1992). The Alfv\u00e9n vortices have an axis that is nearly parallel to the unperturbed magnetic field, along which the shape is generally invariant. Thus, these vortices are quasi-2D structures. The associated perpendicular magnetic fluctuations are linearly related with the perpendicular velocity fluctuations, but their relative amplitudes are not obligatorily equal (as is the case in an Alfv\u00e9n wave): \n\n\n\n\n\n, where \u03be is not necessarily equal to 1. In addition, Alfv\u00e9n vortices do not propagate along \n\n\n\n\n\n in the plasma frame, and they hardly propagate in the perpendicular plane when the axis of the vortex is inclined with respect to \n\n\n\n\n\n that are in contrast with Alfv\u00e9n wave (Wang et al. 2012). After first being reported in the Earth\u2019s magnetosheath (Alexandrova et al. 2006; Alexandrova 2008), multiscale quasi-bidimensional Alfv\u00e9n vortices (with \n\n\n\n\n\n) have been identified in numerous space environments: in slow solar wind (Perrone et al. 2016; Roberts et al. 2016), in fast solar wind (Lion et al. 2016; Perrone et al. 2017), and in Saturn\u2019s magnetosheath (Alexandrova & Saur 2008). It seems that the intermittent structures in fast solar wind are dominated by Alfv\u00e9n vortices (Perrone et al. 2017), which agrees with the 2D MHD turbulence model (Zank et al. 2017).","Citation Text":["Alexandrova 2008"],"Functions Text":["These vortices tend to have a long lifetime and are widely observed in space, laboratory, and numerical simulation of plasma"],"Functions Label":["Background"],"Citation Start End":[[503,519]],"Functions Start End":[[249,373]]}
{"Identifier":"2016ApJ...827...47L__Dmitruk_et_al._2004_Instance_1","Paragraph":"Given the time-consuming nature of simulating energetic particle acceleration in multiple contracting and reconnecting (merging) small-scale magnetic islands produced by magnetohydrodynamic (MHD) turbulence and particle-in-the-cell (PIC) models on large spatial scales and in three dimensions, there has been an increased effort to develop statistical kinetic transport theories that capture the essential physics of particle acceleration in multi-island regions (de Gouveia dal Pino & Lazarian 2005; Drake et al. 2006, 2010, 2013; Bian & Kontar 2013\n). Recently, Zank et al. (2014) developed a comprehensive kinetic transport theory for small-scale flux ropes (helical magnetic field structures) which, for the first time, unified three different mechanisms currently thought to play an important role in the energization of suprathermal particles traversing solar wind regions containing numerous contracting and reconnecting (merging) small-scale flux ropes. It was argued that magnetic flux ropes in the solar wind can be viewed to first order as being quasi-two-dimensional (quasi-2D) helical contracting and merging flux-rope structures based on three-dimensional (3D) simulations of compressible MHD turbulence that included a strong guide magnetic field (Dmitruk et al. 2004). The three acceleration mechanisms include (i) curvature and grad-B drift acceleration in the motional electric fields generated by the plasma flow at the endpoints of contracting magnetic flux ropes (Drake et al. 2006), (ii) curvature and grad-B drift acceleration in the motional electric fields induced by x-point plasma outflows generated near the center of merging neighboring magnetic islands (Drake et al. 2013), and (iii) acceleration due to field-aligned guiding center motion in the secondary reconnection electric field formed in the x-point region at the center of reconnecting islands (Oka et al. 2010). Interestingly, in Zank et al. (2014) contracting flux ropes were modeled to behave in a compressible manner so that betatron acceleration by the increasing flux-rope field strength also contributes to particle energization. Both Drake et al. (2013) and Zank et al. (2014) consider reconnecting (merging) magnetic islands\/flux ropes as incompressible phenomena so that negative betatron acceleration from the decreasing flux-rope field strength reduces particle energization. Zank et al. (2014) found that the mean electric field induced by contracting and merging flux ropes naturally accelerates suprathermal test particles to hard power-law distributions in the supersonic, slow solar wind. The results show that the power-law index depends on the Alfv\u00e9n Mach number and on the ratio of the diffusion timescale to that for island contraction.","Citation Text":["Dmitruk et al. 2004"],"Functions Text":["It was argued that magnetic flux ropes in the solar wind can be viewed to first order as being quasi-two-dimensional (quasi-2D) helical contracting and merging flux-rope structures based on three-dimensional (3D) simulations of compressible MHD turbulence that included a strong guide magnetic field"],"Functions Label":["Background"],"Citation Start End":[[1263,1282]],"Functions Start End":[[962,1261]]}
{"Identifier":"2018ApJ...869...47H__Hoppe_et_al._1994_Instance_1","Paragraph":"As for Mg, one would expect to find a positive correlation between \u03b429Si and initial 18O\/16O ratios for the parent stars of category A grains if these quantities represent GCE. Indeed, there is a correlation between \u03b429Si and initial 18O\/16O, as can be seen from Figure 10. The observed trend is a bit shallower than predicted by the GCE model of Timmes & Clayton (1996), which was inferred and extended to higher metallicities from predicted 18O\/16O ratios at [Fe\/H] = \u22120.3 (18O\/16O = 0.56 \u00d7 solar) and [Fe\/H] = 0 (18O\/16O = 1.26 \u00d7 solar); corresponding \u03b429Si values are \u2212500\u2030 ([Fe\/H] = \u22120.3) and +260\u2030 ([Fe\/H] = 0), respectively. It is interesting to note that astronomical observations of a variety of sources have not detected a variation in Si isotope ratios with galactocentric radius, which stands in contrast to the monotonically decreasing trend for the 18O\/16O ratio (Monson et al. 2017) and our observations for category A grains. Following the approach of Nguyen et al. (2010), we have inferred a GCE relationship between \u03b429Si and initial 18O\/16O, which is based on a tight correlation between 29Si\/28Si and 46Ti\/48Ti ratios in SiC mainstream grains (Figure 11, left; data from Hoppe et al. 1994; Alexander & Nittler 1999; Huss & Smith 2007) and a correlation between 46Ti\/48Ti and initial 18O\/16O ratios in presolar oxide grains (Figure 11, right; data from Choi et al. 1998; Hoppe et al. 2003), interpreted to represent largely GCE. Error-weighted linear regressions yield \u03b446Ti = (0.98 \u00b1 0.08) \u00d7 \u03b429Si-(11 \u00b1 7) for SiC mainstream grains and \n\n\n\n\n\n for Group 1 oxides. Combining the two equations gives \n\n\n\n\n\n, which is labeled as \u201cSiC-Oxides GCE\u201d in Figure 10. Note that this equation is different from the one presented by Nguyen et al. (2010), which was recognized to be in error (L. Nittler 2018, private communication). Gyngard et al. (2018) recently reported a slightly shallower slope for the \u03b446Ti versus \u03b429Si of SiC mainstream grains than inferred here, which would make the \u03b446Ti versus \n\n\n\n\n\n relationship a bit steeper. Our data for category A grains fall to the right of the SiC-Oxides GCE line but are fully compatible with this line if errors for the SiC-Oxides GCE line are taken into account (Figure 10). Finally, the good correlation between \u03b425Mg and \u03b429Si (Figure 12) rounds out the picture emerging from presolar silicate category A grains, suggesting that a subset of presolar Group 1 silicates carries resolvable signatures of the GCE of O, Mg, and Si isotopes.","Citation Text":["Hoppe et al. 1994"],"Functions Text":["Following the approach of Nguyen et al. (2010), we have inferred a GCE relationship between \u03b429Si and initial 18O\/16O, which is based on a tight correlation between 29Si\/28Si and 46Ti\/48Ti ratios in SiC mainstream grains (Figure 11, left; data from"],"Functions Label":["Uses"],"Citation Start End":[[1191,1208]],"Functions Start End":[[942,1190]]}
{"Identifier":"2021MNRAS.503.1319G__Chae_&_Mao_2003_Instance_1","Paragraph":"In the first scenario, we assume that neither the characteristic velocity dispersion (\u03c3*) nor the number density (n*) of galaxies evolves with redshifts (\u03bdn = \u03bdv = 0). Given the redshift coverage of the lensing galaxies in the lens sample (0.06  zl  1.0), if we constrain a non-evolving VDF using the lens data, then, assuming the VDF evolution with redshift is smooth, the fits on the VDF parameters may represent the properties of ETGs at an effective epoch of z \u223c 0.5. Such non-evolving VDF has been extensively applied in the previous studies on lensing statistics (Chae & Mao 2003; Ofek et al. 2003; Capelo & Natarajan 2007; Cao et al. 2012a). By applying the above-mentioned \u03c72 \u2013 minimization procedure to Sample A \u2013 we obtain the best-fitting values and corresponding 1\u03c3 uncertainties (68.3 per cent confidence level): $\\alpha =0.66^{+2.13}_{-0.66}$, $\\beta =2.28^{+0.24}_{-0.18}$. It is obvious that the full sample analysis has yielded improved constraints on the high-velocity exponential cut-off index \u03b2, compared with the previous analysis of using the distribution of image separations observed in CLASS and PANELS to constrain a model VDF of ETGs (Chae 2005). Suffering from the limited size of lens sample, such analysis (Chae 2005) found that neither of the two VDF parameters (\u03b1, \u03b2) can be tightly constrained, due to the broad regions in the \u03b1 \u2212 \u03b2 plane. Consequently, the image separation distribution is consistent with the SDSS measured stellar VDF (Sheth et al. 2003) and the Second Southern Sky Redshift Survey (SSRS2) inferred stellar VDF (Chae & Mao 2003), although the two stellar VDFs are significantly different from each other concerning their corresponding parameter values. We also consider constraints obtained for the Sample B (defined in previous section), with the likelihood is maximized at $\\alpha =1.00^{+2.38}_{-1.00}$ and $\\beta =2.34^{+0.26}_{-0.24}$, from which one could clearly see the marginal consistency between our fits and recent measurements of three stellar VDFs (especially the SDSS DR5 VDF of ETGs).","Citation Text":["Chae & Mao 2003"],"Functions Text":["Such non-evolving VDF has been extensively applied in the previous studies on lensing statistics"],"Functions Label":["Background"],"Citation Start End":[[570,585]],"Functions Start End":[[472,568]]}
{"Identifier":"2022ApJ...934...63O__Werf_et_al._1988_Instance_1","Paragraph":"Cosmic rays (CRs) play a vital role in the chemistry of cold (10\u201330 K), dense (>102 cm\u22123) molecular clouds as they can pierce deep into them, unlike interstellar UV radiation (for a review see Indriolo & McCall 2013). These high-energy interstellar particles primarily consisting of protons can be heavier elements and electrons, and have large energy ranges, up to zetaelectronvolt energies (Blandford et al. 2014). Although the energies can be high, it is the lower energy CRs (\u22641 TeV) that affect the dense interiors (Viti et al. 2013; Padovani et al. 2020). In these regions, CRs have a wide variety of effects, one of the most important is being a producer of atomic hydrogen through the dissociation of H2 (van der Werf et al. 1988; Montgomery et al. 1995; Li & Goldsmith 2003; Goldsmith & Li 2005; Padovani et al. 2018a). Other important effects are being the dominant source of ionization; regulating the degree of coupling of the gas and the magnetic field; having an important role in the dynamics and the collapse timescale of collapsing clouds (e.g., Padovani et al. 2013, 2014); providing heating and energy to dust grains (de Jong & Kamijo 1973; Shingledecker et al. 2018; Kalv\u0101ns & Kalnin 2019; Sipil\u00e4 et al. 2020, 2021; Silsbee et al. 2021) producing internal UV photons (Prasad & Tarafdar 1983); may have a role on the charge distribution on dust grains (Ivlev et al. 2015); influencing disk growth (Kuffmeier et al. 2020); and affecting deuteration (Caselli et al. 2008). For example, each species ionized by a CR releases an electron. This secondary electron can cause further collisions, which in turn, depending on the energy, can induce more ionization and heating (Ivlev et al. 2021). If a secondary electron does not have enough energy to ionize a species, the species may become excited (Shingledecker & Herbst 2018). Excited species produced by CR bombardment have energy levels higher than their base counterparts, allowing these excited species to overcome some reaction barriers that would otherwise be difficult in cold environments. These species have been shown to drive more complex chemistry from reactions that can form interstellar complex organic molecules (Abplanalp et al. 2016).","Citation Text":["van der Werf et al. 1988"],"Functions Text":["n these regions, CRs have a wide variety of effects, one of the most important is being a producer of atomic hydrogen through the dissociation of H2"],"Functions Label":["Background"],"Citation Start End":[[713,737]],"Functions Start End":[[563,711]]}
{"Identifier":"2021MNRAS.502.5935S__Hani_et_al._2018_Instance_1","Paragraph":"We now turn to constraining c1. First, note that $\\mathcal {Z} \\gt 0$ for all x. In practice, we ask that $\\mathcal {Z} \\gt \\mathcal {Z}_{\\rm {min}}$ for some fiducial $\\mathcal {Z}_{\\rm {min}} \\approx 0.01$. For x \u226b 1, this gives\n(42)$$\\begin{eqnarray*}\r\nc_1 \\gt \\left(\\mathcal {Z}_{\\mathrm{min}}-\\frac{\\mathcal {S}}{\\mathcal {A}}\\right)\\, x^{-\\frac{1}{2}\\left[\\sqrt{\\mathcal {P}^2+\\, 4\\mathcal {A}} - \\mathcal {P}\\right]}_{\\mathrm{max}},\r\n\\end{eqnarray*}$$where xmax is the outer radius of the disc at which we apply this condition.3 Secondly, the total metal flux into the disc across the outer boundary cannot exceed that supplied by advection of gas with metallicity $\\mathcal {Z}_{\\rm CGM}$ into the disc, since otherwise this would imply the presence of a metal reservoir external to the disc that is supplying metals to it, which is true only in special circumstances, e.g. during or after a merger (Torrey et al. 2012; Hani et al. 2018), or due to long-term wind recycling through strong galactic fountains (Grand et al. 2019). Mathematically, this condition can be written as\n(43)$$\\begin{eqnarray*}\r\n-\\underbrace{\\frac{\\dot{M} \\mathcal {Z}}{2\\pi x}}_\\text{adv. flux} - \\underbrace{\\kappa \\Sigma _g\\frac{\\partial \\mathcal {Z}}{\\partial x}}_\\text{diff. flux} \\ge -\\underbrace{\\frac{\\dot{M} \\mathcal {Z}_{\\rm {CGM}}}{2\\pi x}}_\\text{CGM flux}.\r\n\\end{eqnarray*}$$For x \u226b 1, this translates to,\n(44)$$\\begin{eqnarray*}\r\nc_1 \\le \\frac{2\\mathcal {P}\\left(\\mathcal {Z}_{\\mathrm{CGM}} - \\mathcal {S}\/\\mathcal {A}\\right)}{\\mathcal {P}+\\sqrt{\\mathcal {P}^2+\\, 4\\mathcal {A}}}\\, x^{-\\frac{1}{2}\\left[\\sqrt{\\mathcal {P}^2+\\, 4\\mathcal {A}} - \\mathcal {P}\\right]}_{\\mathrm{max}}.\r\n\\end{eqnarray*}$$Thus, we find that c1 is bounded within a range dictated by the two conditions above. Given a value of c1, we can also calculate the \u03a3g-weighted and $\\dot{\\Sigma }_{\\star }$-weighted mean metallicity in the model,\n(45)$$\\begin{eqnarray*}\r\n\\overline{\\mathcal {Z}}_{\\Sigma _\\mathrm{ g}} = \\frac{\\int ^{x_\\mathrm{max}}_{x_{\\mathrm{min}}} 2\\pi x \\Sigma _{\\mathrm{ g}0} s_\\mathrm{ g} \\mathcal {Z} dx}{\\int ^{x_\\mathrm{max}}_{x_{\\mathrm{min}}} 2\\pi x \\Sigma _{\\mathrm{ g}0} s_\\mathrm{ g} dx},\r\n\\end{eqnarray*}$$\n (46)$$\\begin{eqnarray*}\r\n\\overline{\\mathcal {Z}}_{\\dot{\\Sigma }_{\\star }} = \\frac{\\int ^{x_\\mathrm{max}}_{x_{\\mathrm{min}}} 2\\pi x \\dot{\\Sigma }_{\\star 0} \\dot{s}_{\\star } \\mathcal {Z} dx}{\\int ^{x_\\mathrm{max}}_{x_{\\mathrm{min}}} 2\\pi x \\dot{\\Sigma }_{\\star 0} \\dot{s}_{\\star } dx}.\r\n\\end{eqnarray*}$$Finding $\\overline{\\mathcal {Z}}$ is helpful because we can use it to produce a mass-metallicity relation (MZR) that can serve as a sanity check for the model. We show in a companion paper that our model can indeed reproduce the MZR (Sharda et al. 2020b).","Citation Text":["Hani et al. 2018"],"Functions Text":["Secondly, the total metal flux into the disc across the outer boundary cannot exceed that supplied by advection of gas with metallicity $\\mathcal {Z}_{\\rm CGM}$ into the disc, since otherwise this would imply the presence of a metal reservoir external to the disc that is supplying metals to it, which is true only in special circumstances, e.g. during or after a merger"],"Functions Label":["Uses"],"Citation Start End":[[928,944]],"Functions Start End":[[536,906]]}
{"Identifier":"2020AandA...644L...7G__Magdis_et_al._2020_Instance_2","Paragraph":"As in G18, we compiled existing constraints on the molecular gas fraction fgas of quiescent and pSB galaxies from recent literature, namely: local QGs consisting of the ATLAS3D (Young et al. 2011; Cappellari et al. 2013; Davis et al. 2014) and HRS (Boselli et al. 2014; Lianou et al. 2016) ETG samples as well as the samples of pSB galaxies (hereafter, the \u201clow-z pSB\u201d sample) of French et al. (2015) and Alatalo et al. (2016); at low and intermediate redshift, the ETG sample of Spilker et al. (2018) and the pSB sample of Suess et al. (2017); at intermediate and high redshift, constraints from Hayashi et al. (2018) on gas in z\u2004\u223c\u20041.46 cluster ETGs, as well as on individual galaxies from Sargent et al. (2015), Bezanson et al. (2019), and Rudnick et al. (2017). Given its size, we divided the ATLAS3D sample into high- and low-mass subsamples, choosing 5\u2005\u00d7\u20051010 M\u2299 as the cut-off mass. In addition, we also included fgas estimates derived from the (median) stacked FIR spectral energy distributions of ETGs at z\u2004\u223c\u20041.8 (G18; 977 galaxies), z\u2004\u223c\u20041.2, z\u2004\u223c\u20040.8, and z\u2004\u223c\u20040.5 (1394, 1536, and 563 galaxies, respectively; Magdis et al. 2020, hereafter M20). Finally, at higher redshift (z\u2004\u223c\u20043), we converted star formation rates (SFR) estimated from spectroscopy (Schreiber et al. 2018a; D\u2019Eugenio et al. 2020) into gas masses assuming the star formation efficiency found by G18. As a consequence of our zmax\u2004=\u20043.5, we did not include higher-redshift quiescent galaxies (Glazebrook et al. 2017; Schreiber et al. 2018b; Tanaka et al. 2019; Valentino et al. 2020) in the analysis and considered z\u2004\u223c\u20043 galaxies as pSB. The dust-based estimates of G18 and M20 (and, by extension, the z\u2004\u223c\u20043 semi-constraints) assume a gas-to-dust ratio (G\/D). It is dependent on metallicity, which is presumed to be solar or higher owing to both the relatively high gas-phase metallicity of MS galaxies at z\u2004\u2272\u20041 (e.g., Mannucci et al. 2010) and the already high stellar metallicities of QGs at z\u2004> \u20041 (Onodera et al. 2015; Estrada-Carpenter et al. 2019). Here we adopted an intermediate value between the solar and supersolar G\/Ds used in M20, and we increased the error bars of these points to include both the solar and supersolar confidence estimates. These various samples, which are summarized with their selection criteria in Table B.1, combine into a nonhomogeneous dataset: some were specifically selected as ETGs, and others were based on varying degrees of quiescence. In particular, pSB galaxies are not necessarily truly quiescent and could, in principle, resume normal star formation. However, as a possible precursor of QGs, they provide useful, though not constraining (see Sect. 4), comparison samples for the model. Here we refer to all equally as either QGs or pSB galaxies, and we make the assumption that, on average, these different samples are not otherwise significantly biased with regard to their gas content compared to the full population, given each mass limit and type.","Citation Text":["M20"],"Functions Text":["The dust-based estimates of G18 and","(and, by extension, the z\u2004\u223c\u20043 semi-constraints) assume a gas-to-dust ratio (G\/D)."],"Functions Label":["Background","Background"],"Citation Start End":[[1647,1650]],"Functions Start End":[[1611,1646],[1651,1732]]}
{"Identifier":"2015ApJ...814....2M__Davies_&_Taylor_1950_Instance_1","Paragraph":"The 3D view provided in Figure 1 can be completed through the time evolution of the cross section of the same tube presented in Figure 2 for times t = 13, 35, 48 and 62. In the figure, magnetic field strength (top row) and vorticity (bottom row) maps are drawn on a vertical cut that coincides with the midplane of the box perpendicular to the initial axis of the tube. We confirm that most of the magnetic flux is concentrated at the very top of the structure. The general appearance and evolution of the structures shown are strongly reminiscent of those described by Moreno-Insertis & Emonet (1996), Emonet & Moreno-Insertis (1998) and Cheung et al. (2006) using 2D simulations. We recall a few salient features obtained in those papers and reproduced here: in the initial stages (two leftmost panels), the rising portion of the flux tube has a mushroom shape, akin to an air bubble rising in water (e.g., Davies & Taylor 1950; Collins 1965; Parlange 1969; Wegener & Parlange 1973; Hnat & Buckmaster 1976; Ryskin & Leal 1984; Christov & Volkov 1985). Within the ascending portion of the tube, vorticity is generated at the boundary layer between the tube and the ambient flow, in fact with opposite sign on the either side of the head (bottom row in the figure). Magnetic flux is dragged along the sides of the tube into a trailing pair of counter-rotating vortices. Later in time the wake is fragmented into smaller and smaller vortices because the Reynolds number increases in time due to the expansion of the tube (Cheung et al. 2006). The AMR device allows to see this fragmentation in much more detail than could be obtained in the original work of Emonet & Moreno-Insertis (1998). Inside the head of the tube, on the other hand, the plasma and magnetic field are executing a twisting oscillation, with small velocity compared to the global rise speed. This description of the time evolution of the flux tube structure generally applies to other simulations with comparable levels of refinement even when the twist and curvature vary. The cases with untwisted flux tubes, instead, follow a somewhat different pattern and their evolution is described in the following.","Citation Text":["Davies & Taylor 1950"],"Functions Text":["We recall a few salient features obtained in those papers and reproduced here: in the initial stages (two leftmost panels), the rising portion of the flux tube has a mushroom shape, akin to an air bubble rising in water (e.g.,"],"Functions Label":["Similarities"],"Citation Start End":[[909,929]],"Functions Start End":[[682,908]]}
{"Identifier":"2017ApJ...844..152L__Famaey_et_al._2008_Instance_1","Paragraph":"Figure 3 shows the metallicity distribution of the LS. We divided the whole region into square bins with size 3 km s\u22121 and used color to represent the mean metallicity of each bin. The most prominent feature is the Hyades\u2013Pleiades structure, which includes the Hyades stream (peak 0) and Pleiades stream (peak 1). It exhibits a very metal-rich and consistent distribution; it is well known that the Hyades cluster has a high metallicity ([Fe\/H] \u223c 0.13, Heiter et al. 2014). We speculate that some stars of the Hyades stream might be remnants of the Hyades cluster, while others may be the dynamical effect of non-axisymmetric components (Famaey et al. 2008). Though not in sharp contrast, there is a valley with lower metallicity between the horizontal V = \u221250 km s\u22121 structure and center. To better present the gap, we picked up a window in the metallicity distribution. We divided the window into parallelogram bins, and the edge parallel to the V axis of each bin equals 1 km s\u22121. The right panel of Figure 4 shows that the mean metallicity of each bin distributes along with the V direction. In Figure 3, the Hercules I (peak 6) and Hercules II (peak 4) have different metallicities, and peak 4 is definitely different from peak 12 (NEW2). That the Hercules stream has a wide metallicity range (Table 1) is in accord with the purported effect of Galactic bar resonances (Kalnajs 1991; Dehnen 2000; Fux 2001). The upper left side of this plot contains stars moving faster than the LSR. In this region the kinematics suggest peak 13 and peak 14 may be part of the thin disk, although their metallicities are quite low. The Sirius stream (peak 2) is a little more metal-poor than its surroundings and has a very small metallicity scatter. This is similar to the results of Klement et al. (2008), and this supports the notion that the Sirius stream (or at least a part of it) comprises remnants evaporated from a cluster. According to De Silva et al. (2007), HR1614 (peak 11) may be dispersed remnants of star-forming events. New1 (peak 9) is close to peak 11 in both kinematics and metallicity; part of its member stars may be remnants of star-forming events, too. New2 (peak 12) has a large metallicity scatter. It is possible to be induced by other dynamical n:1 resonance of the bar. The \u03b3 Leo stream (peaks 7, 10) is more metal-poor than our Sun, and it is rapidly moving toward the center of the Galaxy. Considering that there is a radial metallicity gradient in the Galaxy, we conjecture this structure has drifted inwards from beyond the solar Galactocentric radius.","Citation Text":["Famaey et al. 2008"],"Functions Text":["We speculate that some stars of the Hyades stream might be remnants of the Hyades cluster, while others may be the dynamical effect of non-axisymmetric components"],"Functions Label":["Uses"],"Citation Start End":[[638,656]],"Functions Start End":[[474,636]]}
{"Identifier":"2015ApJ...807...26T__Zucker_&_Mazeh_1994_Instance_1","Paragraph":"The numerous historical radial-velocity (RV) measurements of Capella have been discussed at length in our T09 study, which highlighted how challenging it has been to determine accurate values for the rapidly rotating secondary star (\n\n\n\n\n\n km s\u22121), whereas those of the sharp-lined primary (\n\n\n\n\n\n km s\u22121) have been quite consistent over the last century. T09 presented 162 new RV determinations for both components based on spectra obtained at the CfA covering only a very narrow wavelength range (45 \u212b). The RVs were measured using the two-dimensional cross-correlation algorithm TODCOR (Zucker & Mazeh 1994), with synthetic templates appropriate for each star. Because of the limited wavelength coverage, those measurements are susceptible to systematic errors resulting mostly from lines shifting in and out of the spectral window as a function of orbital phase. Therefore, an effort was made to control those biases by performing numerical simulations to determine corrections to the velocities, which were at the level of the final uncertainties in the individual measurements for the secondary, and slightly larger for the primary. Final errors in the RVs as measured from the scatter in the orbital fit were about 0.44 km s\u22121 for the primary and 0.89 km s\u22121 for the secondary. A sign that systematic errors remained at some level in the CfA velocities was evident in the residuals of the secondary star shown in Figure 2 of T09, in which a pattern can be seen as a function of orbital phase, with a peak semi-amplitude of about twice the typical error. Possible explanations for this, as discussed by T09, include the presence of spots on the active secondary star, or template mismatch.6\n\n6\nIn particular, due to limitations in the available library of synthetic spectra they used, the macroturbulent velocity of the templates (\n\n\n\n\n\n km s\u22121) was not quite as large as appropriate for giant stars. This also resulted in an overestimation of the rotational velocities of the components, as discussed by T09.\n An additional indication of possible biases was the fact that a small offset (0.267 \u00b1 0.079 km s\u22121) was found between the primary and secondary velocities in the global orbital fit of T09 that could not be accounted for by differences in the gravitational redshift between the stars, and was ascribed to similar reasons as the secondary residual pattern.","Citation Text":["Zucker & Mazeh 1994"],"Functions Text":["The RVs were measured using the two-dimensional cross-correlation algorithm TODCOR"],"Functions Label":["Uses"],"Citation Start End":[[590,609]],"Functions Start End":[[506,588]]}
{"Identifier":"2015MNRAS.453.3461S__Shakura_et_al._2012_Instance_1","Paragraph":"The recent detection of X-ray pulsations in PSR J1023+0038 during the accretion active state (Archibald et al. 2015; Bogdanov et al. 2015) implies that channelled accretion, similar to that seen in higher luminosity accreting millisecond X-ray pulsars, is occurring at a much lower accretion rate, implying that the inner edge of the accretion disc lies close to rcor and a propeller does not form. In this case material can accumulate near rcor and non-stationary accretion can occur as matter piles up around the intrinsically unstable magnetospheric boundary. Accretion discs accreting on to the magnetosphere of a rotating star can end up in a trapped state, in which the inner edge of the disc stays near rcor. The captured material can form a quasi-spherical shell (Pringle & Rees 1972; Shakura et al. 2012) or a new disc structure, known as a \u2018dead-disc\u2019 and episodic accretion can occur (D'Angelo & Spruit 2010, 2012). As noted by Patruno et al. (2014), variations in the mass accretion rate of the accretion disc can move the inner disc radius by a factor of 2 or 3. The viscous time-scale defines the time-scale for this drift and to reach a viscous time-scale of 10\u2013100\u2009s requires a region with an annulus of radius \u223c10\u2013100 km. The thermal time-scale is given by ttherm \u223c (H\/R)2\u2009tvis (where H and R are the height and radius of the disc) and for an inner thin disc with H\/R  0.02 implies ttherm few seconds, much lower than the time-scales observed. Thus, in principle, fluctuations in the mass accretion rate can move rm outside rlc ,which allows the radio pulsar to turns on, triggering a transition to the passive state, where the lower X-ray luminosity is produced by a shock between the pulsar wind and innermost accretion flow. An increase in the mass accretion rate pushes the inner edge of the disc back inside the light cylinder and turns off the radio pulsar. The X-ray pulsations are only observed in the active state and not in the passive-state light curves, which suggests that the switching between the passive and active states results in transitions between a non-accreting pure propeller mode and an accreting trapped-disc mode (Archibald et al. 2015).","Citation Text":["Shakura et al. 2012"],"Functions Text":["The captured material can form a quasi-spherical shell"],"Functions Label":["Background"],"Citation Start End":[[793,812]],"Functions Start End":[[716,770]]}
{"Identifier":"2019ApJ...884..132K__Tanihata_et_al._2003_Instance_1","Paragraph":"First, we discuss the discrepancy of the distribution scale of the radio core positions based on the discussions of the internal shock model (Koyama et al. 2015; Niinuma et al. 2015). As is discussed there, the radio cores in Mrk 501 and Mrk 421 observed at 43 or 22 GHz can usually be considered as the internal shocked regions owing to the convex shape of the radio spectrum peaking around 10 GHz (Giroletti et al. 2008; Sokolovsky et al. 2010; Lico et al. 2012; Blasi et al. 2013). The standard internal shock model of blazars considers that the discrete ejecta with higher speeds (with Lorentz factor \u0393f) catch up with the preceding slower ejecta (with Lorentz factor \u0393s), and the collision leads to the nonthermal emission (e.g., Spada et al. 2001; Tanihata et al. 2003; Guetta et al. 2004; Kino et al. 2004; Ghisellini et al. 2005). Based on the model, the distribution scale of the internal shocks (\u0394DIS in Figure 7, defined as the difference between the largest distance between the internal shock and the central engine DIS,max and the closest one DIS,min) can be explained as the variation of the Lorentz factors of the ejecta (Koyama et al. 2015; Niinuma et al. 2015), by assuming the Lorentz factor ratio (\u0393f\/\u0393s) and the initial separation of the ejecta (IIS). The core stable within 200 \u03bcas constrained by the VERA can be explained by Lorentz factors within a factor of two variation for the slower ejecta, i.e., 8 \u2264 \u0393s \u2264 17, by assuming a minimum value of 8 (e.g., Kino et al. 2002), \u0393f\/\u0393s \u2264 1.01 (Tanihata et al. 2003), and IIS \u223c 1 Rs (Koyama et al. 2015). This time we refined the distribution scale of the radio core within 42 \u03bcas along its main jet axis, or 4.6 \u00d7 103 Rs deprojected (see Figure 7). Based on the same assumptions as in Koyama et al. (2015), to explain the further stable distribution scale of the internal shocks, the variation of Lorentz factors of the slower ejecta is constrained to be much smaller, within 30% or 8 \u2264 \u0393s \u2264 10. On the other hand, the radio core wandering of \u0394DIS \u223c 2.6 \u00d7 105 Rs in Mrk 421 can be explained by the maximum value as \u0393s \u223c 60 (with different assumptions; Niinuma et al. 2015). Even by applying the same assumptions to Mrk 421 as those for Mrk 501, the maximum of the slower Lorentz factor is estimated to be \u0393s \u223c 50, which is still a few times as large as that of Mrk 501. Therefore, even during the X-ray and VHE \u03b3-ray active states in 2012, the maximum Lorentz factors that explain the stability of Mrk 501's core are roughly a few times smaller than those for Mrk 421's wandering core, based on the internal shock model.","Citation Text":["Tanihata et al. 2003"],"Functions Text":["The standard internal shock model of blazars considers that the discrete ejecta with higher speeds (with Lorentz factor \u0393f) catch up with the preceding slower ejecta (with Lorentz factor \u0393s), and the collision leads to the nonthermal emission (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[754,774]],"Functions Start End":[[485,734]]}
{"Identifier":"2021MNRAS.508.2743A__Kennedy_et_al._2019_Instance_1","Paragraph":"Star formation from molecular cloud cores collapse is a chaotic process that is expected to result in initially misaligned and warped discs (Bonnell & Bastien 1992; Bate, Lodato & Pringle 2010; Offner et al. 2010; Bate 2018). We expect this to be the general case for circumbinary discs as well as circumstellar ones. Indeed, some misaligned circumbinary discs have already been reported (e.g. KH 15D Chiang & Murray-Clay 2004; Winn et al. 2004; Lodato & Facchini 2013; Smallwood et al. 2019; Fang et al. 2019; Poon, Zanazzi & Zhu 2021, GG Tau A K\u00f6hler 2011; Andrews et al. 2014; Aly, Lodato & Cazzoletti 2018, IRS 43 Brinch et al. 2016, L1551 NE Takakuwa et al. 2017, and HD 98800B Kennedy et al. 2019). Misaligned discs around binaries experience a gravitational torque that leads to radially differential precession, which causes disc warping and twisting. For the gas component, warps are expected to propagate in a wave-like manner in thicker discs with low viscosity that are more relevant in protoplanetary contexts (Papaloizou & Lin 1995; Nelson & Papaloizou 1999; Facchini, Lodato & Price 2013), as opposed to the diffusive warp propagation that occurs in more viscous, thinner discs expected around black holes (Papaloizou & Pringle 1983; Ogilvie 1999; Lodato & Pringle 2007; Lodato & Price 2010). The final state of gas discs precessing around binaries depends on the binary parameters; for circular and low-eccentricity binaries, the disc will align with the binary plane, either in a prograde or retrograde sense depending on the initial misalignment (King et al. 2005; Nixon et al. 2011). For eccentric binaries, discs with high initial misalignments will align in a polar configuration around the binary eccentricity vector (Farago & Laskar 2010; Aly et al. 2015; Martin & Lubow 2017; Zanazzi & Lai 2018). Either way, the disc is susceptible to breaking when the binary torque is larger than the viscous torque and the disc communicates the precession more slowly than it occurs (Nixon, King & Price 2013; Do\u01e7an et al. 2018).","Citation Text":["Kennedy et al. 2019"],"Functions Text":["Indeed, some misaligned circumbinary discs have already been reported","and HD 98800B"],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[683,702]],"Functions Start End":[[318,387],[669,682]]}
{"Identifier":"2015ApJ...800...97T___2012_Instance_1","Paragraph":"The spectral evolution of an ensemble of stars of various masses can be combined through an IMF and followed over time. Such spectral synthesis codes are valuable tools for ultraviolet spectral libraries (Robert et\u00c2 al. 1993; Rix et\u00c2 al. 2004; Leitherer et\u00c2 al. 2014) and stellar population synthesis codes such as Starburst99 (Leitherer et\u00c2 al. 1999). Figure\u00c2 12 shows a series of spectra for a cluster starburst containing 105\u00e2\u0080\u0089M in which the stars follow a Salpeter IMF (0.1 \u00e2\u00a9\u00bd m \u00e2\u00a9\u00bd 120). The two panels show model spectra at times t = 0, 1, 3, 5, and 7\u00e2\u0080\u0089Myr after a coeval burst of star formation, using evolutionary tracks with rotation, at solar metallicity (Z = 0.014 from Ekstr\u00c3\u00b6m et\u00c2 al. 2012) and at sub-solar metallicity (Z = 0.014 from Georgy et\u00c2 al. 2013). These models were created by a Monte Carlo sampling of stars for each mass for which we have an evolutionary track (m = 20, 25, 32, 40, 60, 85, 120). For the Salpeter differential mass distribution, \u00ce\u00be(m) = Km\u00e2\u0088\u0092\u00ce\u00b1 with \u00ce\u00b1 = 2.35, the fraction of stars above mass m is given by\n8The constant K is normalized to the total cluster mass M = 5.862\u00e2\u0080\u0089K, and the total number of stars N = 16.582\u00e2\u0080\u0089K for mass limits mmin = 0.1 and mmax = 120. For M = 105\u00e2\u0080\u0089M, we expect a mean number  stars and mean stellar mass \u00e2\u008c\u00a9m\u00e2\u008c\u00aa = M\/N = 0.354 M. In many of our Monte Carlo samples the numbers fluctuate about this value, with small numbers at the high-mass end of the IMF. The mean numbers of stars at the high-mass end are N(> 60\u00e2\u0080\u0089M) \u00e2\u0089\u0088 30 and N(> 100\u00e2\u0080\u0089M) \u00e2\u0089\u0088 5. If the upper mass was extended to mmax = 200, we would expect 40 stars above 60 M and 15 stars above 100 M. As discussed earlier, the existence of these very massive stars is controversial, owing to resolution effects. Moreover, the LyC from stars at m > 100 M may not escape the embedded cloud from which they formed, or from the dense gas produced in mass-loss episodes or binary mergers (Smith 2014). We therefore do not consider stars above the 120 M track.","Citation Text":["Ekstr\u00c3\u00b6m et\u00c2 al. 2012"],"Functions Text":["The two panels show model spectra at times t = 0, 1, 3, 5, and 7\u00e2\u0080\u0089Myr after a coeval burst of star formation, using evolutionary tracks with rotation, at solar metallicity (Z = 0.014 from"],"Functions Label":["Uses"],"Citation Start End":[[684,705]],"Functions Start End":[[495,683]]}
{"Identifier":"2021ApJ...922...78X__Nan_et_al._2011_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":["Nan et al. 2011"],"Functions Text":["After the discovery of the first FRB","a number of dedicated facilities have been conducted to search FRBs, such as","and the Five-hundred-meter Aperture Spherical radio Telescope (FAST;","All these efforts result in an increasing rate of new FRB detections. Among them, more than 20 repeating FRBs have been reported."],"Functions Label":["Background","Background","Background","Background"],"Citation Start End":[[871,886]],"Functions Start End":[[197,233],[257,333],[802,870],[905,1034]]}
{"Identifier":"2018AandA...615A.148D__Weidner_et_al._(2010)_Instance_2","Paragraph":"The last column in Table 1 reports the number of OB stars minus the \u201cdiffuse\u201d population estimated from their density in the Reference field (22.5 stars per square degree): as is immediately seen, the M-star statistics is much larger than the OB star statistics. This can hardly be considered surprising, if an ordinary IMF (e.g., that from Weidner et al. 2010) is assumed for the star-formation region. In Fig. 19 we show the density ratio between M and OB stars, which provides a consistency test between our results and a plausible IMF: this ratio varies however by a large factor, close to 20, among our subregions. This might reflect differences in the respective IMFs, but also differences in completeness among the stellar samples considered for the various regions. We first note that the ratio between M and OB stars in NGC 6231 is dramatically lower than anywhere else in Sco OB1. We can indeed expect that M stars are detected less efficiently in the inner parts of NGC 6231, where the density of bright stars is very large, and their diffuse glare raises the limiting magnitude locally. As already discussed above in Sect. 4.1, this causes our sample of M stars in NGC 6231 to be highly incomplete. Moreover, we determined above that NGC 6231 is significantly more extincted, by almost half a magnitude in V, than Tr 24, and this implies a higher minimum detectable mass among NGC 6231 M stars compared to Tr 24 (see the MDA diagrams in Fig. 5); this effect reduces the completeness of the M-star sample in NGC 6231 more than in Tr 24. If Tr 24 is also slightly younger than NGC 6231, as we argue below, our M-stars in Tr 24 will reach down to lower masses than in NGC 6231, with a steep increase in the detected M-star population: adopting the IMF from Weidner et al. (2010), the predicted number of cluster M stars doubles considering the mass interval 0.25\u20130.5 M\u2299 rather than 0.35\u20130.5 M\u2299. If Tr 24 is younger than NGC 6231, moreover, its stars in the mass range 2.5\u20133 M\u2299 might not have yet reached their ZAMS position as B stars, and therefore would not be counted among OB stars; this would further raisethe M\/OB star ratio there by up to 30%. Therefore, the proportions of both M and OB stars that are detected in a young cluster will depend on their age and extinction, in accordance with the MDA diagrams, even for a fixed, spatially uniform photometric sensitivity. We estimated using the Weidner et al. (2010) IMF the expected range for the observed M\/OB number ratio. Siess et al. (2000) predict that the latest-type B stars have a mass of ~ 3.5 M\u2299 at 2 Myr, and ~ 2.2 M\u2299 at 10 Myr, that is, in the range of ages expected for Sco OB1 clusters. The MDA diagrams of Fig. 5 predict that the lowest-mass stars we are able to detect using the available Sco-OB1 data have ~ 0.2 M\u2299, even assuming the most favorable (and unlikely) circumstances of an age less than 2 Myr and negligible reddening. The extreme values found for the M\/OB ratio are then ~ 3.8 for a minimum M-star mass as high as 0.35 M\u2299 and an old age of 10 Myr, and ~ 20 for a minimum M-star massas low as 0.2 M\u2299 and age of 2 Myr. These extremes are also shown as horizontal lines in Fig. 19. We note that the M\/OB ratio in NGC 6231 falls well within this range; however, both Tr 24 regions are significantly richer of M stars than expected, by more than a factor of two and well above (statistical) errors. If true, then paradoxically this part of the OB association would form preferentially lower-mass stars. Of course, more detailed studies are needed to confirm this result. In the G345.45+1.50 region the M\/OB ratio is highest, and far above predictions from the IMF: we may tentatively explain this since this region is very young, and some of its most massive members, like IRAS 16562-3959, are still in formation, thus decreasing the number of optically revealed OB stars. The lowest M\/OB ratio in NGC 6231 is unlikely to be real, since as discussed above our M-star sample in this densest subregion is likely incomplete.","Citation Text":["Weidner et al. (2010)"],"Functions Text":["If Tr 24 is also slightly younger than NGC 6231, as we argue below, our M-stars in Tr 24 will reach down to lower masses than in NGC 6231, with a steep increase in the detected M-star population: adopting the IMF from","the predicted number of cluster M stars doubles considering the mass interval 0.25\u20130.5 M\u2299 rather than 0.35\u20130.5 M\u2299."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1766,1787]],"Functions Start End":[[1548,1765],[1789,1903]]}
{"Identifier":"2017MNRAS.470.1442C__Hurley_et_al._2002_Instance_2","Paragraph":"We then allow the synthetic single or binary system to evolve until present time, adopting for our reference model a thin disc age of 10 Gyr (Cojocaru et al. 2014) and a thick disc age of 12 Gyr. This is motivated by the findings of Feltzing & Bensby (2009), who presented a sample of very likely thick disc candidates with ages, on average, well above 10 Gyr and of Ak et al. (2013), who found that thick disc cataclysmic variables have ages up to 13 Gyr. If the synthetic star is single and has time to become a white dwarf, it evolves following the cooling tracks detailed in the following section. If that is the case, the mass of the white dwarf is obtained from the initial-to-final mass relation (IFMR) according to the prescription from Hurley, Tout & Pols (2002). If the object is member of a binary system and the primary star has time to become a white dwarf, then the pair can evolve through two different scenarios. In the first scenario, the binary evolves without mass transfer interactions as a detached system and the primary star evolves into a white dwarf that subsequently cools down following the cooling sequences described in the next section. In this case, the mass of the white dwarf is also calculated from the IMFR of Hurley et al. (2002). The second scenario involves mass transfer episodes and the evolution of the binary is obtained following the prescriptions of the bse package (Hurley et al. 2002), following the parameter assumptions detailed in Camacho et al. (2014). If the system evolves though the common envelope phase, we use the \u03b1-formalism as described in Tout et al. (1997), with \u03b1CE being the efficiency in converting orbital energy into kinetic energy to eject the envelope (assumed to be 0.3 in our reference model). This implementation also takes into account the \u03b1int parameter (assumed to be 0.0 in our reference model), first presented in Han, Podsiadlowski & Eggleton (1995), describing the fraction of the internal energy (thermal, radiation and recombination energy) used to eject the envelope. As described in Camacho et al. (2014), the \u03b1int parameter is used to include the effects of the internal energy in the binding energy parameter \u03bb, which is thus not taken as a constant, but computed using a specific algorithm (Claeys et al. 2014) in bse. In the current version of the code, provided that a positive value is used, the parameter \u03b1int represents the fraction of recombination energy that contributes to eject the envelope. It is important to note that the thermal energy of the envelope is always taken into account (using the virial theorem) even if \u03b1int is set to zero. For a more detailed discussion on how this is implemented in the latest version of BSE and important comments on the correct use of BSE and the notations used in the code itself, we direct the reader to Zorotovic, Schreiber & Parsons (2014a), mentioning that the notations \u03b1int or \u03b1rec are, in our case, equivalent.","Citation Text":["Hurley et al. 2002"],"Functions Text":["The second scenario involves mass transfer episodes and the evolution of the binary is obtained following the prescriptions of the bse package","following the parameter assumptions detailed in Camacho et al. (2014)."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1411,1429]],"Functions Start End":[[1267,1409],[1432,1502]]}
{"Identifier":"2018MNRAS.476..814H__Byun_et_al._2017_Instance_1","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":["Byun et al. 2017"],"Functions Text":["The covariance of the 3PCF, on the other hand, is not diagonal, even for Gaussian fluctuations, which makes the modelling more difficult"],"Functions Label":["Background"],"Citation Start End":[[1403,1419]],"Functions Start End":[[1222,1358]]}
{"Identifier":"2020MNRAS.494.4382S___2010_Instance_1","Paragraph":"It has been thought that QPOs originate from the innermost part of an accretion disc, which is associated with strong gravity, so that we might detect general relativistic effects. Miller et al. (1998) proposed beat-frequency models and estimated the parameters of NSs using this model. Stella & Vietri (1999) developed the relativistic precession model. In the last 20 years, disc-oscillation and resonance models and wave models have been proposed (e.g. Osherovich & Titarchuk 1999; Abramowicz & Klu\u017aniak 2001; Abramowicz et al. 2003; Zhang 2004; Li & Zhang 2005; Erkut, Psaltis & Alpar 2008; Shi & Li 2009, 2010; Shi 2011; Shi, Zhang & Li 2014, 2018; de Avellar et al. 2018). Shi & Li (2009, 2010) obtained the twin modes of MHD waves in LMXBs (including NS LMXBs and black hole LMXBs), which are considered as the sources of high-frequency QPOs. Shi, Zhang & Li (2014, 2018) also considered the waves produced by the two MHD oscillation modes at the magnetosphere radius as the origin of kHz QPOs. A relationship between the frequencies of the twin-peak kHz QPOs and the accretion rate, in which parallel tracks can be explained, was obtained (Shi, Zhang & Li 2018). Recently, many simulations on the oscillations of accreting tori in the accretion process of NS LMXBs (e.g. Kulkarni & Romanova 2013; Parthasarathy, Klu\u017aniak, \u010cemelji\u0107 2017) have been performed, and almost every model can reproduce some of the observed characteristics of QPOs. However, most models cannot fit the observed data perfectly, the observed data. Belloni, M\u00e9ndez & Homan (2005) suggested that the twin kHz QPOs showed no intrinsically preferred frequency ratio, and this weakened support for the resonance models. Morsink & Stella (1999) were able to fit the overall NS data with different masses and spins of NSs using the relativistic precession model; however, Belloni, M\u00e9ndez & Homan (2007) found that there were deviations between the expected and the observed trends. Recently, T\u00f6r\u00f6k et al. (2016b, 2018) identified the observed QPO frequencies with the frequencies of the epicyclic modes of torus oscillations, and suggested that the relationship between the strong modulation of the X-ray flux and high values of QPO frequencies is connected to the orbital motion in the innermost part of an accretion disc. In addition, there are studies that compared a large set of models with the data of many sources in a complex manner (Lin et al. 2011; T\u00f6r\u00f6k et al. 2012, 2016a).","Citation Text":["Shi & Li","2010"],"Functions Text":["In the last 20 years, disc-oscillation and resonance models and wave models have been proposed (e.g."],"Functions Label":["Background"],"Citation Start End":[[595,603],[610,614]],"Functions Start End":[[355,455]]}
{"Identifier":"2018ApJ...854...17L__Larionov_et_al._2016_Instance_1","Paragraph":"Early 15 GHz images using the very long baseline interferometry (VLBI) technique obtained from the very long baseline array (VLBA) indicated a twisted morphology with jet bending on a scale of \u223c20 mas (Kellermann et al. 1998), and multi-epoch 43 GHz observations show jet knots with complex kinematics involving a mixture of apparent superluminal motion and as well as stationary components (Jorstad et al. 2001, 2005). Multi-frequency (15 and 43 GHz) multi-epoch observations during the 2006 radio flare (Fromm et al. 2013a) inferred a possible association between a jet component ejection event at the end of 2005 and a strong radio flare in 2006 April. The authors interpreted the 2006 radio flare as a result of the interaction between a propagating shock and a stationary shock at a de-projected distance of 18 pc from the core. Temporal variability studies during the 2012 September\u2013October multi-band flaring period found a near-simultaneous \u03b3-ray and optical flaring behavior, inferring a co-spatial origin (Cohen et al. 2014; Larionov et al. 2016). The latter study in addition suggested that the measured Stokes parameter variations is consistent with a bright jet knot moving along a helical trajectory. A multi-wavelength polarimetric study during the same flaring phase (Casadio et al. 2015) detected a co-spatial origin from the near-simultaneous variability and identified the passage of a superluminal radio knot coincident with the \u03b3-ray flare. Further evidence, including an intra-day optical polarization variability and clockwise rotation of the electric vector position angle (EVPA, \n\n\n\n\n\n, where U and Q are Stokes components) during the flaring phase, is consistent with a jet knot passing a region hosting helical magnetic fields. A study of \u03b3-ray\u2013optical variability in flux density and polarization between prominent flares during the end of 2016 (Larionov et al. 2017) finds no time lag between the light curves indicating co-spatial origin of synchrotron (optical) and inverse-Compton (\u03b3-ray) flux, and a smaller viewing angle (more energetic, emission closer to jet base) compared to a flare in 2012 when interpreted in terms of a blob or shock wave on a helical trajectory. From the above studies, the \u03b3-ray flares may be associated with optical flux and polarization variability, and the jet kinematics and polarization properties may be described in terms of a helical jet.","Citation Text":["Larionov et al. 2016"],"Functions Text":["Temporal variability studies during the 2012 September\u2013October multi-band flaring period found a near-simultaneous \u03b3-ray and optical flaring behavior, inferring a co-spatial origin"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1035,1055]],"Functions Start End":[[834,1014]]}
{"Identifier":"2021ApJ...914L...6A__Chhiber_et_al._2018_Instance_2","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"],"Functions Text":["As reported in previous works","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"],"Functions Label":["Uses","Uses"],"Citation Start End":[[1020,1039]],"Functions Start End":[[962,991],[1070,1295]]}
{"Identifier":"2017MNRAS.471.4286F__McConnell_&_Ma_2013_Instance_1","Paragraph":"One of the key model ingredients that determines the TDE rates is the distribution of stars in galactic nuclei (Magorrian & Tremaine 1999; Wang & Merritt 2004). Depending on the merger history of the galaxy and the efficiency of feedback on star formation, the stellar density profile can develop either a core or a cusp. For simplicity, we adopt a singular isothermal sphere density profile \u03c1(r) = \u03c32\/2\u03c0GR2 with \u03c3 being the constant velocity dispersion and R the halo virial radius. For a galaxy of halo mass Mh, the relation between the halo mass and the velocity dispersion is simply Mh = 2\u03c32R\/G; while the velocity dispersion can be directly related to the black hole mass using the MBH\u2013\u03c3 relation (Kormendy & Ho 2013; McConnell & Ma 2013; Baldassare et al. 2015; Saglia et al. 2016; Thomas et al. 2016)\n(1)\r\n\\begin{equation}\r\nM_{{\\rm BH}} = 0.309\\times 10^9\\times \\left(\\sigma \/200\\ \\rm{km\\,\\,s}^{-1}\\right)^{4.38}\\ \\mathrm{M}_{\\odot },\r\n\\end{equation}\r\nwhich holds for a wide range of black hole masses from 5 \u00d7 104 M\u2299 (Baldassare et al. 2015) to 1.7 \u00d7 1010 M\u2299 (Thomas et al. 2016) in galaxies with a bulge (Guillochon & Loeb 2015). Assuming the isothermal stellar distribution, Wang & Merritt (2004) derived TDE rates for galaxies with a single central black hole, while Chen et al. (2009) report the rates in a case of a black hole binary. As we discuss in Section 2.2, for MBH with masses in the range MBH \u223c 105\u2013108\u2009M\u2299, the TDE rates per halo computed using the isothermal stellar distribution are similar (within tens of percent) to more realistic estimates based on a large galaxy sample (Stone & Metzger 2016), which justifies our assumption. The error in the rate estimation due to the idealized stellar density profile is small compared to other uncertainties, e.g. introduced by the poorly constrained occupation fraction of IMBHs in low-mass galaxies, which amounts to one\u2013two orders of magnitude uncertainty in the derived volumetric TDE rates.","Citation Text":["McConnell & Ma 2013"],"Functions Text":["For a galaxy of halo mass Mh, the relation between the halo mass and the velocity dispersion is simply Mh = 2\u03c32R\/G; while the velocity dispersion can be directly related to the black hole mass using the MBH\u2013\u03c3 relation"],"Functions Label":["Uses"],"Citation Start End":[[723,742]],"Functions Start End":[[484,701]]}
{"Identifier":"2022AandA...662A..42M__V\u00e1zquez_2007_Instance_1","Paragraph":"A number of fundamental results have been rigorously proved in the mathematical literature concerning the asymptotic behaviour in time of some of the solutions of the porous medium equation and related equations (e.g. Kamin & V\u00e1zquez 1991; Bernis et al. 1993; Hulshof et al. 2001). What is of interest for us here is, primarily, the results that can be applied to the cylindrically symmetric case with diffusion coefficient which is proportional to the square of the dependent variable (n\u2004=\u20042, m\u2004=\u20043 in the notation of Eq. (7)). The most basic result, already mentioned in Sect. 4.4.1, is that initial conditions which have a finite nonzero flux integral (called \u2018the mass\u2019 in the mathematical literature for the PME) converge toward the ZKBP solution with the same flux integral (\u2018mass\u2019) asymptotically in time (V\u00e1zquez 2007, Theorem 18.2); here, allowance is made for either a positive or negative flux integral by globally changing the sign of the ZKBP solution; also, \u2018convergence\u2019 is meant in the sense that the Lp norm of the difference between the actual solution and the ZKBP function tends to zero as t\u2004\u2192\u2004\u221e faster than a negative power of the time with an exponent which is a function of n, m, and p (e.g. \u22121\/3 for n\u2004=\u20042 and m\u2004=\u20043 in the L2 norm; see details in the book by V\u00e1zquez 2007). A complementary result is the following: when the initial condition has positive net flux and its negative part has compact support, then the whole solution evolves into a positive function after a finite time (V\u00e1zquez 2007, Theorem 18.29). Since we are dealing with signed functions which have zero flux integral, these results are of interest mainly because they impose a strict condition on the possible flux imbalance caused by numerical errors (as discussed in Sect. 4.4.1, final paragraph): if it is not small, the numerical solutions will approach the ZKBP solution in a comparatively short time. However, the flux imbalance in all the Bifrost experiments discussed in the present paper is small enough that they have not shown this behaviour even though they have been run until a very long diffusive time.","Citation Text":["V\u00e1zquez 2007"],"Functions Text":["The most basic result, already mentioned in Sect. 4.4.1, is that initial conditions which have a finite nonzero flux integral (called \u2018the mass\u2019 in the mathematical literature for the PME) converge toward the ZKBP solution with the same flux integral (\u2018mass\u2019) asymptotically in time","Theorem 18.2); here, allowance is made for either a positive or negative flux integral by globally changing the sign of the ZKBP solution"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[813,825]],"Functions Start End":[[529,811],[827,964]]}
{"Identifier":"2017ApJ...846...52G__Mauche_et_al._1995_Instance_1","Paragraph":"In addition to the above problem in the UV spectroscopy of disk-dominated systems, X-ray data of CVs have also been in strong disagreement with theoretical expectations for more than three decades (Ferland et al. 1982). The culprit has been the difficult to study boundary layer (BL) between the accretion disk and the WD star. About half of the disk accretion energy (in the form of kinetic energy) is expected to be dissipated in the BL between the Keplerian disk and slowly rotating stellar surface (Pringle 1981). Because of its small size, the BL was predicted to emit in the X-ray band: at low accretion rates (when the WD is dominant in the UV), the BL was expected to be optically thin and emit hard X-rays (Pringle & Savonije 1979; Tylenda 1981; Narayan & Popham 1993); at large accretion rates, typical of NLs in a high state and DNe in outburst, the BL was expected to be optically thick and emit soft X-rays (Pringle 1977; Narayan & Popham 1993; Popham & Narayan 1995). Systems in the low state indeed reveal optically thin hard X-ray emission (Szkody et al. 2002; Pandel et al. 2005; Mukai et al. 2009). However, systems in a state of high mass accretion often do not show an optically thick soft X-ray component; instead, many exhibit optically thin hard X-ray emission (Patterson & Raymond 1985a, 1985b; Mauche et al. 1995; van Teeseling et al. 1996; Baskill et al. 2005; Balman et al. 2014), with an X-ray luminosity much smaller than expected, i.e., much smaller than the disk luminosity. While optically thin hard X-ray emission from high mass accretion rate systems was unexpected, it is, however, not especially inconsistent with the theoretical work: optically thin BLs can occur in high mass accretion rate systems, since the transition to being optically thin depends not only on the mass accretion rate, but also on the WD mass, the WD rotation rate, and the (unknown) alpha viscosity parameter (Popham & Narayan 1995). Simulations of optically thin BLs (Narayan & Popham 1993; Popham 1999) show that the inner edge of the Keplerian (and optically thick) disk starts at an actual radius R0 = Rwd + \u03b4BL, where the size \u03b4BL of the BL is of the order of the stellar radius Rwd: \u03b4BL \u221d Rwd (the BL is actually geometrically thick). The direct consequence of having an optically thin and geometrically thick BL is that the optically thick Keplerian disk will appear to have an inner hole of size \u03b4BL (possibly of the order of the radius of the WD Rwd). Two decades ago, it had already been pointed out that optically thin BL can explain the inner hole observed in circumstellar disks around young stellar objects (T Tauri stars; Godon 1996). In other words, optically thin BLs are consistent not only with the X-ray data, but also with the UV data, as truncated optically thick disks produce a UV continuum with a shallow slope in better agreement with the UV observations than non-truncated disks.","Citation Text":["Mauche et al. 1995"],"Functions Text":["However, systems in a state of high mass accretion often do not show an optically thick soft X-ray component; instead, many exhibit optically thin hard X-ray emission","with an X-ray luminosity much smaller than expected, i.e., much smaller than the disk luminosity."],"Functions Label":["Background","Background"],"Citation Start End":[[1319,1337]],"Functions Start End":[[1117,1283],[1408,1505]]}
{"Identifier":"2021AandA...655A..98A__Agliozzo_et_al._2019_Instance_1","Paragraph":"We have employed a simple grey-body fitting method to model the infrared SED of individual sources. In the LMC, which has the most numerous list of LBVs, large amounts of dust are observed (\u223c10\u22123\u2005\u2212\u200510\u22122\u2006M\u2299), similar to Galactic LBVNe. We stacked the infrared images of the LMC LBVs and extracted the photometry of the resulting source, detected up to 160\u2212250\u2006\u03bcm. The integrated SED from the stacks resembles that of LBVs with a strong ionised stellar wind and an extended dusty nebula. The SED can be fitted with only two components: a power-law describing the free-free spectrum of ionised stellar winds and the stellar photosphere, and a single-component grey-body for the dust. For the grey-body we adopted two different values of the \u03ba parameter, including the value determined by Gordon et al. (2014) to fit the integrated ISM SED of the LMC. A significant contribution to the stack SED comes from a few sources, the most important one is RMC143. This was already identified as a massive nebula (Agliozzo et al. 2019). We obtain an integrated present dust mass of \n\n\n\n0\n.\n\n11\n\n\u2212\n0.03\n\n\n+\n0.06\n\n\n\n\nM\n\u2299\n\n\n\n$ 0.11^{+0.06}_{-0.03}\\,M_{\\odot} $\n\n\n. We have repeated a similar analysis on the sample of AGBs and RSGs by Riebel et al. (2012). We obtain a detection in the stacked images only when considering the \u201cextreme\u201d-AGBs. We find that the integrated 160\u2006\u03bcm emission of 1342 extreme-AGBs is of the same order of magnitude as that of 18 LBVs. The integrated dust mass from these sources is \n\n\n\n1\n.\n\n2\n\n\u2212\n0.4\n\n\n+\n0.3\n\n\n\u00d7\n\n10\n\n\u2212\n4\n\n\n\n\nM\n\u2299\n\n\n\n$ 1.2^{+0.3}_{-0.4}\\times 10^{-4}\\,M_{\\odot} $\n\n\n. We do not find any correlation between the dust masses and the stellar luminosities. This could be due to the fact that such stars have different evolutionary histories or that the dust production mechanism does not depend on the initial mass of the star. Most likely we are also unable to detect the lowest-mass nebulae. To estimate the total dust mass produced by LBVs in the LMC during its full lifetime, we consider two cases: constant number of LBVs across time and a case accounting for IMF and SFH. The uncertainty on the duration of LBV phase in the first case, or on the population of LBVs in the second case, add a significant uncertainty in the total mass produced by LBVs.","Citation Text":["Agliozzo et al. 2019"],"Functions Text":["A significant contribution to the stack SED comes from a few sources, the most important one is RMC143. This was already identified as a massive nebula"],"Functions Label":["Background"],"Citation Start End":[[1001,1021]],"Functions Start End":[[848,999]]}
{"Identifier":"2022MNRAS.513.1459M__Kaviraj,_Martin_&_Silk_2019_Instance_1","Paragraph":"Hierarchical structure formation scenarios (e.g. Fall & Efstathiou 1980; van den Bosch et al. 2002; Agertz, Teyssier & Moore 2011) predict that massive galaxies acquire much of their stellar mass through a combination of continuous cold gas accretion and mergers with smaller objects (e.g. Press & Schechter 1974; Moster, Naab & White 2013; Kaviraj et al. 2015; Rodriguez-Gomez et al. 2016; Martin et al. 2018b; Davison et al. 2020; Martin et al. 2021). As a consequence, mergers are also expected to play a significant role in driving the evolution of galaxy properties, for example, by triggering (Schweizer 1982; Mihos & Hernquist 1996; Duc et al. 1997; Elbaz & Cesarsky 2003; Kaviraj et al. 2011; Lofthouse et al. 2017; Martin et al. 2017) or quenching (Schawinski et al. 2014; Barro et al. 2017; Kawinwanichakij et al. 2017; Pontzen et al. 2017) star formation in the host galaxy or by driving its morphological evolution (e.g. Toomre 1977; Conselice, Yang & Bluck 2009; Dekel, Sari & Ceverino 2009; Taranu, Dubinski & Yee 2013; Naab et al. 2014; Fiacconi, Feldmann & Mayer 2015; Graham, Dullo & Savorgnan 2015; Deeley et al. 2017; G\u00f3mez et al. 2017; Welker et al. 2017; Martin et al. 2018a; Jackson et al. 2019). Signatures of past mergers take the form of faint extended tidal features such as tails (e.g. Pfleiderer 1963; Toomre & Toomre 1972; Peirani et al. 2010; Kaviraj 2014; Kaviraj, Martin & Silk 2019), or plumes (e.g. Lauer 1988) \u2013 which are typically produced by major mergers \u2013 and streams (e.g. Johnston, Sigurdsson & Hernquist 1999; Shipp et al. 2018; Martinez-Delgado et al. 2021) or shells (e.g. Malin & Carter 1983; Quinn 1984) \u2013 which mainly arise from minor interactions \u2013 as well as in the structure of the surrounding diffuse light (e.g. Choi, Guhathakurta & Johnston 2002; Graham 2002; Johnston, Choi & Guhathakurta 2002; Seigar, Graham & Jerjen 2007; Kaviraj et al. 2012; Monachesi et al. 2016, 2019; Iodice et al. 2019; Montes 2019). These features, which arise from many different types of encounter, hold a fossil record of the host galaxy\u2019s past interactions and mergers which can be used to reconstruct its assembly history and dynamical history (Johnston et al. 2008; Mart\u00ednez-Delgado et al. 2009; Belokurov et al. 2017; Montes et al. 2020; Ren et al. 2020; Spavone et al. 2020; Vera-Casanova et al. 2021). However, the majority of tidal features are expected to have surface brightnesses fainter than 30 mag arcsec\u22122 in the r-band (Johnston et al. 2008). Although pushing towards these kinds of limiting surface brightnesses remains extremely challenging, it is nevertheless desirable to do so, being necessary to uncover a more detailed history of local Universe. This is not only vital for our understanding of hierarchical galaxy assembly (e.g. Johnston, Sackett & Bullock 2001; Wang et al. 2012), but also serves as a novel galactic scale probe of more fundamental physics such as theories of gravity (e.g. Gentile et al. 2007; Renaud, Famaey & Kroupa 2016) and dark matter (Dubinski, Mihos & Hernquist 1996; Kesden & Kamionkowski 2006; Dumas et al. 2015; van Dokkum et al. 2018; Montes et al. 2020). In particular, tidal structure is a powerful tracer of the underlying galactic halo potential (e.g. Dubinski, Mihos & Hernquist 1999; Varghese, Ibata & Lewis 2011; Bovy et al. 2016; Ibata et al. 2020; Malhan, Valluri & Freese 2021).","Citation Text":["Kaviraj, Martin & Silk 2019"],"Functions Text":["Signatures of past mergers take the form of faint extended tidal features such as tails (e.g.","\u2013 which are typically produced by major mergers"],"Functions Label":["Background","Background"],"Citation Start End":[[1387,1414]],"Functions Start End":[[1219,1312],[1445,1492]]}
{"Identifier":"2016ApJ...830..156M__Villforth_et_al._2014_Instance_1","Paragraph":"Comparison to a sample of inactive galaxies is also key to demonstrating that an observed merger fraction is actually related to AGN activity. Large samples of inactive galaxies are observed at all merger stages, so it is clear that a major merger alone is not a sufficient condition for quasar activity. Thus, to conclusively demonstrate that mergers are an important channel for quasar fueling, we would need to observe an enhancement to the merger fraction in quasar hosts relative to a matched sample of inactive galaxies. Several studies of lower-luminosity AGN host morphologies with inactive control samples have been conducted in HST extragalactic survey fields (e.g., Grogin et al. 2005; Gabor et al. 2009; Cisternas et al. 2011; Schawinski et al. 2011; Kocevski et al. 2012; B\u00f6hm et al. 2013; Villforth et al. 2014). In particular, we designed our study methodology following Cisternas et al. (2011), who used visual classification to compare strong distortion signatures in moderate-luminosity X-ray selected AGN hosts to a comparison sample of inactive galaxies in the redshift range z = 0.3\u20131.0. They found no significant enhancement to the merger fraction of AGN hosts relative to inactive galaxies, demonstrating that the majority of cosmic black hole mass accretion at \n\n\n\n\n\n, i.e., in AGN with inferred SMBH masses \n\n\n\n\n\n (Vestergaard & Osmer 2009), is not merger driven. How do we then reconcile this result with the results from the red quasar and radio galaxy studies? One possibility is that certain sub-classes of AGN may be preferentially merger driven, even though the bulk of all objects are not. In particular, a downsizing trend has been observed, such that near the peak of quasar activity at z = 2, higher-mass SMBHs dominate the cosmic mass accretion (\n\n\n\n\n\n, Vestergaard & Osmer 2009). It is possible that forming these most massive black holes requires major mergers, as a particularly efficient gas transport mechanism, and that the declining major merger rate of galaxies is one of the driving forces behind this downsizing trend.","Citation Text":["Villforth et al. 2014"],"Functions Text":["Several studies of lower-luminosity AGN host morphologies with inactive control samples have been conducted in HST extragalactic survey fields (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[803,824]],"Functions Start End":[[527,676]]}
{"Identifier":"2019MNRAS.488.5029H__Stacey_et_al._2010_Instance_1","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"],"Functions Text":["[C\u2009ii]\u2009158-\u03bcm luminosity, L[C\u2009II], has been discussed as an indicator of SFR (e.g."],"Functions Label":["Motivation"],"Citation Start End":[[617,635]],"Functions Start End":[[534,616]]}
{"Identifier":"2016MNRAS.459.3161C__Lorimer_&_Kramer_2012_Instance_1","Paragraph":"For an informed opinion on the expected fluxes and durations of short transients, we have to understand coherent emission processes. What we know from incoherent emission physics is that the intrinsic brightness temperature of sources is likely limited to 1011 \u2212 12 K (e.g. Kellermann & Pauliny-Toth 1969; Singal 1986, where brightness temperature is defined as the value of T in the Rayleigh\u2013Jeans formula that yields the correct flux of the source). Sources having brightness temperatures above about 1012 K must emit coherently, have their emission relativistically boosted, or both. However, we understand the processes underlying such sources very poorly, and so in this paper, we shall take the approach of expecting a wide range of known and unknown types of source, and exploring as much of parameter space as our experiment allows. One important thing to note, that is particularly relevant here, is that most of those coherent emitters for which we know the properties of the radio spectrum have quite steep spectra, typically going as \u03bd\u22122 or even \u03bd\u22123 (see e.g. Melrose 2009; Lorimer & Kramer 2012), in contrast with a typical \u03bd\u22120.8 for optically thin synchrotron emission. This means that low-frequency instruments such as LOFAR may be intrinsically at an advantage to find coherent emitters (in addition to having larger fields of view). While known coherent transients have mostly been found in beam-formed searches and last milliseconds to seconds, more recently fast transients have been discovered in low-frequency image plane surveys. For example, the sources GCRT J1745-3009 (Hyman et al. 2005) and GCRT J1746-2757 (Hyman et al. 2002) were detected at 330 MHz with the VLA, while GCRT J1742-3001 (Hyman et al. 2009) was discovered at 235 MHz with the Giant Metrewave Radio Telescope. These sources showed bright flares lasting from minutes to a few hours. More recently, the low-frequency radio transient ILT J225347+862146 (Stewart et al. 2016) was discovered at 60 MHz with LOFAR, lasting about 10 min. The only significant population of transient radio sources previously known in this duration range are relatively nearby and low-luminosity flare stars, having fluxes of about 1 Jy at 1.2 GHz (Osten & Bastian 2006).","Citation Text":["Lorimer & Kramer 2012"],"Functions Text":["One important thing to note, that is particularly relevant here, is that most of those coherent emitters for which we know the properties of the radio spectrum have quite steep spectra, typically going as \u03bd\u22122 or even \u03bd\u22123 (see e.g.",", in contrast with a typical \u03bd\u22120.8 for optically thin synchrotron emission. This means that low-frequency instruments such as LOFAR may be intrinsically at an advantage to find coherent emitters (in addition to having larger fields of view)."],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[1086,1107]],"Functions Start End":[[841,1071],[1108,1349]]}
{"Identifier":"2020ApJ...889..164M__Romani_et_al._2004_Instance_1","Paragraph":"The hard X-ray band has been crucial to studying some of the most powerful blazars (see e.g., Tavecchio et al. 2000; Massaro et al. 2004a, 2004b, 2006; Donato et al. 2005). More recently, the outstanding sensitivity of the Nuclear Spectroscopic Telescope Array (NuSTAR, 3\u201379 keV, Harrison et al. 2013) has enabled us to find and study some of the most distant and luminous ones (e.g., Sbarrato et al. 2013; Tagliaferri et al. 2015; Ajello et al. 2016; Paliya et al. 2016; Sbarrato et al. 2016; Marcotulli et al. 2017). Harboring highly relativistic jets pointed closely at the observer (\u03b8V  1\/\u0393, \u03b8V being the viewing angle and \u0393 the bulk Lorentz factor, \u0393 \u223c 10\u221215, Urry & Padovani 1995), this subclass of the Active Galactic Nuclei (AGNs) is home to some of the most energetic particle acceleration and radiation processes known in astrophysics. The boost in flux ascribed to relativistic beaming, arising from the peculiar orientation of the jets, renders them visible at redshifts well beyond z = 2 (the farthest blazar detected so far is at z = 5.47, Romani et al. 2004), making them extraordinary beacons with which to explore the early universe. Their typical double-hump spectral energy distribution (SED) spans the whole electromagnetic spectrum and is shaped by the nonthermal processes occurring in the jets. Relativistic electrons, spiraling along the magnetic field lines, undergo both synchrotron and inverse Compton (IC) processes. The first produces a peak in the SED located between infrared and X-ray frequencies. The second instead results in a peak located between X- and \u03b3-ray energies. If the electrons interact with a source of low-energy photons within the jet, this is referred to as Synchrotron Self Compton (SSC, e.g., Ghisellini & Maraschi 1989), whereas if the photons are external to the jet (i.e., the accretion disk, the torus, and\/or the broad line region, BLR), it is referred to as External Compton process (EC, e.g., Sikora et al. 1994). Based on their optical spectra, blazars are usually classified either as BL Lacertae objects (BL Lacs) or flat spectrum radio quasars (FSRQs), the first showing weak or no emission lines, the second showing broad (EW > 5 \u212b) ones. Following Abdo et al. (2010a), these sources can also be classified according to the position of the synchrotron peak (\n\n\n\n\n\n), with low-, intermediate-, and high-synchrotron peak (LSP, ISP, HSP) blazars having, respectively, \n\n\n\n\n\n Hz, \n\n\n\n\n\n Hz, and \n\n\n\n\n\n Hz. FSRQs usually belong to the LSP class and at the high-luminosity end of such subclass are the so-called \u201cMeV blazars,\u201d whose high-energy peak falls in (or close to) the MeV band. With bolometric luminosities exceeding 1048 erg s\u22121, these are among the most powerful objects in the universe. In fact, they host powerful relativistic jets (Ghisellini et al. 2014), are usually found at high-redshift (z > 2, e.g., Ajello et al. 2009; Ghisellini et al. 2010; Ackermann et al. 2017; Marcotulli et al. 2017), and typically host billion solar mass black holes (e.g., Ghisellini et al. 2010; Paliya et al. 2017a).","Citation Text":["Romani et al. 2004"],"Functions Text":["the farthest blazar detected so far is at z = 5.47,"],"Functions Label":["Background"],"Citation Start End":[[1054,1072]],"Functions Start End":[[1002,1053]]}
{"Identifier":"2019AandA...629A..54U__Marinucci_et_al._2015_Instance_3","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"],"Functions Text":["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"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1795,1816]],"Functions Start End":[[1639,1793]]}
{"Identifier":"2020MNRAS.499..462S__Quillen_2002_Instance_1","Paragraph":"The physical nature of the X-structures themselves is also an open question, although significant progress has already been made in this direction. The spatial resolution of the models was not very good in the first numerical studies of 3D bar structure (\u2248103\u2212104 particles represented the disc component). X-structures were only distinguished in unsharp masked images constructed from such models. Some authors (Friedli & Pfenniger 1990; Pfenniger & Friedli 1991) therefore suggested that X-structures are akin to an optical illusion due to the tendency of eyes to perceive the intensity gradients instead of actual intensity values. However, with an increase of spatial resolution, it became evident that the X-structures are real density enhancements that can be observed even without unsharp masking processing (see Smirnov & Sotnikova 2018 for many representative examples). The question then is why such density enhancements are observed. Studies of orbit composition of B\/PS bulges and X-structures in different numerical and analytical models (Patsis, Skokos & Athanassoula 2002; Quillen 2002; Quillen et al. 2014; Patsis & Katsanikas 2014a; Parul, Smirnov & Sotnikova 2020) showed that an X-shape is observed in these models due to a tendency of a star to spend more time near turning points of its trajectory. More specifically, 3D bars are constituted by different types of periodic, quasi-periodic, and sticky chaotic orbits (Pfenniger 1984; Pfenniger & Friedli 1991; Patsis et al. 2002; Skokos, Patsis & Athanassoula 2002; Patsis & Katsanikas 2014a,b; Patsis & Harsoula 2018; Patsis & Athanassoula 2019). Stars that move along such orbits spend different periods of time in different parts of their trajectories. For example, stars moving along banana-shaped orbits (Pfenniger & Friedli 1991) spend more time at the highest points of their trajectory (Patsis et al. 2002; Patsis & Katsanikas 2014a). Therefore, the bulk of such orbits produces a density profile with visible density enhancements at the highest points of such orbits. For an X-structure to be observed, these density enhancements should be aligned along an almost straight line for orbits with different apocentric distances. This is indeed the case for the realistic bar potential (Patsis et al. 2002; Quillen 2002; Quillen et al. 2014; Patsis & Katsanikas 2014a). Parul et al. (2020) showed that the orbits of a more complicated morphology than banana-shaped orbits can build an X-structure in a similar manner. In general, cited works showed that X-structures and B\/PS bulges are produced by the same orbits. They are not constituted by different types of orbits like, for example, disc and classical bulge components. The open question that has to be answered in the upcoming studies is what types of orbits are actually presented in real galaxies.","Citation Text":["Quillen 2002","Quillen 2002"],"Functions Text":["Studies of orbit composition of B\/PS bulges and X-structures in different numerical and analytical models","showed that an X-shape is observed in these models due to a tendency of a star to spend more time near turning points of its trajectory.","This is indeed the case for the realistic bar potential"],"Functions Label":["Background","Background","Background"],"Citation Start End":[[1088,1100],[2282,2294]],"Functions Start End":[[945,1050],[1183,1319],[2205,2260]]}
{"Identifier":"2020MNRAS.499.4206B__Johnston_2020_Instance_1","Paragraph":"Lempo\u2013Paha\u2014Hiisi is a trans-Neptunian hierarchical triple system composed of a tight inner binary with components of similar size and an outer companion about half their size orbiting 10 times further away (Trujillo & Brown 2002; Benecchi et al. 2010). All large trans-Neptunian objects like Pluto have multiple small moons, but Lempo\u2019s structure is unique in the Solar system. The place and timing of its origin are still a subject of debate (Nesvorn\u00fd, Youdin & Richardson 2010; Correia 2018). In contrast to the Lempo system, all other known triples in the Solar system have their orbits almost regularly spaced, with one component much smaller than the others, with the most distant component being the largest (Johnston 2020). Unveiling the possible origin of the Lempo system is relevant, as the architecture of multiples holds clues to their formation story and the conditions prevailing in the primitive outer Solar system, but its origin is still unclear. Capture theories proposed so far failed to reproduce the orbital characteristics of observed binaries, especially the distribution of their orbital inclinations (Nesvorn\u00fd et al. 2010; Brunini 2020). Also, triple formation requires multiple captures, a very unlikely event. Gravitational collapse of pebble clouds in a turbulent gas disc would be efficient in producing binaries and, in some particular conditions, can also produce triple systems. However, such triples do not seem to match the orbital structure of Lempo\u2013Paha\u2013Hiisi (Nesvorn\u00fd et al. 2010). The non-detection of triple systems in the cold classical Kuiper belt, where the number of known binaries is much higher than in the resonant populations (Noll et al. 2020), argues against this formation mechanism of triple systems. The fragile dynamical stability of Lempo\u2013Paha\u2013Hiisi (Correia 2018) also casts doubt on the place and time of its origin, leading to speculations about a possible recent formation at the place in which it currently resides.","Citation Text":["Johnston 2020"],"Functions Text":["In contrast to the Lempo system, all other known triples in the Solar system have their orbits almost regularly spaced, with one component much smaller than the others, with the most distant component being the largest"],"Functions Label":["Differences"],"Citation Start End":[[715,728]],"Functions Start End":[[495,713]]}
{"Identifier":"2019AandA...629A..92G__Octau_et_al._2016_Instance_1","Paragraph":"The Nan\u00e7ay Radio Telescope (NRT) is a meridian telescope equivalent to a 94 m parabolic dish located near Orl\u00e9ans (France). Owing to its design, the NRT can track objects with declinations \u03b4\u2004> \u2004\u221239\u00b0 for approximately one hour around culmination, and is thus well suited for the long-term timing of pulsars, for example, for the study of individual objects (see, e.g., Cognard et al. 2017; Octau et al. 2018, for recent examples) or for searching low-frequency gravitational waves from supermassive black hole binaries, using pulsar timing arrays (PTAs; see, e.g., Desvignes et al. 2016). With the goal of identifying new exotic pulsar systems or highly stable MSPs suitable for PTA studies, the SPAN512 pulsar survey (Desvignes et al. 2013; Octau et al. 2016; Desvignes et al., in prep.) was conducted between 2012 and 2018 at the NRT. As part of this survey, new pulsars were searched for at intermediate Galactic latitudes (3.5\u2005\u00b0\u2005 |b|  5\u00b0) and away from the inner Galaxy (Galactic longitudes 74\u2005\u00b0\u2005 l\u2004 \u2004150\u00b0). Observations were conducted at 1.4 GHz with 0.5 MHz frequency channels over a total bandwidth of 512 MHz and a fine time resolution of 64 \u03bcs, to be sensitive to faint and distant MSPs. We used PRESTO pulsar searching routines (Ransom et al. 2002) to search the data for pulsars with dispersion measures (DMs) up to 1800 pc cm\u22123, and a moderate acceleration search in the Fourier domain (the zmax parameter was set to 50 in PRESTO analyses) to be sensitive to pulsars in binary systems. Searches for periodic signals in the data from this survey so far led to the discovery of one \u201cordinary\u201d (i.e., non-millisecond) pulsar, PSR J2048+49, and two MSPs, PSRs J2055+3829 and J2205+6012. Details on the survey, the data analysis, and the discovered pulsars will be reported in Desvignes et al. (in prep.). In the present paper we report on the results from the timing of PSR J2055+3829, an MSP in an eclipsing BW system, and from the analysis of the radio eclipses of the pulsar. In Sect. 2 we describe the radio timing observations and the results from the analysis of the timing data. In Sect. 3 we present observations of eclipses of PSR J2055+3829 at 1.4 GHz, and analyses of the data taken around superior conjunction of the pulsar. In the following section (Sect. 4), we present comparisons of the mass function distributions for eclipsing and non-eclipsing BWs, and for Galactic disk and globular cluster BWs. Finally, Sect. 5 summarizes our findings.","Citation Text":["Octau et al. 2016"],"Functions Text":["With the goal of identifying new exotic pulsar systems or highly stable MSPs suitable for PTA studies, the SPAN512 pulsar survey","was conducted between 2012 and 2018 at the NRT."],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[741,758]],"Functions Start End":[[588,716],[788,835]]}
{"Identifier":"2018MNRAS.480.4154C__Li_et_al._2011a_Instance_1","Paragraph":"In addition to the unconstrained optimization problems of (11) and (12), many CS-based approaches consider constrained forms of the analysis and synthesis models, which are, respectively, given by \n(14)\r\n\\begin{eqnarray*}\r\n\\min _{\\boldsymbol{x}} \\Vert \\boldsymbol {\\mathsf {\\Psi }}^\\dagger {\\boldsymbol{x}}\\Vert _1, \\quad {\\rm s.t.} \\ \\ \\Vert {\\boldsymbol{y}}-\\boldsymbol {\\mathsf {\\Phi }} {\\boldsymbol{x}}\\Vert _2^2 \\le \\epsilon\r\n\\end{eqnarray*}\r\nand \n(15)\r\n\\begin{eqnarray*}\r\n\\min _{\\boldsymbol{a}} \\Vert {\\boldsymbol{a}}\\Vert _1, \\quad {\\rm s.t.} \\ \\ \\Vert {\\boldsymbol{y}}-\\boldsymbol {\\mathsf {\\Phi }}\\boldsymbol {\\mathsf {\\Psi }} {\\boldsymbol{a}}\\Vert _2^2 \\le \\epsilon ,\r\n\\end{eqnarray*}\r\nwhere \u03b5 is an upper-bound related to the noise level present in . CS approaches based on constrained optimization problems, solved via convex optimization techniques, have been applied broadly in RI imaging (Wiaux et al. 2009a,b; McEwen & Wiaux 2011; Li et al. 2011a,b; Carrillo et al. 2012, 2014; Onose et al. 2016; Pratley et al. 2018). These techniques have shown promising results, with improvements in terms of image fidelity and flexibility compared to traditional approaches such as clean-based methods and MEM. For these constrained regularization approaches, parallel implementation structures have also been explored (Carrillo et al. 2014; Onose et al. 2016). Compared with the unconstrained analysis and synthesis models, constrained approaches are parametrized by \u03b5 (related to noise level) which controls the error of the reconstruction explicitly; in contrast, unconstrained models use regularization parameter \u03bc to impose a tradeoff between the prior and data fidelity. The constrained approach therefore avoids the problem of unknown regularization parameter \u03bc, replacing it with the problem of estimating the noise bound \u03b5. The latter can be performed in a principled manner by noting that for Gaussian noise the \u21132 norm data fidelity term follows a \u03c72 distribution with 2M degrees of freedom (see e.g. Carrillo et al. 2012). While constrained problems do not afford a straightforward Bayesian interpretation, the constrained and unconstrained models are closely related (Nikolova 2016).","Citation Text":["Li et al. 2011a"],"Functions Text":["CS approaches based on constrained optimization problems, solved via convex optimization techniques, have been applied broadly in RI imaging","These techniques have shown promising results, with improvements in terms of image fidelity and flexibility compared to traditional approaches such as clean-based methods and MEM."],"Functions Label":["Background","Compare\/Contrast"],"Citation Start End":[[947,962]],"Functions Start End":[[762,902],[1035,1214]]}
{"Identifier":"2018MNRAS.476..814H__Scoccimarro_2000_Instance_1","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"],"Functions Text":["This approximation works well, even for evolved density fields, while deviations from Gaussianity can also be taken into account"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1162,1178]],"Functions Start End":[[1032,1160]]}
{"Identifier":"2018MNRAS.474.2444S__Merloni_et_al._2015_Instance_1","Paragraph":"Other examples of AGN variability are given by the so-called changing-look AGN, which present a change in the AGN type (e.g. from Type 1 to Type 1.9) due to broadening or narrowing of the Balmer lines4 (Denney et al. 2014; LaMassa et al. 2015; Ruan et al. 2015; Gezari et al. 2016; Husemann et al. 2016; McElroy et al. 2016; MacLeod et al. 2016; Runnoe et al. 2016; Stern et al., in preparation). The appearance or disappearance of broad emission lines is often accompanied by a change in luminosity of a factor \u223c10 over \u223c10\u2009yr time-scales. As described above, these time-scales are much shorter compared to time-scales expected for accretion state changes (e.g. Sobolewska, Siemiginowska & Gierli\u0144ski 2011; Hickox et al. 2014), and possible alternative explanations of the changing-look behaviour are variable absorption due to a clumpy torus (e.g. Elitzur 2012), transient events, e.g. tidal disruption of a star by the central black hole (e.g. Eracleous et al. 1995; Merloni et al. 2015), or major changes in the photoionization balance. The magnitude of the drop in luminosity measured in IC\u20092497 is at least a factor of \u223c2 higher than what has been observed in a changing-look AGN. Moreover, the Chandra \u2009and NuSTAR data do not show significant variability, and the upper limits obtained from archival WISE, NEOWISE and IRAS data seem to exclude that the total drop in luminosity happened within the last decades. For these reasons, we argue that the AGN in IC\u20092497 should not be classified as a changing-look AGN. On the other hand, we suggest that a changing-look AGN corresponds to a short-time (\u223c10\u2013100\u2009yr) variability which is superimposed on the long-term (\u223c105\u20136\u2009yr) AGN phases suggested by this work and other observations (e.g. Schawinski et al. 2015), high resolution (sub-kpc) simulations (Hopkins & Quataert 2010; Bournaud et al. 2011; Novak, Ostriker & Ciotti 2011; Gabor & Bournaud 2013; DeGraf et al. 2014; Sijacki et al. 2015) and theoretical models (Siemiginowska & Elvis 1997; Sanders 1981; Di Matteo, Springel & Hernquist 2005; Hopkins et al. 2005; Springel, Di Matteo & Hernquist 2005; King & Pringle 2007; King & Nixon 2015).","Citation Text":["Merloni et al. 2015"],"Functions Text":["and possible alternative explanations of the changing-look behaviour are","transient events, e.g. tidal disruption of a star by the central black hole (e.g."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[970,989]],"Functions Start End":[[729,801],[865,946]]}
{"Identifier":"2022ApJ...927...91V__Koopmann_&_Kenney_2004_Instance_1","Paragraph":"According to simulations, blue galaxies infalling from the field reach pericenter in about 2.4 Gyr (McGee et al. 2009; see also Oman et al. 2013), and we can hypothesize that, by the time they get there, they lose most of their gas. In the first part of this time frame, they can start feeling the action of the ram pressure, and their neutral gas content, which is less bound, starts being stripped. No signs at optical wavelengths are visible yet, but the stripping has started (Tonnesen & Bryan 2010). In some galaxies, ram pressure stripping might be able to unwind the spiral arms (Bellhouse et al. 2021), while in other galaxies it just strips the gas that at some point collapses and starts forming new stars (e.g., Poggianti et al. 2019a). In this phase, tails shine at optical wavelengths, as OB massive stars are very bright. We can assume the visibility of this phase lasts on the order of \u223c6 \u00d7 108 yr (Fumagalli et al. 2011; Poggianti et al. 2019a). Stripping and star formation consume the available gas and galaxies first appear as truncated disks (Koopmann & Kenney 2004; Fritz et al. 2017) and then become passive, showing k+a spectra (Vulcani et al. 2020). Galaxies typically maintain blue colors for 0.5 Gyr after becoming passive and then move to the red sequence (Poggianti et al. 2004). In these phases, tails are not visible anymore at optical wavelengths. So to summarize, an infalling galaxy maintains its blue color for at least 2.4 + 0.5 \u223c3 Gyr since infall and a tail is visible only for 6 \u00d7 108yr. We can assume that the sum of RS+SC+UG samples corresponds to all the noninteracting blue late-type galaxies in the clusters and that the tailed galaxies are either only SC or SC+UG. In the first case, \u223c15% of the blue cluster galaxies currently show signs of stripping, but since the visibility phase is 6 \u00d7 108\/3 \u00d7 109 \u223c 0.2, the total amount of galaxies undergoing stripping is 15%\/0.2 \u223c 75%. In the second case, the incidence of blue cluster galaxies currently showing signs of stripping is even higher. On one hand, this suggests that all blue cluster galaxies undergo a stripping phase during their life in clusters, and on the other hand, it indicates that most likely we are overestimating the number of ram pressure stripped galaxies only using optical imaging.","Citation Text":["Koopmann & Kenney 2004"],"Functions Text":["Stripping and star formation consume the available gas and galaxies first appear as truncated disks"],"Functions Label":["Background"],"Citation Start End":[[1063,1085]],"Functions Start End":[[962,1061]]}
{"Identifier":"2017AandA...607L...7M__Menezes_et_al._(2013)_Instance_1","Paragraph":"Andromeda is a galaxy that lies in the green valley (Mutch et al. 2011; Tempel et al. 2011; Jin et al. 2014). According to Belfiore et al. (2016), it is typically a low-ionisation emission-line region (LIER), as first observed by Rubin & Ford (1971) and discussed by Heckman (1996). Gonz\u00e1lez-Mart\u00edn et al. (2015) discussed that the torus is disappearing in LIER: there is indeed little gas in the inner part of M31 (Melchior et al. 2000; Melchior & Combes 2011, 2013). It is the closest external large galaxy in which we can explore the mechanisms that quenched the star formation activity. Optical ionised gas has been observed by Menezes et al. (2013) next to the black hole in a field of view1 of 5\u2032\u2032 \u00d7 3.5\u2032\u2032, but this emission is weak. Jacoby et al. (1985) estimated the ionised gas mass in the bulge (10\u2032 \u00d7 10\u2032) to be of the order of 1500 M\u2299. It also hosts a very massive black hole of 1.4 \u00d7 108M\u2299 (Bender et al. 2005), but as studied by Li et al. (2011a), it is non-active and only murmurs at a level of 10-10LEdd. It hosts very little star formation of the order of 0.25\u22120.3 M\u2299\u2009yr-1, mainly located in the 10 kpc ring of the disc (e.g. Ford et al. 2013; Rahmani et al. 2016). Inside the central region (10\u2032 \u00d7 10\u2032), no obvious sign of star formation is detected (e.g. Kang et al. 2012; Azimlu et al. 2011; Amiri & Darling 2016), except for a central cluster of A stars formed 200\u2009Myr ago that is located next to the black hole (within 1\u2032\u2032) (Lauer et al. 2012), designated by P3 by Bender et al. (2005). Viaene et al. (2014) estimated the star formation rate (SFR) on a pixel basis with panchromatic spectral energy distribution modelling. This infrared-based SFR estimated in the central pixel (36\u2032\u2032 \u00d7 36\u2032\u2032) is 4 \u00d7 10-5M\u2299\u2009yr-1, while an integration over the central region with a radius of 1 kpc corresponds to 1.25 \u00d7 10-3M\u2299\u2009yr-1. This negligible SFR is much lower than the value predicted by Rimoldi et al. (2016), considering supernovae remnants expected within the sphere of influence of quiescent supermassive black holes. For M31, an SFR of 0.13 M\u2299yr-1 is expected in the sphere of influence (RSOI = 14 pc = 3.7\u201d) of its supermassive black hole. A past AGN activity is also expected, and the associated molecular torus, if it survives, should have a radius RMT = 25 pc = 6.7\u2032\u2032. In parallel, Chang et al. (2007) expected next to the black hole an accumulation of molecular gas (about 104M\u2299) originating from stellar feed-back. Melchior & Combes (2013) estimated a minimum molecular mass of 4.2  \u00d7  104M\u2299 within 30\u2033 from the centre, while about 106M\u2299 of gas is expected from stellar feedback (e.g. Gallagher & Hunter 1981). ","Citation Text":["Menezes et al. (2013)"],"Functions Text":["Optical ionised gas has been observed by","next to the black hole in a field of view1 of 5\u2032\u2032 \u00d7 3.5\u2032\u2032, but this emission is weak."],"Functions Label":["Background","Background"],"Citation Start End":[[632,653]],"Functions Start End":[[591,631],[654,739]]}
{"Identifier":"2021MNRAS.502.4794N__Dullo_&_Graham_2012_Instance_1","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"],"Functions Text":["Initially the core size of a galaxy was determined by fitting the so-called \u2018Nuker-profile\u2019","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."],"Functions Label":["Background","Background"],"Citation Start End":[[964,983]],"Functions Start End":[[624,715],[736,943]]}
{"Identifier":"2020ApJ...895...51M__Rogers_&_McElwaine_2017_Instance_1","Paragraph":"The mechanism(s) driving the transport of angular momentum (e.g., Aerts et al. 2019a) and chemical elements (e.g., Salaris & Cassisi 2017) within stars are still not understood from stellar evolution theory. Discrepancies between observations and theory have been shown for stars with birth masses between 1.3 and 8 M\n\n\n\n\n\n, which comprise a convective core, enshrouded by a radiative envelope (possibly with internal convective shells from partial ionization zones or a thin outer convective envelope for \n\n\n\n\n\n M\n\n\n\n\n\n). In these radiative envelopes, the transport of chemical elements on a macroscopic scale is ascribed to convective core overshooting (e.g., Viallet et al. 2015), rotation (e.g., Maeder 2009), or internal gravity waves (IGWs, e.g., Rogers & McElwaine 2017), whereas the transport of chemical elements on a microscopic scale is the result of atomic diffusion (Michaud et al. 2015). In stellar evolution codes, the description of macroscopic mixing introduces additional free parameters, whereas mixing from atomic diffusion can be derived from first principles. So far, the theory of element transport has been mainly evaluated by measurements of surface abundances. Asteroseismology constitutes a novel technique to empirically assess the conditions deep inside the interior (Aerts et al. 2010) of a star, as well as its evolutionary history (e.g., Bowman et al. 2019). The unprecedented high-quality data from the space-based CoRoT (Auvergne et al. 2009), Kepler  (Borucki et al. 2010), and TESS (Ricker et al. 2015) missions allow for scrutiny of the current stellar evolution models of the stars\u2019 interiors by means of gravito-inertial asteroseismology (Aerts et al. 2018). The current state-of-the-art stellar models and pulsation codes are not capable of reproducing the observed oscillation frequencies of gravity (g) modes in \u03b3 Doradus (\u03b3 Dor; see Kurtz et al. 2014; Saio et al. 2015; Schmid & Aerts 2016; Van Reeth et al. 2016) and slowly pulsating B-type (see P\u00e1pics et al. 2014; Moravveji et al. 2015; Szewczuk & Daszy\u0144ska-Daszkiewicz 2018; Aerts et al. 2019b) stars within the uncertainties of the data. Hence, additional physics is required in order to improve both the current stellar models and the prediction of the g-mode frequencies from these equilibrium models. Studies have already demonstrated the potential of g-modes to distinguish between different near-core mixing profiles (Pedersen et al. 2018) and the temperature gradient close to the convective core interface (Michielsen et al. 2019). The work of Aerts et al. (2018) has evaluated a hierarchy of input physics when modeling g-modes across a wide mass range. In the current work, we investigate to what extent the process of atomic diffusion can improve the theoretically predicted oscillation frequencies in two slowly rotating \u03b3 Dor stars.","Citation Text":["Rogers & McElwaine 2017"],"Functions Text":["In these radiative envelopes, the transport of chemical elements on a macroscopic scale is ascribed to","or internal gravity waves (IGWs, e.g.,"],"Functions Label":["Background","Background"],"Citation Start End":[[753,776]],"Functions Start End":[[523,625],[714,752]]}
{"Identifier":"2018MNRAS.477.4308R__Matteo,_Springel_&_Hernquist_2005_Instance_1","Paragraph":"It is well known that the accretion on to compact objects may influence the nearby ambient around SMBHs in the centre of galaxies (e.g. Salpeter 1964; Fabian 1999; Barai 2008; Germain, Barai & Martel 2009). Together with the outflow phenomena, it is believed to play a major role in the feedback processes invoked by modern cosmological models (i.e. \u039b-Cold Dark Matter) to explain the possible relationship between the SMBH and its host galaxy (e.g. Magorrian et al. 1998; Gebhardt et al. 2000) as well as in the self-regulating growth of the SMBH. The problem of accretion on to an SMBH has been studied via hydrodynamical simulations (e.g. Ciotti & Ostriker 2001; Di Matteo, Springel & Hernquist 2005; Li et al. 2007; Ostriker et al. 2010; Novak, Ostriker & Ciotti 2011). In numerical studies of galaxy formation, spatial resolution permits resolving scales from the kpc to the pc, while the sub-parsec scales of the Bondi radius are not resolved. This is why a prescribed sub-grid physics is employed to solve this lack of resolution. With sufficiently high X-ray luminosities, the falling material will have the correct opacity, developing outflows that originate at sub-parsec scales. Therefore, calculations of the processes involving accretion on to SMBH have become of primary importance (e.g. Proga, Stone & Kallman 2000; Proga 2000, 2003; Proga & Kallman 2004; Proga 2007; Ostriker et al. 2010). Numerical calculations of the accretion of matter on to SMBHs, including the radiative-outflow component, have been mostly performed using Eulerian finite-difference methods (Mo\u015bcibrodzka & Proga 2013) [see also the overviews by Edgar (2004) and Foglizzo et al. (2005) and references therein for earlier work], while only a few calculations have been reported using smoothed particle hydrodynamics (SPH) techniques (Barai 2008; Barai, Proga & Nagamine 2011, 2012; Nagamine, Barai & Proga 2012), where results for the accretion rates, outflow rates, thermal instabilities, and impact of the thermal, mechanical, and X-ray feedbacks have been obtained for evolutions up to 5\u2009Myr and scales from 0.1 to 200\u2009pc.","Citation Text":["Di Matteo, Springel & Hernquist 2005"],"Functions Text":["The problem of accretion on to an SMBH has been studied via hydrodynamical simulations (e.g."],"Functions Label":["Background"],"Citation Start End":[[666,702]],"Functions Start End":[[549,641]]}
{"Identifier":"2021AandA...646A..53P__Roman-Oliveira_et_al._2019_Instance_1","Paragraph":"In addition, the Abell 901\/902 field has deep and extensive photometry in the UV (Galex, Gray et al. 2009), optical (COMBO-17 survey, Wolf et al. 2003), near-infrared (Spitzer 24 \u03bcm, Bell et al. 2007), and X-ray (XMM-Newton, Gilmour et al. 2007) wavelength range. Out of the four main substructures that embody Abell 901\/902, only the two most massive (A901a and A901b) show significant X-ray emission, but this may come from a very bright AGN close to the cluster center in A901a (Gilmour et al. 2007). The evolution of the cluster members\u2019 colors, morphologies, and star-formation activity has been studied over the years, and redder stellar populations, lower star-formation rates, and indications of interactions within the cluster population of galaxies were reported (Wolf et al. 2009; Gallazzi et al. 2009). More recently, the emission-line OMEGA-OSIRIS survey provided similar results (e.g., Chies-Santos et al. 2015; Rodr\u00edguez del Pino et al. 2017; Weinzirl et al. 2017; Roman-Oliveira et al. 2019). Our group carried out a kinematic analysis of a subsample of cluster galaxies using slit spectra taken by VIMOS at the Very Large Telescope (VLT) with the HR-blue grism (R\u2004\u223c\u20042000) to search for indications of ram-pressure stripping (B\u00f6sch et al. 2013a) and study the slope and scatter of the TFR (B\u00f6sch et al. 2013b). Galaxies within three times the velocity dispersion (3\u03c3) of each substructure are considered to be cluster members (B\u00f6sch et al. 2013a). We retrieved a sample of 45 cluster galaxies from B\u00f6sch et al. (2013b) characterized by their high-quality rotation curves, their disk morphology, and stellar masses above log(M*) = 9.5. This sample is used in Sect. 4 to explore the evolution of the TFR between clusters at different redshifts. The distribution of the main parameters (M* and Vmax) and their mean properties are shown in Fig. 1 and listed in Table 3. However, the sizes and exact coordinates of these galaxies within the Abell 901\/902 cluster complex are not provided in the publications mentioned above, and we therefore excluded this sample from the analysis that requires them.","Citation Text":["Roman-Oliveira et al. 2019"],"Functions Text":["More recently, the emission-line OMEGA-OSIRIS survey provided similar results"],"Functions Label":["Background"],"Citation Start End":[[980,1006]],"Functions Start End":[[815,892]]}
{"Identifier":"2020MNRAS.492..686L__Shiokawa_et_al._2015_Instance_2","Paragraph":"After the disruption phase, the star is tidally stretched into a very long thin stream and the evolution of the stream structure in the transverse and longitudinal directions are decoupled (Kochanek 1994). Thus, the system enters the free-fall phase where each stream segment follows its own geodesic like a test particle (Coughlin et al. 2016). Then, after passing the apocentres of the highly eccentric orbits, the bound debris falls back towards the BH at a rate given by the distribution of specific energy (Evans & Kochanek 1989; Phinney 1989). Due to relativistic apsidal precession, the bound debris, after passing the pericentre, collides violently with the still in-falling stream (see Fig. 1). It has been shown that shocks at the self-intersection point is the main cause of orbital energy dissipation and the subsequent formation of an accretion disc (Rees 1988; Kochanek 1994; Hayasaki, Stone & Loeb 2013; Guillochon, Manukian & Ramirez-Ruiz 2014; Shiokawa et al. 2015; Bonnerot et al. 2016). However, the aftermath of the self-intersection is an extremely complex problem, which depends on the interplay among magnetohydrodynamics, radiation, and general relativity in 3D. No numerical simulations to date have been able to provide a deterministic model for TDEs with realistic star-to-BH mass ratio and high eccentricity (see Stone et al. 2018a, for a review). Many simulations consider either an intermediate-mass BH (e.g. Guillochon et al. 2014, Evans, Laguna & Eracleous 2015; Shiokawa et al. 2015; Sa\u0327dowski et al. 2016) or the disruption of a low-eccentricity (initially bound) star (e.g. Bonnerot et al. 2016; Hayasaki, Stone & Loeb 2016). It is unclear how to extrapolate the simulation results to realistic configurations and provide an answer to the following questions: How long does it take for the bound gas to form a circular accretion disc (if at all)? How much radiative energy is released from the system? What fraction of the radiation is emitted in the optical, UV, or X-ray bands?","Citation Text":["Shiokawa et al. 2015"],"Functions Text":["Many simulations consider either an intermediate-mass BH (e.g."],"Functions Label":["Background"],"Citation Start End":[[1495,1515]],"Functions Start End":[[1376,1438]]}
{"Identifier":"2022MNRAS.512.5165A__Uhlemann_et_al._2016_Instance_1","Paragraph":"The number of objects in a specific volume such as a sphere is called counts-in-cells (CIC) (Hubble 1936; Zwicky 1957; Balian & Schaeffer 1989). We know that in small scales, the distribution of CIC or equivalently one-point probability of the matter density field is not Gaussian and follows approximately a lognormal distribution (Coles & Jones 1991; Ueda & Yokoyama 1996). The nearly lognormal behaviour of the cosmic matter density field is a feature of evolving perturbations from the linear Gaussian inflationary initial conditions (Guth 1981; Linde 1982) to the late time non-linear and non-Gaussian distribution. Accordingly, it is informative to investigate the shape of the probability distribution of the matter density field and its evolution (Lam & Sheth 2008; Bernardeau, Pichon & Codis 2014; Uhlemann et al. 2016; Ivanov, Kaurov & Sibiryakov 2019; Mandal & Nadkarni-Ghosh 2020; Repp & Szapudi 2020b). It is worth mentioning that other interesting ideas such as an Edgeworth expansion, skewness, and kurtosis analysis are argued in this field of study (Colombi 1994; Gaztanaga, Fosalba & Elizalde 2000; Shin et al. 2017; Klypin et al. 2018; Einasto et al. 2021). On large scales, the CIC contains the same information as the two-point correlation function. But there is more information encoded in the lognormal shape of the CIC distribution on small scale. Accordingly, the cosmological dependence of the CIC statistics is interesting to be explored in mildly and strongly non-linear regimes from an observational and theoretical point of view. Uhlemann et al. (2020) showed that the measured CIC statistics are sensitive to the cosmological parameters and neutrino mass. Repp & Szapudi (2020a) used the CIC to break the degeneracy between \u03c38 and galaxy bias. Also, it is employed to study the primordial non-Gaussianity in LSS by measuring the fNL parameter (Friedrich et al. 2020). Recently, Jamieson & Loverde (2020) introduced an approach to use a position-dependent one-point distribution of matter density as a cosmological observable. Besides all this progress, there are some limitations. The results depend on the size and shape of the cell, where for large smoothing scales the distribution is Gaussian whereas, for small scales, the distribution is approximately lognormal. So for each cell size, one should calculate the CIC distribution, which is a computationally expensive procedure. Finally, we conclude that although the CIC is a computationally costly process, it is advantageous to explore different regimes and compare the linear and non-linear scales.","Citation Text":["Uhlemann et al. 2016"],"Functions Text":["Accordingly, it is informative to investigate the shape of the probability distribution of the matter density field and its evolution"],"Functions Label":["Motivation"],"Citation Start End":[[807,827]],"Functions Start End":[[621,754]]}
{"Identifier":"2016MNRAS.462.3912Q__Pinnington_et_al._1974_Instance_1","Paragraph":"When comparing the results obtained with our computational approach with the previously published decay rates, we find an overall good agreement. This is illustrated in Figs 2\u20135, where our calculations are compared with the experimental data reported by Salih et al. (1985), Crespo L\u00f3pez-Urrutia et al. (1994), Mullman et al. (1998a,b) and those obtained theoretically by Raassen et al. (1998). More precisely, as shown in Fig. 2, we obtain a mean ratio of 1.036 \u00b1 0.205 when comparing our transition probabilities with those of Salih et al. (1985), who combined lifetime values obtained by TR-LIF with branching fractions measured on spectra recorded with the 1-m Fourier-transform spectrometer at the Kitt Peak National Observatory, to deduce the decay rates for 41 transitions depopulating the 3d74p z5F, z5D and z5G levels. Fig. 3 shows a slightly larger scatter between our calculations and the transition probabilities published by Crespo L\u00f3pez-Urrutia et al. (1994). Here the mean ratio gAThiswork\/gACrespo is equal to 1.098 \u00b1 0.493. However, it is worth noting that the gA-values obtained by the latter authors were deduced from the combination of branching ratio measurements with available experimental lifetimes for the 3d74p z5F, z5D, z5G levels (Pinnington, Lutz & Carriveau 1973; Pinnington et al. 1974; Salih et al. 1985) but also with estimated lifetimes for the z3G, z3F and z3D levels from the measurement of total intensity of all lines of each of these levels under the assumption of almost equal population. The branching fractions were found using intensity measurements with a special hollow electrode radio frequency discharge and using a phase method with a modified wall-stabilized arc in a spectrointerferometric arrangement. Mullman et al. (1998a) reported 28 oscillator strengths for ultraviolet Co\u2009ii lines deduced from laser-induced fluorescence lifetimes and branching fraction measurements using a high-resolution grating spectrometer and an optically thin hollow cathode discharge. As illustrated in Fig. 4 , our gf-values were found to agree in general within 20\u201330 per cent with the results obtained by Mullman et al. (1998a). More particularly, the mean ratio gfThiswork\/gfMullman was found to be equal to 1.339 \u00b1 0.552 when using all the common lines, and to 1.238 \u00b1 0.411 when using the strongest transitions for which log gf > \u22121. Finally, Fig. 5 shows the comparison between our computed oscillator strengths and those published by Raassen et al. (1998) who used the theoretical method of orthogonal operators to determine the log gf-values for (3d8+3d74s)\u20133d74p transitions in Co\u2009ii. In this case, the two sets of data generally agree within about 20 per cent, this percentage difference being reduced to 12 per cent when considering the most intense transitions with log gf > \u22121. We also note that our oscillator strengths are on average larger than those obtained by Raassen et al. This could be due to the fact that these latter authors included explicitly a less extended set of interacting configurations in their model.","Citation Text":["Pinnington et al. 1974"],"Functions Text":["However, it is worth noting that the gA-values obtained by the latter authors were deduced from the combination of branching ratio measurements with available experimental lifetimes for the 3d74p z5F, z5D, z5G levels","but also with estimated lifetimes for the z3G, z3F and z3D levels from the measurement of total intensity of all lines of each of these levels under the assumption of almost equal population."],"Functions Label":["Background","Background"],"Citation Start End":[[1294,1316]],"Functions Start End":[[1041,1257],[1337,1528]]}
{"Identifier":"2020AandA...635A.121M__Casassus_et_al._2018_Instance_1","Paragraph":"As scattered light imaging is sensitive to the stellar irradiation, it allows one to search for misalignments between various disk regions. While studying the morphology of the innermost disk region is challenging due to its very small radial extent, often marginally resolvable by optical interferometry (Lazareff et al. 2017), scattered light imaging of the outer disk can indirectly reveal the presence of a misaligned inner disk. In this scenario, depending on the misalignment angle, the outer disk image will show narrow shadow lanes (e.g., Pinilla et al. 2015; Stolker et al. 2016; Benisty et al. 2017; Casassus et al. 2018), broad extended shadows (Benisty et al. 2018) or low-amplitude azimuthal variations (Debes et al. 2017; Poteet et al. 2018). In some cases, studies of the CO line kinematics support a misalignment between inner and outer disk regions (Loomis et al. 2017; P\u00e9rez et al. 2018). The exact origin of such a misalignment is still unclear. In the case of T Tauri stars, if the stellar magnetic field is inclined, it can warp the innermost edge of the disk, which would then rotate at the stellar period (AA Tau; Bouvier et al. 2007). Alternatively, inner and outer disk regions can have different orientations if the primordial envelope had a different angular momentum vector orientation at the time of the inner\/outer disk formation (Bate 2018). Other scenarios involve the presence of a massive companion\/planet that is inclined with respect to the disk. If the companion is massive enough, the disk can break into two separate inner and outer disk regions, that can then precess differently and result in a significant misalignment between each other (e.g., Nixon et al. 2012; Facchini et al. 2013; Nealon et al. 2018; Zhu 2019). A clear example of such a scenario is the disk around HD 142527, in which an M-star companion was detected (Biller et al. 2012), likely on an inclined and eccentric orbit (Lacour et al. 2016; Claudi et al. 2019). Dedicated hydrodynamical simulations successfully reproduce most of the observed features in this disk (eccentric cavity, spiral arms, misaligned inner disk and shadows; Price et al. 2018).","Citation Text":["Casassus et al. 2018"],"Functions Text":["In this scenario, depending on the misalignment angle, the outer disk image will show narrow shadow lanes (e.g.,"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[610,630]],"Functions Start End":[[434,546]]}
{"Identifier":"2021MNRAS.507.4367C__Mazzali_et_al._2005_Instance_1","Paragraph":"In the previous section, we discussed how SNe Ia with blueshifted Na\u2009i D have broader light curves but this may be driven by their preference for late-type host galaxies and their differing environments. One of the main aims of this study was to investigate, for the first time, if there is a connection between the presence of broad high-velocity Ca\u2009ii features in SNe Ia and the presence of narrow blueshifted Na\u2009i D features. As previously discussed, high-velocity Ca\u2009ii features are ubiquitous in SNe Ia and we confirm this here finding that the vast majority of the sample require a high-velocity Ca\u2009ii component. We also identify high-velocity Si\u2009ii features in 12 SNe Ia in our sample, although these features are not clearly \u2018detached\u2019 from the photospheric component. The reason that high-velocity features are interesting to investigate in connection with blueshifted Na\u2009i D features is that there are suggestions that the high-velocity Ca\u2009ii features may be at least partially due to ejecta\u2013CSM interaction (e.g. Mazzali et al. 2005; Tanaka et al. 2006). Therefore, if a link between these quantities was identified, it would provide evidence that blueshifted Na\u2009i D features are due to CSM also rather than contamination. We remind the reader that the two probes do explore different distances from the SN, with Ca\u2009ii originating at significantly shorter distances than Na\u2009i D. However, when we examine the strength of the high-velocity Ca\u2009ii components (parametrized through the ${\\rm Ca\\, {\\small{\\rm II}}}~R_{\\textrm {HVF}}$ of Childress et al. 2013a) compared to the blueshifted${\\rm Na\\, {\\small{I}}}~\\text{D}_{2}$ pEQW (15), we do not identify any clear trend between them. In particular, we also find that a number of SNe Ia with no Na\u2009i D absorption at all have very high ${\\rm Ca\\, {\\small{\\rm II}}}~R_{\\textrm {HVF}}$ values, which is difficult to interpret in the context of a common CSM origin for both strong blueshifted Na\u2009i D features and strong high-velocity Ca\u2009ii components, as with this interpretation of their origin a strong signature would be expected in both probes rather than singularly. We also find no correlation between the pEQW of the high-velocity Ca\u2009ii features and the pEQW of the blueshifted Na\u2009i D features.","Citation Text":["Mazzali et al. 2005"],"Functions Text":["The reason that high-velocity features are interesting to investigate in connection with blueshifted Na\u2009i D features is that there are suggestions that the high-velocity Ca\u2009ii features may be at least partially due to ejecta\u2013CSM interaction (e.g.","Therefore, if a link between these quantities was identified, it would provide evidence that blueshifted Na\u2009i D features are due to CSM also rather than contamination."],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[1024,1043]],"Functions Start End":[[777,1023],[1066,1233]]}
{"Identifier":"2019MNRAS.486.4671M__Webb_&_Hundhausen_1987_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":["Webb & Hundhausen 1987"],"Functions Text":["CMEs have been observed to occur often having spatial and temporal relation with solar flares, eruptive prominences"],"Functions Label":["Background"],"Citation Start End":[[997,1019]],"Functions Start End":[[861,976]]}
{"Identifier":"2021AandA...646A..67W__Widrow_et_al._2014_Instance_1","Paragraph":"We consider time-varying dynamical effects, breaking the assumption of a steady state, to be the most probable reason for this unexpected result. We do see how local phase-space substructures can bias our result for individual stellar samples, most clearly for those of area cell B8. However, explaining the steep gravitational potential close to the Galactic mid-plane, inferred for almost all stellar samples, requires a phase-space structure that spans the whole spatial volume that is studied in this work. Indeed, spatially large time-varying phase-space structures are present in the Galaxy, for example in the form of phase-space spirals and ridges (Gaia Collaboration 2018c; Antoja et al. 2018), and Galactic plane mirror asymmetries (Bennett & Bovy 2019, especially prominent for heights |z|\u2273400 pc). In order to produce a steep gravitational potential at low |z|, there could be a breathing mode in the stellar disk (Widrow et al. 2014; Monari et al. 2016) which is currently in its most compressed state. Such a configuration would not be detectable by comparing the mean vertical velocities above and below the mid-plane, because the breathing oscillation would be at a turning point between contraction and expansion. Mass estimates close to the mid-plane and under the steady state assumption would be biased towards more massive results, as the stellar disk would have a smaller scale height and larger vertical velocities. In order to explain our results, such a breathing mode would have to be large enough for the stellar number density in the mid-plane to oscillate with a relative amplitude of \u223c5%. Monari et al. (2016 see for example Fig. 4) have shown that a local breathing mode could be created by a spiral arm that passes close enough to the Sun, inducing a net motion away from (towards) the Galactic mid-plane for stars on the outside (inside) of the spiral arm. Furthermore, such a close passage of a spiral arm is indicated by some dynamical models of the horizontal motions within the Galactic disk (for example Siebert et al. 2012). The dynamics of the Galactic disk are very complex and probably contains a combination of breathing and bending modes created by the last impact of the Sagittarius dwarf galaxy (for example Laporte et al. 2019); it is very plausible to think that the famous phase-space spirals in the z\u2013w plane are a local manifestation of these larger scale perturbations. The phase-space spiral structure could have an effect on our analysis, especially for the stellar samples with higher Zlim.. On a larger scale, the bending and breathing modes of the Galaxy have been shown to affect dynamical mass measurements of the Galactic disk, especially at greater heights (Banik et al. 2017; Haines et al. 2019).","Citation Text":["Widrow et al. 2014"],"Functions Text":["In order to produce a steep gravitational potential at low |z|, there could be a breathing mode in the stellar disk","which is currently in its most compressed state."],"Functions Label":["Uses","Uses"],"Citation Start End":[[927,945]],"Functions Start End":[[810,925],[967,1015]]}
{"Identifier":"2016AandA...587A.159G__Tian_et_al._2014_Instance_1","Paragraph":"One has to be sure to rule out cases where inorganic chemistry can mimic the presence of life (\u201cfalse positives\u201d). Potential abiotic ozone production on Venus- and Mars-like planets has been discussed by Schindler & Kasting (2000, and references therein). While this is based on photolysis of e.g., CO2 and H2O and is thus limited in extent, a sustainable production of abiotic O3 which could build up to a detectable level has been suggested by Domagal-Goldman & Meadows (2010) for a planet within the habitable zone of AD Leonis with a specific atmospheric composition. Indeed, other studies confirm that abiotic buildup of ozone is possible (e.g., Hu et al. 2012; Tian et al. 2014); however, detectable levels are unlikely if liquid water is abundant, as e.g. rainout of oxidized species would keep atmospheric O2 and O3 low (Segura et al. 2007), unless the CO2 concentration is high and both H2 and CH4 emissions are low (Hu et al. 2012). False-positive detection of molecules such as CH4 and O3 is discussed by von Paris et al. (2011). Seager et al. (2013) present a biosignature gas classification. Since abiotic processes cannot be ruled out for individual molecules (e.g. for O3), searches for biosignature molecules should search for multiple biosignature species simultaneously. It has been suggested that the simultaneous presence of O2 and CH4 can be used as an indication for life (Sagan et al. 1993, and references therein). Similarly, Selsis et al. (2002) suggest a so-called \u201ctriple signature\u201d, where the combined detection of O3, CO2 and H2O would indicate biological activity. Domagal-Goldman & Meadows (2010) suggest to simultaneously search for the signature of O2, CH4, and C2H6. Of course, care has to be taken to avoid combinations of biosignature molecules which can be generated abiotically together (see e.g. Tian et al. 2014). The detectability of biosignature molecules is discussed, e.g. by von Paris et al. (2011) and Hedelt et al. (2013). In particular, the simulation of the instrumental response to simulated spectra for currently planned or proposed exoplanet characterization missions has shown that the amount of information the retrieval process can provide on the atmospheric composition may not be sufficient (von Paris et al. 2013). ","Citation Text":["Tian et al. 2014"],"Functions Text":["Indeed, other studies confirm that abiotic buildup of ozone is possible (e.g.,"],"Functions Label":["Similarities"],"Citation Start End":[[667,683]],"Functions Start End":[[572,650]]}
{"Identifier":"2017MNRAS.469S..39F__Youdin_&_Goodman_2005_Instance_1","Paragraph":"At the disc age t \u2248 1.5 Myr, the decay of 26Al has sufficiently advanced to not melt forming planetesimals and dust from the inner disc has enriched the outer disc (Ciesla 2011). At rh = 30 au, we get \u03c4 \u2248 0.1 Myr and t0 \u2248 0.5 Myr because \u03c1(1\u2009au)\/\u03c1(30\u2009au) = 170 and 1800 in the MMSN and Nice models, respectively (Zsom et al. 2010). The number of fractals is still larger than that of pebbles of similar sizes:\n(3)\r\n\\begin{equation}\r\nN_{\\rm F} \/ N_0 \\gtrsim \\exp [-(t - t_0)\/\\tau ] \\gtrsim N_{\\rm P} \/ N_0 \\approx \\phi _{\\rm F} \/ \\phi _{\\rm P}\r\n\\end{equation}\r\nwith \u03d5F \u2248 10\u22125, the others being restructured into cm-sized pebbles (Weidling et al. 2009; Zsom et al. 2010). Any further pebble growth would require so longer times at so high relative speeds (Davidsson et al. 2016) to destroy or convert into pebbles all the remaining fractals. The input of dust from the inner disc triggers a streaming instability in the outer disc (Youdin & Goodman 2005; Johansen et al. 2007), lasting \u22480.1 Myr and concentrating pebbles and fractals sticking to pebbles into filaments where the comet nucleus collapses by gravity at speeds 1 m s\u22121 into a randomly packed (Onoda & Liniger 1990; Song, Wang & Makse 2008) aggregate of cm-sized pebbles, with the fractals stored in the voids among the pebbles (Fulle et al. 2016b). The voids fill (40\u2009\u00b1\u20095) per cent of the nucleus volume (Onoda & Liniger 1990; Song, Wang & Makse 2008). This scenario is confirmed by the pebbles observed in the pristine terrain at the final Philae landing site, which have a narrow differential size distribution (Fig. 2) peaked at diameters of $6^{+4}_{-2}$ mm (Poulet et al. 2016). For polydisperse spheres, the void volume occupies a lower percentage of the total, e.g. 32 per cent for a log-normal size distribution with a standard deviation of 0.3 (Baranau & Tallarek 2014). These facts allow us to model the 67P nucleus as composed of pebbles of the same size, with a void volume filling (37\u2009\u00b1\u20095) per cent of the total nucleus volume. A random packing of spheres does not result in the densest possible configuration, i.e. small pebbles will not always occupy available void spaces, because access to these voids is not necessarily available during the formation of the cometary nucleus.","Citation Text":["Youdin & Goodman 2005"],"Functions Text":["The input of dust from the inner disc triggers a streaming instability in the outer disc","lasting \u22480.1 Myr and concentrating pebbles and fractals sticking to pebbles into filaments where the comet nucleus collapses by gravity at speeds 1 m s\u22121 into a randomly packed","aggregate of cm-sized pebbles, with the fractals stored in the voids among the pebbles"],"Functions Label":["Uses","Uses","Uses"],"Citation Start End":[[930,951]],"Functions Start End":[[840,928],[976,1152],[1201,1287]]}
{"Identifier":"2018MNRAS.475.1160H__Zschaechner_et_al._2016_Instance_1","Paragraph":"Given their potentially profound impact on both the stellar and gas properties, galaxy mergers are often invoked to explain the presence of metals and low ionization species in the CGM at large impact parameters (e.g. Farina et al. 2013, 2014; Johnson, Chen & Mulchaey 2015b). Indeed, merger-induced tidal torques can give rise to observed stellar tidal features extending to many tens of kpc (e.g. Hern\u00e1ndez-Toledo et al. 2006; Patton et al. 2011, 2013; Casteels et al. 2014). Although some gas asymmetries coexist with asymmetries in the stellar profiles of interacting galaxies, tidally induced asymmetries can persist even longer in H\u2009i gas (e.g. Lelli, Verheijen & Fraternali 2014a,b; Scott et al. 2014). Such tidal debris can potentially diffuse into the CGM contributing a significant gas and metal mass. In addition to morphological disturbances, galaxy mergers can trigger vigorous galactic outflows associated with feedback from the enhanced star formation (e.g. Martin 2005; Rupke, Veilleux & Sanders 2005a; Strickland & Heckman 2009; Hayward & Hopkins 2017) and AGN activity (e.g. Rupke, Veilleux & Sanders 2005b; Veilleux et al. 2013; Zschaechner et al. 2016; Woo, Son & Bae 2017) which can populate the CGM with metals while giving rise to multiphase absorbers (e.g. Borthakur et al. 2013; Bird et al. 2015; Heckman & Borthakur 2016; Bordoloi et al. 2017; Heckman et al. 2017). Besides, merger-induced shocks and feedback can increase the CGM's internal energy; numerical simulations of galaxy mergers show that CGM gas can be significantly heated (to X-ray emitting temperatures) through shocks and feedback processes (e.g. Cox et al. 2004, 2006b; Sinha & Holley-Bockelmann 2009). Furthermore, studies of galaxies in dense environments (higher merger probability) show different CGM properties when compared to a matched isolated galaxy sample: The CGM of galaxies in groups show distinct kinematics (Pointon et al. 2017) and ionic covering fractions (Johnson, Chen & Mulchaey 2015a; Burchett et al. 2016). Despite the possibly significant influence of galaxy mergers on the CGM, the details of the interplay between galaxy mergers and the CGM remain relatively unexplored and we are currently lacking clear and quantitative predictions of how the CGM will be affected during the merger process. Current observations of the impact of galaxy mergers on the CGM are limited to a few case studies (e.g. Keeney et al. 2011; Johnson et al. 2014). A systematic survey targeting the CGM of kinematic galaxy pairs (i.e. COS-Pairs: Bordoloi et al. in preparation) is needed to place observational constraints on the effect of mergers on the CGM.","Citation Text":["Zschaechner et al. 2016"],"Functions Text":["In addition to morphological disturbances, galaxy mergers can trigger vigorous galactic outflows associated with feedback from the enhanced star formation (e.g."],"Functions Label":["Background"],"Citation Start End":[[1148,1171]],"Functions Start End":[[812,972]]}
{"Identifier":"2016AandA...585A..48G__Mostardi_et_al._(2013)_Instance_1","Paragraph":"On the same SSA22 field, Nestor et al. (2011) and Nestor et al. (2013) show nine LBGs and 20 Lyman-\u03b1 emitters (LAEs) with LyC detection out of a sample of 41 LBGs and 91 LAEs (all spectroscopically confirmed). They started from a different narrowband image centred at ~3640 \u00c5, which is deeper than that used by Iwata et al. (2009), at ~3590 \u00c5. A careful analysis of their LBG detections, however, shows that even in this case the LyC emission for many z ~ 3 sources is offset by 0.4\u20131.0 arcsec. The observed ratio between the 900 and 1500 \u00c5 rest-frame emission is difficult to reconcile with that expected by standard stellar populations (Vanzella et al. 2012a). The HST images in I814W filter available for a few of them show the presence of clearly separated galaxies, sometimes fainter at 1500 \u00c5 than in the ionizing continuum (i.e. their C16 object), with a resulting escape fraction well exceeding 1000% if estimated in the LyC position. For the majority of them, no HST imaging is available but even ground-based images often show the presence of slightly offset emission in LyC, w.r.t. the non-ionizing continuum. Similar conclusions can be reached for the Mostardi et al. (2013) sample where they adopt the same analysis as in Nestor et al. (2013). At z ~ 2.8, they found four LyC emitters out of 49 LBG galaxies and seven LyC emitters out of 91 LAEs. In this case the lack of high spatial resolution data from HST for the majority of the sample prevents any detailed analysis about possible contamination by interlopers\/foregrounds. These conclusions have been strengthened by recent observations by Siana et al. (2015), who found no convincing detection in their deep HST imaging with WFC3-UVIS of five LyC emitters extracted from the sample of Nestor et al. (2011), or by Mostardi et al. (2015), who only found one robust LyC emitter after a reanalysis of a sample of 16 galaxies by Mostardi et al. (2013). More interestingly, the only robust candidate LyC emitter by Mostardi et al. (2015), the galaxy MD5b, has an observed ratio FUV\/FLyC = 4.0 \u00b1 2.0, equivalent to a relative escape fraction of 75%, when assuming complete transmission of the IGM. Instead, if a mean value of \u27e8 exp(\u2212\u03c4IGM) \u27e9  = 0.4 at z ~ 3.1 is adopted, following Inoue et al. (2014), the relative escape fraction of MD5b turns out to be 188%. Imposing the constraint of a physical value for the relative escape fraction of \\hbox{$f^{\\rm rel}_{\\rm esc}<100\\%$}fescrel100%, the LoS of MD5b must be very transparent, exp(\u2212\u03c4IGM) > 0.75, which corresponds to a probability 10-4, according to Inoue et al. (2014). This could be an indication that this galaxy is also a low-z contaminant, similar to the other cases studied by Mostardi et al. (2015). ","Citation Text":["Mostardi et al. (2013)"],"Functions Text":["Similar conclusions can be reached for the","sample where they adopt the same analysis as in Nestor et al. (2013)."],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[1164,1186]],"Functions Start End":[[1121,1163],[1187,1256]]}
{"Identifier":"2022MNRAS.513..232N__Haywood_et_al._2013_Instance_1","Paragraph":"There are a plethora of data available in the form of spectra, astrometric, and photometric information, as well as multiwavelength maps with the advent of large-scale spectroscopic (Apache Point Observatory Galactic Evolution Experiment\/APOGEE: Eisenstein et al. 2011, RAdial Velocity Experiment\/RAVE: Steinmetz et al. 2006, Gaia-ESO: Gilmore et al. 2012, Large Sky Area Multi-Object Fiber Spectroscopic Telescope\/LAMOST: Cui et al. 2012, Galactic Archaeology with HERMES\/GALAH: De Silva et al. 2015, Abundances and Radial velocity Galactic Origins Survey\/ARGOS: Ness et al. 2012), astrometric (Hipparcos: Perryman et al. 1997, Gaia: Gaia Collaboration 2016), and photometric surveys (Two-Micron All Sky Survey\/2MASS: Skrutskie et al. 2006, Sloan Digital Sky Survey\/SDSS: Stoughton et al. 2002, Vista Variables in the V\u00eda L\u00e1ctea\/VVV: Minniti et al. 2010, the SkyMapper Southern Survey : Wolf et al. 2018). These surveys have enabled the chemo-dynamic characterization of stellar populations in the Milky way that constitute different Milky Way components like thin disc, thick disc, halo, bulge, etc. For example, star count observations in the solar neighbourhood (Yoshii 1982; Gilmore & Reid 1983) led to the discovery of the thick disc, followed by its characterization as the old \u03b1-enhanced population in the double sequence exhibited by the solar neighbourhood stars in the [\u03b1\/Fe] versus [Fe\/H] plane (Fuhrmann 1998; Bensby, Feltzing & Lundstr\u00f6m 2003; Reddy, Lambert & Allende Prieto 2006; Adibekyan et al. 2012; Haywood et al. 2013). At present, data from large-scale spectroscopic surveys (Anders et al. 2014; Hayden et al. 2015; Weinberg et al. 2019) have led to the discovery of this trend at different galactocentric radius, R, and average height, |Z|, across the Galaxy shedding light on the disc formation and evolution scenarios. In addition, many age determination methods have been developed that uses these survey data to provide valuable information about the star formation histories and age metallicity relation of disc stellar populations (Casagrande et al. 2011; Bedell et al. 2018; Lin et al. 2020; Nissen et al. 2020). Secular processes such as radial migration (Sellwood & Binney 2002; Sch\u00f6nrich & Binney 2009; Minchev & Famaey 2010), which leads to the mixing of stars across the Galaxy, are also being explored using a combination of accurate phase space information from Gaia (Gaia Collaboration 2018) and chemistry and age information of stars from large-scale spectroscopic surveys (Buder et al. 2019). The discovery of streams and dynamically different stellar populations in the Milky Way halo, considered to be the result of past accretion\/merger events (Belokurov et al. 2018; Helmi et al. 2018; Ibata, Malhan & Martin 2019; Myeong et al. 2019) using the Gaia data and their further exploration with chemistry from large-scale spectroscopic surveys (Buder et al. in preparation) is another example. Multiple components in the Bulge metallicity distribution function discovered by multiple individual and large-scale spectroscopic observations, are being studied in detail to understand the origin of the Bulge and its connection with the Milky Way bar and Galaxy evolution (Ness et al. 2013; Rojas-Arriagada et al. 2017, 2020). There are many upcoming surveys [4-metre Multi-Object Spectroscopic Telescope\/4MOST: de Jong et al. (2019), Sloan Digital Sky Survey\/SDSS-V: Kollmeier et al. (2017), WEAVE: Dalton et al. (2018)] that will further improve our understanding of the formation and evolution of the Milky Way and its components.","Citation Text":["Haywood et al. 2013"],"Functions Text":["For example, star count observations in the solar neighbourhood","followed by its characterization as the old \u03b1-enhanced population in the double sequence exhibited by the solar neighbourhood stars in the [\u03b1\/Fe] versus [Fe\/H] plane"],"Functions Label":["Background","Background"],"Citation Start End":[[1519,1538]],"Functions Start End":[[1102,1165],[1241,1406]]}
{"Identifier":"2017ApJ...840...98J__Nataf_et_al._2010_Instance_1","Paragraph":"Figure 1 shows our synthetic color\u2013magnitude diagrams (CMDs) for the two RCs at four different metallicity regimes, which are almost identical to those presented in Figure 1 of Paper I, but here (\n\n\n\n\n\n) CMDs are added in the right panels. We recall that in our models, the fRC and bRC are produced by G1 and G2, respectively, where G1 follows the standard helium-enrichment parameter (i.e., \n\n\n\n\n\n), while G2 is substantially enhanced in helium abundance (Y = 0.406). Following Paper I, we have also adopted a 0.2 dex difference in metallicity between G2 and G1, and assumed 12 and 10 Gyrs for the ages of G1 and G2, respectively, with the same population ratio for the two RCs (Nataf et al. 2010; Paper I and references therein). It is evident from these models that in the metal-rich population like the bulge, highly helium-enhanced stars \n\n\n\n\n\n are not placed on the very blue HB as in the metal-poor GCs, but are instead placed on the bRC. We refer the reader to Section 2 of Paper I for a detailed description of the RC features of our models at four different metallicity regimes in Figure 1. Here, we note from the (\n\n\n\n\n\n) CMDs that the variation of the overall RC features on metallicity is similar to that in (\n\n\n\n\n\n) CMDs. Panel (f) in (\n\n\n\n\n\n) CMD can also naturally reproduce the observed double RC feature, i.e., \u223c0.5 mag difference with almost negligible color difference between the two RCs, as is the case in panel (b) in (\n\n\n\n\n\n) CMD. Hereafter, we refer to these models in panels (b) and (f) constructed at \n\n\n\n\n\n as \u201c\n\n\n\n\n\n,\u201d to distinguish them from the models with different input parameters presented below. The CMDs for these reference models, down to the MS luminosity level, are further presented in Figure 2. This figure confirms that, in the position and width of the lower red-giant-branch (RGB) and main-sequence turn-off (MSTO), our models are not inconsistent with the observed CMDs of bulge population by Clarkson et al. (2008, 2011) and Brown et al. (2010).","Citation Text":["Nataf et al. 2010"],"Functions Text":["Following Paper I, we have also adopted a 0.2 dex difference in metallicity between G2 and G1, and assumed 12 and 10 Gyrs for the ages of G1 and G2, respectively, with the same population ratio for the two RCs"],"Functions Label":["Uses"],"Citation Start End":[[680,697]],"Functions Start End":[[469,678]]}
{"Identifier":"2015ApJ...799...42D__Skemer_et_al._2014_Instance_1","Paragraph":"The Large Binocular Telescope (LBT) consists of two 8.4\u00c2 m aperture optical telescopes on a single ALT-AZ mount installed on Mount Graham in southeastern Arizona (at an elevation of 3192\u00c2 m) and operated by an international collaboration among institutions in the United States, Italy, and Germany (Hill et al. 2014; Veillet et al. 2014). Both telescopes are equipped with deformable secondary mirrors which are driven with the LBT's adaptive optics (AO) system to correct atmospheric turbulence at 1\u00c2 kHz (Esposito et al. 2010; Bailey et al. 2014). Each deformable mirror uses 672 actuators that routinely correct 400 modes and provide Strehl ratios exceeding 80%, 95%, and 99% at 1.6\u00e2\u0080\u0089\u00ce\u00bcm, 3.8\u00e2\u0080\u0089\u00ce\u00bcm, and 10\u00e2\u0080\u0089\u00ce\u00bcm, respectively (Esposito et al. 2012; Skemer et al. 2014). The LBTI is an interferometric instrument designed to coherently combine the beams from the two 8.4\u00c2 m primary mirrors of the LBT for high-angular resolution imaging at infrared wavelengths (1.5\u00e2\u0080\u009313\u00e2\u0080\u0089\u00ce\u00bcm; Hinz et al. 2012). It is developed and operated by the University of Arizona and based on the heritage of the Bracewell Infrared Nulling Cryostat on the MMT Hinz et al. (2000). The overall LBTI system architecture and performance will be presented in full detail in a forthcoming publication (P. M. Hinz et al., in preparation). In brief, the LBTI consists of a universal beam combiner (UBC) located at the bent center Gregorian focal station and a cryogenic Nulling Infrared Camera (NIC). The UBC provides a combined focal plane for the two LBT apertures while the precise overlapping of the beams is done in the NIC cryostat. Nulling interferometry, a technique proposed 36\u00c2 yr ago to image extra-solar planets (Bracewell 1978), is used to suppress the stellar light and improve the dynamic range of the observations. The basic principle is to combine the beams in phase opposition in order to strongly reduce the on-axis stellar light while transmitting the flux of off-axis sources located at angular spacings which are odd multiples of 0.5\u00ce\u00bb\/B (where B = 14.4\u00c2 m is the distance between the telescope centers and \u00ce\u00bb is the wavelength of observation). Beam combination is done in the pupil plane on a 50\/50 beamsplitter which can be translated to equalize the pathlengths between the two sides of the interferometer. One output of the interferometer is reflected on a short-pass dichroic and focused on the Nulling Optimized Mid-Infrared Camera (NOMIC) (Hoffmann et al. 2014). NOMIC uses a 1024 \u00c3\u0097 1024 Raytheon Aquarius detector split into two columns of eight contiguous channels. The optics provides a field of view (FOV) of 12\u00c2 arcsec with a plate-scale of 0.018\u00c2 arcsec. Tip\/tilt and phase variations between the LBT apertures are measured using a fast-readout (1\u00e2\u0080\u0089Hz) K-band PICNIC detector (PHASECam) which receives the near-infrared light from both outputs of the interferometer. Closed-loop correction uses a fast pathlength corrector installed in the UBC (see more details in Defr\u00c3\u00a8re et al. 2014).","Citation Text":["Skemer et al. 2014"],"Functions Text":["Each deformable mirror uses 672 actuators that routinely correct 400 modes and provide Strehl ratios exceeding 80%, 95%, and 99% at 1.6\u00e2\u0080\u0089\u00ce\u00bcm, 3.8\u00e2\u0080\u0089\u00ce\u00bcm, and 10\u00e2\u0080\u0089\u00ce\u00bcm, respectively"],"Functions Label":["Background"],"Citation Start End":[[754,772]],"Functions Start End":[[550,730]]}
{"Identifier":"2021AandA...652A..30S__Santini_et_al._2019_Instance_1","Paragraph":"This paper is the fourth in a series. In our first work (Merlin et al. 2018) we presented an accurate and conservative technique to single out passive galaxies at high redshift by means of spectral energy distribution (SED) fitting with a probabilistic approach. We selected 30 z\u2004>\u20043 candidates in the GOODS-S field. Passive galaxy candidates, while being relatively easy to select from photometric surveys once the technique is established, need to be confirmed by other means. This is usually achieved through spectroscopic observations. However, spectroscopy becomes particularly difficult and time consuming at z\u2004>\u20043, where only few candidates have been confirmed so far (Glazebrook et al. 2017; Schreiber et al. 2018a; Tanaka et al. 2019; Valentino et al. 2020; Forrest et al. 2020a,b; Saracco et al. 2020; D\u2019Eugenio et al. 2020). In our second work (Santini et al. 2019, S19 hereafter), we used a complementary approach and looked for evidence of the lack of star formation as seen in the sub-millimetre regime to confirm the passive classification of the high-z candidates selected in GOODS-S. At that time, we could confirm 35% of the targets on an individual basis adopting conservative assumptions, and we validated the sample as a whole in a statistical sense. In our third work (Merlin et al. 2019, M19 hereafter), we extended the search for passive galaxies to the entire CANDELS sample and selected 102 z\u2004>\u20043 candidates over the five fields. In the present work, we first confirm the passive nature of these candidates, adopting the method presented in S19 and taking advantage of the richer ALMA archive, which includes observations that were still proprietary at the time of our previous work. We then analyse the emergence and mass growth of this peculiar class of galaxies by means of two powerful statistical tools: the stellar mass function (SMF) and the stellar mass density (SMD). Very few studies so far have pushed the analysis of the SMF of passive galaxies at z\u2004>\u20043 (Muzzin et al. 2013; Davidzon et al. 2017; Ichikawa & Matsuoka 2017; Girelli et al. 2019) because of the difficulty in assembling statistically meaningful samples of candidates at such high redshift.","Citation Text":["Santini et al. 2019"],"Functions Text":["In our second work","S19 hereafter), we used a complementary approach and looked for evidence of the lack of star formation as seen in the sub-millimetre regime to confirm the passive classification of the high-z candidates selected in GOODS-S."],"Functions Label":["Background","Background"],"Citation Start End":[[856,875]],"Functions Start End":[[836,854],[877,1100]]}
{"Identifier":"2015AandA...584A.103S__Chamel_et_al._2007_Instance_1","Paragraph":"Before leaving this section, in Fig. 8 we display the spatial dependence of the self-consistent neutron and proton density profiles for the optimal solutions in spherical WS cells with average baryon densities nb = 0.0475 fm-3, 0.065 fm-3, and 0.076 fm-3. It is observed that in denser matter the size of the WS cell decreases, as we discussed previously, and that the amount of free neutrons in the gas increases, as expected. It can be seen that the nuclear surface is progressively washed out with increasing average baryon density as the nucleon distributions become more uniform. At high nb the density profile inside the WS cell extends towards the edge of the cell, pointing out that the WS approximation may be close to its limits of validity (Negele & Vautherin 1973; Chamel et al. 2007; Baldo et al. 2007; Pastore et al. 2011; G\u00f6gelein & M\u00fcther 2007; Newton & Stone 2009). Although the proton number Z is similar for the three average baryon densities of Fig. 8, the local distribution of the protons is very different in the three cases. In Fig. 8c the proton density profile extends more than 3 fm farther from the origin than in Fig. 8a, while the central value of the proton density has decreased by more than a factor 2, hinting at the fact that the neutrons have a strong drag effect on the protons. Figure 9 presents the nucleon density profiles obtained for cylindrical and planar geometries at the same average density nb = 0.076 fm-3 as in Fig. 8c. From Figs. 8c (droplets),  9a (rods), and 9b (slabs) we see that the size of the WS cells decreases with decreasing dimensionality, i.e. Rc,droplet>Rc,rod>Rc,slab. At high average densities near the crust-core transition, nucleons inside the WS cell can arrange themselves in such a way that the region of higher density is concentrated at the edge of the cell, leaving the uniform region of lower density in the inner part of the cell. This distribution of nucleons corresponds to the cylindrical tube and spherical bubble configurations. In Figs. 9c and d, we plot the neutron and proton density profiles of the optimal solution for tubes and bubbles at nb = 0.076 fm-3. At equal average density, the size of the cells containing tubes and bubbles is larger than the size of the cells accommodating rods and droplets, respectively, as can be appreciated by comparing Fig. 9a for rods with Fig. 9c for tubes, and Fig. 8c for droplets with Fig. 9d for bubbles. As a consequence of this fact and of the effectively larger value of the integration factors 2\u03c0r and 4\u03c0r2 when the densities are accumulated near the edge of the cell, the total number of nucleons and the atomic number in the tube and bubble cells is about 1.5\u22122 times larger than in their rod and droplet counterparts. The proton fraction xp = Z\/A is, however, practically the same for all geometries. ","Citation Text":["Chamel et al. 2007"],"Functions Text":["At high nb the density profile inside the WS cell extends towards the edge of the cell, pointing out that the WS approximation may be close to its limits of validity"],"Functions Label":["Uses"],"Citation Start End":[[777,795]],"Functions Start End":[[585,750]]}
{"Identifier":"2022AandA...660A..56A__Myers_2009_Instance_1","Paragraph":"Recent theoretical models clearly demonstrate the role of expanding H\u202fI shells for the formation of molecular clouds (e.g., Hennebelle et al. 2008; Heitsch et al. 2009; Inoue & Inutsuka 2009). In particular, these models highlight the importance of multiple compressions for the formation of magnetized filamentary molecular clouds (e.g., Inutsuka et al. 2015; Iwasaki et al. 2019). The typical timescale of such compressions is estimated in the Galactic disk to be on average ~1 Myr (McKee & Ostriker 1977; Inutsuka et al. 2015). Thus, the formation of molecular clouds may last from a few million years to ~10 Myr or more (see, e.g., Kobayashi et al. 2017). These successive compressions may continuously alter the density, velocity, and magnetic field structures of the clouds, and also have a strong impact on the formation of new generations of filaments and consequently that of stars. While the first generation of stars form and impact their (local) surroundings (through outflow, jets, winds, and ionizing radiation), new cold matter is continuously assembled replenishing the sites of star formation (i.e., filaments and hubs). This matter replenishment may be channeled from within the cloud itself through molecular filaments toward dense ridge-like main filaments (Schneider et al. 2010; Palmeirim et al. 2013) or toward hubs (Myers 2009; Peretto et al. 2013, 2014; Trevi\u00f1o-Morales et al. 2019). Matter can also be brought into the system (the cloud) by a new event of external collision (e.g., Fukui et al. 2018b). Arzoumanian et al. (2018) identified, in position-velocity (PV) diagrams, extended structures with mean line-of-sight (LOS) velocities offset with respect to, and connected to the velocity of, a filament. They suggested a multi-interaction scenario where sheet-like extended structures interact, in space and time, with a star-forming filament, and are responsible for its compression or disruption. Arzoumanian et al. (2018) also identified a bent velocity structure in the PV space. They showed that such a V- or \u039b-shaped velocity structure can result from the filament formation process by shock compression as proposed by the theoretical model of Inoue et al. (2018). In this model, a filament is formed due to convergence of a flow of matter generated by the bending of the ambient magnetic field structure induced by an interstellar shock compression (see also Inoue & Fukui 2013; Vaidya et al. 2013). This velocity structure has also been observed toward another filament, the Musca filament (Bonne et al. 2020). More recently, in a theoretical study, Abe et al. (2021) proposed a classification of filament formation mechanisms resulting from the variation in the relative importance between the shock velocity, the turbulence, and the magnetic field strength (see also the theoretical study by Chen et al. 2020a).","Citation Text":["Myers 2009"],"Functions Text":["This matter replenishment may be channeled from within the cloud itself through molecular filaments","or toward hubs"],"Functions Label":["Background","Background"],"Citation Start End":[[1340,1350]],"Functions Start End":[[1138,1237],[1324,1338]]}
{"Identifier":"2022MNRAS.510.5676I__Faisst_et_al._2017_Instance_1","Paragraph":"The above finds an explanation within the context of the two-phases formation scenario of ETGs (Oser et al. 2010). According to this scenario, in the early (z \u2273 2) dissipative phase driven by wet major mergers (e.g. Hopkins et al. 2008) and\/or violent disc instability (e.g. Barro et al. 2013; Dekel & Burkert 2014; Zolotov et al. 2015; Tacchella et al. 2016), the proto-ellipticals are compact objects that form stars in situ in a very intense regime until this process is stopped abruptly by gas shock heated in massive haloes, and\/or rapid gas exhaustion and Supernova\/AGN feedback. As the compact quiescent galaxy ages and reddens (red nugget), during the second non-dissipative phase (z \u2272 1\u22122), more mass is assembled by the accretion of ex situ stars through merging with other galaxies, which are also mostly quiescent. Dry minor or intermediate mergers contribute little to the mass growth but promote substantial size growth (the accreted stars tend to be deposited mostly in the external regions), making spheroidal galaxies less compact and likely also less concentrated (Bezanson et al. 2009; Naab, Johansson & Ostriker 2009; Oser et al. 2010; Trujillo, Ferreras & de La Rosa 2011; Bluck et al. 2012; Johansson, Naab & Ostriker 2012; Hilz, Naab & Ostriker 2013; Shankar et al. 2013; van Dokkum et al. 2015; Wellons et al. 2016; Faisst et al. 2017; Furlong et al. 2017; Hill et al. 2017; Genel et al. 2018, among others). Dry mergers can also be major, though they are less frequent and relevant only to the most massive galaxies (e.g. Bundy et al. 2009; L\u00f3pez-Sanjuan et al. 2012; Rodr\u00edguez-Puebla et al. 2017). These mergers increase substantially the mass, while the size increases approximately proportional to the mass increase (e.g. Nipoti, Londrillo & Ciotti 2003; Johansson et al. 2012; Nipoti et al. 2012; Hilz et al. 2013) in a such way that the shift in the mass\u2013size relation is small, affecting in a lesser degree the compactness and concentration of the merged galaxies. Although the second phase of ETG formation, driven by dry mergers, is a reliable explanation for the growth in size and the puffing-up of massive ETGs, there is still an intense debate as to whether or not this mechanism is enough to describe observational inferences (e.g. L\u00f3pez-Sanjuan et al. 2012; Newman et al. 2012; Nipoti et al. 2012; Sonnenfeld, Nipoti & Treu 2014; Man, Zirm & Toft 2016; Frigo & Balcells 2017, see for a discussion Zanisi et al. 2021 and more references therein). An alternative or complementary mechanism suggested for the apparent strong growth in size of massive ETGs is quasar feedback, which removes huge amounts of cold gas from the central regions, inducing an expansion of the stellar distribution (Fan et al. 2008, but see Trujillo et al. 2011).","Citation Text":["Faisst et al. 2017"],"Functions Text":["Dry minor or intermediate mergers contribute little to the mass growth but promote substantial size growth (the accreted stars tend to be deposited mostly in the external regions), making spheroidal galaxies less compact and likely also less concentrated"],"Functions Label":["Background"],"Citation Start End":[[1340,1358]],"Functions Start End":[[827,1081]]}
{"Identifier":"2021MNRAS.500.4042S__Snellen_et_al._2010_Instance_1","Paragraph":"To highlight the evidence for CO2 in the transmission spectrum, in Fig. 18 we show three atmo model atmospheres: our best-fitting model from the free-chemistry retrieval; a model with all of the same parameters as the best fit, except the CO2 abundance, which is set to zero; and a third model with both CO and CO2 abundances set to zero. The strong absorption feature centred on the 4.5 $\\, \\mu$m Spitzer channel disappears in the latter two models. With only Spitzer photometry, however, the contribution of CO to the 4.5 $\\, \\mu$m point complicates the interpretation of the C\/O ratio. Theoretical models have found that CO should be the dominant carbon-bearing molecule for hydrogen-dominated atmospheres above 1000 K (e.g. Lodders & Fegley 2002; Heng & Lyons 2016), and CO has been detected at high resolution in hot Jupiter atmospheres (e.g. Snellen et al. 2010). In our equilibrium chemistry retrieval, CO is at least 100\u00d7 more abundant than CO2 (see Fig. 12). However, at 4.5 $\\, \\mu$m the CO2 opacity is much stronger and dominates over the CO contribution even though CO has much higher VMR concentrations. In the free-chemistry retrieval, the CO VMR is not constrained by the data \u2013 only an upper limit to the CO is found, as very high values affect the mean molecular weight, and the data are consistent with no CO contribution. The lack of a CO feature in the WFC3 data further pushes the free-chemistry retrieval to prefer CO2 over CO. With CO constrained through chemistry in one retrieval and unconstrained in the free case, the C\/O ratios obtained are vastly different. In the chemical-equilibrium case, a supersolar C\/O is found (see Fig. 13) while in the free case a subsolar C\/O ratio is found (Fig. 14). This finding highlights the extreme sensitivity and degeneracies of measuring the C\/O ratio with a free-chemistry retrieval model, as all major molecular species have to be well constrained by the data. For a hot Jupiter such as WASP-127b, we consider a scenario with all of the carbon found in CO2 and little to none in CO to be thermochemically implausible, as no obvious non-equilibrium mechanism would deplete CO by many orders of magnitude while enhancing CO2. This situation is unlike CH4, where dynamical mixing and vertical quenching can dramatically enhance CH4 (e.g. Cooper & Showman 2006; Moses et al. 2011; Tsai et al. 2017; Drummond et al. 2018a, b). With only one photometric data point at 4.5 $\\, \\mu$m, it is currently impossible to fully disentangle the contribution of both CO and CO2 in a model-independent way. Further transmission spectroscopy observations of WASP-127b at high resolution with the JWST will clarify which is the dominant carbon-bearing molecule in WASP-127b\u2019s atmosphere, and allow stronger constraints to be placed on its carbon-to-oxygen ratio.","Citation Text":["Snellen et al. 2010"],"Functions Text":["Theoretical models have found that CO should be the dominant carbon-bearing molecule for hydrogen-dominated atmospheres above 1000 K","and CO has been detected at high resolution in hot Jupiter atmospheres (e.g."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[848,867]],"Functions Start End":[[589,721],[771,847]]}
{"Identifier":"2021ApJ...922..268J__Bonoli_et_al._2021_Instance_1","Paragraph":"In the last decade, several projects have focused on collecting the deepest images of clusters and groups of galaxies. Among the most recent examples, we can mention the Frontier Fields initiative (Lotz et al. 2017), the Hyper Suprime-Cam Subaru Strategic Program (Aihara et al. 2018a, 2018b, 2019), the Fornax Deep Survey (Iodice et al. 2016), or the upcoming Vera C. Rubin Observatory (Ivezic et al. 2008). In addition, technical development offers a new and barely explored perspective: a panchromatic, more detailed view of the universe through the use of narrowband filter systems, such as those from the J-PLUS and J-PAS surveys (Benitez et al. 2014; Cenarro et al. 2019; Bonoli et al. 2021). In this work we intend to show how the combination of these two characteristics, high-quality data and information at different wavelengths, can provide strong clues on the clusters\u2019 formation and evolution through the analysis of the ICL. It is known that the ICL presents different morphologies, substructures, and ICL fractions depending on the wavelength used (Tang et al. 2018). Whereas the ICL fraction measured in single bands can be useful for establishing qualitative relations (e. g., evolution with redshift or halo mass), a panchromatic view of the ICL can give insights into its stellar populations or the dynamical stage of the cluster (e.g., Jim\u00e9nez-Teja et al. 2019). With this aim, we use the recently acquired data from the Reionization Lensing Cluster Survey (RELICS; Coe et al. 2019), a Hubble Space Telescope (HST) Treasure program that has observed 41 very massive, strong-lensing clusters in three optical (ACS) and four infrared (WFC3) bands. In this work we present the potential of the RELICS data for ICL purposes, focused on the cluster WHL J013719.8\u201308284 (WHL0137 hereafter). Future work will apply the CICLE algorithm, which is free of a priori assumptions, extensively to the whole RELICS sample with two main aims: (1) analyze the optical and infrared ICL fractions to extract information about the stellar populations of the diffuse light, and (2) infer the dynamical stage of the clusters.","Citation Text":["Bonoli et al. 2021"],"Functions Text":["In addition, technical development offers a new and barely explored perspective: a panchromatic, more detailed view of the universe through the use of narrowband filter systems, such as those from the J-PLUS and J-PAS surveys"],"Functions Label":["Background"],"Citation Start End":[[678,696]],"Functions Start End":[[409,634]]}
{"Identifier":"2020ApJ...891...28T__Antoja_et_al._2018_Instance_1","Paragraph":"With the release of Gaia DR2 (Gaia Collaboration et al. 2018; Katz et al. 2019), the proper motions of billions of stars are now available to the astronomical community. Combining with radial velocities from large spectroscopic surveys, like the Sloan Digital Sky Survey (Eisenstein et al. 2011; Blanton et al. 2017), Gaia-ESO survey (Gilmore et al. 2012; Randich et al. 2013), and LAMOST Galactic spectroscopic survey (Deng et al. 2012; Zhao et al. 2012), the wealth of 6D information of billions of stars have challenged and even overthrown many aspects of our understanding of the Milky Way (MW). The discovery of snail shells in the phase-space distribution of MW disk stars (Antoja et al. 2018) has inspired debates about their origin: whether they are generated by the passage of a dwarf galaxy (probably the Sagittarius dwarf galaxy) through the MW disk (Binney & Sch\u00f6nrich 2018) or are the echo of the MW bar buckling (Khoperskov et al. 2019). Meanwhile, major accretion events begin to unveil themselves when the stellar distribution in various energy\u2013momentum spaces (e.g., Myeong et al. 2019) is investigated. These major accretion events injected most of the materials from the progenitor dwarf galaxies into our MW, including globular clusters (GCs). As GCs are one of the oldest objects in our Galaxy, identifying and studying accreted GCs help us trace back the accretion history of our Galaxy. Though details of major accretion events, e.g., the number of accretion events and the GCs associated with each event, are still under debate (e.g., Helmi et al. 2018; Massari et al. 2019; Myeong et al. 2019), it is widely accepted that a substantial number of halo stars and GCs were accreted (e.g., Ostdiek et al. 2019). Along the same line, more and more substructures, e.g., stellar streams, are identified inside the MW (e.g., Malhan et al. 2018; Ibata et al. 2019b). An increasing number of stellar streams are suggested to be related to the debris of (inner-halo) GCs (e.g., Ibata et al. 2019a). Aside from these GC destruction events under the influence of Galactic potential, the dynamical relaxation of GCs (e.g., Weinberg 1994; Vesperini & Heggie 1997) also ejects member stars into the field. It would be of great interest to find such GC-ejected stars in order to estimate the mass loss from GCs to better understand the formation and evolution of our MW. To help achieve this goal, another characteristic of GCs is very helpful.","Citation Text":["Antoja et al. 2018"],"Functions Text":["The discovery of snail shells in the phase-space distribution of MW disk stars","has inspired debates about their origin:"],"Functions Label":["Background","Background"],"Citation Start End":[[680,698]],"Functions Start End":[[600,678],[700,740]]}
{"Identifier":"2016MNRAS.461.3982B__Scheeres_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":["Scheeres 2002"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[916,929]],"Functions Start End":[[635,873]]}
{"Identifier":"2021MNRAS.505.2111L__Melia_et_al._2017_Instance_1","Paragraph":"Recently, quasars observed with multiple measurements, another potential cosmological probe with a higher redshift range that reaches to z \u223c 5, is becoming popular to constrain cosmological models in the largely unexplored portion of redshift range from z \u223c 2 to z \u223c 5. A sample that contains 120 angular size measurements in intermediate-luminosity quasars from the very long baseline interferometry (VLBI) observations (Cao et al. 2017a,b) has become an effective standard ruler, which have been extensively applied to test cosmological models (Li et al. 2017; Melia et al. 2017; Qi et al. 2017; Zheng et al. 2017; Xu et al. 2018; Ryan, Chen & Ratra 2019), measuring the speed of light (Cao et al. 2017a, 2020a), exploring cosmic curvature at different redshifts (Cao et al. 2019; Qi et al. 2019), and the validity of cosmic distance duality relation (Zheng et al. 2020). Then, Risaliti & Lusso (2019) put forward a new compilation of quasars containing 1598 quasi-stellar object (QSO) X-ray and ultraviolet (UV) flux measurements in the redshift range of 0.036 \u2264 z \u2264 5.1003, which have been used to constrain cosmological models (Khadka & Ratra 2020b) and cosmic curvature at high redshifts (Liu et al. 2020a,c), as well as test the cosmic opacity (Geng et al. 2020; Liu et al. 2020b). Making use of this data to explore cosmological researches mainly depends on the empirical relationship between the X-ray and UV luminosity of these high-redshift quasars proposed by Avni & Tananbaum (1986), which leads to the Hubble diagram constructed by quasars (Risaliti & Lusso 2015, 2017; Lusso & Risaliti 2016; Bisogni, Risaliti & Lusso 2017). In general, the advantage of these two QSO measurements over other traditional cosmological probes is that QSO has a larger redshift range, which may be rewarding in exploring the behaviour of the non-standard cosmological models at high redshifts, providing an important supplement to other astrophysical observations and also demonstrating the ability of QSO as an additional cosmological probe (Zheng et al. 2021).","Citation Text":["Melia et al. 2017"],"Functions Text":["A sample that contains 120 angular size measurements in intermediate-luminosity quasars from the very long baseline interferometry (VLBI) observations","has become an effective standard ruler, which have been extensively applied to test cosmological models"],"Functions Label":["Background","Background"],"Citation Start End":[[563,580]],"Functions Start End":[[270,420],[442,545]]}
{"Identifier":"2017AandA...599A...8V__Carbone_&_Veltri_1990_Instance_1","Paragraph":"At scales comparable with dp, a variety of observations in the solar wind have suggested that fluctuations may consist primarily of kinetic Alfv\u00e9n waves (KAWs; Bale et al. 2005; Sahraoui et al. 2009). In the linear fluctuation terminology, KAWs are waves belonging to the Alfv\u00e9n branch, at wavevectors k almost perpendicular to the ambient magnetic field B0, with \\hbox{$k \\sim d_{\\rm p}^{-1}$}k~dp-1. A detailed discussion of the properties of KAWs can be found in Hollweg (1999; see also references therein for a more complete view of the subject), for example. In the last decades, KAWs have received considerable attention due to their possible role in a normal mode description of turbulence. Indeed, theoretical studies (e.g., Shebalin et al. 1983; Carbone & Veltri 1990; Oughton et al. 1994) have shown that the turbulent cascade in magnetized plasma tends to develop mainly in directions perpendicular to B0. Anisotropic spectra have been commonly observed in space plasmas, showing the presence of a significant population of quasi-perpendicular wavevectors (Matthaeus et al. 1986, 1990). The above considerations suggest that fluctuations with characteristics similar to KAWs are naturally generated by a turbulent cascade at scales close to dp. Many solar wind observational studies (Bale et al. 2005; Sahraoui et al. 2009; Podesta & TenBarge 2012; Salem et al. 2012; Chen et al. 2013; Kiyani et al. 2013), theoretical works (Howes et al. 2008a; Schekochihin et al. 2009; Sahraoui et al. 2012) as well as numerical simulations (Gary & Nishimura 2004; Howes et al. 2008b; TenBarge & Howes 2012) have suggested that fluctuations near the end of the magnetohydrodynamics inertial cascade range may consist primarily of KAWs, and that such fluctuations can play an important role in the dissipation of turbulent energy. Due to a non-vanishing parallel component of the electric field associated with KAWs, these waves have also been considered in the problem of particle acceleration (Voitenko & Goossens 2004; D\u00e9camp & Malara 2006). Particle acceleration in Alfv\u00e9n waves in a dispersive regime has been studied both in 2D (Tsiklauri et al. 2005; Tsiklauri 2011) and 3D (Tsiklauri 2012) configurations. Recently, Vasconez et al. (2014) have studied collisionless Landau damping and wave-particle resonant interactions in KAWs. ","Citation Text":["Carbone & Veltri 1990"],"Functions Text":["In the last decades, KAWs have received considerable attention due to their possible role in a normal mode description of turbulence. Indeed, theoretical studies (e.g.,","have shown that the turbulent cascade in magnetized plasma tends to develop mainly in directions perpendicular to B0."],"Functions Label":["Background","Background"],"Citation Start End":[[755,776]],"Functions Start End":[[564,732],[799,916]]}
{"Identifier":"2021MNRAS.500.2209Z__Miyama,_Narita_&_Hayashi_1987_Instance_1","Paragraph":"There is growing observational and numerical evidence that star-forming regions may be in a state of global gravitational contraction; see V\u00e1zquez-Semadeni et al. (2019). The supersonic collisions of flows of warm diffuse atomic gas simulated with both self-gravity and cooling exhibit the hierarchical collapse of the turbulent medium, as was shown by V\u00e1zquez-Semadeni et al. (2007) and Naranjo-Romero, V\u00e1zquez-Semadeni & Loughnane (2015) for example. This implies that gravitational instability (GI hereafter) manifests itself in a wide range of sufficiently large scales during the evolution of molecular clouds. Theoretical work has revealed that flattened dense structures form as a result of large collisions of diffuse matter. Later on, they give birth to filaments, which then fragment into multiple cores. This scenario is provided by the dynamical instability of self-gravitating layers, cylinders, and spheres, respectively. The linear stability analysis of these idealized configurations (e.g. Ledoux 1951; Chandrasekhar & Fermi 1953; Bonnor 1956; Elmegreen & Elmegreen 1978; Nagasawa 1987; Fiege & Pudritz 2000) as well as the corresponding non-linear solutions (e.g. Larson 1969; Penston 1969; Miyama, Narita & Hayashi 1987; Inutsuka & Miyama 1997; Masunaga & Inutsuka 2000, and many others) confirms this view. At the same time, as was noted by Larson (1985), the specific geometry of self-gravitating objects is not crucial for the instability condition, which does not differ much from the basic one derived for the unbounded uniform medium. In the latter case, the study of GI goes back to Jeans (1902), who established that plane-wave perturbations on such a background having finite pressure are heavy sound waves propagating at the subsonic velocity, which vanishes as the wavelength approaches the value now referred to as the Jeans length. Perturbations with scale larger than the Jeans length are the growing and damping static waves. Thus, the critical scales for GI of realistic configurations mentioned above are always similar to the Jeans scale, which includes typical speed of sound and density chosen appropriately for the corresponding configuration. However, the most unstable scale for realistic configurations has a finite value in contrast to the Jeans result, when the largest growth rate (corresponding to the inverse free-fall time) manifests at the infinitely large scale. The largest growth rates for GI of realistic configurations are commonly the fractions of the inverse free-fall time.","Citation Text":["Miyama, Narita & Hayashi 1987"],"Functions Text":["The linear stability analysis of these idealized configurations","as well as the corresponding non-linear solutions (e.g.","confirms this view."],"Functions Label":["Similarities","Similarities","Similarities"],"Citation Start End":[[1208,1237]],"Functions Start End":[[936,999],[1125,1180],[1306,1325]]}
{"Identifier":"2019ApJ...872...97L__Skokos_et_al._2002_Instance_1","Paragraph":"We find a similar dichotomy of bars in previous studies (Elmegreen & Elmegreen 1985, 1989; Baumgart & Peterson 1986; Elmegreen et al. 1996; Regan & Elmegreen 1997; Kim et al. 2015). They investigated two types of bars: flat and exponential profiles in surface brightness. Flat bars have nearly constant light distributions along the bar, whereas those of exponential bars decrease exponentially. They also differ in their structures and the intensity contrast between the bar and the disk: flat bars are longer, wider, and stronger than exponential bars and have a higher contrast than exponential bars. Besides, Athanassoula (1992) explained that flat bars could have roughly rectangular orbits around the end of the bar through the stellar orbit calculation. These properties can also be explained by the locations of resonances (Lynden-Bell 1979; Contopoulos & Papayannopoulos 1980; Sellwood 1981; Contopoulos et al. 1989; Athanassoula 1992; Skokos et al. 2002). A flat density profile develops in crowding stellar orbits between the inner 4:1 resonance and corotation radius (Combes & Elmegreen 1993; Elmegreen & Elmegreen 1985; Elmegreen et al. 1996). Exponential bars end near the inner Lindblad resonance and do not have such crowding orbits (Lynden-Bell 1979; Elmegreen & Elmegreen 1985; Elmegreen et al. 1996). Therefore, flat bars and exponential bars may be expected to have different pattern speeds based on the value of \n\n\n\n\n\n, where Rcr and Rbar are the radius of the corotation resonance and the bar, respectively (Debattista & Sellwood 2000; Valenzuela & Klypin 2003). Although some studies reported the observational lack of slow bars (Debattista & Sellwood 2000; Aguerri et al. 2015), others showed that the pattern speed of bars roughly depends on the Hubble type: fast bars in early-type spirals and slow bars in late-type spirals (Aguerri et al. 1998; Rautiainen et al. 2008). More observational data will help us understand the relation between the density profile and the pattern speed of bars.","Citation Text":["Skokos et al. 2002"],"Functions Text":["These properties can also be explained by the locations of resonances"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[945,963]],"Functions Start End":[[761,830]]}
{"Identifier":"2019AandA...624A..92K___2017_Instance_1","Paragraph":"In Sects. 4.3 and 4.4 we showed that our image simulations match the actual data very well and that the re-weighting factors are close to unity for subsets of galaxies split by size and S\/N, with fluctuations of the order of 10\u22124 or less. Consequently, the estimates of the shear biases presented in Table 2 should be accurate. Nonetheless it is worthwhile to explore the robustness of our results and attempt to quantify the potential systematic uncertainties that may still be present. For instance, we simplified the galaxy morphologies by representing them with S\u00e9rsic profiles. Our input catalogue is incomplete at the faintest magnitudes, but the missing galaxies may still affect the estimate of the multiplicative bias. Varying star densities can affect the results (e.g. Hoekstra et al. 2015, 2017). In this section we therefore explore the sensitivity of the shear measurement bias to various assumptions and simplifications made in the simulations. These results help to assess the robustness of the calibration presented in the previous section. In the language described in Sect. 2.2, these tests correspond to various evaluations of the \u0394 terms from different simulations. As these tests can become computationally expensive very quickly, we use only 5 of the 13 PSF sets. This results in only minor changes in the mean residual bias values ( 0.005; see lower right panel of Fig. 13), and hence they remain a good representation of the data. The smaller volume of simulations naturally results in a larger statistical uncertainty, but we note that we are interested in determining the change in the mean values of the multiplicative biases of the different tomographic samples when we vary the inputs. The errors are tightly correlated among the different simulations within a tomographic bin, as there is a significant overlap among the input samples of galaxies in the different simulations and because the noise realisations in the images are identical in many of these simulations, unless explicitly mentioned otherwise. Hence, the shifts are not driven by noise. Our main objective is to ensure that uncertainties in the input quantities do not change the bias by more than 0.02. The results in this section indicate that they are indeed controlled under 0.02, but it appears that we cannot impose a tighter limit on the overall uncertainty at this time.","Citation Text":["Hoekstra et al.","2017"],"Functions Text":["Varying star densities can affect the results (e.g.","In this section we therefore explore the sensitivity of the shear measurement bias to various assumptions and simplifications made in the simulations."],"Functions Label":["Uses","Uses"],"Citation Start End":[[780,795],[802,806]],"Functions Start End":[[728,779],[809,959]]}
{"Identifier":"2022ApJ...930...32C__Pilbratt_et_al._2010_Instance_1","Paragraph":"Most DSFGs are discovered in wide-area observations using single-dish far-IR (FIR)\/millimeter instruments such as the Submillimeter Common-User Bolometer Array (SCUBA\/SCUBA-2) on the James Clerk Maxwell Telescope (JCMT; e.g., Smail et al. 1997; Hughes et al. 1998; Chapman et al. 2005; Koprowski et al. 2017; Simpson et al. 2019), the Herschel Space Observatory\n10\n\n\n10\nHerschel is an ESA space observatory with science instruments provided by European-led Principal Investigator consortia and with important participation from NASA. (Pilbratt et al. 2010; Eales et al. 2010; Oliver et al. 2012), and the AzTEC instrument (Scott et al. 2008; Aretxaga et al. 2011). Multi-wavelength follow-up of hundreds of survey-identified DSFGs reveals that most sit between 1  z  3 (Casey et al. 2012a; Magnelli et al. 2011, 2013; Gruppioni et al. 2013; Le Floc\u2019h et al. 2005). While there exists a handful of individually studied DSFGs at redshifts as high as z \u223c 5\u20137 (e.g., Cooray et al. 2014; Strandet et al. 2017; Marrone et al. 2018; Zavala et al. 2018b; Casey et al. 2019; Reuter et al. 2020), such high-z systems have proven difficult to both identify and spectroscopically confirm. This is because (i) DSFGs at z > 4 are outnumbered by the dominant DSFG population at z \u2248 1\u20133 and (ii) there are serious physical and evolutionary degeneracies that make DSFG photometric redshifts highly uncertain (with precision often \u03c3\n\u0394z\/1+z\n \u2273 1; Casey 2020). This latter point is often seen as a benefit: their strongly negative k-correction means that the flux density of DSFGs at z > 1 remains constant with increasing z for \u03bb\nobs \u2273 850 \u03bcm, meaning a DSFG at z \u223c 10 can be observed as readily as a DSFG at z \u223c 1 (Blain et al. 2002). However, when searching for high-z DSFGs, this negative k-correction is also a hindrance as it becomes difficult to identify redshifts for galaxies with only long-wavelength emission. This effect is further exacerbated by the (sub)millimeter color degeneracy between dust temperature and redshift. Thus, barring clear identification at other wavelengths, it is easy to confuse z \u223c 2 and z \u223c 6 DSFGs with solely submillimeter observations.","Citation Text":["Pilbratt et al. 2010"],"Functions Text":["Most DSFGs are discovered in wide-area observations using single-dish far-IR (FIR)\/millimeter instruments such as","the Herschel Space Observatory"],"Functions Label":["Background","Background"],"Citation Start End":[[535,555]],"Functions Start End":[[0,113],[331,361]]}
{"Identifier":"2020AandA...644L...7G__Magdis_et_al._2020_Instance_1","Paragraph":"As in G18, we compiled existing constraints on the molecular gas fraction fgas of quiescent and pSB galaxies from recent literature, namely: local QGs consisting of the ATLAS3D (Young et al. 2011; Cappellari et al. 2013; Davis et al. 2014) and HRS (Boselli et al. 2014; Lianou et al. 2016) ETG samples as well as the samples of pSB galaxies (hereafter, the \u201clow-z pSB\u201d sample) of French et al. (2015) and Alatalo et al. (2016); at low and intermediate redshift, the ETG sample of Spilker et al. (2018) and the pSB sample of Suess et al. (2017); at intermediate and high redshift, constraints from Hayashi et al. (2018) on gas in z\u2004\u223c\u20041.46 cluster ETGs, as well as on individual galaxies from Sargent et al. (2015), Bezanson et al. (2019), and Rudnick et al. (2017). Given its size, we divided the ATLAS3D sample into high- and low-mass subsamples, choosing 5\u2005\u00d7\u20051010 M\u2299 as the cut-off mass. In addition, we also included fgas estimates derived from the (median) stacked FIR spectral energy distributions of ETGs at z\u2004\u223c\u20041.8 (G18; 977 galaxies), z\u2004\u223c\u20041.2, z\u2004\u223c\u20040.8, and z\u2004\u223c\u20040.5 (1394, 1536, and 563 galaxies, respectively; Magdis et al. 2020, hereafter M20). Finally, at higher redshift (z\u2004\u223c\u20043), we converted star formation rates (SFR) estimated from spectroscopy (Schreiber et al. 2018a; D\u2019Eugenio et al. 2020) into gas masses assuming the star formation efficiency found by G18. As a consequence of our zmax\u2004=\u20043.5, we did not include higher-redshift quiescent galaxies (Glazebrook et al. 2017; Schreiber et al. 2018b; Tanaka et al. 2019; Valentino et al. 2020) in the analysis and considered z\u2004\u223c\u20043 galaxies as pSB. The dust-based estimates of G18 and M20 (and, by extension, the z\u2004\u223c\u20043 semi-constraints) assume a gas-to-dust ratio (G\/D). It is dependent on metallicity, which is presumed to be solar or higher owing to both the relatively high gas-phase metallicity of MS galaxies at z\u2004\u2272\u20041 (e.g., Mannucci et al. 2010) and the already high stellar metallicities of QGs at z\u2004> \u20041 (Onodera et al. 2015; Estrada-Carpenter et al. 2019). Here we adopted an intermediate value between the solar and supersolar G\/Ds used in M20, and we increased the error bars of these points to include both the solar and supersolar confidence estimates. These various samples, which are summarized with their selection criteria in Table B.1, combine into a nonhomogeneous dataset: some were specifically selected as ETGs, and others were based on varying degrees of quiescence. In particular, pSB galaxies are not necessarily truly quiescent and could, in principle, resume normal star formation. However, as a possible precursor of QGs, they provide useful, though not constraining (see Sect. 4), comparison samples for the model. Here we refer to all equally as either QGs or pSB galaxies, and we make the assumption that, on average, these different samples are not otherwise significantly biased with regard to their gas content compared to the full population, given each mass limit and type.","Citation Text":["Magdis et al. 2020"],"Functions Text":["In addition, we also included fgas estimates derived from the (median) stacked FIR spectral energy distributions of ETGs at","z\u2004\u223c\u20041.2, z\u2004\u223c\u20040.8, and z\u2004\u223c\u20040.5 (1394, 1536, and 563 galaxies, respectively;","hereafter M20"],"Functions Label":["Uses","Uses","Uses"],"Citation Start End":[[1117,1135]],"Functions Start End":[[889,1012],[1042,1116],[1137,1150]]}
{"Identifier":"2018AandA...614A..31B__Nakajima_et_al._2012_Instance_1","Paragraph":"High-redshift z \u2273 2 galaxies with prominent Lyman-\u03b1 emission (Partridge & Peebles 1967), referred to as Lyman-\u03b1 emitters (LAEs), have become the subject of intense research over the last two decades. Since LAEs are believed to be powered at least partially by ongoing star formation, they are expected to belong to a relatively low-mass and young class of actively star-forming galaxies. This suggests that the number density of LAEs is large enough to map out the large-scale structure of the high-redshift universe. Increasing numbers of LAEs are being observed by surveys including narrow-band imaging (e.g., Nakajima et al. 2012) and integral field unit spectroscopy such as MUSE (e.g., Bacon et al. 2006), amounting to about 104 emitters known to date since the first detections (e.g., Cowie & Hu 1998; Hu & McMahon 1996). These observations enable us to study galaxy clustering at somewhat small scales (e.g., Diener et al. 2017; Ouchi et al. 2010, 2018) as well as evolution of the Lyman-\u03b1 luminosity function (e.g., Konno et al. 2016, 2018; Ouchi et al. 2008). A remarkable example for such surveys is the Hobby-Eberly Telescope Dark Energy Experiment (HETDEX; Adams et al. 2011; Hill et al. 2008). HETDEX will map out the three-dimensional distribution of nearly one million LAEs (Leung et al. 2017) at 1.9  z  3.5 over 400 deg2. This allows us to precisely measure the large-scale (\u227310 Mpc) clustering of LAEs (Agrawal et al. 2017; Chiang et al. 2013) and of the Lyman-\u03b1 intensity map (Saito et al., in prep.) with the primary scientific goal being to determine the cosmic expansion history via baryon acoustic oscillations (BAOs) and the growth of structure via redshift-space distortions (RSD). We refer the reader to Alam et al. (2017) and references therein for recent efforts of BAO and RSD measurements. In addition, these surveys offer exciting opportunities to study the connection between LAEs and their environment, including the cross-correlation between LAEs and the Lyman-\u03b1 forest or other galaxy populations.","Citation Text":["Nakajima et al. 2012"],"Functions Text":["Increasing numbers of LAEs are being observed by surveys including narrow-band imaging (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[612,632]],"Functions Start End":[[518,611]]}
{"Identifier":"2021MNRAS.505.2561C__Michtchenko_et_al._2018_Instance_1","Paragraph":"Several mechanisms have been suggested to explain the formation of moving groups. A common explanation is that these velocity structures are the remnants of open clusters, or formed by interactions with a bar (Eggen 1965; Dehnen 2000). One problem with the cluster formation idea is that stars in moving groups can have a variety of different ages and compositions, so it is unlikely that they all came from the same cluster (Eggen 1965; Kushniruk et al. 2020). Analysis of GALAH data (Quillen et al. 2018) indicates that some moving groups, such as the Hercules moving group, may be due to a resonant bar. It has also been suggested that moving groups could have been formed from perturbations due to the Magellanic Clouds via gravitational interactions (Dehnen 1998). Recent work also finds that transient spiral structure (Hunt et al. 2018) may lead to the formation of moving groups, as well as perturbations due to spiral arms in the MW (Michtchenko et al. 2018). Moving groups in Gaia data have also been identified and analysed in action space. In the (JR, Jz) plane there are at least seven overdensities that follow lines of constant slope in this plane, which correspond to known moving groups in the solar neighbourhood (Trick, Coronado & Rix 2019). It is likely that there may be multiple causal mechanisms at play in the formation of moving groups in the Milky Way. The analysis of Gaia DR2 data has revealed many facets of a Galaxy that are clearly out of equilibrium, including the so-called phase-space spiral (Antoja et al. 2018), and the Enceladus merger (Helmi et al. 2018), that have been interpreted as arising from interactions with dwarf galaxies. Analysis of Gaia DR2 data also led to the discovery of a new dwarf galaxy (Torrealba et al. 2019) that likely interacted with the Milky Way (Chakrabarti et al. 2019). However, the formation of moving groups due to dwarf galaxy interactions has not yet been studied with full N-body simulations. Motivated by these earlier works that indicate that the MW may has been perturbed by dwarf galaxies, we focus our study here in trying to understand if some of the moving groups in the Galaxy may have arisen from dwarf galaxy interactions.","Citation Text":["Michtchenko et al. 2018"],"Functions Text":["Recent work also finds that transient spiral structure","may lead to the formation of moving groups, as well as perturbations due to spiral arms in the MW"],"Functions Label":["Background","Background"],"Citation Start End":[[943,966]],"Functions Start End":[[770,824],[844,941]]}
{"Identifier":"2021AandA...650A.205V__Jones_et_al._2021_Instance_1","Paragraph":"The question of the evolution of exoplanet systems after the main sequence of their host is generally addressed by studying exoplanets around subgiants, RGB stars, and normal HB (RC) stars (hereafter the \u2019classical\u2019 evolved stars). These classical evolved stars are typically very large stars, with radii ranging from ~ 5\u2212 10 R\u2299 to more than 1000 R\u2299. This is much larger than hot subdwarfs, which have radii in the range ~ 0.1\u22120.3 R\u2299 (Heber 2016). Their mass is typically higher than ~ 1.5 M\u2299, compared to~0.47 M\u2299 for hot subdwarfs. The transit and radial velocity (RV) methods are both challenging for these classical evolved stars because the transit depth is diluted and there are additional noise sources (Van Eylen et al. 2016). Another difficulty forthe question of the fate of exoplanet systems after the RGB phase itself is the difficulty of distinguishing RGB and RC stars based on their spectroscopic parameters alone, which is sometimes hard even with help of asteroseismology (Campante et al. 2019). As a consequence, only large or massive planets are detected around the classical evolved stars (Jones et al. 2021, and references therein). A dearth of close-in giant planets is observed around these evolved stars compared to solar-type main-sequence stars (Sato et al. 2008; D\u00f6llinger et al. 2009). This may be caused by planet engulfment by the host star, but current technologies do not allow us to determine whether smaller planets and remnants (such as the dense cores of former giant planets) are present. The lack of close-in giant planets may also be explained by the intrinsically different planetary formation for these intermediate-mass stars (see the discussion in Jones et al. 2021). Ultimately, the very existence of planet remnants may be linked to the ejection of most of the envelope on the RGB that occurs for hot subdwarfs, while for classical evolved stars, nothing stops the in-spiraling planet inside the host star, and in all cases, the planet finally merges with the star, is fully tidally disrupted, or is totally ablated by heating or by the strong stellar wind. In other words, the ejection of the envelope not only enables the detection of small objects as remnants, but most importantly, it may even be the reason for the existence of these remnants by stopping the spiraling-in in the host star.","Citation Text":["Jones et al. 2021"],"Functions Text":["As a consequence, only large or massive planets are detected around the classical evolved stars","and references therein"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1109,1126]],"Functions Start End":[[1012,1107],[1128,1150]]}
{"Identifier":"2022MNRAS.515.4520F__Rodriguez-Franco_&_Cuevas_2013_Instance_1","Paragraph":"The atmosphere above the Canary Islands, and especially the western islands Tenerife and La Palma, has been extensively characterized over the past 50 yr (Kiepenheuer 1972; McInnes & Walker 1974; Brandt & Righini 1985; Murdin 1985; Stickland et al. 1987; Whittet, Bode & Murdin 1987; Men\u00e9ndez-Vald\u00e9s & Blanco 1992; J\u00edmenez, Gonz\u00e1lez Jorge & Rabello-Soares 1998; Mahoney, Mu\u00f1oz-Tu\u00f1\u00f3n & Varela 1998; Maring et al. 2000; Torres et al. 2002; Rodr\u00edguez et al. 2004; Alonso-P\u00e9rez et al. 2007; Basart et al. 2009; Lombardi, Zitelli & Ortolani 2009; Delgado et al. 2010; Vernin et al. 2011; Cuevas et al. 2013; Rodriguez-Franco & Cuevas 2013; Varela et al. 2014; Laken et al. 2016; Vogiatzis et al. 2018; Hidalgo et al. 2021). Due to the combination of large-scale atmospheric circulation on the descending branch of the Hadley cell (see e.g. Palm\u00e9n & Newton 1969; Rodr\u00edguez et al. 2004), and the \u2018Trade\u2019 or \u2018Alisios\u2019 winds coming from the Azores high area, a stable and strong temperature inversion layer (TIL) appears (Font 1956; Huetz-de-Lemps 1969), whose top is typically found at heights around 1200\u2009m a.s.l. in summer and 1800\u2009m a.s.l. in winter (Torres et al. 2002; Carrillo et al. 2016). Whenever the temperature inversion is able to separate two well-defined regimes, the moist marine boundary layer (MBL) and above it, the dry free troposphere (FT), the phenomenon is called an \u2018Alisio\u2019 inversion. This happens about 80 per\u2009cent of the time (Torres et al. 2002). Under these conditions, the FT can be characterized as \u2018ultraclean\u2019, i.\u2009e., the concentration of accumulation mode particulate matter is lower than 50\u2009cm\u22123 (Pennypacker, Diamond & Wood 2020). The clean atmospheric conditions are one of the reasons why the Canarian observatories \u2018Observatorio del Teide (OT)\u2019, located at \u223c2400\u2009m a.s.l., and \u2018Observatorio del Roque de los Muchachos (ORM)\u2019, located between \u223c2100 and \u223c2400\u2009m a.s.l., are known to belong to the best astronomical sites worldwide.","Citation Text":["Rodriguez-Franco & Cuevas 2013"],"Functions Text":["The atmosphere above the Canary Islands, and especially the western islands Tenerife and La Palma, has been extensively characterized over the past 50 yr"],"Functions Label":["Background"],"Citation Start End":[[603,633]],"Functions Start End":[[0,153]]}
{"Identifier":"2021AandA...656A.126K__Lowry_et_al._2007_Instance_1","Paragraph":"Very small asteroids (VSAs) are objects with diameters D  150 m. They often rotate with periods shorter than 2 h enabling us to study their internal structure by comparing the centrifugal force with the material forces holding them together (Holsapple 2007). Because of their small sizes, VSAs are sensitive to the YORP effect (Rubincam 2000), which is a torque induced on the rotating body by the thermal radiation emitted by its surface complemented by a torque produced by scattered sunlight. YORP can either spin up or slow down the asteroid rotation as well as change the obliquity of its spin axis \u03f5, which is an angle between the normal to the asteroid orbital plane and its rotation axis. While the fast rotation has been observed for many VSAs, their spin axes were not determined except for one object1, (54 509) YORP (whose name is the same as the name of the effect itself). (54 509) was the first asteroid for which the effect of YORP has been observed (Lowry et al. 2007; Taylor et al. 2007). The obliquity of the (54 509) spin axis is \u03f5 = 173\u00b0, which means it is nearly perpendicular to the asteroid orbital plane. Such an orientation of the spin axis was found as an end state of the YORP evolution in the simulations performed by \u010capek & Vokrouhlick\u00fd (2004) for objects with finite surface thermal conductivity. If their prediction is true, then for the VSAs that experienced a strong YORP effect for a long time (and the fastest rotating VSAs are such objects), we should observe spin axis obliquities close to 0\u00b0 or 180\u00b0. Recently Golubov et al. (2021) have shown that for very small objects of a highly irregular shape, the transverse heat conduction (TYORP) can add new asymptotic states for the obliquity. For this reason it would be interesting to verify those predictions with observationsof VSAs. To do that, we should observe their lightcurves at many different positions on the sky to be able to determine their spin axes. For near-Earth objects (NEAs) this condition is met either by the Earth co-orbital asteroids \u2013 and (54 509) is an example of such objects \u2013 or by objects for which their close encounter with the Earth can be observed along a long arc on the sky.","Citation Text":["Lowry et al. 2007"],"Functions Text":["(54 509) was the first asteroid for which the effect of YORP has been observed"],"Functions Label":["Background"],"Citation Start End":[[967,984]],"Functions Start End":[[887,965]]}
{"Identifier":"2020MNRAS.499..710G__Assef_et_al._2015_Instance_1","Paragraph":"An alternative possibility to reduce the fraction of X-ray sources in the WISE R75 wedge is to allow a large fraction of heavily obscured AGN. This can be achieved for example, by increasing the fraction of Compton-thick AGN in the X-ray luminosity function above the current assumption of 34 per cent, or by increasing the scatter in the relation that links X-ray obscuration to optical extinction beyond the adopted value of 0.5\u2009dex. Relaxing the above model assumptions would allow heavily obscured and hence, X-ray faint, sources to be selected by the WISE R75 criteria. There is indeed evidence for a potentially large population of heavily obscured, possibly Compton-thick AGN, among the WISE population (e.g. Assef et al. 2015; Mountrichas et al. 2017; Yan et al. 2019). The SDSS spectroscopic follow-up programme presented by LaMassa et al. (2019) also revealed a non-negligible number of WISE R75 sources that are optically faint ($r\\gtrsim 22$\u2009mag), lie at redshifts $z\\lesssim 1$ and are spectroscopically identified by their prominent [O\u2009ii]\u20093727 emission lines. These sources are prime candidates for heavily obscured AGN. Our baseline model predicts that the most heavily obscured, Compton-thick, WISE AGN are at low redshift, $z\\lesssim 0.6$ (see Fig. 19) and relatively optically bright. The top panel of Fig. 22 demonstrates the latter point by plotting the r-band distribution of the Compton-thick AGN population predicted by the model. Observationally, the identification of such AGN needs to account for the relatively high level of contamination of the WISE R75 AGN selection by star-forming galaxies at redshifts $z\\lesssim 0.6$ (s1ee Fig. 14). One approach to achieve this is via diagnostic optical emission-line ratios (e.g. Kewley et al. 2001) to separate star-forming galaxies from Seyfert-2s. Observations at hard X-rays can also provide useful information on the nuclear activity of a galaxy and the level of line-of-sight obscuration. The bottom panel of Fig. 22 shows the expected 2\u201310 keV X-ray flux distribution of the Compton-thick AGN predicted by the model. The expected fluxes have already been reached in deep X-ray survey fields, for example, COSMOS-Legacy (Civano et al. 2016). Study of the X-ray spectral properties of WISE-selected AGN in such fields can test the baseline model predictions for the demographics of Compton-thick AGN in the WISE R75 wedge.","Citation Text":["Assef et al. 2015"],"Functions Text":["There is indeed evidence for a potentially large population of heavily obscured, possibly Compton-thick AGN, among the WISE population (e.g."],"Functions Label":["Background"],"Citation Start End":[[716,733]],"Functions Start End":[[575,715]]}
{"Identifier":"2018AandA...613L...1D__Salafia_et_al._2015_Instance_1","Paragraph":"Similarly to long bursts, short GRBs are thought to be produced by a relativistic jet with a typical half-opening angle \u03b8jet ~ 5\u201315 deg (Fong et al. 2016). However, whether or not BNS mergers can always efficiently produce a relativistic jet is still debated (Paschalidis et al. 2015; Ruiz et al. 2016; Kawamura et al. 2016). Given the small probability that our line of sight was within the jet half-opening angle, 1 \u2212 cos(\u03b8jet), it is unlikelythat the first short GRB associated to a GW event had a jet pointing towards the Earth. The extremely low \u03b3-ray luminosity of GRB 170817A has been interpreted as being due to (i) the debeamed radiation of a jet observed off-beam (i.e. viewing angle \u03b8view > \u03b8jet), provided that the jet bulk Lorentz factor is significantly smaller than usually assumed (see, e.g. Pian et al. 2017). Alternatively, the jet could be (ii) structured, with a fast and energetic inner core surrounded by a slower, less energetic layer\/sheath\/cocoon (first proposed for long GRBs \u2013 Lipunov et al. 2001; Rossi et al. 2002; Salafia et al. 2015 \u2013 and only recently extended to short GRBs \u2013 Kathirgamaraju et al. 2018; Lazzati et al. 2017a; Gottlieb et al. 2017; Lazzati et al. 2017b; Lyman et al. 2018; Margutti et al. 2018; Troja et al. 2018a). In this scenario the faint, off-beam emission is due to the slower component, which originates from the interaction of the jet with the merger dynamical ejecta or the post-merger winds. Recently, Mooley et al. (2018) suggested the possibility that (iii) the jet was not successful in excavating its way through the ambient medium and that GRB 170817A was due to its vestige, a quasi-isotropic cocoon with a velocity profile. Last but not least, (iv) a jet-less interpretation of GRB 17017A could still be viable: an isotropic fireball, expanding ahead of the kilonova ejecta, which could account for both the low luminosity of the \u03b3-ray event and the properties of the EM component in the radio and X-ray bands (Salafia et al. 2017). In this case, all BNS mergers should have this kind of faint, hard X-ray counterpart. All of the above scenarios have relatively clear predictions for the temporal and spectral evolution of the electromagnetic emission from X-rays to the radioband. A comprehensive discussion of the possible physical scenarios leading to the observed broad-band emission of GW 170817\/GRB 170817A can be found in Nakar & Piran (2018). Recent radio and X-ray observations(Mooley et al. 2018; Margutti et al. 2018; Ruan et al. 2018; Troja et al. 2018a), carried out until ~ 110\u2212115 d after the event, indicate that the source flux is steadily rising and that the spectral energy distribution (SED) over these bands is consistent with a single power-law component. These results disfavour interpretation (i) reported above (an off-beam homogeneous jet).","Citation Text":["Salafia et al. 2015"],"Functions Text":["Alternatively, the jet could be (ii) structured, with a fast and energetic inner core surrounded by a slower, less energetic layer\/sheath\/cocoon (first proposed for long GRBs"],"Functions Label":["Uses"],"Citation Start End":[[1044,1063]],"Functions Start End":[[827,1001]]}
{"Identifier":"2015MNRAS.454.1644L__Kotze_&_Charles_2012_Instance_2","Paragraph":"The period candidates of other three ULXs may range from \u223c100 to \u223c600 d. Apart from noise and artefacts, all the candidate periods are only significant in a specific epoch. This suggests that they are not associated with any stable mechanism such as orbital motion. Instead, such long-term (> 100 d) X-ray quasi-periodic variations are likely related to superorbital periods that are thought to be due to radiation-driven warping of accretion discs (Ogilvie & Dubus 2001) or tidal interaction-induced disc precession (Whitehurst & King 1991). Alternatively, mass transfer rate-related events such as X-ray state changes and disc instability can also cause long-term modulations (Kotze & Charles 2012). In particular, there are two intermittent quasi-periodicities for both NGC 5408 X-1 and M81 X-6, suggesting that the quasi-periods are changing or evolving. They resemble some Galactic X-ray binaries that show similar behaviour (e.g. Cyg X\u20132 and SMC X\u20131; Kotze & Charles 2012) and it has been suggested that a warped disc could lead to an unstable steadily precessing disc, causing quasi-periodic behaviour (Ogilvie & Dubus 2001). We note that there are many uncertainties on the physical parameters of ULXs. To determine the origin of superorbital periods of ULXs, one has to know at least the mass ratio between the companion and the compact star (q = MC\/MX) and the binary separation. Unfortunately, it is very difficult to get these parameters for ULXs. For the three ULXs discussed here (i.e. excluding ESO 243-49 HLX-1), only M81 X-6 has better constraints on the black hole mass and the nature of the companion. The masses of the black hole and companion star are estimated (MX = 18\u2009M\u2299, MC = 23\u2009M\u2299) such that q can be derived. In this case, we can rule out a tidal interaction-induced disc precession scenario that requires q  0.25\u20130.33 (Whitehurst & King 1991). For a warped disc, the binary separation and the mass ratio suggest that M81 X-6 lies in the intermediate instability zone for radiation-driven warping in X-ray binaries (see fig. 1 of Kotze & Charles 2012). The quasi-periodic variability may represent the switching time-scale between a warped disc and a flat disc.","Citation Text":["Kotze & Charles 2012"],"Functions Text":["In particular, there are two intermittent quasi-periodicities for both NGC 5408 X-1 and M81 X-6, suggesting that the quasi-periods are changing or evolving. They resemble some Galactic X-ray binaries that show similar behaviour (e.g. Cyg X\u20132 and SMC X\u20131;"],"Functions Label":["Similarities"],"Citation Start End":[[957,977]],"Functions Start End":[[702,956]]}
{"Identifier":"2020AandA...644A..59K__Pastorello_et_al._2019_Instance_1","Paragraph":"The year 2020 marks the 350 yr anniversary of the discovery of the eruption of Nova 1670 (or CK Vul) made by European astronomers (Shara et al. 1985). Their observations, predominantly performed with a naked eye, traced the object\u2019s evolution on the sky in 1670\u22121672. From the archive records, we know that the eruption was rather unusual, in particular it was very much unlike classical novae. The light curve of CK Vul displayed three peaks and the star was described as reddish in the later stages of the eruption (Hevelius 1671; Shara et al. 1985). These characteristics resemble closest the behavior often observed in (luminous) red novae (Kato 2003; Tylenda et al. 2013), a modern category of eruptive stars known from our and other galaxies (e.g. Pastorello et al. 2019). Red novae are recognized as manifestations of on-going mergers of non-compact stars such as main-sequence dwarfs, sub-giants, or red giants (Soker & Tylenda 2003; Tylenda & Soker 2006; Tylenda et al. 2011; Pastorello et al. 2019). While the number of known red novae, mainly extragalactic ones, is quickly rising (e.g., Stritzinger et al. 2020), we know only a few red-nova remnants that are decades old (Kami\u0144ski et al. 2018a). The remnant of the 1670 eruption of CK Vul, as a candidate post-merger site, could be the oldest (counting from the onset of the eruption) known object of this type and as such offers the opportunity to investigate a merger aftermath centuries after the stellar coalescence. The nature of the progenitor system of CK Vul has been debated. Eyres et al. (2018) proposed that the seventeenth-century merger took place between a white dwarf and a brown dwarf, but there is little quantitative evidence to support this. Based on the analysis of the source\u2019s chemical and isotopic composition, including the unique presence of the radioactive isotope of 26Al, Kami\u0144ski et al. (2018b) found that the progenitor system of CK Vul included at least one red-giant-branch (RGB) star with a fully developed helium core.","Citation Text":["Pastorello et al. 2019"],"Functions Text":["These characteristics resemble closest the behavior often observed in (luminous) red novae","a modern category of eruptive stars known from our and other galaxies (e.g."],"Functions Label":["Background","Background"],"Citation Start End":[[754,776]],"Functions Start End":[[553,643],[678,753]]}
{"Identifier":"2015AandA...582A..22L__Todorov_et_al._2014_Instance_1","Paragraph":"Dust distribution:We employed the standard flared disk model with well-mixed gas and dust, which has been successfully used to explain the observed SEDs of a large sample of young stellar objects and BDs (e.g., Wolf et al. 2003; Sauter et al. 2009; Harvey et al. 2012a; Joergens et al. 2013; Liu et al. 2015). The structure of the dust density is assumed with a Gaussian vertical profile (1)\\begin{equation} \\rho_{\\rm{dust}}=\\rho_{0}\\left(\\frac{R_{*}}{\\varpi}\\right)^{\\alpha}\\exp\\left[-\\frac{1}{2}\\left(\\frac{z}{h(\\varpi)}\\right)^2\\right], \\label{dust_density} \\end{equation}\u03c1dust=\u03c10R\u2217\u03d6\u03b1exp\u221212zh(\u03d6)2,and the surface density is described as a power-law function (2)\\begin{equation} \\Sigma(\\varpi)=\\Sigma_{0}\\left(\\frac{R_{*}}{\\varpi}\\right)^p, \\end{equation}\u03a3(\u03d6)=\u03a30R\u2217\u03d6p,where \u03d6 is the radial distance from the central star measured in the disk midplane, and h(\u03d6) is the scale height of the disk. The disk extends from an inner radius Rin to an outer radius Rout. To the best of our knowledge, among our sample, there are five objects that have been identified as binary systems so far. They are 2M1207 (a~55 AU, Chauvin et al. 2004), J04221332+1934392 (a~7 AU, Todorov et al. 2014), J04414489+2301513 (a~15 AU, Todorov et al. 2014), USD161833 (a~134 AU, Bouy et al. 2006), and USD161939 (a ~ 26 AU, Bouy et al. 2006), where a refers to the separation within the system. The disks around individual components in binary systems are expected to have truncation radii of the order of (0.3 \u2212 0.5)a (Papaloizou & Pringle 1977). We adopted 0.5 a as the disk outer radii for 2M1207, USD161833, and USD161939. For the close pairs (a \u2272 15 AU, J04221332+1934392 and J04414489+2301513), dynamical simulations of star-disk interactions suggest that individual disks are unlikely to survive (e.g., Artymowicz & Lubow 1994). Disk modeling is complicated in those close multiple systems. For simplicity, we assume that the emission is associated with circumbinary disks of 100 AU in size. For other objects, we fix Rout = 100 AU in the modeling because the choice of this parameter value makes essentially no difference to the synthetic SEDs in the simulated wavelength range (Harvey et al. 2012a). The scale height follows the power-law distribution(3)\\begin{equation} h(\\varpi) = H_{100}\\left(\\frac{\\varpi}{100\\,\\rm{AU}}\\right)^\\beta,\\\\ \\end{equation}h(\u03d6)=H100\u03d6100\u2009AU\u03b2,with the exponent \u03b2 characterizing the degree of flaring and H100 representing the scale height at a distance of 100 AU from the central star. The indices \u03b1, p, and \u03b2 are codependent through p = \u03b1 \u2212 \u03b2. We fix p = 1, which is the typical value found for T Tauri disks in the sub-millimeter (e.g., Isella et al. 2009; Guilloteau et al. 2011), since only spatially resolved data can place constraints on this parameter (e.g., Ricci et al. 2013, 2014).               Dust properties:We assume the dust grains to be a homogeneous mixture of 75% amorphous silicate and 25% carbon with a mean density of \u03c1grain = 2.5 g cm-3 and the complex refractive indices given by J\u00e4ger et al. (1994, 1998), and Dorschner et al. (1995). Porous grains are not considered because the fluxes at wavelengths beyond ~ 2 \u03bcm are almost independent of the degree of grain porosity in low-mass disks, as shown by Kirchschlager & Wolf (2014). The grain size distribution is given by the standard power law dn(a) \u221d a-3.5da with minimum and maximum grain sizes amin = 0.1 \u03bcm and amax = 100 \u03bcm, respectively. The choice of the minimum value for the grain size, amin, ensures that its exact value has a negligible impact on the synthetic SEDs. Since there is no information about the maximum grain sizes of our target disks, as provided, e.g., by the (sub)millimeter spectral index, we adopt amax = 100 \u03bcm. The Herschel\/PACS far-IR observations are sensitive to dust grains with this assumed sizes. Strong grain growth up to millimeter sizes as detected in some BD disks (e.g., Ricci et al. 2012, 2013, 2014; Broekhoven-Fiene et al. 2014) would remain undetected in our data and could affect the disk mass. Our prescription for the dust properties is identical to those used in Liu et al. (2015). ","Citation Text":["Todorov et al. 2014"],"Functions Text":["J04221332+1934392 (a~7 AU,"],"Functions Label":["Uses"],"Citation Start End":[[1160,1179]],"Functions Start End":[[1133,1159]]}
{"Identifier":"2021AandA...649A..58L__Bemporad_et_al._(2018)_Instance_3","Paragraph":"The leading edges of the transients normally leave bright traces in the images of visible light, inspiring many methods that were developed to derive their locations and velocities, such as the icecream cone model (Fisher & Munro 1984), the graduated cylindrical shell (GCS) model (Thernisien 2011), geometric triangulation methods (Liu et al. 2010), mask-fitting methods (Feng et al. 2012), and trace-fitting methods including the point-p, fixed-\u03a6, harmonic mean, and self-similar expansion fitting methods (e.g., Sheeley et al. 1999; Howard et al. 2006; Davies et al. 2012; M\u00f6stl & Davies 2013). To derive the velocity distribution inside one transient rather than only at its leading edge, some other techniques have been proposed. Colaninno & Vourlidas (2006) applied an optical flow tool to extract the velocity vector of a coronal mass ejection (CME) in digital images. Feng et al. (2015) derived the radial velocity profiles of the whole CME from the spatial distribution of its density given by the mass continuum equation. A cross-correlation method was applied to derive continuous 2D speed maps of a CME from coronagraphic images by Bemporad et al. (2018). In their work, the radial shift pixel by pixel is determined by maximizing the cross correlation between the signal in a radial window at one frame and the signal in a radial shifted window at the previous frame, and the radial speed just equals the radial shift over the time interval between the two frames. Ying et al. (2019) improved this cross-correlation method by analyzing data in three steps: forward step (FS), backward step (BS), and average step (AS). In the FS (BS), the 2D velocity map between the current and the previous (next) frame is constructed with almost the same method as Bemporad et al. (2018). In the AS the average, velocity is obtained from the FS and BS. The velocities derived by all these methods are the component of the flow velocity vector projected onto the POS. This may underestimate the velocity especially for transients that do not propagate in the POS. Methods such as the polarizaition ratio technique (Moran & Davila 2004; DeForest et al. 2017) or the local correlation tracking (LCT) method (Mierla et al. 2009) can derive the 3D geometric information of the whole transients, but not the velocity distribution. Bemporad et al. (2018) chose the propagating direction averaged over the whole CME derived by the polarization ratio technique to correct the radial speed in the 2D maps, but the key information along the LOS is still lacking.","Citation Text":["Bemporad et al. (2018)"],"Functions Text":["chose the propagating direction averaged over the whole CME derived by the polarization ratio technique to correct the radial speed in the 2D maps, but the key information along the LOS is still lacking."],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[2324,2346]],"Functions Start End":[[2347,2550]]}
{"Identifier":"2020ApJ...898..143S__Shao_&_Li_2018_Instance_1","Paragraph":"In Figure 6, we plot the histogram distributions of the calculated number of Galactic BH\u2013He XRBs as a function of the companion mass, BH mass, orbital period, and orbital eccentricity for Models A\u2013D. Since the progenitor systems containing a BH and an OB star possess similar distributions of the OB star masses and the orbital periods in all models, the BH\u2013He XRBs also have similar distributions of helium masses and orbital periods under the assumptions of different models. The masses of the helium stars are mainly distributed in the range of \n\n\n\n\n\n, and the orbital periods are distributed in a wide range (\u223c0.1\u20131000 days) with a peak at \u223c30\u2013100 days. Some recent investigations showed that the mass-transfer process in the lobe-filling BH binaries is more stable than previously expected (e.g., Pavlovskii et al. 2017; Shao & Li 2018); the maximal mass ratio of the donor star to the BH for stable mass transfer can reach as high as \u223c6. Hence only a few BH\u2013He XRBs with orbital periods of \u22721 day are produced in our calculations. Remarkably, such wind-fed XRBs in close orbits may be bright enough to be detected. In nearby galaxies, there also exist a couple such close XRBs (Esposito et al. 2015, and references therein). During the evolution of the progenitor binaries, mass accretion of the BH can increase its mass. We see that the BH masses in the BH\u2013He XRBs are distributed in a broad range of \n\n\n\n\n\n. Model A predicts that BHs are more likely to possess masses of \n\n\n\n\n\n since more binaries survive from BH formation via direct collapse without natal kicks, while Models B\u2013D tend to produce light BHs with mass distributions that peak at \n\n\n\n\n\n due to the IMF. Differently from the BH systems with an OB star, almost all BH\u2013He XRBs have relatively low eccentricities of \u22720.4, whose distribution has two distinct peaks at \u22720.1 and \u223c0.3. Tides and mass transfer between binary components tend to circularize the orbits during the progenitor system evolution. Our obtained low eccentricities can coincide with the observations of Galactic W-R\u2212O binaries (van der Hucht 2001) that are also post-mass-transfer systems; although, we do not include a detailed treatment for the orbital evolution of mass-transferring eccentric binaries (e.g., Sepinsky et al. 2009; Dosopoulou & Kalogera 2016).","Citation Text":["Shao & Li 2018"],"Functions Text":["Some recent investigations showed that the mass-transfer process in the lobe-filling BH binaries is more stable than previously expected (e.g.,","; the maximal mass ratio of the donor star to the BH for stable mass transfer can reach as high as \u223c6. Hence only a few BH\u2013He XRBs with orbital periods of \u22721 day are produced in our calculations."],"Functions Label":["Uses","Uses"],"Citation Start End":[[826,840]],"Functions Start End":[[658,801],[841,1036]]}
{"Identifier":"2017MNRAS.470..626M__Robitaille_et_al._2012_Instance_1","Paragraph":"The opacity file used for the PAH\/VSG population (draine_opac_new.dat) was computed by Bruce Draine (Draine & Li 2007) and is available at https:\/\/github.com\/hyperion-rt\/hyperion-pah\/tree\/master\/PAH-Legacy\/input. It was computed taking two lognormal size distributions, one with grain radii below 20 \u00c5 and another for grains with radii from 20 to 200 \u00c5. PAHs constitute a mass fraction of 4.5 per cent of this grain population and they have the size-dependent fractional ionization of the diffuse ISM (Wood et al. 2008). The presence of the 8.6-\u03bcm feature in the Spitzer IRS spectrum of MPI13 implies the existence of ionized PAHs. The PAH\/VSG grain population is not in thermal equilibrium with the radiation field. However, their emissivities can be computed based on specific energy absorbed in each grid cell and assuming that the emissivity is the function of the mean intensity of the radiation field, which approximates the spectral shape to the first order (Robitaille et al. 2012; Whitney et al. 2013). The code uses the pre-computed emissivity tables of Draine & Li (2007) for different specific energy. As a Monte Carlo energy pocket is absorbed by a PAH\/VSG, the pocket is reprocessed, sampling a new frequency from pre-computed emissivity tables. A detailed discussion on this can be found in Wood et al. (2008). The updated version of Hochunk3D is hence quite useful in modelling the circumstellar shells around evolved stars, which have thermally fluctuating dust. The code uses the method of Lucy (1999) to calculate the dust temperature very efficiently by summing up photon path-lengths. The temperature of the grid cell is updated when an iteration of simulation is completed, and the temperature converges in three to five iterations; this efficiency is due to the fact that flux is conserved exactly across all surfaces (Lucy 1999). By constructing Monte Carlo radiation fields that are rigorously divergence-free, rapid convergence is achieved with the temperature correction procedure in this method. We have performed the Lucy temperature correction of dust with five iterations.","Citation Text":["Robitaille et al. 2012"],"Functions Text":["However, their emissivities can be computed based on specific energy absorbed in each grid cell and assuming that the emissivity is the function of the mean intensity of the radiation field, which approximates the spectral shape to the first order"],"Functions Label":["Uses"],"Citation Start End":[[966,988]],"Functions Start End":[[717,964]]}
{"Identifier":"2022ApJ...929..186L__Ferraro_et_al._2018b_Instance_2","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"],"Functions Text":["In this context we promoted the ESO-VLT Multi-Instrument Kinematic Survey (hereafter the MIKiS survey;","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."],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[1118,1138]],"Functions Start End":[[1015,1117],[1148,1415]]}
{"Identifier":"2022AandA...667A.131B__Izumi_et_al._(2016)_Instance_2","Paragraph":"Molecular line ratio diagnostics are often used to investigate the physics and chemistry of the ISM in all of these environments. For example, as the gas chemistry located in the central, nuclear regions of galaxies is believed to be dominated by X-rays produced by the AGN, in so-called X-ray dominated regions (XDRs), the molecular content of the ISM surrounding such nuclei will greatly differ from that in starburst regions (Usero et al. 2004; Garc\u00eda-Burillo et al. 2010). Hence, line ratios of specific molecules have been proposed as indicators of certain energetic or physical processes, for example HCN\/HCO+ as a tracer of AGNs (Loenen et al. 2007), HCN\/HNC as a mechanical heating tracer (Hacar et al. 2020), and HCN\/CO as a density tracer (Leroy et al. 2017). In particular, the \u201csubmillimeter-HCN diagram\u201d, first proposed in Izumi et al. (2013) and later expanded upon in Izumi et al. (2016), is a very notable example of the use of molecular line ratios as a probe of AGN-galaxies; this diagram makes use of two line ratios, HCN(4\u22123)\/HCO+(4\u22123) and HCN(4\u22123)\/CS(7\u22126), where all of the molecules involved are considered tracers of dense gas. Izumi et al. (2016) observed a clear trend that AGNs, including the Seyfert composite galaxy NGC 1068, tend to show higher HCN\/HCO+ and\/or HCN\/CS than in SB galaxies as long as the observations were at high enough spatial resolutions to separate energetically discrete regions. Izumi et al. (2016) propose a scenario where it is the high temperature that is responsible for the HCN enhancement, whereby neutral-neutral reactions with high reaction barriers are enhanced (Harada et al. 2010), thus leading to the possible enhancement of HCN and the depletion of HCO+ via newly available formation and destruction paths, respectively. However, while of course higher gas temperatures are expected in AGN-dominated regions, these are not unique to these environments, as starburst regions and\/or regions where outflows dominate can also harbour high enough temperatures for such enhancement to occur. Additionally, the higher temperatures could increase HCN excitation, relative to HCO+ and CS, without necessarily changing their relative abundances (Imanishi et al. 2018a). Finally, infrared radiative pumping is also a possible explanation of the HCN intensity enhancement relative to HCO+ and CS. Infrared pumping is a result of the emission of 14 \u03bcm infrared photons due to the presence of hot dust around AGN. These photons vibrationally excite HCN to the \u03bd2\u2004=\u20041 state. Upon de-exciting from this state back to the vibrational ground state, \u03bd\u2004=\u20040, the HCN line intensities are thus pumped to higher fluxes (Imanishi et al. 2018a). However, we note that it is also not unlikely that the 12 \u03bcm infrared photons can similarly vibrationally excite HCO+, thus nullifying the extent of this effect (Imanishi et al. 2016).","Citation Text":["Izumi et al. (2016)"],"Functions Text":["observed a clear trend that AGNs, including the Seyfert composite galaxy NGC 1068, tend to show higher HCN\/HCO+ and\/or HCN\/CS than in SB galaxies as long as the observations were at high enough spatial resolutions to separate energetically discrete regions."],"Functions Label":["Background"],"Citation Start End":[[1151,1170]],"Functions Start End":[[1171,1428]]}
{"Identifier":"2020MNRAS.499.1788W__Malhotra_et_al._2001_Instance_2","Paragraph":"Owing to their brightness at rest-frame FIR wavelengths, the ionized and neutral species of Carbon, Nitrogen, and Oxygen are powerful diagnostic lines for tracing the ISM of nearby and distant galaxies. When combined with photodissociation region (PDR) models (Tielens & Hollenbach 1985), measurement of the emission from different lines provide a means to constrain quantities such as the ionization rate and metallicity of the ISM of galaxies. Multilevel FIR transition lines have now been widely surveyed in local galaxies, originally by ISO (e.g. Luhman et al. 1998; Malhotra et al. 2001) and more recently using the PACS spectrometer (Poglitsch et al. 2010) on the ESA Herschel Space Observatory (Pilbratt et al. 2010). The [C\u2009ii]158\u2009\u03bcm line is typically the brightest in star-forming galaxies, arising from ionized, and even neutral gas where it is the main coolant (Wolfire et al. 2003). Another commonly observed FIR line tracer of the ionized ISM is [N\u2009ii]122\u2009\u03bcm. It has the advantage that it can be found associated with lower excitation gas, close to that observed in our own Galaxy (e.g. Goldsmith et al. 2015; Herrera-Camus et al. 2016). Another major coolant of the ISM is [O\u2009i]63\u2009\u03bcm (Wolfire et al. 2003). Owing to its high excitation temperature and critical density, it can dominate the cooling in regions of starburst activity (Kaufman et al. 1999; Kaufman, Wolfire & Hollenbach 2006; Brauher, Dale & Helou 2008; Croxall et al. 2012; Narayanan & Krumholz 2017). When combined with measurements of the [C\u2009ii]158\u2009\u03bcm line intensity and FIR luminosity, the [O\u2009i]63\u2009\u03bcm line intensity can constrain the FUV field, G, and the gas density using PDR models. The luminosity in these FIR lines generally exhibits a deficit in the most FIR luminous galaxies compared to the trend expected from lower luminosity galaxies (e.g. Malhotra et al. 2001; Graci\u00e1-Carpio et al. 2011; D\u00edaz-Santos et al. 2017). This has made the emission from lines like [O\u2009i]63\u2009\u03bcm more challenging to detect at high-redshifts.","Citation Text":["Malhotra et al. 2001"],"Functions Text":["The luminosity in these FIR lines generally exhibits a deficit in the most FIR luminous galaxies compared to the trend expected from lower luminosity galaxies (e.g."],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1832,1852]],"Functions Start End":[[1667,1831]]}
{"Identifier":"2022AandA...661A..10B__Santos_et_al._2008_Instance_1","Paragraph":"It is also possible that these clusters have a smaller extent and can just be missed by our extent selection as our detection algorithm sets the extent to zero if it is smaller than 6 (Brunner et al. 2022). Following the method presented in Ghirardini et al. (2021a), we estimated several dynamical properties of the clusters in the point source sample and compared them with the extent-selected sample presented in Ghirardini et al. (2021a). In Fig. 5 we compare the distributions of the core radii (Rcore) constrained by the V06 model and the concentration parameter (cSB) between these two samples. The concentration parameter is defined as the ratio of the surface brightness within 0.1 R500 to the surface brightness within R500 (Ghirardini et al. 2021a; Santos et al. 2008; Maughan et al. 2012). Intuitively, the expectation is that the smaller the core radius, the more compact the cluster. The left panel of Fig. 5 clearly shows that the clusters in the point source sample have relatively smaller core radii, hence the emission is more concentrated in a smaller area. Consistently, the concentration of the point source sample shows a clear excess in higher values than the extent-selected sample, indicating that a significantly larger fraction of cool-core clusters and clusters host a central AGN. We performed the same experiment by applying cuts in flux 1.5 \u00d7 10\u221214 ergs s\u22121 cm\u22122 and in detection likelihood to test whether the clusters are missed by the extent selection because they are fainter and\/or more compact than the extent-selected clusters. The distribution of number density, core radius, and concentration parameter remains the same, indicating that the population of clusters in the point source catalog is more compact than the extent-selected sample. The extent-selected sample does not show a clear bias toward cool-core clusters or clusters with a central AGN, but contains the fraction of cool-cores is similar to that of SZ surveys (Ghirardini et al. 2021a). In this sample, we observe the opposite trend.","Citation Text":["Santos et al. 2008"],"Functions Text":["The concentration parameter is defined as the ratio of the surface brightness within 0.1 R500 to the surface brightness within R500"],"Functions Label":["Uses"],"Citation Start End":[[760,778]],"Functions Start End":[[602,733]]}
{"Identifier":"2019MNRAS.482.2731Z__Perley_et_al._2008_Instance_1","Paragraph":"We searched the literature carefully and selected all published GRB-DLAs and QSO-DLAs sightlines conforming to our requirements which are as follows. The object must have spectral energy distributions (SEDs) and optical spectroscopic data available with measurements of AV, column densities of Zn ii and Fe ii, or of S ii and Si ii. The GRBs are selected only if they had their optical extinction derived from simultaneous SED fitting to X-ray-to-optical\/NIR data using either a single or broken power law (see Greiner et al. 2011; Zafar et al. 2011; Schady et al. 2012; Covino et al. 2013; Bolmer et al. 2018; Zafar et al. 2018a for discussion on AV determination). This is a reliable method to determine extinctions at higher redshifts where the intrinsic slopes are constrained by the X-ray data. Note that there is some degeneracy between broken power-law break frequency (see De Cia et al. 2011) and extinction, which could lead to inference of grey dust for some instances (Watson et al. 2006; Perley et al. 2008; Friis et al. 2015). However, overall a fixed spectral break change (\u0394\u03b2 = 0.5) between the optical and X-ray slopes is preferred for GRBs (Greiner et al. 2011; Zafar et al. 2011; Japelj et al. 2015). For QSO-DLAs, reddening must be determined either from QSO colours or extinctions through template fitting to the QSO SED. Those methods are less robust than the X-ray-supported GRB fits, but are widely adapted. We refer the reader to Zafar et al. (2015) and Krogager et al. (2015, 2016) for more discussions on AV determination for QSO. The requirement for the pairs of elements (Zn ii and Fe ii, or of S ii and Si ii) are in order to be able to derive depletions. Here, Zn and S are volatile elements and Fe and Si are refractory elements (e.g. Ledoux, Bergeron & Petitjean 2002; Draine 2003; Vladilo et al. 2011; De Cia et al. 2016). Defined this way our initial sample consists of 28 GRBs and 32 QSO-DLAs with the required measurements available. To this we add sources where only part of the required measurements are complete but where limits have been determined for the rest. The vast majority of the elemental abundances have been determined via detailed spectral line fitting, for eight GRBs elemental abundances (or limits) were derived from rest-frame equivalent widths of non-saturated lines provided by Fynbo et al. (2009) as described in Laskar, Berger & Chary (2011) and Zafar & Watson (2013, see references to Table 1).","Citation Text":["Perley et al. 2008"],"Functions Text":["Note that there is some degeneracy between broken power-law break frequency","and extinction, which could lead to inference of grey dust for some instances"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1000,1018]],"Functions Start End":[[800,875],[901,978]]}
{"Identifier":"2017AandA...607A.103G__Gradie_&_Tedesco_(1982)_Instance_1","Paragraph":"In Fig. 4, 29 observations of 15B-type asteroids are presented. As shown by Gil-Hutton & Ca\u00f1ada-Assandri (2012) and Gil-Hutton et al. (2014), in any plot of the observations of B-type objects a dispersion always appears in the whole phase angle range and it seems that these objects do not follow a single phase-polarization curve. The reason is probably that the asteroids classified by Bus & Binzel (2002) as members of this taxonomic class include objects that belong to the old F-class, originally proposed by Gradie & Tedesco (1982) and included in the Tholen taxonomy, due to its flat spectrum and low albedo; but several F-class objects have phase-polarization curves that are characterized by a comparatively low value of the inversion angle (Belskaya et al. 2005). Therefore, we also include in Fig. 4 nine observations of five asteroids classified as F-class by Tholen (1989) but included in other taxonomic types by Lazzaro et al.: (426) Hippo (X-type), (530) Turandot (C-type), (762) Pulcova (Cb-type), (778) Theobalda (C-type), and (877) Walkure (without a Bus taxonomy classification). The group of objects formed by (142) Polana, (213) Lilaea, (877) Walkure, and (1021) Flammario, all of them belonging to the old F-class, have observations indicating a low value of inversion angle, and the single observation of (314) Rosalia also indicates a similar polarimetric behavior even though it was not classified using the taxonomy of Tholen. On the other hand, the asteroids (62) Erato, (372) Palma, (635) Vundtia, and (981) Martina have measurements well below the theoretical phase-polarization curve for this taxonomic class, so it would seem to have a polarimetric behavior that corresponds more to a C-type than to a B-type. The apparent discrepancies might be indicative of some heterogeneity in the polarimetric properties within this particular taxonomic type, or alternatively, there might be some misclassification of the objects as could be the case for (635) Vundtia and (981) Martina, which were classified as C- and CFU-type by Tholen (1984). ","Citation Text":["Gradie & Tedesco (1982)"],"Functions Text":["The reason is probably that the asteroids classified by Bus & Binzel (2002) as members of this taxonomic class include objects that belong to the old F-class, originally proposed by","and included in the Tholen taxonomy, due to its flat spectrum and low albedo;"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[514,537]],"Functions Start End":[[332,513],[538,615]]}
{"Identifier":"2018ApJ...864..158L__Drake_et_al._2006_Instance_1","Paragraph":"However, if the magnetic curvature contraction term (last term in Equation (15)) is dominated by the compression term (second-to-last term in Equation (15)) so that \n\n\n\n\n\n, the contraction or merging is strongly compressible (\u2207 \u00b7  VE  0), resulting in an increase in flux-rope magnetic energy instead. In this case the plasma environment is doing positive mechanical work on the flux-rope structure, pushing magnetic field lines together to enhance flux-rope magnetic energy. Although it appears that small-scale flux ropes tend to contract or merge predominantly incompressibly in discussions of particle simulations (e.g., Drake et al. 2006; Dahlin et al. 2016), and it is also thought of as intrinsically incompressible in its manifestation as the quasi-2D turbulence component in N i MHD theory of solar wind turbulence (Zank et al. 2017), there is observational evidence to the contrary. For example, when primary current sheets associated with interplanetary coronal mass ejections (ICMEs) interact with the heliospheric current sheet, these structures are disturbed and several small-scale flux-rope structures may be formed when turbulent magnetic reconnection occurs in these structures. The flux ropes, being trapped between the converging heliospheric current sheet and the primary current sheets of ICMEs, experience compression, which may lead to efficient particle acceleration (e.g., Khabarova et al. 2015). However, it is possible that the particles are bounded in space because they cannot escape easily the region filled with small-scale flux ropes, which implies more efficient acceleration. Furthermore, in N i MHD theory of quasi-2D magnetic island turbulence, incompressible flux ropes can be compressed by large-scale density and flow velocity gradients in the nonuniform solar wind (Zank et al. 2017; see also discussion of Equation (69) in Section 8.2). Closer to the Sun, Guidoni et al. (2016) discuss the possibility of strong plasma compression during magnetic island contraction for islands propagating sunward during a solar flare event.","Citation Text":["Drake et al. 2006"],"Functions Text":["Although it appears that small-scale flux ropes tend to contract or merge predominantly incompressibly in discussions of particle simulations (e.g.,"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[625,642]],"Functions Start End":[[476,624]]}
{"Identifier":"2015ApJ...806..152S__Ransom_et_al._2005_Instance_1","Paragraph":"One of the most astonishing characteristics of Liller 1 is the extremely large value of the collision rate parameter. Verbunt & Hut (1987) showed that Liller 1 has the second-highest value of stellar encounter rate (after Terzan 5; see also Lanzoni et al. 2010) among all star clusters in the Galaxy, thus suggesting that it represents an ideal environment where exotic objects, generated by collisions, can form. In fact, it is commonly believed that dynamical interactions in GCs facilitate the formation of close binary systems and exotic objects such as cataclysmic variables (CVs), low-mass X-ray binaries (LMXBs), millisecond pulsars (MSPs), and blue straggler stars (BSSs) (e.g., Bailyn 1992; Paresce et al. 1992; Ferraro et al. 2001, 2009a, 2012; Ransom et al. 2005; Pooley & Hut 2006). Moreover, Hui et al. (2010) found that clusters with large collisional parameters and high metallicity (see also Bellazzini et al. 1995) usually host more MSPs. Indeed, Terzan 5 hosts the largest population of MSPs among all Galactic GCs (Ransom et al. 2005).\n6\n\n\n\n6\n\nNote that Terzan 5 is suspected to not be a genuine GC, because it harbors at least three stellar populations with different iron abundances (Ferraro et al. 2009a; Origlia et al. 2011, 2013; Massari et al. 2014).\n A strong \u03b3-ray emission has been recently detected in the direction of Liller 1 by the Large Area Telescope (LAT) on board Fermi (Tam et al. 2011). This is the most intense emission detected so far from a Galactic GC, again suggesting the presence of a large number of MSPs. However, no direct radio detection of these objects has been obtained so far in this system (Ransom et al. 2005). The only exotic object identified in the cluster is the rapid burster MXB 1730-335, an LMXB observed to emit radio waves and type I and type II X-ray bursts (Hoffman et al. 1978). It seems to be located in the central region of Liller 1, but no optical\/IR counterpart of this object has been found so far (Homer et al. 2001).","Citation Text":["Ransom et al. 2005"],"Functions Text":["In fact, it is commonly believed that dynamical interactions in GCs facilitate the formation of close binary systems and exotic objects such as cataclysmic variables (CVs), low-mass X-ray binaries (LMXBs), millisecond pulsars (MSPs), and blue straggler stars (BSSs) (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[755,773]],"Functions Start End":[[414,686]]}
{"Identifier":"2019ApJ...872..143B__Seckel_et_al._1991_Instance_2","Paragraph":"The gamma-ray emission from the solar disk due to CR cascades in the solar atmosphere is denoted as a disk component. This secondary gamma-ray produced by the hadronic interaction of cosmic ray with the solar surface was first proposed by Dolan & Fazio (1965). While only upper limits were obtained by early measurements over the range 20 keV\u201310 MeV (Peterson et al. 1966). A detailed theoretical model for gamma-rays from the collision of cosmic ray with the solar atmosphere was presented by Seckel et al. (1991). The predicted gamma-ray flux at energies from 10 MeV to 10 GeV has a large uncertainty, being sensitive to the assumptions about the cosmic-ray transport in the magnetic field near the Sun. Gamma rays from the Sun were first detected by the Energetic Gamma-ray Experiment Telescope (Orlando & Strong 2008). The measured flux from 100 MeV to 2 GeV was within the range of the theoretical predictions. The Fermi collaboration (Abdo et al. 2011) reported the detection of high energy gamma-rays at 0.1\u201310 GeV from the quiescent Sun using the first 1.5 yr data. However, the measured solar disk emission flux was about a factor of seven higher than that predicted about this disk component by a \u201cnominal\u201d model (Seckel et al. 1991). This mismatch motivated Ng et al. (2016) to analyze 6 yr of public Fermi-LAT data. The obtained gamma-ray spectrum follows a simple power-law shape (\u03b1 = \u22122.3) in 1\u2013100 GeV without any evident high energy cutoff. For the flux in 1\u201310 GeV, a significant time variation of the solar disk gamma-ray flux that anticorrelates with solar activity was discovered, suggesting that the solar magnetic field would play an important role. An updated analysis with 9 yr of Fermi-LAT data, from 2008 August 7 to 2017 July 27, was performed, and Tang et al. (2018) confirmed these results and extended the gamma-ray spectrum up to >200 GeV. Notably, the bright gamma-ray flux above 100 GeV is dominant only during solar minimum at the end of Cycle 23 (Linden et al. 2018). The HAWC measurements in periods of high solar activity may support these findings (Albert et al. 2018a). Data collected from 2014 November to 2017 December, the second half of solar cycle 24, have been used to set strong upper limits on the flux of 1\u2013100 TeV gamma-rays from the solar disk, about 10% of the maximum gamma-ray flux estimated by Linden et al. (2018). The HAWC 95% upper limit at 1 TeV is about 13% of the flux extrapolated from the solar minimum Fermi-LAT gamma-ray spectrum.","Citation Text":["Seckel et al. 1991"],"Functions Text":["However, the measured solar disk emission flux was about a factor of seven higher than that predicted about this disk component by a \u201cnominal\u201d model"],"Functions Label":["Differences"],"Citation Start End":[[1224,1242]],"Functions Start End":[[1074,1222]]}
{"Identifier":"2019ApJ...887..118C__Joshi_et_al._2011_Instance_1","Paragraph":"Generally, solar eruptions release their prestored magnetic energy via three phases, namely the precursor, impulsive, and gradual phases (Zhang et al. 2001; Zhou et al. 2016). Thereinto, the latter two, jointly termed as the main phase, correspond to the impulsive acceleration of the erupting CME flux rope, while the less-studied precursor phase includes key information on the eruptive structure and its trigger process. In the past, the limited observations showed that precursor activities of solar eruptions can come in various forms, in which preceding flux emergence (e.g., Palacios et al. 2015; Yan et al. 2017; Yang & Chen 2019) or cancellation (e.g., Green et al. 2011; Yang et al. 2016; Chen et al. 2018) in magnetograms, pre-eruption brightenings in H\u03b1\/EUV images (e.g., Bi et al. 2012; Dud\u00edk et al. 2016; Wang et al. 2017; Awasthi et al. 2018; Chen et al. 2018), as well as nonthermal processes in microwave or hard X-ray wavelength (Joshi et al. 2011; Altyntsev et al. 2012; Chen et al. 2017) are most common ones. In recent years, in the precursor phase of many CME\/flare eruptive events, a type of new progenitor of CME flux ropes, namely hot channels (HCs, Cheng et al. 2011, 2013; Zhang et al. 2012), have been detected in the AIA high-temperature passbands (e.g., 131 and 94 \u212b). These newfound features commonly exist in the low corona with obvious helical\/twisted fine fields or distinct writhed elbows (Cheng et al. 2013; Zhang et al. 2015), and they are found to be directly related to the occurrence of CMEs (Li & Zhang 2013; Patsourakos et al. 2013) and major eruptive flares (Nindos et al. 2015). To date, HCs have been evidenced as MFRs by some joint remote-sensing and in situ observations (e.g., Song et al. 2015), but for many HC eruption events, two problems are still elusive: (1) whether a corresponding MFR already exists before the HC eruption or is newly\/partially built up during the eruption; and (2) whether its loss-of-equilibrium is initially facilitated by the preflare reconnection beneath\/above the HC (Antiochos et al. 1999; Moore et al. 2001) or directly triggered the ideal magnetohydrodynamic (MHD) instabilities (T\u00f6r\u00f6k et al. 2004; Kliem & T\u00f6r\u00f6k 2006).","Citation Text":["Joshi et al. 2011"],"Functions Text":["In the past, the limited observations showed that precursor activities of solar eruptions can come in various forms,","as well as nonthermal processes in microwave or hard X-ray wavelength","are most common ones."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[948,965]],"Functions Start End":[[424,540],[877,946],[1008,1029]]}
{"Identifier":"2015MNRAS.454.1468K__Winckel_2003_Instance_3","Paragraph":"Owing to their dusty circumstellar environments, a large mid-infrared (mid-IR) excess is a characteristic feature of post-AGB stars and a detection of cold circumstellar material using mid-IR photometry can be used to identify these objects. The first extensive search for these objects was initiated in the mid-80's using results from the Infrared Astronomical Satellite (Neugebauer et al. 1984) which enabled the identification of post-AGB stars in our Galaxy (Kwok 1993). The Toru$\\acute{\\rm n}$ catalogue (Szczerba et al. 2007) for Galactic post-AGB stars lists around 391 very likely post-AGB objects. The Galactic sample of post-AGB stars have been found to be a very diverse group of objects (Van Winckel 2003). Studies showed that the majority of the optically visible Galactic post-AGB stars could be classified based on their spectral energy distributions (SEDs) into two groups: shell-sources and disc-sources (Van Winckel 2003). The shell-sources show a double-peaked SED with the hot central star peaking at shorter wavelengths while the cold, detached, expanding dust shell peaks at longer wavelengths. This type of SED is considered to be characteristic of objects that follow the single-star evolution scenario mentioned above. The disc-sources do not show two distinct flux peaks in the mid-IR but they do display a clear near-infrared (near-IR) excess indicating that circumstellar dust must be close to the central star, near sublimation temperature. It is now well established that this feature in the SED indicates the presence of a stable compact circumbinary disc, and therefore these sources are referred to as disc-sources (de Ruyter et al. 2006; Deroo et al. 2007; Gielen et al. 2011a; Hillen et al. 2013). The rotation of the disc was resolved with the ALMA array (Bujarrabal et al. 2013a) in one object and using single dish observations Bujarrabal et al. (2013b) confirmed that disc rotation is indeed widespread. Moreover, these disc-sources are confirmed to be binaries and show orbital periods between 100 and 2000 d (Van Winckel et al. 2009; Gorlova et al. 2014). In contrast, for the Galactic shell-sources long-term radial velocity monitoring efforts have not yet resulted in any clear detected binary orbit (Hrivnak et al. 2011), which either confirms the single-star nature of these objects or introduces a possibility that these systems can have companions on very wide orbits.","Citation Text":["Van Winckel et al. 2009"],"Functions Text":["Moreover, these disc-sources are confirmed to be binaries and show orbital periods between 100 and 2000 d"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[2050,2073]],"Functions Start End":[[1943,2048]]}
{"Identifier":"2019AandA...623A.140G__Dong_et_al._2018_Instance_1","Paragraph":"Planet formation occurs in disks around young stellar objects. Interactions between planets and disks are very complex. Young planets are expected to cause rings, cavities, spirals, and disturbances in the velocity field and other features in the disk, which in turn may be used to infer the presence of these young planets. In the past few years, much evidence about this phase of planet formation has been accumulated because high-resolution images in the millimeter and sub-millimeter wavelength ranges have been provided by the Very Large Array (VLA) and the Atacama Large Millimeter Array (ALMA; see e.g. the case of HL Tau; ALMA Partnership 2015), and by high-contrast imagers such as the Gemini Planet Imager (GPI, Macintosh et al. 2014) and SPHERE (Spectro- Polarimetic High contrast imager for Exoplanets REsearch, Beuzit et al. 2008; see, e.g., Avenhaus et al. 2018). The literature on indirect evidence of the presence of planets is now becoming very rich, and nearby young stars surrounded by gas-rich disks are intensively studied for this purpose. In most cases, available data cannot fully eliminate alternative hypotheses, or the data have ambiguous interpretations (see, e.g., Bae et al. 2018 and Dong et al. 2018), although strong indirect evidence of the presence of planets from local disturbances of the velocity field have recently been considered for the case of HD 163296 (Pinte et al. 2018; Teague et al. 2018). In general, small grains are thought to be more strongly coupled with gas and are thus less sensitive to radial drift and concentration that can strongly affect large grains (see the discussion in Dipierro et al. 2018). For this reason, observations at short wavelengths provide an important complementary view of what can be seen with ALMA. On the other hand, a direct detection of still-forming planets embedded within primordial gas-rich disks, which is expected to be possible with high-contrast imaging in the near infrared (NIR), is still scarce; remarkable cases are LkCa-15 (Kraus & Ireland 2012; Sallum et al. 2015) and PDS-70 (Keppler et al. 2018; M\u00fcller et al. 2018; Wagner et al. 2018). In particular, in this second case, a clear detection of an accreting planet in the cavity between the inner and outer ring was obtained, making it an archetype for planet formation and planet-disk interactions. However, many cases remain ambiguous; a classical example is HD 100546 (see, e.g., Quanz et al. 2013a, 2015; Currie et al. 2014, 2015; Rameau et al. 2017; Sissa et al. 2018).","Citation Text":["Dong et al. 2018"],"Functions Text":["The literature on indirect evidence of the presence of planets is now becoming very rich, and nearby young stars surrounded by gas-rich disks are intensively studied for this purpose. In most cases, available data cannot fully eliminate alternative hypotheses, or the data have ambiguous interpretations (see, e.g.,"],"Functions Label":["Background"],"Citation Start End":[[1214,1230]],"Functions Start End":[[878,1193]]}
{"Identifier":"2019AandA...623A..16S__Olano_&_Poeppel_1987_Instance_1","Paragraph":"Figure 10 (see also Figs. A.4 and A.5) compare the spatial distributions of the H\u03b1 and 857 GHz emission in the Taurus\u2013California\u2013Perseus region (e.g., Taurus, Auriga, California, and Perseus). The 857 GHz dust emission traces each molecular cloud and exhibits a hole-like structure. This hole-like structure can also be seen in HI emission, as shown in Figs. A.4 and A.5. The H\u03b1 emission fills the hole-like structure seen in the 857 GHz dust emission near the center of the field. The Taurus, California, and Perseus molecular complexes traced by the 857 GHz dust emission are distributed at the edge of the hole-likestructure. Lim et al. (2013) also found evidence of a shell-like structure using dust extinction and 12 CO (1\u20130) maps. The hole-like structure may result from the expansion of a large-scale supershell produced by a supernova in the Per OB2 association that compresses the Taurus cloud from the far side (Olano & Poeppel 1987; Bally et al. 2008). An H\u03b1 absorption feature is detected toward the Taurus cloud (see Figs. 10 and A.6), suggesting that the Taurus cloud lies at the front surface of the large-scale supershell produced by the Per OB2 association. The distance to the Per OB2 association is estimated to be 340 pc from the Sun (Cernis 1993), while the distance to the Taurus cloud is ~140 pc (Elias 1978). These distances are consistent with the Taurus cloud lying in front of the Per OB2 association. The B211\/B213 filament also appears to be in front of the HI shell (see Fig. 10 in Chapman et al. 2011). This morphology suggests that the B211\/B213 filament may have formed as a result of an expanding supershell. This may provide another reason for the different initial gas velocities for the northeastern and southwestern sheet components in addition to large-scale acceleration by the gravitational potential of the B211\/B213 cloud (see Sect. 4.2.3). The Local Bubble surrounding the Sun might also compress the Taurus cloud from the opposite direction. The Local Bubble surrounding the Sun was produced by supernovae (Snowden et al. 1998; Sfeir et al. 1999), and the wall of the Local Bubble is located close to the Taurus cloud (K\u00f6nyves et al. 2007; Lallement et al. 2014).","Citation Text":["Olano & Poeppel 1987"],"Functions Text":["The hole-like structure may result from the expansion of a large-scale supershell produced by a supernova in the Per OB2 association that compresses the Taurus cloud from the far side"],"Functions Label":["Background"],"Citation Start End":[[922,942]],"Functions Start End":[[737,920]]}
{"Identifier":"2019ApJ...883...73C__Ruffolo_et_al._2012_Instance_1","Paragraph":"Assuming that the force-field approach to the solution of the Parker (1965) cosmic-ray transport equation is valid, the connection between historic cosmic-ray intensities and the solar properties they encountered lies in the effective diffusion coefficient that is assumed in this approximation. Establishing such a connection, however, is no simple task. Many theories have been proposed to describe the scattering of cosmic rays in the heliosphere. The most likely candidates for this task, given their reasonable agreement with observations and numerical simulations of cosmic-ray diffusion coefficients, are the quasilinear theory (QLT) of Jokipii (1966), the weakly nonlinear theory (WNLT) of Shalchi et al. (2004b), and the nonlinear guiding center theory (NLGC; or one of its variants; see, e.g., Matthaeus et al. 2003; Shalchi 2006, 2009, 2010; Ruffolo et al. 2012). Shalchi (2009) provides in depth theoretical treatments of most the abovementioned theories. These scattering theories all require as a key input an expression for the power spectrum of the turbulent fluctuations of the HMF. These spectra depend upon basic turbulence quantities, such as the magnetic variance, and various correlation scales. Turbulence power spectra are discussed in detail by, e.g., Batchelor (1970) and Matthaeus et al. (2007), whereas more background on the abovementioned turbulence quantities can be found in, e.g., Matthaeus & Goldstein (1982), Petrosyan et al. (2010), Matthaeus & Velli (2011), and Bruno & Carbone (2013). These basic turbulence quantities have been observed to show a marked dependence on the solar cycle at Earth (see, e.g., Smith et al. 2006b; Burger et al. 2014; Zhao et al. 2018). It follows then that mean free paths derived from these scattering theories would be expected to depend on the solar cycle as well, and several studies have reported such a dependence. Chen & Bieber (1993) find from an analysis of cosmic-ray anisotropies and gradients as observed by means of NMs, that larger mean free paths are associated with solar minima, and smaller mean free paths with solar maxima. The authors also report a mean free path dependence on solar magnetic polarity. Nel (2016) and Zhao et al. (2018) both extensively analyze spacecraft observations, using the turbulence quantities so calculated as inputs for expressions for diffusion coefficients derived from the QLT and NLGC theories. Both authors report that the resulting mean free paths display solar cycle dependences.","Citation Text":["Ruffolo et al. 2012"],"Functions Text":["The most likely candidates for this task, given their reasonable agreement with observations and numerical simulations of cosmic-ray diffusion coefficients, are the quasilinear theory (QLT) of Jokipii (1966), the weakly nonlinear theory (WNLT) of Shalchi et al. (2004b), and the nonlinear guiding center theory (NLGC; or one of its variants; see, e.g.,"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[853,872]],"Functions Start End":[[451,803]]}
{"Identifier":"2018AandA...610A..44M__Kr\u00fcger_&_Dreizler_(1992)_Instance_3","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)"],"Functions Text":["The global fits included","29 lines from"],"Functions Label":["Uses","Uses"],"Citation Start End":[[1652,1676]],"Functions Start End":[[1569,1593],[1638,1651]]}
{"Identifier":"2018ApJ...852L..20A__Prieto_et_al._2015_Instance_1","Paragraph":"J0815+4729 is a main-sequence star (T\n\n\n\n\n\n\n\neff\n\n\n=\n6215\n\u00b1\n82\n\n\n K, \n\n\n\n\nlog\ng\n\n\n = 4.7\u00b10.5) with a metallicity of [Fe\/H] \u2264 \u22125.8 dex. Finding unevolved stars at this extremely low metallicity is very important since their stellar surface composition is not expected to be significantly modified by any internal mixing processes as in giant stars (Spite et al. 2005). J0815+4729 is similar to HE 1327\u20132326 in regard to its carbon enhancement, effective temperature, and metallicity. HE 1327\u20132326 is considered a turn-off\/subgiant star, while J0815+4729 appears to be a dwarf. The ISIS spectrum of HE 1327\u20132326 indicates a metallicity of [Fe\/H] \u223c \u22124.9 since the stellar Ca line is blended in that spectrum with the ISM features (Aguado et al. 2017b). However, the authors proposed a simple analysis taking into account the ISM absorption based on the UVES spectrum of HE 1327\u20132326. For J0815+4729, we require a high-resolution spectrum to clearly isolate the stellar Ca feature from possible additional ISM lines, and thus together with the detection of Fe lines, to establish the metallicity of this star. There are two other confirmed dwarf stars in this metallicity regime: one without any detectable carbon, J1029+1729 (Caffau et al. 2011), and another carbon-enhanced unevolved star, J1035+0641 (Bonifacio et al. 2015). The majority of extremely metal-poor stars shows overabundances of carbon, [C\/Fe] > 0.7, and it appears that carbon-enhanced metal-poor (CEMP) stars split into two groups, with dramatically different carbon abundances (see, e.g., Beers & Christlieb 2005; Allende Prieto et al. 2015; Bonifacio et al. 2015 and references therein). The two carbon bands (high and low) studied have different origins. On the one hand, CEMP stars in the high-carbon band (A(C) \u223c 8.2) are probably produced by mass transfer from a binary companion, most likely an AGB star (Starkenburg et al. 2014). On the other hand, objects lying in the low-carbon band (A(C) \u223c 6.8) are thought to show the original carbon abundance inherited by the star from the ISM (Stancliffe 2009; Bonifacio et al. 2015; Abate et al. 2016). J0815+4729 has an abundance ratio of \n\n\n\n\n\n\n[\n\nC\n\n\/\n\nFe\n\n]\n\n\n\u2265\n+\n5.0\n\n\n dex  corresponding to A(C) \u223c 7.7 dex (adopting [Fe\/H] \u2264 \u22125.8). In Figure 4 (bottom panel), we show the carbon abundance ratio [C\/Fe] for all stars at [Fe\/H]  \u22124.5. All stars in this metallicity regime are considered to belong to the low-carbon band (Bonifacio et al. 2015), except for J0815+4729, which appears to be in between the low- and high-carbon bands. Both metallicity and carbon abundance are considered upper and lower limits, respectively. High-resolution spectra would be very useful to measure other elemental abundances and investigate the properties of the first supernovae. In particular, the barium abundance, or that of any other s-element, is not measurable from ISIS or OSIRIS spectra, and this is required to determine whether J0815+4729 is a CEMP-s, CEMP-r, or i-process star (Hampel et al. 2016). If we establish the abundance pattern, we will learn about the progenitor properties. Finally, the radial velocity accuracy from medium-resolution data is not enough to discard variations among different exposures, which would be indicative of binarity.","Citation Text":["Allende Prieto et al. 2015"],"Functions Text":["The majority of extremely metal-poor stars shows overabundances of carbon, [C\/Fe] > 0.7, and it appears that carbon-enhanced metal-poor (CEMP) stars split into two groups, with dramatically different carbon abundances (see, e.g.,"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1579,1605]],"Functions Start End":[[1324,1553]]}
{"Identifier":"2022MNRAS.515.1086L__Naoz,_Farr_&_Rasio_2012_Instance_1","Paragraph":"Regarding non-restricted hierarchical three-body systems, Krymolowski & Mazeh (1999) and Ford, Kozinsky & Rasio (2000) presented secular equations of motion (or Hamiltonian) up to the octupole order in semimajor axial ratio by using Hamiltonian perturbation techniques. Lee & Peale (2003) adopted both the octupole-level perturbation theory and direct numerical integrations to investigate the dynamical evolution for coplanar hierarchical planetary systems. In particular, the dynamics of apsidal resonance with critical argument of \u03c3 = \u03d61 \u2212 \u03d62 (\u03d61, 2 are the longitudes of pericentre) is studied and applied to some representative exoplanetary systems (Lee & Peale 2003). In a hierarchical planetary system with two comparable-mass planets orbiting a central star, Naoz et al. (2011) showed that orbits of the inner planet could flip from prograde to retrograde and back again due to the secular planet\u2013planet interaction. Based on this behaviour, it becomes possible to form hot Jupiters on retrograde orbits by combining the eccentric ZLK effect and tidal friction (Naoz et al. 2011; Naoz, Farr & Rasio 2012; Teyssandier et al. 2013; Petrovich 2015; Petrovich & Tremaine 2016; Dawson & Johnson 2018). Naoz et al. (2013) re-derived the secular evolution equations for hierarchical three-body systems at the octupole-level approximation and found that orbital flips of inner planet are possible even at the quadrupole-level approximation. They pointed out that the relation h1 \u2212 h2 = \u03c0 can be used to simplify the expression of Hamiltonian but the evolutions of H1 and H2 should be derived from the conservation of the total angular momentum rather than from the Hamiltonian canonical relations. Tan et al. (2020) explored the secular resonances with critical arguments arising in the Hamiltonian under the resonant Hamiltonian model, which is obtained by directly removing those terms involving short-period angles from the octupole-level Hamiltonian (i.e. only the secular and resonant terms are retained). It is of no problem when dealing with the quadrupole-order resonance (the ZLK resonance) because in this case the omitting terms are of octupole order. However, it may be inadequate to formulate the resonant model by directly removing those quadrupole-order periodic terms from the Hamiltonian when studying the octupole-order resonances. Hamers (2021) performed a semianalytic study about the properties of the ZLK oscillations at the quadrupole-level approximation, including the maximum eccentricities, time-scales of eccentricity\/inclination oscillation and orbit flips. Naoz (2016) and Shevchenko (2016) reviewed various applications of the eccentric ZLK effect to a broad range of astrophysical systems, such as planetary and exoplanetary systems, stellar systems, and galaxies.","Citation Text":["Naoz, Farr & Rasio 2012"],"Functions Text":["Based on this behaviour, it becomes possible to form hot Jupiters on retrograde orbits by combining the eccentric ZLK effect and tidal friction"],"Functions Label":["Background"],"Citation Start End":[[1088,1111]],"Functions Start End":[[925,1068]]}
{"Identifier":"2018ApJ...854...26L___2015a_Instance_1","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"],"Functions Text":["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.,"],"Functions Label":["Background"],"Citation Start End":[[1023,1038]],"Functions Start End":[[694,1022]]}
{"Identifier":"2022ApJ...925...62K__Cohen_et_al._2020_Instance_1","Paragraph":"We noticed that the fast wind stream had He++ beams moving ahead of the protons, low number densities, and low heavy-ion charge state ratios, which are characteristic of coronal-hole-originated solar wind. It should be noted that although the solar wind stream originates from the same coronal hole, the speed and the stream arrival time to 1 au varied from one Carrington rotation to another as was seen in Figure 2. The timing of the coronal-hole-originated solar wind observation aligned well with the PSP orbit. PSP started passing through the leading edge of the high-speed stream, continued as the high-speed stream corotates over the spacecraft, and ended passing the trailing edge of the stream as seen in Figure 3, standing as a perfect classic CIR\/SIR event (Cohen et al. 2020). The decrease in crossing width and steepening of the increasing velocity profile at L1 versus at PSP is unlikely to be a product of the evolution of the coronal hole and is most likely a product of the faster wind stream overtaking the slower plasma at the interface during propagation (Burlaga 1974). Analyzing the SIR observed at ACE on 2018 November 4 to the corresponding SIR observed at PSP around 2018 November 15 indicated the enhancement of low-energy (keV) suprathermal ions (Figure 6) at the stream interface as the keV range ions are enhanced by local acceleration. The keV ions were enhanced again after the interface region, but for a short period of time compared to the PSP observation (Figure 8), while high-energy (MeV) suprathermal ions propagate along field lines to the inner heliosphere (Filwett et al. 2019) and are enhanced after the interface region in which articles accelerated at distant shocks dominated (see Desai et al. 2020; Joyce et al. 2021) and lasted longer in time in PSP. Minor enhancements in MeV particles were inconsistently observed over various Carrington rotations at ACE (not shown), but these were often at or just above the instrument noise floor. We cannot make any conclusive inferences as to the evolution of higher-energy ions in the coronal wind from these observations. In the future, as we approach solar maximum, it is likely that there will be more equatorial coronal holes observed during upcoming PSP orbits, which will allow us to compare more in depth the evolution of the solar wind structure at 1 au with PSP.","Citation Text":["Cohen et al. 2020"],"Functions Text":["The timing of the coronal-hole-originated solar wind observation aligned well with the PSP orbit. PSP started passing through the leading edge of the high-speed stream, continued as the high-speed stream corotates over the spacecraft, and ended passing the trailing edge of the stream as seen in Figure 3, standing as a perfect classic CIR\/SIR event"],"Functions Label":["Motivation"],"Citation Start End":[[769,786]],"Functions Start End":[[418,767]]}
{"Identifier":"2015AandA...580A...5L__Voit_1991_Instance_1","Paragraph":"The solutions in Sect. 3 show that it is possible for a high X-ray flux in galactic nuclei to alter the carbon ionization balance and reduce the C+ abundance and correspondingly the [C\u2009ii] luminosity. In the presence of a high flux of soft X-rays above 1 keV, a condition encountered in many AGNs (Stacey et al. 2010; Ebrero et al. 2009), the abundance of singly ionized carbon is reduced and converted into higher ionization states. This reduction occurs primarily in the hot highly ionized gas that fills most of the galactic central zone, to some degree in the dense ionized skins surrounding molecular clouds, less so in diffuse atomic hydrogen clouds, and very little, if at all, in the dense PDRs at the edge of the CO molecular cores. Furthermore, in galactic nuclei with a high X-ray flux we would expect to see an increase in the dust temperature and infrared luminosity (Voit 1991). Therefore in galactic nuclei, and in particular in AGNs, we expect a reduction in [C\u2009ii] emission relative to FIR\/IR emission depending on the relative contribution of different ISM components to the [C\u2009ii] luminosity. Unfortunately the next most abundant ion, C2+, does not have fine structure FIR emission lines as its 2s2 ground state has spin zero, and C3+, although it has spin angular momentum due to its unpaired electron in the 2S level, does not have a nuclear spin or orbital angular momentum to break the degeneracy of the two electron spin states. Instead to test the effect of X-ray ionization on the carbon balance we would need studies of their UV emission. For example the [C\u2009iv] UV resonance lines are detected in extragalactic sources and have been used to trace the formation rate of massive stars (Leitherer & Lamers 1991; Robert et al. 1993). [C\u2009iv] UV absorption lines have also been used to study the properties of the Galactic Halo (Savage et al. 2000) where carbon is presumed to be collisionally ionized in hot coronal gas (Gnat & Sternberg 2007). ","Citation Text":["Voit 1991"],"Functions Text":["Furthermore, in galactic nuclei with a high X-ray flux we would expect to see an increase in the dust temperature and infrared luminosity","Therefore in galactic nuclei, and in particular in AGNs, we expect a reduction in [C\u2009ii] emission relative to FIR\/IR emission depending on the relative contribution of different ISM components to the [C\u2009ii] luminosity."],"Functions Label":["Uses","Uses"],"Citation Start End":[[881,890]],"Functions Start End":[[742,879],[893,1111]]}
{"Identifier":"2022AandA...663A.172M___2012_Instance_1","Paragraph":"We note that Pavesi et al. (2019) derived a lower \u03bas parameter log(\u03bas)\u2004\u2248\u2004\u22121 for HZ10, by observing the CO(2-1) line, which implies a low star formation efficiency for this source. The conflict between the two results can be explained by the fact that Pavesi et al. (2019) estimated the gas mass by adopting a large CO-to-Mgas conversion factor \u03b1CO = 4.5 M\u2299 (K km s\u22121 pc2)\u22121, a value that is close to the Galactic conversion factor \u03b1CO = 4.36 M\u2299 (K km s\u22121 pc2)\u22121 (Bolatto et al. 2013). Although the Galactic conversion factor is a derived value for Milky Way and normal, star-forming galaxies in the local Universe, it may not be applicable for more extreme environments of starburst galaxies at high-z (see Carilli & Walter 2013 for a review). The conversion factor depends on the physical conditions of the gas in the ISM (temperature, surface density, dynamics, and metallicity), as well as the star formation and associated feedback (Narayanan et al. 2011, 2012; Genzel et al. 2012; Feldmann et al. 2012; Renaud et al. 2019; see, e.g., Bolatto et al. 2013 for a review). It is typically in the range between 0.8 and 4.36 M\u2299 (K km s\u22121 pc2)\u22121 (see, e.g., Bolatto et al. 2013; Carilli & Walter 2013, and Combes 2018 for reviews). Low metallicities (Z\u2004=\u20040.6\u2006Z\u2299 for HZ10) will drive \u03b1CO towards values higher than the Galactic value (Narayanan et al. 2012; Genzel et al. 2012; Popping et al. 2014), although \u03b1CO spans a broad range of values of \u03b1CO \u223c 0.4 \u2212 11 M\u2299 (K km s\u22121 pc2)\u22121, due to large uncertainties (see, e.g., Fig. 9 of Bolatto et al. 2013). On the other hand, high values of temperature, surface density, and velocity dispersion in a turbulent ISM of starbursts and merging systems will shift \u03b1CO towards lower values (Narayanan et al. 2011, 2012; Vallini et al. 2018). HZ10 has an extremely high value of the burstiness parameter log(\u03bas)\u2004\u223c\u20041.4 and high total density of the [C\u202fII] emitting gas log(n) \u223c 3.35 cm\u22123, and it is also a multi-component system (Jones et al. 2017, Carniani et al. 2018a). Thus, for this source we assumed \u03b1CO = 0.8 M\u2299 (K km s\u22121 pc2)\u22121, usually adopted for starburst galaxies (e.g., Downes & Solomon 1998; Bolatto et al. 2013). We obtained log(\u03bas)\u2004=\u20040.53\u2005\u00b1\u20050.34, which is within the 2\u03c3 uncertainties of the \n\n\n\nlog\n\n(\n\n\u03ba\ns\n\n)\n\n=\n1\n.\n\n43\n\n\u2212\n0.53\n\n\n+\n0.38\n\n\n\n\n$ \\log{(\\kappa_s)} = 1.43_{-0.53}^{+0.38} $\n\n\n, estimated exploiting the C\u202fIII] emission and the Vallini et al. (2020) model.","Citation Text":["Narayanan et al.","2012"],"Functions Text":["The conversion factor depends on the physical conditions of the gas in the ISM (temperature, surface density, dynamics, and metallicity), as well as the star formation and associated feedback"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[937,953],[960,964]],"Functions Start End":[[744,935]]}
{"Identifier":"2019ApJ...875...61M__Jones_&_Boffin_2017_Instance_1","Paragraph":"A substantial fraction of metal-poor stars that have recently evolved off the MS, e.g., giants and planetary nebulae (PNe), have been influenced by binary interactions. The IMF is significantly weighted toward low-mass stars (Bastian et al. 2010; Kroupa et al. 2013), and the MW star formation rate was \u22483 times larger \u224810 Gyr ago than it is now (Governato et al. 2007; De Lucia et al. 2014). Based on the measured IMF and modeled galactic star formation history, we estimate that \u224855% of MW giants and PNe have old, solar-type progenitors (\u03c4* > 7 Gyr, M \u2248 0.8\u20131.2 \n\n\n\n\n\n). Such old, low-mass giants tend to be metal-poor (Ratnatunga & Yoss 1991; Carollo et al. 2010; Mackereth et al. 2017). The metallicity trend therefore dramatically affects the properties of low-mass evolved stars. For example, the enhanced close binary fraction of metal-poor solar-type stars substantially strengthens the conclusion that the shaping of PN morphologies is the result of binary interactions (Moe & De Marco 2006; De Marco 2009; Jones & Boffin 2017). Providing further corroboration, Badenes et al. (2015) measured the delay-time distribution of bright PNe in the LMC and discovered two distinct populations of PN progenitors: an old channel (\u03c4* = 5\u20138 Gyr) deriving from solar-type stars (M \u2248 1.0\u20131.2 \n\n\n\n\n\n) and a young channel (35\u2013800 Myr) evolving from late-B\/early-A stars (\u22482\u20138 \n\n\n\n\n\n). According to the measured age\u2013metallicity relation of the LMC (Olszewski et al. 1991; Pagel & Tautvaisiene 1998; Cole et al. 2005; Carrera et al. 2011; Piatti & Geisler 2013), the old, solar-type progenitors are metal-poor ([Fe\/H] \u2272 \u22121.0) and hence have a large close binary fraction of Fclose = 40%\u201350%. The young progenitors have a higher metallicity of [Fe\/H] \u2248 \u22120.4 but are sufficiently massive so that they also have a large close binary fraction of Fclose = 40%\u201360%. Meanwhile, evolved stars with intermediate masses (M \u2248 1.2\u20132.0 \n\n\n\n\n\n) in the LMC have intermediate metallicities and therefore a smaller close binary fraction of Fclose \u2248 30%. If PNe derive from interactions in close binaries, then the variations in Fclose with respect to mass and metallicity can explain the observed bimodal mass\/age distribution of PN progenitors in the LMC.","Citation Text":["Jones & Boffin 2017"],"Functions Text":["For example, the enhanced close binary fraction of metal-poor solar-type stars substantially strengthens the conclusion that the shaping of PN morphologies is the result of binary interactions"],"Functions Label":["Similarities"],"Citation Start End":[[1017,1036]],"Functions Start End":[[787,979]]}
{"Identifier":"2019ApJ...875L..31H__Leary_et_al._2006_Instance_1","Paragraph":"The recent detection of gravitational-wave (GW) emission from a merging neutron star binary (Abbott et al. 2017d) and merging black hole binaries (BHBs; Abbott et al. 2016a, 2016b, 2017a, 2017b, 2017c; The LIGO Scientific Collaboration & The Virgo Collaboration 2018) by the Laser Interferometer Gravitational-Wave Observatory (LIGO)\/Virgo have ushered in an exciting new era of GW astrophysics. The astrophysical origin of the detected mergers is currently under debate, with numerous explanations proposed. These explanations can be very roughly divided into two main categories: mergers due to isolated binary evolution (e.g., Belczynski et al. 2016; de Mink & Mandel 2016; Mandel & de Mink 2016; Marchant et al. 2016), and mergers due to dynamical interactions (e.g., Portegies Zwart & McMillan 2000; Wen 2003; O\u2019Leary et al. 2006, 2009, 2016; Antonini & Perets 2012; Kocsis & Levin 2012; Antonini et al. 2014; Antonini & Rasio 2016; Rodriguez et al. 2016; VanLandingham et al. 2016; Askar et al. 2017; Arca-Sedda & Gualandris 2018; Fragione & Kocsis 2018; Hoang et al. 2018; Randall & Xianyu 2018; Arca-Sedda & Capuzzo-Dolcetta 2019). Orbital eccentricity has been explored as a way to distinguish between these merger channels in both the LIGO\/Virgo and Laser Interferometer Space Antenna (LISA) frequency bands. In contrast to mergers from isolated binary evolution, merging binaries from dynamical channels have been shown to have measurable eccentricities when they enter the LISA and\/or LIGO\/Virgo band, and can potentially be used as a way to distinguish between channels (e.g., O\u2019Leary et al. 2009; Cholis et al. 2016; Gond\u00e1n et al. 2018; Lower et al. 2018; Randall & Xianyu 2018; Rodriguez et al. 2018; Samsing 2018; Zevin et al. 2019). Unlike LIGO\/Virgo, which can only detect merging BHBs in the final inspiral phase before merger, LISA will be able to detect eccentric stellar-mass BHBs for long timescales before they merge in the LIGO\/Virgo band (e.g., O\u2019Leary et al. 2006; Breivik et al. 2016; Nishizawa et al. 2016; Chen & Amaro-Seoane 2017; Nishizawa et al. 2017; D\u2019Orazio & Samsing 2018; Kremer et al. 2019; Samsing & D\u2019Orazio 2018). This provides us with invaluable insight into the dynamical evolution of eccentric binaries leading up to the merger, which has important implications about the astrophysical context in which merging binaries evolve.","Citation Text":["O\u2019Leary et al. 2006"],"Functions Text":["These explanations can be very roughly divided into two main categories:","and mergers due to dynamical interactions (e.g.,"],"Functions Label":["Background","Background"],"Citation Start End":[[815,834]],"Functions Start End":[[509,581],[723,771]]}
{"Identifier":"2018MNRAS.479.4509R__Kingma_&_Ba_2014_Instance_2","Paragraph":"After each step of calculations, the network should optimize the model based on its current and previous states to improve the subsequent mapping. Our model utilizes a computationally memory efficient optimization due to its dependence to only the first-order gradients, namely the \u2018adaptive moment estimation\u2019 (or Adam). For more details, we refer the readers to Kingma & Ba (2014). Adam optimization, compared to other gradient-based optimization, is very suitable for noisy and sparse gradients, and for simulated data that show very large scatter with respect to a given quantity of parameter (Kingma & Ba 2014). With this optimizer, we have to decide few parameters in advance. The learning step \u03b1 and the parameters controlling the moving averages of the first- and second-order moments, namely \u03b21 and \u03b22 (both \u2208[0,1)), respectively. For this purpose, we chose to minimize the MSE between the target and the prediction from the model: in what follows, we will alternatively call the MSE the \u2018objective function\u2019 f($\\bf x$): with ${\\bf x}$ the parameters of the model to be updated, such as weights and biases. At a given time t \u2264 T, where T is the maximal learning time-step, we can update the parameters of the model as shown in the following:\n(11)\r\n\\begin{eqnarray*}\r\ng_t &amp;=\\nabla _\\mathrm{ \\text{$x$}} f(\\mathrm{\\text{$x$}}_{t-1}),\r\n\\end{eqnarray*}\r\n(12)\r\n\\begin{eqnarray*}\r\n\\mu _{1,t} &amp;=\\beta _1 \\times \\mu _{1,t-1} + (1-\\beta _1)\\times g_t,\r\n\\end{eqnarray*}\r\n(13)\r\n\\begin{eqnarray*}\r\n\\bar{\\mu }_{1,t} &amp;=\\mu _{1,t}\/(1-\\beta _1^t),\r\n\\end{eqnarray*}\r\n(14)\r\n\\begin{eqnarray*}\r\n\\mu _{2,t} &amp;=\\beta _2 \\times \\mu _{2,t-1} + (1-\\beta _2)\\times g_t^2,\r\n\\end{eqnarray*}\r\n(15)\r\n\\begin{eqnarray*}\r\n\\bar{\\mu }_{2,t} &amp;=\\mu _{2,t}\/(1-\\beta _2^t),\r\n\\end{eqnarray*}\r\n(16)\r\n\\begin{eqnarray*}\r\n\\mathrm{\\text{$x$}}_t &amp;=\\mathrm{\\text{$x$}}_{t-1} - \\alpha _t \\times \\bar{\\mu }_{1,t}\/ (\\sqrt{\\bar{\\mu }_{2,t}} + \\epsilon),\r\n\\end{eqnarray*}\r\nwhere $\\alpha _t=\\alpha \\sqrt{1-\\beta _2^t}\/(1-\\beta _1^t)$ is the time-step at t. Equation (11) shows the gradients of the objective function at t with respect to the model parameters. Equations (12) and (14) update the estimations of the first and second moments. Our moments are biased towards the initial values; thus, we require equations (13) and (15) to account for the corrections. Finally, we update the model parameters with equation (16).","Citation Text":["Kingma & Ba 2014"],"Functions Text":["Adam optimization, compared to other gradient-based optimization, is very suitable for noisy and sparse gradients, and for simulated data that show very large scatter with respect to a given quantity of parameter"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[598,614]],"Functions Start End":[[384,596]]}
{"Identifier":"2019ApJ...882..168P__Hildebrand_1983_Instance_1","Paragraph":"In order to provide constraints on the gas masses in these galaxies independently from the CO measurements, we can use the Rayleigh\u2013Jeans dust continuum emission. This will provide the first constraints to the \u03b1CO conversion factor in \u201cnormal\u201d galaxies at z > 3 in the following. The Rayleigh\u2013Jeans dust continuum emission has been used to estimate dust and gas masses, assuming an average emissivity and dust temperature for the dominant cold dust component and a constant dust-to-gas ratio (Hildebrand 1983; Eales et al. 2012; Bourne et al. 2013; Scoville 2013; Scoville et al. 2013, 2016, 2017; Groves et al. 2015). The dependence on cold dust temperature and dust-to-gas ratio may make the Rayleigh\u2013Jeans method less reliable than at lower redshifts (e.g., Pavesi et al. 2018a). On the other hand, the opposing effects of increasing dust temperatures and decreasing dust-to-gas ratios that may occur in \u201cnormal\u201d galaxies at high redshift may partially compensate for each other, as also found in recent simulations that are consistent with this approach to gas mass measurement (e.g., Liang et al. 2018; Privon et al. 2018). We here adopt Equations (10) and (13) of Scoville et al. (2016) to derive gas mass estimates based on our continuum flux measurements through the same assumptions that were used in those lower-redshift samples (Scoville et al. 2016, 2017). The 34 GHz upper limits imply 3\u03c3 gas mass limits of 2.8 \u00d7 1011 M\u2299 for HZ10 and 1.6 \u00d7 1011 M\u2299 for LBG-1, adopting the relation derived by Scoville et al. (2016, 2017). We also use the \u223c230 GHz continuum fluxes to derive approximate estimates, although these measurements may not lie on the Rayleigh\u2013Jeans tail and therefore may not accurately trace the cold dust component. These continuum measurements would imply gas masses of \u223c1.3 \u00d7 1010 M\u2299 for HZ4, \u223c2.5 \u00d7 1010 M\u2299 for LBG-1, \u223c4.4 \u00d7 1010 M\u2299 for HZ9, and \u223c1.1 \u00d7 1011 M\u2299 for HZ10, with dominant systematic uncertainties due to the extrapolation of the method to very high redshift.","Citation Text":["Hildebrand 1983"],"Functions Text":["The Rayleigh\u2013Jeans dust continuum emission has been used to estimate dust and gas masses, assuming an average emissivity and dust temperature for the dominant cold dust component and a constant dust-to-gas ratio"],"Functions Label":["Background"],"Citation Start End":[[493,508]],"Functions Start End":[[280,491]]}
{"Identifier":"2022ApJ...935..135B__Mathur_1990_Instance_2","Paragraph":"All responses calculated in this paper only account for the direct response to a perturbing potential. In general, though, the response also has an indirect component that arises from the fact that neighboring regions in the disk interact with each other gravitationally. This self-gravity of the response, which we have ignored, triggers long-lived normal-mode oscillations of the slab that are not accounted for in our treatment. Several simulation-based studies have argued that including self-gravity is important for a realistic treatment of phase spirals (e.g., Darling & Widrow 2019a; Bennett & Bovy 2021). Using the Kalnajs matrix method (Kalnajs 1977; Binney & Tremaine 2008), we have made some initial attempts to include the self-gravity of the response in our perturbative analysis, along the lines of Weinberg (1991). Our preliminary analysis shows that the self-gravitating response is a linear superposition of two terms: (i) a continuum of modes given in Equation (12), dressed by self-gravity, that undergo phase mixing and give rise to the phase spiral; and (ii) a discrete set of modes called point modes or normal modes (see Mathur 1990; Weinberg 1991) that follow a dispersion relation. The continuum response can be amplified by self-gravity when the continuum frequencies, n\u03a9\nz\n + kv\n\nx\n, are close to the point-mode frequencies, \u03bd. Depending on the value of k, the normal modes can be either stable or unstable. Araki (1985) finds that in an isothermal slab the bending normal mode undergoes fire hose instability below a certain critical wavelength if \u03c3\n\nz\n\/\u03c3 \u2272 0.3, while the breathing normal mode becomes unstable above the Jeans scale. In the regime of stability, the normal modes are undamped oscillation modes in absence of lateral streaming (Mathur 1990) but are Landau damped otherwise (Weinberg 1991). For an isothermal slab with typical MW-like parameter values, the point modes are strongly damped since their damping timescale (inverse of the imaginary part of \u03bd) is of order their oscillation period (inverse of the real part of \u03bd), which turns out to be of order the vertical dynamical time, h\n\nz\n\/\u03c3\n\nz\n. Moreover, the normal-mode oscillations are coherent oscillations of the entire system, independent of the vertical actions of the stars, and are decoupled from the phase spiral in linear theory since the full response is a linear superposition of the two. Based on the above arguments, we conclude that self-gravity has little impact on the evolution of phase spirals in the isothermal slab, at least in the linear regime. We emphasize that Darling & Widrow (2019a), who found their phase spirals to be significantly affected by the inclusion of self-gravity, assumed a perturber-induced velocity impulse with magnitude comparable to the local velocity dispersion in the solar neighborhood; hence, their results are likely to have been impacted by nonlinear effects. Moreover, the self-gravitating response of an inhomogeneous disk embedded in a DM halo, as in the simulations of Darling & Widrow (2019a), can be substantially different from that of the isothermal slab. We intend to include a formal treatment of self-gravity along the lines of Weinberg (1991) in future work.","Citation Text":["Mathur 1990"],"Functions Text":["In the regime of stability, the normal modes are undamped oscillation modes in absence of lateral streaming"],"Functions Label":["Uses"],"Citation Start End":[[1773,1784]],"Functions Start End":[[1664,1771]]}
{"Identifier":"2016ApJ...821...19F__Fulle_et_al._2010_Instance_1","Paragraph":"GIADA characterizes individual dust particles by means of two independent sensors. At the instrument entrance the particle crosses a laser curtain, and is detected by photoelectric sensors (GDS, grain detection system) registering a signal (proportional to the particle cross section times the albedo) and the time at which the laser curtain is crossed. Then the particle hits the impact sensor (IS, with the same GDS cross section, A = 10\u22122 m2), which registers the individual particle impact momentum and its travel time from GDS to IS. The combination of GDS and IS measurements (GDS+IS particles) provides the particle mass and velocity, and constrains the particle bulk density by means of calibration curves (Della Corte et al. 2016) derived on the ground using cometary analogues (Ferrari et al. 2014). If the particle is too small to be detected by the GDS system, it may be detected by the IS sensor only (IS particles): in this case the particle momentum is converted to the mass assuming the mean value of the velocities of the GDS+IS particles in the same momentum bin, or assuming the velocities predicted by tail models (Fulle et al. 2010) if Ngds+is = 0 in that mass bin. The spacecraft velocities listed in Table 1 are always much lower than the dust velocities measured by GIADA. In this condition, in the Sun-facing coma (assumed to have uniform and R-dependent space density \u03c1), the dust flux from the nucleus surface corresponds to the dust flux at nadir-pointing GIADA scaled by the factor 2\u03c0R2\/A. The dust number loss rate at the nucleus surface per GIADA detection is Qn = 2\u03c0R2(A\u0394t)\u22121, where \u0394t is the total dust collection time (Tables 5\u20138). In the same Tables, we show the mass loss rates Qm and the mean dust velocities already integrated in each mass bin, corresponding to the four GIADA collection periods considered in this paper: from 2015 February 19 to 28 (Table 5), from 2015 March 13 to 17 (Table 6), on 2015 March 28 (Table 7), and from 2015 August 23 to September 3 (Table 8). In Table 9 we show the data obtained during the first post-perihelion excursion at low phase angles (60  \u03b1  64\u00b0, 125  R  290 km). The R-values are too small to use the NAC DUST-MON sequences. The uncertainty affecting Af\u03c1 and the loss rates measured by GIADA and OSIRIS depends on the number of detections in each mass bin: an estimate of the relative error is given by \n\n\n\n\n\n\nN\n\n\np\n\n\n\u2212\n1\n\n\/\n\n2\n\n\n\n\n and by \n\n\n\n\n\n\n(\n\n\nN\n\n\ngds\n+\nis\n\n\n+\n\n\nN\n\n\nis\n\n\n)\n\n\n\u2212\n1\n\n\/\n\n2\n\n\n\n\n. The dispersion of the dust velocities in Tables 2\u20134 provides the error affecting the dust velocities measured by OSIRIS, close to 30%. The relative error of the dust velocities provided by each GDS+IS detection is below 10%.","Citation Text":["Fulle et al. 2010"],"Functions Text":["If the particle is too small to be detected by the GDS system, it may be detected by the IS sensor only (IS particles): in this case the particle momentum is converted to the mass assuming the mean value of the velocities of the GDS+IS particles in the same momentum bin, or assuming the velocities predicted by tail models","if Ngds+is = 0 in that mass bin."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1135,1152]],"Functions Start End":[[810,1133],[1154,1186]]}
{"Identifier":"2022MNRAS.512..186K__Dutta_&_Bharadwaj_2013_Instance_1","Paragraph":"A widely used statistical property of the sky brightness distribution is its power spectrum (Lazarian 1995; Bharadwaj & Sethi 2001, and others). As the redshifted 21-cm signal is expected to be faint and hard to detect with imaging, estimating its power spectrum or equivalently intensity mapping gives a possible probe of the evolution of the baryonic matter distribution over cosmic time. Bharadwaj & Sethi (2001) show that visibility correlation directly measures the power spectrum. This method and its variants (Datta, Choudhury & Bharadwaj 2007; Choudhuri et al. 2014; Choudhuri et al. 2016; Bharadwaj et al. 2019; Choudhuri et al.2019, and others) have been used to estimate the angular power spectrum of the diffused galactic foreground (Ghosh et al. 2012; Choudhuri et al. 2017b; Chakraborty et al. 2019a; Choudhuri et al. 2020) as well as the power spectrum of H\u2009i distribution in nearby galaxies (Dutta et al. 2009; Dutta & Bharadwaj 2013; Nandakumar & Dutta 2020). These works propagate the uncertainties in each visibility estimate and combine that with the sample variance error in measuring the power spectrum to quote uncertainties in the power spectrum estimates. In this work, we use the estimator discussed in Choudhuri et al. (2014), where visibilities are gridded before estimating the power spectrum. Given an angular field of view of \u03b80 to which the telescope is sensitive, it has been shown (Bharadwaj & Sethi 2001; Bharadwaj & Ali 2005; Choudhuri et al. 2014) that the visibilities in the nearby baselines remain correlated to a baseline separation of $\\Delta U \\lt \\frac{1}{\\pi \\theta _0}$. The size of the uv-grids is chosen such that they are large enough to include a sufficient number of baselines in a given uv-grid and small enough to have all visibilities in the uv-grid correlated. In each uv-grid, they estimate the power spectrum by correlating visibilities only in nearby baselines, omitting the visibility autocorrelations. This drastically reduces the noise bias in estimates of the power spectrum in uv-grids. The contribution from each uv-grid within a given annulus in $U = \\mid \\vec{U} \\mid$ is then combined, and the real part of it is used to quote the value of the isotropic power spectrum for the baseline separation U. We may schematically write it as\n(4)$$\\begin{eqnarray}\r\n\\mathcal {E} \\lbrace P(U)\\rbrace = \\mathcal {R} [\\langle \\tilde{V}(\\vec{U})^{*} \\tilde{V}(\\vec{U}+\\Delta \\vec{U}) \\rangle].\r\n\\end{eqnarray}$$Here, the average is taken over the uv-grid first and then within the annulus, as explained above. Note that the power spectrum estimator here assumes that a perfect calibration is done and the gains are all unity. In such a case, the power spectrum estimate has no bias arising from instrumental noise, and its uncertainties can be written as (Ali et al. 2008; Dutta 2011)\n(5)$$\\begin{eqnarray}\r\n\\sigma _P^2 = \\frac{P^2(U)}{N_\\mathrm{ G}} + 2\\frac{P(U)\\sigma _N^2}{N_\\mathrm{ B}} + 2\\frac{\\sigma _N^4}{N_\\mathrm{ B}},\r\n\\end{eqnarray}$$where NG is the number of independent estimates of the power spectrum in a given annulus bin at U, NB is the total number of visibility pairs in the bin.","Citation Text":["Dutta & Bharadwaj 2013"],"Functions Text":["This method and its variants","have been used to estimate","as well as the power spectrum of H\u2009i distribution in nearby galaxies","These works propagate the uncertainties in each visibility estimate and combine that with the sample variance error in measuring the power spectrum to quote uncertainties in the power spectrum estimates."],"Functions Label":["Background","Background","Background","Background"],"Citation Start End":[[927,949]],"Functions Start End":[[487,515],[655,681],[838,906],[977,1180]]}
{"Identifier":"2022ApJ...929...19G__Yu_et_al._2019_Instance_1","Paragraph":"From the theoretical point of view, we normalize the disk size to a fixed black hole mass and luminosity (thus to a fixed accretion rate; see Equations (14) and 17); therefore, we expect that the disk size \u03c4\n0 is independent of AGN properties. However, there are two possibilities that may cause the observed dependence in Figure 12. First, we adopt a constant virial factor f\nBLR when estimating the black hole mass using Equation (1). Dynamical modeling of broad-line regions (BLRs) generally showed that f\nBLR might depend on AGN properties and change from object to object (e.g., Pancoast et al. 2011; Li et al. 2013; Pancoast et al. 2014; Grier et al. 2017; Li et al. 2018; Williams et al. 2018). Observational calibrations of the virial factor also tend to support this conclusion (e.g., Ho & Kim 2014; Mej\u00eda-Restrepo et al. 2018; Yu et al. 2019). If there are systematic correlations between f\nBLR and luminosity L\n5100, black hole mass M\n\u2022, or accretion rate \n\n\n\nM\u0307\n\n, using a constant f\nBLR will lead to the apparent dependence of \u03c4\n0 on AGN properties. Such a bias can be eliminated once we have a solid understanding of the virial factor in the future. Second, it is possible that the disk\u2019s temperature profile does change with the accretion rate. For example, the standard and slim-disk models predict different temperature profiles\n12\n\n\n12\nHowever, we note that the temperature profiles of slim disks (T \u221d R\n\u22121\/2) become different from those of standard disks (T \u221d R\n\u22123\/4) only within the photon-trapping radius (e.g., Wang & Zhou 1999).  (Shakura & Sunyaev 1973; Abramowicz et al. 1988). There are also other physical processes that may contribute to the correlation between \u03b2 and \n\n\n\nM\u0307\n\n (see, e.g., Li et al. 2021; Kammoun et al. 2021). Figure 7 implies that \u03c4\n0 and \u03b2 are highly anticorrelated. As a result, a shallower temperature profile (larger \u03b2) will lead to a smaller disk size \u03c4\n0 and vice versa. In our sample, we find that high-luminosity and high-mass black holes generally have low accretion rates, giving rise to the dependence of \u03c4\n0 on luminosity and black hole mass.","Citation Text":["Yu et al. 2019"],"Functions Text":["Observational calibrations of the virial factor also tend to support this conclusion (e.g.,"],"Functions Label":["Similarities"],"Citation Start End":[[837,851]],"Functions Start End":[[702,793]]}
{"Identifier":"2022ApJ...939...43E__Protheroe_1999_Instance_1","Paragraph":"The multimessenger observations of NGC 1068 dictate the need for multiple emission zones. In this work, we will capture its inner microparsecs referring to a spherically symmetric structure for the corona of the AGN as well as an outer starburst ring with a radius of \u223c1 kpc (see Figure 1). Note, that in between these two emission sites NGC 1068 shows strong indications of a jet structure on scales of up to about 1 kpc (e.g., Wilson & Ulvestad 1982; Gallimore et al. 2004, 2006), which however is not included in this work. Due to mathematical convenience we treat both spatial regions as homogeneous. For particle acceleration processes that take place on considerably shorter timescales than the energy loss in these zones, we can disentangle these processes and only describe the steady-state transport of nonthermal, accelerated electrons and protons. Hereby, we suppose that in both zones some acceleration mechanism yields a differential source rate q(T) of relativistic protons and primary electrons that can be described by a power-law distribution in momentum space up to a certain maximal kinetic energy \n\n\n\nT\u02c6\n\n, which depends on the competing energy-loss timescales in these zones. In the case of the starburst zone, we suppose that a certain fraction \n\n\n\nfSN\n\n of the total energy of the SN that releases about 1051 erg and occurs with an approximate rate\n5\n\n\n5\nSupposing a SN rate \n\n\n\n\u03bdSN\u22430.02[SFR\/(1M\u2299\/yr)]yr\u22121\n\n (note that Condon1992 suggested a value of 0.04 instead of 0.02 for normal galaxies) where the star formation rate SFR \u2243 17[L\nIR\/(1011\nL\n\u2299)] M\n\u2299yr\u22121. (Veilleux et al. 2005) \n\n\n\n\u03bdSN\u22430.34[LIR\/(1011L\u2299)]yr\u22121\n\n dependent on the IR luminosity L\nIR gets accelerated into CRs according to diffusive shock acceleration (DSA; e.g., Drury 1983; Protheroe 1999) by individual supernova remnants (SNR; e.g., Bell 2014, and references therein). In general, many starburst galaxies\u2014NGC 253 is a prominent example\u2014show a galactic superwind (e.g., Veilleux et al. 2005, and references therein) as a result of a large number of core-collapse SNe. These winds introduce another source of acceleration\n6\n\n\n6\nNote that in these phenomena also stochastic diffuse acceleration may become relevant due to the presence of a turbulent plasma within the wind bubbles. (e.g., Anchordoqui et al. 1999; Romero et al. 2018), however, we are not aware of any observational indications of such a superwind in the starburst ring of NGC 1068. For the AGN corona, we suppose that a fraction f\ninj \u226a 1 of the mass accretion rate \n\n\n\nM\u0307=Lbol\/(\u03b7radc2)\n\n, with a radiation efficiency of \u03b7\nrad = 0.1 (Kato et al. 2008), goes into relativistic protons via stochastic diffuse acceleration (SDA; e.g., Lemoine & Malkov 2020, and references therein). For both zones, the nonthermal primary electrons are normalized by the nonthermal proton rates due to the requested quasi-neutral total charge number of the injection spectra of primary CRs above a characteristic kinetic energy of \n\n\n\nT\u02c7\u224310keV\n\n (Schlickeiser 2002; Eichmann & Becker Tjus2016; Merten et al. 2017). Note that this corresponds to a quasi-neutral acceleration site; however, CR transport can subsequently remove CR electrons and protons in different amounts from the nonthermal energy regime. Nevertheless their charge stays conserved. Transforming the source rates from momentum space into kinetic energy T, we obtain\n1\n\n\n\nqp(T)\u2261dNdVdTdt=qp,0T\u02c72+2T\u02c7Ep,0\u00d7T+Ep,0T2+2TEp,0T2+2TEp,0T\u02c72+2T\u02c7Ep,0\u2212s\/2exp[\u2212T\/T\u02c6],\n\nfor the injected nonthermal protons, and\n2\n\n\n\nqe\u00b1(T)=qe\u2212,0T\u02c72+2T\u02c7Ee,0T+Ee,0T2+2TEe,0\u00d7T2+2TEe,0T\u02c72+2T\u02c7Ee,0\u2212s\/2exp[\u2212T\/T\u02c6]+qe\u00b12nd(T),\n\nfor the nonthermal electrons (e\n\u2212) and positrons (e\n+). Here, the latter term \n\n\n\nqe\u00b12nd(T)\n\n introduces the source rate of secondary electrons and positrons that are generated by hadronic interaction processes, as discussed in the following. Thus, the steady-state behavior of the differential nonthermal electron and proton density n(T) in the AGN corona and the starburst zone, respectively, can be approximated by\n3\n\n\n\n\u2212\u2202\u2202TTn(T)\u03c4cool(T)=q(T)\u2212n(T)\u03c4esc(T).\n\nHere, \u03c4\n\ncool\n refers to the total continuous energy-loss timescale, which in the case of the relativistic electrons is given by the inverse of the sum of the synchrotron (syn), inverse Compton (IC), nonthermal Bremsstrahlung (brems), and Coulomb (C) loss rates, according to\n4\n\n\n\n\u03c4cool(e)=[\u03c4syn(e)\u22121+\u03c4ic\u22121+\u03c4brems\u22121+\u03c4C(e)\u22121]\u22121,\n\nand in the case of the relativistic protons we use\n5\n\n\n\n\u03c4cool(p)=[\u03c4syn(p)\u22121+\u03c4C(p)\u22121+\u03c4p\u03b3\u03c0\u22121+\u03c4BH\u22121+\u03c4pp\u22121]\u22121,\n\nincluding the photopion (\u03c0), Bethe\u2013Heitler pairs (BH), and hadronic pion (pp) production loss rates. Proton synchrotron losses\u2014as well as the associated radiation\u2014are negligible for the considered environments. Note that these processes require additional information on the associated interaction medium, which is one of the following targets:(i)A magnetic field, which is assumed to be uniform on small scales (with respect to the particles\u2019 gyro radius) and randomly oriented on significantly larger scales (due to isotropic Alfv\u00e9nic turbulence).(ii)A photon target, which is in the case of the starburst zone dominated by the thermal IR emission due to the rescattered starlight by dust grains with a temperature \u03b8\ndust and can be described by an isotropic, diluted modified blackbody radiation field\n6\n\n\n\nnIR(E)=Cdil\u03c02\u210fc3E2exp(E\/(kB\u03b8dust))\u22121EE0,\n\nwhere the dust clouds become optically thick above a critical energy E\n0 = 8.2 meV (Yun & Carilli 2002). The constant dilution factor C\n\ndil\n is determined from the observed IR luminosity L\nIR according to the relation \n\n\n\nLIR\/(\u03c0Rstr2c)=\u222bdEEnIR(E)\n\n.\n7\n\n\n7\nA more accurate approach to the IR photon spectrum has been proposed by Casey (2012), where a coupled modified greybody plus a mid-infrared power law has been used, but these modifications have no impact on our results. In case of the coronal region we used a parametrized model (Ho 2008) above 1 eV that accounts for the optical and UV emission by the disk as well as the Comptonized X-ray emission by hot thermal electrons in the corona. Hereby, the parameterization depends on the Eddington ratio (L\nbol\/L\nEdd), i.e., the ratio of the bolometric over the Eddington luminosity, and we adopt the relation of Hopkins et al. (2007) to determine L\nbol based on the intrinsic X-ray luminosity L\nX between (2 and 10)keV.(iii)A thermal gas target with a given temperature \u03b8, which is due to mathematical convenience assumed to be homogeneously distributed in both regions. For the starburst ring, Spinoglio et al. (2012) determine \u03b8 = 127 K, a gas density of n(H2) = 102.9 cm\u22123, and a molecular hydrogen mass of M(H2) \u223c 3.5 \u00d7 108\nM\n\u2299.More details on the individual energy-loss timescales can be found in Appendix A.","Citation Text":["Protheroe 1999"],"Functions Text":["according to diffusive shock acceleration (DSA; e.g.,"],"Functions Label":["Uses"],"Citation Start End":[[1765,1779]],"Functions Start End":[[1699,1752]]}
{"Identifier":"2017MNRAS.471.3057M__Bovy_et_al._2016b_Instance_2","Paragraph":"We have performed the first detailed dissection of the stellar populations of the Milky Way disc in age, [Fe\/H] and $\\mathrm{[ \\alpha \\mathrm{\/Fe]}}$ space, bridging the gap between the detailed observational understanding of MAPs (e.g. Bovy et al. 2012b, 2016b) and the plethora of studies of co-eval stellar populations in simulated galaxies (e.g Bird et al. 2013; Stinson et al. 2013; Martig et al. 2014a). We have placed novel constraints on models for the formation of the Milky Way disc by combining detailed density models fit to the mono-age, mono-[Fe\/H] populations of the low and high $\\mathrm{[ \\alpha \\mathrm{\/Fe]}}$ disc, with surface mass density contributions calculated on the basis of these density fits and stellar evolution models. We summarize our key results as follows:\nRadial and vertical profiles: The mono-age, mono-[Fe\/H] populations of the $\\mathrm{[ \\alpha \\mathrm{\/Fe]}}$ poor disc are well fitted by a radially broken exponential, with a peak radius, Rpeak, that varies as a function of age and [Fe\/H]. We find that the distance between Rpeak's of the low and high [Fe\/H] populations increases with age, which we interpret as evidence for a decreasing [Fe\/H] gradient with time (e.g. Anders et al. 2017). The radial variation of the stellar surface density of the high $\\mathrm{[ \\alpha \\mathrm{\/Fe]}}$ mono-age populations is found to have insignificant breaks, and they are better fit by a single exponential in this disc region. As these populations are the oldest, this may be a sign of the disc evolution washing out the density peak over time, or may point to a different formation scenario for high $\\mathrm{[ \\alpha \\mathrm{\/Fe]}}$ stars, where no density peak ever existed. These findings are in good agreement with earlier studies of MAPs (Bovy et al. 2016b). We measure an average high $\\mathrm{[ \\alpha \\mathrm{\/Fe]}}$ population scalelength of hR, in = 1.9 \u00b1 0.1\u2009kpc, and find scaleheights between 600 and 1000\u2009pc, in good agreement with current measures of the $\\mathrm{[ \\alpha \\mathrm{\/Fe]}}$ rich disc scalelength and scaleheight (e.g. those outlined in Bland-Hawthorn & Gerhard 2016).Profile broadening: We show that the radial surface density profile of the low $\\mathrm{[ \\alpha \\mathrm{\/Fe]}}$ populations broadens with age in a given [Fe\/H] bin, which we interpret as evidence of the gradual dispersal of mono-[Fe\/H] populations, presumably due to radial migration and radial heating. The variation in shape of the broken exponential profile changes differently depending on the population [Fe\/H], with low [Fe\/H] populations inner profiles flattening faster, whereas the high [Fe\/H] outer profiles flatten faster. We interpret this effect as tentative evidence for [Fe\/H] dependent radial migration arising from pre-existing [Fe\/H] gradients in the star-forming disc. We showed that our results qualitatively reproduce those of Hayden et al. (2015), finding a skewed MDF that varies as a function of R.Flaring: We find that flaring seems to be present in almost all mono-age populations, at differing levels. We have shown that the inverse flaring scalelength Rflare\u2212 1 increases with age, meaning that the youngest populations flare most strongly. This finding appears inconsistent with that above, under the assumption that flaring is the result of radial migration. However, these results may be reconciled by invoking a more active accretion history in the early life of the disc, which could have suppressed flaring (e.g. Minchev et al. 2014b).The surface-mass density at R0: We have measured the surface mass density at the solar radius for each mono-age, mono-[Fe\/H] population, finding a total surface mass density of $\\Sigma _{R_0, {\\rm tot}} = 20.0_{-2.9}^{+2.4}\\mathrm{(stat.)}_{-2.4}^{+5.0}\\mathrm{(syst.)}\\ \\mathrm{M_{{\\odot }} \\ pc^{-2}}$. Before allowing for systematics, this value is less than current estimates (e.g. Flynn et al. 2006; Bovy et al. 2012a; McKee et al. 2015), however, the systematic uncertainties are large, mainly due to a mismatch between the log\u2009g scales in APOGEE and the PARSEC models, and as such, we find our value to be consistent within the uncertainties. The relative contribution of high to low $\\mathrm{[ \\alpha \\mathrm{\/Fe]}}$ populations, $f_\\Sigma$, is 18\u2009per\u2009cent \u00b1 5\u2009per\u2009cent, which is consistent with existing measurements (e.g. Bland-Hawthorn & Gerhard 2016).The hZ distribution at R0: The shape of the mass-weighted hZ distribution found by this study is in good agreement with that of Bovy et al. (2012a), calling into question the existence of a vertical structural discontinuity in the Milky Way disc. The reconciliation of this finding with the discontinuity in chemical space (e.g. the bimodality in $\\mathrm{[ \\alpha \\mathrm{\/Fe]}}$ at fixed [Fe\/H]: Nidever et al. 2014; Hayden et al. 2015) may shed new light on our understanding of the formation of the Galactic disc.The surface-mass density profile of the Milky Way: We have found the combined (from mono-age, mono-[Fe\/H] populations at low and high $\\mathrm{[ \\alpha \\mathrm{\/Fe]}}$) surface-mass density-weighted profiles of the Milky Way disc as a function of $\\mathrm{[ \\alpha \\mathrm{\/Fe]}}$, age and [Fe\/H], and found that the total surface density is also described by a broken exponential. We find that our results fail to determine the sign of the inner exponential to high significance out to \u223c10\u2009kpc, but detect a turnover to a declining exponential, at high significance, thereafter. We find evidence of a radial mean age and [Fe\/H] gradient driven by the changing dominant population as a function of radius. A detailed comparison of these findings with numerical simulations is necessary for a proper interpretation. Our finding of a decline in stellar density may be consistent with that found in other studies (e.g. Reyl\u00e9 et al. 2009; Sale et al. 2010), albeit at shorter radii.","Citation Text":["Bovy et al. 2016b"],"Functions Text":["The radial variation of the stellar surface density of the high $\\mathrm{[ \\alpha \\mathrm{\/Fe]}}$ mono-age populations is found to have insignificant breaks, and they are better fit by a single exponential in this disc region. As these populations are the oldest, this may be a sign of the disc evolution washing out the density peak over time, or may point to a different formation scenario for high $\\mathrm{[ \\alpha \\mathrm{\/Fe]}}$ stars, where no density peak ever existed. These findings are in good agreement with earlier studies of MAPs"],"Functions Label":["Similarities"],"Citation Start End":[[1780,1797]],"Functions Start End":[[1235,1778]]}
{"Identifier":"2019MNRAS.487..364L__Kruijssen_et_al._2018_Instance_1","Paragraph":"It is well known that star formation is inefficient on galactic scales. The observed linear correlation between molecular gas surface density and star formation rate (SFR) surface density in normal star-forming galaxies suggests a nearly constant gas depletion time-scale around \u223c2\u2009Gyr, much longer than the dynamical time-scale of galactic discs (Kennicutt 1989; Bigiel et al. 2008; Saintonge et al. 2011; Leroy et al. 2013; Genzel et al. 2015; Tacconi et al. 2018). In contrast, the SFE on GMC scales shows a large variation ranging from less than a few per\u2009cent to nearly unity (Zuckerman & Evans 1974; Krumholz & Tan 2007; Wu et al. 2010; Evans, Heiderman & Vutisalchavakul 2014; Heyer et al. 2016; Lee, Miville-Desch\u00eanes & Murray 2016; Vutisalchavakul, Evans & Heyer 2016). The origin of this large scatter is usually explained as a combination of the time variability of the SFR during the course of cloud evolution and intrinsic scatter of SFEs due to the diversity of GMC properties (Feldmann & Gnedin 2011; Kruijssen & Longmore 2014; Lee et al. 2016; Grudi\u0107 et al. 2018a; Kruijssen et al. 2018). For example, recent theoretical models and high-resolution magneto-hydrodynamics simulations suggest that the SFE depends on the local virial parameters of the cloud controlled by large-scale turbulence (e.g. Krumholz & McKee 2005; Padoan, Haugb\u00f8lle & Nordlund 2012). However, it has recently been recognized that large-scale turbulence can only account for an \u223c0.3\u2009dex scatter, which is not sufficient to explain the observed SFE variations (Lee et al. 2016). Another source of variation comes from different stellar feedback channels that alter the dynamical states of the GMCs (Fall, Krumholz & Matzner 2010; Murray, Quataert & Thompson 2010; Dale et al. 2014; Myers et al. 2014; Raskutti, Ostriker & Skinner 2016; Kim et al. 2017; Grudi\u0107 et al. 2018b). Previous studies found that GMC simulations adopting different stellar feedback mechanisms (stellar winds, ionizing radiation, or supernovae) lead to dramatically different final SFEs. The problem has recently been recognized to be more subtle than previously thought, since even small differences in numerical treatments, such as different radiative transfer schemes, massive star sampling, and momentum and energy deposition algorithms, can lead to drastic changes for the final SFE (Dale et al. 2005; Ro\u0161kar et al. 2014; Raskutti et al. 2016; Grudi\u0107 et al. 2018b; Kim, Kim & Ostriker 2018). Therefore, how the SFE depends on GMC properties and the strength of stellar feedback remains an open question.","Citation Text":["Kruijssen et al. 2018"],"Functions Text":["The origin of this large scatter is usually explained as a combination of the time variability of the SFR during the course of cloud evolution and intrinsic scatter of SFEs due to the diversity of GMC properties"],"Functions Label":["Background"],"Citation Start End":[[1081,1102]],"Functions Start End":[[779,990]]}
{"Identifier":"2021ApJ...914L..38U__Dullemond_et_al._2001_Instance_1","Paragraph":"Snow lines of abundant volatile species are one of the possible origins of the observed substructures in dust continuum emission (Zhang et al. 2015; Okuzumi et al. 2016; Pinilla et al. 2017). Recent disk surveys have shown that the locations of the observed gap\/ring structures do not seem to be related to the radial locations of the snow lines (Huang et al. 2018; Long et al. 2018; van der Marel et al. 2019). To estimate the positions of the snow lines in observed disks, the temperature profile is often assumed to be a simple power law (Huang et al. 2018; Long et al. 2018), which is broken if the disk has shadows on the disk surface (Dullemond et al. 2001; Bailli\u00e9 & Charnoz 2014; Ueda et al. 2019). Figure 5 shows the radial positions of snow lines of H2O, NH3, CO2, H2S, C2H6, CH4, and CO at \n\n\n\n\n\n\nt\n\n\n\n\n\n=\n0\n\n\n and 11. For simplicity, we define the snow lines as the radial location where the midplane temperature reaches 150, 70, 55, 50, 40, 25, and 20 K for H2O, NH3, CO2, H2S, C2H6, CH4, and CO, respectively. We clearly see that the radial positions of snow lines move with time and multiple snow lines emerge even for a single species. This means that it is necessary to determine the disk temperature precisely when we evaluate the snow line locations. These snow lines would induce additional ring and gap structures and some of them would overlap with the TWI-induced substructures. It should be noted that the oscillation timescale of the TWI is much shorter than the dust radial drift timescale. Therefore, sintering-induced substructures would not coincide with the locations of the snow lines if the TWI is present. However, the dust-size variation can be induced by sublimation and recondensation, which potentially produces millimeter substructures. Since the radial locations of the snow lines are important not only for the substructure formation but also for the chemical composition of forming planets (e.g., Sato et al. 2016; \u00d6berg & Wordsworth 2019), we should investigate how the TWI evolves in planet forming disks.","Citation Text":["Dullemond et al. 2001"],"Functions Text":["To estimate the positions of the snow lines in observed disks, the temperature profile is often assumed to be a simple power law","which is broken if the disk has shadows on the disk surface"],"Functions Label":["Uses","Uses"],"Citation Start End":[[641,662]],"Functions Start End":[[412,540],[580,639]]}
{"Identifier":"2018AandA...615A.148D__Weidner_et_al._(2010)_Instance_1","Paragraph":"The last column in Table 1 reports the number of OB stars minus the \u201cdiffuse\u201d population estimated from their density in the Reference field (22.5 stars per square degree): as is immediately seen, the M-star statistics is much larger than the OB star statistics. This can hardly be considered surprising, if an ordinary IMF (e.g., that from Weidner et al. 2010) is assumed for the star-formation region. In Fig. 19 we show the density ratio between M and OB stars, which provides a consistency test between our results and a plausible IMF: this ratio varies however by a large factor, close to 20, among our subregions. This might reflect differences in the respective IMFs, but also differences in completeness among the stellar samples considered for the various regions. We first note that the ratio between M and OB stars in NGC 6231 is dramatically lower than anywhere else in Sco OB1. We can indeed expect that M stars are detected less efficiently in the inner parts of NGC 6231, where the density of bright stars is very large, and their diffuse glare raises the limiting magnitude locally. As already discussed above in Sect. 4.1, this causes our sample of M stars in NGC 6231 to be highly incomplete. Moreover, we determined above that NGC 6231 is significantly more extincted, by almost half a magnitude in V, than Tr 24, and this implies a higher minimum detectable mass among NGC 6231 M stars compared to Tr 24 (see the MDA diagrams in Fig. 5); this effect reduces the completeness of the M-star sample in NGC 6231 more than in Tr 24. If Tr 24 is also slightly younger than NGC 6231, as we argue below, our M-stars in Tr 24 will reach down to lower masses than in NGC 6231, with a steep increase in the detected M-star population: adopting the IMF from Weidner et al. (2010), the predicted number of cluster M stars doubles considering the mass interval 0.25\u20130.5 M\u2299 rather than 0.35\u20130.5 M\u2299. If Tr 24 is younger than NGC 6231, moreover, its stars in the mass range 2.5\u20133 M\u2299 might not have yet reached their ZAMS position as B stars, and therefore would not be counted among OB stars; this would further raisethe M\/OB star ratio there by up to 30%. Therefore, the proportions of both M and OB stars that are detected in a young cluster will depend on their age and extinction, in accordance with the MDA diagrams, even for a fixed, spatially uniform photometric sensitivity. We estimated using the Weidner et al. (2010) IMF the expected range for the observed M\/OB number ratio. Siess et al. (2000) predict that the latest-type B stars have a mass of ~ 3.5 M\u2299 at 2 Myr, and ~ 2.2 M\u2299 at 10 Myr, that is, in the range of ages expected for Sco OB1 clusters. The MDA diagrams of Fig. 5 predict that the lowest-mass stars we are able to detect using the available Sco-OB1 data have ~ 0.2 M\u2299, even assuming the most favorable (and unlikely) circumstances of an age less than 2 Myr and negligible reddening. The extreme values found for the M\/OB ratio are then ~ 3.8 for a minimum M-star mass as high as 0.35 M\u2299 and an old age of 10 Myr, and ~ 20 for a minimum M-star massas low as 0.2 M\u2299 and age of 2 Myr. These extremes are also shown as horizontal lines in Fig. 19. We note that the M\/OB ratio in NGC 6231 falls well within this range; however, both Tr 24 regions are significantly richer of M stars than expected, by more than a factor of two and well above (statistical) errors. If true, then paradoxically this part of the OB association would form preferentially lower-mass stars. Of course, more detailed studies are needed to confirm this result. In the G345.45+1.50 region the M\/OB ratio is highest, and far above predictions from the IMF: we may tentatively explain this since this region is very young, and some of its most massive members, like IRAS 16562-3959, are still in formation, thus decreasing the number of optically revealed OB stars. The lowest M\/OB ratio in NGC 6231 is unlikely to be real, since as discussed above our M-star sample in this densest subregion is likely incomplete.","Citation Text":["Weidner et al. 2010"],"Functions Text":["The last column in Table 1 reports the number of OB stars minus the \u201cdiffuse\u201d population estimated from their density in the Reference field (22.5 stars per square degree): as is immediately seen, the M-star statistics is much larger than the OB star statistics. This can hardly be considered surprising, if an ordinary IMF (e.g., that from","is assumed for the star-formation region."],"Functions Label":["Uses","Uses"],"Citation Start End":[[341,360]],"Functions Start End":[[0,340],[362,403]]}
{"Identifier":"2022MNRAS.511.6218Z__Malkov,_Diamond_&_Sagdeev_2012_Instance_1","Paragraph":"To further reduce the number of fitting parameters, we assume that no source is unique in terms of the acceleration spectral shape, i.e. all the source components have the same H spectral index \u03b1, He spectral index, maximum rigidity Rm, and He-to-H Galactic injection flux ratio \u03c7. The spectrum of He seems harder than that of H as indicated by many spectral studies, implying, e.g. time-dependent particle acceleration (Zhang et al. 2017), or more efficient He injection in the shock acceleration process (Malkov, Diamond & Sagdeev 2012). It is thus assumed the He spectral index to be \u03b1 \u2212 0.077 according to the sub-TeV data fit (Aguilar et al. 2015b). To be consistent with the conventional scenario, we also consider that the nearby SNR can inject a total CR energy of 1050 erg into the ISM, i.e.\n(15)$$\\begin{eqnarray}\r\n\\int _{\\text{GeV}}^{\\infty }{\\left[ \\frac{{\\rm d}R_{\\text{H}}}{{\\rm d}E}+\\chi \\left(E \\right) \\frac{{\\rm d}R_{\\text{He}}}{{\\rm d}E} \\right] N\\left(E \\right) E{\\rm d}E}=10^{50}\\text{ erg},\r\n\\end{eqnarray}$$where E is the particle kinetic energy, N refers to the injection spectrum of H with $N\\left( \\text{GV}\\right) \\approx 4.7\\times 10^{52}\\text{ GV}^{-1}$ (for the fitted parameters in Table 1). For brevity, we hereafter treat all of N and Q as quantities for H components. With these respects, in the naive approach, the spectral fitting parameters are $\\chi \\left( \\text{GV}\\right)$, \u03b1, Rm, \u03d5, \u03b8, MA, $\\left| z\/h \\right|$, $\\mu Q_{\\text{d}}\\left( \\text{GV}\\right)$, $Q_{\\text{c}}\\left( \\text{GV}\\right)$, tc, where $\\theta =\\arcsin \\left(\\rho \/r_{\\text{n}} \\right)$ is the angle between the regular magnetic field line and the line-of-sight vector towards the nearby source. In the quasi-local approach equation (11), the alignment of the field line must uniquely be determined, leading to an additional angular parameter, which can be chosen as the angle \u03c6 around the line-of-sight vector (towards the nearby source with \u03c6 = 0\u00b0 shown in Fig. 2d). The fitting results of the CR spectrum for models with the naive and quasi-local approach are shown with Figs 1(a) and 2(a), respectively.","Citation Text":["Malkov, Diamond & Sagdeev 2012"],"Functions Text":["The spectrum of He seems harder than that of H as indicated by many spectral studies, implying, e.g. time-dependent particle acceleration","or more efficient He injection in the shock acceleration process"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[507,537]],"Functions Start End":[[282,419],[441,505]]}
{"Identifier":"2021AandA...655A..99D__Carigi_et_al._2005_Instance_3","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"],"Functions Text":["The turning point of increased relative production of carbon from massive stars takes place at A(O) ~ 8.7 dex (see Fig. 2 of","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)."],"Functions Label":["Uses","Uses","Differences"],"Citation Start End":[[2266,2284]],"Functions Start End":[[2141,2265],[2286,2318],[2319,2469]]}
{"Identifier":"2021MNRAS.507.2766S__Sumiyoshi_et_al._2005_Instance_1","Paragraph":"In order to make a linear analysis, first we have to prepare the PNS models as a background. The PNS properties depend on not only the density and pressure profiles but also the distributions of temperature (or entropy per baryon) and electron fraction inside the PNS, while such profiles can be determined only via the numerical simulation of the core-collapse supernova explosion. In this study, as in Sotani & Sumiyoshi (2019), we particularly adopt the profiles obtained via the numerical simulations performed by solving the general relativistic neutrino-radiation hydrodynamics under the spherical symmetry. In the simulations, hydrodynamics and neutrino transfer in general relativity are solved simultaneously (Yamada 1997; Yamada, Janka & Suzuki 1999; Sumiyoshi et al. 2005). To describe the neutrino transfer, the Boltzmann equation is directly solved with the multi-angle and multi-energy neutrino distributions for four species, \u03bde, $\\bar{\\nu }_\\mathrm{ e}$, \u03bd\u03bc\/\u03c4, and $\\bar{\\nu }_{\\mu \/\\tau }$, i.e. we implement six species of neutrinos by assuming \u03bc-type and \u03c4-type (anti-)neutrinos have identical distributions. For the collision term associated with neutrino emission, absorption, and scattering with leptons, nucleons, and nuclei handles, the basic neutrino reactions are adopted (Bruenn 1985; Sumiyoshi et al. 2005). The metric adopted in the numerical code is given by\n(1)$$\\begin{eqnarray*}\r\n\\mathrm{ d}s^2 = -\\mathrm{ e}^{2\\Phi (t,m_\\mathrm{ b})}\\mathrm{ d}t^2 + \\mathrm{ e}^{2\\Lambda (t,m_\\mathrm{ b})}\\mathrm{ d}m_\\mathrm{ b}^2 + r^2(t,m_\\mathrm{ b})(\\mathrm{ d}\\theta ^2 + \\sin ^2\\theta \\mathrm{ d}\\mathrm{ }\\phi ^2), \\nonumber\\\\\r\n\\end{eqnarray*}$$where t and mb denote the coordinate time and the baryon mass coordinate, respectively (Misner & Sharp 1964). In addition, mb is related to the circumference radius (r) via the baryon mass conservation, while the metric functions, \u03a6(t, mb) and \u039b(t, mb), are evolved together with hydrodynamical variables in the numerical simulations (Yamada 1997). The numerical simulations for core-collapse supernovae have been done with 255 grid points in the radial mass coordinate, 6 grid points in the neutrino angle, and 14 grid points in the neutrino energy. The rezoning of radial mesh is made during the simulations to resolve the accreting matter. We remark that the radial grids of mass coordinate are non-uniformly arranged to cover not only the dense region inside the central object but also the region for accreting matter.","Citation Text":["Sumiyoshi et al. 2005"],"Functions Text":["In the simulations, hydrodynamics and neutrino transfer in general relativity are solved simultaneously"],"Functions Label":["Uses"],"Citation Start End":[[761,782]],"Functions Start End":[[614,717]]}
{"Identifier":"2019MNRAS.487.1210T__McNamara_&_Nulsen_2007_Instance_3","Paragraph":"On larger scales, the clusters in which BCGs reside can generally be divided into two categories: cool core clusters, which exhibit very peaked surface brightness distributions at X-ray wavelengths, and non cool core clusters, with similar overall X-ray luminosities but with smoother, less peaked X-ray surface brightness distributions. Some authors (e.g. Hudson et al. 2010; Santos et al. 2010) define an intermediate category called moderate or weak cool core clusters. Since cool core clusters have short radiative cooling time-scales on the order of 108 yr in their centres (e.g. Voigt & Fabian 2004; McNamara & Nulsen 2007, 2012; Hlavacek-Larrondo et al. 2012), starbursts are expected to be common at the centre of such clusters. Indeed, the central cool gas in these clusters should condense onto the BCG, forming stars at rates of hundreds of solar masses per year (e.g. Fabian 1994). However, most BCGs are relatively quiescent and those that do show evidence of star formation generally tend to have star formation rates 1 order of magnitude smaller, on the order of $1-150 \\, \\mathrm{M_{\\odot }\\, {yr}^{-1}}$ (e.g. Donahue et al. 2007; Bildfell et al. 2008; O\u2019Dea et al. 2008, 2010; Rawle et al. 2012). This mismatch between expected and observed star-forming rates, known as the cooling flow problem, is thought to be caused by active galactic nuclei (AGNs) feedback processes from the BCG. AGNs can release copious amounts of energy into the intracluster medium (ICM) through many ways, including: jetted outflows that inflate cavities, weak shocks, sound waves, or turbulence in the ICM (e.g. Markevitch & Vikhlinin 2007; McNamara & Nulsen 2007, 2012; Zhuravleva et al. 2014; Fabian et al. 2017). Alone, the energy released by jetted outflows appears to be on the same order as the energy needed to offset cooling (e.g. Rafferty et al. 2006; McNamara & Nulsen 2007; Hlavacek-Larrondo et al. 2012), therefore suggesting that AGN feedback is a good candidate for solving the cooling flow problem.","Citation Text":["McNamara & Nulsen 2007"],"Functions Text":["Alone, the energy released by jetted outflows appears to be on the same order as the energy needed to offset cooling (e.g.","therefore suggesting that AGN feedback is a good candidate for solving the cooling flow problem."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1857,1879]],"Functions Start End":[[1712,1834],[1913,2009]]}
{"Identifier":"2020MNRAS.499.2575E__Pontzen_&_Governato_2014_Instance_1","Paragraph":"We note in passing that recent studies address improved satellite modellimg that ameliorates many of these issues, including the core\u2013cusp issue via non-sphericity of the stellar velocity distribution (Hayashi, Chiba & Ishiyama 2020) and the detectability of MWG satellites (Nadler et al. 2020). Other proposed solutions include those invoking baryonic physics, ranging from inclusion of baryon-contraction-induced diversity (Lazar et al. 2020), through dynamical friction-mediated coupling with baryonic clumps (El-Zant, Shlosman & Hoffman 2001; El-Zant et al. 2004; Tonini, Lapi & Salucci 2006; Romano-D\u00edaz et al. 2008; Goerdt et al. 2010; Cole, Dehnen & Wilkinson 2011; Del Popolo et al. 2014; Nipoti & Binney 2015), or through dynamical feedback driven by starbursts or active galactic nuclei (AGNs; Read & Gilmore 2005; Mashchenko, Couchman & Wadsley 2006; Mashchenko, Wadsley & Couchman 2008; Peirani, Kay & Silk 2008; Governato et al. 2012; Pontzen & Governato 2012; Zolotov et al. 2012; Martizzi, Teyssier & Moore 2013; Teyssier et al. 2013; Madau, Shen & Governato 2014; Ogiya & Mori 2014; Pontzen & Governato 2014; El-Zant, Freundlich & Combes 2016; Silk 2017; Freundlich et al. 2020). Alternatively, modifications to the particle physics model of the dark matter have been proposed. Such proposals include \u2018pre-heated\u2019 warm dark matter (e.g. Col\u00edn, Avila-Reese & Valenzuela 2000; Bode, Ostriker & Turok 2001; Macci\u00f2 et al. 2012; Schneider et al. 2012; Shao et al. 2013; Lovell et al. 2014; El-Zant, Khalil & Sil 2015) and self-interacting dark matter, whereby energy flows into the central cores of haloes through conduction (e.g. Burkert 2000; Kochanek & White 2000; Spergel & Steinhardt 2000; Miralda-Escud\u00e9 2002; Peter et al. 2013; Zavala, Vogelsberger & Walker 2013; Elbert et al. 2015). Ultralight axions, with boson mass \u223c10\u221222 eV, have also been considered as dark matter candidates in connection with these same small (sub)galactic scale problems (e.g. Peebles 2000; Hu, Barkana & Gruzinov 2000; Peebles 2000; Marsh & Silk 2014; Schive et al. 2014b; Hui et al. 2017; Mocz et al. 2019; Nori et al. 2019; see Niemeyer 2019 for recent review). Here the zero-point momentum associated with a long de Broglie wavelength corresponding to the small mass comes along with \u2018fuzziness\u2019 in particle positions. This in turn leads to a hotter halo core with non-diverging central density and a cut-off in halo mass. Such axion fields can also be relevant for inflationary scenarios or late dark energy models. The non-thermal production implies that the axions are present with the required abundance for dark matter; they behave as cold dark matter on larger scales despite the tiny masses (Marsh 2016, 2017).","Citation Text":["Pontzen & Governato 2014"],"Functions Text":["Other proposed solutions include","or through dynamical feedback driven by starbursts or active galactic nuclei"],"Functions Label":["Background","Background"],"Citation Start End":[[1099,1123]],"Functions Start End":[[296,328],[720,796]]}
{"Identifier":"2022ApJ...940L..18Z__Perna_et_al._2022_Instance_1","Paragraph":"In addition to some fraction of BNS mergers masquerading as long GRBs, our sample used to constrain the DTD may suffer from other issues of incompleteness. As we rely on the modeling of host galaxies when constraining the DTD, we do not consider short GRBs that do not have a confident host association. Though the properties of short GRB hosts do not seem to deviate strongly as a function of the host-association confidence (see, e.g., Figure 4 of Nugent et al. 2022), neglecting these events may have a potential impact on both the low and high ends of our inferred DTD. For example, GRBs that are highly offset from their hosts may have afterglows with much lower luminosities, making precise localization (and therefore host identification) difficult (Perna et al. 2022). Such systems may have migrated over longer timescales to reach the highly offset locations of the burst and therefore may have longer delay times than the general population. Furthermore, the P\ncc method for host identification may incorrectly associate a GRB with a faint underlying host rather than a bright host at a larger offset, though Fong et al. (2022) predicted this to be an effect only at the \u22727% level. On the other hand, if such poorly associated GRBs are instead truly associated with faint galaxies that are below detection limits, we may be excluding additional systems with short delay times as these faint, low-mass galaxies are typically star-forming. Furthermore, though Swift can detect GRBs out to z \u223c 3, there is likely some fraction of short GRBs that occur beyond this horizon, when the universe was \u22722 Gyr old. Short GRBs that occur at these early stages in the history of the universe must have short delay times, and this selection effect may bias the general population in our analysis to longer delay times. This would lead to a larger inferred \n\n\n\ntmin\n\n, and, due to the correlation between \n\n\n\ntmin\n\n and \u03b1, more negative values of \u03b1. However, this population of high-redshift short GRBs is likely small; assuming the SFH from Madau & Fragos (2017), 10% of stars are born beyond z = 3, and the fraction of compact object binary mergers beyond this redshift will be even smaller due to the delay time between formation and merger.","Citation Text":["Perna et al. 2022"],"Functions Text":["For example, GRBs that are highly offset from their hosts may have afterglows with much lower luminosities, making precise localization (and therefore host identification) difficult"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[757,774]],"Functions Start End":[[574,755]]}
{"Identifier":"2021AandA...653A.156F__Valentini_et_al._2007_Instance_1","Paragraph":"In general, hybrid-kinetic models have proved to be capable of satisfactorily catching the main kinetic physics at play for a large number of problems, ranging from fluid and kinetic instabilities (Hellinger & Matsumoto 2000; Matteini et al. 2006; Califano et al. 2008; Henri et al. 2013; Kunz et al. 2014) to collisionless shocks (Lemb\u00e8ge et al. 2009; Caprioli & Spitkovsky 2013; Weidl et al. 2016), dynamo effects (Rincon et al. 2016; St-Onge & Kunz 2018), MR (Le et al. 2016; Palmroth et al. 2017; Cerri & Califano 2017; Franci et al. 2017; Wang et al. 2019; Califano et al. 2020), and kinetic-scale turbulence (Servidio et al. 2015; Gro\u0161elj et al. 2017; Cerri et al. 2018, 2019; Hellinger et al. 2019; Wang et al. 2019). The main goal of our project is to investigate the possibility of improving the electron description in hybrid-kinetic models. The starting point is the \u201chybrid Vlasov\u2013Maxwell\u201d (HVM) code (Mangeney et al. 2002; Valentini et al. 2007), which is already equipped with a fluid model in which electrons are described as an isotropic, isothermal fluid with finite inertia. The HVM code has recently been upgraded in order to implement a more sophisticated model for the electron fluid. This includes evolution equations for the anisotropic (gyrotropic) electron pressures, p\u2225,\u2006e and p\u22a5,\u2006e (where \u2225 and \u22a5 refer to the local magnetic-field direction, b\u2004=\u2004B\/|B|), and a Landau-fluid (LF) closure for the parallel transport of the gyrotropic electron thermal energy along field lines (i.e., parallel heat fluxes, q\u2225,\u2006e and q\u22a5,\u2006e). Hereafter, we refer to this new model as \u201chybrid Vlasov\u2013Landau-fluid\u201d (HVLF). The idea is to include within a hybrid description the relevant electron pressure-anisotropy effects and a fluid model for the electron-kinetic response that still holds in a nonlinear regime. Therefore, the LF model implemented in the HVLF code goes beyond the early attempts to include these effects in simplified settings (e.g., Hammett & Perkins 1990; Snyder et al. 1997; Passot & Sulem 2007) and is based on the approach presented by Sulem & Passot (2015).","Citation Text":["Valentini et al. 2007"],"Functions Text":["The starting point is the \u201chybrid Vlasov\u2013Maxwell\u201d (HVM) code","which is already equipped with a fluid model in which electrons are described as an isotropic, isothermal fluid with finite inertia. The HVM code has recently been upgraded in order to implement a more sophisticated model for the electron fluid. This includes evolution equations for the anisotropic (gyrotropic) electron pressures, p\u2225,\u2006e and p\u22a5,\u2006e (where \u2225 and \u22a5 refer to the local magnetic-field direction, b\u2004=\u2004B\/|B|), and a Landau-fluid (LF) closure for the parallel transport of the gyrotropic electron thermal energy along field lines (i.e., parallel heat fluxes, q\u2225,\u2006e and q\u22a5,\u2006e)."],"Functions Label":["Uses","Background"],"Citation Start End":[[936,957]],"Functions Start End":[[852,912],[960,1546]]}
{"Identifier":"2021AandA...655A..72S___2019_Instance_2","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"],"Functions Text":["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","their Fig. 5)."],"Functions Label":["Background","Background"],"Citation Start End":[[1287,1304]],"Functions Start End":[[1089,1285],[1306,1320]]}
{"Identifier":"2022ApJ...937...62L__Xiao_et_al._2022_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":["Xiao et al. 2022"],"Functions Text":["while short GRBs are consistent with null or negligible intrinsic spectral lag and are therefore an ideal tool to measure the LIV effect"],"Functions Label":["Motivation"],"Citation Start End":[[803,819]],"Functions Start End":[[612,748]]}
{"Identifier":"2022MNRAS.510.4943S__Murray_&_Dermott_1999_Instance_1","Paragraph":"The gravitational potential of an eccentric companion at the quadrupole order can be decomposed as a sum over circular orbits (e.g. Storch & Lai 2013; Vick, Lai & Fuller 2017):\n(5)$$\\begin{eqnarray*}\r\nU\\left(\\boldsymbol{\\mathbf {r}}, t\\right) = \\sum \\limits _{m=-2}^2 U_{2m} \\left(\\boldsymbol{\\mathbf {r}}, t\\right) ,\r\n\\end{eqnarray*}$$(6)$$\\begin{eqnarray*}\r\nU_{2m}\\left(\\boldsymbol{\\mathbf {r}}, t\\right) &amp;=&amp; -\\frac{GM_2 W_{2m} r^2}{D(t)^3} Y_{2m}(\\theta , \\phi) e^{-imf\\!\\!\\!\\:(t)}, \\\\\r\n&amp;=&amp; -\\frac{GM_2W_{2m} r^2}{a^3}Y_{2m}\\left(\\theta , \\phi \\right) \\sum \\limits _{N = -\\infty }^\\infty \\!\\!F_{Nm}e^{-iN\\Omega t} .\r\n\\end{eqnarray*}$$Here, the coordinate system is centered on the MS star, (r, \u03b8, \u03d5) are the radial, polar, and azimuthal coordinates of $\\boldsymbol{\\mathbf {r}}$ respectively, $W_{2 \\pm 2} = \\sqrt{3\\pi \/10}$, W2 \u00b1 1 = 0, $W_{20} = -\\sqrt{\\pi \/ 5}$, D(t) is the instantaneous distance to the companion, f is the true anomaly, and Ylm denote the spherical harmonics. FNm denote the Hansen coefficients for l = 2 (also denoted $X^N_{2m}$ in Murray & Dermott 1999), which are the Fourier coefficients of the perturbing function, i.e.\n(7)$$\\begin{eqnarray*}\r\n\\frac{a^3}{D(t)^3} e^{-imf\\!\\!\\!\\:(t)} = \\sum \\limits _{N = -\\infty }^\\infty \\!\\!F_{Nm} e^{-iN\\Omega t}.\r\n\\end{eqnarray*}$$The FNm can be written explicitly as an integral over the eccentric anomaly (Murray & Dermott 1999; Storch & Lai 2013):\n(8)$$\\begin{eqnarray*}\r\nF_{Nm} = \\frac{1}{\\pi }\\int \\limits _{0}^{\\pi } \\frac{\\cos \\left[N\\left(E - e\\sin E\\right) - mf(E)\\right]}{\\left(1 - e\\cos E\\right)^2}\\,\\,\\mathrm{d}E.\r\n\\end{eqnarray*}$$By considering the effect of each summand in equation (5), the total torque on the star, energy transfer in the inertial frame, and energy transfer in the star\u2019s corotating frame (which is also the tidal heating rate) can be obtained (Storch & Lai 2013; Vick et al. 2017):\n(9)$$\\begin{eqnarray*}\r\nT = \\sum \\limits _{N = -\\infty }^\\infty F_{N2}^2 T_{\\rm circ}\\left(N\\Omega - 2\\Omega _{\\rm s}\\right),\r\n\\end{eqnarray*}$$(10)$$\\begin{eqnarray*}\r\n\\dot{E}_{\\rm in} &amp;=&amp; \\frac{1}{2}\\sum \\limits _{N = -\\infty }^\\infty \\Bigg [ \\left(\\frac{W_{20}}{W_{22}}\\right)^2 N\\Omega F_{N0}^2 T_{\\rm circ}\\left(N\\Omega \\right) \\\\\r\n&amp;&amp;+\\, N\\Omega F_{N2}^2 T_{\\rm circ}\\left(N\\Omega - 2\\Omega _{\\rm s}\\right) \\Bigg ] ,\r\n\\end{eqnarray*}$$(11)$$\\begin{eqnarray*}\r\n\\dot{E}_{\\rm rot} = \\dot{E}_{\\rm in} - \\Omega _{\\rm s} T .\r\n\\end{eqnarray*}$$Here, dots indicate time derivatives.","Citation Text":["Murray & Dermott 1999"],"Functions Text":["also denoted $X^N_{2m}$ in"],"Functions Label":["Background"],"Citation Start End":[[1074,1095]],"Functions Start End":[[1047,1073]]}
{"Identifier":"2022MNRAS.514.2010M__Feng_&_Holder_2018_Instance_1","Paragraph":"In the last few years, several experiments have reported upper limits on the power spectrum of 21-cm fluctuations during reionization (Parsons et al. 2014; Patil et al. 2017; Barry et al. 2019; Mertens et al. 2020; The HERA Collaboration 2021b) and the earlier cosmic-dawn era (Eastwood et al. 2019; Gehlot et al. 2019, 2020; Garsden et al. 2021; Yoshiura et al. 2021). Scenarios in which the bulk IGM is still colder than the cosmic microwave background (CMB) during reionization give rise to the strongest fluctuations and so will be the first models to be tested as upper limits continue to improve (e.g. Parsons et al. 2014; Pober et al. 2015; Greig, Mesinger & Pober 2016). Similarly, stronger-than-expected 21-cm signals can arise if the cosmic radio background has contributions other than the CMB (Feng & Holder 2018), e.g. synchrotron emission from accreting black holes (Ewall-Wice et al. 2018), star-forming galaxies (Mirocha & Furlanetto 2019), or from decaying particles (Fraser et al. 2018; Pospelov et al. 2018). Indeed, constraints from MWA, HERA, and LoFAR disfavour models with negligible X-ray heating at z \u223c 8\u20139 or very strong radio backgrounds (Ghara et al. 2020, 2021; Mondal et al. 2020; Greig et al. 2021a, b; The HERA Collaboration 2021a). Of course, the recent report of an absorption signal in the sky-averaged spectrum at z \u223c 17 from EDGES (Bowman et al. 2018) requires an even colder IGM (Barkana 2018; Boddy et al. 2018; Fialkov, Barkana & Cohen 2018; Kovetz et al. 2018; Mu\u00f1oz & Loeb 2018) or a brighter background (Ewall-Wice et al. 2018; Feng & Holder 2018; Fialkov & Barkana 2019; Mirocha & Furlanetto 2019) than models in \u039bCDM cosmologies generally predict. However, the most stringent power spectrum upper limits from The HERA Collaboration (2021b) are derived at sufficiently low redshifts relative to EDGES (z \u2272 10 versus z \u2243 18) that they cannot yet directly address the EDGES controversy (Hills et al. 2018; Bradley et al. 2019; Singh & Subrahmanyan 2019; Sims & Pober 2020; Tauscher, Rapetti & Burns 2020; Singh et al. 2021).","Citation Text":["Feng & Holder 2018"],"Functions Text":["Similarly, stronger-than-expected 21-cm signals can arise if the cosmic radio background has contributions other than the CMB"],"Functions Label":["Motivation"],"Citation Start End":[[806,824]],"Functions Start End":[[679,804]]}
{"Identifier":"2021ApJ...908..220T__Milone_et_al._2020_Instance_1","Paragraph":"Recently, our knowledge of the MW formation and evolution has been revolutionized by the massive amount of data products from the Gaia mission and large spectroscopic surveys. Several dwarf galaxies are suggested to have been accreted by the MW since its formation (e.g., Helmi et al. 2018; Belokurov et al. 2018; Myeong et al. 2019). It is expected that GCs formed in these dwarf galaxies mix up with GCs formed in situ, i.e., in the main MW progenitor, to form the current GC system. This formation history can be revealed by precise kinematics and dynamical modeling. Massari et al. (2019) placed Pal 13 into a group of GCs related to a dwarf galaxy, named Sequoia, which was likely accreted 9 Gyr ago (Myeong et al. 2019). The existence of MPs in Pal 13 supports the statement that this phenomenon is not unique in Galactic GCs formed in situ (e.g., Li & de Grijs 2019; Milone et al. 2020). Besides its proposed accreted origin, Pal 13 is suggested to be experiencing tidal stripping (Yepez et al. 2019; Piatti & Fern\u00e1ndez-Trincado 2020). Hamren et al. (2013) found that Pal 13 has lost a considerable amount of mass, which is related to its low present-day cluster mass. Recently, the discovery of tidal tails in Pal 13 was reported by Shipp et al. (2020). Using the RR Lyrae stars, these authors estimated the initial luminosity to be \n\n\n\n\n\n L\u2299, which is significantly larger than the current luminosity estimated by B11, \n\n\n\n\n\n L\u2299. Due to its large distance, several detailed features of Pal 13 have been discovered just recently (e.g., Piatti & Fern\u00e1ndez-Trincado 2020; Shipp et al. 2020) and more are still waiting for further investigation, the estimation of total mass loss or initial mass of Pal 13 should be illuminative for future studies. Besides Pal 13, several remote, low mass GCs, including Whiting 1 and Eridanus GC, also show evidence of tidal tails or extra-tidal structures(Carballo-Bello et al. 2014; Myeong et al. 2017). These features agree with the predictions of dynamical evolution of low mass GCs: it is difficult to keep stars in a shallow potential well. The relatively large stellar mass lost from low mass GCs also complicates the discussion of MPs and especially of the lower mass limit needed to maintain MPs. On the other hand, the N-rich stars located in these GCs are lost to the field, contributing to the rare N-rich field stars (e.g., Martell & Grebel 2010; Fern\u00e1ndez-Trincado et al. 2017; Tang et al. 2019; Fern\u00e1ndez-Trincado et al. 2019; Tang et al. 2020).","Citation Text":["Milone et al. 2020"],"Functions Text":["The existence of MPs in Pal 13 supports the statement that this phenomenon is not unique in Galactic GCs formed in situ (e.g.,"],"Functions Label":["Similarities"],"Citation Start End":[[874,892]],"Functions Start End":[[727,853]]}
{"Identifier":"2022MNRAS.515.2256V__Drury_1983_Instance_1","Paragraph":"SBs are complex environments. The hot rarefied plasma is spanned by primary shocks, which decay into turbulent motions and MHD waves. Collective effects such as re-acceleration processes contribute to the acceleration of particles. SBs are also delimited by a forward shock which expands into the ISM. While the size of this shock can be very large (R \u2273 100 pc, Weaver et al. 1977), it is too slow (u \u223c 30 km\u2009s\u22121) to accelerate PeV protons, even assuming large magnetic fields in the ISM. More generally, in the case of strong shocks, the velocity u is identified as the velocity of the shock, as it is precisely the velocity jump at the shock discontinuity which drives the acceleration of the particles via the DSA mechanism (e.g. Drury 1983). Inside SBs, there are two types of strong primary shocks: the time-dependent SNR shocks which expand after a supernova (SN) explosion, and the wind termination shocks (WTS) which surround the individual massive stars, or the entire stellar cluster if it is compact enough. SNRs and WTSs have a typical velocity of several 1000 km\u2009s\u22121. In the low-density SB interior, SNRs expand to a typical radius of a few tens of pc before reaching the Sedov\u2013Taylor phase, which is generally larger than the radius of the WTS, even in the case of a WTS powered by a very massive compact cluster. The size of the latter shock depends on the mechanical power of the stellar cluster and is typically of the order of 10 pc (Weaver et al. 1977). It is believed that, due to in situ CR acceleration, efficient magnetic field amplification takes place upstream of SNR shocks, leading to a magnetic field of up to several tens of \u03bcG. On the other hand, it is less clear if the magnetic field can be as efficiently amplified in stellar winds. In both cases, the maximum energy is generally inferred to be only \u223c1 PeV (e.g. Gupta et al. 2020; Morlino et al. 2021; Vieu et al. 2022b). While this is slightly larger than the maximum energy achieved at isolated SNR, atypical conditions would be required for protons to be accelerated well beyond PeV by primary shocks embedded in SBs. A promising situation combining the advantages of both WTS and SNRs might be that of an SNR expanding in a wind profile close to a compact cluster. A powerful cluster may indeed convert a substantial amount of its mechanical energy into turbulence, amplifying magnetic fields up to hundreds of \u03bcG in its vicinity, such that particles could be accelerated up to 10 PeV by a powerful SNR propagating in the wind, even in the absence of additional magnetic-field amplification. Considering a similar scenario of SNR shock \u2013 cluster wind interaction with efficient turbulence generation, Bykov et al. (2015) found that proton energies up to 40 PeV could be reached in the case of fast shocks (U = 104 km\u2009s\u22121).","Citation Text":["Drury 1983"],"Functions Text":["More generally, in the case of strong shocks, the velocity u is identified as the velocity of the shock, as it is precisely the velocity jump at the shock discontinuity which drives the acceleration of the particles via the DSA mechanism (e.g."],"Functions Label":["Background"],"Citation Start End":[[733,743]],"Functions Start End":[[489,732]]}
{"Identifier":"2016AandA...588A..44Y__Lehtinen_&_Mattila_(1996)_Instance_1","Paragraph":"All coreshine observations were obtained with the IRAC instrument on board the Spitzer observatory and are gathered in Paladini (2014) and Lef\u00e8vre et al. (2014), from which we selected 21 starless cores in the Taurus-Perseus, Chamaeleon, Cepheus, and L183\/L134 regions (see their Fig. 9 and Table 1). Lef\u00e8vre et al. (2014) summarised their results in two figures that we reproduce here: the 4.5 to 3.6 \u03bcm ratio, that they name \u201ccoreshine ratio\u201d, as a function of the 3.6 \u03bcm intensity and the 2.2 to 3.6 \u03bcm ratio, that they name the \u201cnear-IR to mid-IR ratio\u201d, as a function of the coreshine ratio. For the model cloud, we use the parameters of the control cloud defined in Sect. 3. After convolving our models with a 10\u2033 FWHM Gaussian kernel to simulate the data analysis presented in Lef\u00e8vre et al. (2014) and following the Lehtinen & Mattila (1996) prescription to take into account the part of the ISRF+CM light that can be transmitted through the cloud3, we compute the synthetic photometry for each pixel along a radial cut through our model clouds. The results are shown in Figs. 10a and b, which present the coreshine ratio as a function of the 3.6 \u03bcm intensity for the 21 starless cores of Lef\u00e8vre et al. (2014) and in Fig. 10c and d, which displays the near- to mid-IR ratio as a function of the coreshine ratio for the Taurus-Perseus and L183\/L134 regions. The CMM model has to be ruled out to explain coreshine since it only marginally fits the coreshine ratio and fails to reproduce the near- to mid-IR ratio (green areas in Fig. 10). On the contrary, the CMM+AMM and CMM+AMMI models can explain the coreshine observations (magenta and blue areas in Fig. 10, respectively). Based on the results presented in Sect. 3 and Fig. 7, the cloud parameters are as important as the dust model for explaining the dispersion in the observations. Varying the cloud external radius from Rout = 0.3 to 0.1 pc (Figs. 10a and c) and the central density from \u03c1C = 104 to 5 \u00d7 105 H\/cm3 (Figs. 10b and d) seem enough to explain most of the observed scatter. ","Citation Text":["Lehtinen & Mattila (1996)"],"Functions Text":["After convolving our models with a 10\u2033 FWHM Gaussian kernel to simulate the data analysis presented in Lef\u00e8vre et al. (2014) and following the","prescription to take into account the part of the ISRF+CM light that can be transmitted through the cloud3, we compute the synthetic photometry for each pixel along a radial cut through our model clouds. The results are shown in Figs. 10a and b, which present the coreshine ratio as a function of the 3.6 \u03bcm intensity for the 21 starless cores of Lef\u00e8vre et al. (2014) and in Fig. 10c and d, which displays the near- to mid-IR ratio as a function of the coreshine ratio for the Taurus-Perseus and L183\/L134 regions."],"Functions Label":["Uses","Uses"],"Citation Start End":[[824,849]],"Functions Start End":[[681,823],[850,1365]]}
{"Identifier":"2022ApJ...933..243F__Hajela_et_al._2022_Instance_1","Paragraph":"Recently, Fraija et al. (2021a) presented the afterglow light curves generated by the deceleration of sub-relativistic masses ejected from the merger of BCOs and the death of massive stars. The authors assumed that a PL velocity distribution describes the isotropic-equivalent kinetic energy of these masses and that the sub-relativistic ejected masses were decelerated, in turn, by a stratified-density environment. As a particular case, to explain the multiwavelength observations of the gravitational event GW170817\/GRB 170817A at \u223c900 days, they constrained the parameter space of the synchrotron light curves of a sub-relativistic mass ejected during the merger of two NSs and decelerated in a constant-density environment. The synchrotron radiation of the sub-relativistic material was consistent with the faster blue KN afterglow. Inspired by the new observations of this GW event at 3.3 yr after the initial merger (Hajela et al. 2022), in this paper, we extend the synchrotron model presented in Fraija et al. (2021a), including the continuous energy injection from the central engine (either a spinning magnetized NS or BH remnant) into the blast wave through a numerical approach and analytic arguments. In addition, we apply the current model to potential candidates of sGRB events with evidence of a KN. The paper is organized as follows: Section 2 presents the dynamical evolution of the afterglow when the central engine continuously injects energy into the blast wave. We show an analytical solution and numerical approach. In Section 3, we show a synchrotron model with energy injection from a spinning magnetized NS and BH remnants. Section 4 shows the analysis of the multiwavelength light curves using typical values of the GRB afterglow. In Section 5, we apply our model to several potential candidates, including GW170817\/GRB 170817A, and finally, in Section 6, we summarize. We consider the convention \n\n\n\nQx=Q10x\n\n in cgs units and assume for the cosmological constants a spatially flat universe \u039b cold dark matter model with H\n0 = 69.6 km s\u22121 Mpc\u22121, \u03a9M = 0.286, and \u03a9\u039b = 0.714 (Planck Collaboration et al. 2016).","Citation Text":["Hajela et al. 2022"],"Functions Text":["Inspired by the new observations of this GW event at 3.3 yr after the initial merger","in this paper, we extend the synchrotron model presented in Fraija et al. (2021a), including the continuous energy injection from the central engine (either a spinning magnetized NS or BH remnant) into the blast wave through a numerical approach and analytic arguments."],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[924,942]],"Functions Start End":[[838,922],[945,1214]]}
{"Identifier":"2019AandA...622A..60C__Drake_et_al._2013a_Instance_1","Paragraph":"In order to establish the completeness and purity of the RR Lyrae stars confirmed by the SOS Cep&RRL pipeline and to estimate the number of new discoveries by Gaia, we performed a deep and careful comparison with the literature. As a first step, the catalogue of 140 784 confirmed sources was cross-matched against all major catalogues of known RR Lyrae stars that are available. We primarily used the OGLE catalogues for RR Lyrae stars (version IV of the survey, Soszy\u0144ski et al. 2014, 2016), but we also used RR Lyrae stars by CTRS (Drake et al. 2013a,b, 2014, 2017; Torrealba et al. 2015), ASAS (Pojmanski 1997; Richards et al. 2012), ASAS-SN (Jayasinghe et al. 2018), ATLAS (Tonry et al. 2018), IOMC (Alfonso-Garz\u00f3n et al. 2012), LINEAR (Palaversa et al. 2013), NSVS (Kinemuchi et al. 2006), Pann-Stars (PS1 Sesar et al. 2017), and from the works based on Kepler\/K2 (Debosscher et al. 2011; Nemec et al. 2011; Moln\u00e1r et al. 2015a,b, 2016) and on the Simbad database (Wenger et al. 2000). These cross-matches returned a list of 88 578 known RR Lyrae stars in our sample of 140 784 stars. The SOS Cep&RRL confirmed RR Lyrae stars were also cross-matched against catalogues of candidate RR Lyrae stars discovered by the VVV survey (Gran et al. 2016; Minniti et al. 2017; D. Minniti, priv. comm.) in the MW disc and bulge. This returned 319 VVV cross-identified sources in the MW disc and 222 in the MW bulge. We thus confirm these VVV candidates. For known RR Lyrae stars in GCs, the main reference was the catalogue of Clement et al. (2001), which was updated to the latest literature as described in Garofalo et al. (in prep.). For variables in dSphs, we used the following references: Kaluzny et al. (1995), Clementini et al. (2005), Kinemuchi et al. (2008), Dall\u2019Ora et al. (2012) and Garofalo et al. (2013). These latter cross-matches returned a list of 1986 further known RR Lyrae stars. At the end of this cross-match procedure, of the 140 784 RR Lyrae stars that are confirmed by the SOS Cep&RRL pipeline, 90 564 were shown to be known previously, and 50 220 are new discoveries by Gaia.","Citation Text":["Drake et al. 2013a"],"Functions Text":["We primarily used the OGLE catalogues for RR Lyrae stars","but we also used RR Lyrae stars by CTRS","These cross-matches returned a list of 88 578 known RR Lyrae stars in our sample of 140 784 stars."],"Functions Label":["Uses","Uses","Uses"],"Citation Start End":[[535,553]],"Functions Start End":[[380,436],[494,533],[992,1090]]}
{"Identifier":"2021ApJ...909..172Z__Read_&_Lebonnois_2018_Instance_2","Paragraph":"Atmospheric superrotation is characterized by eastward wind at the equator, which means the atmosphere there has a higher angular momentum than the solid surface. Atmospheric superrotation is a common phenomenon across the universe. In the solar system, superrotation exists in the atmospheres of Venus, Titan, Saturn, and Jupiter, as well as the stratospheric atmosphere of Earth during the westerly phase of the quasi-biennial oscillation (e.g., Kraucunas & Hartmann 2005; Schneider & Liu 2009; Lutsko 2018; Read & Lebonnois 2018). In order to maintain atmospheric superrotation, there must be momentum transports from higher latitudes to the equator against friction or other processes, according to angular momentum conservation (Hide 1969; Held 1999; Showman et al. 2013). This up-gradient transport into the jet can result from Rossby waves, coupled Rossby\u2013Kelvin waves, mixed Rossby\u2013gravity waves, wave\u2013jet resonance, barotropic instability, or baroclinic instability (Suarez & Duffy 1992; Del Genio & Zhou 1996; Joshi et al. 1997; Lee 1999; Williams 2003; Kraucunas & Hartmann 2005; Schneider & Liu 2009; Caballero & Huber 2010; Mitchell & Vallis 2010; Showman & Polvani 2010, 2011; Showman et al. 2010; Liu & Schneider 2011; Arnold et al. 2012; Pinto & Mitchell 2014; Tsai et al. 2014; Wang & Mitchell 2014; Laraia & Schneider 2015; Lutsko 2018; Read & Lebonnois 2018; Pierrehumbert & Hammond 2019). For example, Kraucunas & Hartmann (2005) suggested that in an Earth-like atmosphere, equatorial superrotation can be generated by equatorward stationary eddy momentum convergence, which is associated with zonal variations in the diabatic heating at low latitudes. Mitchell & Vallis (2010) studied the transition from current Earth-like atmospheric circulation to an equatorial superrotation state. They found that during the spin-up period, superrotation is generated by equatorward momentum convergence associated with both barotropic and baroclinic instabilities.","Citation Text":["Read & Lebonnois 2018"],"Functions Text":["This up-gradient transport into the jet can result from Rossby waves, coupled Rossby\u2013Kelvin waves, mixed Rossby\u2013gravity waves, wave\u2013jet resonance, barotropic instability, or baroclinic instability"],"Functions Label":["Background"],"Citation Start End":[[1355,1376]],"Functions Start End":[[778,974]]}
{"Identifier":"2021AandA...649A.168D__Oh_&_Escuti_2008_Instance_1","Paragraph":"A critical property of the phase mask is that it needs to be able to image subapertures in off-axis interferograms. The off-axis interferograms are rather large, with size scaling with \u03bb\/Dsub, where \u03bb is the wavelength and Dsub is the diameter of the subaperture. Therefore, imaging multiple interferograms onto separate locations (so as to avoid overlap) on the detector requires large phase tilts. This makes it difficult to manufacture classical phase implementations of a HAM phase mask for transmissive pupil planes. A solution is offered by liquid-crystal diffractive phase masks as they have an unbounded continuous phase (Escuti et al. 2016). This property enables the creation of steep phase ramps that efficiently diffract light into a single order without scattering (Oh & Escuti 2008). In Fig. 1, no noticeable second-order diffraction is seen for any off-axis interferogram. Another advantage of liquid-crystal masks is that it is possible to manufacture almost any phase pattern (Kim et al. 2015), meaning there is more design freedom. We exploited this by combining phase ramps into a single phase pattern that images a single subaperture onto multiple locations in the focal plane. This was done through multiplexing the phase ramps, and the mathematical description of multiplexing can be found in Doelman et al. (2018). An example of multiplexed subapertures is seen in the third column of Fig. 1, where multiple baselines connected to a single aperture are imaged onto different interferograms. In addition, liquid-crystal masks are diffractive because they apply a different kind of phase delay to incoming light that is independent of wavelength. These phase delays are called \u201cgeometric phase delays\u201d and are discussed in greater detail in Sect. 3. Due to this diffraction, the location of an imaged subaperture changes with wavelength. The advantage of the diffractive nature is that, together with the right subaperture combination and fringe orientation, the wavelength smearing enables low-resolution spectroscopy. However, each interferogram can then only consist of 1D combinations of subapertures. This limits the design freedom significantly and also greatly increases the number of off-axis interferograms. As shown in Fig. 1, the fringe direction of all off-axis interferograms is orthogonal to the smearing direction for 1D combinations of subapertures. Lastly, a specific property of liquid-crystal diffractive phase masks is that they produce two off-axis interferograms for a single phase ramp with opposite location in the focal plane. This can be seen in Fig. 1, where all interferograms have an identical counterpart. The aforementioned properties have a large impact on the design of the HAM mask, which we discuss next.","Citation Text":["Oh & Escuti 2008"],"Functions Text":["This property enables the creation of steep phase ramps that efficiently diffract light into a single order without scattering"],"Functions Label":["Uses"],"Citation Start End":[[779,795]],"Functions Start End":[[651,777]]}
{"Identifier":"2018MNRAS.473.2144B__Nakamura_&_Li_2007_Instance_1","Paragraph":"The problem of the balance between driving and decay for magnetized turbulence is most acute in molecular clouds. Since these have linewidths indicating the presence of supersonic flow, the fast dissipation of turbulence found by these simulations necessitates a mechanism to reinject the energy equally quickly. A number of candidates have been proposed, including internal feedback from H\u2009ii regions (Matzner 2002; Krumholz, Matzner & McKee 2006; Goldbaum et al. 2011) or protostellar outflows (Li & Nakamura 2006; Nakamura & Li 2007; Wang et al. 2010; Federrath et al. 2014a), driving of turbulence by ongoing accretion (Klessen & Hennebelle 2010; Goldbaum et al. 2011; Lee & Hennebelle 2016) or gravitational contraction on small scales (Federrath et al. 2011b; Sur et al. 2012), thermal instability driving (Koyama & Inutsuka 2002; Hennebelle & Inutsuka 2006) and injection of energy from external supernova shocks (Mac Low & Klessen 2004; Padoan et al. 2016a,b; Pan et al. 2016). Alternately, it is possible that the linewidths do not reflect turbulent motion at all, and instead indicate global gravitational collapse (Ballesteros-Paredes et al. 2011; Zamora-Avil\u00e9s & V\u00e1zquez-Semadeni 2014). Each of these proposals, however, faces challenges \u2013 internal feedback must maintain large linewidths without destroying the clouds in which they occur, driving by accretion faces the problem of what happens when the accretion eventually ends, thermal instability seems unlikely to be a viable mechanism in molecule-dominated galaxies that lack a significant warm phase, and external driving requires efficient coupling between the low-density external medium and the dense clouds. The view that clouds are in global collapse is hard to reconcile with the observed very low rates of star formation found even in gas at densities \u2273 105\u2009cm\u22123 (Krumholz & Tan 2007; Federrath & Klessen 2012; Krumholz, Dekel & McKee 2012; Evans, Heiderman & Vutisalchavakul 2014; Padoan et al. 2014; Salim, Federrath & Kewley 2015; Usero et al. 2015; Heyer et al. 2016; Vutisalchavakul, Evans & Heyer 2016).","Citation Text":["Nakamura & Li 2007"],"Functions Text":["The problem of the balance between driving and decay for magnetized turbulence is most acute in molecular clouds. Since these have linewidths indicating the presence of supersonic flow, the fast dissipation of turbulence found by these simulations necessitates a mechanism to reinject the energy equally quickly. A number of candidates have been proposed, including","or protostellar outflows"],"Functions Label":["Background","Background"],"Citation Start End":[[517,535]],"Functions Start End":[[0,365],[471,495]]}
{"Identifier":"2021MNRAS.507.5053E__Johnston_et_al._2006_Instance_1","Paragraph":"Multiwavelength observations of the GC indicate that the number of pulsars in the central few parsecs should be high (Wharton et al. 2012) and conditions are highly favourable for relativistic binaries (Faucher-Gigu\u00e8re & Loeb 2011). The dense nuclear star cluster surrounding Sgr A* (see e.g. Genzel, Eisenhauer & Gillessen 2010, for a review) contains a majority of older late-type stars, but contrary to expectations, massive young main-sequence stars (Ghez et al. 2003) and possible neutron star progenitors such as Wolf\u2013Rayet stars (Paumard et al. 2001). The presence of neutron stars is further indicated by large numbers of X-ray binaries, possible pulsar wind nebulae, X-ray features such as the \u2018cannonball\u2019 and compact radio variables (Muno et al. 2005; Wang, Lu & Gotthelf 2006; Zhao, Morris & Goss 2013, 2020). Despite this only six radio pulsars have been discovered within half a degree of Sgr A* (Johnston et al. 2006; Deneva, Cordes & Lazio 2009; Eatough et al. 2013c; Shannon & Johnston 2013) even after many dedicated searches at multiple wavelengths (Kramer et al. 1996a, 2000; Klein et al. 2004; Klein 2005; Deneva 2010; Macquart et al. 2010; Eatough et al. 2013a; Siemion et al. 2013). Hyperstrong scattering of radio waves in the GC has been the principal explanation for the scarcity of detected pulsars (Cordes & Lazio 1997, 2002; Lazio & Cordes 1998a,b), however, scatter broadening measurements of PSR J1745\u22122900 in Spitler et al. (2014) and Bower et al. (2014) appear to contest this.1 Other authors have noted that the lack of GC pulsars is expected under a certain set of conditions and considering the sensitivity limits of existing pulsar surveys (Chennamangalam & Lorimer 2014; Liu & Eatough 2017; Rajwade, Lorimer & Anderson 2017). Alternatively, the scarcity of detected pulsars might be caused by a more complex scattering structure towards the GC (Cordes & Lazio 1997; Lazio & Cordes 1998a, b; Johnston et al. 2006; Schnitzeler et al. 2016; Dexter et al. 2017).","Citation Text":["Johnston et al. 2006"],"Functions Text":["Despite this only six radio pulsars have been discovered within half a degree of Sgr A*"],"Functions Label":["Background"],"Citation Start End":[[911,931]],"Functions Start End":[[822,909]]}
{"Identifier":"2021MNRAS.503..324M__Zhao_et_al._2019_Instance_3","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"],"Functions Text":["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"],"Functions Label":["Similarities"],"Citation Start End":[[1148,1164]],"Functions Start End":[[988,1146]]}
{"Identifier":"2022ApJ...935..137K__Tegler_et_al._1995_Instance_1","Paragraph":"In Figure 2, the reduced AKARI IRC spectra of all protostars are presented; the absorption features of the H2O, CO2, and CO ices are clearly detected. All of our targets show deep and broad absorption features of H2O ice in the wavelength range 2.7\u20133.4 \u03bcm. In the case of AFGL 7009S, strong extinction toward the source saturates the absorption features throughout the wavelengths from 2.7 to 3.6 \u03bcm, including H2O ice. Other ice features, such as CH4 (Lacy et al. 1991; Boogert et al. 2004) and CH3OH (Grim et al. 1991; Brooke et al. 1996), were observed at 3.3\u22123.5 \u03bcm, but it is difficult to extract their absorption profiles from the blended features due to the low spectral resolution of AKARI IRC. An absorption feature of the CO2 ice around 4.27 \u03bcm was clearly detected toward all targets. At the wavelength around 4.6 \u03bcm for Perseus 1 and 3, RNO 91, and AFGL 7009S, there is a hint for another ice component overlapping with the CO absorption feature at 4.67 \u03bcm. Many near-infrared observations have revealed the same feature on the blue wing part of the CO absorption feature (Tegler et al. 1995; Chiar et al. 1998; Whittet et al. 2001; van Broekhuizen et al. 2005; Aikawa et al. 2012), which was suggested as an absorption feature of XCN ice. Lacy et al. (1984) and Pendleton et al. (1999) reported that the 4.62 \u03bcm absorption feature of XCN ice consists of a nitrile group and an unknown component \u201cX\u201d. Many laboratory studies have suggested that UV photolysis or cosmic ray irradiation of ice mantle could make the solid state OCN\u2212 on grain surfaces (Lacy et al. 1984; Grim & Greenberg 1987; Bernstein et al. 2000; Palumbo et al. 2000; Hudson et al. 2001; van Broekhuizen et al. 2004). In addition to these ice components, there are some absorption features around 4.8 and 4.9 \u03bcm. For Perseus 3 and the background star, the absorption features with a peak position around 4.78 \u03bcm are likely associated with 13CO ice (Boogert et al. 2002; Pontoppidan et al. 2003). We also detected another absorption feature at 4.83 \u03bcm toward the low-luminosity targets. However, we could not find any corresponding ice features from previous studies. The 4.9 \u03bcm absorption feature detected toward all targets was identified as solid carbonyl sulfide (OCS) ice. OCS ice can be produced when the interstellar ices containing CO and CO2 are exposed to UV photons or cosmic rays (Palumbo et al. 1997).","Citation Text":["Tegler et al. 1995"],"Functions Text":["Many near-infrared observations have revealed the same feature on the blue wing part of the CO absorption feature","which was suggested as an absorption feature of XCN ice."],"Functions Label":["Similarities","Background"],"Citation Start End":[[1085,1103]],"Functions Start End":[[970,1083],[1195,1251]]}
{"Identifier":"2020AandA...643A..35P__Irwin_et_al._2007_Instance_1","Paragraph":"In order to achieve high photometric accuracy and be sensitive to low amplitude undulations, we adopted techniques from the exoplanet community, with the purpose of eliminating the systematic errors. When performing differential photometry (Sect. 3), accurate bias-subtraction and flat-fielding are of major importance. According to Irwin et al. (2007), the Poisson noise is 200 e\u2212 for a typical detector with a gain of a few e\u2212 ADU\u22121 and the flat illumination level is of 20 000 ADU pixel\u22121 = 40 000 e\u2212 pixel\u22121. Thus, a typical photometric aperture with a radius of 3 pixels contributes \u223c1 mmag photon noise. For this reason, we obtained a considerable amount of biases (150\u2212300 frames) and twilight flat-fields (25\u2212100 frames) each night to reduce the Poisson noise to less than 0.2 mmag (Irwin et al. 2007). The bias frames were averaged together using the minmax in the reject option of the zerocombine task in IRAF with a view of keeping radiation events out of the master bias frame. The master flat frame was the result of combining all the frames using a median mode. The median value is an excellent way of removing the effects of hot pixels and cosmic rays, so these extreme values do not affect the calculation, as they would, if they would averaged. The reject option was set to avsigclip, in which case the \u201ctypical\u201d sigma would have been determined from the data itself rather than an a priori knowledge of the noise characteristics of the CCD. Other related issues that can limit the photometric precision are: (i) the positioning of the telescope, (ii) fringing issues, and (iii) the differential variations on the quantum efficiency of the pixels. With the aim of minimizing the contribution of these effects, we repositioned each star almost on the same pixel of the detector using the autoguiding system of each telescope. The read-out-noise of the detectors are insignificant, as it can be as low as a few e\u2212 ( 10\u2006e\u2212 for RISE2 and Andor Zyla cameras).","Citation Text":["Irwin et al. (2007)"],"Functions Text":["According to",", the Poisson noise is 200 e\u2212 for a typical detector with a gain of a few e\u2212 ADU\u22121 and the flat illumination level is of 20 000 ADU pixel\u22121 = 40 000 e\u2212 pixel\u22121. Thus, a typical photometric aperture with a radius of 3 pixels contributes \u223c1 mmag photon noise."],"Functions Label":["Uses","Uses"],"Citation Start End":[[333,352]],"Functions Start End":[[320,332],[352,609]]}
{"Identifier":"2021AandA...654A..88W__Bordoloi_et_al._2014_Instance_1","Paragraph":"The CGM (see Tumlinson et al. 2017, for a detailed review) is now understood to be a key component in disentangling the feedback processes in active galaxies. It links the smaller-scale interstellar medium (ISM) of the galaxy to the larger-scale intergalactic medium (IGM), not only in a geometrical way but also by acting as the reservoir fueling star formation and the central black hole, where the feedback interacts with the galactic environment and where the gas recycling during galaxy evolution is controlled. This complex environment is multiphase and has been observed in numerous surveys (e.g., Tumlinson et al. 2013; Bordoloi et al. 2014; Peek et al. 2015; Borthakur et al. 2015) at low redshift. A prominent feature of the CGM around active galaxies is the Ly\u03b1 (Lyman-\u03b1) emission line, which is also ubiquitously observed at high redshift (e.g., Haiman & Rees 2001; Reuland et al. 2003; van Breugel et al. 2006; Villar-Mart\u00edn 2007; Humphrey et al. 2013; Cantalupo et al. 2014; Wisotzki et al. 2016, 2018; Arrigoni Battaia et al. 2018, 2019; Nielsen et al. 2020). Ly\u03b1 is the transition of the hydrogen electron from the 2p orbit to its ground state. It can happen primarily through collisional excitation and recombination (see Dijkstra 2014, 2017, for a detailed review of Ly\u03b1 emission mechanisms and radiative transfer). In extragalactic studies, the recombination production of Ly\u03b1 emission can be generated by photoionization by young stars and\/or AGN (fluorescence). This fluorescence emission on larger scales (CGM and IGM) can also be due to UV background radiation. Additionally, collisional excitation can play an important role in the emission seen in outflows and infalling gas (Ouchi et al. 2020). The bright Ly\u03b1 emission line, along with other UV lines excited by the central or background sources, provides a useful tool for studying the galactic environments in the early Universe. Additionally, H\u202fI and metal absorption features observed in the CGM are powerful tracers of feedback signatures as well as tracers of infalling pristine gas (e.g., low metallicity absorption in a z\u2004\u223c\u20042.7 submillimeter galaxy; Fu et al. 2021). The sensitive integral field spectrographs on the largest ground-based telescopes, such as MUSE (Multi-Unit Spectroscopic Explorer; Bacon et al. 2010, 2014) and KCWI (Keck Cosmic Web Imager; Morrissey et al. 2012; see Cai et al. 2019 for observation of Ly\u03b1 halos with KCWI), are perfectly suited for mapping these UV features as they move into the optical band for high-redshift sources.","Citation Text":["Bordoloi et al. 2014"],"Functions Text":["This complex environment is multiphase and has been observed in numerous surveys","at low redshift."],"Functions Label":["Background","Background"],"Citation Start End":[[628,648]],"Functions Start End":[[517,597],[691,707]]}
{"Identifier":"2015AandA...584A..75V__Essen_et_al._(2014)_Instance_2","Paragraph":"The data presented here comprise quasi-simultaneous observations during secondary eclipse of WASP-33 b around the V and Y bands. The predicted planet-star flux ratio in the V-band is 0.2 ppt, four times lower than the accuracy of our measurements. Therefore, we can neglect the planet imprint and use this band to measure the stellar pulsations, and most specifically to tune their current phases (see phase shifts in von Essen et al. 2014). Particularly, our model for the light contribution of the stellar pulsations consists of eight sinusoidal pulsation frequencies with corresponding amplitudes and phases. Hence, to reduce the number of 24 free parameters and the values they can take, we use prior knowledge about the pulsation spectrum of the star that was acquired during von Essen et al. (2014). As the frequency resolution is 1\/\u0394T (Kurtz 1983), 3.5 h of data are not sufficient to determine the pulsations frequencies. Therefore, during our fitting procedure we use the frequencies determined in von Essen et al. (2014) as starting values plus their derived errors as Gaussian priors. As pointed out in von Essen et al. (2014), we found clear evidences of pulsation phase variability with a maximum change of 2 \u00d7 10-3 c\/d. In other words, as an example after one year time a phase-constant model would appear to have the correct shape with respect to the pulsation pattern of the star, but shifted several minutes in time. To account for this, the eight phases were considered as fitting parameters. The von Essen et al. (2014) photometric follow-up started in August, 2010, and ended in October, 2012, coinciding with these LBT data. We then used the phases determined in von Essen et al. (2014) during our last observing season as starting values, and we restricted them to the limiting cases derived in Sect. 3.5 of von Essen et al. (2014), rather than allowing them to take arbitrary values. The pulsation amplitudes in \u03b4 Scuti stars are expected to be wavelength-dependent (see e.g. Daszy\u0144ska-Daszkiewicz 2008). Our follow-up campaign and these data were acquired in the blue wavelength range. Therefore the amplitudes estimated in von Essen et al. (2014), listed in Table 1, are used in this work as fixed values. This approach would be incorrect if the pulsation amplitudes would be significantly variable (see e.g., Breger et al. 2005). Nonetheless, the short time span of LBT data, and the achieved photometric precision compared to the intrinsically low values of WASP-33\u2019s amplitudes, make the detection of any amplitude variability impossible. ","Citation Text":["von Essen et al. (2014)"],"Functions Text":["Hence, to reduce the number of 24 free parameters and the values they can take, we use prior knowledge about the pulsation spectrum of the star that was acquired during"],"Functions Label":["Uses"],"Citation Start End":[[781,804]],"Functions Start End":[[612,780]]}
{"Identifier":"2019MNRAS.485.5453S__Sahal-Brechot_et_al._1996_Instance_1","Paragraph":"Here we give a simple estimate for the population of 2s-state hydrogen atoms by considering a three-level system (1s, 2s, and 3p). We set the rate equation for the 2s population as\n(1)\r\n\\begin{eqnarray*}\r\nn_{\\rm H,1s}C_{\\rm 1s,2s}+n_{\\rm H,3p}A_{\\rm 3p,2s}-n_{\\rm H,2s}A_{\\rm 2s,1s}=0,\r\n\\end{eqnarray*}\r\nwhere nH, j is the number density of hydrogen atoms in the state j; Cj, k and Aj, k are the collisional excitation rate and spontaneous decay rate for the transition from j to k, respectively. For the bound states, we use the notation j = njlj, where nj is the principal quantum number of the state j. Similarly, lj = 0, 1, 2, 3, ..., nj \u2212 1 (equivalently: s, p, d, f,...) is the orbital angular-momentum quantum number of the state j. Here we suppose that the depopulation term of 2s-state atoms is dominated by the spontaneous transition at the rate of A2s, 1s \u2243 8.2\u2009s\u22121. In reality, the collisional transition from 2s to 2p can be a subdominant process for depopulation. For a collision at a velocity \u223c108\u2009cm\u2009s\u22121, which is a typical velocity scale for young SNR shocks, the cross-section is \u223c10\u221213\u2009cm2, giving a reaction rate \u223c10\u22125\u2009cm3 s\u22121 (e.g. Janev, Langer & Evans 1987; Sahal-Brechot et al. 1996). Thus, if the density is \u223c106\u2009cm\u22123, the collisional depopulation becomes important. Note that we assume no strong radiation field inducing the radiative transition from 2s to any other state.1 The occupation number of 3p, nH, 3p, depends on the absorption of Ly\u2009\u03b2. Here we assume an isotropic radiation field for Ly\u2009\u03b2. Then, we obtain the rate equation for 3p as\n(2)\r\n\\begin{eqnarray*}\r\nn_{\\rm H,1s} \\left(C_{\\rm 1s,3p}{+}\\int _0^{\\infty } \\frac{4\\pi \\sigma _\\nu ^{\\rm 1s,3p}}{h\\nu }I_\\nu {\\rm d}\\nu \\right) {-}n_{\\rm H,3p} \\left(A_{\\rm 3p,1s}+A_{\\rm 3p,2s}\\right){=}0, \\nonumber \\\\\r\n\\end{eqnarray*}\r\nwhere h, \u03bd, $\\sigma _\\nu ^{\\rm 1s,3p}$, and I\u03bd are the Planck constant, frequency, absorption cross-section for the transition from 1s to 3p and the specific intensity, respectively. The intensity is set to be\n(3)\r\n\\begin{eqnarray*}\r\nI_\\nu =S_\\nu (1-{\\rm e}^{-\\tau _\\nu }) =\\frac{ \\frac{h\\nu }{4\\pi } A_{\\rm 3p,1s} n_{\\rm H,3p} }{\\sigma ^{\\prime } n_{\\rm H,1s}} (1-{\\rm e}^{-\\tau _\\nu }),\r\n\\end{eqnarray*}\r\nwhere S\u03bd and \u03c4\u03bd are the source function and optical depth, respectively. \u03c3\u2032 is a combination of physical constants relevant to the radiative absorption cross-section. Thus, we derive the occupation number of 2s as\n(4)\r\n\\begin{eqnarray*}\r\nn_{\\rm H,2s}=\\frac{C_{\\rm 1s,2s}}{A_{\\rm 2s,1s}} \\left[ 1+\\frac{A_{\\rm 3p,2s}}{ {\\rm e}^{-\\tau _0}A_{\\rm 3p,1s}+A_{\\rm 3p,2s} } \\frac{ C_{\\rm 1s,3p} }{ C_{\\rm 1s,2s} } \\right] n_{\\rm H,1s},\r\n\\end{eqnarray*}\r\nwhere \u03c40 is the optical depth at the line centre. Here we assume a narrow line profile function \u03d5\u03bd for which we can approximate as $\\int _0^{\\infty }(1-{\\rm e}^{-\\tau _\\nu }){\\rm d}\\nu \\approx 1-{\\rm e}^{-\\tau _0}$. The terms in the brackets [...] indicate the contribution of the combination of the absorption and cascades. Note that roughly say, the ratios are A3p, 2s\/(A3p, 1s + A3p, 2s) \u2243 0.118, C1s, 3p\/C1s, 2s \u223c 2\u201310, and C1s, 2s\/A2s, 1s \u223c 10\u22129np, where np is the proton number density. Thus, if Ly\u2009\u03b2 is in the optically thick limit, nH, 2s is enhanced roughly at most 10 times compared with the optically thin case. The absorption coefficient of H\u2009\u03b1 at the line centre becomes\n(5)\r\n\\begin{eqnarray*}\r\nk_0({\\rm H\\alpha }) &amp;=&amp; \\sigma _0({\\rm H\\alpha }) n_{\\rm H,2s} \\nonumber \\\\\r\n&amp;&amp;{\\sim} 10^{-23}\\!-\\!10^{-22}{\\rm cm^{-1}} \\!\\left(\\frac{T_0}{6000{\\rm K}}\\right)^{-\\frac{1}{2}}\\! \\left(\\frac{n_{\\rm H,1s}}{ {\\rm 1cm^{-3} } }\\right)\\! \\left(\\frac{n_{\\rm p}}{ {\\rm 1cm^{-3} } }\\right)\\!, \\nonumber \\\\\r\n\\end{eqnarray*}\r\nwhere \u03c30(H\u03b1) is the radiative cross-section of H\u2009\u03b1 at the line centre for given temperature T0. Thus, if the SNR shock interacts with somewhat dense clump with a density of \u223c30\u2009cm\u22123 and a size of \u223c1\u2009pc, the H\u2009\u03b1 emission can be scattered. Note that the H\u2009\u03b2 emission can also be scattered but its absorption coefficient is about quarter of the H\u2009\u03b1 coefficient. The interaction between the shock and a dense clump is implied by the ripple of an SNR shock with a length-scale of \u223c10 per\u2009cent of SNR radius (e.g. Ishihara et al. 2010; Williams et al. 2013, 2016; Miceli et al. 2014; Sano et al. 2017; Tsubone et al. 2017, and see the discussion of Shimoda et al. 2015). Note that according to magnetohydrodynamic simulations performed by Inoue, Yamazaki & Inutsuka (2009) and Inoue et al. (2012), even if the shock propagates into a simulated ISM having density contrast ranging in \u223c1\u201330\u2009cm\u22123 as a consequence of thermal instability, the scale length of rippling is \u223c10\u2009per\u2009cent of the length of sides of simulation box.","Citation Text":["Sahal-Brechot et al. 1996"],"Functions Text":["For a collision at a velocity \u223c108\u2009cm\u2009s\u22121, which is a typical velocity scale for young SNR shocks, the cross-section is \u223c10\u221213\u2009cm2, giving a reaction rate \u223c10\u22125\u2009cm3 s\u22121 (e.g.","Thus, if the density is \u223c106\u2009cm\u22123, the collisional depopulation becomes important."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1181,1206]],"Functions Start End":[[978,1152],[1209,1291]]}
{"Identifier":"2019AandA...622A..91G__Genzel_&_Stutzki_1989_Instance_1","Paragraph":"Among the species studied in this work, the detection of broad (\u0394v > 30 km s\u22121) line-wing H2O (312-221) emission implies the presence of shocked gas activity (e.g., van Dishoeck et al. 2011, 2013). Indeed, we only detect H2 O and CH3OH rotationally excited lines toward Orion S and BN\/KL (Figs. 2 and B.1). Both species are abundant in the ice mantles that coat grains in cold dark clouds (e.g., Gibb et al. 2004). After the onset of protostellar outflows, high-velocity shocks sputter these grain mantles and heat the gas to high temperatures. Both effects enhance the abundance of gas-phase H2 O and CH3OH (e.g., Draine 1995; Jim\u00e9nez-Serra et al. 2008). In OMC-1, the low- and high-velocity outflows from BN\/KL plunge into the ambient molecular cloud (Genzel & Stutzki 1989) producing hot (from Tk \u2243 200\u20132000 K) and dense (up to n(H2) \u2243 106\u2013107 cm\u22123) post-shocked gas in H2 Peaks 1 and 2 (e.g., Gonz\u00e1lez-Alfonso et al. 2002; Goicoechea et al. 2015b). In our HIFI maps, these extreme conditions can be inferred from the moderately extended emission of the HCN J = 13\u201312 line (Fig. B.2), a rotational transition with a critical density close to 1010 cm\u22123 (ncr = Aul\u2215\u03b3ul(Tk), where \u03b3ul is the collisional de-excitation rate coefficient in cm3 s\u22121), around BN\/KL outflows. Interestingly, the observed HCN to HCO+ J = 6\u20135 line intensity ratio is \u22652 toward BN\/KL, and 1 almost elsewhere (see Fig. B.4, right). This may reflect a change in the chemistry between the extended PDR cloud component and the shocked gas in BN\/KL outflows. It may also reflect the much stronger mid-IR (MIR) radiation from the BN\/KL region that favors the radiative pumping of HCN through its vibrational levels and enhances the high-J rotational emission (e.g, Carroll & Goldsmith 1981; Ziurys & Turner 1986). Finally, the maps show that both HCO+ and HCN J = 6\u20135 lines display widespread emission outside the main star-forming sites (see Fig. B.1). This suggests that the gas density of the extended cloud layers traced by HCO+ and HCN J = 6\u20135 lines is moderately high.","Citation Text":["Genzel & Stutzki 1989"],"Functions Text":["In OMC-1, the low- and high-velocity outflows from BN\/KL plunge into the ambient molecular cloud","producing hot (from Tk \u2243 200\u20132000 K) and dense (up to n(H2) \u2243 106\u2013107 cm\u22123) post-shocked gas in H2 Peaks 1 and 2"],"Functions Label":["Uses","Uses"],"Citation Start End":[[754,775]],"Functions Start End":[[656,752],[777,889]]}
{"Identifier":"2017MNRAS.470..755H__Cox_et_al._2008_Instance_1","Paragraph":"Supermassive black holes (SMBHs) are believed to exist in the centres of all massive galaxies (Kormendy & Richstone 1995). A small proportion of these are growing, with gas accretion rates ranging from \u223c10\u22124 to 10 M\u2299 yr\u22121 and a proportionately wide range of bolometric luminosities (\u223c1042\u20131047 erg s\u22121). These are active galactic nuclei (AGNs) and may accrete large fractions of their mass in bursts of rapid accretion (Croton et al. 2006), requiring rapid inflow of gas from galaxy length-scales. Stripping the gas of enough angular momentum to allow for such rapid accretion, thereby powering the most luminous AGN, proves extremely challenging. Theoretical work suggests major mergers can provide the torque to displace such an overwhelming fraction of the angular momentum of the gas, allowing for the highest accretion rates on to the central black hole whilst transforming the galaxy morphology (Toomre & Toomre 1972; Barnes 1988; BarnesBarnes & Hernquist 1991; Di Matteo, Springel & Hernquist 2005; Cox et al. 2008). Gas rich mergers may trigger nuclear and global starbursts (Mihos & Hernquist 1994, 1996; Hopkins et al. 2006) and major mergers disrupt the morphologies of the colliding galaxies, often exhibiting long tidal tails or shells of expelled gas and stars soon after the merger has begun. Detecting this can be challenging however, since the single new galaxy has a relaxation time-scale after which morphological features of mergers fade (Tinsley 1978; Kennicutt et al. 1987; Ellison et al. 2013). Observational evidence suggesting a link between major mergers and SMBH accretion has been mixed (e.g. Gabor et al. 2009; Cisternas et al. 2011; Schawinski et al. 2011; Kocevski et al. 2012; Treister et al. 2012; Ellison et al. 2013; Villforth et al. 2014; Kocevski et al. 2015; Villforth et al. 2017). Alternatively, AGNs may be triggered secularly through, for example, disc instabilities (Bournaud et al. 2011), bars (Knapen, Shlosman & Peletier 2000; Oh, Oh & Yi 2012) or otherwise by minor mergers (Kaviraj 2013). It remains unclear whether alternatives to major merger triggering can drive several M\u2299 yr\u22121 of gas to the central SMBHs, as is necessary to power the most luminous AGN.","Citation Text":["Cox et al. 2008"],"Functions Text":["Theoretical work suggests major mergers can provide the torque to displace such an overwhelming fraction of the angular momentum of the gas, allowing for the highest accretion rates on to the central black hole whilst transforming the galaxy morphology"],"Functions Label":["Background"],"Citation Start End":[[1006,1021]],"Functions Start End":[[648,900]]}
{"Identifier":"2018MNRAS.473.1633K__Fremling_et_al._2014_Instance_1","Paragraph":"Kochanek (2009) examined the statistical properties expected for surviving binary companions to SNe assuming passively evolving systems (i.e. no binary interactions). As already noted, the companions are generally significantly fainter than the exploding star, although this is frequently not the case for stripped SN progenitors \u2013 for Type Ibc SNe, it should not be surprising to find that the binary companion is more visually luminous than the SN progenitor. This point is of considerable importance for the one candidate Type Ib progenitor iPTF13bvn (Cao et al. 2013, Groh, Georgy & Ekstr\u00f6m 2013, Bersten et al. 2014, Fremling et al. 2014, Eldridge et al. 2015, Eldridge & Maund 2016, Folatelli et al. 2016). If the initial binary fraction is F, then the fraction of passively evolving binaries that are in stellar binaries at death is\n(1)\r\n\\begin{eqnarray}\r\nf_{\\rm b} = { F \\over 1 + F f_q } \\quad {\\rm where}\\quad f_q = \\int _{q_{{\\rm min}}}^{q_{{\\rm max}}} q^{x-1} P(q) {\\rm d}q,\r\n\\end{eqnarray}\r\nx \u2243 2.35 is the slope of the initial mass function (IMF), qmin \u2264 q = M2\/M1 \u2264 qmax \u2264 1 is the mass ratio and P(q) with \u222bdqP(q) \u2261 1 is the distribution of mass ratios. For a Salpeter IMF and a flat P(q) distribution extending over 0 \u2264 q \u2264 1, fq = 0.426 and the fraction of SNe in stellar binaries at death is 23\u2009per\u2009cent, 41\u2009per\u2009cent, 57\u2009per\u2009cent and 70\u2009per\u2009cent for initial binary fractions of F = 25\u2009per\u2009cent, 50\u2009per\u2009cent, 75\u2009per\u2009cent and 100\u2009per\u2009cent, respectively. Such a flat f(q) distribution is commonly found for massive star binaries, although there is evidence that the widest binaries show a different distribution (e.g. Sana et al. 2012, Kobulnicky et al. 2014, Moe & Di Stefano 2016). Essentially, only the explosions of primaries occur in stellar binaries, so the fraction of SNe in stellar binaries is less than the initial fraction of binaries because some of the SNe are the explosions of secondaries. Binary evolution, particularly stellar mergers then adds further complications, but changes the rough statistics little (e.g. Sana et al. 2012).","Citation Text":["Fremling et al. 2014"],"Functions Text":["As already noted, the companions are generally significantly fainter than the exploding star, although this is frequently not the case for stripped SN progenitors \u2013 for Type Ibc SNe, it should not be surprising to find that the binary companion is more visually luminous than the SN progenitor. This point is of considerable importance for the one candidate Type Ib progenitor iPTF13bvn"],"Functions Label":["Background"],"Citation Start End":[[622,642]],"Functions Start End":[[167,553]]}
{"Identifier":"2022MNRAS.509..693R__Foreman-Mackey_et_al._2013_Instance_1","Paragraph":"We detect 1.33 mm continuum extended emission originating from HD 36546 (see Fig. 3). Peak and integrated flux, size, and inclination are reported in Table 2. Once deconvolved from the beam, the major axis of the disc spans 180 au, compatible with the size reported from scattered light observations (semimajor axis of 85 au, Currie et al. 2017). The inclination is directly provided by the casa tool, as obtained from the ratio of the major and minor axis deconvolved from the beam. The total mass can be estimated by assuming an optically thin dust disc as $M_d = \\frac{F_{\\nu } d^2}{B_{\\nu }(T_{d,c}) \\kappa _\\nu }$, where F\u03bd is the measured flux at 1.33 mm, d is the distance to the source, B\u03bd is the Planck function at the corresponding dust temperature Td, c, and \u03ba\u03bd is the mass absorption coefficient that we take as \u03ba\u03bd = 2 cm2 g\u22121 following Nilsson et al. (2010). In order to estimate a temperature for the dust, we have fitted a modified blackbody to the available photometry at wavelengths longer than 10 \u00b5m (AKARI, WISE, IRAS, Herschel)3 and the new ALMA photometry using the emcee Affine Invariant Markov chain Monte Carlo Ensemble sampler implementation (Foreman-Mackey et al. 2013). This wavelength range was chosen to avoid the silicate emission observed in the mid-IR spectra reported by Lisse et al. (2017) around \u223c10 \u00b5m. Additionally, there is some discrepancy between the photometric data of WISE, Herschel, and ALMA, and those of IRAS and AKARI. Given their larger point spread functions and lower sensitivites, we decided to exclude the latter two from the fitting. The modified blackbody model assumes all dust grains have the same composition and size, and accounts for changes in the dust emission efficiency (Q\u03bb) via two additional free parameters: \u03b2 and a reference wavelength \u03bb0, such that the blackbody emission is modified by a factor Q\u03bb = 1 \u2212 exp\u2009[(\u03bb0\/\u03bb)\u03b2] (e.g. Williams et al. 2004) Therefore, our model has four free parameters: a scaling factor that controls the disc luminosity, the dust temperature, \u03b2, and \u03bb0. Uniform priors were used for all parameters within reasonable ranges: scaling values that produce disc luminosities consistent with the observed photometry, dust temperatures between 20 and 250\u2009K, \u03b2 values between 0 and 1.5, and \u03bb0 between 0.3 and 300\u2009\u00b5m (the latter was explored in log scale). This analysis yields a dust temperature value of 153 \u00b1 3 K, and a \u03b2 value of 0.24 $^{+0.07}_{-0.05}$ as shown in Fig. 1. \u03bb0 is unconstrained. Adopting this temperature value results in a dust mass of (9.0 \u00b1 1.0) \u00d7 10\u22122 M\u2295. Additionally, the resulting models yield Ldisc\/L* = (4.43 \u00b1 0.15) \u00d7 10\u22123, compatible with previous results (see Table 1).","Citation Text":["Foreman-Mackey et al. 2013"],"Functions Text":["In order to estimate a temperature for the dust, we have fitted a modified blackbody to the available photometry at wavelengths longer than 10 \u00b5m (AKARI, WISE, IRAS, Herschel)3 and the new ALMA photometry using the emcee Affine Invariant Markov chain Monte Carlo Ensemble sampler implementation"],"Functions Label":["Uses"],"Citation Start End":[[1168,1194]],"Functions Start End":[[872,1166]]}
{"Identifier":"2016AandA...591L...7B__Kruit_(1994)_Instance_1","Paragraph":"We created 10\u2009000 3D models of galaxies, each with an exponential disc plus a S\u00e9rsic bulge. We adopted the following functional form for the exponential disc (L\u00f3pez-Corredoira et al. 2002): (1)\\begin{equation} \\label{eq:Corredoira}  \\rho(R,z)=\\rho_{\\mathrm{0}}\\cdot\\exp{\\Bigg(\\frac{-R}{h_{R}}\\Bigg)}\\cdot \\exp{\\Bigg(\\frac{-|z|}{h_{z}(R)}\\Bigg)}\\cdot \\frac{h_{z}(R)}{h_{z}(0)} \\cdot \\end{equation}\u03c1(R,z)=\u03c10\u00b7exp\u2212RhR\u00b7exp\u2212|z|hz(R)\u00b7hz(R)hz(0)\u00b7Following the observations, we explored the effect of a linear increase in the vertical scale height, as follows: (2)\\begin{equation} \\label{eq:Flares_linear} h_{z}(R) = \\begin{cases} h_{z}(0) &amp; \\text{if } R \\leq R_{\\mathrm{flare}} \\\\ h_{z}(0) + \\dhzdR \\cdot R &amp; \\text{if } R > R_{\\mathrm{flare}}. \\end{cases} \\end{equation}hz(R)=ifR\u2264RflareifR>Rflare.Graham (2001) analysed a sample of 86 face-on disc-dominated galaxies previously selected by de Jong & van der Kruit (1994). This author performed a bulge + disc decomposition for 69 galaxies in the I-band, correcting for the effects of the internal extinction, Galactic extinction, inclination, and cosmological dimming (Graham 2003), that we used as a reference for our models. We estimated stellar masses using the relationship between the V-band mass-to-light ratio of galaxies and their dust-corrected rest-frame colours derived by Wilkins et al. (2013). According to these authors, for z 0.1 the optimal observed colour is (B \u2212 V), so we have estimated this colour for the 69 galaxies of Graham (2003) from HyperLeda data1, and estimated the stellar mass of each galaxy using the relations in Wilkins et al. For those objects without (B \u2212 V) available in HyperLeda, we estimated them from their SDSS (g \u2212 r) colour following the transformations published in Jester et al. (2005). To simulate realistic images of the disc galaxies, we adopted the observational I-band distributions of Graham (2001, 2003) for the photometric parameters (re, S\u00e9rsic index n, hR, \u03bc0, \u03bce, B\/T, the absolute magnitudes of the disc Mabs,disc and the bulge Mabs,bulge), and four morphological type bins (S0\u2013Sa, Sb\u2013Sbc, Sc\u2013Scd and Sd\u2013Sdm), in three mass bins (10  log\u200910M\/M\u2299 10.7,10.7  log\u200910M\/M\u2299 11 and log\u200910M\/M\u2299> 11) in order to explore realistic mass distributions for the morphological type bins. In Fig. 1 we represent the distributions of the structural and photometric parameters from which we created the models and compared them to the observations they are based on. For each morphological type bin we randomly chose the ratio of scale height to scale length (hz\/hR) from the observational range of values corresponding to each type reported by Kregel et al. (2002) and Mosenkov et al. (2015) in the I-band, as shown in Fig.\u20091. ","Citation Text":["de Jong & van der Kruit (1994)"],"Functions Text":["Graham (2001) analysed a sample of 86 face-on disc-dominated galaxies previously selected by"],"Functions Label":["Uses"],"Citation Start End":[[890,920]],"Functions Start End":[[797,889]]}
{"Identifier":"2017MNRAS.464..183N__Biviano_&_Katgert_2004_Instance_1","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"],"Functions Text":["hus, 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."],"Functions Label":["Uses"],"Citation Start End":[[830,852]],"Functions Start End":[[610,829]]}
{"Identifier":"2021ApJ...914L..19Z__McKernan_et_al._2012_Instance_1","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"],"Functions Text":["The disk of an AGN provides a natural environment for stars and compact objects to accrete materials and to migrate within it (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[1050,1070]],"Functions Start End":[[917,1049]]}
{"Identifier":"2015MNRAS.453.3414A__the_1999_Instance_1","Paragraph":"Filippenko & Chornock (2001) first presented the dynamical estimate of mass of the source to be around 7.4 \u00b1 1.1\u2009M\u2299. Recently, Radhika & Nandi (2014) claimed that the mass of XTE J1859+226 is perhaps in between 6.58 and 8.84\u2009M\u2299 which is similar to the prediction of Shaposhnikov & Titarchuk (2009), although the lower mass limit is estimated as 5.4\u2009M\u2299 by Corral-Santana et al. (2011). However, we consider the typical mass of the source as 7\u2009M\u2299. The distance of this source is around d \u223c 11\u2009kpc (Filippenko & Chornock 2001). Steiner et al. (2013) measured the spin as ak \u223c 0.4; however, Motta et al. (2014b) recently reported that the spin of the source is ak \u223c 0.34. Since the spin predictions are quite close, we use ak \u223c 0.4 for this analysis. We estimate the fluxes Fx (see Table 1) of LHS and HIMS of the 1999 outburst of the source (Radhika & Nandi 2014). The corresponding disc luminosities are calculated as $L_{\\rm disc}^{{\\rm LHS}}=8.26 \\times 10^{37}\\ {\\rm erg\\ s^{-1}}$ and $L_{\\rm disc}^{{\\rm HIMS}}=1.85 \\times 10^{38}\\ {\\rm erg\\ s^{-1}}$, respectively. Now, it is reasonable to assume the accretion efficiency for rotating BH as \u03b7 = 0.3 which corresponds to the accretion rate of the inflowing matter as ${\\dot{M}}_{{\\rm acc}}^{{\\rm LHS}} = 0.304 {\\dot{M}}_{{\\rm Edd}}$ in LHS and ${\\dot{M}}_{{\\rm acc}}^{{\\rm HIMS}} = 0.680 {\\dot{M}}_{{\\rm Edd}}$ in HIMS. For LHS, we use $R_{\\dot{m}}=9.83$\u2009per\u2009cent following our theoretical estimate where xs = 64.6rg for ak = 0.4, ${\\mathcal {E}}=0.001\\,98$ and \u03bb = 3.18. Incorporating these inputs in equation (15), we obtain the jet kinetic power as $L^{{\\rm LHS}}_{{\\rm jet}} = 2.52\\times 10^{37}\\ {\\rm erg\\ s^{-1}}$. The maximum mass outflow rate for HIMS corresponding to ak = 0.4 is obtained from Fig. 9 as $R^{\\rm max}_{\\dot{m}}=17.5$\u2009per\u2009cent for ${\\mathcal {E}}=0.005\\,47$ and \u03bb = 3.1, where the shock transition occurs at 21.9rg. Using these values in equation (15), we obtain the maximum jet kinetic power as $L^{{\\rm HIMS}}_{{\\rm jet}} = 1.08\\times 10^{38}\\ {\\rm erg\\ s^{-1}}$ which we regard to be associated with the HIMS of this source.","Citation Text":["Radhika & Nandi (2014)"],"Functions Text":["Recently,","claimed that the mass of XTE J1859+226 is perhaps in between 6.58 and 8.84\u2009M\u2299 which is similar to the prediction of Shaposhnikov & Titarchuk (2009)","although the lower mass limit is estimated as 5.4\u2009M\u2299 by Corral-Santana et al. (2011)."],"Functions Label":["Similarities","Similarities","Differences"],"Citation Start End":[[127,149]],"Functions Start End":[[117,126],[150,297],[299,384]]}
{"Identifier":"2017AandA...607A..71G__Hansen_&_Oh_(2006)_Instance_1","Paragraph":"An implication of the respective escape fractions of the two regimes is visible in Fig. 12. Here we show several values of NHI,cl for the static setup using \u03c4d,cl = 10-4 (empty symbols) and \u03c4d,cl = 1 (filled symbols), which correspond to metallicities of \\hbox{$Z\/Z_\\odot = 0.63\\left(\\tau_{\\rm d}\/10^{-4}\\right)\\left(10^{17}\\cm^{-2}\/N_{\\HI,\\cl}\\right)$}Z\/Z\u2299=0.63\u03c4d(\/10-4)(1017\u2009cm-2\/NHI,cl) (Pei 1992; Laursen et al. 2009); this reaches clearly unrealistic values. However, as in this paper we are interested in the fundamental impact of the individual parameters, we also study these extreme values. Also shown in Fig. 12 (with a black [gray] solid line for the low [high] dust content) is the proposed analytic solution for fesc by Hansen & Oh (2006)(18)\\begin{equation} f_{\\rm esc}^{\\rm HO06} = 1\/{\\rm cosh}(\\!\\sqrt{2 N_{\\cl}\\epsilon}), \\label{eq:fescHO06} \\end{equation}fescHO06=1\/cosh(2Ncl\u03f5),where for Ncl we used Eq. (8)(with (a1,b1) = (3\/2, 2) as found in Sect. 4.1) and for the clump albedo (i.e., the fraction of incoming photons that are reflected) \u03f5, we adopted a value of c1(1\u2212e\u2212 \u03c4d,cl) with c1 = 1.6 [c1 = 0.06] to match the NHI,cl = 1022 cm-2 data points for \u03c4d,cl = 10-4 [\u03c4d,cl = 1]. The behavior for the low- and high-dust contents is quite different. On the one hand, the escape fractions versus NHI,cl scales for \u03c4d,cl = 1 (filled symbols in Fig. 12) as predicted by Hansen & Oh (2006) in their \u201csurface scatter\u201d approximation, that is, a larger clump hydrogen column density \u201cshields\u201d the dust better from the Ly\u03b1 photons and thus prevents their destruction more efficiently. On the other hand, however, this is not the case for the low-dust scenario presented in Fig. 12 (with unfilled symbols) where a larger value of NHI,cl implies a lower fesc. This is because here the dust optical depth through all the clumps (shown in the black dotted line in Fig. 12) is lower than the accumulated dust optical depth through the subsequent random-walk clump encounters (black solid line), i.e., \\hbox{$\\exp(-4\/3 \\fc \\tau_{\\rm d, cl}) \\lesssim f_{\\rm esc}^{\\rm HO06}$}exp(\u22124\/3fc\u03c4d,cl)\u2272fescHO06. Consequently, configurations in the \u201cfree-streaming\u201d regime can possess enhanced Ly\u03b1 escape fractions compared to the \u201crandom walk\u201d regime (see Sect. 5.2 for a more detailed discussion). Still, both cases possess (much) larger escape fractions than a homogeneous slab, which is shown in Fig. 12 with a black dashed line. Here, we use the derived escape fraction by Neufeld (1990) with NHI = 4\/3 \u00d7 fc1022 cm-2 and \u03c4d = 4\/3fc\u03c4d,cl, i.e., with equal column densities as in the NHI,cl = 1022 cm-2 case. ","Citation Text":["Hansen & Oh (2006)"],"Functions Text":["Also shown in Fig. 12 (with a black [gray] solid line for the low [high] dust content) is the proposed analytic solution for fesc by","(18)\\begin{equation} f_{\\rm esc}^{\\rm HO06} = 1\/{\\rm cosh}(\\!\\sqrt{2 N_{\\cl}\\epsilon}), \\label{eq:fescHO06} \\end{equation}fescHO06=1\/cosh(2Ncl\u03f5),where for Ncl we used Eq. (8)(with (a1,b1) = (3\/2, 2) as found in Sect. 4.1) and for the clump albedo (i.e., the fraction of incoming photons that are reflected) \u03f5, we adopted a value of c1(1\u2212e\u2212 \u03c4d,cl) with c1 = 1.6 [c1 = 0.06] to match the NHI,cl = 1022 cm-2 data points for \u03c4d,cl = 10-4 [\u03c4d,cl = 1]."],"Functions Label":["Uses","Uses"],"Citation Start End":[[733,751]],"Functions Start End":[[600,732],[751,1197]]}
{"Identifier":"2020MNRAS.498.5299M__Bernardeau,_Waerbeke_&_Mellier_2003_Instance_1","Paragraph":"From the early days of detection the weak lensing (Munshi et al. 2008) studies have matured to a point where weak lensing results from Euclid are expected to constrain the cosmological parameters to sub-per\u2009cent accuracy. However, weak lensing at smaller angular scales probes the non-linear regime of gravitational clustering, and is thus key to understanding the non-Gaussianity induced by the non-linearity and fullly exploiting in the weak lensing maps. The higher order statistics are also useful for the breaking of parameter degeneracies in studies involving the power spectrum analysis alone and they are also important in understanding the variance or error of estimation of lower order statistics. These higher order statistics including the cumulants (Bernardeau 1994b) and their correlators (Bernardeau 1996a; Calabrese et al. 2010; Munshi et al. 2011; Riquelme & Spergel 2012) are among the best-known diagnostics of the deviation from Gaussianity that characterizes the non-linear regime (Bartolo et al. 2004), with a long history analytical modelling (Bernardeau et al. 2002a). Most of these studies use extensions of perturbative results in the quasi-linear regime valid at large smoothing angular scales or variants of halo models (Cooray & Sheth 2002). Early studies concentrated on measurements of higher order correlation hierarchy in the angular space due to small survey size (Bernardeau, Mellier & van Waerbeke 2002b; Bernardeau, Waerbeke & Mellier 2003). However, the near all-sky coverage of future surveys such as Euclid will let us estimate higher order statistics in the harmonic domain with unprecedented accuracy (Amendola et al. 2013). While measurements of real space correlations are much simpler in the presence of complicated survey design, the measurements for different angular scales can be highly correlated (Munshi 2000; Munshi & Jain 2000). In comparison measurements in the harmonic domain are less correlated and each mode contains (nearly) independent information in the limit of all-sky coverage. The primary motivation of this study is to develop analytical predictions for one such statistic called the skew-spectrum, and test them against numerical simulations. We will borrow the concepts developed for constructing skew-spectra for the study of non-Gaussianity in the context of CMBR observations (Planck Collaboration XVII2016b). However, we also include gravity-induced secondary non-Gaussianity. The skew-spectrum is the lowest order member in the family of higher order spectra (Munshi et al. 2011a, 2020). In a series of papers, the one-point statistics such as the skewness and kurtosis were generalized to two-point cumulant correlator, e.g. the two-to-one correlator and its higher order generalizations. These can be represented in the harmonic domain by their associated spectra such as the skew-spectrum (Munshi & Heavens 2010) and its higher order generalizations (Munshi et al. 2011a, 2020). These spectra have already been used to analyse WMAP9 (Smidt et al. 2010) as well as Planck data (Planck Collaboration XVII2016b). They are useful tools to separate individual contributions and estimate systematics. In this paper, we will concentrate on the projected skew-spectrum and kurt-spectrum in the context of weak lensing surveys (Munshi et al. 2011c).","Citation Text":["Bernardeau, Waerbeke & Mellier 2003"],"Functions Text":["Early studies concentrated on measurements of higher order correlation hierarchy in the angular space due to small survey size"],"Functions Label":["Background"],"Citation Start End":[[1441,1476]],"Functions Start End":[[1271,1397]]}
{"Identifier":"2019MNRAS.485.2235B__Schmalzing_&_Buchert_1997_Instance_1","Paragraph":"Analyses of 21-cm signals are mainly based on traditional N-point correlation statistics. Beyond the simplest two-point function (power spectrum), higher-order correlations are quite non-trivial to calculate and sometimes they suffer from conceptual challenges. On the other hand, the Minkowski functionals (MFs) are extremely useful tools in quantitatively describing the morphology because, in principle, they contain information on all the higher-order moments. The MFs were first introduced in cosmology by Mecke, Buchert & Wagner (1994). Since then they have been extensively employed to study the morphology of the large-scale structure of the universe and the cosmic web (Schmalzing & Buchert 1997; Sahni, Sathyaprakash & Shandarin 1998; Sathyaprakash, Sahni & Shandarin 1998; Bharadwaj et al. 2000; Hikage et al. 2003; Bharadwaj, Bhavsar & Sheth 2004; Pandey & Bharadwaj 2008; Einasto et al. 2011; Wiegand & Eisenstein 2017) as well as the CMB (Schmalzing & Gorski 1998; Novikov, Feldman & Shandarin 1999; Novikov, Schmalzing & Mukhanov 2000; Hikage, Komatsu & Matsubara 2006). Since the reionization landscape is similarly rich in geometrical properties because of growth and overlap of ionized \u2018bubbles\u2019, studying the morphology of reionization using MFs is highly compelling and feasible. The physics underlying the reionization process is expected to be manifested in the geometry and morphology of H\u2009i and H\u2009ii regions. The ratios of MFs are introduced in Sahni et al. (1998) as Shapefinders that precisely assess the shape of an object by directly estimating its physical dimensions. Therefore, using MFs and Shapefinders of the ionization field, it should be possible to probe the physics of the high-redshift universe. For instance, if reionization is driven in a non-standard manner through large energy output from multiple quasar jets, the very first ionized bubbles might be filamentary and not spherical (as they would be for point-like sources like stars in galaxies). Clearly the 3D structure of cosmological reionization in this scenario would be very different from the standard mechanism of point-like sources.","Citation Text":["Schmalzing & Buchert 1997"],"Functions Text":["On the other hand, the Minkowski functionals (MFs) are extremely useful tools in quantitatively describing the morphology because, in principle, they contain information on all the higher-order moments.","Since then they have been extensively employed to study the morphology of the large-scale structure of the universe and the cosmic web"],"Functions Label":["Background","Background"],"Citation Start End":[[679,704]],"Functions Start End":[[262,464],[543,677]]}
{"Identifier":"2018ApJ...867...55X__Acero_et_al._2015_Instance_1","Paragraph":"During the likelihood analysis, the normalizations and spectral parameters of all sources within 5\u00b0 of 3FGL J1640.4-4634c, as well as the normalizations of the two diffuse backgrounds, are left free. First, we create a Test Statistic (TS) map by subtracting the emission from the sources and backgrounds in the best-fit model with gttsmap, which is shown in the top panel of Figure 1. Some residual emission are shown in this TS map. Then we add additional point sources with power-law spectra in the model. The accurate positions of these sources obtained using the gtfindsrc tool, together with their TS values, are listed in Table 1. Next, we adopt the position of 3FGL J1640.4-4634c provisionally provided by the 3FGL catalog (Acero et al. 2015) and investigate the spectrum of 3FGL J1640.4-4634c. We bin the data into 12 equal logarithmic energy bins from 1 to 500 GeV and repeat the same likelihood fitting for each energy bin. In the model, the normalization parameters of sources within 5\u00b0 around 3FGL J1640.4-4634c and the two diffuse backgrounds are left free, while all spectral indices, except 3FGL J1640.4-4634c, are fixed. The 95% upper limits are calculated for energy bins with TS values smaller than 4. The resulting spectral energy distribution (SED) is shown in the bottom panel of Figure 1. An obvious spectral upturn is shown in the SED of 3FGL J1640.4-4634c at an energy of about 10 GeV. To test whether the upturn spectrum is intrinsic or is due to two overlapping sources, we do the same likelihood fitting using the events with energies of 1\u201310 GeV and 10\u2013500 GeV, respectively. For each analysis, we create a TS map with all sources (except 3FGL J1640.4-4634c) included in the model, which are shown in the Figure 2. The TS maps show a clear difference between two energy bands, and the centroids of emission in both energy bands deviate from that of 3FGL J1640.4-4634c. We thus expect that 3FGL J1640.4-4634c should consist of two different sources (labeled as \u201cSource A\u201d for the 1\u201310 GeV source and \u201cSource B\u201d for the 10\u2013500 GeV source), and the source in the 3FGL catalog is simply the sum of these two sources.","Citation Text":["Acero et al. 2015"],"Functions Text":["Next, we adopt the position of 3FGL J1640.4-4634c provisionally provided by the 3FGL catalog"],"Functions Label":["Uses"],"Citation Start End":[[731,748]],"Functions Start End":[[637,729]]}
{"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)"],"Functions Text":["This is illustrated in Fig. 5, where thick blue and red data points show mild high-energy excesses of the electron spectra derived by","Kerszberg et al. (2017), Kerszberg (2017) over broken power-law models derived from the fits to lower energy data."],"Functions Label":["Uses","Uses"],"Citation Start End":[[620,632]],"Functions Start End":[[486,619],[634,748]]}
{"Identifier":"2020MNRAS.495.4845C__Bernardi_et_al._2016_Instance_1","Paragraph":"The exploration of the Universe out to times earlier than the point of complete reionization is rapidly advancing. One of the most informative probes of these epochs is the 21-cm line produced by hydrogen atoms in the neutral intergalactic medium (IGM) at redshifts z > 6. This line redshifts to frequencies below 200 MHz and can be detected by low-frequency radio telescopes. Global 21-cm experiments measure the spectrum of this line averaged over the sky. The first tentative detection of the Cosmic Dawn signal was recently made by the Low-Band implementation of the Experiment to Detect the Global EoR Signature (EDGES, Bowman et al. 2018). Other global 21-cm experiments, including the Large-Aperture Experiment to Detect the Dark Ages (Bernardi et al. 2016; Price et al. 2018), the EDGES High-Band (Bowman & Rogers 2010; Monsalve et al. 2017, 2018, 2019), and the Shaped Antenna measurement of the background RAdio Spectrum (Singh et al. 2017, 2018), provide upper limits on the signal from Cosmic Dawn and the Epoch of Reionization (EoR), ruling out some extreme astrophysical scenarios. A parallel effort is being made by interferometric radio arrays that are placing upper limits on the fluctuations of the 21-cm signal, including the Donald C. Backer Precision Array for Probing the Epoch of Reionization (Kolopanis et al. 2019), the Low Frequency Array (LOFAR, Patil et al. 2017; Gehlot et al. 2019; Mertens et al. 2020), the Giant Metrewave Radio Telescope (Paciga et al. 2013), the Murchison Widefield Array (Beardsley et al. 2016; Barry et al. 2019; Li et al. 2019; Trott et al. 2020), and the Owens Valley Radio Observatory Long Wavelength Array (Eastwood et al. 2019). The most recent upper limit reported by LOFAR (Mertens et al. 2020) made it possible to place (weak) upper limits on the temperature of the neutral gas and ionization state of the Universe at z = 9.1 (Ghara et al. 2020; Mondal et al. 2020). Upcoming arrays, including the Hydrogen Epoch of Reionization Array (DeBoer et al. 2017), the Square Kilometer Array (Koopmans et al. 2015), and the New Extension in Nancay Upgrading LOFAR (Zarka et al. 2012), will provide measurements of the power spectrum over a wide range of scales and redshifts.","Citation Text":["Bernardi et al. 2016"],"Functions Text":["Other global 21-cm experiments, including the Large-Aperture Experiment to Detect the Dark Ages","provide upper limits on the signal from Cosmic Dawn and the Epoch of Reionization (EoR), ruling out some extreme astrophysical scenarios."],"Functions Label":["Background","Background"],"Citation Start End":[[743,763]],"Functions Start End":[[646,741],[958,1095]]}
{"Identifier":"2018AandA...610A..61P__Glampedakis_et_al._2006_Instance_1","Paragraph":"It was soon realized that these frequencies are probably related to oscillations of the neutron star, and several groups tried to identify them as elastic oscillations of the crust (Duncan 1998; Messios et al. 2001; Strohmayer & Watts 2005; Piro 2005; Sotani et al. 2007; 2013, 2016; Samuelsson & Andersson 2007, 2009; Steiner & Watts 2009; Deibel et al. 2014), Alfv\u00e9n oscillations (Cerd\u00e1-Dur\u00e1n et al. 2009; Sotani et al. 2008; Colaiuda et al. 2009), or coupled magneto-elastic oscillations (Levin 2006; 2007; Glampedakis et al. 2006; Gabler et al. 2011; 2012; Colaiuda & Kokkotas 2011; van Hoven & Levin 2011, 2012). The theoretical models based on the observed frequencies are very elaborate and may be able to constrain properties of high-density matter as found in the interior of neutron stars. Some of the models, for instance, require a superfluid component in the core of the star (van Hoven & Levin 2011, 2012; Glampedakis et al. 2011; Passamonti & Lander 2013, 2014; Gabler et al. 2013; 2016, and in prep.). Different models depend sensitively on the identification of the fundamental oscillation frequency, and may not explain all of the observed frequencies. Even when the fundamental frequency is identified, the interpretation and parameter estimation is not yet straightforward because of degeneracies in the parameter space. However, keeping other stellar parameters fixed, some general trends of the fundamental oscillation frequency can be summarized as follows (see Gabler et al. 2016, and in prep. for a detailed discussion): i) The frequency scales linearly with the magnetic field strength. ii) It decreases with increasing compactness (Sotani et al. 2008). The compactness is related to the hardness of the equation of state (EOS): Material with a stiff equation of state is harder to compress, leading to larger radii and hence lower compactnesses. iii) It can only reach the surface for significantly strong magnetic fields \n$\\bar B\\gtrsim\\bar\nB_\\text{outbreak}(\\sqrt{c_s})$\n\nB\u0304\u2273B\u0304outbreak(cs)\n\n\n, whose thresholds depend on the square root of the shear velocity (Gabler et al., in prep.).","Citation Text":["Glampedakis et al. 2006"],"Functions Text":["It was soon realized that these frequencies are probably related to oscillations of the neutron star, and several groups tried to identify them as","Alfv\u00e9n oscillations"],"Functions Label":["Background","Background"],"Citation Start End":[[510,533]],"Functions Start End":[[0,146],[362,381]]}
{"Identifier":"2020ApJ...893...54Y___2021_Instance_1","Paragraph":"The polarimetric observations with ALMA at 0.87 mm toward B335 were conducted during 2016 to 2018, consisting of 13 successful executions (project code: 2015.1.01018.S). In the observations, 40\u201347 antennae were used in the configurations with baseline lengths ranging from 15 to 1400 m. The pointing center was \u03b1(J2000) = 19h37d00.s89, \u03b4(J2000) = +7\u00b034\u2032096. The on-source integration time was 7.4 hr. The observations were conducted with the full polarization mode and at the frequency ranges of 335.5\u2013339.5 GHz and 347.5\u2013351.5 GHz with a total bandwidth of 8 GHz. In these observations, J1751\u22120939 was observed as the bandpass calibrator, J1938 + 0448 or J1935 + 2021 as the gain calibrators, and J1924\u22122914 or J2000\u22121748 for polarization calibration. The flux calibration was performed with the observations of quasars or the asteroid, Pallas. The data were manually calibrated by the EA ARC node using the Common Astronomy Software Applications (CASA) of the version 5.1.1 (McMullin et al. 2007). We additionally performed self-calibration of the phase using the Stokes I data. Then the calibrated visibility data were Fourier-transformed with the Briggs weighting with a robust parameter of +0.5 to generate Stokes IQU images, and the images were cleaned using the CASA task tclean. The achieved synthesized beam is 019 \u00d7 017. The noise level in the Stokes I image is 40 \u03bcJy beam\u22121, and it is 9 \u03bcJy beam\u22121 in the Stokes Q and U images. When we generated the polarized intensity (Ip) map, we debiased the polarized intensity (Ip) with \n\n\n\n\n\n, where \u03c3Q, U is the noise level in Stokes Q and U (Wardle & Kronberg 1974; Simmons & Stewart 1985). To extract polarization detections, we first binned up the Stokes IQU and Ip maps to have a pixel size of 01, which is approximately half of the beam size, and computed polarization orientations and fractions. Thus, the minimal separation between two polarization detections is 01. The Stokes I and Ip maps with their original pixel size of 002 are presented below. The polarization detections are extracted when the signal-to-noise ratios (S\/N) of both Stokes I and Ip are larger than three, and thus the expected uncertainties in the polarization orientations are \u22729\u00b0.","Citation Text":["McMullin et al. 2007"],"Functions Text":["The data were manually calibrated by the EA ARC node using the Common Astronomy Software Applications (CASA) of the version 5.1.1"],"Functions Label":["Uses"],"Citation Start End":[[977,997]],"Functions Start End":[[846,975]]}
{"Identifier":"2021ApJ...922..140R__Hughes_et_al._1998_Instance_1","Paragraph":"We fit the total foreground absorption column by applying two-component multiplicative absorption models, tbabs and tbnew.\n8\n\n\n8\n\nhttps:\/\/pulsar.sternwarte.uni-erlangen.de\/wilms\/research\/tbabs\/\n The absorption component tbabs accounts for the Galactic absorption column N\nH,Gal, and it is fixed at N\nH,Gal = 6 \u00d7 1020 cm\u22122 (Dickey & Lockman 1990). The second absorption component, tbnew, accounts for the LMC absorption column, N\nH,LMC and is varied in the fits. We use tbnew as the absorption column for LMC as it allows us to set individual elemental abundances associated with N\nH,LMC at their respective LMC values. Recent measurements of the LMC abundances based on the X-ray spectral analysis of SNRs (Maggi et al. 2016; Schenck et al. 2016) suggest \u223c50% lower LMC abundance values for O, Ne, Mg, and Fe than the previously estimated values (Hughes et al. 1998). Our derived shock parameters are consistent (within statistical uncertainties), assuming either set of these LMC abundances in tbnew. We tie N\nH,LMC among all the fitted spectra, assuming no temporal variation in the LMC column for SNR 1987A. We obtain a best-fit total absorption column density, \n\n\n\nNH=NH,Gal+NH,LMC=2.17\u22120.22+0.22\u00d71021cm\u22122\n\n. This value is comparable to estimates of N\nH obtained by other X-ray analyses: \n\n\n\n2.35\u22120.08+0.09\u00d71021cm\u22122\n\n (Park et al. 2006), \n\n\n\n1.44\u22120.12+0.16\u00d71021cm\u22122\n\n(Zhekov et al. 2009), and \n\n\n\n2.60\u22120.05+0.05\u00d71021cm\u22122\n\n (Alp et al. 2021). We note that more recent H i surveys have suggested much higher N\nH,Gal values toward the LMC, i.e., \u223c(2.5\u20134) \u00d7 1021 cm\u22122 (Kalberla et al. 2005; Willingale et al. 2013; HI4PI Collaboration et al. 2016). By adopting these high Galactic columns, the overall fits are equally good, but this results in a negligible LMC column. The implied negligible LMC column does not appear to be reasonable because optical extinction estimates show that the LMC contribution is greater than the Milky-Way contribution (Fitzpatrick & Walborn 1990; France et al. 2011). A detailed analysis of the contribution of the Galactic N\nH toward total absorption column density is beyond the scope of our work. While this issue was similarly outlined by Alp et al. (2021) in relation to their analysis of the XMM-Newton data of SNR 1987A, we note that the total columns (Galactic + LMC) are generally consistent between either values of the Galactic column, and thus that the best-fit parameters in our spectral model fits (i.e., electron temperatures, ionization age, abundances, and volume-emission measures) are not significantly affected (within statistical uncertainties).","Citation Text":["Hughes et al. 1998"],"Functions Text":["Recent measurements of the LMC abundances based on the X-ray spectral analysis of SNRs","suggest \u223c50% lower LMC abundance values for O, Ne, Mg, and Fe than the previously estimated values","Our derived shock parameters are consistent (within statistical uncertainties), assuming either set of these LMC abundances in tbnew."],"Functions Label":["Differences","Differences","Differences"],"Citation Start End":[[847,865]],"Functions Start End":[[619,705],[747,845],[868,1001]]}
{"Identifier":"2021MNRAS.506.3313G__Bluck_et_al._2014_Instance_1","Paragraph":"The B\/T ratio of a galaxy is the fraction of total luminosity contributed by the bulge component of the galaxy. Bulges are the central component of disc galaxies which appear as central bright cores in galaxy images or as excess-light over the disc light in the inner region of galaxy light profiles. The B\/T ratio is an important indicator of galaxy structure and is well correlated with several other physical quantities of interest in studies of galaxy evolution, such as galaxy morphology (Graham & Worley 2008), kinematics (Cappellari et al. 2013), stellar mass, and star formation rate (Bluck et al. 2014). The B\/T ratio also directly determines the bulge luminosity which in turn correlates with the mass of the central supermassive black hole of the galaxy (Kormendy & Richstone 1995; Marconi & Hunt 2003; Kormendy & Ho 2013). Galaxies having different B\/T ratios are thought to have undergone evolution along different evolutionary paths. A high value of B\/T of a galaxy generally indicates its evolutionary history dominated by galaxy mergers (Hopkins et al. 2010). On the other hand, majority of low B\/T systems are often pseudobulge hosting galaxies (Fisher & Drory 2008; Gadotti 2009) which are thought to undergo slow evolution through internal, secular evolution processes (Kormendy & Kennicutt 2004). It must be noted that unlike visual morphology which is a qualitative measure, the B\/T ratio is a quantitative measure of the galaxy morphology. It is the only reliable way to separate ellipticals and disc galaxies when morphological features such as spiral arms cannot be resolved by the telescope (due to resolution or sensitivity limitations, particularly at high redshifts). The quantitative nature of the B\/T ratio also allows direct comparison of galaxies with their counterparts in cosmological simulations. One can study these simulated galaxies to understand the origin of various properties of real galaxies. All these factors make the B\/T ratio an important parameter to measure in studies of galaxy formation and evolution.","Citation Text":["Bluck et al. 2014"],"Functions Text":["The B\/T ratio is an important indicator of galaxy structure and is well correlated with several other physical quantities of interest in studies of galaxy evolution, such as","stellar mass, and star formation rate"],"Functions Label":["Background","Background"],"Citation Start End":[[593,610]],"Functions Start End":[[301,474],[554,591]]}
{"Identifier":"2019ApJ...887..174V__Barkana_&_Loeb_2001_Instance_1","Paragraph":"Supermassive black holes (SMBH) with masses M\u2022 \u223c 109\u20131010\n\n\n\n\n\n are recognized to be present at redshifts as high as z \u2243 6\u20137.5, when the universe was 650\u2013800 Myr young (Fan et al. 2003; Willott et al. 2010; Mortlock et al. 2011; Ba\u00f1ados et al. 2014, 2018; Wu et al. 2015; Decarli et al. 2018; Izumi et al. 2019)\u2014in total more than 150 such SMBHs are already known (see, e.g., Fan et al. 2019). Their origin remains elusive\u2014it is unclear how massive their seeds were, how efficient their growth rate was, and what was the mass reservoir for their growth. To fit the existence of the quasars J0100 + 2802 (z = 6.33), J1120 + 0641 (z = 7.09), and J1342 + 0928 (z = 7.54) with SMBH masses M\u2022 = 1.2 \u00d7 1010,  2 \u00d7 109, and 7.8 \u00d7 108\n\n\n\n\n\n respectively, one has to assume that their masses grow as M\u2022 = M\u2022,0 exp [t\/(47 Myr)], corresponding to the standard Eddington limit with a 10% radiative efficiency , the Salpeter growth time tS = c\u03c3T\/(4\u03c0Gmp) = 47 Myr, and the seed mass M\u2022,0 \u2265 103\n\n\n\n\n\n at z \u2265 40 (Ba\u00f1ados et al. 2018). In this scenario SMBHs have to begin growing even earlier than the very first stars are assumed to have appeared (see discussion in Barkana & Loeb 2001). Moreover, it suggests that the accretion is tightly tuned to the Eddington rate, which seems physically unlikely (see discussion in Haiman & Loeb 2001; Volonteri & Rees 2005; Haiman 2013; Alexander & Natarajan 2014; Madau et al. 2014). Lower-mass black holes with M\u2022 \u223c 100 \n\n\n\n\n\n, originated from Population III stars, are apparently unlikely to serve as seeds for growing SMBHs, because photoionization and photoheating from their massive progenitors strongly suppress further supply of cold mass onto the BH (Johnson & Bromm 2007), and would require even more time for the black hole to grow. Note, however, that this channel for SMBH seeds is currently widely discussed (see references in Natarajan et al. 2019). Also note that this problem of the presence of such enormously massive BHs in a younger than 1 Gyr universe can be to a certain extent eased when possible magnification of z > 6 SMBHs due to gravitational lensing is considered (Fan et al. 2019; Pacucci & Loeb 2019a, 2019b). Recent millimeter observations of the quasar J0100 + 2802 (z = 6.33) with the most massive BH, M\u2022 \u223c 1.2 \u00d7 1010\n\n\n\n\n\n, known at z > 6, indicate strong lensing with a magnification factor of \u223c450 (Fujimoto et al. 2019). As a result, the estimate of the SMBH mass may be reduced by more than an order of magnitude, though this still remains exceedingly large for a 1 Gyr universe \u223c109\n\n\n\n\n\n (Fujimoto et al. 2019). The fraction of so strongly magnified quasars is fairly low considering the very small typical angular size of such lenses, as a rule \u226a1\u2033 (see, e.g., Pei 1995; Bolton et al. 2008; Pacucci & Loeb 2019b).","Citation Text":["Barkana & Loeb 2001"],"Functions Text":["In this scenario SMBHs have to begin growing even earlier than the very first stars are assumed to have appeared (see discussion in"],"Functions Label":["Uses"],"Citation Start End":[[1150,1169]],"Functions Start End":[[1018,1149]]}
{"Identifier":"2018MNRAS.476.2591V__Patton_et_al._2016_Instance_1","Paragraph":"Galaxy interactions represent a fundamental component of our current view of hierarchical galaxy evolution. Studies based on both observations and simulations have shown that galaxy collisions and mergers can dramatically affect the galaxies undergoing the interaction, by, e.g. triggering nuclear activity (e.g. Kennicutt 1984; Kennicutt et al. 1987; Ellison et al. 2011, 2013a; Silvermann et al. 2011; Satyapal et al. 2014), producing colour changes (e.g. Larson & Tinsley 1978; Darg et al. 2010; Patton et al. 2011), disrupting morphologies (e.g. Kaviraj et al. 2011; Patton et al. 2016; Lofthouse et al. 2017), and altering the metallicities (e.g. Rupke et al. 2010; Perez, Michel-Dansac & Tissera 2011; Scudder et al. 2012; Torrey et al. 2012). The most evident effect driven by galaxy encounters is probably the triggering of new episodes of star formation, which can occur both in the pre-merger regime between first pericentre and coalescence (e.g. Nikolic, Cullen & Alexander 2004; Ellison et al. 2008, 2013b; Patton et al. 2011; Scudder et al. 2012), and in the post-merger phase, where the two nuclei of the interacting galaxies have merged together (e.g. Kaviraj et al. 2012; Kaviraj 2014; Ellison et al. 2013a). The idea that galaxy mergers have a strong impact on the star formation activity is supported by studies of Ultra-Luminous InfraRed Galaxies (ULIRGs), i.e. galaxies with IR luminosities exceeding 1012 L\u2299 and characterized by star formation rates (SFRs) up to \u223c1000 M\u2299 yr\u22121 (e.g. Barnes & Hernquist 1991; Mihos & Hernquist 1994; Daddi et al. 2010; Scoville et al. 2015). Observations have revealed that the majority of ULIRGs reside in interacting systems (e.g. Sanders & Mirabel 1996; Veilleux, Kim & Sanders 2002; Kartaltepe et al. 2010, 2012; Haan et al. 2011). Nevertheless, ULIRGs are rare and extreme examples of highly star-forming galaxies. Most galaxy\u2013galaxy interactions result in SFR increases of at most a factor of a few, as shown in both numerical simulations (e.g. Di Matteo et al. 2008) and observations of galaxy pairs and post-mergers (Ellison et al. 2008; Martig & Bournaud 2008; Jogee et al. 2009; Robaina et al. 2009; Scudder et al. 2012).","Citation Text":["Patton et al. 2016"],"Functions Text":["Studies based on both observations and simulations have shown that galaxy collisions and mergers can dramatically affect the galaxies undergoing the interaction, by, e.g.","disrupting morphologies (e.g."],"Functions Label":["Background","Background"],"Citation Start End":[[571,589]],"Functions Start End":[[108,278],[520,549]]}
{"Identifier":"2021MNRAS.505..435S___2020_Instance_1","Paragraph":"Detections of ionic, atomic, and molecular species in exoplanetary atmospheres serve as a unique and strong diagnostic of those chemical and dynamical processes driving their formation and evolution. Their detection and abundance measurements could act as indicators of planetary formation scenarios and reveal connections to the primordial protoplanetary disc and the host star (Williams & Cieza 2011; Mordasini et al. 2016; Madhusudhan et al. 2017). Furthermore, discoveries of atmospheric chemical species allow us to better understand various thermodynamical processes and chemistry, winds in the upper atmosphere (Goodman 2009; Snellen et al. 2010; Brogi et al. 2016; Madhusudhan et al. 2016; Wyttenbach et al. 2020), and to probe planetary interiors and various bulk properties through their abundances (Kite et al. 2016; Thorngren & Fortney 2019; Madhusudhan et al. 2020). A whole host of ions, atoms, and molecules have been detected through a variety of, often complementary, techniques, such as differential spectrophotometry using low-to-mid resolution spectroscopy (e.g. Gibson et al. 2012, 2017; Deming et al. 2013; Kreidberg et al. 2014; Kirk et al. 2016; Nortmann et al. 2016), and high resolution spectroscopic techniques (e.g. Redfield et al. 2008; Snellen et al. 2008; Rodler, Lopez-Morales & Ribas 2012; Birkby et al. 2013; Hoeijmakers et al. 2015, 2018, 2020; Brogi et al. 2016; Birkby et al. 2017; \u017d\u00e1k et al. 2019; Ehrenreich et al. 2020). To date, ionic species such as Fe\u2009ii and Ti\u2009ii (Hoeijmakers et al. 2019), atomic absorption from Na, K, H \u03b1, and He (e.g. Redfield et al. 2008; Sedaghati et al. 2016; Casasayas-Barris et al. 2017; Spake et al. 2018; Chen et al. 2020; Seidel et al. 2020), and molecules such as H2O, CH4, and CO (e.g. Konopacky et al. 2013; Brogi et al. 2014; Fraine et al. 2014; Barman et al. 2015; Sing et al. 2016) have been detected through the aforementioned techniques. Needless to say that this list of detected constituents is by no means exhaustive, nor that of methods employed to detect exoplanetary atmospheres. For instance, high-resolution imaging instruments such as SPHERE (Beuzit et al. 2019) and GRAVITY (Gravity Collaboration et al. 2017), both at the VLT (ESO\u2019s Very Large Telescope), through combination with low-dispersion spectroscopy, have facilitated direct measurements of exoplanetary atmospheres (Samland et al. 2017; Gravity Collaboration et al. 2020).","Citation Text":["Hoeijmakers et al.","2020"],"Functions Text":["A whole host of ions, atoms, and molecules have been detected through a variety of, often complementary, techniques,","and high resolution spectroscopic techniques (e.g."],"Functions Label":["Background","Background"],"Citation Start End":[[1343,1361],[1374,1378]],"Functions Start End":[[880,996],[1193,1243]]}
{"Identifier":"2020MNRAS.497.3504T__Migliori_et_al._2017_Instance_1","Paragraph":"Galactic BHXBs offer excellent laboratories in which to study jet interaction zones, as their jets evolve on day to month time-scales, they are located at nearby distances, and they are thought to be good analogues for AGNs. Given their incredible diagnostic potential, over the last couple of decades there have been many observational campaigns searching for these highly sought after interaction sites near BHXBs. To date, several candidate jet\u2013ISM interaction sites have been identified; SS 433 (Dubner et al. 1998), Cygnus X\u20131 (Gallo et al. 2005; Russell et al. 2007), 1E 1740\u20132942 (Mirabel et al. 1992), GRS 1758\u2212258 (Mart\u00ed et al. 2002), GRS 1915+105 (Rodr\u00edguez & Mirabel 1998; Chaty et al. 2001; Kaiser et al. 2004), H1743\u2013322 (Corbel et al. 2005), XTE J1550\u2013564 (Corbel et al. 2002; Kaaret et al. 2003; Migliori et al. 2017), XTE J1748\u2013288 (Brocksopp et al. 2007), GRO J1655\u201340 (Hjellming & Rupen 1995; Hannikainen et al. 2000), GX 339-4 (Gallo et al. 2004), 4U 1755\u201333 (Kaaret et al. 2006), XTE J1752\u2013223 (Yang et al. 2010; Miller-Jones et al. 2011; Yang et al. 2011; Ratti et al. 2012), XTE J1650\u2013500 (Corbel et al. 2004), XTE J1908+094 (Rushton et al. 2017), 4U 1630\u201347 (Neilsen et al. 2014; Kalemci, Maccarone & Tomsick 2018), LMC X\u20131 (Russell et al. 2006; Cooke et al. 2007; Hyde et al. 2017), and GRS 1009\u201345 (Russell et al. 2006). From these past works, it is clear that finding and confirming interaction sites can be incredibly difficult and often observationally expensive (e.g. requiring deep, wide-field radio continuum observations). This difficulty mainly results from the fact that interaction sites can manifest with a wide variety of morphologies and emission properties, likely dependent on the properties of the BHXB (e.g. space velocity; Miller-Jones et al. 2008; Heinz et al. 2008; Wiersema et al. 2009) and\/or local ISM properties (e.g. density; Heinz 2002; Kaiser et al. 2004). Once identified, detailed calorimetric calculations of interaction sites are highly sensitive to the properties of the ISM (i.e. density, kinetic temperature, shock velocity; Russell et al. 2007; Sell et al. 2015), and thus require accurate constraints on the physical conditions in the interacting gas, which cannot be derived with continuum observations alone. Therefore, developing and implementing new methods that allow us to identify and place improved observational constraints on these parameters at multiple interaction sites, is crucial for taking full advantage of the diagnostic potential of these regions.","Citation Text":["Migliori et al. 2017"],"Functions Text":["To date, several candidate jet\u2013ISM interaction sites have been identified;","XTE J1550\u2013564","From these past works, it is clear that finding and confirming interaction sites can be incredibly difficult and often observationally expensive (e.g. requiring deep, wide-field radio continuum observations)."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[811,831]],"Functions Start End":[[417,491],[756,769],[1346,1554]]}
{"Identifier":"2020MNRAS.499.4158G__Fialkov_&_Loeb_2016_Instance_1","Paragraph":"In 2018, a deep spectral feature centred at 78 MHz was reported by the EDGES collaboration (Bowman et al. 2018). The feature was presented as the long sought-after 21-cm absorption feature seen against the CMB during the CD at z \u223c 17. The location of this putative absorption trough is consistent with redshift predictions from theoretical models and simulations of the CD (Furlanetto et al. 2006; Pritchard & Loeb 2010; Mesinger, Ferrara & Spiegel 2013; Cohen et al. 2017). However, the depth of the feature is \u0394T21 \u223c 0.5 K ($99{{\\ \\rm per\\ cent}}$ confidence level), which is two to three times stronger and considerably wider (\u0394\u03bd \u223c 19 MHz) than that predicted by the most optimistic astrophysical models (e.g. Pritchard & Loeb 2010; Fialkov, Barkana & Visbal 2014; Fialkov & Loeb 2016; Cohen et al. 2017). Moreover, the observed feature is flat-bottomed instead of a smooth Gaussian-like shape. Several \u2018exotic\u2019 theoretical models have already been proposed which might explain the depth of the feature, such as a considerably colder IGM due to interaction between baryons and dark matter particles causing a lower spin-temperature and therefore a deeper absorption feature (e.g. Barkana 2018; Fialkov, Barkana & Cohen 2018), or a stronger radiation background against which the absorption is taking place (e.g. Dowell & Taylor 2018; Ewall-Wice et al. 2018; Feng & Holder 2018; Fialkov & Barkana 2019). Although the 21-cm signal is expected to be stronger at these redshifts, the foreground emission is several times brighter at these frequencies compared to EoR 21-cm signal observations at 150 MHz (Bernardi et al. 2009, 2010). Moreover, ionospheric effects are amplified at lower frequencies (de Gasperin et al. 2018; Gehlot et al. 2018), rendering the measurement of the signal equally (or even more) challenging than in EoR experiments. As of now, Ewall-Wice et al. (2016) reported a systematics-limited power spectrum upper limit of $\\Delta _{21}^2 \\lt (10^4~\\text{mK})^2$ on co-moving scales $k\\lesssim 0.5~h\\, \\text{cMpc}^{-1}$ (in 3 h of integration) on the 21-cm signal brightness temperature in the redshift range 12 \u2272 z \u2272 18 using MWA. This overlaps with the low-redshift edge of the 21-cm absorption feature (Bowman et al. 2018). Gehlot et al. (2019) provided a 2\u03c3 upper limit of $\\Delta _{21}^2 \\lt (1.4\\times 10^4~\\text{mK})^2$ on the 21-cm signal power spectrum at $k = 0.038~h\\, \\text{cMpc}^{-1}$ (in 14 h of integration) using the LOFAR-Low Band Antenna (LBA) system in the redshift range 19.8 \u2272 z \u2272 25.2, which corresponds to the high-redshift edge of the absorption feature. More recently, Eastwood et al. (2019) used OVRO-LWA observations to report a 2\u03c3 upper limit of $\\Delta _{21}^2 \\lt (10^4~\\text{mK})^2$ at $k \\approx 0.1~h\\, \\text{cMpc}^{-1}$ (in 28 h of integration) at redshift z \u2248 18.4.","Citation Text":["Fialkov & Loeb 2016"],"Functions Text":["However, the depth of the feature is \u0394T21 \u223c 0.5 K ($99{{\\ \\rm per\\ cent}}$ confidence level), which is two to three times stronger and considerably wider (\u0394\u03bd \u223c 19 MHz) than that predicted by the most optimistic astrophysical models (e.g."],"Functions Label":["Differences"],"Citation Start End":[[768,787]],"Functions Start End":[[475,712]]}
{"Identifier":"2018MNRAS.479.3254V___2000_Instance_2","Paragraph":"The lifetime of molecular clouds (MCs) remains an active research topic in the study of the interstellar medium and star formation, and most recent studies, both observational and theoretical, place this lifetime at a few times 107 yr for clouds in the 105\u2013106M\u2299 mass range (e.g. Blitz & Shu 1980; Kawamura et al. 2009; Zamora-Avil\u00e9s, V\u00e1zquez-Semadeni & Col\u00edn 2012; Zamora-Avil\u00e9s & V\u00e1zquez-Semadeni 2014; Lee, Miville-Desch\u00eanes & Murray 2016). In addition, several observational studies have suggested that the star formation rate (SFR) of the clouds appears to increase over their lifetimes. For example, studies of young clusters embedded in moderate-mass MCs (\u223c104M\u2299) (e.g. Palla & Stahler 1999, 2000; Da Rio et al. 2010) have shown that their age histograms contain a large majority of young (1\u20132 Myr) objects, but also a tail of older (up to several Myr) ones suggesting an accelerating star formation activity, sometimes followed by a subsequent decline (see also Povich et al. 2016; Schneider et al. 2018). In addition, Kawamura et al. (2009) reported a clear evolutionary process over the lifetime of giant molecular clouds (GMCs; of masses \u223c105\u2013106M\u2299) in the Large Magellanic Cloud, evidenced by the increasing number of massive stars across the sequence of GMC \u2018classes\u2019 proposed by those authors. Finally, on the basis of the large scatter in the observed star formation efficiency (SFE) in Milky Way GMCs, Lee et al. (2016) have concluded that the SFR in those clouds must also be time variable. Numerical simulations of MC formation and evolution also exhibit time varying, increasing SFRs during their early stages (e.g. V\u00e1zquez-Semadeni et al. 2007; Hartmann, Ballesteros-Paredes & Heitsch 2012). Also, in the presence of stellar feedback, at late times the SFRs reach a maximum and begin to decrease again (e.g. V\u00e1zquez-Semadeni et al. 2010; Col\u00edn, V\u00e1zquez-Semadeni & G\u00f3mez 2013). V\u00e1zquez-Semadeni, Gonz\u00e1lez-Samaniego & Col\u00edn (2017) have recently shown that the simulations of Col\u00edn et al. (2013) in fact produce stellar age histograms highly resemblant of the observed ones (Palla & Stahler 1999, 2000; Da Rio et al. 2010), and reproduce observed radial age gradients in clusters (Getman et al. 2014) as well as bottom-heavy stellar initial mass functions (IMFs) in scattered regions of massive star formation (Povich et al. 2016).","Citation Text":["Palla & Stahler","2000"],"Functions Text":["V\u00e1zquez-Semadeni, Gonz\u00e1lez-Samaniego & Col\u00edn (2017) have recently shown that the simulations of Col\u00edn et al. (2013) in fact produce stellar age histograms highly resemblant of the observed ones"],"Functions Label":["Similarities"],"Citation Start End":[[2092,2107],[2114,2118]],"Functions Start End":[[1897,2090]]}
{"Identifier":"2017MNRAS.469.3108C__Diamond-Stanic_et_al._2012_Instance_1","Paragraph":"Still unconstrained are the mechanisms that are able to modify both SF and morphology and their corresponding time-scales. Several hypotheses have been proposed to explain the quenching of the SF in blue galaxies, such as gas stripping (e.g. Gunn & Gott 1972), morphological or gravitational quenching (Martig et al. 2009; Genzel et al. 2014), shock heating of infalling cold gas by the hot halo (Dekel & Birnboim 2006) or an exhaustion of the gas supply (e.g. Larson, Tinsley & Caldwell 1980). Moreover, in massive galaxies, the role of AGNs in influencing galaxy evolution and quenching SF is supported by several observations (Hopkins et al. 2005; Kaviraj et al. 2007; Diamond-Stanic et al. 2012; Fabian 2012, and references therein, Cimatti et al. 2013; Cicone et al. 2014; F\u00f6rster Schreiber et al. 2014) and corroborated by the theoretical results obtained combining N-body simulations of dark matter haloes evolution (Springel et al. 2005; Boylan-Kolchin et al. 2009) with semi-analytic models for galaxy formation (White & Frenk 1991; Springel et al. 2005; Lu et al. 2011; Benson 2012). However, other models are capable of forming rapidly quiescent galaxies without invoking the AGN feedback (e.g. Khochfar & Silk 2006; Naab, Khochfar & Burkert 2006, 2009; Johansson, Thomas & Maraston 2012). Stellar or supernova (SN) feedback is most likely channel for the SF quenching in low-mass galaxies (e.g. 1010\u2009M\u2299; Kaviraj et al. 2007). Several mechanisms have also been invoked to explain the morphological transformation. Numerical simulations have shown that major merging can give rise to elliptical and S0 galaxies (Bekki 1998) and that also minor merging can play an important role in spheroid and bulge growth (Bournaud, Jog & Combes 2007; Naab et al. 2007). From an observational point of view, evidence that the morphological transformation can also be induced by environmental mechanisms (Larson et al. 1980; Farouki & Shapiro 1981; Moore et al. 1999; Quilis, Moore & Bower 2000; ; Bekki, Couch & Shioya 2002) or by the secular growth of pseudo-bulges (Courteau, de Jong & Broeils 1996; Norman, Sellwood & Hasan 1996; MacArthur, Courteau & Holtzman 2003; Kormendy & Kennicutt 2004; Debattista et al. 2006) has been found.","Citation Text":["Diamond-Stanic et al. 2012"],"Functions Text":["Moreover, in massive galaxies, the role of AGNs in influencing galaxy evolution and quenching SF is supported by several observations"],"Functions Label":["Background"],"Citation Start End":[[672,698]],"Functions Start End":[[495,628]]}
{"Identifier":"2022MNRAS.517.1218L__Barnes_&_Hernquist_1996_Instance_1","Paragraph":"Several consortia have been actively approaching the question based on large photometric and spectroscopic surveys such as 2dF and SDSS, covering large sky areas and redshift ranges (Lewis et al. 2002; Balogh et al. 2004; Kauffmann et al. 2004; Poggianti et al. 2017). All these efforts have found indisputable evidence for the impact of environment in galaxy evolution (and references therein; Boselli & Gavazzi 2006, 2014; Cortese et al. 2021). However, understanding the role played by different physical mechanisms exerted under diverse environment conditions constitutes the matter of a very active debate. The involved physical mechanisms are classified in two types: hydrodynamic and gravitational. Hydrodynamic effects concern the stripping of cold\/warm interstellar gas (H\u2009i and H2) by the hot intracluster medium (ICM). The ram-pressure stripping (RPS, Gunn & Gott 1972) and the viscous stripping (Nulsen 1982) are the most studied cases. On the other side, we have the tidal (gravitational) mechanisms occurring between a galaxy and the cluster potential (Byrd & Valtonen 1990) or among neighbour galaxies (Merritt 1983; Barnes & Hernquist 1996; Walker, Mihos & Hernquist 1996). These include major mergers, accretion of low-mass satellites, and the accumulation of fast speed encounters between galaxies (the galaxy harassment, Moore, Katz & Lake 1996). The removal of the galaxy halo gas, known as galaxy starvation (e.g. Larson, Tinsley & Caldwell 1980), is predicted to occur either by hydrodynamic or gravitational interactions. Most of these mechanisms are predicted to transform a spiral galaxy into an S0, and it is known that more than one process might act simultaneously on a single galaxy. The pre-processing of galaxies occuring within groups infalling towards clusters and seems to be particularly important (Donnari et al. 2021). Groups of galaxies are known to have lower velocity dispersions than clusters, allowing slower and deeper tidal interactions among their members. Several authors (Fadda et al. 2008; Poggianti et al. 2009b) provided substantial evidence that strong galaxy evolution is occurring in low-mass systems at a large distance from the cluster core. However, the debate pre-processingversuscluster effects remains open because many variables are involved, such as the infalling orbits, initial gas\/stellar masses, the group\/cluster properties, and even the surrounding large-scale structure (Rhee et al. 2020; Salerno et al. 2020).","Citation Text":["Barnes & Hernquist 1996"],"Functions Text":["On the other side, we have the tidal (gravitational) mechanisms occurring","or among neighbour galaxies"],"Functions Label":["Background","Background"],"Citation Start End":[[1132,1155]],"Functions Start End":[[949,1022],[1089,1116]]}
{"Identifier":"2021ApJ...920L..31N__Sterling_et_al._2017_Instance_1","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"],"Functions Text":["In fact, homologous jets, continuing for hours at a time, have been commonly observed (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[719,739]],"Functions Start End":[[557,649]]}
{"Identifier":"2017ApJ...839...56T___2013_Instance_2","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"],"Functions Text":["More evolved disks do not show a degree of linear polarization larger than 0.5%"],"Functions Label":["Background"],"Citation Start End":[[1143,1156],[1163,1167]],"Functions Start End":[[1062,1141]]}
{"Identifier":"2019MNRAS.487..475C__Ward-Thompson_et_al._2009_Instance_2","Paragraph":"We attempted to find a correlation between the mean magnetic field and the outflow and minor axis of the cloud CB 17. Relative orientations between various quantities of CB 17 are presented in the first row of Table 6, along with a comparative study of the same for some dark clouds. The first column gives the cloud ID and columns 2\u20136 give the position angles of the mean magnetic field at the envelope ($\\lt \\theta ^{\\rm env}_B\\gt $), mean magnetic field at the core ($\\lt \\theta ^{\\rm core}_B\\gt $), outflow axis (\u03b8out), minor axis (\u03b8min) of the core of the cloud and Galactic plane (\u03b8GP), respectively. $\\lt \\theta ^{\\rm env}_B\\gt $ of CB 17 is found to be almost aligned along the Galactic plane over that region of the sky (column 7), which indicates the dominance of the Galactic magnetic field over the envelope magnetic field of the cloud and thus we cannot infer much about the magnetic field structure from the optical study. A similar feature has also been observed in the cases of CB 34 (Das et al. 2016), L328, L673-7 (Soam et al. 2015), CB 26 (Halder et al. 2019), CB 3 and CB 246 (Ward-Thompson et al. 2009). However, $\\lt \\theta ^{\\rm core}_B\\gt $ of CB 17 (obtained by submm polarimetry) turned out to be perpendicular to $\\lt \\theta ^{\\rm env}_B\\gt $ (column 8); a similar phenomenon has been observed in the case of CB 34-C1 (Das et al. 2016), IRAM 04191 (Soam et al. 2015) and CB 54 (Wolf et al. 2003). Since, in the case of CB 17, $\\lt \\theta ^{\\rm env}_B\\gt $ is along the Galactic plane orientation, this implies that only $\\lt \\theta ^{\\rm core}_B\\gt $ (denser region) is linked with the ongoing physical phenomena in the cloud. $\\lt \\theta ^{\\rm core}_B\\gt $ is oriented perpendicular to \u03b8GP as well (column 9) and a similar orientation has been observed in the case of CB 34-C1 (Das et al. 2016), IRAM 04191 (Soam et al. 2015), CB 230 and CB 244 (Wolf et al. 2003) as well. Moreover, $\\lt \\theta ^{\\rm core}_B\\gt $ is found to be almost aligned along the minor axis of the core of the cloud; the angular offset is nearly 5.9\u00b0 (column 10). The alignment of $\\lt \\theta ^{\\rm core}_B\\gt $ with the minor axis of the cloud fits the magnetically regulated star formation model, in which the magnetic field should lie along the minor axis of the cloud (Mouschovias & Morton 1991; Li 1998), and the same feature has also been observed for the clouds CB 34-C1 (Das et al. 2016) and IRAM 04191 (Soam et al. 2015). The angular offset between $\\lt \\theta ^{\\rm core}_B\\gt $ and the outflow axis is found to be 80.9\u00b0 (column 11), that is, the core-scale magnetic field is oriented almost perpendicular to the outflow direction and a similar phenomenon has also been observed in the case of CB 34 (Das et al. 2016), CB 68 (Bertrang et al. 2014), B335, CB 230, CB 244 (Wolf et al. 2003) and CB 3 (Ward-Thompson et al. 2009). The angular offset between \u03b8out and \u03b8min is found to be 75\u00b0 and the same feature has been observed for CB 34-C1 (Das et al. 2016).","Citation Text":["Ward-Thompson et al. 2009"],"Functions Text":["and a similar phenomenon has also been observed in the case of","CB 3"],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[2811,2836]],"Functions Start End":[[2643,2705],[2805,2809]]}
{"Identifier":"2019MNRAS.487.1210T__McNamara_&_Nulsen_2007_Instance_1","Paragraph":"On larger scales, the clusters in which BCGs reside can generally be divided into two categories: cool core clusters, which exhibit very peaked surface brightness distributions at X-ray wavelengths, and non cool core clusters, with similar overall X-ray luminosities but with smoother, less peaked X-ray surface brightness distributions. Some authors (e.g. Hudson et al. 2010; Santos et al. 2010) define an intermediate category called moderate or weak cool core clusters. Since cool core clusters have short radiative cooling time-scales on the order of 108 yr in their centres (e.g. Voigt & Fabian 2004; McNamara & Nulsen 2007, 2012; Hlavacek-Larrondo et al. 2012), starbursts are expected to be common at the centre of such clusters. Indeed, the central cool gas in these clusters should condense onto the BCG, forming stars at rates of hundreds of solar masses per year (e.g. Fabian 1994). However, most BCGs are relatively quiescent and those that do show evidence of star formation generally tend to have star formation rates 1 order of magnitude smaller, on the order of $1-150 \\, \\mathrm{M_{\\odot }\\, {yr}^{-1}}$ (e.g. Donahue et al. 2007; Bildfell et al. 2008; O\u2019Dea et al. 2008, 2010; Rawle et al. 2012). This mismatch between expected and observed star-forming rates, known as the cooling flow problem, is thought to be caused by active galactic nuclei (AGNs) feedback processes from the BCG. AGNs can release copious amounts of energy into the intracluster medium (ICM) through many ways, including: jetted outflows that inflate cavities, weak shocks, sound waves, or turbulence in the ICM (e.g. Markevitch & Vikhlinin 2007; McNamara & Nulsen 2007, 2012; Zhuravleva et al. 2014; Fabian et al. 2017). Alone, the energy released by jetted outflows appears to be on the same order as the energy needed to offset cooling (e.g. Rafferty et al. 2006; McNamara & Nulsen 2007; Hlavacek-Larrondo et al. 2012), therefore suggesting that AGN feedback is a good candidate for solving the cooling flow problem.","Citation Text":["McNamara & Nulsen 2007"],"Functions Text":["Since cool core clusters have short radiative cooling time-scales on the order of 108 yr in their centres (e.g.","starbursts are expected to be common at the centre of such clusters."],"Functions Label":["Uses","Uses"],"Citation Start End":[[606,628]],"Functions Start End":[[473,584],[668,736]]}
{"Identifier":"2022MNRAS.515.1942D___2022b_Instance_1","Paragraph":"Comparison of 1D marginal posterior distributions over the parameters S8 \u2261 \u03c38(\u03a9m\/0.3)0.5, \u03c38 and \u03a9m, from DES Y3 data as well as other experiments, and consistency tests for this work (in blue). (a) Constraints obtained from the harmonic (this work) and real (Amon et al. 2022; Secco et al. 2022) space analyses of DES Y3 data are shown in blue and green (see also Fig. 14), both with and without shear ratio information (SR; S\u00e1nchez et al. 2021). (b) Constraints from other weak lensing surveys, namely HSC Y1 (Hikage et al. 2019; Hamana et al. 2020, 2022b), KiDS-1000 (Asgari et al. 2021), and KiDS-450 (Hildebrandt et al. 2017; K\u00f6hlinger et al. 2017) are shown in grey, and constraints from cosmic microwave background observations from Planck 2018 are shown in yellow (Planck Collaboration VI 2020). (c) Constraints from four weak lensing analyses of DES Y3 data are compared, including the analysis of mass map moments (Gatti et al. 2021b) and peaks (Z\u00fcrcher et al. 2022), and illustrating a high level of consistency (see also Fig. 15). (d) Consistency tests where redshift bins are removed one at a time (first four) and where the data vector is split into its large- and small-scale data points (last two) (see also Appendix C). (e) Various other consistency tests: removing autopower spectra, swapping the covariance matrix, and marginalizing over redshift distribution uncertainties with HyperRank and MultiRank(see also Appendix C). (f) Modelling robustness test for intrinsic alignment (IA), including B-mode power spectra, or replacing TATT by NLA, or removing IA contributions altogether (see also Section 6.2, Fig. 12). (g) Other robustness test, freeing the dark energy equation-of-state w or fixing the neutrino mass to 0.06 eV. (h) Baryonic feedback tests where the matter power spectrum is computed with HMCode instead of HaloFit, and fiducial scale cuts are replaced with kmax = 1, 3, and 5 $h\\, {\\rm Mpc}^{-1}$ scale cuts (see also Section 6.3 and Fig. 13).","Citation Text":["Hamana et al.","2022b"],"Functions Text":["Comparison of 1D marginal posterior distributions over the parameters S8 \u2261 \u03c38(\u03a9m\/0.3)0.5, \u03c38 and \u03a9m, from DES Y3 data as well as other experiments, and consistency tests for this work (in blue).","Constraints from other weak lensing surveys, namely HSC Y1","are shown in grey"],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[532,545],[552,557]],"Functions Start End":[[0,194],[452,510],[654,671]]}
{"Identifier":"2018AandA...614A..56B__Chung_et_al._2007_Instance_1","Paragraph":"The studies of emission lines, however, which for wide-field cameras require specific and expensive narrow-band (NB) filters of large physical size, have thus far been limited to pointed observations. Very deep H\u03b1 observations of a few galaxies in nearby clusters, including our recent observations with MegaCam, have led to several intriguing discoveries. They have shown that the ionised phase appears to be an ideal tracer of stripped gas in dense regions: approximately 50% of late-type galaxies show extended (~50 kpc) tails of ionised gas with surface brightness \u03a3(H\u03b1) of approximately a few 10\u221218 erg s\u22121 cm\u22122 arcsec\u22122 (Boselli & Gavazzi 2014), while only a handful of galaxies have extended cold or hot gaseous tails (Chung et al. 2007; Sun et al. 2006, 2007, 2010; Scott et al. 2012; Sivanandam et al. 2014; J\u00e1chym et al. 2014). In some objects, the cometary shape of the tails indicates that the gas has been stripped by the interaction with the hot ICM (Gavazzi et al. 2001; Yoshida et al. 2002; Yagi et al. 2010; Fossati et al. 2012, 2016, 2018 \u2013 paper II; Zhang et al. 2013; Boselli et al. 2016a); in other systems, bridges of ionised gas linking different nearby galaxies are associated with tidal tails in the stellar component, suggesting gravitational perturbations with nearby companions or within infalling groups (i.e. pre-processing; Kenney et al. 2008; Sakai et al. 2002; Gavazzi et al. 2003a; Cortese et al. 2006). They have also shown that within the tails of stripped gas, star formation in compact HII regions occurs in some but not in all objects (Gavazzi et al. 2001; Yoshida et al. 2008; Hester et al. 2010; Fumagalli et al. 2011a; Fossati et al. 2012; Boissier et al. 2012; Yagi et al. 2013; Kenney et al. 2014; Boselli et al. 2016a, 2018 \u2013 paper III). The removal of the gas affects the activity of star formation of galaxies on different timescales that depend on the perturbing mechanism (Larson et al. 1980; Boselli et al. 2006, 2016b; Bekki 2009, 2014; McGee et al. 2009; Cen 2014; Fillingham et al. 2015; Rafieferantsoa et al. 2015). The distribution and the morphology of the star-forming regions within galaxies is also tightly connected to the perturbing mechanisms (increases in the nuclear star formation activity and asymmetric distributions of star-forming regions are typical in gravitational interactions, radially truncated star-forming discs in interactions with the ICM, and fainter star formingdiscs in starvation, Kennicutt & Keel 1984; Barton et al. 2000; Boselli et al. 2006; Ellison et al. 2008; Woods et al. 2010; Scudder et al. 2012; Patton et al. 2011, 2013). All these pieces of evidence underline the power of NB H\u03b1 imaging data in identifying the dominant perturbing mechanism in dense environments.","Citation Text":["Chung et al. 2007"],"Functions Text":["They have shown that the ionised phase appears to be an ideal tracer of stripped gas in dense regions: approximately 50% of late-type galaxies show extended (~50 kpc) tails of ionised gas with surface brightness \u03a3(H\u03b1) of approximately a few 10\u221218 erg s\u22121 cm\u22122 arcsec\u22122","while only a handful of galaxies have extended cold or hot gaseous tails"],"Functions Label":["Background","Background"],"Citation Start End":[[726,743]],"Functions Start End":[[357,625],[652,724]]}
{"Identifier":"2022MNRAS.513.4361M__Zdziarski,_Johnson_&_Magdziarz_1996_Instance_1","Paragraph":"The relxillDCp model combines the ionized disc reflection code xillverDCp (Garc\u00eda et al. 2013, 2016) with the convolution model relconv (Dauser et al. 2013). The relconv model determines the relativistic effects in the reflection spectrum and assumes a broken power-law emissivity profile for the illumination of the disc by an X-ray corona. The emissivity profile has the following form: $\\epsilon (r)\\propto r^{-q_{{\\rm in}}}$ for rin \u2264 r \u2264 rbr, and $\\epsilon (r)\\propto r^{-q_{{\\rm out}}}$ for $r_{\\rm br}\\le r\\le r_{{\\rm {\\rm out}}}$, where qin and qout represent inner and outer emissivity indices, respectively, rbr is the break radius, rin and rout are the inner and outer disc radii, respectively. The other parameters that characterize the disc reflection model are: black hole spin (a), the inclination angle (\u03b8\u25cb) of the disc to the observer, reflection fraction (refl_frac), iron abundance (AFe), ionization state (log\u2009\u03be), and number density (ne) of electrons in the disc atmosphere. The disc irradiation profile is considered Newtonian over the outer regions of the disc, and hence we fixed the outer emissivity index at 3. We fixed the break radius at 6rg, which is a typical AGN coronal radius (e.g. Wilkins & Fabian 2011; Mallick et al. 2021). We assume the solar abundance of iron and fixed the inner and outer radii of the accretion disc at the innermost stable circular orbit (risco) and 1000rg (rg = GMBH\/c2), respectively. We set refl_frac=\u22121 in the relxillDCp model to solely describe the disc reflection. Since relxillDCp considers the thermal Comptonization model nthComp (Zdziarski, Johnson & Magdziarz 1996; \u0179ycki, Done & Smith 1999) as the irradiating (primary) continuum, we replaced the zpowerlw continuum with nthComp. The slope of the nthComp and relxillDCp components are tied. The seed photon temperature in the nthComp model was set at the maximum possible disc temperature (see column 8 of Table 2) for each source. We fixed the electron temperature of the hot coronal plasma at a typical value of 100 keV. The relative strength of reflection was measured as a ratio of the disc reflected flux to the irradiating primary source flux in the 0.3\u201310 keV band. We find that the broad-band best-fitting spectral model is Tbabs\u00d7(relxillDCp+nthComp) for all the sources in the sample, except for J1559 and POX 52. The source, J1559, showed an absorption feature at \u223c0.7 keV, which was modeled by a Gaussian absorption line gabs. The gabs component improves the fit statistic by \u0394\u03c72 = 92.9 for 3 free parameters. The expression for the broad-band best-fit model of J1559 is Tbabs\u00d7gabs\u00d7(zgauss+relxillDCp+nthComp). To model the absorption curvature present in the 1\u20132 keV band of POX 52, we included an ionized partial covering absorption component (zxipcf; Reeves et al. 2008). The zxipcf model has three free parameters: column density (NH,wa), ionization state (log\u2009\u03bewa) of the absorbing medium and covering fraction (Cf). The inclusion of the ionized absorption component provided an improvement of \u0394\u03c72 = 237.5 for three additional free parameters. The maximum-likelihood ratio (MLR) test suggests that the zxipcf model component is \u226599.99 per cent significant. The best-fitting model for POX 52 has the following expression: Tbabs\u00d7zxipcf\u00d7(relxillDCp+nthComp). We find that the relativistic reflection from an ionized, higher density accretion disc can self-consistently explain the soft X-ray excess emission for the low-mass AGN sample. The best-fitting unfolded spectral models, model components, and data-to-model ratio plots are shown in Fig. 1. We show the broad-band photon count spectra, best-fitting models, and residuals plots in Fig. A2.","Citation Text":["Zdziarski, Johnson & Magdziarz 1996"],"Functions Text":["Since relxillDCp considers the thermal Comptonization model nthComp","as the irradiating (primary) continuum, we replaced the zpowerlw continuum with nthComp."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1596,1631]],"Functions Start End":[[1527,1594],[1659,1747]]}
{"Identifier":"2016ApJ...827..107P__Padoan_&_Nordlund_1999_Instance_1","Paragraph":"To estimate the total integrated intensity of a shock-excited molecular line coming from a GMC, the shock models are scaled up so that the total energy dissipated in shocks is equal to the expected turbulent energy dissipation rate of the molecular cloud, as done in Basu & Murali (2001) and Pon et al. (2012). The dissipation rate of the turbulent energy of a molecular cloud, Lturb, is\n2\n\n\n\n\n\nwhere \u03c1 is the gas density, \u03c3 is the one-dimensional velocity dispersion, R is the radius of the (spherical) cloud, and \u03ba is the ratio of the dissipation timescale to the flow crossing timescale of the cloud. For this paper, \u03ba is taken to be equal to 1, in agreement with numerical simulations of decaying turbulence (Gammie & Ostriker 1996; Mac Low et al. 1998; Stone et al. 1998; Mac Low 1999; Padoan & Nordlund 1999; Ostriker et al. 2001). This value of \u03ba, however, is relatively uncertain, such that the predicted integrated intensities should only be considered to be accurate to a factor of a few. The integrated intensity of any shock-powered line, assuming a mean mass per hydrogen nuclei of 2.77 amu, is thus\n3\n\n\n\n\n\nwhere  is the fraction of the shock energy being emitted in the line. While the total turbulent energy of the cloud depends on the cube of the radius, the conversion of a luminosity to an intensity introduces an r\u22122 dependence, and setting the dissipation timescale to be equal to the turbulent crossing time introduces a further r\u22121 dependence, such that this predicted integrated intensity is independent of the size of the cloud. Pon et al. (2012) show that if the velocity distribution of gas particles in a molecular cloud is Gaussian and isotropic, then the shock velocity at which the peak energy dissipation occurs is approximately 3.2 times larger than the one-dimensional velocity dispersion of the gas, since the energy dissipation rate of a shock scales with the third power of the shock speed, but the probability of gas particles having a particular velocity difference decreases with increasing velocity. We assume that all of the turbulent energy in a cloud is dissipated at a shock speed of 3 km s\u22121, which would be the peak energy dissipation velocity for a velocity dispersion of roughly 1 km s\u22121 that would lead to observed FWHMs of 2.3 km s\u22121. While this velocity dispersion is on the lower end for what is usually associated with IRDCs (e.g., Paper II), slightly larger velocity dispersions should create larger peak temperatures and larger energy dissipation rates, such that our shock models can be considered to be conservative estimates for the shock emission. See Pon et al. (2012) for a more detailed discussion about this method of scaling the shock models.","Citation Text":["Padoan & Nordlund 1999"],"Functions Text":["For this paper, \u03ba is taken to be equal to 1, in agreement with numerical simulations of decaying turbulence"],"Functions Label":["Uses"],"Citation Start End":[[791,813]],"Functions Start End":[[604,711]]}
{"Identifier":"2019MNRAS.487.1626Q__Fender_2006_Instance_1","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"],"Functions Text":["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","and \u0393jet \u2272 1.2"],"Functions Label":["Background","Background"],"Citation Start End":[[801,812]],"Functions Start End":[[484,763],[785,799]]}
{"Identifier":"2016ApJ...817..117S__Brisken_&_Zirin_1997_Instance_1","Paragraph":"Running penumbral waves in velocity and intensity observations were first reported by Giovanelli (1972) and Zirin & Stein (1972). Later, they were found in the photosphere as well (Musman et al. 1976), but there they appear to be more intermittent and to have higher radial phase velocity (40\u201390 km s\u22121) than the waves in H\u03b1. Whereas the velocity amplitudes are less in the photosphere than in the chromosphere, the density is very low there and most of the wave energy lies in the photosphere and subphotosphere. Larger amplitudes on the disk-side penumbra demonstrate an alignment of the oscillations along the magnetic field. Running waves are also detected in the umbra, but the waves were believed to be unrelated to those in the penumbra (Kobanov & Makarchik 2004). In the chromosphere, the frequency of travelling waves decreases as they propagate from the umbra into the outer penumbra (e.g., Lites 1988). A similar effect is also found in measurements of the propagation velocity of travelling waves (Brisken & Zirin 1997; Sigwarth & Mattig 1997; Alissandrakis et al. 1998; Kobanov & Makarchik 2004; Tziotziou et al. 2006, 2007). Generally, the waves decelerate from 40 km s\u22121 near the inner part of the penumbra to 10 km s\u22121 or less near the outer edge of the penumbra. More recently, from a multi-wavelength study including the coronal channels of the Atmospheric Imaging Assembly (AIA) on board the Solar Dynamics Observatory (SDO), Jess et al. (2015) revealed the presence of a wide range of frequencies, with longer periodicities preferentially occurring at increasing distance from the umbra. The phase speeds also tend to decrease with increasing periodicity as the waves propagate away from the umbral barycenter. These observations also suggest that these slow waves are driven by a regular coherent source. The physical nature of running penumbral waves has been controversial. Some researchers have regarded them as trans-sunspot waves originating from umbral oscillations since they detected waves starting from the umbra and propagating through the penumbra (e.g., Alissandrakis et al. 1992; Tsiropoula et al. 1996, 2000). However, others suggest that the trans-sunspot (i.e., outward) motion is apparent to a given line of sight, and that these oscillations actually represent the upward propagation of field-guided magnetoacoustic waves from the photosphere (e.g., Christopoulou et al. 2000, 2001; Georgakilas et al. 2000; Rouppe van der Voort et al. 2003; Bogdan & Judge 2006; Kobanov et al. 2006; Bloomfield et al. 2007; Jess et al. 2013, 2015). The gradual change in the inclination of the penumbral field lines is responsible for changes in the oscillation periods and phase speeds.","Citation Text":["Brisken & Zirin 1997"],"Functions Text":["A similar effect is also found in measurements of the propagation velocity of travelling waves"],"Functions Label":["Background"],"Citation Start End":[[1010,1030]],"Functions Start End":[[914,1008]]}
{"Identifier":"2017ApJ...837..127G__Kasliwal_et_al._2015_Instance_1","Paragraph":"Even though the physical processes giving rise to the flaring emission of blazars remain debatable, considerable progress has been made in characterizing the statistical properties of blazar variability at different wavelengths, and in different time domains. In particular, it has been demonstrated repeatedly that the power spectral density (PSD) of blazar light curves is, in general, of a power-law form (Simonetti et al. 1985; Kataoka et al. 2001; Brinkmann et al. 2003; Papadakis et al. 2003; Aharonian et al. 2007; Ciprini et al. 2007; Chatterjee et al. 2008, 2012; Abdo et al. 2010a; Carini et al. 2011; Kastendieck et al. 2011; Edelson et al. 2013; Nakagawa & Mori 2013; Max-Moerbeck et al. 2014; Park & Trippe 2014; Revalski et al. 2014; Sobolewska et al. 2014; Aleksi\u0107 et al. 2015; Isobe et al. 2015; Kasliwal et al. 2015). A physical process with such a variability power spectrum, denoted hereafter as \n\n\n\n\n\n, where \u03bdk is the temporal frequency (corresponding to the timescale 1\/\u03bdk), A is the normalization constant, and \u03b2 is the spectral slope, is called white noise when \u03b2 = 0, flicker (pink) noise when \u03b2 = 1, and Brownian (red) noise when \u03b2 = 2 (Press 1978). The PSD integrated over some variability frequency range is then a measure of the variance of the underlying signal in the time series within the corresponding range of variability timescales. Breaks in the slope or in the normalization of a PSD may appear, signaling characteristic\/critical variability timescales in the system. In the case of blazars, various segments of radio, optical, X-ray, and \u03b3-ray power spectra within the variability time domains from years to days (and in some instances, even sub-hour timescales), are characterized by spectral slopes 1 \u2264 \u03b2  3, meaning that the variability amplitude increases with increasing variability timescale. Rarely, however, have blazar PSDs been analyzed in a systematic way at different wavelengths across the electromagnetic spectrum and over a truly broad range of temporal frequencies. It is important to note that colored-noise-type power spectra are expected to flatten on longer variability timescales (to preserve the total finite variance), and to cut-off at frequencies corresponding to the shortest variability timescale in a system. The detection of such cutoffs in blazar periodograms would be of primary importance for constraining the physics of blazar jets; however, such detections may be hampered by the finite duration of available monitoring blazar data on the one hand, and statistical fluctuations resulting from the measurement errors on the other hand.","Citation Text":["Kasliwal et al. 2015"],"Functions Text":["In particular, it has been demonstrated repeatedly that the power spectral density (PSD) of blazar light curves is, in general, of a power-law form"],"Functions Label":["Background"],"Citation Start End":[[812,832]],"Functions Start End":[[260,407]]}
{"Identifier":"2021MNRAS.502..915C__Ogilvie_2014_Instance_1","Paragraph":"Under the Applegate model, the change in orbital period is directly related to the change in the companion star\u2019s gravitational quadrupole moment Q (Applegate & Patterson 1987),\n(8)$$\\begin{eqnarray*}\r\n\\frac{\\Delta P_{\\rm orb}}{P_{\\rm orb}} = -9\\frac{\\Delta Q}{M_{\\rm c} A^2},\r\n\\end{eqnarray*}$$where A = x(1 + q)\/sin\u2009i is the orbital separation. For comparison, the total quadrupole moment induced by the spin of the companion star and the tidal distortion in the pulsar\u2019s gravitational field is (Voisin, Breton & Summers 2020a)\n(9)$$\\begin{eqnarray*}\r\n\\frac{Q}{M_{\\rm c} A^2} = -\\frac{2}{9} k_2 \\left(\\frac{R_{\\rm c}}{A}\\right)^5 \\left(4 q + 1\\right),\r\n\\end{eqnarray*}$$where Rc is the radius of the companion star and k2 is the apsidal motion constant, a parameter describing the deformability of the companion star (Sterne 1939). For solar-type stars k2 \u223c 0.035 (Ogilvie 2014), while if we assume that redback companions are akin to the companions in CV systems whose outer envelopes have also been stripped through accretion, then we may expect a smaller value k2 \u223c 10\u22123 (Cisneros-Parra 1970). For J2039, the hyperparameter $h = 3.9^{+2.2}_{-1.2}$\u2009s corresponds to the typical fractional amplitude for the variations in orbital phase. Taking the simpler squared exponential covariance function of equation (4) corresponding to n \u2192 \u221e, then the deviations in orbital period have covariance function,\n(10)$$\\begin{eqnarray*}\r\nK_{\\Delta P_{\\rm orb}\/P_{\\rm orb}}(t_1,t_2) &amp;=&amp; \\frac{\\partial ^2 K}{\\partial t_1 \\partial t_2} \\nonumber\\\\\r\n&amp;=&amp; \\frac{h^2}{l^2} \\exp \\left(\\!-\\frac{(t_1 - t_2)^2}{2\\ell ^2}\\!\\right) \\left(\\!1 - \\frac{(t_1 - t_2)^2}{\\ell ^4}\\!\\right).\r\n\\end{eqnarray*}$$The typical (fractional) amplitude of the orbital period variations is therefore \u0394Porb\/Porb \u223c h\/\u2113 = (3 \u00b1 1) \u00d7 10\u22127, corresponding to $\\Delta Q \/ Q \\sim 3\\times 10^{-5} k_2^{-1}$. The time-varying component to the gravitational quadrupole moment is therefore required to be of order a few per\u2009cent of the total expected quadrupole moment at most to explain the observed orbital period variations. From this, it seems plausible that the observed period variations can be powered by quadrupole moment changes, without requiring that a large fraction of the star be involved in the process. The required fractional quadrupole moment changes are very similar to those recently calculated for the companion to the black widow PSR J2051\u20130827 by Voisin et al. (2020b), despite the large difference in their masses.","Citation Text":["Ogilvie 2014"],"Functions Text":["For solar-type stars k2 \u223c 0.035"],"Functions Label":["Uses"],"Citation Start End":[[867,879]],"Functions Start End":[[834,865]]}
{"Identifier":"2015AandA...576A..26K__Hathaway_&_Rightmire_2010_Instance_1","Paragraph":"We have seen that some of our runs show clear activity cycles. Therefore we expect to see a corresponding modulation of the flow. In Fig. 13, we show for Run A the temporal variation of the mean large-scale magnetic field (\\hbox{$\\mean B$}B) normalized by Beq, the latitudinal component of the meridional circulation u\u03b8(r, \u00b1 32\u00b0) at r \u2248 0.95 R\u2299 and r \u2248 0.73 R\u2299, the mean rotation rate \\hbox{$\\mean\\Omega(0.95~R_\\odot,\\pm32^\\circ)$}\u03a9(0.95R\u2299,\u00b132\u25e6), \\hbox{$\\overline\\Omega(r,0^\\circ)$}\u03a9(r,0\u25e6) at r = 0.73 R\u2299 and r = 0.95 R\u2299, as well as the latitudinal and radial differential rotation \\hbox{$\\Delta_\\Omega^{(r)}$}\u0394\u03a9(r) and \\hbox{$\\Delta_\\Omega^{(\\theta)}$}\u0394\u03a9(\u03b8), defined in Eq. (13). We see that the meridional circulation varies with the magnetic field, becoming weaker during maximum and stronger during minimum, the overall temporal variation being about 50% in this case. (The linear correlation coefficient between \\hbox{$\\mean B$}B and u\u03b8(0.95 R\u2299, \u00b1 32\u00b0) \u2248  \u22120.36, \u2212 0.38.) This kind of weak anti-correlation between the activity cycle and the meridional flow has been found in solar observations (Chou & Dai 2001; Hathaway & Rightmire 2010) and is believed to arise at least in part from the Lorentz force of the dynamo-generated magnetic fields (see, e.g., Rempel 2006, Karak & Choudhuri 2012, Passos et al. 2012).The meridional circulation at the bottom is also weakly correlated with the activity cycle (correlation coefficients between \\hbox{$\\mean B$}B and u\u03b8(0.73 R\u2299, \u00b1 32\u00b0) \u2248  \u22120.22, \u2212 0.47). We see that \\hbox{$\\mean\\Omega(0.95~R_\\odot,\\pm32^\\circ)$}\u03a9(0.95R\u2299,\u00b132\u25e6) (Fig. 13c) also shows a weak anti-correlation with the magnetic variations (having correlation coefficient \u2248 \u2212 0.25). The strong magnetic fields during maxima change \\hbox{$\\mean\\Omega$}\u03a9 by a few per cent (\u2248 6%). However \\hbox{$\\mean\\Omega(0.95~R_\\odot,0^\\circ)$}\u03a9(0.95R\u2299,0\u25e6) (Fig. 13d) shows positive correlation (correlation coefficient \u2248 0.36) and the overall variation is larger (\u2248 12%). Because of this variation of \\hbox{$\\mean\\Omega$}\u03a9 at the equator, the values of \\hbox{$\\Delta_\\Omega^{(r)}$}\u0394\u03a9(r) and \\hbox{$\\Delta_\\Omega^{(\\theta)}$}\u0394\u03a9(\u03b8) (Figs. 13e, f) show a positive correlation with the magnetic field (correlation coefficients 0.36, 0.21) with the overall variation being ~ 75% and 166%, respectively.","Citation Text":["Hathaway & Rightmire 2010"],"Functions Text":["This kind of weak anti-correlation between the activity cycle and the meridional flow has been found in solar observations"],"Functions Label":["Similarities"],"Citation Start End":[[1118,1143]],"Functions Start End":[[977,1099]]}
{"Identifier":"2022AandA...659A.180G__Stangalini_et_al._2014_Instance_1","Paragraph":"The reasons for studying the dynamic properties of small-scale magnetic fields in the quiet photosphere are manifold. Firstly, magnetic fields provide an opportunity to probe some aspects inherent to turbulent convection (see, e.g. Abramenko et al. 2011; Lepreti et al. 2012; Giannattasio et al. 2013, 2014a,b; Giannattasio et al. 2019; Abramenko 2014, 2017; Del Moro et al. 2015; Caroli et al. 2015; Chian et al. 2019; Giannattasio & Consolini 2021) and energy propagation to the upper atmospheric layers (see, e.g. Viticchi\u00e9 et al. 2006; Jefferies et al. 2006; Tomczyk et al. 2007; De Pontieu et al. 2007; Chae & Sakurai 2008; Stangalini et al. 2014, 2015, 2017; Stangalini 2014; Rouppe van der Voort et al. 2016; Go\u0161i\u0107 et al. 2018; Bellot Rubio & Orozco Su\u00e1rez 2019; Rajaguru et al. 2019; Keys et al. 2019, 2020; Guevara G\u00f3mez et al. 2021; Jess et al. 2021) that cannot be addressed otherwise from an observational point of view. Secondly, they allow us to constrain the available energy in the quiet Sun and infer some details of the processes that amplify and organise them from subgranular to supergranular scales (see, e.g. November 1980; Roudier et al. 1998; Berrilli et al. 1999, 2002, 2004; Berrilli et al. 2005, 2013, 2014; Consolini et al. 1999; Getling & Brandt 2002; Rast 2002; Del Moro 2004; Del Moro et al. 2004; Getling 2006; Nesis et al. 2006; Centeno et al. 2007; Brandt & Getling 2008; de Wijn et al. 2008; Yelles Chaouche et al. 2011; Orozco Su\u00e1rez et al. 2012; Giannattasio et al. 2018, 2020; Requerey et al. 2018). Thirdly, their study provides important constraints for theoretical models and\/or models implemented in simulations of the photospheric layer (see, e.g. Stein & Nordlund 1998, 2001; Cattaneo et al. 2003; V\u00f6gler et al. 2005; Rempel et al. 2009; Shelyag et al. 2011; Beeck et al. 2012; Rempel 2014; Danilovic et al. 2015; Khomenko et al. 2017). All these aspects coalesce, advancing our knowledge of the processes capable of accumulating, transporting, and eventually dissipating the enormous amount of energy available in the photosphere of the quiet Sun.","Citation Text":["Stangalini et al. 2014"],"Functions Text":["The reasons for studying the dynamic properties of small-scale magnetic fields in the quiet photosphere are manifold. Firstly, magnetic fields provide an opportunity to probe","and energy propagation to the upper atmospheric layers (see, e.g.","that cannot be addressed otherwise from an observational point of view."],"Functions Label":["Motivation","Motivation","Motivation"],"Citation Start End":[[629,651]],"Functions Start End":[[0,174],[451,516],[861,932]]}
{"Identifier":"2021MNRAS.504.5840F__Eriksen_et_al._2007_Instance_1","Paragraph":"The standard cosmological model stands on the shoulders of a fundamental assumption: that the universe is statistically homogeneous and isotropic on the largest scales. This assumption has been thoroughly tested over the last years both with cosmic microwave background (CMB) and Large-scale structure data. In particular, the analysis of CMB data, most notably from the Wilkinson Microwave Anisotropy Probe (WMAP; Bennett et al. 2013) and Planck (Planck Collaboration I 2020) experiments, has not yet provided conclusive evidence for the hypothesis of Cosmological Isotropy (Eriksen et al. 2004, 2007; Hajian, Souradeep & Cornish 2005; Land & Magueijo 2007; Hansen et al. 2009; Samal et al. 2009; see also Planck Collaboration VII 2020 and references therein). Moreover, Galactic foreground contamination or known systematic effects in the data alone can not explain the observed CMB \u2018anomalies\u2019, i.e. large-scale deviations from the concordance Lambda cold dark matter (\u039bCDM) model (see e.g. Rassat et al. 2014; see Planck Collaboration VII 2020 for a recent overview). Power asymmetry from CMB data has also been a matter of intense debate and scrutiny (Gazta\u00f1aga, Fosalba & Elizalde 1998; Eriksen et al. 2007; Lew 2008; Hoftuft et al. 2009; Paci et al. 2010; Axelsson et al. 2013; Shaikh et al. 2019, see also Dai et al. 2013 for a comprehensive discussion and references therein), and evidence has been reported that this could source deviations from isotropy on cosmological scales (Hansen et al. 2009). However, a more recent analysis based on Planck data finds no evidence for such power asymmetry when all scales are taken into account (Quartin & Notari 2015). This is in qualitative agreement with the latest results from the Planck Collaboration analysis (Planck Collaboration VII 2020) where they conclude that the observed power asymmetry is not robust to foreground contamination or systematic residuals. It is important to note that previous analysis have concentrated on quantifying potential deviations from statistical isotropy using a statistical prior. First analyses using WMAPdata looked for the direction of maximal asymmetry in the sky, thus quantifying anisotropy for a given preferred direction (Hansen et al. 2009). In turn, this led to proposing a particular angular distribution of power in the sky to simply capture the observed anisotropy, such as the so-called \u2018dipole anisotropy\u2019 modulation (Prunet et al. 2005; Gordon 2007). This same model has been further constrained with Planck data (Planck Collaboration XXIII 2014; Planck Collaboration XVI 2016; Planck Collaboration VII 2020; Aiola et al. 2015; Mukherjee et al. 2016). Alternatively, a recent analysis (Ho & Chiang 2018) focuses on quantifying possible CMB peak shifts across the sky, finding significant variations, but they attribute this behaviour to possible systematic effects or the solar dipole. Complementary evidence for cosmological anisotropy has been investigated using probes of the low-redshift universe (see Colin et al. 2011; Secrest et al. 2021 and references therein).","Citation Text":["Eriksen et al.","2007"],"Functions Text":["In particular, the analysis of CMB data, most notably from the Wilkinson Microwave Anisotropy Probe","and Planck","experiments, has not yet provided conclusive evidence for the hypothesis of Cosmological Isotropy"],"Functions Label":["Background","Background","Background"],"Citation Start End":[[576,590],[597,601]],"Functions Start End":[[308,407],[436,446],[477,574]]}
{"Identifier":"2018MNRAS.478.3890B__Heckman_et_al._2017_Instance_1","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"],"Functions Text":["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","giving the neutral material time to propagate out to the impact parameters probed by COS-AGN"],"Functions Label":["Background","Background"],"Citation Start End":[[514,533]],"Functions Start End":[[148,331],[420,512]]}
{"Identifier":"2019ApJ...871...58T__Charbonnel_&_Lagarde_2010_Instance_2","Paragraph":"We derived stellar parameters, [C\/M], and [N\/M] using SLAM. To avoid bad fits at the edges of the parameter space, we exclude stars with spectral S\/N in the g band less than 50, and metallicity less than \u22121.4. The derived C and N abundances are shown in Figure 8. Clearly, in the top panel, the CH-strong, CH-normal, and metal-poor field stars are separated, and their relative distribution in the N\u2013C parameter space is similar to the case of APOGEE abundances (left panel of Figure 7): (1) metal-poor field stars form a sequence in the lower left of the top panel. As evolved stars ascend the RGB, C and N abundances may be changed by first dredge-up (Iben 1964, 1967) and extra mixing (Gratton et al. 2000; Charbonnel & Lagarde 2010). Given that for a typical halo\/thick-disk star of 1 M\u2299, the first dredge-up occurs around Teff = 5200 K (Boothroyd & Sackmann 1999), and most of our sample stars have \n\n\n\n\n\n K and log g  2.5, we infer that most stars have already undergone first dredge-up. On the other hand, the C and N abundances of these stars could be altered by extra mixing. Stars with brighter K-band absolute magnitudes tend to have higher [N\/Fe] and lower [C\/Fe] (middle and bottom panels of Figure 8), which is consistent with extra-mixing theory and observation (Gratton et al. 2000; Charbonnel & Lagarde 2010); (2) CH-normal stars show an enhanced median N abundance and slightly depleted median C abundance. Clearly, the median N abundance of CH-normal stars is enhanced compared to that of normal metal-poor field stars with similar C abundances. In other words, the enhanced N abundances in CH-normal stars cannot be explained by the extra-mixing effect alone. We notice that a few CH-normal stars may have low N abundances, probably due to large uncertainties when spectra of a particular type are scarce in the training set, i.e., high-N metal-poor stars. The statistical similarity between APOGEE C and N abundances and LAMOST-derived C and N abundances further strengthens our statement above.","Citation Text":["Charbonnel & Lagarde 2010"],"Functions Text":["Stars with brighter K-band absolute magnitudes tend to have higher [N\/Fe] and lower [C\/Fe] (middle and bottom panels of Figure 8), which is consistent with extra-mixing theory and observation"],"Functions Label":["Similarities"],"Citation Start End":[[1299,1324]],"Functions Start End":[[1085,1276]]}
{"Identifier":"2016AandA...589A..44G__within_2000_Instance_2","Paragraph":"W51e2 is the strongest and best-studied HC HII region in the W51 Main cluster, and it is believed to be powered by an O8-type young star (e.g., Shi et al. 2010a). A number of interferometric studies conducted with varying angular resolutions, at centimetre (cm) and (sub)millimetre (mm) bands, identified molecular and ionized gas undergoing infall and rotation toward W51e2. VLA observations of the NH3 inversion lines (1,\u20091) and (2,\u20092) seen in absorption (1\\hbox{$\\farcs$}.\u030b1 beamsize) revealed radial infall on scales larger than 5000 AU toward the W51e2 core (Zhang & Ho 1997). Higher angular resolution observations of the (3,\u20093) NH3\u2009absorption line (0\\hbox{$\\farcs$}.\u030b3 beamsize) showed signatures of rotation within 2000 AU based on a position-velocity (pv) diagram (Zhang & Ho 1997). Zhang et al. (1998) identified a velocity gradient in a CH3CN transition at 2 mm, deriving a position angle (PA) of 20 \u00b1 20\u00b0. Keto & Klaassen (2008) imaged the H53\u03b1 radio recombination line (RL) with the VLA (0\\hbox{$\\farcs$}.\u030b45 beamsize) and they claimed rotation in the ionized gas along the axis of a molecular bipolar outflow (oriented NW-SE) imaged with the SMA in the CO (2\u22121) line (1\u2032\u2032 beamsize), suggesting a simple inflow\/outflow picture in a single high-mass young stellar object (YSO). However, higher resolution observations, using the SMA at the wavelengths of 0.85 mm (0\\hbox{$\\farcs$}.\u030b3 beamsize) and 1.3 mm (0\\hbox{$\\farcs$}.\u030b7 beamsize), revealed a more complex picture, by resolving W51e2 into three subcores (Shi et al. 2010a): W51e2-W, corresponding to the HC HII region, W51e2-E, located about 1\u2032\u2032\u2009east of the HC HII region and corresponding to the brightest dust continuum source, and W51e2-NW, the weakest continuum component, located about 1\u2032\u2032\u2009 NW of the HC HII region. Shi et al. (2010b) imaged the CO (3\u22122) line (with a 0\\hbox{$\\farcs$}.\u030b7 beamsize) and established that the driving source of the powerful molecular outflow in this region is the protostellar core W51e2-E, and not the HC HII region W51e2-W, challenging the scenario proposed by Keto & Klaassen (2008). Etoka et al. (2012) used MERLIN to image the Class II 6.7 GHz CH3OH\u2009masers (typical signpost of HMSF), and found that the bulk of maser emission is indeed concentrated toward W51e2-E, and not the HC HII region W51e2-W. This further supports the scenario proposed by Shi et al. (2010a), where the ongoing star formation activity in the region is not concentrated on the HC HII region but toward its companion 1\u2032\u2032 to the east. ","Citation Text":["Zhang & Ho 1997"],"Functions Text":["Higher angular resolution observations of the (3,\u20093) NH3\u2009absorption line (0\\hbox{$\\farcs$}.\u030b3 beamsize) showed signatures of rotation within 2000 AU based on a position-velocity (pv) diagram"],"Functions Label":["Background"],"Citation Start End":[[774,789]],"Functions Start End":[[582,772]]}
{"Identifier":"2022AandA...658A..78S__Rosen_&_Krumholz_2020_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":["Rosen & Krumholz 2020"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[749,770]],"Functions Start End":[[543,722]]}
{"Identifier":"2020AandA...644L...7G__Magdis_et_al._2020_Instance_3","Paragraph":"As in G18, we compiled existing constraints on the molecular gas fraction fgas of quiescent and pSB galaxies from recent literature, namely: local QGs consisting of the ATLAS3D (Young et al. 2011; Cappellari et al. 2013; Davis et al. 2014) and HRS (Boselli et al. 2014; Lianou et al. 2016) ETG samples as well as the samples of pSB galaxies (hereafter, the \u201clow-z pSB\u201d sample) of French et al. (2015) and Alatalo et al. (2016); at low and intermediate redshift, the ETG sample of Spilker et al. (2018) and the pSB sample of Suess et al. (2017); at intermediate and high redshift, constraints from Hayashi et al. (2018) on gas in z\u2004\u223c\u20041.46 cluster ETGs, as well as on individual galaxies from Sargent et al. (2015), Bezanson et al. (2019), and Rudnick et al. (2017). Given its size, we divided the ATLAS3D sample into high- and low-mass subsamples, choosing 5\u2005\u00d7\u20051010 M\u2299 as the cut-off mass. In addition, we also included fgas estimates derived from the (median) stacked FIR spectral energy distributions of ETGs at z\u2004\u223c\u20041.8 (G18; 977 galaxies), z\u2004\u223c\u20041.2, z\u2004\u223c\u20040.8, and z\u2004\u223c\u20040.5 (1394, 1536, and 563 galaxies, respectively; Magdis et al. 2020, hereafter M20). Finally, at higher redshift (z\u2004\u223c\u20043), we converted star formation rates (SFR) estimated from spectroscopy (Schreiber et al. 2018a; D\u2019Eugenio et al. 2020) into gas masses assuming the star formation efficiency found by G18. As a consequence of our zmax\u2004=\u20043.5, we did not include higher-redshift quiescent galaxies (Glazebrook et al. 2017; Schreiber et al. 2018b; Tanaka et al. 2019; Valentino et al. 2020) in the analysis and considered z\u2004\u223c\u20043 galaxies as pSB. The dust-based estimates of G18 and M20 (and, by extension, the z\u2004\u223c\u20043 semi-constraints) assume a gas-to-dust ratio (G\/D). It is dependent on metallicity, which is presumed to be solar or higher owing to both the relatively high gas-phase metallicity of MS galaxies at z\u2004\u2272\u20041 (e.g., Mannucci et al. 2010) and the already high stellar metallicities of QGs at z\u2004> \u20041 (Onodera et al. 2015; Estrada-Carpenter et al. 2019). Here we adopted an intermediate value between the solar and supersolar G\/Ds used in M20, and we increased the error bars of these points to include both the solar and supersolar confidence estimates. These various samples, which are summarized with their selection criteria in Table B.1, combine into a nonhomogeneous dataset: some were specifically selected as ETGs, and others were based on varying degrees of quiescence. In particular, pSB galaxies are not necessarily truly quiescent and could, in principle, resume normal star formation. However, as a possible precursor of QGs, they provide useful, though not constraining (see Sect. 4), comparison samples for the model. Here we refer to all equally as either QGs or pSB galaxies, and we make the assumption that, on average, these different samples are not otherwise significantly biased with regard to their gas content compared to the full population, given each mass limit and type.","Citation Text":["M20"],"Functions Text":["Here we adopted an intermediate value between the solar and supersolar G\/Ds used in","and we increased the error bars of these points to include both the solar and supersolar confidence estimates."],"Functions Label":["Uses","Uses"],"Citation Start End":[[2112,2115]],"Functions Start End":[[2028,2111],[2117,2227]]}
{"Identifier":"2018ApJ...853L..34B__K\u00f3sp\u00e1l_et_al._2014_Instance_1","Paragraph":"With the identification of the \u223c6.6 day stellar rotation period, the question remains, what is origin of the \u223c9.0 day signal in the periodogram? Because of K2\u2019s large pixel size, we considered the possibility of contamination from additional sources in the field and found that images from DSS, Galex, 2MASS, WISE, and PanSTARRS show that there are no objects of comparable brightness within 1\u2032 of CI Tau. We also examined the potential for multiplicity in the system. High-resolution observations reaching 5\u03c3 contrast at 025 separation (Uyama et al. 2017) provide no evidence for a companion down to \u0394H = 6.8 that could contribute to the photometric signal. However, there is evidence in support of a substellar body. Johns-Krull et al. (2016) reported the detection of a planet orbiting CI Tau using data from an extensive optical and infrared RV survey. The planet mass they derived is Mp = 11.29 \u00b1 2.16 MJup and the orbital period is Porb = 8.9891 \u00b1 0.0202 days, consistent with the \u223c9.0 day period shown in both of the Lomb\u2013Scargle periodograms in Figure 1. Our current understanding of the CI Tau system is that the planet does not transit. This is supported by evidence that the disk is inclined i \u223c 45\u00b0 (Guilloteau et al. 2014), though it is possible that a planet could be in an orbit misaligned with disk mid-plane (e.g., Kepler-63; Sanchis-Ojeda et al. 2013). However, the \u223c9.0 day periodic signal may be the result of a planet\u2013disk interaction because the presence of a massive planet in an actively accreting disk should show both spectroscopic and photometric variability. Indeed, Johns-Krull et al. (2016) find evidence in the H\u03b1 profile variations of CI Tau that the planet may be modulating the accretion of disk material onto the star. Although hotspots located at the foot of accretion streams can affect RV measurements mimicking the signal of an orbiting body (K\u00f3sp\u00e1l et al. 2014; Sicilia-Aguilar et al. 2015), Johns-Krull et al. (2016) specifically looked for these signals and found no evidence that hotspots produced the RV signals seen in photospheric absorption lines. Given the stellar mass of 0.80 M\u2299 (Guilloteau et al. 2014), Johns-Krull et al. (2016) determined a semimajor axis of 0.079 au for the planetary orbit. This would place the planet inside the inner edge of the disk at 0.12 au (McClure et al. 2013). It is probable that an 11.3 MJup planet so close to the inner edge of the disk would stimulate the accretion of material and possibly modify accretion onto the star, creating non-axisymmetric accretion flows (Tofflemire et al. 2017a, 2017b). As the planet orbits the star, this interaction could produce a periodic variation in the H\u03b1 emission, a tracer of accretion. The impact on stellar accretion could produce photometric variability on the \u223c9.0 day period of the planet\u2019s orbit, which K2 data show as a periodic component in the system brightness. Additionally, the strength of this periodic component can be expected to fluctuate because of the sporadic nature of accretion on short timescales (Herbst et al. 2007). This is observed in the analysis of the first and second halves of the data and is in contrast to the relatively consistent signal strength of the \u223c6.6 day periodic component in brightness, which is caused by cold starspots that are long-lived in young stars (e.g., Stelzer et al. 2003) and produce relatively consistent fluctuations in brightness as the star rotates (Herbst et al. 2007; Bradshaw & Hartigan 2014). If the 9.0 day periodic signal is the result of stellar rotation, it would call into question the legitimacy of the planet. On the other hand, the similar RV amplitudes seen in the optical and the IR and the null results on tests for RV variations produced by an accretion hotspot (Johns-Krull et al. 2016) attest to the significance of planet\u2019s detection. In addition, we would then need to explain the source of the 6.6 day signal, which is the most persistent signal observed in Figures 1 and 2.","Citation Text":["K\u00f3sp\u00e1l et al. 2014"],"Functions Text":["Although hotspots located at the foot of accretion streams can affect RV measurements mimicking the signal of an orbiting body","Johns-Krull et al. (2016) specifically looked for these signals and found no evidence that hotspots produced the RV signals seen in photospheric absorption lines."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1882,1900]],"Functions Start End":[[1754,1880],[1932,2094]]}
{"Identifier":"2020MNRAS.499.4394M__Bate_et_al._2014_Instance_2","Paragraph":"For these four remaining FHSC candidates (L1451-mm, MC35-mm, SM1N, and Oph A-N6) that have been observed at intermediate scales (few 100 au to few 1000 au) a final confirmation of their true evolutionary state requires higher resolution observations. For L1451-mm, the compact outflow needs to be resolved to investigate its morphology and kinematics, as a higher velocity component (an indication of protostellar nature) could be revealed by observations with a beam smaller than 100 au, similar to the case of B1b-N (Hirano 2019). An additional goal of high-resolution observations for L1451-mm and the remaining youngest candidates should be to investigate the temperature and density profiles of the envelope at scales from few au to 100 au. This is because simulations show that the temperature remains lower than \u223c30 K even at several tens of au up to 100 au from the centre (Bate et al. 2014; Tomida et al. 2015; Hincelin et al. 2016; Young et al. 2019) during the FHSC stage. On the other hand, Class 0 sources show temperatures of 20\u201330 K or higher at scales of several 100 au (sufficient for thermal evaporation of CO) that results in the inner envelope and disc being easily detected using C18O observations (Yen et al. 2015, 2017; Stephens et al. 2018). This holds even in very low luminosity objects, for which the unexpected large extent of C18O is interpreted as evidence of a previous burst of accretion (Frimann et al. 2017; Hsieh et al. 2018). As for the density profile, simulations of the FHSC stage show a flat inner region, corresponding to the FHSC structure and extending up to \u223c10 au (Tomida et al. 2013; Bate et al. 2014). For a protostar, on the other hand, the density profile should increase towards the central \u223c1 au region (Young et al. 2019). Observations of the continuum emission with a resolution better than a few tens of au are likely required to model the emission and provide a density and temperature profile that can probe the relevant scales. Additional line observations with a similar resolution can also help to further distinguish between the different models. We note that, as pointed out in Young et al. (2019), distinguishing a dense core with only an FHSC and one that has recently formed a protostar but in which the FHSC structure is still present is likely not possible, even with high-resolution observations. Given the optically thick nature of the FHSC core, it is difficult to probe the physical properties within the FHSC structure. Despite this, finding a source with density and temperature profiles as well as with outflow properties consistent with the theoretical predictions will provide convincing evidence in support of a bona fide FHSC.","Citation Text":["Bate et al. 2014"],"Functions Text":["As for the density profile, simulations of the FHSC stage show a flat inner region, corresponding to the FHSC structure and extending up to \u223c10 au"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1630,1646]],"Functions Start End":[[1462,1608]]}
{"Identifier":"2017MNRAS.466.3961S__Grcevich_&_Putman_2009_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":["Grcevich & Putman 2009"],"Functions Text":["Some of the recent observations and studies","indicate that the CGM can account for a good fraction of the baryons in these galaxies."],"Functions Label":["Background","Background"],"Citation Start End":[[836,858]],"Functions Start End":[[791,834],[1001,1088]]}
{"Identifier":"2021MNRAS.508.4429C__Hobbs_et_al._2010_Instance_1","Paragraph":"We explore another mechanism that can possibly result in torque reversals in the neutron star on long time-scales. It has been known for a long time that radio pulsars (which are usually isolated systems) show unexplained stochastic deviations in their spin-down behaviour (known as \u2018timing noise\u2019) on varied time-scales of a few hundred days to a few tens of years. This manifestation is akin to the random-walk behaviour in spin frequency observed in wind-fed accretion powered pulsars like Vela X-1. In an interesting study of timing irregularities of a sample of 366 pulsars, Hobbs, Lyne & Kramer (2010) found some radio pulsars showing quasiperiodic structures in their long-term timing residuals. From power spectrum analysis, significant periodicities ranging from about 1.4 to 10 yr were found in PSR B1540\u221206, PSR B1642\u221203, PSR B1818\u221204, PSR B1826\u221217, PSR B1828\u221211, and PSR B2148+63 (Hobbs et al. 2010). Interestingly, we also find quasiperiodic variations in the long-term spin evolution of Vela X-1 on time-scales of about 5.9 yr that is comparable to those inferred for radio pulsars showing quasiperiodic changes in their timing residuals. The underlying phenomena causing quasiperiodic structures in timing noise of radio pulsars is elusive. However, it has been suggested that these changes are driven by changes in the magnetosphere of the neutron star (Lyne et al. 2010). In this \u2018state-switching model\u2019, the magnetosphere of the neutron star is suggested to harbour two or more magnetospheric states which can be stable on time-scales of years but the pulsar can switch abruptly between these states driven by changes in the parameters regulating the spin-down (Lyne et al. 2010). This can possibly happen in Vela X-1 where the coupling between the dynamic magnetosphere and the neutron star can change in a quasiperiodic fashion. Interestingly, the disc-fed pulsar LMC X-4 has been found showing a near cyclic spin period evolution on time-scales of about 6.8 yr (Molkov et al. 2016) which is within a factor of 1.2 of the inferred time-scale in wind-fed pulsar Vela X-1. Recent observations of transient pulsar V0332+53 suggests switching of coupling between the accretion disc and the neutron star magnetosphere in a disc-fed pulsar (Doroshenko et al. 2017).","Citation Text":["Hobbs et al. 2010"],"Functions Text":["From power spectrum analysis, significant periodicities ranging from about 1.4 to 10 yr were found in PSR B1540\u221206, PSR B1642\u221203, PSR B1818\u221204, PSR B1826\u221217, PSR B1828\u221211, and PSR B2148+63"],"Functions Label":["Uses"],"Citation Start End":[[893,910]],"Functions Start End":[[703,891]]}
{"Identifier":"2016AandA...592A..19C__Maraston_et_al._(2009)_Instance_2","Paragraph":"Since the star-formation histories of galaxies (ETGs included, e.g. De Lucia et al. 2006; Maraston et al. 2009) can be stochastic and include multiple bursts, we also verify the full-spectrum fitting capabilities to retrieve more complex SFHs. In particular, we take an 11 Gyr old composite stellar population with an exponentially delayed SF (\u03c4 = 0.3 Gyr) as the main SF episode (this age is compatible with the age of the Universe at z ~ 0.15, which is the median redshift of our sample, see Sect. 3). We then define more complex SFHs by combining this single CSP with a burst of SF at different ages (5, 6, 7 Gyr) and with different mass contributions (3, 5, 10 %). In all cases, we consider a solar metallicity for the main SF episode and, according to the results of Maraston et al. (2009), a subsolar metallicity (Z = 0.004) for the later one. We do not mask any spectral feature of the input spectra, we assume AV  =  0.1 mag for the two components and apply a velocity dispersion of 200 km s-1. We show the results for a S\/N of 80, which matches the typical S\/N of the SDSS median stacked spectra analyzed in the following (see Sect. 3). Useful information can be derived from the comparison between the output SFH obtained from these input simulated spectra and the one provided when the single CSP alone is taken as input SFH. Fig. 5 shows that the single CSP alone is well recovered by the full spectrum fitting. In particular, ~80% of the stellar mass is retrieved within ~1 Gyr from the SFH peak. When a burst is added to this major episode of SF, the full-spectrum fitting is able to recognize the presence of a more complex SFH, as indicated by the tail appearing at smaller ages, and the total mass percentage of the later burst is retrieved within 1 Gyr from the expected age. However, we note that the main episode of SF is spread on a time interval longer than expected, and 50% of the stellar mass is retrieved around ~1 Gyr from the SFR peak. We also find that, in this case, the mean properties of the global stellar population are well retrieved, with a percentage accuracy larger than 10% starting from S\/N ~ 15 for age, ~7 for metallicity, ~20 for AV, ~8 for \u03c3 and that the metallicities of the two SF episodes are separately recovered. These S\/Ns are well below those typical of the stacked spectra analyzed in the following sections. ","Citation Text":["Maraston et al. (2009)"],"Functions Text":["In all cases, we consider a solar metallicity for the main SF episode and, according to the results of","a subsolar metallicity (Z = 0.004) for the later one."],"Functions Label":["Uses","Uses"],"Citation Start End":[[772,794]],"Functions Start End":[[669,771],[796,849]]}
{"Identifier":"2016ApJ...819...97T__Gordon_et_al._1987_Instance_1","Paragraph":"The ReTOF results from irradiated phosphine and deuterated methane ice provide crucial information regarding the mechanism of formation for methylphosphanes by analyzing the intensities of various isotopologues. Figure 8 shows the possible formation routes that would lead to each of the three observed isotopologues of methylphosphine (CH3PH2). To obtain m\/z = 50 (CHD2PH2), CD4 has to decompose via the loss of molecular hydrogen or two deuterium atoms to form carbene (CD2), which has been observed in previous irradiated ice studies (Holtom et al. 2005; Bennett & Kaiser 2007), and then insert into a phosphorus\u2013hydrogen bond of phosphine (reaction (5)). If the carbene is formed in its first excited singlet state (a1A1), the insertion is barrierless (Gordon et al. 1987). For m\/z = 51 (CD3PH2), methane and phosphine each lost a hydrogen or deuterium atom, and the resulting methyl (CD3) (Kaiser et al. 1997) and phosphino (PH2) radicals recombined barrierlessly (reaction (6)). Finally, the formation of m\/z = 52 (CD3PHD) mirrors that for CHD2PH2 but in this case phosphine lost two hydrogen atoms or molecular hydrogen to create the phosphinidene (PH) radical and then inserted into a carbon\u2013deuterium bond of methane (reaction (7)). Phosphinidene, like imidogen (NH) (Fueno et al. 1983), is expected to insert barrierlessly in its first excited singlet state (a1\u0394):\n5a\n\n\n\n\n\n\n\n5b\n\n\n\n\n\n\n\n6a\n\n\n\n\n\n\n\n6b\n\n\n\n\n\n\n\n6c\n\n\n\n\n\n\n\n7a\n\n\n\n\n\n\n\n7b\n\n\n\n\n\nTherefore, our results provide compelling evidence that methane decomposes not only to the methyl radical, but also to carbene. Likewise, phosphine was found to fragment to the phosphino radical and also to phosphinidene. The ratio of ion intensities for m\/z = 50:51:52 is 2:10:1, indicating that radical recombination was the preferred formation pathway with CD3PH2 as the most abundant isotopologue. This could either be a result of the methyl and phosphino radicals reacting quickly or that more of these radicals were produced than carbene and phosphinidene.","Citation Text":["Gordon et al. 1987"],"Functions Text":["If the carbene is formed in its first excited singlet state (a1A1), the insertion is barrierless"],"Functions Label":["Uses"],"Citation Start End":[[757,775]],"Functions Start End":[[659,755]]}
{"Identifier":"2021ApJ...911...59L__Kere\u0161_et_al._2005_Instance_1","Paragraph":"In addition to the galaxy colors (e.g., Baldry et al. 2004; Bell et al. 2004; Borch et al. 2006; Xue et al. 2010; Salim 2014; Lee et al. 2015; Wang et al. 2017), there are other indicators characterizing the star formation nature of galaxies, e.g., the morphology (Strateva et al. 2001; Barro et al. 2013, 2014; Pan et al. 2013) and some spectral features such as the Balmer absorption line (e.g., Kuntschner et al. 2002; Kim et al. 2013; Kim & Yoon 2017) and the 4000 \u212b break (e.g., Kauffmann et al. 2003; Lambas et al. 2012; Rowlands et al. 2018; Angthopo et al. 2019). Using these indicators, a number of studies have found that the color or SFR transition of galaxies is associated with the consumption of their gas content (e.g., Kere\u0161 et al. 2005; Dekel & Birnboim 2006; Kruijssen 2015; Nelson et al. 2018). If this \u201cquenching\u201d scenario is correct, the timescales of the galaxy color\/SFR transition and their gas depletion must be consistent. Some authors have studied galaxy quenching by quantifying the evolution of SMFs (or luminosity functions) of SFGs and QGs at high redshifts (e.g., Fritz et al. 2014; Rowlands et al. 2018). The others, with the help of some infrared or submillimeter surveys (e.g., Hershel and ALMA), have managed to study the cosmic evolution of the gas content as well as the gas-depletion rate of galaxies with different SFRs (e.g., Geach et al. 2011; Lagos et al. 2011; Tacconi et al. 2018; Liu et al. 2019; Castignani et al. 2020; Magnelli et al. 2020). These studies have revealed that both the galaxy quenching and gas-depletion timescales increase with decreasing redshifts. The typical quenching timescale of GV galaxies can be as long as several billion years in the local universe (Rowlands et al. 2018; Correa et al. 2019; Phillipps et al. 2019). Methods such as fitting the spectral energy distributions (SEDs) of galaxies (e.g., Belfiore et al. 2018; Phillipps et al. 2019; Zick et al. 2018) and using cosmological simulations (e.g., Feldmann et al. 2017; Nelson et al. 2018; Correa et al. 2019; Donnari et al. 2019) have been widely used to derive the star formation histories of galaxies and their lifetimes.","Citation Text":["Kere\u0161 et al. 2005"],"Functions Text":["Using these indicators, a number of studies have found that the color or SFR transition of galaxies is associated with the consumption of their gas content (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[735,752]],"Functions Start End":[[572,734]]}
{"Identifier":"2017MNRAS.469.3738S__Wargelin_etal._2004_Instance_1","Paragraph":"For the spectral fitting, the astrophysical background components are determined from a simultaneous fit to data from the RASS4 (Snowden etal. 1997). The extraction region for the RASS data is an annulus from 0.7 to 1 around the cluster centre (for NGC4636, NGC1399 and A3526 r500 is larger than 0.7, so the RASS data were extracted from 1.5 to 2). The particle background was directly subtracted from the Chandra spectra using the stow events files from an epoch close to the observation date. The stow events files are created when the ACIS detector is in a position where it is not exposed to the sky and the HRC-I camera is in the field of view. This configuration is also called event histogram mode. As shown by comparisons to dark moon observations only particle events are recorded in the stow position (Markevitch etal. 2003; Wargelin etal. 2004). For each annulus, the same detector region was used to extract the particle background spectra. These background spectra are normalized by the ratio of the (9.512)keV band count rate of the observation and the stow events file to account for variations of the quiescent particle background component. The cluster emission is modelled by an absorbed thermal model (phabs*apec), where all parameters apart from the redshift and the NH are left free to vary. Following Willingale etal. (2013), the hydrogen columns density used as a tracer for the X-ray absorption,\n\n(16)\n\r\n\\begin{eqnarray}\r\nN_\\mathrm{H\\, tot} = N_\\mathrm{H\\,\\small {I}} + 2\\cdot N_{\\mathrm{H}_2\\mathrm{m}} \\cdot \\left[ 1-\\exp \\left(-N_\\mathrm{{\\rm H\\,\\small {I}}}\\cdot \\frac{ E\\left(B-V\\right) }{ N_\\mathrm{c}} \\right) \\right]^\\alpha , \\nonumber\\\\\r\n\\end{eqnarray}\r\n\nwhere the parameters NH2m  7.2  0.3  1020cm 2, Nc3.0  0.31020cm2 and 1.1  0.1 were calibrated using X-ray afterglows of Gamma-ray bursts in the aforementioned reference. Both, the absorption E(BV) from the IRAS and COBE\/DIRBE infrared dust maps (Schlegel, Finkbeiner  Davis 1998) and the $N_\\mathrm{H\\,\\small {I}}$ from Kalberla etal. (2005) are computed at each cluster position. The combined effect of the uncertainties of these parameters, the scatter of this scaling relation (0.087) plus accounting for a 10percent uncertainty on $N_\\mathrm{H\\,\\small {I}}$ and E(BV) has only an 11percent effect on NHtot, which typically affects best-fitting temperatures by 1percent. Since this is much smaller than the typical statistical uncertainties, any statistical uncertainty of NHtot is neglected. For the relative abundance of heavy elements the Asplund etal. (2009) abundance table was used. For each observation, all spectra from the different regions and chips5 are fit simultaneously. The temperature and metallicity of spectra from the same region but different chips are linked together, while the normalizations are not because of spatial variations of the density distribution. For all observations, the steppar command was run on the temperatures. This task calculates the 2 for the parameter within a given range of values in order not to get best-fitting parameters of a local minimum of the likelihood distribution. The reduced 2 of all spectral fits was on average 1.03, while in 95percent of the cases it was below 1.17. This gives a hint that the spectral modelling is appropriate.","Citation Text":["Wargelin etal. 2004"],"Functions Text":["As shown by comparisons to dark moon observations only particle events are recorded in the stow position"],"Functions Label":["Background"],"Citation Start End":[[835,854]],"Functions Start End":[[706,810]]}
{"Identifier":"2020ApJ...894..107I__Gibb_et_al._2004_Instance_1","Paragraph":"AFGL 2136 IRS 1 (also referred to as CRL 2136, G17.64+0.16, and IRAS 18196\u22121331) is a luminous (1.0 \u00d7 105 L; Lumsden et al. 2013), high-mass (45 \u00b1 10 M; Maud et al. 2019) protostar that is believed to be in the latter stages of its evolution due to a variety of observed characteristics (Boonman & van Dishoeck 2003; Maud et al. 2018 and references therein). It is located at a distance of 2.2 kpc away from the Sun (Urquhart et al. 2014), and has been extensively observed from centimeter to micron wavelengths, at low and high angular resolution, and low and high spectral resolution. The myriad observations paint a picture where a single, isolated massive protostar is driving a wide-angle bipolar outflow through its natal cloud. The large scale outflow is observed in CO emission at millimeter wavelengths, with both the red and blue lobes being about 100\u2033 in extent (Kastner et al. 1994; Maud et al. 2018). Closer to the central source (2\u2033\u201310\u2033), the outflow cavity walls are seen in scattered light at near-infrared wavelengths (Kastner et al. 1992; Murakawa et al. 2008; Maud et al. 2018). The cool molecular envelope exhibits ice and dust absorption bands (Willner et al. 1982; Keane et al. 2001b; Dartois et al. 2002; Gibb et al. 2004), as well as molecular emission at millimeter wavelengths (van der Tak et al. 2000a, 2000b), but a much warmer component is also inferred from several different molecules seen in absorption in the near- to mid-infrared (Mitchell et al. 1990; Lahuis & van Dishoeck 2000; Keane et al. 2001a; Boonman et al. 2003; Boonman & van Dishoeck 2003; Goto et al. 2013, 2019; Indriolo et al. 2013a). The presence of a dust disk on small spatial scales was suggested by near-infrared polarization imaging (Minchin et al. 1991; Murakawa et al. 2008) and by mid-infrared interferometric observations (de Wit et al. 2011; Boley et al. 2013). A compact source was marginally resolved at centimeter wavelengths, along with a cluster of nearby 22 GHz H2O masers (Menten & van der Tak 2004), but only with the recent ALMA 1.3 mm continuum observations has the 93 \u00d7 71 mas dust disk been fully resolved (Maud et al. 2019). Thermal line emission at 232.687 GHz from the H2O \u03bd2 = 1\u20131, 55,0\u201364,3 transition has the same spatial extent as the dust emission, and the H2O gas velocities indicate Keplerian rotation within the disk (Maud et al. 2019). It is ideal that the reader has a clear picture of the AFGL 2136 region in mind to best understand the discussion throughout this paper. In particular, Figure 10 of Maud et al. (2018) provides an up-to-date schematic diagram of the AFGL 2136 region, and Figures 1 and 2 of Maud et al. (2019) present the compact disk observed in dust and gas emission, respectively.","Citation Text":["Gibb et al. 2004"],"Functions Text":["The cool molecular envelope exhibits ice and dust absorption bands"],"Functions Label":["Background"],"Citation Start End":[[1228,1244]],"Functions Start End":[[1098,1164]]}
{"Identifier":"2018MNRAS.476.1412I__Yuan,_Quataert_&_Narayan_2003_Instance_1","Paragraph":"There have been a number of analytical solutions and numerical studies for rotating flows with viscous angular momentum transport (e.g. Shakura & Sunyaev 1973; see reviews by Pringle 1981; Kato, Fukue & Mineshige 2008, references therein). Among those, we here focus on accretion flows which cannot lose internal energy via radiative cooling because of very low gas density. Such radiatively inefficient accretion flows are quite interesting since many observed BHs accrete at rates of only a small fraction of the Bondi accretion rate and their radiation luminosity is as low as \u223c10\u22121 to 10\u22129\u2009LEdd (Ho 2008, 2009). Sagittarius A* (Sgr A*) is inferred to be accreting at a rate of 10\u22123 to $10^{-2}\\ \\dot{M}_{\\rm B}$ (e.g. Yuan, Quataert & Narayan 2003; Quataert 2004), where $\\dot{M}_{\\rm B}\\simeq 10^{-5}\\ {\\rm M}_{\\odot }\\, {\\rm yr}^{-1}$ is measured from the temperature and density near the Bondi radius with X-ray observations (Baganoff et al. 2003). Because of such a low accretion rate, the bolometric luminosity of Sgr A* (M\u2022 \u2243 4 \u00d7 106\u2009M\u2299; Ghez et al. 2003) is as small as Lbol \u223c 1036\u2009erg\u2009s\u22121 \u223c 2 \u00d7 10\u22129\u2009LEdd. The second example is a BH at the centre of the giant elliptical galaxy M87. The gas accretion rate at the vicinity of the BH is estimated as \u2272 9.2 \u00d7 10\u2212 4\u2009M\u2299\u2009yr\u2212 1 (Kuo et al. 2014), which is lower than ${\\sim } 10^{-2}\\ \\dot{M}_{\\rm B}$ (Russell et al. 2015). Since the BH mass is estimated as $M_\\bullet = 6.6^{+0.4}_{-0.4}\\times 10^9\\ {\\rm M}_{\\odot }$ (Gebhardt et al. 2011) and $M_\\bullet = 3.5^{0.9}_{-0.7}\\times 10^9\\ {\\rm M}_{\\odot }$ (Walsh et al. 2013), the bolometric luminosity of Lbol \u223c 2 \u00d7 1041\u2009erg\u2009s\u22121 is \u223c3 \u00d7 10\u22127\u2009LEdd. The third example is a BH at the centre of the Andromeda Galaxy (M31). The estimated BH mass is $M_\\bullet \\simeq 1.4^{+0.7}_{-0.3}\\times 10^8\\ {\\rm M}_{\\odot }$ (Bender et al. 2005). The Bondi accretion rate and the X-ray luminosity are estimated as $\\dot{M}_{\\rm B}\\simeq 7\\times 10^{-5}\\ {\\rm M}_{\\odot }\\, {\\rm yr}^{-1}$ and LX \u2243 2 \u00d7 1036\u2009erg\u2009s\u22121 \u2243 10\u221210\u2009LEdd, respectively (Garcia et al. 2010). The corresponding bolometric luminosity is inferred as \u224310\u22129\u2009LEdd by assuming the bolometric correction factor to be \u224310 (Hopkins, Richards & Hernquist 2007).","Citation Text":["Yuan, Quataert & Narayan 2003"],"Functions Text":["Sagittarius A* (Sgr A*) is inferred to be accreting at a rate of 10\u22123 to $10^{-2}\\ \\dot{M}_{\\rm B}$ (e.g."],"Functions Label":["Uses"],"Citation Start End":[[722,751]],"Functions Start End":[[616,721]]}
{"Identifier":"2021AandA...654A.124W__Tanvir_et_al._2017_Instance_2","Paragraph":"The first multi-messenger GW event was discovered on 17 August, 2017. About 1.7 s after the GW170817 signal detected by LIGO and Virgo (Abbott et al. 2017a), the Fermi Gamma-ray Burst Monitor was successfully triggered by GRB 170817A (Abbott et al. 2017b; Goldstein et al. 2017; Zhang et al. 2018) and subsequently a large number of follow-up observations monitored the afterglow emission in different electromagnetic bands from the radio to X-rays (Alexander et al. 2017; Hallinan et al. 2017; Margutti et al. 2017; Troja et al. 2017; D\u2019Avanzo et al. 2018; Ghirlanda et al. 2019; Lazzati et al. 2018; Lyman et al. 2018) and the kilonova AT 2017gfo in the ultraviolet\u2013optical\u2013infrared band (Abbott et al. 2017c; Andreoni et al. 2017; Arcavi et al. 2017; Chornock et al. 2017; Coulter et al. 2017; Covino et al. 2017; Cowperthwaite et al. 2017; Evans et al. 2017; Hu et al. 2017; Kilpatrick et al. 2017; Lipunov et al. 2017; Nicholl et al. 2017; Smartt et al. 2017; Soares-Santos et al. 2017; Tanvir et al. 2017). The observations of GRB 170817A and its afterglows robustly confirmed the long-standing hypothesis that SGRBs can originate from compact binary mergers. Moreover, it became possible to explore the angular structure of the SGRB jet from an off-axis view (Lamb & Kobayashi 2017; Granot et al. 2018; Lazzati et al. 2018; Mooley et al. 2018a,b; Li et al. 2019). Meanwhile, the observations of AT 2017gfo indicated the existence of the merger ejecta, which suggests that the progenitor binary should at least contain one NS. In more detail, the existence of a \u201cblue\u201d and possibly also a \u201cpurple\u201d component in the AT 2017gfo emission further indicated that the merger product of the GW170817 event is very likely to be a hypermassive NS, which lasted for at least a few hundred milliseconds, as an immediately formed black hole can only be associated with a \u201cred\u201d kilonova1 (Cowperthwaite et al. 2017; Perego et al. 2017; Tanaka et al. 2017; Tanvir et al. 2017; Villar et al. 2017; Kawaguchi et al. 2018). Therefore, in summary, the progenitor of the GW170817 event can be identified as a DNS system, which is consistent with the result of the GW analysis.","Citation Text":["Tanvir et al. 2017"],"Functions Text":["In more detail, the existence of a \u201cblue\u201d and possibly also a \u201cpurple\u201d component in the AT 2017gfo emission further indicated that the merger product of the GW170817 event is very likely to be a hypermassive NS, which lasted for at least a few hundred milliseconds, as an immediately formed black hole can only be associated with a \u201cred\u201d kilonova1"],"Functions Label":["Background"],"Citation Start End":[[1949,1967]],"Functions Start End":[[1533,1880]]}
{"Identifier":"2022MNRAS.509.1010R__Laughlin_&_Adams_1998_Instance_1","Paragraph":"Recent work by Longmore, Chevance & Kruijssen (2021) has revealed an intriguing correlation between stellar phase-space density and the architecture of planetary systems, in particular the multiplicity. This work followed a similar analysis by Winter et al. (2020), which uncovered a correlation between stellar phase-space density and the occurrence of hot Jupiters. Using Gaia DR2 data (Gaia Collaboration et al. 2018), Longmore et al. (2021) computed the local stellar phase-space density of planet-hosting stars and their neighbours (within 40 pc) to determine whether the exoplanet host was in a relatively low or high phase-space density zone compared to its neighbours. They hypothesized that stars in current stellar overdensities were previously part of dense stellar clusters, from which only local residual overdensities remain. They showed that Kepler systems in local stellar phase-space overdensities have a significantly larger single-to-multiple ratio compared to those in the low phase-space density environment. The origin of this correlation is puzzling, as stellar clustering is expected to affect mostly the outer part of planetary systems in very dense environments (Laughlin & Adams 1998; Malmberg, Davies & Heggie 2011; Parker & Quanz 2012; Cai et al. 2017; Li, Mustill & Davies 2020a). Recent works have also suggested that the correlation is weaker when using a smaller unbiased stellar sample (Adibekyan et al. 2021), and that the current stellar overdensities could be associated with stellar age or to galactic-scale ripples as opposed to dense birth clusters (Mustill, Lambrechts & Davies 2021; Kruijssen et al. 2021). Alternatively, new studies have suggested that stellar flybys can excite the eccentricities and inclinations of outer planets\/companions, which then trigger the formation of hot Jupiters from cold Jupiters via high-eccentricity migration (Wang et al. 2020; Rodet, Su & Lai 2021). For this flyby scenario to be effective, certain requirements (derived analytically in Rodet et al. 2021) on the companion property (mass and semimajor axis) and the cluster property (such as stellar density and age) must be satisfied. In this paper, we will examine a similar \u2018outside\u2013in\u2019 effect of stellar flybys on the SE systems. Earlier, Zakamska & Tremaine (2004) examined the excitation and inward propagation of eccentricity disturbances in planetary systems. Our work focuses on inclination disturbances, as they are most relevant in determining the co-transit geometry of multiplanet systems.","Citation Text":["Laughlin & Adams 1998"],"Functions Text":["The origin of this correlation is puzzling, as stellar clustering is expected to affect mostly the outer part of planetary systems in very dense environments"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1189,1210]],"Functions Start End":[[1030,1187]]}
{"Identifier":"2015MNRAS.448.1644S__Rines_&_Diaferio_2006_Instance_1","Paragraph":"Despite a number of attempts of employing aspherical models for the matter distribution in the analysis of X-ray and lensing observations of galaxy clusters (see e.g. Corless, King & Clowe 2009; Samsing, Skielboe & Hansen 2012; Sereno et al. 2013), all dynamical methods are based on spherical symmetry. As the first step in addressing the problem of asphericity in mass measurements based on kinematics of galaxies in clusters, we study how dynamical mass estimators assuming spherical symmetry depend on the orientation of galaxy clusters with respect to the line of sight. We assess this effect by studying dynamical masses inferred from mock kinematic data of galaxy clusters generated from cosmological simulations. We restrict our analysis to dynamical masses measured with the so-called caustic technique (Diaferio 1999), which is one of the commonly used methods of mass determination in galaxy clusters (see e.g. Biviano & Girardi 2003; Rines & Diaferio 2006; Lemze et al. 2009; Geller et al. 2013; Rines et al. 2013). The caustic technique does not explicitly assume dynamical equilibrium beyond the virial radius, therefore it can be used to measure masses of galaxy clusters at distances larger than their virial radius and allows us to study the mass bias in a wide range of radii. As all dynamical methods currently applied to cluster data, it assumes spherical symmetry and testing this assumption is an objective of this work. Dependence of the mass measurement on the orientation of galaxy clusters with respect to the sight line is expected not only due to clusters\u2019 intrinsic phase-space shapes, but also due to co-alignment of the surrounding large-scale structures. The latter effect has been clearly shown both in cosmological simulations (Libeskind et al. 2013) and observations (Paz et al. 2011), and it is expected to modulate the contribution of background galaxies in kinematic samples and thus to affect the final estimate of dynamical mass.","Citation Text":["Rines & Diaferio 2006"],"Functions Text":["which is one of the commonly used methods of mass determination in galaxy clusters (see e.g."],"Functions Label":["Background"],"Citation Start End":[[946,967]],"Functions Start End":[[829,921]]}
{"Identifier":"2015AandA...576A...5C__J\u00f8rgensen_et_al._2012_Instance_3","Paragraph":"The relative abundances of the three species are derived from the column densities in Table 2 and are compared with other star-forming regions and comets in Table 3. The (CH2OH)2\/CH2OHCHO abundance ratio of ~0.3\u20130.5 previously derived in IRAS 16293 by J\u00f8rgensen et al. (2012) was revised. Indeed, the assignment in J\u00f8rgensen et al. (2012) was based on only one line of the gGg\u2032 conformer of ethylene glycol about 200 cm-1 (~290 K, M\u00fcller & Christen 2004) above the lowest-energy aGg\u2032 conformer \u2013 and thus tentative. An analysis from observations of six transitions of the lower energy conformer from ALMA Cycle 1 observations at 3 mm (four spectral windows at 89.48\u201389.73, 92.77\u201393.03, 102.48\u2013102.73, and 103.18\u2013103.42 GHz; J\u00f8rgensen et al., in prep.) results in a higher ethylene glycol-to-glycolaldehyde abundance ratio of 1.0\u2009\u00b1\u20090.3. This new estimate is consistent with the ratio expected between the aGg\u2032 and gGg\u2032 conformers under thermal equilibrium conditions at 300\u2009K, the excitation temperature of glycolaldehyde derived in IRAS 16293 (J\u00f8rgensen et al. 2012). The (CH2OH)2\/CH2OHCHO abundance ratio in IRAS2A is estimated at 5.5 \u00b1 1.0 if we consider the column densities derived from the rotational diagrams. It is slightly lower (4.6), however, if we use the column density of ethylene glycol of 1.1 \u00d7 1016 cm-2 that does not overproduce the peak intensities of a few lines (see Fig. 3). The (CH2OH)2\/CH2OHCHO abundance ratio consequently is a factor ~5 higher than in the Class 0 protostar IRAS 16293. It is also higher than in the other star-forming regions (see Table 3), but similar to the lower limits derived in comets (\u22733\u20136). This indicates that the glycolaldehyde chemistry may in general vary among hot corinos. It is possible that like IRAS2A, other very young low-mass protostars show high (CH2OH)2\/CH2OHCHO abundance ratios, in agreement with the cometary values. The CH3OCHO\/CH2OHCHO column density ratio found in IRAS2A (~20) ranges between the values derived in the molecular clouds from the Galactic center (~3.3\u20135.2) and the high-mass star-forming regions (~40\u201352). A lower limit of 2 was derived for comet Hale-Bopp. ","Citation Text":["J\u00f8rgensen et al. 2012"],"Functions Text":["This new estimate is consistent with the ratio expected between the aGg\u2032 and gGg\u2032 conformers under thermal equilibrium conditions at 300\u2009K, the excitation temperature of glycolaldehyde derived in IRAS 16293"],"Functions Label":["Similarities"],"Citation Start End":[[1044,1065]],"Functions Start End":[[836,1042]]}
{"Identifier":"2016MNRAS.462.3441D__Namouni_1999_Instance_2","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"],"Functions Text":["see equations 3 in","these are based on the vector eccentricity and the vector inclination."],"Functions Label":["Background","Background"],"Citation Start End":[[1126,1138]],"Functions Start End":[[1107,1125],[1141,1211]]}
{"Identifier":"2016ApJ...832...57P__Parashar_et_al._2009_Instance_1","Paragraph":"We employ two types of kinetic codes, hybrid particle-in-cell (PIC) and full PIC simulations. Both types make use of the P3D family of codes (Zeiler et al. 2002), in hybrid PIC (e.g., Parashar et al. 2011) mode, and fully kinetic PIC mode (e.g., Wu et al. 2013b). All simulations discussed here are performed in the 2.5D geometry (two-dimensional (2D) grid and all three components of field vectors). The hybrid simulation has \n\n\n\n\n\n (where \n\n\n\n\n\n is the ion inertial length, with c the speed of light and \n\n\n\n\n\n the proton plasma frequency), \n\n\n\n\n\n, 200 particles per cell, \n\n\n\n\n\n, cold isothermal electrons with \n\n\n\n\n\n. The simulation is initialized with energy only in wavevectors \n\n\n\n\n\n that have \n\n\n\n\n\n. v and b fluctuations are chosen with a specified initial spectral shape, Gaussian random phases, and only in essentially incompressive modes of the system. This simulation was also used in a recent study of variance anisotropy in kinetic plasmas (Parashar et al. 2016). The first full PIC simulation has \n\n\n\n\n\n, \n\n\n\n\n\n, 200 particles per cell, \n\n\n\n\n\n, \n\n\n\n\n\n. The initial condition is the Orszag\u2013Tang vortex (OTV) (e.g., Orszag & Tang 1979; Dahlburg & Picone 1989; Parashar et al. 2009; Vasquez & Markovskii 2012). This simulation was performed for a recent study of transition from kinetic to MHD-like behavior (Parashar et al. 2015). The final PIC simulation (Turb812) has \n\n\n\n\n\n, \n\n\n\n\n\n, 400 particles per cell, \n\n\n\n\n\n, \n\n\n\n\n\n. The initial condition is MHD-like, and more \u201cturbulent,\u201d with v and b fluctuations excited in a band of wave-vectors with \n\n\n\n\n\n with a specified initial spectrum. This simulation was done as part of a recent study that discussed the relation of timescales at the proton gyroscale and their relation to relative proton\u2013electron heating (Matthaeus et al. 2016). PIC codes have an inherent noise associated with them due to the finite number of particles per cell. While performing these simulations, the two most important numerical criteria that we paid attention to were: (i) excellent conservation of total energy (less than a few percent change in any fluctuation energy), and (ii) the particle noise in the spectrum was significant only at scales much smaller than the scales of interest (Debye length \n\n\n\n\n\n for PIC and di for hybrid PIC). On this basis, the modest number of particles employed here was considered adequate. As an additional measure, we employed filtering (e.g., Wan et al. 2012) to remove particle noise at grid scales prior to computing gradients (e.g., vorticity).","Citation Text":["Parashar et al. 2009"],"Functions Text":["The initial condition is the Orszag\u2013Tang vortex (OTV) (e.g.,"],"Functions Label":["Uses"],"Citation Start End":[[1174,1194]],"Functions Start End":[[1069,1129]]}
{"Identifier":"2022AandA...662A..42M__V\u00e1zquez_2007_Instance_2","Paragraph":"A number of fundamental results have been rigorously proved in the mathematical literature concerning the asymptotic behaviour in time of some of the solutions of the porous medium equation and related equations (e.g. Kamin & V\u00e1zquez 1991; Bernis et al. 1993; Hulshof et al. 2001). What is of interest for us here is, primarily, the results that can be applied to the cylindrically symmetric case with diffusion coefficient which is proportional to the square of the dependent variable (n\u2004=\u20042, m\u2004=\u20043 in the notation of Eq. (7)). The most basic result, already mentioned in Sect. 4.4.1, is that initial conditions which have a finite nonzero flux integral (called \u2018the mass\u2019 in the mathematical literature for the PME) converge toward the ZKBP solution with the same flux integral (\u2018mass\u2019) asymptotically in time (V\u00e1zquez 2007, Theorem 18.2); here, allowance is made for either a positive or negative flux integral by globally changing the sign of the ZKBP solution; also, \u2018convergence\u2019 is meant in the sense that the Lp norm of the difference between the actual solution and the ZKBP function tends to zero as t\u2004\u2192\u2004\u221e faster than a negative power of the time with an exponent which is a function of n, m, and p (e.g. \u22121\/3 for n\u2004=\u20042 and m\u2004=\u20043 in the L2 norm; see details in the book by V\u00e1zquez 2007). A complementary result is the following: when the initial condition has positive net flux and its negative part has compact support, then the whole solution evolves into a positive function after a finite time (V\u00e1zquez 2007, Theorem 18.29). Since we are dealing with signed functions which have zero flux integral, these results are of interest mainly because they impose a strict condition on the possible flux imbalance caused by numerical errors (as discussed in Sect. 4.4.1, final paragraph): if it is not small, the numerical solutions will approach the ZKBP solution in a comparatively short time. However, the flux imbalance in all the Bifrost experiments discussed in the present paper is small enough that they have not shown this behaviour even though they have been run until a very long diffusive time.","Citation Text":["V\u00e1zquez 2007"],"Functions Text":["also, \u2018convergence\u2019 is meant in the sense that the Lp norm of the difference between the actual solution and the ZKBP function tends to zero as t\u2004\u2192\u2004\u221e faster than a negative power of the time with an exponent which is a function of n, m, and p (e.g. \u22121\/3 for n\u2004=\u20042 and m\u2004=\u20043 in the L2 norm; see details in the book by"],"Functions Label":["Background"],"Citation Start End":[[1283,1295]],"Functions Start End":[[966,1282]]}
{"Identifier":"2022MNRAS.517.5744G__Caro_et_al._2016_Instance_1","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"],"Functions Text":["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","The explanation for this phenomenon motivated further research."],"Functions Label":["Background","Motivation"],"Citation Start End":[[260,277],[284,288]],"Functions Start End":[[0,220],[308,371]]}
{"Identifier":"2017ApJ...835..169O___2016_Instance_1","Paragraph":"Recent MIR studies have provided us with a much more realistic view of the central part of the AGNs. Spitzer studies of nearby Compton-thick AGNs have shown that even Compton-thick AGNs, especially low-luminosity ones, often show only modest to moderate silicate absorption at \n\n\n\n\n\n \u03bcm (e.g., Hao et al. 2007; Goulding et al. 2012). A classical smooth torus model, such as that of Pier & Krolik (1992), predicts deeper absorption in proportion to the X-ray absorption column density. On the other hand, if the torus is made of a collection of clouds, each cloud is heated to \u223c300 K to emit MIR emission while absorbing the background light when the foreground cloud is cooler than the one behind. The radiation transfer effect among the clouds significantly reduces the net silicate absorption even when the torus is seen edge-on (Nenkova et al. 2002, 2008a, 2008b; H\u00f6nig et al. 2006; H\u00f6nig & Kishimoto 2010; Stalevski et al. 2012, 2016). Meanwhile, recent MIR interferometric studies of nearby AGNs have started to directly reveal the dust distribution in the vicinity of the AGNs at parsec scales. In some best-studied AGNs, extended optically thin dust emission elongated toward the system\u2019s polar direction (e.g., direction of the extended narrow-line region or outflow) is typically found in addition to the compact disk-like component (e.g., Raban et al. 2009; H\u00f6nig et al. 2012, 2013; Tristram et al. 2012, 2014; L\u00f3pez-Gonzaga et al. 2016; see also Asmus et al. 2016 for the single-dish study; see Netzer 2015 for a review). Such extended polar emission is clearly inconsistent with the classical idea of the dusty torus in the unification theory, and its nature is under debate. Some proposed ideas are that it originates from the inner funnel of an extended dust distribution above and below the torus and\/or the dusty outflow within the ionizing cone that is radiatively driven from the inner wall of the compact dusty torus (e.g., H\u00f6nig et al. 2012, 2013; Keating et al. 2012; Roth et al. 2012; Tristram et al. 2014).","Citation Text":["Stalevski et al.","2016"],"Functions Text":["On the other hand, if the torus is made of a collection of clouds, each cloud is heated to \u223c300 K to emit MIR emission while absorbing the background light when the foreground cloud is cooler than the one behind. The radiation transfer effect among the clouds significantly reduces the net silicate absorption even when the torus is seen edge-on"],"Functions Label":["Differences"],"Citation Start End":[[910,926],[933,937]],"Functions Start End":[[485,830]]}
{"Identifier":"2021AandA...648A...3K__Lonsdale_et_al._2003_Instance_1","Paragraph":"LoTSS is currently mapping all of the northern sky to a high sensitivity and resolution (S150MHz ~ 0.1 mJy beam\u22121 and FWHM ~ 6\u2032\u2032) at the relatively unexplored 120\u2013168 MHz frequencies. In parallel with this, LOFAR is also undertaking deep observations of best studied multi-wavelength, degree scale fields in the northern sky, as part of the deep tier of LoTSS: the LoTSS Deep Fields (Tasse et al. 2021 and Sabater et al. 2021; hereafter Paper I and Paper II). The first three LoTSS Deep Fields are the European Large-Area ISO Survey-North 1 (ELAIS-N1; Oliver et al. 2000), Lockman Hole, and Bo\u00f6tes (Jannuzi & Dey 1999); these were chosen to have extensive multi-wavelength coverage from past and ongoing deep, wide-area surveys sampling the X-ray (e.g. Brandt et al. 2001; Hasinger et al. 2001; Manners et al. 2003; Murray et al. 2005), ultra-violet (UV; e.g. Martin et al. 2005; Morrissey et al. 2007) to optical (e.g. Jannuzi & Dey 1999; Cool 2007; Muzzin et al. 2009; Wilson et al. 2009; Chambers et al. 2016; Huber et al. 2017; Aihara et al. 2018) and to infrared (IR; e.g. Lonsdale et al. 2003; Lawrence et al. 2007; Ashby et al. 2009; Whitaker et al. 2011; Mauduit et al. 2012; Oliver et al. 2012) wavelengths; this is ideal for a wide range of our scientific objectives. These fields also benefit from additional radio observations at higher frequencies from the Giant Metrewave Radio Telescope (GMRT; e.g. Garn et al. 2008a,b; Sirothia et al. 2009; Intema et al. 2011; Ocran et al. 2019; Ishwara-Chandra et al. 2020) and the VLA (e.g. Ciliegi et al. 1999; Ibar et al. 2009). The current LoTSS Deep Fields dataset, covering ~ 26 deg2 (including multi-wavelength coverage) and reaching an unprecedented depth of S150MHz ~20 \u03bcJy beam\u22121, is comparable in depth to the deepest existing radio continuum surveys (e.g. VLA-COSMOS) but with more than an order of magnitude larger sky-area coverage. With this combination of deep, high-quality radio and multi-wavelength data over tens of square degrees, and along multiple sight-lines, the LoTSS Deep Fields are now able to probe a cosmological volume large enough to sample all galaxy environments to beyond z ~ 1, minimise the effects of cosmic variance (to an estimated level of ~4% for 0.5  z  1.0; Driver & Robotham 2010), and build statistical radio-selected samples of AGN and star-forming galaxies, even when simultaneously split by various physical parameters.","Citation Text":["Lonsdale et al. 2003"],"Functions Text":["The first three LoTSS Deep Fields are the European Large-Area ISO Survey-North 1 (ELAIS-N1","Lockman Hole, and Bo\u00f6tes","these were chosen to have extensive multi-wavelength coverage from past and ongoing deep, wide-area surveys sampling the X-ray","and to infrared (IR; e.g.","wavelengths"],"Functions Label":["Background","Background","Background","Background","Background"],"Citation Start End":[[1078,1098]],"Functions Start End":[[460,550],[573,597],[620,746],[1052,1077],[1204,1215]]}
{"Identifier":"2020MNRAS.493.4950S__Haines_et_al._2015_Instance_1","Paragraph":"In the framework of the hierarchical formation of structures, clusters of galaxies are continuously accreting galaxies. It has been suggested that in this process of falling, galaxies could undergo different physical processes that could affect the star formation even before they reach the cluster. Consequently, to fully understand what the cluster environment produces in galaxies, it is of key importance to have a throughout characterisation of the population of galaxies in the outskirts of clusters. Several observations have shown that properties of galaxies such as star formation, gas content, and colour are affected by the cluster environment at large clustercentric distances (e.g. Lewis et al. 2002; Solanes et al. 2002; G\u00f3mez et al. 2003; Braglia et al. 2009; Hansen et al. 2009; Park & Hwang 2009; von der Linden et al. 2010; Haines et al. 2015; Rhee et al. 2017). In particular, spiral galaxies with low star formation rates were found in the outskirts of clusters in early studies such as Couch et al. (1998) or Dressler et al. (1999). In recent years, a deficit of star-forming galaxies in the infalling region of clusters has been reported (e.g. Wetzel et al. 2013; Haines et al. 2015; Bianconi et al. 2018). This has been reproduced in simulations by Bah\u00e9 et al. (2013). These results can be explained by the presence of environmental effects accelerating the consumption of the gas reservoir before galaxies enter in a cluster, a process known as pre-processing (e.g. Fujita 2004; Mihos 2004). An important fraction of the cluster galaxies has spent time in groups or filaments before they fall into the cluster (e.g. McGee et al. 2009; De Lucia et al. 2012; Wetzel et al. 2013; Hou, Parker & Harris 2014). The population of galaxies in the outskirts of clusters includes not only galaxies that have not yet entered the cluster but also backsplash galaxies, i.e. galaxies that have passed close to the centre of the cluster since their infall and are now beyond the virial radius (e.g. Mamon et al. 2004; Gill, Knebe & Gibson 2005; Mahajan, Mamon & Raychaudhury 2011). For an adequate characterisation of the properties of galaxies that are falling into clusters, it is important to take into account the contamination by backsplash galaxies, which, having orbited through the inner regions of a cluster, could have been affected by the physical processes present in that extreme environment. The backsplash scenario in the evolution of galaxies has also been explored in Rines & Diaferio (2005), Pimbblet et al. (2006), Aguerri & S\u00e1nchez-Janssen (2010), and Muriel & Coenda (2014).","Citation Text":["Haines et al. 2015"],"Functions Text":["Several observations have shown that properties of galaxies such as star formation, gas content, and colour are affected by the cluster environment at large clustercentric distances (e.g."],"Functions Label":["Background"],"Citation Start End":[[842,860]],"Functions Start End":[[507,694]]}
{"Identifier":"2020AandA...641A.155V__Puglisi_et_al._2019_Instance_1","Paragraph":"The scenario presented above has been formulated in various flavors to individually explain several of the properties reported here. The main addition of this work, namely the excitation of CO in distant main-sequence and starburst galaxies, fits in the general picture that we sketched. The ensemble of properties and correlations that we reported here can be also used to revisit the definition of what a starburst is. A standard operational classification is based on the distance from the observed empirical M\u22c6-SFR correlation, the main sequence. This proved to be a useful distinction and an excellent predictor of several trends (e.g., Sargent et al. 2014), but recent results, including our present and previous analysis (Puglisi et al. 2019), show that the demarcation between starburst and main-sequence galaxies is more blurred that we previously considered. We do detect starburst-like behaviors in galaxies on the main sequence (Elbaz et al. 2018), likely linked to the existence of transitional objects (Popping et al. 2017; Barro et al. 2017b; G\u00f3mez-Guijarro et al. 2019; Puglisi et al. 2019, and in prep. to limit the references to recent works based on submillimeter observations). Such transition might well imply an imminent increase of the SFR, driving the object in the realm of starbursts (e.g., Barro et al. 2017b), or its cessation, bringing the system back onto or even below the main sequence (G\u00f3mez-Guijarro et al. 2019; Puglisi et al. 2019), with the CO properties potentially able to distinguish between these two scenarios. Regardless of these transitional objects, a definition of starburst based on \u03a3SFR, rather than \u0394MS, would naturally better account for the observed molecular gas excitation properties, dust temperatures and opacities, or SFE (see also Elbaz et al. 2011; Rujopakarn et al. 2011; Jim\u00e9nez-Andrade et al. 2018; Tacconi et al. 2020). As an example, in Fig. 8 we show the mean SLED of the subsample of galaxies with both CO (2\u2005\u2212\u20051) and CO (5\u2005\u2212\u20054) coverage, split at its median \u03a3SFR. While only tentative at this stage, this suggests a trend of increasing CO excitation with \u03a3SFR, consistently with Fig. 7 and what mentioned above.","Citation Text":["Puglisi et al. 2019"],"Functions Text":["This proved to be a useful distinction and an excellent predictor of several trends","but recent results, including our present and previous analysis","show that the demarcation between starburst and main-sequence galaxies is more blurred that we previously considered."],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[729,748]],"Functions Start End":[[551,634],[664,727],[751,868]]}
{"Identifier":"2017AandA...604A.118T__Zaritsky_et_al._1994_Instance_1","Paragraph":"In the Local Universe, most of our current knowledge of chemical distributions in disc galaxies comes from studies of the ISM via observations of HII regions or young stellar populations. Current observations are consistent with the metallicity profiles of the ISM having negative metallicity gradients, on average. Metallicity gradients in units of dex kpc-1 determine tight correlations with the global properties of galaxies such as the stellar mass or the size, which are erased when normalised by a characteristics radius such as the half-mass radius (e.g. Zaritsky et al. 1994; S\u00e1nchez et al. 2013). There are few indirect estimations of the evolution of the stellar metallicity gradients using planetary nebulae (PN; e.g. Henry et al. 2010; Stanghellini et al. 2010). Maciel et al. (2003) calculated the gradients of SPs with different ages in the Milky Way disc, finding a signal of increasingly negative metallicity gradients for older stars. Recently Magrini et al. (2016) analysed the metallicity gradients determined by PN for four nearby galaxies, finding them to be flatter than those detected using HII regions. Surveys such as CALIFA provide detailed information on the properties of the ISM and the SPs, including their chemical abundances and age distributions on a variety of galaxies (S\u00e1nchez-Bl\u00e1zquez et al. 2014; Gonz\u00e1lez Delgado et al. 2015). The SDDS-IV MaNGA survey (Bundy et al. 2015) also investigate spatially resolved SPs and radial age and metallicity gradients for nearby galaxies (Li et al. 2015; Wilkinson et al. 2015) and will provide a large statistical sample to confront with models. The available observations of HII regions for high-z galaxies do not allow the formation of a robust conclusion on the evolution of the metallicity gradients (e.g. Yuan et al. 2011; Queyrel et al. 2012; Stott et al. 2014; Jones et al. 2015). In fact, high-redshift observations show a complex situation with gas-phase components showing a variety of metallicity gradients that could respond to the action of different physical processes (e.g. Cresci et al. 2010). ","Citation Text":["Zaritsky et al. 1994"],"Functions Text":["Metallicity gradients in units of dex kpc-1 determine tight correlations with the global properties of galaxies such as the stellar mass or the size, which are erased when normalised by a characteristics radius such as the half-mass radius (e.g."],"Functions Label":["Background"],"Citation Start End":[[562,582]],"Functions Start End":[[316,561]]}
{"Identifier":"2020ApJ...897...73M__Coburn_2001_Instance_1","Paragraph":"Some of the accretion-powered X-ray pulsars showed additional features in emission between the 10 and 20 keV energy bands and more rarely in absorption between 8 and 10 keV in their respective residuals when fitted with a variety of continuum models (Coburn et al. 2002). Coburn et al. (2002) argued that such features may be caused by inadequacies of the continuum model rather than cyclotron resonance features. For example, it was observed that a single emission line at around 14 keV can fit two features around 10 and 20 keV for Vela X-1 (Kreykenbohm et al. 2002), Her X-1 (Coburn 2001), and Cep X-4 (McBride et al. 2007), and an absorption line between 8 and 10 keV can fit the features for 4U 1907+09, 4U 1538\u201352, and 4U 0352+309 as in Figure 6 of Coburn et al. (2002). In the case of GRO J2058+42, an introduction of a single Gaussian emission line in this range of energies could not appropriately fit the spectrum. The results described in Section 3.3 using ratios of spectral counts with respect to the Crab spectrum derived for four different phases of the pulsar (Figure 6) clearly indicate the presence of prominent depressions around 10 keV and 20 keV for phase 1 in particular, and its presence in other phases as well. Therefore, it confirms the presence of such absorption features in the spectral data associated with a physical origin and not due to any discrepancies of the continuum model as discussed above. Additionally, it also excludes the possibility of any uncertainty in the response matrix of the detector as the response matrix was not used to deconvolve the spectrum to calculate these ratios. The relative significance of these absorption features that was subsequently estimated after modeling the data and results is shown in Table 2 for the phase-averaged case and in Table 3 for the four different pulse phases. These observations strongly favor the presence and detection of these absorption features in their respective pulsar spectra.","Citation Text":["Coburn 2001"],"Functions Text":["For example, it was observed that a single emission line at around 14 keV can fit two features around 10 and 20 keV for","Her X-1","In the case of GRO J2058+42, an introduction of a single Gaussian emission line in this range of energies could not appropriately fit the spectrum."],"Functions Label":["Uses","Uses","Compare\/Contrast"],"Citation Start End":[[579,590]],"Functions Start End":[[414,533],[570,577],[777,924]]}
{"Identifier":"2017MNRAS.470.4099B__Smith_et_al._2003_Instance_1","Paragraph":"Numerical simulations using the N-body method are the primary instrument used to probe the non-linear regime of structure formation in cosmology and provide the basis for all theoretical predictions for the distribution of dark matter at the corresponding physical scales. Over the last few decades, such simulations have gained in refinement and complexity and have allowed the exploration of an ever larger range of scales (for a review see e.g. Bertschinger 1998; Springel et al. 2005; Dehnen & Reed 2011). Nevertheless, the understanding of their precision and their convergence towards the continuum limit remains, at very least, incomplete, in particular for smaller scales (see e.g. Splinter et al. 1998; Knebe et al. 2000; Romeo et al. 2008; Joyce et al. 2009; Power et al. 2016). In this context \u2018scale-free\u2019 cosmological models, in which both the expansion law and the power spectrum characterizing the initial fluctuations are simple power laws, have the advantage of relative simplicity, and they have for this reason been studied quite extensively in the literature (see e.g. Efstathiou et al. 1988; Colombi et al. 1996; Bertschinger 1998; Jain & Bertschinger 1998; Smith et al. 2003; Knollmann et al. 2008; Widrow et al. 2009; Orban 2013; Diemer & Kravtsov 2015). More specifically these models provide a testing ground for the numerical method through the predicted \u2018self-similarity\u2019 of the clustering: the temporal evolution of the clustering statistics must be equivalent to a rescaling of the distances. This follows from the fact that there is only one characteristic length scale (derived from the amplitude of the fluctuations) and one characteristic time-scale in the model. Further the exact rescaling function can be determined from the evolution in the linear regime of arbitrarily small fluctuations. However, discreteness and numerical effects typically introduce additional characteristic scales (e.g. force regularization at small scales, particle density, finite box size, etc.) which lead directly to a breaking of such self-similarity. Thus the self-similarity of clustering provides a potentially powerful tool to separate the scales affected by such non-physical effects from the physical results representing the continuum limit. The focus of this study is to exploit self-similar models to better understand the resolution at small scales of N-body simulations. In particular, we will use simulations with a very small force smoothing which allow us to follow carefully the propagation of self-similarity to small scales in the course of a simulation.","Citation Text":["Smith et al. 2003"],"Functions Text":["In this context \u2018scale-free\u2019 cosmological models, in which both the expansion law and the power spectrum characterizing the initial fluctuations are simple power laws, have the advantage of relative simplicity, and they have for this reason been studied quite extensively in the literature (see e.g."],"Functions Label":["Background"],"Citation Start End":[[1179,1196]],"Functions Start End":[[789,1088]]}
{"Identifier":"2022ApJ...928...18Z__Eastwood_et_al._2017_Instance_1","Paragraph":"The intensity of a geomagnetic storm is the result of the sustained interaction between the solar wind and the magnetosphere during the main phase of a storm. The most outstanding advantage of empirical formulae is their simplicity for space weather forecast: By inputting the solar wind parameters responsible for the main phase of a geomagnetic storm into an empirical formula, one can quickly get the estimated intensity of this geomagnetic storm. Of all empirical formulae, the Burton equation and OM equation have been used mostly to predict extreme geomagnetic storms. For example, the Burton equation was used to estimate the intensity of the storm that occurred on 1859 September 1\u20132 by Tsurutani et al. (2003), and Liu et al. (2014) used the OM equation to estimate the intensity of the storm on 2014 July 24. Extreme geomagnetic storms can cause widespread interference and damage to technological systems (Love 2021 and references therein) and then lead to significant economic loss (e.g., Council 2008; Schulte in den B\u00e4umen et al. 2014; Eastwood et al. 2017; Riley & Love 2017). Hence, it is very important to ensure the accuracy of extreme geomagnetic storm forecasts. Now the question is whether the Burton equation or the OM equation can estimate the intensities of very large geomagnetic storms correctly. To answer this question, we compare the performance of the three models mentioned above using 15 great geomagnetic storms (\u0394SYM-H \u2264 \u2212200 nT) that occurred during solar cycle 23. The solar wind parameters responsible for the main phases of the 15 storms are inputted into the three models to get the estimated intensities of the great geomagnetic storms (GGSs), which are compared with the observed intensities. To avoid any possible confusion, we use \u0394SYM-H\nb\n, \u0394SYM-Hom, and \u0394SYM-H\nw\n to indicate the intensities of a geomagnetic storm estimated by the Burton equation, the OM equation, and the WCL equation, respectively, while we use \u0394SYM-Hob to represent the observed intensity, namely the real variation of the SYM-H index during the main phase of a geomagnetic storm. The rest of this article is organized as follows. The data source and method are presented in Section 2, the results are presented in Section 3, and the discussion and summary are presented in Section 4.","Citation Text":["Eastwood et al. 2017"],"Functions Text":["Extreme geomagnetic storms can cause widespread interference and damage to technological systems","and then lead to significant economic loss (e.g.,"],"Functions Label":["Background","Background"],"Citation Start End":[[1050,1070]],"Functions Start End":[[819,915],[951,1000]]}
{"Identifier":"2016MNRAS.461.4317M__Zharikov_et_al._2008_Instance_1","Paragraph":"Finally, we compared the extinction-corrected upper limit on the optical flux of the PSR J0633+0632 PWN (Section 2.1) with its unabsorbed 0.3\u201310 keV X-ray flux. This is $F_{\\rm X}^{{\\rm pwn}} = 2.92^{+0.79}_{-0.81} \\times 10^{-13}$ erg cm\u22122 s\u22121 (Abdo et al. 2013), computed by fitting the PWN area with an ellipse of semimajor and semiminor axis of 0.58 and 0.54 arcmin, respectively, oriented 130\u00b0 due east (Marelli 2012). We subtracted the flux contribution of the point-like X-ray source south-west of the pulsar position (Fig. 2), which is only spatially coincident with the PWN. The extinction-corrected optical flux of the PWN in the g\u2032 band is $F_{\\rm opt}^{{\\rm pwn}} \\lesssim 9.8 \\times 10^{-13}$ erg cm\u22122 s\u22121, integrated over the same area as used to compute the PWN X-ray flux. As done for the pulsars, we assumed the most conservative value of the interstellar extinction. This yields an optical-to-X-ray flux ratio of Fopt\/Fx \u2272 4.6. PWNe have been detected both in the optical and X-rays around the Crab pulsar, PSR J0205+6449, PSR B0540\u221269, and PSR J1124\u22125916. Our upper limit on the Fopt\/Fx for the PSR J0633+0632 PWN is above the values obtained for the other PWNe, which are typically \u223c 0.02\u20130.04, apart from the Crab PWN which has an Fopt\/Fx \u223c 2 (Zharikov et al. 2008). This means that, owing to the faintness of the PSR J0633+0632 PWN in the X-rays, much deeper optical observations are needed to set similar constraints on its optical emission. We also compared the extinction-corrected optical spectral flux upper limit on the PWN in the g\u2032 and r\u2032-bands with the extrapolation of its X-ray spectrum in the optical domain. Like in Abdo et al. (2013), we used the best-fitting spectral index of the PWN, $\\Gamma _{\\rm X}^{{\\rm pwn}} = 1.19^{+0.59}_{-0.22}$. The PWN SED is shown in Fig. 5. As seen, we cannot rule the presence of a spectral break between the optical and the X-ray energy range. A break in the optical\/X-ray SED has been observed in other PWNe. For instance, the PWN around PSR B0540\u221269 features a clear break, with the optical fluxes being fainter than expected from the extrapolation of the X-ray PWN spectrum (Mignani et al. 2012). This is also the case for the PSR J1124\u22125916 PWN (Zharikov et al. 2008). A break in the opposite direction is observed in the SED of the PSR J1833\u22121034 PWN (Zajczyk et al. 2012), where the infrared fluxes (the PWN is not yet detected in the optical) are about two orders of magnitude above the extrapolation of the PWN X-ray spectrum. Only in the case of the Crab and PSR J0205+6449 PWNe, the PWN spectrum is compatible with a single PL, extending from the X-rays to the optical (Hester 2008; Shibanov et al. 2008). Optical detections of more PWNe through dedicated observing campaigns can allow one to relate the differences in the SEDs to the characteristics of the PWN.","Citation Text":["Zharikov et al. 2008","Zharikov et al. 2008"],"Functions Text":["Our upper limit on the Fopt\/Fx for the PSR J0633+0632 PWN is above the values obtained for the other PWNe, which are typically \u223c 0.02\u20130.04, apart from the Crab PWN which has an Fopt\/Fx \u223c 2","This is also the case for the PSR J1124\u22125916 PWN"],"Functions Label":["Differences","Compare\/Contrast"],"Citation Start End":[[1265,1285],[2220,2240]],"Functions Start End":[[1075,1263],[2170,2218]]}
{"Identifier":"2015MNRAS.446.4168R__Kroupa_2014_Instance_1","Paragraph":"The IGIMF theory (Kroupa et al. 2013) predicts a coupling between some properties of a galaxy (the SFR and the metallicity) and the IMF. Since the IMF in turn strongly affects the dynamical evolution of the galaxy, the feedback between galaxy evolution and IMF is difficult and the fully complexity of a variable IMF has not been yet included in hydrodynamical simulations (but see Bekki 2013; Ploeckinger et al. 2014; Recchi 2014). Even the here treated approach based on the so-called simple model of chemical evolution leads to implicit integral equations that must be solved iteratively. We note in passing that the dependence of the IMF on the metallicity is well established theoretically, as low-metallicity gas cools less efficiently and self-gravitating clumps are more resilient to fragmentation. This leads to the formation of dense cores of higher mass. For this reason, the IMF is supposed to be extremely biased towards massive stars if the metallicity is smaller than \u223c10\u22124 Z\u2299 (Schneider et al. 2002). Also the dependence of the IMF on the SFR is nowadays observationally well established (Hoversten & Glazebrook 2008; Meurer et al. 2009; Gunawardhana et al. 2011; Kroupa 2014), and the IGIMF theory is the only existing computable access accounting for these observations (Weidner et al. 2013). Thus, in spite of the inherent complexity, detailed simulations of galaxy evolution based on variable IMFs need to be performed. With this paper, we aimed at showing in a simple setting how to take into account IMF variations in models of the chemical evolution of galaxies (see also Martinelli & Matteucci 2000; Calura et al. 2010). The method outlined here can be used in order to extend analytical studies of the evolution of galaxies based on simple models of chemical evolution (e.g. Spitoni et al. 2010; Lilly et al. 2013; Pipino et al. 2014). We described how the solution of the simple model differs from the standard, textbook analytical solutions of the simple model. We also showed (building on previous results of K\u00f6ppen et al. 2007) that the IGIMF theory naturally leads to a MZ relation. In fact, low-mass galaxies are characterized on average by smaller SFRs. According to the IGIMF theory, a low SFR leads to a steep, top-light IMF, in which the production of heavy elements by massive stars is extremely limited. More massive galaxies instead produce many more massive stars because of the higher level of SFR, hence the attained present-day metallicity is larger.","Citation Text":["Kroupa 2014"],"Functions Text":["Also the dependence of the IMF on the SFR is nowadays observationally well established"],"Functions Label":["Background"],"Citation Start End":[[1180,1191]],"Functions Start End":[[1017,1103]]}
{"Identifier":"2019MNRAS.484.1487E__Roca-Fabrega_et_al._2013_Instance_2","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"],"Functions Text":["but see also a noticeable exception in"],"Functions Label":["Differences"],"Citation Start End":[[1081,1105]],"Functions Start End":[[1042,1080]]}
{"Identifier":"2021ApJ...919..140S__Bartos_et_al._2017_Instance_2","Paragraph":"Resonant dynamical friction may have applications beyond the relaxation of IMBHs examined in this paper. It may affect all objects in stellar clusters much more massive than the individual constituents of the disk, if present, including massive stars, stellar mass black holes (BHs), or the center of mass of massive binaries. Furthermore, it is also expected to operate in any type of disk with a high number of particles, including active galactic nucleus (AGN) accretion disks. Previously, it has been argued that stars and BHs crossing the disk on low-inclination orbits get captured by Chandrasekhar dynamical friction into the disk (Bartos et al. 2017; Panamarev et al. 2018; Tagawa et al. 2020). An interesting implication is that, if BHs settle into the disk, they interact dynamically and form BH\u2013BH binaries efficiently, and frequent dynamical interactions and gas effects drive the BHs to merger, producing gravitational waves (GWs) detectable by LIGO, VIRGO, and KAGRA (McKernan et al. 2014, 2018; Bartos et al. 2017; Leigh et al. 2018; Yang et al. 2019; Tagawa et al. 2020, 2021; Samsing et al. 2020). Mergers are also facilitated by Lidov\u2013Kozai oscillations in anisotropic systems (Heisler & Tremaine 1986; Petrovich & Antonini 2017; Hamilton & Rafikov 2019). The results in this paper show that resonant dynamical friction may accelerate the capture of objects in the accretion disks by a factor proportional to the SMBH mass over the local disk mass for large orbital inclinations. Pressure and viscosity in a gaseous disk do not inhibit the orbit-averaged torque from the IMBH, which leads to realignment and the warping of the disk (Bregman & Alexander 2012). Thus, RDF may efficiently catalyze the alignment of the orbital planes of BHs even in low-luminosity AGN or Seyfert galaxies with relatively small disk masses, which may not be possible for Chandrasekhar dynamical friction. In fact, this mechanism extends the scope of the \u201cAGN merger channel\u201d for GW source populations even beyond low-luminosity AGN and Seyfert galaxies, as it may organize BHs into disks also in nonactive galaxies with nuclear stellar disks.","Citation Text":["Bartos et al. 2017"],"Functions Text":["An interesting implication is that, if BHs settle into the disk, they interact dynamically and form BH\u2013BH binaries efficiently, and frequent dynamical interactions and gas effects drive the BHs to merger, producing gravitational waves (GWs) detectable by LIGO, VIRGO, and KAGRA"],"Functions Label":["Background"],"Citation Start End":[[1010,1028]],"Functions Start End":[[703,980]]}
{"Identifier":"2021AandA...645A.141V__Metcalfe_et_al._2016_Instance_1","Paragraph":"This paper presents the results of our study of the effect of a dynamically adapting heat-conduction prescription, based on Kramers opacity law, in conjunction with semi-global MHD simulations. The main aim is to determine the effect of this prescription on the two major transitions reported in numerical studies (e.g., Gastine et al. 2014; Viviani et al. 2018). One concerns the rotation profiles, and is the transition from accelerated poles and decelerated equator to a solar-like profile, with a faster equator. The other involves the large-scale magnetic field, and is the transition from an axisymmetric magnetic field, as in the Sun, to a nonaxisymmetric magnetic field found in more rapid rotators. Previous studies (Viviani et al. 2018) found these transitions to occur at the same rotation rate, in contrast with the current interpretation of observations. The fact that simulations usually produce anti-solar differential rotation for the solar rotation rate could indicate that the Sun is in a transitional regime (e.g., K\u00e4pyl\u00e4 et al. 2014; Metcalfe et al. 2016), or could also mean that simulations still cannot fully capture the true rotational influence on turbulent convection in the Sun. Lehtinen et al. (2016) reported on the existence of nonaxisymmetric structures in stars with varying rotation rates, and were therefore able to determine quite a sharp transition point in terms of the rotation period, when fields turn from axi- to nonaxisymmetric configurations. According to dynamo theory, these two modes can compete, and there can be a transition region, where both dynamo modes co-exist, as is also clearly demonstrated by the models presented in this paper and those of Viviani et al. (2018). Therefore, the observational transition point must be regarded as a lower limit for the transition in terms of the rotation period, as it could be that the sensitivity of the current instruments is insufficient to detect the very weak nonaxisymmetric components. However, since active longitudes have not been detected on the Sun (Pelt et al. 2006), these two transitions should not be located at the same, nearly solar rotation rate.","Citation Text":["Metcalfe et al. 2016"],"Functions Text":["The fact that simulations usually produce anti-solar differential rotation for the solar rotation rate could indicate that the Sun is in a transitional regime (e.g.,","or could also mean that simulations still cannot fully capture the true rotational influence on turbulent convection in the Sun."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1054,1074]],"Functions Start End":[[868,1033],[1077,1205]]}
{"Identifier":"2020ApJ...889...15Y__Yang_et_al._2016b_Instance_2","Paragraph":"Although each one of the four aforementioned mechanisms has some observational support in certain systems, there is not a single mechanism that can explain all observed polarization in protoplanetary disks. Alignment with respect to the local radiation anisotropy (\u201ck-RAT alignment\u201d thereafter) is best supported by the azimuthal polarization pattern observed at ALMA Band 3 in the HL Tau system (Kataoka et al. 2017). However, it predicts a strong azimuthal variation of polarization and circular pattern (rather than elliptical pattern) (Yang et al. 2019). There is some tentative evidences for alignment with respect to the magnetic field, through either Radiative Alignment Torques (\u201cB-RAT alignment\u201d; Lazarian & Hoang 2007), or recently proposed Mechanical Alignment Torques (Hoang et al. 2018), in, e.g., the IRAS 4A system at cm wavelengths (Cox et al. 2015; Yang et al. 2016b) and BHB07-11 (Alves et al. 2018) at (sub)millimeter wavelengths. But there is no well-resolved system that matches the theoretical expectations (see, e.g., Cho & Lazarian 2007; Yang et al. 2016b; Bertrang et al. 2017) assuming the widely expected disk toroidal magnetic field yet (Flock et al. 2015). Mechanical alignment has recently received some attention. Hoang et al. (2018) claims that under MATs, grains can be aligned with respect to local dust-gas streaming direction, in the case of a weak or zero magnetic field, even if the velocity difference is sub-sonic. Within this picture, Kataoka et al. (2019) investigated the direction of streaming velocities for dust grains with different Stokes numbers, and the resulting polarization orientations. They found that their polarization pattern in the order-of-unity Stokes number case resembles that observed by Alves et al. (2018) in BHB07-11. The BHB07-11, however, is a binary system, and we expect more complicated velocity fields than the simple one assumed in Kataoka et al. (2019). Yang et al. (2019) investigated the observational features of another mechanical alignment mechanism, the Gold mechanism (Gold 1952), to address the circular versus elliptical pattern problem in the ALMA Band 3 polarization observations of HL Tau disk. However, they failed to explain the nonexistence of strong azimuthal variation, and suggested the scattering by dust grains aligned under the Gold mechanism may be the origin of the polarization at ALMA Band 3 in the HL Tau system.","Citation Text":["Yang et al. 2016b"],"Functions Text":["But there is no well-resolved system that matches the theoretical expectations (see, e.g.,"],"Functions Label":["Motivation"],"Citation Start End":[[1062,1079]],"Functions Start End":[[950,1040]]}
{"Identifier":"2016MNRAS.455..449H__McGaugh_2012_Instance_1","Paragraph":"With only six free parameters, the standard \u039b cold dark matter (\u039bCDM) cosmological model fits no less than 2500 multipoles in the cosmic microwave background (CMB) angular power spectrum (Planck Collaboration XVI 2014), the Hubble diagram of Type Ia supernovae, the large-scale structure matter power spectrum, and even the detailed scale of baryonic acoustic oscillations. It thus provides the current basis for simulations of structure formation, and is extremely successful down to the scale of galaxy clusters and groups. Nevertheless, it still faces numerous challenges on galaxy scales. Among these, the most important ones are the too-big-to-fail problem (Boylan-Kolchin, Bullock & Kaplinghat 2011) and the satellite-plane problem (e.g. Pawlowski, Pflamm-Altenburg & Kroupa 2012; Ibata et al. 2014) for dwarf galaxies, the tightness of the baryonic Tully\u2013Fisher relation (McGaugh 2012; Vogelsberger et al. 2014), or the unexpected diversity of rotation curve shapes at a given mass scale (Oman et al. 2015). The latter problem is actually a subset of a more general problem, i.e. that the shapes of rotation curves indeed do not depend on the Dark Matter (DM) halo mass, contrary to what would be expected in \u039bCDM, but rather on the baryonic surface density, as has long been noted (e.g. Zwaan et al. 1995). This makes the problem even worse, since the rotation curve shapes are not only diverse at a given mass scale, but uniform at a given baryonic surface density scale, implying a completely ununderstood fine-tuning of putative feedback mechanisms. On the other hand, this behaviour of rotation curves is an a priori prediction of the formula proposed by Milgrom more than 30 yr ago (Milgrom 1983a,b), relating the total gravitational field to the Newtonian field generated by baryons alone, and which can be interpreted as a modification of Newtonian dynamics on galaxy scales below a characteristic acceleration (Modified Newtonian Dynamics (MOND), for a review see Famaey & McGaugh 2012; Milgrom 2014). With this simple formula, high surface brightness (HSB) galaxies are predicted to have rotation curves that rise steeply before becoming essentially flat, or even falling somewhat to the not-yet-reached asymptotic circular velocity, while low surface brightness (LSB) galaxies are predicted to have rotation curves that rise slowly to the asymptotic velocity. This is precisely what is observed, and was predicted by Milgrom long before LSB galaxies were even known to exist. The formula also predicts the tightness of the baryonic Tully\u2013Fisher relation.","Citation Text":["McGaugh 2012"],"Functions Text":["Nevertheless, it still faces numerous challenges on galaxy scales. Among these, the most important ones are","for dwarf galaxies, the tightness of the baryonic Tully\u2013Fisher relation"],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[879,891]],"Functions Start End":[[526,633],[806,877]]}
{"Identifier":"2019MNRAS.488.5029H__Stacey_et_al._2010_Instance_3","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"],"Functions Text":["The comparison sample is compiled from the literature up to z \u223c 3"],"Functions Label":["Uses"],"Citation Start End":[[1479,1497]],"Functions Start End":[[1343,1408]]}
{"Identifier":"2017MNRAS.472.1152R__Cenko_et_al._2010_Instance_1","Paragraph":"Alternatively, if a magnetar is the central engine powering GRBs, we might expect to see periodic features in the emission. Known magnetars have clear periodic signals in their emission caused by their rotation periods (e.g. Mazets et al. 1979; Kouveliotou et al. 1998). The X-ray pulsations typically contribute to 30\u2009per\u2009cent of the signal, with a range of 10\u201380\u2009per\u2009cent (Israel et al. 1999; Kargaltsev et al. 2012; Kaspi & Beloborodov 2017). There is an energy dependence on the pulsed fraction of the signal, where low energies tend to have smaller pulsed fractions (Vogel et al. 2014). Detection of a periodic signal during the plateau phase in the X-ray light curve would provide excellent supporting evidence for the magnetar central engine model. There have been searches for a periodic signal in the prompt emission of GRBs with a number of instruments with no success, for example: Burst And Transient Source Experiment (BATSE) GRBs ( Deng & Schaefer 1997), INTErnational Gamma-Ray Astrophysics Laboratory (INTEGRAL) GRBs (Ryde et al. 2003), GRB 051103 (an extragalactic Soft Gamma-ray Repeater giant flare candidate detected by the Inter Planetary Network; Hurley et al. 2010) and Burst Alert Telescope (BAT) GRBs (Cenko et al. 2010; de Luca et al. 2010; Guidorzi et al. 2012). Dichiara et al. (2013) searched the prompt emission of a number of short GRBs for evidence of a precessing jet (predicted by Stone, Loeb & Berger 2013). However, these searches typically target the prompt emission and have not probed the regime where we might expect periodic signals from a magnetar central engine (i.e. during the plateau phase). Only two GRBs have been searched for periodic emission during the X-ray observations when the magnetar central engine may dominate the emission, GRB 060218 (Mirabal & Gotthelf 2010) and GRB 090709A (Mirabal & Gotthelf 2009; de Luca et al. 2010). The prompt emission of GRB 090709A possibly showed evidence of a periodic signal (Golenetskii et al. 2009; Gotz et al. 2009; Markwardt et al. 2009; Ohno et al. 2009), however this was ruled out with a more careful analysis of the prompt data from BAT, X-ray Telescope (XRT) and X-ray Multi-mirror Mission (XMM) observations of the X-ray afterglow (Cenko et al. 2010; de Luca et al. 2010). However, in the majority of these studies, the authors have targeted a constant spin period whereas a magnetar central engine is expected to have a rapidly decelerating spin period which would be very difficult to detect in standard searches for periodic signals. Dichiara et al. (2013) did conduct a deceleration search, however they were targeting signals in the prompt emission where we do not expect the signal from a spinning down magnetar.","Citation Text":["Cenko et al. 2010"],"Functions Text":["There have been searches for a periodic signal in the prompt emission of GRBs with a number of instruments with no success, for example:","and Burst Alert Telescope (BAT) GRBs"],"Functions Label":["Background","Background"],"Citation Start End":[[1227,1244]],"Functions Start End":[[756,892],[1189,1225]]}
{"Identifier":"2020ApJ...889L..10M__McKay_et_al._2018_Instance_2","Paragraph":"As stated earlier, during review of this manuscript Croviser et al. announced in a CBET a tentative water production rate approximately five times larger than our reported value. While the brief nature of the CBET precludes a detailed comparison, we discuss some possible reasons for this discrepancy. At the high airmass of these observations and the small dimensions of the ARCES slit, differential refraction can result in wavelength-dependent slit loss, which can skew flux measurements. However, this is not expected for [O i] 6300 \u212b emission because this feature is close to the guiding wavelength (\u223c5500 \u212b). We confirmed that this is indeed negligible for [O i] 6300 \u212b emission based on observations of comet C\/2012 S1 (ISON) that were performed at a similarly high airmass with ARCES, and found that the production rates derived from the ISON [O i] 6300 \u212b measurements were consistent with values determined using other methods (McKay et al. 2018). Therefore, we do not consider this or other airmass-dependent phenomena as the reason for the discrepancy. At certain geocentric velocities the cometary [O i] 6300 \u212b emission sits on top of a strong telluric absorption, and at high airmass inaccurate removal of this feature can result in a decrease in the measured flux and therefore production rate. This was observed for C\/2012 S1 (ISON) (McKay et al. 2018). However, the geocentric velocity of 2I\/Borisov during our observations was \u223c\u221235 km s\u22121, while the effect on observed [O i] 6300 \u212b line fluxes in comet ISON was only observed at geocentric velocities of \u223c\u221250 km s\u22121. Therefore, this is also not a likely candidate to explain the discrepancy. It is also possible that the activity is highly variable, and we observed Borisov at a minimum in activity, while the Nan\u00e7ay observations, which were coadded over three weeks of observations, provide a long-term average production rate. However, no such variability is observed for CN, with the CN production rate being fairly constant over a several week period (Kareta et al. 2019; Opitom et al. 2019).","Citation Text":["McKay et al. 2018"],"Functions Text":["At certain geocentric velocities the cometary [O i] 6300 \u212b emission sits on top of a strong telluric absorption, and at high airmass inaccurate removal of this feature can result in a decrease in the measured flux and therefore production rate. This was observed for C\/2012 S1 (ISON)"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1349,1366]],"Functions Start End":[[1064,1347]]}
{"Identifier":"2020MNRAS.493.4950S__Linden_et_al._2010_Instance_1","Paragraph":"In the framework of the hierarchical formation of structures, clusters of galaxies are continuously accreting galaxies. It has been suggested that in this process of falling, galaxies could undergo different physical processes that could affect the star formation even before they reach the cluster. Consequently, to fully understand what the cluster environment produces in galaxies, it is of key importance to have a throughout characterisation of the population of galaxies in the outskirts of clusters. Several observations have shown that properties of galaxies such as star formation, gas content, and colour are affected by the cluster environment at large clustercentric distances (e.g. Lewis et al. 2002; Solanes et al. 2002; G\u00f3mez et al. 2003; Braglia et al. 2009; Hansen et al. 2009; Park & Hwang 2009; von der Linden et al. 2010; Haines et al. 2015; Rhee et al. 2017). In particular, spiral galaxies with low star formation rates were found in the outskirts of clusters in early studies such as Couch et al. (1998) or Dressler et al. (1999). In recent years, a deficit of star-forming galaxies in the infalling region of clusters has been reported (e.g. Wetzel et al. 2013; Haines et al. 2015; Bianconi et al. 2018). This has been reproduced in simulations by Bah\u00e9 et al. (2013). These results can be explained by the presence of environmental effects accelerating the consumption of the gas reservoir before galaxies enter in a cluster, a process known as pre-processing (e.g. Fujita 2004; Mihos 2004). An important fraction of the cluster galaxies has spent time in groups or filaments before they fall into the cluster (e.g. McGee et al. 2009; De Lucia et al. 2012; Wetzel et al. 2013; Hou, Parker & Harris 2014). The population of galaxies in the outskirts of clusters includes not only galaxies that have not yet entered the cluster but also backsplash galaxies, i.e. galaxies that have passed close to the centre of the cluster since their infall and are now beyond the virial radius (e.g. Mamon et al. 2004; Gill, Knebe & Gibson 2005; Mahajan, Mamon & Raychaudhury 2011). For an adequate characterisation of the properties of galaxies that are falling into clusters, it is important to take into account the contamination by backsplash galaxies, which, having orbited through the inner regions of a cluster, could have been affected by the physical processes present in that extreme environment. The backsplash scenario in the evolution of galaxies has also been explored in Rines & Diaferio (2005), Pimbblet et al. (2006), Aguerri & S\u00e1nchez-Janssen (2010), and Muriel & Coenda (2014).","Citation Text":["von der Linden et al. 2010"],"Functions Text":["Several observations have shown that properties of galaxies such as star formation, gas content, and colour are affected by the cluster environment at large clustercentric distances (e.g."],"Functions Label":["Motivation"],"Citation Start End":[[814,840]],"Functions Start End":[[507,694]]}
{"Identifier":"2020MNRAS.496..152R__Springel_&_Hernquist_2002_Instance_1","Paragraph":"From observations it is known that the interstellar gas has a complex structure with different phases \u2013 a hot volume-filling phase and a dense cold gas phase \u2013 both of which should be represented in galaxy formation simulations (see Naab & Ostriker 2017; Tumlinson et al. 2017). To model this, we treat the gas as a multiphase medium with many co-existing phases (see Marri & White 2003). We let two SPH particles i and j decouple into separate phases, if the following two conditions apply (Aumer et al. 2013):\n(1)$$\\begin{eqnarray*}\r\n\\max \\left(\\frac{A_i}{A_j}, \\frac{A_j}{A_i} \\right) \\gt 50 , \\quad -\\mu _{ij} \\lt c_{ij} .\r\n\\end{eqnarray*}$$Here, Ai,j are the entropic functions of the particles (Springel & Hernquist 2002), $\\mu _{ij} := (\\boldsymbol{v}_i - \\boldsymbol{v}_j) \\cdot \\frac{\\boldsymbol{r}_i - \\boldsymbol{r}_j}{|\\boldsymbol{r}_i - \\boldsymbol{r}_j|}$ is the relative velocity of the particles along their vector of separation, and cij is the pair-averaged sound speed. Two SPH particles decouple if their entropy (actually their entropic functions2) are very different unless they approach faster than with the local sound speed. The velocity restriction (equation 1) is required to capture shocks properly (Marri & White 2003). This multiphase treatment results in a continuum of phases from cold to hot and the results are not very sensitive to the exact ratio in equation (1). This model is aimed at preventing overcooling, i.e. artificially short cooling times (see e.g. Naab & Ostriker 2017, for a review) and allows for the simultaneous representation of, and energy injection into, a hot and a cold phase on the resolution scale (see Scannapieco et al. 2006, for a detailed discussion). Such multiphase ISM structure naturally arises in much higher resolution simulations of the supernova-driven multiphase ISM (see e.g. Walch et al. 2015). For further details on the multiphase model, see Marri & White (2003) and Aumer et al. (2013, 2014).","Citation Text":["Springel & Hernquist 2002"],"Functions Text":["Here, Ai,j are the entropic functions of the particles"],"Functions Label":["Uses"],"Citation Start End":[[701,726]],"Functions Start End":[[645,699]]}
{"Identifier":"2021MNRAS.504.5702W__Werk_et_al._2014_Instance_1","Paragraph":"Another notable accretion-regulated halo property is gas temperature, with the CGM of haloes nominally having both a hot and cold-phase. The hot coronal gas phase (at \u2248Tvir) originates from the virial shock-heating of gas accreting high-mass haloes ($M_{\\rm halo}\\gtrsim 10^{12}\\, \\mathrm{M}_{\\odot }$; e.g. Rees & Ostriker 1977). A cold-phase of the CGM at \u2248104 K has also been observed, but its origins are less clear (e.g. Adelberger et al. (e.g. Adelberger et al. 2003; Stocke et al. 2006; Lehner & Howk 2011; Prochaska, Hennawi & Simcoe 2013; Werk et al. 2014; Zhu et al. 2014; Heckman et al. 2017; Zahedy et al. 2019). Several origins of this cool CGM phase have been proposed, namely pristine IGM accretion (e.g. simulation-based findings in van de Voort & Schaye 2012; Afruni, Fraternali & Pezzulli 2019, 2021), the condensation of hot halo gas (e.g. the empirical arguments of Voit 2018 and illustris-TNG findings in Nelson et al. 2020), feedback-driven outflows (e.g. Bouch\u00e9 et al. 2013; Borthakur et al. 2015; Angl\u00e9s-Alc\u00e1zar et al. 2017; Oppenheimer et al. 2018; Hafen et al. 2019), and the stripping of satellite galaxies in larger systems (e.g. Hafen et al. 2019 using the fire-2 simulations). Afruni et al. (2019, 2021), using semi-analytic models and results from the COS-Halos and COS-GASS surveys, argue that star-formation-driven outflows cannot account for the amount of cool gas in the CGM of observed haloes, pointing towards IGM accretion as the origin of this gas. Additionally, the radial variation of CGM properties was explored in Fielding et al. (2020) using a number of hydrodynamical simulations (as part of the smaug project). They find that the properties of the outer-CGM (at \u22730.5R200,crit) are shaped by larger scale processes, such as cosmological accretion, rather than galactic feedback that dominates the inner regions, \u22720.5R200,crit. In any case, the wide range of observed metallicities and temperatures observed implies that the CGM is a diverse, multi-phase gas reservoir, making it an ideal laboratory to study the influence of cosmological inflows.","Citation Text":["Werk et al. 2014"],"Functions Text":["A cold-phase of the CGM at \u2248104 K has also been observed, but its origins are less clear"],"Functions Label":["Motivation"],"Citation Start End":[[548,564]],"Functions Start End":[[331,419]]}
{"Identifier":"2021MNRAS.508.2583Z__Sch\u00f6ier_et_al._2002_Instance_1","Paragraph":"Located in the star-forming region \u03c1-Ophiuchi, inside the dark cloud L1689N and at a distance of 141 pc (Dzib et al. 2018), IRAS16293\u22122422 is a well-studied Young Stellar Object (YSO) classified as a Class 0 source with less than 104 yr (Andre, Ward-Thompson & Barsony 1993), and represents one of the very early stages of low-mass star formation. It was the first source identified as a hot corino (Blake et al. 1994; van Dishoeck et al. 1995) based on the detection of Complex Organic Molecules (COMs) in the source, which was later supported by follow-up studies (Ceccarelli et al. 1998, 2000; Sch\u00f6ier et al. 2002; Crimier et al. 2010; J\u00f8rgensen et al. 2011, 2016; Pineda et al. 2012; Oya et al. 2016; Jacobsen et al. 2018; van der Wiel et al. 2019). Higher resolution observations revealed that IRAS16293\u22122422 is in fact a triple system, composed of sources A1 and A2, separated by 54 au from each other (Maureira et al. 2020) and source B, 738 au (5 arcse; Wootten 1989) away from source A. Due to this larger separation, tidal truncation between the three protostars is discarded and therefore source B is considered to have evolved as an isolated source (Rodr\u00edguez et al. 2005). It was initially proposed to be either an evolved T Tauri star (Stark et al. 2004; Takakuwa et al. 2007) or a very young object (Chandler et al. 2005), however, Chandler et al. (2005) suggested that source B has large-scale infalls based on SO line emission. Pineda et al. (2012) confirmed the infall of an inner envelope, with mass accretion rates of 4.5 \u00d7 10\u22125 M\u2299yr\u22121, based on ALMA detections of inverse P-Cygni profiles in CH3OCHO-E, CH3OCHO-E-A and H2CCO, ruling out the possibility of it being a T Tauri star. The interpretations of infall from these profiles was also suggested by J\u00f8rgensen et al. (2012) and Zapata et al. (2013). Unlike the A1 and A2 protostars, source B has not shown clear signs of outflow launching, explained by the lack of free\u2013free emission at low frequencies (Chandler et al. 2005; Rodr\u00edguez et al. 2005; Loinard et al. 2007; Rao et al. 2009; Liu et al. 2018; Hern\u00e1ndez-G\u00f3mez et al. 2019b) and also based on molecular lines (Loinard et al. 2002; van der Wiel et al. 2019).","Citation Text":["Sch\u00f6ier et al. 2002"],"Functions Text":["It was the first source identified as a hot corino","based on the detection of Complex Organic Molecules (COMs) in the source, which was later supported by follow-up studies"],"Functions Label":["Background","Background"],"Citation Start End":[[597,616]],"Functions Start End":[[348,398],[445,565]]}
{"Identifier":"2020ApJ...898...52M__Elmegreen_&_Scalo_2004_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":["Elmegreen & Scalo 2004"],"Functions Text":["Turbulence on scales less than the scale height of the warm\u2013cold ISM likely originates primarily due to the feedback from young stars"],"Functions Label":["Background"],"Citation Start End":[[843,865]],"Functions Start End":[[708,841]]}
{"Identifier":"2020MNRAS.497..302T__Watarai_2006_Instance_1","Paragraph":"In this paper, we explore the conditions required for hyper-Eddington accretion on to a BH when both radiative and mechanical feedback operate simultaneously, performing two-dimensional (2D) hydrodynamical simulations with multifrequency radiation transfer. We conduct a comprehensive survey on the parameter dependence of outflow models, varying the outflow opening angle, mass loading degree into outflows, velocity of outflows, and density of gas surrounding the BH. To model mechanical feedback, we adopt a phenomenological model proposed by Ostriker et al. (2010), while radiative feedback is treated by adopting the standard and slim disc model (Shakura & Sunyaev 1973; Abramowicz et al. 1988; Watarai 2006) as in Takeo et al. (2019). We find that the flow structure consists of two distinct parts in the early sub-Eddington phase; the bipolar outflowing region heated up to T \u223c 106\u22127 K due to strong shock and the equatorial inflowing region where ionized gas is mildly heated to T \u223c 105 K due to photoionization. When the ambient gas density exceeds a critical threshold, as in the cases where only radiative feedback is included (Inayoshi et al. 2016; Takeo et al. 2018), the mass accretion rate on to the nuclear region rises to a hyper-Eddington value. Since mechanical power of outflows driven by the rapidly accreting BH is sufficiently strong, bipolar outflows completely evacuate the surrounding gas in the polar region and reduce the mass inflow (BH accretion) rate by a factor of \u22483\u201313 (\u22486\u201326, respectively) from the case without mechanical feedback. Furthermore, we find that the critical gas density required for hyper-Eddington accretion is reduced by a factor of \u223c3 and the transition occurs in a shorter dynamical time-scale when mechanical feedback is modelled in the simulations. In fact, the effects that alleviate the transition to rapid accretion tend to be more prominent as the outflow is stronger, i.e. a wider opening angle, higher mass loading factor, and higher outflow velocity. This is because suppression of BH accretion owing to outflows reduces the radiative output from the nuclear BH, leading to hyper-Eddington accretion.","Citation Text":["Watarai 2006"],"Functions Text":["To model mechanical feedback, we adopt a phenomenological model proposed by Ostriker et al. (2010), while radiative feedback is treated by adopting the standard and slim disc model"],"Functions Label":["Uses"],"Citation Start End":[[700,712]],"Functions Start End":[[470,650]]}
{"Identifier":"2019AandA...630A..30L__H\u00e4ssig_et_al._(2015)_Instance_3","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)"],"Functions Text":["As reported in",", 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."],"Functions Label":["Background","Background","Motivation"],"Citation Start End":[[1070,1090]],"Functions Start End":[[1055,1069],[1090,1227],[1228,1546]]}
{"Identifier":"2021AandA...645A..99C__analysis,_Uttley_et_al._(2011)_Instance_2","Paragraph":"X-ray reverberation in black hole X-ray binaries was first robustly detected in GX 339\u20134 by Uttley et al. (2011) when the source was in its hard state. Previous studies of GX 339\u20134 pointed to the approximate central mass being \u22656\u2006M\u2299 (e.g. Hynes et al. 2003) and a small disc inclination angle (De Marco et al. 2015). Miller et al. (2008) fitted the Suzaku spectra and found that the central black hole has a very high spin, a\u2004\u223c\u20040.998. The X-ray spectroscopic analysis of the hard state spectra from the RXTE archive carried out by Garc\u00eda et al. (2015) suggested the black hole spin to be a\u2004\u223c\u20040.95. Spectral fitting of GX 339\u20134 during its very high flux state using NuSTAR and Swift also suggested a high spin of a\u2004\u223c\u20040.95 (Parker et al. 2016). According to the time-lag analysis, Uttley et al. (2011) found that the disc thermal emission (\u223c0.3\u20130.7 keV, soft band) leads the power-law variations (\u223c0.7\u20131.5 keV, hard band) on long timescales (> 1s). Mahmoud et al. (2019) assumed that the soft component that leads the power-law emission is a soft Comptontized component. Rapisarda et al. (2016) and Rapisarda et al. (2017) instead modelled it as a variable inner region of the thin disc. However, the disc black-body variations lag behind the power-law variations by a few milliseconds on short timescales ( 1s). This switch from low-frequency hard to high-frequency soft lags is thought to be produced by two distinct mechanisms. While the hard lags are likely due to inward propagating fluctuations (e.g. Kotov et al. 2001; Ar\u00e9valo & Uttley 2006), the soft lags can be explained by thermal reverberation associated with the longer light-travel time the hard photons take from the central power-law X-ray source to the disc where they are reprocessed into relatively soft black-body emission. The thermal reverberation lags then provide clues to the geometry of the X-ray source and the inner accretion flow close to the event horizon of the central black hole.","Citation Text":["Uttley et al. (2011)"],"Functions Text":["According to the time-lag analysis,","found that the disc thermal emission (\u223c0.3\u20130.7 keV, soft band) leads the power-law variations (\u223c0.7\u20131.5 keV, hard band) on long timescales (> 1s)."],"Functions Label":["Background","Background"],"Citation Start End":[[779,799]],"Functions Start End":[[743,778],[800,946]]}
{"Identifier":"2022ApJ...929...65I__Ikhsanov_2002_Instance_1","Paragraph":"Our parameter range reaches into the \u201cpropeller\u201d regime, where the azimuthal velocity of the star\u2019s outer magnetosphere exceeds the Keplerian velocity at disk truncation (Romanova et al. 2005), i.e., where the disk is truncated near and outside of the corotation radius. Observations of the propeller regime have been discussed in relation to various astrophysical systems, such as rapidly rotating neutron stars, white dwarfs in cataclysmic variables, and CTTS (see, e.g., Stella et al. 1986; Treves et al. 1993; Cui 1997; Alpar 2001; Ek\u015f\u0131 & Alpar 2003; Mori & Ruderman 2003). The propeller regime has been investigated both analytically (Davies et al. 1979; Li & Wickramasinghe 1997; Lovelace et al. 1999; Ikhsanov 2002; Rappaport et al. 2004; Ek\u015fi et al. 2005), and via MHD simulations (Wang & Robertson 1985; Romanova et al. 2003, 2004, 2005, 2009; Ustyugova et al. 2006). It has been shown that, for a rapidly rotating star with a strong magnetic field, a significant proportion of accreting material is in fact centrifugally expelled as it reaches the inner region of the disk and is redirected as a propeller-driven outflow, allowing the star to rapidly spin down (Davidson & Ostriker 1973; Illarionov & Sunyaev 1975; Lipunov 1992). For example, MHD simulations performed by Romanova et al. (2005) and Zanni & Ferreira (2013) demonstrate the following quasi-periodic behavior: (1) disk material accumulates at the inner region of the disk, which moves the truncation radius closer to the star; (2) some material can then accrete onto the star, reducing the mass of material at the disk inner edge; (3) remaining inner-disk material gains angular momentum and is ejected at the centrifugal barrier, which moves the truncation radius further away from the star, for the cycle to repeat. In this regime, accretion becomes intermittent, or can even be inhibited completely, and the net effect acts to remove angular momentum from the star. Therefore, it is possible that the propeller mechanism is responsible for the slow rotation rates observed in CTTS.","Citation Text":["Ikhsanov 2002"],"Functions Text":["The propeller regime has been investigated both analytically"],"Functions Label":["Background"],"Citation Start End":[[708,721]],"Functions Start End":[[578,638]]}
{"Identifier":"2016ApJ...833....7Y__Owen_&_Wu_2013_Instance_3","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"],"Functions Text":["In this paper, we adopt Mpi = 20 M\u2295 and Mpf = 10 M\u2295 nominally (close to the standard model adopted in"],"Functions Label":["Similarities"],"Citation Start End":[[2228,2242]],"Functions Start End":[[2126,2227]]}
{"Identifier":"2021MNRAS.504.5992M__Ho_et_al._2015_Instance_1","Paragraph":"Characterizing and understanding the distribution and transport of chemical elements inside galaxies is a critical aspect of galaxy evolution. Successive generations of star formation enrich the interstellar medium (ISM) with metals. Therefore, the spatial distribution of chemical abundances in galaxies is a powerful tracer of the history of gas flows, star formation, accretion, and mergers throughout their assembly (e.g. Edmunds & Greenhow 1995; Kewley et al. 2010; Torrey et al. 2012; Finlator 2017; Ma et al. 2017; Bresolin 2019; Hemler et al. 2020). Over the last few decades, our understanding of chemical inhomogeneities in galaxies has advanced dramatically, largely due to the advent of integral field unit (IFU) spectroscopy (see Maiolino & Mannucci 2019 for a review). Negative radial metallicity gradients have been widely observed in the low-redshift galaxy population (e.g. Searle 1971; Vila-Costas & Edmunds 1992; Berg et al. 2013, 2020; Ho et al. 2015; Belfiore et al. 2017; Poetrodjojo et al. 2018) with other studies additionally observing azimuthal variations from these radial metallicity trends (Li, Bresolin & Kennicutt 2013; Vogt et al. 2017; Ho et al. 2018, 2019; Kreckel et al. 2020). Observational limitations mean that metallicities are more challenging to measure at high redshift. In the absence of gravitational lensing, spatial resolution is reduced. Furthermore, fainter targets mean that metallicities are typically derived from fewer emission lines, limiting our ability to control for possible redshift evolution in the ISM conditions of galaxies when constructing an abundance scale (see Appendix A for a more comprehensive discussion). Existing determinations of radial gradients in high-redshift galaxies are generally limited to modest samples of lensed galaxies, or samples of the largest disc galaxies, and show substantial amounts of scatter from steep negative gradients to positive gradients (Yuan et al. 2011; Swinbank et al. 2012; Jones et al. 2013; Leethochawalit et al. 2016; Wuyts et al. 2016; Carton et al. 2018; Wang et al. 2019a,b; Curti et al. 2020b; Gillman et al. 2021). Extending observations to smaller, fainter galaxies remains a challenge.","Citation Text":["Ho et al. 2015"],"Functions Text":["Negative radial metallicity gradients have been widely observed in the low-redshift galaxy population (e.g."],"Functions Label":["Background"],"Citation Start End":[[956,970]],"Functions Start End":[[783,890]]}
{"Identifier":"2017MNRAS.464..635M__Dekel_et_al._2009_Instance_3","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"],"Functions Text":["Gravitational interactions in the perturbed disc drive turbulence causing the disc to self-regulate in a marginally stable state with Q \u2272 1","that can last for more than a Gyr so long as the accretion is not interrupted."],"Functions Label":["Background","Background"],"Citation Start End":[[1368,1385]],"Functions Start End":[[1227,1366],[1517,1595]]}
{"Identifier":"2019MNRAS.490.5722W__Remus_et_al._2013_Instance_2","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"],"Functions Text":["Apart from cosmological simulations, dedicated zoom-in simulations","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"],"Functions Label":["Background","Background"],"Citation Start End":[[1836,1853]],"Functions Start End":[[1713,1779],[1855,2004]]}
{"Identifier":"2021MNRAS.503.2776Y__Ajith_et_al._2007_Instance_1","Paragraph":"In order to investigate the signal-to-noise ratio (SNR), \u03c1 of NS\u2013WD binaries for LISA-type space GW detectors, we calculate the averaged square SNR $\\overline{\\rho ^{2}}$ over the sky location, inclination, and polarization as \n(30)$$\\begin{eqnarray*}\r\n\\overline{\\rho ^{2}} = \\int _{f_{1}}^{f_{2}}\\frac{4\\cdot \\frac{4}{5}fA^{2}(f)}{(P_{\\rm n}(f)\/R(f))} \\rm d (\\ln \\it f),\r\n\\end{eqnarray*}$$(Moore, Cole & Berry 2015; Robson, Cornish & Liu 2019), where f1 and f2 are the lower and upper limits of the integral, respectively. The factor 4 in the numerator of the integrand comes from the addition of strain noise in the detector arms and the two-way noise in each arm (Larson, Hiscock & Hellings 2000). We calculate the GW amplitudes A(f) of NS\u2013WD binaries using the phenomenological (PhenomA) waveform model in the Fourier domain (Ajith et al. 2007; Robson et al. 2019). A(f) is expressed as \n(31)$$\\begin{eqnarray*}\r\nA(f) = \\sqrt{\\frac{5}{24}}\\frac{G^{5\/6}\\mathcal {M}^{5\/6}}{\\pi ^{2\/3}c^{3\/2}R_{\\rm b}}f^{-7\/6}\\, {\\rm Hz}^{-1},\\,\\,\\,\\it f\\lt f_{\\rm m},\r\n\\end{eqnarray*}$$\n (32)$$\\begin{eqnarray*}\r\nf_{\\rm m} = \\frac{0.2974\\zeta ^{2}+0.04481\\zeta +0.09556}{\\pi (GM\/c^{3})}\\, {\\rm Hz},\r\n\\end{eqnarray*}$$\n (33)$$\\begin{eqnarray*}\r\n\\zeta = m_{1}m_{2}\/M^{2},\r\n\\end{eqnarray*}$$where $\\mathcal {M} \\equiv m_1^{3\/5} m_2^{3\/5}\/(m_1+m_2)^{1\/5}$, and fm is the GW frequency at the point of merging. If f > fm, the index of the power-law relation between A(f) and f changes (Ajith et al. 2007) and is beyond the scope of this study. The power spectral density of total detector noise $P_{\\rm n}=\\frac{1}{L^{2}}\\left[P_{\\rm o}+2(1+\\cos ^{2}(f\/f_{\\ast }))\\frac{P_{\\rm a}}{(2\\pi f)^{4}}\\right]$, where f* = c\/(2\u03c0L), L = 2.5 \u00d7 109\u2009m is the armlength of the detector, $P_{\\rm o}=2.25\\times 10^{-22} \\,\\rm m^{2}\\left(1+(\\frac{2\\,mHz}{\\it f})^{4}\\right) \\,\\, \\rm Hz^{-1}$ is the single-link optical metrology noise, and $P_{\\rm a}=9.0\\times 10^{-30} \\,\\rm (m\\,s^{-2})^{2}\\left(1+(\\frac{0.4\\,mHz}{\\it f})^{2}\\right)\\left(1+(\\frac{\\it f}{8\\,\\rm mHz})^{4}\\right) \\,\\,Hz^{-1}$ is the single test mass acceleration noise (LISA Science Study Team 2018; Robson et al. 2019). R(f) is the transfer function numerically calculated from Larson et al. (2000). The effective noise power spectral density can be defined as Sn(f) = Pn(f)\/R(f). For Taiji and Tianqin, we use the sensitivity curve data in Ruan et al. (2020) and Wang et al. (2019), respectively.","Citation Text":["Ajith et al. 2007"],"Functions Text":["We calculate the GW amplitudes A(f) of NS\u2013WD binaries using the phenomenological (PhenomA) waveform model in the Fourier domain"],"Functions Label":["Uses"],"Citation Start End":[[830,847]],"Functions Start End":[[701,828]]}
{"Identifier":"2022ApJ...924...42N__Torres_et_al._2013_Instance_1","Paragraph":"In this model, the pulsar associated with the PWN loses its rotational energy via a pulsar wind composed of magnetic and high-energy particles to power the high-energy physical process inside the nebula (Atoyan & Aharonian 1996; Fang & Zhang 2010). The relativistic wind of particles driven by the pulsar is blown into the ambient medium and generates a termination shock wave, which accelerates the electrons to relativistic energy. These relativistic electrons interact with the magnetic field and low-energy background photons (the synchrotron, thermal, FIR, and microwave background radiation), and generate the multiband nonthermal photons with energies ranging from radio to high-energy gamma-ray bands (Zhang et al. 2008; Fang & Zhang 2010; Lu et al. 2017). According to the review of the leptonic model (see, e.g., Zhang et al. 2008; Venter & de Jager 2007; Torres et al. 2013), the electrons injected into PWNe are accelerated by the pulsar magnetosphere and the termination shock. Therefore, the relativistic particles injected into the PWNe are also assumed as two different power-law components from the pulsar magnetosphere and shock acceleration, respectively. The injected spectrum of relativistic electrons inside PWNe is described as\n1\n\n\n\nQ(Ee,t)=Q0(t)Ee\/Ecut\u2212\u03b11ifEeEcutQ0(t)Ee\/Ecut\u2212\u03b12ifEe\u2265Ecut,\n\nwhere the Q\n0 can be determined by the \u222bQ(E\n\ne\n, t)E\n\ne\n\ndE\n\ne\n =\u03b7\nL(t); \u03b7 is the conversion efficiency from spin-down power into electron luminosity. The maximum energy of the electrons was express as \n\n\n\nEmax(t)\u2248\u03b50L(t)L0\n\n, and L\n0 is initial spin-down power. The electron energy distribution was given by\n2\n\n\n\ndN(Ee,Tage)dt=\u222b0TageQ(Ee,t)exp\u2212Tage\u2212t\u03c4effdt,\n\nwhere the \n\n\n\n\u03c4eff\u22121=\u03c4esc(t)\u22121+\u03c4syn(t)\u22121\n\n, \u03c4\n\nesc\n(t) is the escape timescale, and \u03c4\nsyn(t) is the lifetime of the relativistic electron of the synchrotron emission loss. The details of temporal evolution about the electron in the PWNe are discussed by the Zhang et al. (2008; also see the version of Fang & Zhang 2010; Lu et al. 2017).","Citation Text":["Torres et al. 2013"],"Functions Text":["According to the review of the leptonic model (see, e.g.,","the electrons injected into PWNe are accelerated by the pulsar magnetosphere and the termination shock. Therefore, the relativistic particles injected into the PWNe are also assumed as two different power-law components from the pulsar magnetosphere and shock acceleration, respectively."],"Functions Label":["Uses","Uses"],"Citation Start End":[[866,884]],"Functions Start End":[[765,822],[887,1174]]}
{"Identifier":"2015AandA...584A.103S__Chamel_et_al._2011_Instance_3","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"],"Functions Text":["and a comparison with the other EoSs of the BSk family","and the RMF family","shall be left for future study."],"Functions Label":["Future Work","Future Work","Future Work"],"Citation Start End":[[1963,1981]],"Functions Start End":[[1907,1961],[2047,2065],[2088,2119]]}
{"Identifier":"2021ApJ...922L..11H__Castaing_et_al._1990_Instance_1","Paragraph":"Most of these observations are consistent with the general understanding that, for increasing distance from the Sun, the turbulent power-law spectrum expands toward larger scales (Bavassano et al. 1982; Bruno & Carbone 2013; Chen et al. 2020). However, spectral properties may be misleading, as it is not possible to unequivocally ascribe the Kolmogorov-like power-law scaling to the presence of a turbulent cascade. For example, it is universally observed that turbulence is associated with intermittency (Kolmogorov 1962), as is routinely observed in solar wind measurements (Sorriso-Valvo et al. 1999; Bruno & Carbone 2013). A standard way for characterizing the intermittency of a field \u03d5 (\u03d5 being, for example, a velocity or magnetic field component) is by means of the scale-dependent increments \u0394\u03d5 = \u03d5(t + \u0394t) \u2212 \u03d5(t), which account for the presence of gradients on a timescale \u0394t (Anselmet et al. 1984). Intermittency is related to the scale-dependent shape of the probability distribution function of the increments \u0394\u03d5 (Castaing et al. 1990). This, additionally, implies the existence of nonvanishing odd moments. In particular, a scaling law can be derived for the third-order moment directly from the dynamical MHD equations, as the conservation law of the appropriate inviscid invariants (de Karman & Howarth 1938; Danaila et al. 2001). Such a relation, known in the MHD description as the Politano\u2013Pouquet (PP) law (Politano & Pouquet 1998; Carbone et al. 2009), establishes that under the hypothesis of homogeneity, stationarity, local isotropy, and incompressibility, in the turbulent inertial range the mixed third-order moment of the increments of velocity (v) and magnetic field (in velocity units, \n\n\n\n\nb\n=\nB\n\n\/\n\n\n\n4\n\u03c0\n\u03c1\n\n\n\n\n, with B the magnetic field vector and \u03c1 the plasma mass density) is a linear function of the scale \u0394t. Moreover, the proportionality coefficient is related to the mean energy transfer rate of the turbulent cascade. This can be written as\n1\n\n\n\n\n\n\n\nY\n(\n\u0394\nt\n)\n\n\n\u2261\n\n\n\u2329\n\u0394\n\n\nv\n\n\nR\n\n\n(\n\u2223\n\u0394\nv\n\n\n\u2223\n\n\n2\n\n\n+\n\u2223\n\u0394\nb\n\n\n\u2223\n\n\n2\n\n\n)\n\n\n\n\n\n\n\n\n\u2212\n2\n\u0394\n\n\nb\n\n\nR\n\n\n(\n\u0394\nv\n\u00b7\n\u0394\nb\n)\n\u232a\n=\n\n\n\n4\n\n\n3\n\n\n\n\u03b5\n\n\nV\n\n\nsw\n\n\n\u0394\nt\n.\n\n\n\n\n\nHere \u0394v and \u0394b are scale-dependent vector increments of the plasma velocity and magnetic field, as defined above; \u0394vR = vR(t + \u0394t) \u2212 vR(t) and \u0394bR = bR(t + \u0394t) \u2212 bR(t) are velocity and magnetic field longitudinal increments measured in the sampling direction R; \u03b5 is the mean energy transfer rate; and brackets indicate ensemble average. Note that the solar wind speed Vsw is used for switching between space scales, \u2113, and timescales, \u0394t, through the Taylor hypothesis, \u2113 = Vsw\u0394t (Taylor 1938). This also results in the reversal of the sign in the left-hand side of Equation (1) with respect to the traditional formulation in terms of spatial increments. The PP law is a fundamental relation for MHD turbulence, since it describes the energy cascade and ultimately allows us to estimate the energy that will be dissipated at small scales. It is particularly relevant for solar wind turbulence, where the collisionless processes responsible for removing the energy at the bottom of the nonlinear cascade are not yet fully understood (Chen et al. 2019; Sorriso-Valvo et al. 2019; Matthaeus et al. 2020; Smith & Vasquez 2021). The linear relation (1) has been observed in various regions of the heliosphere, for different conditions of the space plasma, confirming the turbulent nature of their dynamics and providing an estimate of the energy transfer rate (MacBride et al. 2005; Sorriso-Valvo et al. 2007; Marino et al. 2008; Smith et al. 2009; Stawarz et al. 2010; Coburn et al. 2012; Bandyopadhyay et al. 2018; Hadid et al. 2018; Andr\u00e9s et al. 2019; Bandyopadhyay et al. 2020; Sorriso-Valvo et al. 2021).","Citation Text":["Castaing et al. 1990"],"Functions Text":["Intermittency is related to the scale-dependent shape of the probability distribution function of the increments \u0394\u03d5","This, additionally, implies the existence of nonvanishing odd moments."],"Functions Label":["Background","Background"],"Citation Start End":[[1028,1048]],"Functions Start End":[[911,1026],[1051,1121]]}
{"Identifier":"2016MNRAS.459..277S__Ensslin_et_al._1998_Instance_1","Paragraph":"Diffuse synchrotron emission associated with ultrarelativistic particles and magnetic fields in the intracluster medium (ICM) primarily consists of radio haloes and radio relics (see Ferrari et al. 2008; Br\u00fcggen et al. 2012; Feretti et al. 2012; Brunetti & Jones 2014 for recent reviews). It is thought that radio haloes are caused by cluster-wide post-merger turbulence (see e.g. Brunetti et al. 2001; Petrosian 2001), secondary electrons from proton\u2013proton interactions (see e.g. Dennison 1980; Blasi & Colafrancesco 1999) or a combination of the two mechanisms (see Brunetti & Blasi 2005; Brunetti & Lazarian 2011; Pinzke, Oh & Pfrommer 2015), whereas radio relics are apparently associated with localized, post-merger shock-fronts (Ensslin et al. 1998). However, the non-detection of gamma-ray emission by the Fermi satellite (see e.g. Brunetti et al. 2012; Zandanel & Ando 2014; Ackermann et al. 2015) disfavours a purely hadronic model for the origin of radio haloes and challenges standard diffuse shock acceleration to explain radio relics (see e.g. Brunetti & Jones 2014; Vazza & Br\u00fcggen 2014). Additionally, the variety of the observed properties of cluster-scale radio emission is becoming increasingly difficult to describe within the current theoretical picture. For example: whilst in the \u2018Sausage\u2019 cluster (CIZA J2242.8+5301) van Weeren et al. (2010) observe a textbook example of an arc-like radio relic related to a shock; in ZwCl 2341.1+000 Ogrean et al. (2014) observe no X-ray shock at the position of a relic, and in the Bullet cluster (Shimwell et al. 2015), PLCKG287.0+32.9 (Bonafede et al. 2014b) and the Coma cluster (Ensslin et al. 1998) an apparent link is seen between radio galaxies and radio relics. Furthermore, whilst radio haloes are statistically detected in merging clusters (e.g. Cassano et al. 2013), the role of the cluster mass and its dynamical state is difficult to disentangle with current observations (see Cuciti et al. 2015). For example, in CL1821+643 Bonafede et al. (2014a) detect a giant radio halo in a cool core cluster with no obvious merging activity and Russell et al. (2011) observe no diffuse radio emission in Abell 2146, which is a less massive cluster, but a clear merging system. Recently upgraded and new facilities have significantly improved sensitivity to diffuse radio emission from the ICM and are already beginning to reveal increasingly complex phenomena (e.g. Owen et al. 2014) which may shed light on the connection between haloes and relics and further challenge a univocal interpretation of these sources. One such instrument is the Low-Frequency Array (LOFAR; van Haarlem et al. 2013) which can produce deep, high-resolution, high fidelity, low-frequency radio images.","Citation Text":["Ensslin et al. 1998","Ensslin et al. 1998"],"Functions Text":["whereas radio relics are apparently associated with localized, post-merger shock-fronts","and the Coma cluster","an apparent link is seen between radio galaxies and radio relics."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[736,755],[1643,1662]],"Functions Start End":[[647,734],[1621,1641],[1664,1729]]}
{"Identifier":"2015MNRAS.450.3458C__Cichowolski_et_al._2001_Instance_2","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)"],"Functions Text":["Although","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."],"Functions Label":["Differences","Differences"],"Citation Start End":[[599,624]],"Functions Start End":[[590,598],[625,851]]}
{"Identifier":"2015MNRAS.446..330W__Milgrom_2009c_Instance_1","Paragraph":"Consider a space\u2013time scale invariance of the equations of motion under the consideration of the transformation in Minkowsky space (Milgrom 2009a; see also Milgrom 2014a,b):\n(3)\n\n\\begin{equation}\n(t,{\\boldsymbol r}) \\rightarrow (\\lambda t,\\lambda {\\boldsymbol r}),\n\\end{equation}\n\nwhere t and ${\\boldsymbol r}=(x,y,z)$ are time and Cartesian coordinates, respectively, and \u03bb is a positive number. The Newtonian gravitational acceleration for a spherically symmetric system,\n\n(4)\n\n\\begin{equation}\ng_{\\rm N} = \\frac{GM_{\\rm b}}{r^2},\n\\end{equation}\n\nthen transforms as gN \u2192 \u03bb\u22122gN, whereas the kinematical acceleration, $g\\equiv {\\rm d} \\dot{x} \/ {\\rm d} t$, scales as g \u2192 \u03bb\u22121g. Here Mb(  r) is the enclosed baryonic mass within r. As a result, the Newtonian gravitational acceleration and the kinematical acceleration scale differently under equation (3). Linking purely gravitational interactions to symmetries such as defined in equation (3) suggests deeper physics and constitutes a motivation for viewing MOND as much more than a mere phenomenological description of galactic dynamics (Milgrom 2009c). In order to assure that both the gravitational and the kinematical accelerations scale symmetrically under equation (3), i.e. in order to maintain the invariant symmetry, the gravitational acceleration, g, has to scale proportionally to $g_{\\rm N}^{1\/2}$. In order to obtain the correct dimension, a constant with the unit of acceleration, needs to be introduced. This constant is referred to as a0, such that\n\n(5)\n\n\\begin{equation}\ng=(a_0g_{\\rm N})^{1\/2},\n\\end{equation}\n\ni.e. g2 = a0gN. Thus g = (GMba0)1\/2\/r, and the circular velocity, which follows from the centrifugal acceleration g = v2\/r, is\n\n(6)\n\n\\begin{equation}\nv = (GM_{\\rm b}a_0)^{1\/4}={\\rm constant},\n\\end{equation}\n\nwhich is exactly the BTFR (Milgrom 1983a, 2009a, 2014b; McGaugh et al. 2000; Famaey & McGaugh 2012). We refer to gravitational dynamics which thus conforms to low-acceleration scale invariance (equation 3) as low-acceleration SID. SID beautifully reproduces the deep MOND equations of motion. It is rather remarkable that such a simple principle as SID and discovered by Milgrom leads to one of the most important scaling relations which real galaxies are observed to obey. Note that equation (6) implies that each baryonic galaxy is surrounded by a logarithmic non-particle (and thus phantom) DM halo potential, which is however not a real halo as it is only evident if Newtonian dynamics is applied to the galaxy. If SID is true, then a Newtonian observer would thus deduce that each baryonic galaxy is surrounded by a PDMH the mass of which is proportional to radial distance (equation 27 below).","Citation Text":["Milgrom 2009c"],"Functions Text":["Linking purely gravitational interactions to symmetries such as defined in equation (3) suggests deeper physics and constitutes a motivation for viewing MOND as much more than a mere phenomenological description of galactic dynamics"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1089,1102]],"Functions Start End":[[855,1087]]}
{"Identifier":"2021MNRAS.505..515Z__Naul_et_al._2018_Instance_2","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"],"Functions Text":["Although period-folding improves performance","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."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[953,969]],"Functions Start End":[[907,951],[972,1125]]}
{"Identifier":"2021AandA...646A..96C__Brusa_et_al._2018_Instance_2","Paragraph":"AGN-driven outflows. Another possible effect of the AGN activity on the molecular gas is through outflows. This possibility is supported by observations of individual objects: For example, Carniani et al. (2017), Brusa et al. (2018) and Loiacono et al. (2019) find low gas fractions in powerful AGN at cosmic noon hosting high-velocity molecular and ionized outflows (but see also Herrera-Camus et al. 2019). AGN feedback in action in these targets could be depleting the molecular gas reservoir (Brusa et al. 2015). F\u00f6rster Schreiber et al. (2019), studying outflows in a large sample of 0.6\u2004\u2004z\u2004\u20042.7 galaxies through integral field spectroscopy of the H\u03b1 emission line, find that incidence, strength, and velocity of AGN-driven winds are strongly correlated with the stellar mass. In particular, they find that high-velocity (\u223c1000\u20132000 km s\u22121) AGN-driven outflows are commonly detected at masses above log(M*\/M\u2299) = 10.7, and present in up to 75% of the population for log(M*\/M\u2299) > 11.2. Interestingly, above this stellar-mass threshold we find a significant CO luminosity deficit in our AGN sample with respect to inactive galaxies (Fig. 3, bottom). Moreover, our AGN show on average gas fractions 0.57 dex (by using uniform assumptions, Sect. 4) lower than inactive galaxies at the 2.2\u03c3 level. Quantitatively, this translates into Mgas,\u2006mol\/M*\u2004\u2248\u20040.3 for AGN (0.16 if we use r31\u2004=\u20040.92; Kirkpatrick et al. 2019) and \u22481 in inactive galaxies. This representative value for our AGN is in line but not as low as previous work targeting extremely powerful sources (e.g., Mgas,\u2006mol\/M*\u2004\u20040.05 in Brusa et al. 2018). Our team is performing a systematic investigation of ionized gas outflows with SINFONI as part of the SUPER survey, and 11 targets of our ALMA sample have complementary good quality SINFONI data (Kakkad et al. 2020; Perna et al., in prep.). For some of them we measured [O\u202fIII] line widths larger than 600 km s\u22121, interpreted as a clear signature of the presence of an AGN-driven outflow in these objects (Kakkad et al. 2020). A detailed comparison between outflow and CO properties for these targets will be presented in a future work. Distinguishing among the scenarios described above is challenging with the current dataset. AGN feedback could proceed in different ways and different mechanisms likely overlap in shaping the properties of the molecular gas reservoir. For example, AGN radiation could both heat and\/or dissociate CO molecules. In this case, AGN would produce a feedback mechanism that does not require outflows but would potentially work toward inhibiting further star formation. As for AGN-driven outflows, they could impact the gas content by ejecting material out of the galaxy (e.g., Travascio et al. 2020), or they could produce CO heating or dissociation due to shocks. Additionally, numerical simulations predict that AGN-driven outflows may heat via shocks a significant quantity of the gas in the ISM, reaching the high temperatures required for the excitation of high-J CO transitions (Costa et al. 2018). To reach a deeper understanding of the impact of AGN on the molecular gas reservoir, also on longer timescales, predictions from simulations providing the spatial scales and effects of AGN activity on CO properties as a function of cosmic time are needed.","Citation Text":["Brusa et al. 2018"],"Functions Text":["This representative value for our AGN is in line but not as low as previous work targeting extremely powerful sources (e.g., Mgas,\u2006mol\/M*\u2004\u20040.05 in"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1590,1607]],"Functions Start End":[[1443,1589]]}
{"Identifier":"2022MNRAS.512..186K__Datta_et_al._2010_Instance_1","Paragraph":"This work is the second in a series of works aimed at understanding the effect of residual gain errors in different power spectrum estimations in the presence of strong foreground and exploring potential mitigation techniques. In this work, we have not investigated the different possibilities for the presence of residual gain errors and chosen the values of \u03c3R and \u03c3I same as \u03c3g for most of the discussions. In general, the standard deviation of the real and imaginary parts can be different. We observe here that it is essential to assess the time dependence of the gain accurately, as its inaccurate estimation leads to the time-correlated residual gain errors. In this work, we do not consider the effect of frequency correlation in gain error, and all our estimations are done for correlating visibilities in the same frequency channel. Furthermore, this work also uses the foreground subtraction technique, where we expect to have accurate knowledge of the foreground emissions (Jeli\u0107 et al. 2008; Ghosh et al. 2012). An alternative method, more regularly exercised in literature, is foreground avoidance. It has been established that the foregrounds to the redshifted 21-cm emissions remain correlated across relatively larger bandwidth (Platania et al. 1998; Santos et al. 2005; Ali et al. 2008; Jeli\u0107 et al. 2008; Chakraborty et al. 2019b), whereas the H\u2009i signal decorrelates faster (Bharadwaj & Pandey 2003; Bharadwaj & Ali 2005). As a result, when the power spectrum is observed as a function of (k\u2225, k\u22a5), the foreground emission remains concentrated near the low k\u2225, inside the \u2018wedge\u2019 (Datta et al. 2010; Morales et al. 2012; Vedantham et al. 2012). Note that the smaller frequency separation in multifrequency angular power spectrum contributes to larger k\u2225 modes of the power spectrum. Hence, the effect of frequency-independent residual gain we see here at zero frequency separation may contribute to bias in the power spectrum beyond the wedge. Moreover, the antenna-based gains are functions of both time and frequency; the residual gain is expected to have correlated frequency dependence. Such frequency-correlated calibration errors couple the foreground power beyond the foreground wedge into the EoR window region of the 2D power spectrum space (Barry et al. 2016; Ewall-Wice et al. 2017; Byrne et al. 2019; Pal et al. 2021). At present, we are working towards expanding the formalism presented in this paper to estimate bias and variance in power spectrum estimate when visibility correlation in different frequencies is considered. Here, we also consider that the gain errors arising from the different antennae are uncorrelated. Though this is a fairly good assumption for the gain arising from electronics in the antenna system itself, the ionospheric effects may introduce correlated gains across the antenna. Furthermore, as the calibration procedure uses baseline-dependent gains to solve for the antenna dependent gains, calibration errors can lead to correlated residual gain errors across the antenna. Moreover, asymmetry in the telescope aperture, mechanical fatigue of telescope structure, etc., can lead to parts of the gain errors correlated across different antennae and even across different days of observations. We are investigating these effects, and the result will be presented in future work. Though the demonstrations here are done with the uGMRT as a model for the interferometer, a similar analysis can be carried out for any telescope of concern, and a prior assessment of the effect of the gain errors can be made using the analytical expression presented here with minimum computation cost. Furthermore, this work emphasizes the importance of estimating and establishing the gain statistics for a given interferometer. Though we use a simple model for the residual gain error here, the calculations that lead to the analytical expression can be readily expanded for a more complicated gain error model. We believe that this work provides significant direction in understanding and planning observations to detect redshifted 21-cm power spectrum.","Citation Text":["Datta et al. 2010"],"Functions Text":["As a result, when the power spectrum is observed as a function of (k\u2225, k\u22a5), the foreground emission remains concentrated near the low k\u2225, inside the \u2018wedge\u2019"],"Functions Label":["Uses"],"Citation Start End":[[1601,1618]],"Functions Start End":[[1443,1599]]}
{"Identifier":"2021AandA...656A..95P__Houdek_et_al._2017_Instance_1","Paragraph":"Many efforts have thus been devoted to the correction of surface effects, either from theoretical modelling (e.g., Gabriel et al. 1975; Balmforth 1992b; Houdek 1996; Rosenthal et al. 1999; Grigahc\u00e8ne et al. 2005) or through empirical formulae (e.g., Kjeldsen et al. 2008; Christensen-Dalsgaard 2012; Ball & Gizon 2014; Sonoi et al. 2015). Some aspects, however, are very complicated to model, and existing models make use of assumptions that can barely be justified, if at all. For instance, turbulent pressure modulations are usually described in the Gas-\u03931 (GGM) or reduced-\u03931 (RGM) approximations (Rosenthal et al. 1999), which amounts to neglecting the effects of turbulent dissipation and buoyancy on the mode (Belkacem et al. 2021). Another problem is the use of time-dependent mixing-length formalisms (Unno 1967; Gough 1977) to account for modal surface effects (e.g., Gabriel et al. 1975; Houdek 1996; Grigahc\u00e8ne et al. 2005; Sonoi et al. 2017; Houdek et al. 2017, 2019). While useful for the bulk of the convective region, the mixing-length hypothesis is no longer valid in the superadiabatic region just beneath the surface of the star, as shown by 3D hydrodynamic simulations of stellar atmospheres (see Nordlund et al. 2009, for a review). Finally, such formalisms require that the oscillations be separated from the convective motions, thus yielding separate equations. This is done either by assuming a cut-off in wavelength space, with oscillations having much shorter wavelengths than turbulent convection (Grigahc\u00e8ne et al. 2005), or by using 3D hydrodynamic simulations and separating the oscillations from convection though horizontal averaging (Nordlund & Stein 2001). The necessity to separate the equations for oscillations and convection is fundamentally problematic as there is no rigorous way to disentangle the two components, mainly because, in solar-type stars, they have the same characteristic lengths and timescales (Samadi et al. 2015). This is even truer if one wishes to model their mutual coupling.","Citation Text":["Houdek et al. 2017"],"Functions Text":["Another problem is the use of time-dependent mixing-length formalisms","to account for modal surface effects (e.g.,"],"Functions Label":["Background","Background"],"Citation Start End":[[954,972]],"Functions Start End":[[739,808],[833,876]]}
{"Identifier":"2022ApJ...935...49G__Hezaveh_et_al._2016_Instance_1","Paragraph":"Strong gravitational lensing systems are a powerful tool for cosmology. They have been used to study how dark matter is distributed in galaxies and clusters (e.g., Kochanek 1991; Hogg & Blandford 1994; Broadhurst et al. 2000; Koopmans & Treu 2002; Bolton et al. 2006; Koopmans et al. 2006; Brada\u010d et al. 2008; Huang et al. 2009; Vegetti & Koopmans 2009; Jullo et al. 2010; Grillo et al. 2015; Shu et al. 2015, 2016, 2017; Meneghetti et al. 2020) and are uniquely suited to probe the low end of the dark matter mass function and test the prediction of the cold dark matter (CDM) model beyond the local universe (e.g., Vegetti et al. 2010, 2012; Hezaveh et al. 2016; Ritondale et al. 2019; Diaz Rivero & Dvorkin 2020; Ca\u01e7an Seng\u00fcl et al. 2020, 2021; Gilman et al. 2021). Multiply lensed supernovae (SNe) are ideal for measuring time delays and H\n0 because of their well-characterized light curves, and in the case of Type Ia, with the added benefit of standardizable luminosity (Refsdal 1964; Treu 2010; Oguri & Marshall 2010), provided microlensing can be accurately characterized (Yahalomi et al. 2017; Foxley-Marrable et al. 2018). Furthermore, SNe have the benefit of fading, so for these systems lens models can be validated using images that are uncontaminated by bright point sources (Ding et al. 2021). In recent years, strongly lensed SNe, both core-collapse (Kelly et al. 2015; Rodney et al. 2016) and Type Ia (Quimby et al. 2014; Goobar et al. 2017; Rodney et al. 2021), have been discovered. Time-delay H\n0 measurements from multiply imaged SNe (e.g., Goldstein & Nugent 2017; Shu et al. 2018; Goldstein et al. 2018, 2019; Pierel & Rodney 2019; Suyu et al. 2020; Huber et al. 2022), combined with measurements from distance ladders (e.g., Freedman et al. 2019, 2020; Riess et al. 2019, 2022) and lensed quasars (e.g., Suyu et al. 2010, 2013; Treu & Marshall 2016; Bonvin et al. 2017; Birrer et al. 2020; Millon et al. 2020; Wong et al. 2020), can be an important test of the tension between H\n0 measured locally and the value inferred from the cosmic microwave background (CMB; Planck Collaboration et al. 2020).","Citation Text":["Hezaveh et al. 2016"],"Functions Text":["Strong gravitational lensing systems are a powerful tool for cosmology. They have been used to study how dark matter is distributed in galaxies and clusters","and are uniquely suited to probe the low end of the dark matter mass function and test the prediction of the cold dark matter (CDM) model beyond the local universe (e.g.,"],"Functions Label":["Background","Background"],"Citation Start End":[[644,663]],"Functions Start End":[[0,156],[446,616]]}
{"Identifier":"2020ApJ...901L..10L__White_et_al._1992_Instance_1","Paragraph":"The 17 GHz polarization reversal during this CRF can be a yet unknown feature inherent to the fan\u2013spine structure, where magnetic polarity around null point (NP) varies so rapidly as to affect the propagation of microwave polarization. A way to possibly explain this polarization change is to view it as a mode-coupling phenomenon, the process by which the rays reverse their original sense of polarization while passing through a quasi-transverse field region along the line of sight from the radiation source to the observer, depending on the degree of mode coupling there (Cohen 1960; Zheleznyakov 1970; Melrose 1975; White et al. 1992). This is an attractive scenario for a fan\u2013spine structure, because the fan surface may well act as a quasi-transverse layer for the rays emitted underneath. To think about an ideal fan\u2013spine structure with a flux rope inside, in this configuration, the magnetic fields above the fan surface are all in the negative magnetic polarity, and the rays emitted from either magnetic polarity underneath will be observed as LHCP everywhere. Therefore, the LHCP observed everywhere before the flare can simply be due to the fan\u2013spine structure, without any strong mode-coupling phenomenon. On the other hand, if a magnetic flux rope rises to reconnect with the overlying fan field, the fan surface may partially open up to let the flux rope erupt out. Such a change of magnetic field structure can explain the instant reversal of the 17 GHz polarization at t3 more naturally. The reconnection between the magnetic fields inside and outside of the fan will occur across a current sheet, the so-called breakout current sheet (BCS), and the newly open field lines amount to the lower part of the rising and expanding BCS (see, e.g., Lynch et al. 2016; Karpen et al. 2017). A sustained BCS over the active region might affect the microwave polarization, as mode coupling across a current sheet is still a debatable issue (Zheleznyakov et al. 1996; Lee et al. 1998; Lee 2007). We here offer only the simplest interpretation, according to which the change from LHCP to RHCP of the 17 GHz emission over the inner ribbon is not just a signature for any magnetic field perturbation, but may indicate a specific form of a breakout eruption out of the closed fan structure. The implied magnetic field reconfiguration is in line with the recently reported decay of the coronal magnetic field at the flare site by Fleishman et al. (2020).","Citation Text":["White et al. 1992"],"Functions Text":["A way to possibly explain this polarization change is to view it as a mode-coupling phenomenon, the process by which the rays reverse their original sense of polarization while passing through a quasi-transverse field region along the line of sight from the radiation source to the observer, depending on the degree of mode coupling there"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[621,638]],"Functions Start End":[[236,574]]}
{"Identifier":"2017MNRAS.464.4495G__Levine_et_al._2006_Instance_1","Paragraph":"To emphasize it again, one may suggest accordingly that several low-m modes of moderately growing collective oscillation with different amplitudes, number of spirals and pitch angles (radial wavelengths), including the most interesting m = 1 mode, may co-exist in the solar neighbourhood. The so-called lopsidedness, or the m = 1 azimuthal asymmetry, is often seen in the distribution of stars and gas in the outer discs of many disc galaxies (Zaritsky & Rix 1997; van Eymeren et al. 2011). Generally, there is no azimuthal symmetry, at least far from their centres, in multi-arm galaxies (Efremov 2011). As for us the spiral pattern of the Galaxy is not a single m model but a superposition of modes as suggested theoretically in the late 1970s (Bertin et al. 1977; Lin & Lau 1979; Bertin 1980; see also Griv & Wang 2014). The overall pattern is thus basically multiple armed. In addition, we argue that the Galaxy is the azimuthally asymmetric spiral system (cf. Levine et al. 2006). The system therefore may exhibit a complicated asymmetric, multi-arm, not well-defined spiral structure as seen, for instance, in the face-on spiral galaxy M101 (Fig. 13). Interestingly, many spiral structures in galaxies do not appear to be well-organized grand designs. Galaxies dominated by a single and symmetric pattern are exceedingly rare.\n\nThe pitch angle of the spiral pattern seems to be relatively small for all models considered, |tan\u2009p| \u226a 1. Thus, the original Lin\u2013Shu approximation of tightly wound gravity perturbations used throughout the theory does not fail.\nThe relative amplitude of the surface density $\\widetilde{\\Sigma }_0\/\\Sigma _{\\mathrm{basic}} \\gtrsim 1$ for all models considered. The latter means that the non-axisymmetric variation found in the relative amplitude of the surface density does not represent a small perturbation in the basic equilibrium state of the Galaxy that is axisymmetric in the mean.\nThe above mentioned results are in agreement with our previous determinations (Griv et al. 2013, 2014, 2015a,b,c).\n","Citation Text":["Levine et al. 2006"],"Functions Text":["In addition, we argue that the Galaxy is the azimuthally asymmetric spiral system (cf."],"Functions Label":["Differences"],"Citation Start End":[[965,983]],"Functions Start End":[[878,964]]}
{"Identifier":"2017MNRAS.469.3252P__Anderson_&_Bedin_2010_Instance_1","Paragraph":"Tol1247 was imaged with the HST (see Fig. 3) in the optical using the Wide Field Ultraviolet-Visible Channel (UVIS) of its Wide Field Camera 3 (WFC3). For UV imaging, the Advanced Camera for Surveys' Solar Blind Channel (ACS\/SBC) was used. Seven filters were utilized in total, allowing to apply laxs \u2013 the Lyman alpha eXtraction software (Hayes et al. 2009) \u2013 to produce continuum-subtracted Ly\u03b1, H\u03b1 and H\u03b2 images, corrected for underlying stellar absorption and contamination from [N\u2009ii] \u03bb6548, 6584. The latter one is based on the spectroscopic line ratio ${[\\mathrm{N}\\,\\small {II}]}\/{\\mathrm{H}\\alpha }=0.0605$ published by Terlevich et al. (1993). The imaging strategy and data reduction methodology for Tol1247 is very similar to that of the Lyman Alpha Reference Sample (LARS; Hayes et al. 2014; \u00d6stlin et al. 2014; Duval et al. 2015) and the basic data reduction for this data set is done in the same way as for LARS. Flat-field-corrected frames were obtained from the Mikulski Archive for Space Telescope. The charge transfer inefficiency (CTE) correction for the ACS data was performed by the pipeline, whereas CTE losses in WFC3\/UVIS (Anderson & Bedin 2010) were treated manually using the tools supplied by STScI.3 We then stacked the individual data frames and drizzled them to a pixel scale of 0.04 arcsec pixel\u22121 using DrizzlePac version 1.1.16 (Gonzaga et al. 2012). Further pre-processing of the data includes additional masking of cosmic rays in the drizzled frames and matching the point spread functions (PSFs) for the different filters. In order to match the PSF, we first construct PSF models for all of the filters used in the study. For the optical filters, we use TinyTim models (Krist, Hook & Stoehr 2011), re-sampled to a pixel scale of 0.04 arcsec pixel\u22121. However, for the FUV filters, TinyTim is not accurate enough, particularly in the wings. Therefore, the PSF models for the FUV filters are instead built from stacks of stars obtained in calibration observations; see Hayes et al. (2016) for details. All of the PSF models are then normalized by peak flux and stacked by maximum pixel value. We then proceed to calculate convolution kernels that match the PSFs for all of the filters to the maximum width model. Each kernel is built up from a set of delta functions and we find the optimum matching kernel by least-squares optimization; see also Becker et al. (2012) and Melinder et al. (in preparation). The drizzled and registered images are convolved with the kernel found for each filter, which result in a set of images matched to a common PSF.","Citation Text":["Anderson & Bedin 2010"],"Functions Text":["The charge transfer inefficiency (CTE) correction for the ACS data was performed by the pipeline, whereas CTE losses in WFC3\/UVIS","were treated manually using the tools supplied by STScI."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1147,1168]],"Functions Start End":[[1016,1145],[1170,1226]]}
{"Identifier":"2020ApJ...902...77O__Wuyts_et_al._2012_Instance_1","Paragraph":"Guo et al. (2015) used an automated \u201cblob finder\u201d to identify star-forming regions in the HST\/ACS images for 3,239 log M*\/M\u2299  10.6 galaxies in CANDELS (GOODS-S and UDS fields) at 0.5  z  3. They defined clumps as blobs which contribute more than 8% of the total UV light of their host galaxies. In contrast to our results, they find that a much higher fraction of SFGs are clumpy (as much as \u223c60% at z \u223c 2), and also that higher mass bins have lower clumpy fractions. It is worth pointing out that clumpy galaxy fractions are highly sensitive to methodology and clump definition and vary widely in the literature (e.g., Ravindranath et al. 2006; Elmegreen et al. 2007; Wuyts et al. 2012; Guo et al. 2015), so we should not expect complete agreement. Comparing their observed clumpy galaxy fractions (as a function of redshift) to fractions derived from the K15 classification scheme, Guo et al. (2015) find that their results agree best with clumpy fractions derived using both the clumpiness and patchiness flags from the K15 data release (see their Appendix A) rather than either the clumpy or patchy flags alone. This is because the blob finder does not account for the light concentration of the blobs. The inclusion of patches may help explain why their clumpy fractions are generally higher than ours. Guo et al. (2015) also exclude very small (01) and elongated (axis ratio 0.5) galaxies from their sample. Our inclusion of such galaxies could easily lead to lower clumpy fractions given that many galaxies with half-light radius 01 are Spheroids (see Figure 13), which rarely possess clumps. We also include edge-on disks whose clumps may be obscured by dust. The inclusion of galaxies with unresolved or obscured clumps may imply that we are underestimating the clumpy fractions. However, we do include the contribution from non-disky compact or irregular SFGs which would be excluded by the Guo et al. (2015) cuts. The contribution from such galaxies is not insignificant, especially at z \u223c 2, so our looser selection is not without merit. Even if we do underestimate our clumpy fractions, our consistent sample selection and methodology ensures that they should be similarly underestimated at all redshifts, preserving the general evolutionary trends.","Citation Text":["Wuyts et al. 2012"],"Functions Text":["t is worth pointing out that clumpy galaxy fractions are highly sensitive to methodology and clump definition and vary widely in the literature (e.g.,","so we should not expect complete agreement."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[669,686]],"Functions Start End":[[469,619],[706,749]]}
{"Identifier":"2016AandA...588A..48C__Hur\u00e9_(2013)_Instance_1","Paragraph":"The major difference between the axisymmetric case and the asymmetric approach is that we need to derive the 3D, asymmetric gravitational potential of luminous baryons beforehand. We thus computed the potential in Cartesian coordinates (x,y,z), which enables us to derive both radial and azimuthal forces at any desired z. This can be derived independently for each stellar or gaseous contribution. The gravitationtal potential \u03a6 of the mass distribution is in principle deduced from the convolution of the volume mass density by the Green function, namely (11)\\begin{equation} \\Phi = -G \\iiint\\frac{{\\rm d}\\rho'}{|\\vec{r}-\\vec{r}'|}, \\label{eq:pot} \\end{equation}\u03a6=\u2212G$d\u03c1\u2032|r\u2212r\u2032|,where G is the gravitational constant. However, the Green function written by 1\/ | r\u2212r\u2032 |  =  [(x\u2212x\u2032)2 + (y\u2212y\u2032)2 + (z\u2212z\u2032)2] \u2212 1\/2, is well known to diverge at each point where x = x\u2032, y = y\u2032 and z = z\u2032. This function renders any direct estimate of \u03a6 inaccurate and generally encourages modelers to incorporate a softening length to bypass the divergence. Here, we use the new formalism presented in Hur\u00e9 (2013) who showed that the Newtonian potential is exactly reproduced by using an intermediate scalar function \u210b, namely \\hbox{$\\Phi = \\partial_{xy}^2 {\\cal H}$}\u03a6=\u2202xy2\u210b. In 3D, this hyperpotential is written as (12)\\begin{equation} {\\cal H}(x,y,z) = -G \\iiint_{\\Omega'}{\\rho(x',y',z')} \\kappa^{xy}(X,Y,Z){\\rm d}x'{\\rm d}y'{\\rm d}z' , \\end{equation}\u210b(x,y,z)=\u2212G$\u03a9\u2032\u03c1(x\u2032,y\u2032,z\u2032)\u03baxy(X,Y,Z)dx\u2032dy\u2032dz\u2032,with X = x\u2212x\u2032, Y = y\u2212y\u2032 and Z = z\u2212z\u2032. The \u03ba function is a hyperkernel defined by (13)\\begin{equation} \\kappa^{xy}(X,Y,Z) = -Z \\arctan \\frac{XY}{Z |\\vec{r}-\\vec{r}'|}+Y \\ln \\frac{X+|\\vec{r}-\\vec{r}'|}{\\sqrt{Y^2+Z^2}}\\cdot \\label{eq:k} \\end{equation}\u03baxy(X,Y,Z)=\u2212ZarctanXYZ|r\u2212r\u2032|+YlnX+|r\u2212r\u2032|Y2+Z2\u00b7This approach is particularly simple and efficient for 2D or 3D distributions since \u210b is, in contrast to \u03a6, the convolution of the surface or volume density with a regular, finite amplitude kernel. The methodology thus does not make use of a softening length in the derivation of the potential. In practice, this convolution is performed using the second-order trapezoidal rule and the mixed derivatives are estimated at the same order from centered finite differences. Furthermore, the volume density of the tracer are deduced from a surface density map, considering that the vertical density follows a sech-squared or exponential law of constant scaleheight with radius. The precision of these schemes is sufficient for the present purpose. ","Citation Text":["Hur\u00e9 (2013)"],"Functions Text":["Here, we use the new formalism presented in","who showed that the Newtonian potential is exactly reproduced by using an intermediate scalar function \u210b, namely \\hbox{$\\Phi = \\partial_{xy}^2 {\\cal H}$}\u03a6=\u2202xy2\u210b. In 3D, this hyperpotential is written as (12)\\begin{equation} {\\cal H}(x,y,z) = -G \\iiint_{\\Omega'}{\\rho(x',y',z')} \\kappa^{xy}(X,Y,Z){\\rm d}x'{\\rm d}y'{\\rm d}z' , \\end{equation}\u210b(x,y,z)=\u2212G$\u03a9\u2032\u03c1(x\u2032,y\u2032,z\u2032)\u03baxy(X,Y,Z)dx\u2032dy\u2032dz\u2032,with X = x\u2212x\u2032, Y = y\u2212y\u2032 and Z = z\u2212z\u2032."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1078,1089]],"Functions Start End":[[1034,1077],[1090,1512]]}
{"Identifier":"2015ApJ...804..130C___2013_Instance_2","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)"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[1347,1370]],"Functions Start End":[[1140,1346]]}
{"Identifier":"2016AandA...587A.159G__Segura_et_al._2007_Instance_1","Paragraph":"One has to be sure to rule out cases where inorganic chemistry can mimic the presence of life (\u201cfalse positives\u201d). Potential abiotic ozone production on Venus- and Mars-like planets has been discussed by Schindler & Kasting (2000, and references therein). While this is based on photolysis of e.g., CO2 and H2O and is thus limited in extent, a sustainable production of abiotic O3 which could build up to a detectable level has been suggested by Domagal-Goldman & Meadows (2010) for a planet within the habitable zone of AD Leonis with a specific atmospheric composition. Indeed, other studies confirm that abiotic buildup of ozone is possible (e.g., Hu et al. 2012; Tian et al. 2014); however, detectable levels are unlikely if liquid water is abundant, as e.g. rainout of oxidized species would keep atmospheric O2 and O3 low (Segura et al. 2007), unless the CO2 concentration is high and both H2 and CH4 emissions are low (Hu et al. 2012). False-positive detection of molecules such as CH4 and O3 is discussed by von Paris et al. (2011). Seager et al. (2013) present a biosignature gas classification. Since abiotic processes cannot be ruled out for individual molecules (e.g. for O3), searches for biosignature molecules should search for multiple biosignature species simultaneously. It has been suggested that the simultaneous presence of O2 and CH4 can be used as an indication for life (Sagan et al. 1993, and references therein). Similarly, Selsis et al. (2002) suggest a so-called \u201ctriple signature\u201d, where the combined detection of O3, CO2 and H2O would indicate biological activity. Domagal-Goldman & Meadows (2010) suggest to simultaneously search for the signature of O2, CH4, and C2H6. Of course, care has to be taken to avoid combinations of biosignature molecules which can be generated abiotically together (see e.g. Tian et al. 2014). The detectability of biosignature molecules is discussed, e.g. by von Paris et al. (2011) and Hedelt et al. (2013). In particular, the simulation of the instrumental response to simulated spectra for currently planned or proposed exoplanet characterization missions has shown that the amount of information the retrieval process can provide on the atmospheric composition may not be sufficient (von Paris et al. 2013). ","Citation Text":["Segura et al. 2007"],"Functions Text":["Indeed, other studies confirm that abiotic buildup of ozone is possible","however, detectable levels are unlikely if liquid water is abundant, as e.g. rainout of oxidized species would keep atmospheric O2 and O3 low",", unless the CO2 concentration is high and both H2 and CH4 emissions are low"],"Functions Label":["Background","Background","Background"],"Citation Start End":[[829,847]],"Functions Start End":[[572,643],[686,827],[848,924]]}
{"Identifier":"2022ApJ...931...70B___2022a_Instance_1","Paragraph":"RFs can propagate from the magnetotail to Earth over a long distance more than 10 R\nE together with BBFs behind them (Runov et al. 2009; Cao et al. 2010). Studies have suggested that RFs are crucial regions for particle acceleration, pitch-angle evolution, wave\u2013particle interactions, and electromagnetic energy conversion during their Earthward propagation. For instance, rapid increases in energy fluxes of electrons and ions from tens to hundreds of keV are a typical feature of RF events (Khotyaintsev et al. 2011; Liu et al. 2013, 2018c, 2021a, 2022b; Zhou et al. 2018; Liu & Fu 2019; Gabrielse et al. 2021), pitch-angle distribution of suprathermal electrons can evolve dramatically around RFs (Runov et al. 2013; Liu et al. 2020), strong particle and wave activity can occur in the vicinity of RFs (Ono et al. 2009; Zhou et al. 2009, 2014; Fu et al. 2014; Breuillard et al. 2016; Greco et al. 2017; Yang et al. 2017), and RFs are associated with energy conversion from electromagnetic fields to particles (Sitnov et al. 2009; Huang et al. 2015; Khotyaintsev et al. 2017; Liu et al. 2018a, 2022a). The energetic plasma in the vicinity of RFs plays a key role in connecting the magnetotail with the inner magnetosphere because they carry a large amount of energy and can be injected into the inner magnetosphere to affect the ring current and radiation belt (Gabrielse et al. 2012; Duan et al. 2014; Turner et al. 2014). Possible mechanisms responsible for the energization of particles around RFs have been widely investigated based on both spacecraft observations and numerical simulations during the past decade. The strong convection electric field induced by the strong magnetic field gradient of RFs provides significant adiabatic acceleration of the ambient particles (Birn et al. 2004, 2013, 2015; Gabrielse et al. 2012, 2014, 2016; Ganushkina et al. 2013; Liu et al. 2016; Turner et al. 2016). Nonadiabatic effects, caused by particle reflection ahead of the RFs (Zhou et al. 2018), resonance with RFs (Ukhorskiy et al. 2013, 2017), and scattering by wave emissions (Zhou et al. 2009; Greco et al. 2017), are also significant for particle energization. These above studies usually assumed that the RF surface has a planar boundary at a typical thickness comparable to the ion gyroradius and below (Nakamura et al. 2002; Sergeev et al. 2009; Zhou et al. 2009; Schmid et al. 2011; Liu et al. 2013; Vapirev et al. 2013). Divin et al. (2015b) revealed that the RF surface is unstable to instabilities ranging from electron scales to ion scales. Simulation studies found that RFs can be unstable to interchange instability and that finger-like structures on ion\u2013electron hybrid scales can develop at the RF (Vapirev et al. 2013). Such finger-like structures are found to play a role in modulating the electron acceleration process (Wu et al. 2018). Bai et al. (2022) also reported significant ion trapping acceleration at the RF with ion-scale ripples. Unlike these surface structures with ion or ion\u2013electron hybrid scales, Liu et al. (2018b) recently reported that the RF layer has electron-scale density gradients, currents, and electric fields, based on the MMS mission, which consists of four spacecraft separated by 30 km. Such electron-scale ripple structure can be generated by lower hybrid drift instability (Divin et al. 2015b; Pan et al. 2018). Liu et al. (2021c) presented a detailed investigation of energy flux densities at two RFs with\/without the electron-scale surface ripples and indicated that surface ripples may play an important role in the particle dynamics. But how such electron-scale RF structure impacts the electron energization and transport still remains unknown. In this paper, with the aid of observation-based test-particle simulation, we aim to investigate in detail the effect of the front surface ripples on the local electron dynamics.","Citation Text":["Liu et al.","2022a"],"Functions Text":["For instance,","and RFs are associated with energy conversion from electromagnetic fields to particles"],"Functions Label":["Background","Background"],"Citation Start End":[[1078,1088],[1096,1101]],"Functions Start End":[[359,372],[925,1011]]}
{"Identifier":"2016ApJ...817....9K__Chen_et_al._2006_Instance_2","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"],"Functions Text":["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"],"Functions Label":["Uses"],"Citation Start End":[[945,961]],"Functions Start End":[[760,944]]}
{"Identifier":"2016ApJ...820...12P__Camenzind_&_Krockenberger_1992_Instance_1","Paragraph":"Variability in observed emission can be considered a defining characteristic of active galactic nuclei (AGNs), and for the roughly 10% of AGNs that are radio-loud (e.g., Jiang et al. 2007) the majority of this variable emission is understood to arise from the relativistic flows of plasma along two oppositely directed jets (e.g., Urry & Padovani 1995). When viewed at small angles to the jet direction the Doppler boosting makes the emission from the approaching jet appear dramatically brighter and also shortens the observed timescales with respect to those in the emitted frame, thereby explaining many of the properties of blazars (e.g., Blandford & Rees 1978; Lister 2001; Gopal-Krishna et al. 2003). Multiband radio studies and very long baseline interferometry (VLBI), coupled with theoretical models, have provided extremely strong evidence for the presence of both moving and standing shocks in these jets (e.g., Marscher & Gear 1985; Hughes et al. 1991; Lister et al. 2001, 2009; Marscher et al. 2008, 2010), indicating significant changes in the fluid flow and\/or the density of material ejected into the jets; it is now accepted that the largest flares arise from the production and relativistic propagation of new components seen as radio knots. Changes in the overall direction of the inner portions of the jet, or at least its brightest portions, have also been demonstrated via VLBI (e.g., Biretta et al. 1986; Ros et al. 2000; Piner et al. 2003; Caproni & Abraham 2004). Even modest changes in direction (e.g., Camenzind & Krockenberger 1992; Gopal-Krishna & Wiita 1992) have long been recognized as one way to produce significant changes in the flux and polarization (e.g., Gopal-Krishna & Wiita 1993; Piner et al. 2008). It is to be expected that turbulence is produced within these jets, at least in the vicinity of shocks, and thus some of the variations should arise from such smaller scale motions as has been suggested theoretically (Marscher & Travis 1991; Marscher 2014, Calafut & Wiita 2015) and strongly supported by observations of blazars (e.g., Marscher et al. 2008, 2010; Bhatta et al. 2013). In addition, there is the possibility that portions of the jet are moving much faster than other portions, and such misaligned mini-jets could also produce some extremely rapid fluctuations (Giannios et al. 2009; Biteau & Giebels 2012). Variability on a wide range of timescales can be produced within the accretion disks feeding the central black holes (e.g., Czerny 2006). This presumably dominates the variations from radio-quiet AGNs though not those of radio-loud ones because the special relativistic boosting of the jet emission is so important in the latter (Urry & Padovani 1995). Some of the variations in the jet emission might be traced to plasma fluctuations in the disk being advected into the base of the jets (e.g., Wiita 2006) but the exact origins of the initial fluctuations are not addressed in this work. Here we address the question of whether variations in the bulk flow of a propagating relativistic hydrodynamic (RHD) jet along with sub-grid mildly relativistic turbulence can produce light curves and power spectra resembling those of radio-loud AGNs.","Citation Text":["Camenzind & Krockenberger 1992"],"Functions Text":["Even modest changes in direction (e.g.,","have long been recognized as one way to produce significant changes in the flux and polarization"],"Functions Label":["Background","Background"],"Citation Start End":[[1529,1559]],"Functions Start End":[[1489,1528],[1589,1685]]}
{"Identifier":"2017MNRAS.465..492M__Taverna_et_al._2015_Instance_2","Paragraph":"Given the quite strong surface magnetic field of the M7, thermal radiation is expected to be polarized, either if emission is from a bare surface or from an atmosphere (see Turolla et al. 2004; Potekhin 2014). The polarization properties are quite different in the two cases, although there are still uncertainties, especially at optical\/ultraviolet (UV) wavelengths. One of the first predictions of quantum electrodynamics (QED), even before it was properly formulated, was vacuum birefringence, and, in particular, that a strong magnetic field affects the propagation of light through it (Heisenberg & Euler 1936; Weisskopf 1936). In thermally emitting INSs, radiation comes from a region comparable with the entire star surface, over which the magnetic field direction changes substantially. In the absence of QED vacuum polarization effects, this would produce a drastic depolarization of the radiation collected at infinity (Heyl, Shaviv & Lloyd 2003, see also Taverna et al. 2015; Gonz\u00e1lez Caniulef et al. 2016 and references therein). Vacuum birefringence dramatically increases the linear polarization of the observed radiation, from a level of a few \u2009per\u2009cent up to even \u223c100\u2009per\u2009cent, depending on the viewing geometry and the surface emission mechanism (Heyl & Shaviv 2000, 2002; Heyl et al. 2003; Taverna et al. 2015; Gonz\u00e1lez Caniulef et al. 2016). Detecting polarization in the thermal emission from the surface of an INS will be therefore extremely valuable. First, and independently on the physical conditions of the emitting region, the detection of a large degree of linear polarization in the signal would constitute the observational evidence of QED vacuum polarization effects in the strong-field regime. Secondly, the polarization observables can be compared with emission models and help to uncover the physical conditions of INS surfaces and atmospheres, ideally complementing spectral observations (Taverna et al. 2014; Gonz\u00e1lez Caniulef et al. 2016).","Citation Text":["Taverna et al. 2015"],"Functions Text":["Vacuum birefringence dramatically increases the linear polarization of the observed radiation, from a level of a few \u2009per\u2009cent up to even \u223c100\u2009per\u2009cent, depending on the viewing geometry and the surface emission mechanism"],"Functions Label":["Background"],"Citation Start End":[[1309,1328]],"Functions Start End":[[1042,1263]]}
{"Identifier":"2016AandA...594A..64P__Judge_(2015)_Instance_1","Paragraph":"There is now renewed interest in the literature concerning these transitions, because some of the O\u2009iv and S\u2009iv intercombination lines, together with the Si\u2009iv resonance lines, are routinely observed with the Interface Region Imaging Spectrograph (IRIS; De Pontieu et al. 2014) at much higher spectral, spatial and temporal resolution than previously. For example, Peter et al. (2014) used the intensities of the O\u2009iv vs. Si\u2009iv lines to propose that very high densities, on the order of 1013 cm-3 or higher, are present in the so-called IRIS plasma \u201cbombs\u201d. Line ratios involving an O\u2009iv forbidden transition and a Si\u2009iv allowed transition have been used in the past to provide electron densities during solar flares and transient brightenings (e.g. Cheng et al. 1981; Hanssen 1981). However, the validity of using O\u2009iv to Si\u2009iv ratios has been hotly debated because these ratios gave very high densities compared to the more reliable ones obtained from the O\u2009iv ratios alone (see, e.g. Hayes & Shine 1987). In addition, Judge (2015) recalled several issues that should be taken into account when considering the Si\u2009iv\/O\u2009iv density diagnostic. The main ones were: (1) O\u2009iv and Si\u2009iv ions are formed at quite different temperatures in equilibrium and hence a change in the O\u2009iv\/Si\u2009iv ratio could imply a change in the temperature rather than in the plasma density (2) the chemical abundances of O and Si are not known with any great accuracy and could be varying during the observed events (3) density effects on the ion populations could increase the Si\u2009iv\/O\u2009iv relative intensities by a factor of roughly three to four. Judge (2015) has also mentioned the well known problem of the \u201canomalous ions\u201d, that is, the observed high intensities of the Li- and Na-like (as Si\u2009iv) ions (see also Del Zanna et al. 2002). Another important aspect to take into account is the effect of non-equilibrium conditions on the observed plasma diagnostics. It is well known that strong variations in the line intensities are obtained when non-equilibrium ionisation is included in the numerical calculations (see, e.g. Shen et al. 2013; Raymond & Dupree 1978; Mewe & Schrijver 1980; Bradshaw et al. 2004). In particular, Doyle et al. (2013) and Olluri et al. (2013) investigated the consequences of time-dependent ionization on the formation of the O\u2009iv and Si\u2009iv transition region lines observed by IRIS. In addition, Dud\u00edk et al. (2014) showed that non-Maxwellian electron distributions in the plasma can substantially affect the formation temperatures and intensity ratios of the IRIS Si\u2009iv and O\u2009iv lines. These authors also suggested that the observing window used by IRIS should be extended to include S\u2009iv. Recent IRIS observation sequences have indeed included the S\u2009iv line near 1406 \u00c5. The S\u2009iv line ratios have a higher limit for density sensitivity than the O\u2009iv line ratios and are thus particularly useful for diagnosing high densities which might occur in flares. Previous flare studies have in fact reported line ratios involving O ions which lay above the density sensitivity range, indicating an electron density in the excess of 1012 cm-3 (e.g. Cook et al. 1995; Polito et al. 2016). ","Citation Text":["Judge (2015)"],"Functions Text":["In addition,","recalled several issues that should be taken into account when considering the Si\u2009iv\/O\u2009iv density diagnostic. The main ones were: (1) O\u2009iv and Si\u2009iv ions are formed at quite different temperatures in equilibrium and hence a change in the O\u2009iv\/Si\u2009iv ratio could imply a change in the temperature rather than in the plasma density (2) the chemical abundances of O and Si are not known with any great accuracy and could be varying during the observed events (3) density effects on the ion populations could increase the Si\u2009iv\/O\u2009iv relative intensities by a factor of roughly three to four."],"Functions Label":["Background","Background"],"Citation Start End":[[1021,1033]],"Functions Start End":[[1008,1020],[1034,1620]]}
{"Identifier":"2016MNRAS.458.3655V__Casella,_Belloni_&_Stella_2005_Instance_1","Paragraph":"Accreting stellar-mass black holes in binary systems regularly display quasi-periodic oscillations (QPOs) in their X-ray flux with frequencies drifting from \u223c0.1\u201310 Hz (e.g. Van der Klis 1989). Three main components can be identified in the spectrum of these sources: disc blackbody emission, power-law emission from the inner accretion flow, and a reflection spectrum from photons reflected off the disc (Done, Gierlinski & Kubota 2007). So-called Type-C low-frequency QPOs (Casella, Belloni & Stella 2005) are believed to originate from the inner accretion flow\/corona that is associated with the Comptonized power-law component of the X-ray spectrum, as this component shows a much larger variability amplitude than the blackbody disc component (Sobolewska & Zycki 2006; Axelsson, Hjalmarsdotter & C. 2013). Since currently no consensus on the origin of QPOs exists, we can generally divide QPO models into two broad categories: geometric and intrinsic models. In the former, the X-ray emission is constant but an oscillating accretion geometry quasi-periodically alters the observed flux. A possible origin for these geometric oscillations could be Lense\u2013Thirring precession of the Comptonizing medium, due to misalignment of the black hole spin and the binary orbit (Stella & Vietri 1997; Stella, Vietri & Morsink 1999; Ingram, Done & Fragile 2009). Alternatively, in intrinsic models the emitted luminosity itself varies, for example due to changes in mass accretion rate (Tagger & Pellat 1999; Cabanac et al. 2010) or due to a standing shock in the accretion flow (Chakrabarti & Molteni 1993). Recently, Heil, Uttley & Klein-Wolt (2015) and Motta et al. (2015) confirmed that the QPO amplitude depends on the inclination of the binary orbit, strongly suggesting a geometric origin (Schnittman, Homan & Miller 2006). Ingram & Van der Klis (2015) found that the iron line equivalent width changes over a QPO cycle in GRS 1915+105, also strongly pointing towards a geometric origin.","Citation Text":["Casella, Belloni & Stella 2005"],"Functions Text":["So-called Type-C low-frequency QPOs","are believed to originate from the inner accretion flow\/corona that is associated with the Comptonized power-law component of the X-ray spectrum, as this component shows a much larger variability amplitude than the blackbody disc component"],"Functions Label":["Background","Background"],"Citation Start End":[[476,506]],"Functions Start End":[[439,474],[508,747]]}
{"Identifier":"2020ApJ...894..121M__Bruno_&_Carbone_2013_Instance_1","Paragraph":"Current sheet drift is implemented following the approach proposed by Burger (2012) to calculate the drift velocity Vd in Equation (7). The heliospheric current sheet angle is modeled as by K\u00f3ta & Jokipii (1983), but now with a time-dependent tilt angle from Equation (8) such that\n12\n\n\n\n\n\nwhere, after Burger et al. (2008), we employ\n13\n\n\n\n\n\nwith \u03d50 = 0\u00b0, thereby allowing for the inclusion in the model of the effects of a fully time-dependent heliospheric current sheet. Diffusion coefficients are modeled as by Moloto et al. (2018). A composite slab\/2D model for transverse magnetostatic turbulence is assumed (see, e.g., Matthaeus et al. 1995), as well as slab\/2D turbulence power spectra with wavenumber-independent energy-containing ranges, and Kolmogorov inertial ranges. This latter assumption does not perfectly reflect spacecraft observations of the same (see, e.g., Bieber et al. 1993; Matthaeus et al. 2007; Bruno & Carbone 2013), but leads to relatively simple, tractable expressions for the parallel and perpendicular MFPs. For more detail on diffusion coefficients derived assuming more realistic forms for the turbulence power spectra, see Shalchi et al. (2010), Engelbrecht & Burger (2015b), and Strauss et al. (2016). In brief, then, the parallel MFP expression used here is that constructed by Burger et al. (2008) from the quasilinear theory results presented by Teufel & Schlickeiser (2003), and given by\n14\n\n\n\n\n\nwhere s = 5\/3, R = RLkm, and km = 1\/\u03bbsl the wavenumber at which the slab spectrum inertial range commences, and RL is the maximal proton Larmor radius. To model the perpendicular MFP, we use the Nonlinear Guiding Center result of Shalchi et al. (2004), as modified by Burger et al. (2008) to take into account a general ratio of slab to 2D energies (see, e.g., Bieber et al. 1994, 1996). This expression is similar to that derived by Zank et al. (2004), and is given by\n15\n\n\n\n\n\nwhere \u03bd = 5\/6 denotes half the Kolmogorov inertial range spectral index, and we assume that \u03b12 = 1\/3 (Matthaeus et al. 2003). Drift coefficients, reduced in the presence of turbulence, are modeled following the approach of Engelbrecht et al. (2017), so that the lengthscale corresponding to the drift coefficient is given by\n16\n\n\n\n\n\nwith \n\n\n\n\n\n being the total magnetic variance. This expression is chosen as it provides results in reasonable agreement with numerical test particle simulations done by Minnie et al. (2007) and Tautz & Shalchi (2012) for the range of turbulence conditions expected in the supersonic solar wind, and due to the fact that its use in a CR modulation code has been shown by Moloto et al. (2018) to lead to computed CR intensities in reasonable agreement with spacecraft observations.","Citation Text":["Bruno & Carbone 2013"],"Functions Text":["This latter assumption does not perfectly reflect spacecraft observations of the same",", but leads to relatively simple, tractable expressions for the parallel and perpendicular MFPs."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[921,941]],"Functions Start End":[[780,865],[942,1038]]}
{"Identifier":"2022MNRAS.515L..39Z__Weiss_et_al._2008_Instance_1","Paragraph":"Rocky planets are the only harbour for life form in the Solar system, so unravelling their origin and history is fundamental for understanding the habitability of planets other than Earth (Lineweaver & Chopra 2012; Cockell et al. 2016). For example, mantle\u2013crust differentiation on Earth has set boundary conditions through redox conditions and early degassing processes that liquid water can occur on the surface of Earth. Geochemical studies, mainly elemental and isotopic compositions, on the specimens from these planets provide significant information about differentiation processes and their timing. Sample-return missions represent one way to obtain these specimens from the differentiated planets (e.g. Anand et al. 2020), but at great time and expense. Non-chondrite meteorites (including achondrites) also originate from differentiated asteroids and planets, e.g. from Moon, Mars, and Vesta, and the angrite and ureilite parent bodies (Binzel & Xu 1993; Weiss et al. 2008; Agee et al. 2013; Marchi et al. 2013; Bischoff et al. 2014). Some non-chondritic meteorites have a unique mineralogy and bulk composition, indicative of core, mantle, and crustal domains of their parent bodies, and thus, these samples record large-scale early planetary differentiation events. For instance, ureilites (Mg-rich, dominated by olivine and pyroxene) and iron meteorites (Fe\u2013Ni metal) are from the mantle and core of asteroids, respectively, and record planetary mantle differentiation and core formation (Goodrich, Scott & Fioretti 2004; Goldstein, Scott & Chabot 2009). In contrast, shergottite and howardite\u2013eucrite\u2013diogenite meteorites, which are inferred to have derived from Mars, 4 Vesta, and related bodies, reflect a variety of crustal compositions and processes (Mezger, Debaille & Kleine 2013; Mittlefehldt 2015). In addition to the well-known achondrite groups with numerous members, ungrouped achondrites, e.g. Northwest Africa (NWA) 011 (Yamaguchi et al. 2002), Graves Nunatak (GRA) 06128\/06129 (Day et al. 2009), NWA 11119 (Srinivasan et al. 2018), and NWA 7325 (Koefoed et al. 2016) expand the compositional range of achondrites towards chemically more evolved compositions (e.g. higher SiO2 contents), and thus, showcase the diversity of planetary and asteroidal crusts in the Solar system. Some achondrites, e.g. NWA 11119 and NWA 7325, yield evidence for their accretion and differentiation within the first \u223c5\u2009Myr after the formation of Ca\u2013Al-rich inclusions (CAIs); (Koefoed et al. 2016; Srinivasan et al. 2018; Zhu et al. 2019b; Barrat et al. 2021). Hence, dating more achondrites is beneficial to map the early history of Solar system.","Citation Text":["Weiss et al. 2008"],"Functions Text":["Non-chondrite meteorites (including achondrites) also originate from differentiated asteroids and planets, e.g. from Moon, Mars, and Vesta, and the angrite and ureilite parent bodies"],"Functions Label":["Background"],"Citation Start End":[[965,982]],"Functions Start End":[[763,945]]}
{"Identifier":"2018AandA...613A..35K__Anderson_et_al._2010_Instance_1","Paragraph":"As shown in Fig. 3, the differences in metallicity between different SN subclasses are not significant. This is in contradiction with what is expected from single-star evolution theory, where metallicity-driven winds are crucial: type Ic SNe, which are the most highly stripped, would show the highest metallicity, followed by type Ib, and finally the H-rich type II SNe. The observations, on the other hand, reveals that this is not the case. Some SNe Ic are even located in the low-metallicity part of the distribution in the current sample. This result strengthens the notion that metallicity may not play an important role in deciding theresulting SN type, in agreement with other works based on SN environments (Anderson et al. 2010, 2015; Leloudas et al. 2011; Galbany et al. 2016a). The environments of broad-lined SNe IcBL are found to be relatively metal poor compared to the normal CCSNe, in agreement with previous studies (Modjaz et al. 2011; Galbany et al. 2016a). However, we note that there are only two such SNe in the current sample. The explosion site of SN 1998bw (the first SN to be associated with a GRB: 980425; Galama et al. 1998; Kr\u00fchler et al. 2017) in this study shows a lower metallicity of 12 + log(O\/H) = 8.30 dex compared to the GRB-less SN 2009bb (Pignata et al. 2011), 12 + log(O\/H) = 8.49 dex. Levesque et al. (2010a), using slit spectroscopy of the explosion site, concluded that the high metallicity of SN 2009bb site is consistent with typical GRB-less SNe IcBL and not with GRB hosts. Their metallicity value recalculated on the Marino et al. (2013) N2 scale is 12 + log(O\/H) = 8.52 dex. These two different cases illustrate the importance of metallicity in deciding whether an SN IcBL progenitor would also produce GRB or not (Modjaz et al. 2008; Levesque et al. 2010b). Progenitors with higher metallicity are not able to spin fast enough and thus produce high angular momentum essential for GRB jet production, eventually producing a GRB-less SN IcBL (Woosley & Bloom 2006).","Citation Text":["Anderson et al. 2010"],"Functions Text":["This result strengthens the notion that metallicity may not play an important role in deciding theresulting SN type, in agreement with other works based on SN environments"],"Functions Label":["Similarities"],"Citation Start End":[[717,737]],"Functions Start End":[[544,715]]}
{"Identifier":"2021MNRAS.503..354G__Griv_et_al._2020_Instance_1","Paragraph":"The Lindblad\u2013Oort idea of a rotation of the Galaxy around the Galactic Centre (GC) proposes that for any type of Galactic object there is at each point in the plane a mean circular, differential motion (Oort 1927a,b). In the conventional lowest order approach to the determination of the rotation parameters of the system, a strictly circular model is adopted (Mihalas & Binney 1998). Small non-axisymmetric spiral-like perturbations of the basic axisymmetric gravitational potential induce non-circular variations of stellar velocities proportional to $\\widetilde{v}_r \\cos \\phi$ and $\\widetilde{v}_\\varphi \\sin \\phi$, and we are looking for these systematic motions in our study. Here, $\\widetilde{v}_r$ and $\\widetilde{v}_\\varphi$ are the amplitudes of the radial and tangential wave motions that depend weakly on the Galactocentric distance r, and \u03d5 is the phase of the wave (Lin et al. 1969; Yuan 1969; Rohlfs 1977; Griv et al. 2020, annexure B therein). The method of analysis used for this purpose follows the development of Cr\u00e9z\u00e9 & Mennessier (1973), Byl & Ovenden (1978), Mishurov, Pavlovskaya & Suchkov (1979), and Grivnev (1981). Specifically, in addition to the ordinary circular rotation of the system that is axisymmetric in the mean, Galactic objects have small streaming wave motions. The radial and tangential components of the velocity of a star in the plane are represented as\n(1)$$\\begin{eqnarray*}\r\nV_r = \\widetilde{v}_r \\cos \\phi ,\r\n\\end{eqnarray*}$$(2)$$\\begin{eqnarray*}\r\nV_\\varphi = r \\Omega + \\widetilde{v}_\\varphi \\sin \\phi ,\r\n\\end{eqnarray*}$$where \u03a9(r) is the angular velocity of the mean motion at the star\u2019s distance r from the GC for the type of object considered, and the amplitudes $\\widetilde{v}_r$, $\\widetilde{v}_\\varphi$ and the phase \u03d5 are to be found. Distinct populations of Galactic objects will have different average velocities V\u03c6 at r, so the value of \u03a9 should vary from population to population. An observed value of heliocentric line-of-sight velocity of a star $v$los corrected for solar motion towards the apex may be modeled in the following form:\n(3)$$\\begin{eqnarray*}\r\nv_{\\mathrm{los}} &amp;=&amp; \\left\\lbrace r_0 \\left[ \\Omega - \\Omega _0 \\right] \\sin \\ell - \\widetilde{v}_r \\cos \\phi \\cos (\\ell +\\varphi) \\right. \\nonumber \\\\\r\n&amp;&amp;\\left. \\quad + \\widetilde{v}_\\varphi \\sin \\phi \\sin (\\ell +\\varphi) + u_0 \\cos \\ell - v_0 \\sin \\ell \\right\\rbrace \\cos b \\nonumber \\\\\r\n&amp;&amp;\\quad - w_0 \\sin b ,\r\n\\end{eqnarray*}$$where \u2113 and b are the Galactic coordinates, the angle \u03c6 is measured clockwise in the direction of overall rotation from the radius passing through the location of the Sun, \u03a90(r0) is the angular velocity of the mean motion at the Sun\u2019s distance r0, \u03d5 = \u03d50 \u2212 m[\u03c6 \u2212 (1\/tan\u2009p)ln\u2009(r\/r0)], \u03d50 is the phase of the wave at the Sun\u2019s location of the different Fourier m-modes, and the constants u0, $v$0, $w$0 are the components of solar peculiar motion relative to the mean linear speed of rotation at the Sun\u2019s distance r0\u03a90 (Mihalas & Binney 1998). To reiterate, the terms \u221d(\u03a9 \u2212 \u03a90) and u0, $v$0, $w$0 describe the mean rotation of the system under study and the peculiar motion of the Sun. The deviation of the motion of objects from the circular motion due to a wave perturbation is characterized by terms ${\\propto} \\widetilde{v}_r \\cos \\phi$ and $\\widetilde{v}_\\varphi \\sin \\phi$.","Citation Text":["Griv et al. 2020"],"Functions Text":["Here, $\\widetilde{v}_r$ and $\\widetilde{v}_\\varphi$ are the amplitudes of the radial and tangential wave motions that depend weakly on the Galactocentric distance r, and \u03d5 is the phase of the wave",", annexure B therein"],"Functions Label":["Uses","Uses"],"Citation Start End":[[921,937]],"Functions Start End":[[682,878],[937,957]]}
{"Identifier":"2015ApJ...806..152S__Ransom_et_al._2005_Instance_2","Paragraph":"One of the most astonishing characteristics of Liller 1 is the extremely large value of the collision rate parameter. Verbunt & Hut (1987) showed that Liller 1 has the second-highest value of stellar encounter rate (after Terzan 5; see also Lanzoni et al. 2010) among all star clusters in the Galaxy, thus suggesting that it represents an ideal environment where exotic objects, generated by collisions, can form. In fact, it is commonly believed that dynamical interactions in GCs facilitate the formation of close binary systems and exotic objects such as cataclysmic variables (CVs), low-mass X-ray binaries (LMXBs), millisecond pulsars (MSPs), and blue straggler stars (BSSs) (e.g., Bailyn 1992; Paresce et al. 1992; Ferraro et al. 2001, 2009a, 2012; Ransom et al. 2005; Pooley & Hut 2006). Moreover, Hui et al. (2010) found that clusters with large collisional parameters and high metallicity (see also Bellazzini et al. 1995) usually host more MSPs. Indeed, Terzan 5 hosts the largest population of MSPs among all Galactic GCs (Ransom et al. 2005).\n6\n\n\n\n6\n\nNote that Terzan 5 is suspected to not be a genuine GC, because it harbors at least three stellar populations with different iron abundances (Ferraro et al. 2009a; Origlia et al. 2011, 2013; Massari et al. 2014).\n A strong \u03b3-ray emission has been recently detected in the direction of Liller 1 by the Large Area Telescope (LAT) on board Fermi (Tam et al. 2011). This is the most intense emission detected so far from a Galactic GC, again suggesting the presence of a large number of MSPs. However, no direct radio detection of these objects has been obtained so far in this system (Ransom et al. 2005). The only exotic object identified in the cluster is the rapid burster MXB 1730-335, an LMXB observed to emit radio waves and type I and type II X-ray bursts (Hoffman et al. 1978). It seems to be located in the central region of Liller 1, but no optical\/IR counterpart of this object has been found so far (Homer et al. 2001).","Citation Text":["Ransom et al. 2005"],"Functions Text":["Indeed, Terzan 5 hosts the largest population of MSPs among all Galactic GCs"],"Functions Label":["Motivation"],"Citation Start End":[[1034,1052]],"Functions Start End":[[956,1032]]}
{"Identifier":"2021ApJ...910...52C__Kleihaus_et_al._2004_Instance_1","Paragraph":"Half a century ago, a series of theorems laid the ground for the Kerr hypothesis (Israel 1967; Carter 1971; Robinson 1975); according to these no-hair theorems, the only stationary, axisymmetric, asymptotically flat, regular outside of the horizon solution to four-dimensional GR when the matter fields feature the same isometries as the spacetime is the Kerr BH. Notwithstanding their significance, there are many ways with which to circumvent them and discover different solutions. Still, in four dimensions, hairy BHs have been described in different theories of gravity, such as Einstein\u2013Yang\u2013Mills (Bizon 1990; K\u00fcnzle & Masood-ul-Alam 1990; Volkov & Galtsov 1990; Breitenlohner et al. 1992; Kleihaus & Kunz 1998, 2001; Kleihaus et al. 2004), scalar-tensor (Bocharova et al. 1970; Bekenstein 1974; Kleihaus et al. 2015; Collodel et al. 2020a), and Gauss\u2013Bonnet theories (Kanti et al. 1996; Kleihaus et al. 2011, 2016; Antoniou et al. 2018; Doneva & Yazadjiev 2018; Silva et al. 2018; Cunha et al. 2019; Collodel et al. 2020b; Herdeiro et al. 2021; Berti et al. 2021). Remarkably, by dropping the assumption that the matter fields must be stationary and axisymmetric, Herdeiro and Radu found solutions in the context of GR where BHs have hair (Herdeiro & Radu 2014b, 2015), by minimally coupling to gravity a complex scalar field that depends on time and on the axial coordinate while its energy-momentum tensor still possesses the respective isometries; see Herdeiro et al. (2015, 2016a, 2016b), Brihaye et al. (2016), and Delgado et al. (2016) for generalizations. These are known as scalarized Kerr black holes (KBHsSH) and they are the object of study of this paper. In their domain of existence, they connect Kerr BHs (that is, with no hair) with pure solitonic solutions, also known as boson stars (BS), which are regular everywhere and feature no horizons. In this sense, one can think of the KBHsSH indeed as a combined system of a BS with a horizon at its center, and therefore it shares traits of both objects.","Citation Text":["Kleihaus et al. 2004"],"Functions Text":["Still, in four dimensions, hairy BHs have been described in different theories of gravity, such as Einstein\u2013Yang\u2013Mills"],"Functions Label":["Background"],"Citation Start End":[[724,744]],"Functions Start End":[[484,602]]}
{"Identifier":"2018MNRAS.476.1889T__Kulkarni,_Hut_&_McMillan_1993_Instance_1","Paragraph":"The high interaction rates in globular clusters, especially at their centres, lead to the efficient formation of exotic binaries like accreting compact objects \u2013 cataclysmic variables (white dwarfs) and X-ray binaries (neutron stars and black holes; Clark, Markert & Li 1975). Soon after the first X-ray missions, it was recognized that X-ray transients in globular clusters were disproportionately associated with neutron stars, with no confirmed black holes (Verbunt et al. 1995), a fact that remains true to this day (Bahramian et al. 2014). In contrast, about a third of the X-ray transients in the rest of the Galaxy contain black holes (Miller-Jones et al. 2015; Tetarenko et al. 2016). Originally it was proposed that due to mutual gravitational interactions that all black holes would have been ejected from globular clusters (Kulkarni, Hut & McMillan 1993; Sigurdsson & Hernquist 1993). However, there are several recent indications that a number of black holes may still exist within clusters, based on X-ray observations of extragalactic clusters (Maccarone et al. 2007; Brassington et al. 2010; Irwin et al. 2010; Shih et al. 2010; Maccarone et al. 2011), radio\/X-ray observations of Galactic clusters (Maccarone & Knigge 2007; Strader et al. 2012; Chomiuk et al. 2013; Miller-Jones et al. 2015), as well as theoretical simulations (Sippel & Hurley 2013; Morscher et al. 2015; Peuten et al. 2016). The reason why these black holes are so elusive could perhaps be due to the nature of the X-ray binaries they form, rather than their number. If the majority of accreting black holes in globular clusters have very low-mass, possibly degenerate, donors, they will have short, faint outbursts, making them undetectable by current and past all-sky X-ray surveys (Knevitt et al. 2014). Indeed, the dynamics in globular clusters are thought to effectively produce ultracompact X-ray binaries (Verbunt 1987; Deutsch et al. 1996; Deutsch, Margon & Anderson 2000; Ivanova et al. 2005, 2010), in which a black hole or neutron star accretes matter from an H-poor donor in a system with a very short orbital period (\u22721 \u2009h; Nelson, Rappaport & Joss 1986).","Citation Text":["Kulkarni, Hut & McMillan 1993"],"Functions Text":["Originally it was proposed that due to mutual gravitational interactions that all black holes would have been ejected from globular clusters"],"Functions Label":["Background"],"Citation Start End":[[835,864]],"Functions Start End":[[693,833]]}
{"Identifier":"2020AandA...644A..97C__Leroy_et_al._2013_Instance_4","Paragraph":"Major nearby galaxy cold gas mapping surveys (Regan et al. 2001; Wilson et al. 2009; Rahman et al. 2011; Leroy et al. 2009; Donovan Meyer et al. 2013; Bolatto et al. 2017; Sorai et al. 2019; Sun et al. 2018) have focused on observations of the molecular gas (through CO lines). Despite a few notable exceptions (e.g. Alatalo et al. 2013; Saintonge et al. 2017), these surveys observed mainly spiral or infrared-bright galaxies (i.e. galaxies with significant star formation) and have furthered our understanding of how star formation happens, rather than how it stops. This boils down to quantifying the relation between molecular gas and star formation rate (SFR), which appears nearly linear in nearby discs (Kennicutt 1998; Bigiel et al. 2008; Leroy et al. 2013; Lin et al. 2019). This relationship is often parametrised via the ratio between the SFR and the molecular gas mass (Mmol), which is called the molecular star formation efficiency (SFE = SFR\/Mmol = 1\u2215\u03c4dep), where the inverse of the SFE is the depletion time, \u03c4dep. The depletion time indicates how much time is necessary to convert all the available molecular gas into stars at the current star formation rate. On kpc scales and in the discs of nearby star-forming galaxies, \u03c4dep is approximately constant around 1\u20132 Gyr (Bigiel et al. 2011; Rahman et al. 2012; Leroy et al. 2013; Utomo et al. 2017), and it appears to weakly correlate with many galactic properties such as stellar mass surface density or environmental hydrostatic pressure (Leroy et al. 2008; Rahman et al. 2012). Nevertheless, small but important deviations for a constant SFE have been noticed, which can be the first hints of star formation quenching. In some galaxies, the depletion time in the centres appear shorter (Leroy et al. 2013; Utomo et al. 2017) or longer (Utomo et al. 2017) with respectto their discs. These differences may correlate with the presence of a bar or with galaxy mergers (Utomo et al. 2017; see also Muraoka et al. 2019) and do not seem to be related to unaccounted variation in the CO-to-H2 conversion factor (Leroy et al. 2013; Utomo et al. 2017). Spiral arm streaming motions have also been observed to lengthen depletion times (Meidt et al. 2013; Leroy et al. 2015).","Citation Text":["Leroy et al. 2013"],"Functions Text":["These differences","do not seem to be related to unaccounted variation in the CO-to-H2 conversion factor"],"Functions Label":["Background","Background"],"Citation Start End":[[2074,2091]],"Functions Start End":[[1852,1869],[1988,2072]]}
{"Identifier":"2020AandA...637A..59A__Massalkhi_et_al._2019_Instance_3","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"],"Functions Text":["but has been interpreted as evidence that SiO disappears from the gas phase at high densities to be incorporated into dust grains"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1114,1135]],"Functions Start End":[[910,1039]]}
{"Identifier":"2021ApJ...916...70Z__Liu_et_al._2020_Instance_1","Paragraph":"At this juncture in cosmology, independent and complementary probes with considerable precision are very helpful for providing clues about the origin of the abovementioned crises. Strong-lensing systems are one of the most promising probes for investigating these issues. The time delay between images of strongly lensed time-variable sources was proposed to directly determine H0 (Refsdal 1964; Treu 2010). Recently, with elaborate time-delay measurements and lens modeling for five selected strongly lensed quasar systems, the H0 Lenses in COSMOGRAIL\u2019s Wellspring (H0LiCOW) team yielded an estimation of H0 that is also 2.3\u03c3 higher than the Planck-calibrated value (Birrer et al. 2019). It should be mentioned that this result is completely independent of all rungs of the distance ladder. Moreover, on the basis of the distance sum rule (DSR), distance ratios derived from angular separation measurements of images in strong-lensing systems were proposed to test the FLRW metric and model-independently determine \u03a9K (R\u00e4s\u00e4nen et al. 2015). This method has been intensively implemented with updated observations (Xia et al. 2017; Li et al. 2018a; Liu et al. 2020; Zhou & Li 2020). In order to reduce the systematics resulting from oversimplified lensing modeling in the distance ratio method, Liao et al. (2017b) reformulated the DSR in terms of the time-delay distance, which is a combination of three angular diameter distances in a strong-lensing system. With the latest strongly lensed quasar observations and type Ia supernovae (SNe Ia) luminosity distance measurements, H0 and \u03a9K were simultaneously and model-independently determined to a precision of \u223c6% and \u223c0.3 (Collett et al. 2019), respectively. Meanwhile, strongly lensed gravitational waves (GWs) and their electromagnetic (EM) counterparts from the binary of compact object coalescence are proposed as a powerful tool for precision cosmology since the time delay between images in these systems can be precisely measured (Liao et al. 2017a; Li et al. 2019). However, both traditional strongly lensed quasar and proposed strongly lensed GW systems face shortages. For lensed quasars, the precision of time-delay measurements is limited to \u223c3% and lens modeling is difficult to improve because of the bright active-galactic nuclei (AGNs) contamination in the source host galaxy. For the expected lensed GWs, the main challenges may be the event rate and localization ability for images of GW signals, which is crucial for lens modeling.","Citation Text":["Liu et al. 2020"],"Functions Text":["This method has been intensively implemented with updated observations"],"Functions Label":["Background"],"Citation Start End":[[1148,1163]],"Functions Start End":[[1042,1112]]}
{"Identifier":"2015AandA...579A..56B__Lebrun_et_al._2003_Instance_1","Paragraph":"Swift\u2009J1734.5-3027 was detected by INTEGRAL for the first time during the observations performed toward the Galactic bulge in the satellite revolution 1329, i.e., about half a day before the BAT discovery (from 56\u2009535.85950 MJD to 56\u2009536.01338 MJD; see Table 1). It remained within the field of view (FoV) of the instruments on-board INTEGRAL until satellite revolution 1348 (from 56\u2009592.54664 MJD to 56\u2009593.31612 MJD), when the window of seasonal visibility toward the Galactic center closed. We analyzed all INTEGRAL data by using version 10.1 of the Off-line Scientific Analysis software (OSA) distributed by the ISDC (Courvoisier et al. 2003). INTEGRAL observations are divided into science windows (SCWs), i.e., pointings with typical durations of ~2\u22123 ks. To limit the ISGRI calibration uncertainties (Lebrun et al. 2003), we made use of all public SCWs from the Galactic bulge, Scutum\/Sagittarium, and Perseus\/Norma monitoring programs during which the source was located to within 12 deg from the center of the IBIS FoV (Ubertini et al. 2003). We also included in our dataset the observations of the Galactic center and 4U\u20091728-34 for which our group was awarded data rights in revolutions 1329\u22121348. A summary of the total exposure-time available in each revolution is provided in Table 1 for IBIS\/ISGRI and the two JEM-X telescopes (Lund et al. 2003). We extracted the mosaics and spectra in each revolution for the two JEM-X and ISGRI. All JEM-X spectra were extracted by using the standard 16-channel response matrix, while a customized 37 energy bin response matrix was created for ISGRI in order to optimize the signal-to-noise ratio (S\/N) in the energy range (20\u221250\u2009keV). Only in revolution 1329, was the ISGRI spectrum extracted with a reduced energy binning (8 channels), as the source was relatively faint (see Sect. 3). JEM-X lightcurves with a time resolution of 2 s were extracted from all observations to search for Type-I X-ray bursts, but none was found. We did not perform other timing analyses of the INTEGRAL data as the source was too faint to extract meaningful power spectra. In Fig. 1 we show the ISGRI mosaic realized by using all available data. ","Citation Text":["Lebrun et al. 2003"],"Functions Text":["To limit the ISGRI calibration uncertainties"],"Functions Label":["Uses"],"Citation Start End":[[808,826]],"Functions Start End":[[762,806]]}
{"Identifier":"2020AandA...641A..85S__Orienti_&_Dallacasa_2008_Instance_2","Paragraph":"To derive the equipartition magnetic field of J1146+4037, we predict the rest-frame 8.4 GHz (redshifted to 1.4 GHz at z\u2004=\u20045.0059) flux density from our spectral model. However, there is no source size measurement at 1.4 GHz. We make use of the full width at half maximum (FWHM) source size of 0.74\u2005\u00b1\u20050.01 mas derived by the Gaussian fit from 5 GHz VLBI mas angular resolution observations (Frey et al. 2010). We note that in our calculations, we assume a source size that is 1.8 times larger than the FWHM, following the approach of Readhead (1994) and Orienti & Dallacasa (2008). The derived equipartition magnetic field is \n\n\n\n34\n\n\u2212\n7\n\n\n+\n8\n\n\n\n$ 34^{+8}_{-7} $\n\n\n mG. This is within the range of the equipartition magnetic fields of 17 HFP radio sources (7\u201360 mG; quasars and galaxies at 0.22\u2004 \u2004z\u2004 \u20042.91; Orienti & Dallacasa 2012) and 5 HFPs at 0.084\u2004 \u2004z\u2004 \u20041.887 (18\u2013160 mG; Orienti & Dallacasa 2008). The magnetic field calculated from the turnover information listed in Table 4 is \n\n\n\n1\n.\n\n8\n\n\u2212\n2.7\n\n\n+\n2.3\n\n\n\n\n$ 1.8^{+2.3}_{-2.7} $\n\n\n G assuming an SSA origin with Eq. (3), however the uncertainty is very large. The large uncertainty is caused by the fact that we only have four data points to constrain the turnover information and we do not have source size measurements at the turnover frequency, but rather we assume the source size measured at another frequency. More data taken in other wavelength bands are needed to meaningfully constrain the turnover peak, and mas resolution observations at the peak frequency are needed to give reliable magnetic field strength measurements. This may indicate that the turnover is not caused by SSA, by comparing the large magnetic field strength measured from the spectral turnover (\n\n\n\n1\n.\n\n8\n\n\u2212\n2.7\n\n\n+\n2.3\n\n\n\n\n$ 1.8^{+2.3}_{-2.7} $\n\n\n G) with the equipartition magnetic field strength (\n\n\n\n34\n\n\u2212\n7\n\n\n+\n8\n\n\n\n$ 34^{+8}_{-7} $\n\n\n mG). As J1146+4037 is a strong blazar, the turnover may be caused by its strong variability. Another possible explanation for the spectral turnover is that high-density plasma in the nuclear region attenuates the radio emission from the central active BH. High-resolution, interstellar medium observations of the nuclear region of this target may address the latter issue.","Citation Text":["Orienti & Dallacasa 2008"],"Functions Text":["This is within the range of the equipartition magnetic fields of 17 HFP radio sources","and 5 HFPs at 0.084\u2004 \u2004z\u2004 \u20041.887 (18\u2013160 mG"],"Functions Label":["Similarities","Similarities"],"Citation Start End":[[877,901]],"Functions Start End":[[670,755],[833,875]]}
{"Identifier":"2021AandA...654A..89P__Vincentelli_et_al._2021_Instance_1","Paragraph":"The spectral analysis above 3 keV allows us to characterise the hot corona and the relativistic reflection emission properties. The deep January 2017 NuSTAR observation was also included in the simultaneous fit (Ezhikode et al. 2020; Panagiotou & Walter 2020). From the RELXILL reflection model, and assuming a primary exponential cutoff power-law continuum, we find moderate reflection strengths, \u211b \u223c 0.1\u22120.2, and high cutoff energies, Ecut \u223c 110\u2212120 keV. These values are in very good agreement with those measured from the average long-term Swift\/BAT (Burst Alert Telescope) spectrum (Vincentelli et al. 2021). Applying relativistic reflection models that assume a primary Comptonisation continuum, we infer the hot corona temperature to be kThot \u223c 26\u221231 keV (kThot \u223c 21\u221222 keV) and the optical depth to be \u03c4hot \u223c 2 (\u03c4hot \u223c 6\u22127) for the slab (or spherical) geometry. From the spectral analysis, it is not possible to discriminate between either of the hot corona geometries, although the slab geometry provides a better fit. In the near future, X-ray polarimetry is expected to play an important role within such a framework (e.g. Schnittman & Krolik 2010; Beheshtipour et al. 2017; Tamborra et al. 2018; Marinucci et al. 2019) thanks to, for example, IXPE (Imaging X-Ray Polarimetry Explorer; Weisskopf et al. 2016) and eXTP (Enhanced X-ray Timing and Polarimetry observatory; Zhang et al. 2016). While the corona temperatures found for Mrk 110 are broadly consistent with the average ones found by Middei et al. (2019) from a sample of 26 AGN (with \u27e8kThot\u27e9 = 50\u2005\u00b1\u200521 keV and \u27e8kThot\u27e9 = 53\u2005\u00b1\u200523 keV for the slab and spherical geometry, respectively), it is likely to be located in the lower range of this distribution. However, its hot coronal temperature is not as low as the temperatures inferred for some AGN with much lower high-energy cutoffs, such as GRS 1734\u2013292 (Tortosa et al. 2017), Ark 564 (Kara et al. 2017), and PDS 456 (Reeves et al. 2021b), where kT could be as low as 15 keV.","Citation Text":["Vincentelli et al. 2021"],"Functions Text":["These values are in very good agreement with those measured from the average long-term Swift\/BAT (Burst Alert Telescope) spectrum"],"Functions Label":["Similarities"],"Citation Start End":[[588,611]],"Functions Start End":[[457,586]]}
{"Identifier":"2020AandA...639A.104S__Kitaura_et_al._2006_Instance_1","Paragraph":"Today, at least three known populations of gap transients are recognized in the luminoisty gap. These include classical luminous blue variable (LBV) outbursts, intermediate-luminosity red transients (ILRTs), and luminous red novae (LRNe). LBVs are thought to be related to eruptions of massive luminous stars (see Smith et al. 2011). As shown in greater detail in Stritzinger et al. (2020, hereafter Paper I), the ILRT subtype is well represented by NGC 300-2008-OT and SN 2008S, and has been linked to asymptotic giant branch (S-AGB) stars (Prieto et al. 2008, 2009; Thompson et al. 2009; Botticella et al. 2009; Kochanek 2011; Adams et al. 2016; Doherty et al. 2017) that die as electron-capture supernovae (Miyaji et al. 1980; Nomoto 1984; Miyaji & Nomoto 1987; Hashimoto et al. 1993; Kitaura et al. 2006; Poelarends et al. 2008). Other models appearing in the literature for ILRTs consist of moderately massive stars experiencing super-Eddington winds and\/or giant outbursts (e.g., Smith et al. 2009; Humphreys et al. 2011), or massive stars donating material to a main-sequence star, leading to the release of gravitational energy (e.g., Kashi et al. 2010). Finally, a leading model for the origins of LRNe, which all display a ubiquitous double-humped light curve (Pastorello et al. 2019a), consists of the ejection of a common envelope by a massive binary system (e.g., Blagorodnova et al. 2017) upon coalescence (Smith et al. 2016; Metzger & Pejcha 2017; Lipunov et al. 2017; Mauerhan et al. 2018). However, other models have also been proposed in the past to account for LRNe, particularly within articles that have studied the Galactic LRN archetype V838 Mon. These include, among others, outbursts from massive stars (Tylenda 2005), accretion of low-mass stars onto solar-mass main-sequence companions (Soker & Tylenda 2003; Tylenda & Soker 2006; Kashi et al. 2010; Kashi & Soker 2016; Soker 2020), or even giant stars that accrete relatively massive planets (Retter & Marom 2003).","Citation Text":["Kitaura et al. 2006"],"Functions Text":["As shown in greater detail in Stritzinger et al. (2020, hereafter Paper I), the ILRT subtype is well represented by NGC 300-2008-OT and SN 2008S, and has been linked to asymptotic giant branch (S-AGB) stars","that die as electron-capture supernovae"],"Functions Label":["Background","Background"],"Citation Start End":[[788,807]],"Functions Start End":[[334,540],[669,708]]}
{"Identifier":"2021ApJ...911...89M__Mozer_et_al._2020a_Instance_2","Paragraph":"Time domain structures (TDSs; electrostatic or electromagnetic electron holes, ion holes, solitary waves, double layers, nonlinear whistlers, etc.) are \u223c1 ms pulses having significant electric fields parallel to the background magnetic field (Mozer et al. 2015). They are abundant through space, occurring along auroral zone magnetic field lines (Temerin et al. 1982; Mozer et al. 1997; Ergun et al. 1998), in the magnetospheric tail and plasma sheet (Cattell et al. 2005; Tong et al. 2018; Lotekar et al. 2020), at reconnection sites (Cattell et al. 2005; Steinvall et al. 2019; Lotekar et al. 2020), in the solar wind (Mangeney et al. 1999; Malaspina et al. 2013), in collisionless shocks (Wilson et al. 2010; Vasko et al. 2020; Wang et al. 2020), and in the magnetospheres of other planets (Pickett et al. 2015). TDSs are also expected along the Parker Solar Probe orbit (Mozer et al. 2020a). According to theoretical estimates and simulations (Cranmer & van Ballegooijen 2003; Valentini et al. 2011, 2014), these nonlinear structures can provide thermalization of electron and ion beams produced in the course of the turbulence cascade development at scales the order of the electron inertial length and down to the Debye length. This paper discusses such observations at a heliocentric distance of 35 solar radii. The electric field experiment on the Parker Solar Probe measures electric fields from DC to 20 MHz. A general description of the instrument and its electronics appears elsewhere (Bale et al. 2016). In this paper, the data from DC to 2 MHz are discussed. These measurements are obtained from the potentials of the four antennas, V1 through V4, that are located in the plane perpendicular to the Sun\u2013satellite line. They produce E12 = (V1\u2212V2)\/3.5 and E34 = (V3\u2212V4)\/3.5, which are then rotated into the spacecraft coordinate system to produce EX and EY. The direction, X, is perpendicular to the Sun\u2013spacecraft line, in the ecliptic plane, and pointing in the direction of solar rotation (against the ram direction), Y is perpendicular to the ecliptic plane, pointing southward, and Z points toward the Sun. The numerical factor, 3.5, is the effective antenna length (Mozer et al. 2020a). The uncertainties of the amplitudes of the waves and time domain structures reported in this paper are estimated to be about a factor of 2. These amplitudes are underestimated because the capacitive divider that couples the antennas to the electronics decreases the measured electric field relative to that on the antennas. They are often overestimated because short antennas produce overestimates of the electric field by factors of 2\u20134, as was observed during antenna deployment on the Cluster satellite and as is observed on the Parker Solar Probe from the ratio of the electric field to the magnetic field in whistlers.","Citation Text":["Mozer et al. 2020a"],"Functions Text":["In this paper, the data from DC to 2 MHz are discussed. These measurements are obtained from the potentials of the four antennas, V1 through V4, that are located in the plane perpendicular to the Sun\u2013satellite line. They produce E12 = (V1\u2212V2)\/3.5 and E34 = (V3\u2212V4)\/3.5, which are then rotated into the spacecraft coordinate system to produce EX and EY. The direction, X, is perpendicular to the Sun\u2013spacecraft line, in the ecliptic plane, and pointing in the direction of solar rotation (against the ram direction), Y is perpendicular to the ecliptic plane, pointing southward, and Z points toward the Sun. The numerical factor, 3.5, is the effective antenna length"],"Functions Label":["Uses"],"Citation Start End":[[2184,2202]],"Functions Start End":[[1517,2182]]}
{"Identifier":"2019AandA...629A..54U__Marinucci_et_al._2015_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":["Marinucci et al. 2015"],"Functions Text":["No Compton reflection hump has been detected with","or NuSTAR","despite the presence of a complex Fe K\u03b1 line."],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[966,987]],"Functions Start End":[[877,926],[955,964],[990,1035]]}
{"Identifier":"2021AandA...651L...8D__Just_et_al._2007_Instance_2","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"],"Functions Text":["The initial dispersion of the LUV\u2005\u2212\u2005LX relation is relatively large (\u03b4\u2004\u223c\u20040.35\u22120.4,","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."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1639,1655]],"Functions Start End":[[1556,1638],[1677,1838]]}
{"Identifier":"2017AandA...599A..55B__Shakura_1972_Instance_2","Paragraph":"When the characteristic time of variability of the mass flux along the accretion disk is longer than the relaxation time of the local disk equilibrium, it is possible to use the approximation of local equilibrium Shakura (1972), see also Bisnovatyi-Kogan (2011), to calculate the transient disk structure. The equilibrium along a radius of the accretion disk around a star with a mass M is determined by the Keplerian rotational velocity \u03a9K(1)\\begin{equation} \\Omega=\\Omega_{\\rm K}=\\left(\\frac{GM}{r^3}\\right)^{1\/2}\\cdot \\label{omega} \\end{equation}\u03a9=\u03a9K=GMr31\/2\u00b7Writing the equation of the vertical equilibrium in approximate algebraic form, we obtain (2)\\begin{equation} h=\\sqrt 2 \\frac{v_{\\rm s}}{\\Omega}, \\label{h} \\end{equation}h=2vs\u03a9,where \\hbox{$v_{\\rm s}=\\sqrt{P\/\\rho}$}vs=P\/\u03c1 is the speed proportional to the sound velocity, P and \u03c1 are the (gas + radiation) pressure and density at the symmetry plane of the accretion disk, and h is the semi-thickness of the accretion disk. The specific angular momentum l of the matter in the accretion disk is connected to the rotation velocity as(3)\\begin{equation} l=r\\,v_\\phi=r^2\\Omega. \\label{l} \\end{equation}l=r\u2009v\u03c6=r2\u03a9.The mass flux through the disk at radius r is connected to the radial velocity vr as (4)\\begin{equation} \\dot M=-4\\pi h\\rho r v_r, \\quad \\dot M>0,\\quad v_r<0. \\label{mflux} \\end{equation}M\u0307=\u22124\u03c0h\u03c1rvr,\u2001M\u0307>0,\u2001vr0.We use an \u03b1 approximation for the turbulent viscosity (Shakura 1972) when the (r\u03c6) component of the stress tensor tr\u03c6 is written as (5)\\begin{equation} t_{r\\phi}=\\alpha\\, P, \\label{trphi} \\end{equation}tr\u03c6=\u03b1\u2009P,where the phenomenological non-dimensional parameter \u03b1 \u2264 1. The condition of stationarity of the angular momentum, in which the outward viscous radial flux of the angular momentum is balanced by the angular momentum of the inward flux of the mass, is written as (see, e.g., Bisnovatyi-Kogan 2011) (6)\\begin{equation} r^2 h \\alpha P=\\frac{\\dot M}{4\\pi}(l-l_{\\rm in}). \\label{angmom} \\end{equation}r2h\u03b1P=M\u03074\u03c0(l\u2212lin).The main input into the time lag comes from the outer regions of the disk with l \u226b lin. Then we have from Eqs. (4) and (6) the expression for the radial velocity in the form (7)\\begin{equation} v_r=-\\alpha\\frac{v_{\\rm s}^2}{v_\\phi}\\cdot \\label{vr} \\end{equation}vr=\u2212\u03b1vs2v\u03c6\u00b7We also define the surface density \u03a3, and write Eq. (6) in light of Eq. (3), using condition l \u226b lin, in the form (8)\\begin{equation} \\Sigma=2\\rho h,\\quad \\dot M \\Omega=4\\pi \\alpha P h. \\label{sigma} \\end{equation}\u03a3=2\u03c1h,\u2001M\u0307\u03a9=4\u03c0\u03b1Ph.The equation of the local thermal balance in the accretion disk, when the heat produced by viscosity Q+ is entirely emitted through the sites of the optically thick accretion disk with a total flux Q\u2212, at l \u226b lin is written as (see, e.g., Bisnovatyi-Kogan 2011)(9)\\begin{equation} \\frac{3}{2}\\dot M \\Omega^2=\\frac{16\\pi ac T^4}{3\\varkappa \\Sigma}\\cdot \\label{theq} \\end{equation}32M\u0307\u03a92=16\u03c0acT43\u03f0\u03a3\u00b7Here T is the temperature in the symmetry plane of the accretion disk, a is the constant of the radiation energy density, c is the speed of light, and \u03f0 is the Thompson (scattering) opacity of the matter. ","Citation Text":["Shakura 1972"],"Functions Text":["We use an \u03b1 approximation for the turbulent viscosity","when the (r\u03c6) component of the stress tensor tr\u03c6 is written as (5)\\begin{equation} t_{r\\phi}=\\alpha\\, P, \\label{trphi} \\end{equation}tr\u03c6=\u03b1\u2009P,where the phenomenological non-dimensional parameter \u03b1 \u2264 1."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1435,1447]],"Functions Start End":[[1380,1433],[1449,1649]]}
{"Identifier":"2022MNRAS.516.5618P__Muzahid_et_al._2020_Instance_1","Paragraph":"Recent technological advances related to 3D Integral Field Spectroscopy (IFS), which produces data cubes where each pixel on the image has a spectrum, have opened a new window for examining the CGM gas. This approach combines the information gathered in absorption against background sources (whose lines of sight pass through a galaxy\u2019s CGM) with traditional emission-based properties of galaxies. Following at least two decades of limited success in identifying the galaxies associated with quasar absorbers, IFS have open a new era in establishing the relation between absorption and emission with high success rates. Early efforts with near-infrared IFS VLT\/SINFONI (Bouch\u00e9 et al. 2007; P\u00e9roux et al. 2011; P\u00e9roux et al. 2013, 2016) led to efficient discoveries of star-forming galaxies associated with Mg ii and H i absorbers at z \u223c 2 (see also Rudie, Newman & Murphy 2017; Joshi et al. 2021). The optical IFS VLT\/MUSE (Bacon et al. 2010) has proved to be a true game-changer in the field. Early on, the MUSE Guaranteed Time Observations (GTO) team established surveys including MUSE-QuBES (Muzahid et al. 2020) and MEGAFLOW (Bouch\u00e9 et al. 2016; Schroetter et al. 2016; Schroetter et al. 2019; Zabl et al. 2019) to relate gas traced by absorbers to galaxies. In a parallel effort, the MAGG survey targets higher redshift galaxies (Fumagalli, O\u2019Meara & Prochaska 2016; Dutta et al. 2020; Lofthouse et al. 2020). The Cosmic Ultraviolet Baryon Survey CUBS instead is absorption-blind and uncovers new quasar absorbers in a wide range of column densities (ranging from few times 16.0  $\\rm{log}\\,\\,{\\it N}(\\rm{H\\,\\,{\\small I}})$20.1) at z1 (Chen et al. 2020; Boettcher et al. 2021; Cooper et al. 2021; Zahedy et al. 2021). By extending to bluer wavelengths, the optical IFS Keck\/KCWI (Martin et al. 2010) has enabled similar studies at higher spectral resolution (Martin et al. 2019; Nielsen et al. 2020). BlueMUSE, a blue-optimized, medium spectral resolution IFS based on the MUSE concept and proposed for the Very Large Telescope is also under planning (Richard et al. 2019). Contemporary to these works, ALMA - which can be viewed as an IFS at mm-wavelengths - has enabled the detections of both CO and [CII] emission in galaxies associated with strong quasar absorbers at intermediate and high redshifts, respectively (Neeleman et al. 2016; Kanekar et al. 2018; Klitsch et al. 2018; Neeleman et al. 2018; Neeleman et al. 2019; P\u00e9roux et al. 2019; Klitsch et al. 2021; Szakacs et al. 2021a). These lines enable us to trace the colder (\u223c100K) and denser phase of the neutral gas: the molecular hydrogen, H2. The molecular gas constitutes the ultimate phase of the gas reservoir from which stars form and hence is an essential link to the baryon cycle. Together, these IFS observations have provided unique information on the resolved galaxy kinematics which can then be combined with the gas dynamics to probe gas flows in the CGM regions (Bouch\u00e9 et al. 2013; Rahmani et al. 2018a; Schroetter et al. 2019; Zabl et al. 2019; Neeleman et al. 2020; Szakacs et al. 2021a).","Citation Text":["Muzahid et al. 2020"],"Functions Text":["Early on, the MUSE Guaranteed Time Observations (GTO) team established surveys including MUSE-QuBES","to relate gas traced by absorbers to galaxies."],"Functions Label":["Background","Background"],"Citation Start End":[[1096,1115]],"Functions Start End":[[995,1094],[1217,1263]]}
{"Identifier":"2021MNRAS.507.1421M__Novikov,_Schmalzing_&_Mukhanov_2000_Instance_1","Paragraph":"Topological estimators such as the Minkowski functionals (MFs) are also important diagnostics in this direction as they carry information at all orders. The MFs have been extensively developed as a statistical tool in a cosmological setting for both two-dimensional (2D; projected) and 3D (redshift) surveys. The MFs have analytically known results for a Gaussian random field, making them suitable for studies of non-Gaussianity. Examples of such studies include CMB data (Schmalzing & G\u00f3rski 1998; Novikov, Schmalzing & Mukhanov 2000; Hikage et al. 2008; Natoli et al. 2010; Ducout et al. 2013; Planck Collaboration 2016b; Planck Collaboration IX, 2020a), LSS (Gott, Mellot & Dickinson 1986; Coles 1988; Gott et al. 1989, 1992; Melott 1990; Moore et al. 1992; Canavezes et al. 1998; Schmalzing & Diaferio 2000; Kerscher et al. 2001; Hikage et al. 2002, 2008 Park et al. 2005; Hikage, Komatsu & Mastubara 2006), weak lensing (Matsubara & Jain 2001; Sato et al. 2001; Taruya et al. 2002; Munshi et al. 2011d), Sunyaev\u2013Zel\u2019dovich maps (Munshi et al. 2011c), 21cm (Gleser et al. 2006), and N-body simulations (Schmalzing & Diaferio 2000; Kerscher et al. 2001). Note that this is an incomplete list of references and we have selected a sample of representative papers from the literature. The MFs are spatially defined topological statistics and, by definition, contain statistical information of all orders. This makes them complementary to the polyspectra methods that are defined in Fourier space. It is also possible that the two approaches will be sensitive to different aspects of non-Gaussianity and systematic effects, although in the weakly non-Gaussian limit it has been shown that the MFs reduce to a weighted probe of the bispectrum (Hikage et al. 2006). In addition to providing cosmological information, MFs can also be useful diagnostics of any unknown systematics as well as baryonic contamination which are expected to affect weak lensing observables (Herenois-Deraps et al. 2016).","Citation Text":["Novikov, Schmalzing & Mukhanov 2000"],"Functions Text":["Topological estimators such as the Minkowski functionals (MFs) are also important diagnostics in this direction as they carry information at all orders. The MFs have been extensively developed as a statistical tool in a cosmological setting for both two-dimensional (2D; projected) and 3D (redshift) surveys. The MFs have analytically known results for a Gaussian random field, making them suitable for studies of non-Gaussianity. Examples of such studies include CMB data"],"Functions Label":["Background"],"Citation Start End":[[500,535]],"Functions Start End":[[0,472]]}
{"Identifier":"2015ApJ...799..149J___2013_Instance_1","Paragraph":"We use the microlensing magnification estimates for 27 quasar image pairs in 19 lens systems from MED09. In order to have the largest possible sample, but with a similar range of observed rest wavelengths, we include all of the objects from MED09 with magnifications measured in the wavelength range between Ly\u00ce\u00b1 (1216\u00e2\u0080\u0089\u00c3\u0085) and Mg\u00e2\u0080\u0089ii (2798\u00e2\u0080\u0089\u00c3\u0085). With this choice, the average rest wavelength is \u00ce\u00bb = 1736 \u00c2\u00b1 373\u00e2\u0080\u0089\u00c3\u0085, but we still keep 27 out of 29 image pairs from 19 out of 20 lensed quasars. Only the system RXS J1131\u00e2\u0088\u00921231 is excluded, as it was observed in [O\u00e2\u0080\u0089iii] at a much larger wavelength of \u00e2\u0088\u00bc5000\u00e2\u0080\u0089\u00c3\u0085. These microlensing magnification estimates are calculated after subtracting the emission line flux ratios, which are little affected by microlensing (see, e.g., Guerras et\u00c2 al. 2013), from the continuum flux ratios, and are therefore virtually free from extinction, substructure, and macro model effects (as these affect the line and continuum flux ratios equally). Our strategy is to compare the observed microlensing magnification for a given image pair  with a statistical sample of simulated values for that measurement as a function of the source size (rs) and the fraction of surface mass density in stars (\u00ce\u00b1). This will allow us to calculate the likelihood of the parameters (rs, \u00ce\u00b1) given the observations . The procedure is repeated for each of the 27 image pairs. We calculate magnification maps for each image using a grid with 11 values for the fraction of the surface mass density in stars, \u00ce\u00b1, logarithmically distributed between 0.025 and 0.8 as \u00ce\u00b1j = 0.025 \u00c3\u0097 2j\/2 with j = 0, \u00e2\u0080\u00a6, 10. The 517 magnification maps were created using the Inverse Polygon Mapping algorithm described by Mediavilla et\u00c2 al. (2006, 2011a). We used equal mass microlenses of 1\u00e2\u0080\u0089M. All of the linear sizes can be scaled for a different microlens mass as . The maps have a size of 2000 \u00c3\u0097 2000 pixels with a pixel size of 0.5 light-days. The maps therefore span 1000 lt-days. The individual sizes of maps and pixels in (more natural) units of Einstein radii for microlenses of 1\u00e2\u0080\u0089M are given in Table\u00c2 1. On average, the maps span approximately 50 Einstein radii with a pixel scale of roughly 0.025 Einstein radii.","Citation Text":["Guerras et\u00c2 al. 2013"],"Functions Text":["These microlensing magnification estimates are calculated after subtracting the emission line flux ratios, which are little affected by microlensing (see, e.g.,","from the continuum flux ratios, and are therefore virtually free from extinction, substructure, and macro model effects (as these affect the line and continuum flux ratios equally)."],"Functions Label":["Uses","Uses"],"Citation Start End":[[783,803]],"Functions Start End":[[622,782],[806,987]]}
{"Identifier":"2015MNRAS.448.2260G__Debes_et_al._2011_Instance_1","Paragraph":"Furthermore, well-defined large samples of white dwarfs are an extremely useful starting point for identifying rare white dwarf types like magnetic white dwarfs (G\u00e4nsicke, Euchner & Jordan 2002; Schmidt et al. 2003; K\u00fclebi et al. 2009; Kepler et al. 2013), pulsating white dwarfs (Castanheira et al. 2004; Greiss et al. 2014), high\/low-mass white dwarfs (Vennes & Kawka 2008; Brown et al. 2010; Hermes et al. 2014), white dwarfs with unresolved low-mass companions (Farihi, Becklin & Zuckerman 2005; Girven et al. 2011; Steele et al. 2013), white dwarfs with rare atmospheric composition (Schmidt et al. 1999; Dufour et al. 2010; G\u00e4nsicke et al. 2010), close white dwarf binaries (Marsh, Nelemans & Steeghs 2004; Parsons et al. 2011), metal polluted white dwarfs (Sion, Leckenby & Szkody 1990; Zuckerman & Reid 1998; Dufour et al. 2007; Koester, G\u00e4nsicke & Farihi 2014) or white dwarfs with dusty or gaseous planetary debris discs (G\u00e4nsicke et al. 2006; Farihi, Jura & Zuckerman 2009; Debes et al. 2011; Wilson et al. 2014). Because of their intrinsic low luminosities identifying a large, complete and well-defined sample of white dwarfs still remains a challenge. Much progress has been made in recent years thanks to large area surveys, first and foremost the Sloan Digital Sky Survey (SDSS; York et al. 2000; Harris et al. 2003; Eisenstein et al. 2006; Kleinman et al. 2013). The largest published catalogue of white dwarfs to date (Kleinman et al. 2013) fully exploited the spectroscopic data available at the time of SDSS Data Release 7 (DR7) and contains over 20 000 white dwarfs (of which 7424 with g \u2264 19). However not only is DR7 now outdated, but SDSS spectroscopy is only available for less than 0.01\u2009per\u2009cent of all SDSS photometric sources. Furthermore most of SDSS's white dwarfs are only serendipitous spectroscopic targets. The true potential of SDSS's vast multiband photometric coverage still remains to be fully mined for white dwarf research, but this requires a reliable method able to select white dwarfs candidates without recourse to spectroscopy. Proper motion has been traditionally used to distinguish white dwarfs from other stellar populations. In particular many studies that contributed to the census of white dwarfs in the solar neighbourhood specifically targeted high proper motion objects (Holberg et al. 2002; Sayres et al. 2012; Limoges, L\u00e9pine & Bergeron 2013). In this paper we present a novel method which makes use of photometric data and proper motions to calculate a probability of being a white dwarf (PWD) for any photometric source within a broad region in colour space. Unlike any previous similar work, our method does not use a specific cut in colour or proper motion to generate a list of white dwarf candidates; instead it provides a catalogue of sources with an associated PWD. These PWD can then be used to create samples of white dwarf candidates best suited for different specific uses. By applying our method to the full photometric footprint of SDSS Data Release 10 (DR10), we created a catalogue which includes \u223c23\u2009000 bright (g \u2264 19) high-fidelity white dwarfs candidates (Table 1). Using this catalogue, we assess the spectroscopic completeness of the SDSS white dwarf sample.","Citation Text":["Debes et al. 2011"],"Functions Text":["Furthermore, well-defined large samples of white dwarfs are an extremely useful starting point for identifying rare white dwarf types like","or white dwarfs with dusty or gaseous planetary debris discs"],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[985,1002]],"Functions Start End":[[0,138],[870,930]]}
{"Identifier":"2015ApJ...813..103M__Koss_et_al._2012_Instance_1","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"],"Functions Text":["Observations have shown that the AGN fraction does increase with decreasing distance between two merging galaxies","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."],"Functions Label":["Background","Motivation"],"Citation Start End":[[616,632]],"Functions Start End":[[480,593],[656,837]]}
{"Identifier":"2019AandA...625A.147A__Chantzos_et_al._2018_Instance_1","Paragraph":"Hydrocarbons also show a variety of deuterated molecules in L483. We detected C4D, with an isotopic ratio of 1.9 \u00b1 0.6%, which is slightly above that found in TMC-1 (0.43; Turner 1989) and similar to that derived in L1527 (1.8; Sakai et al. 2009b). We also detected the singly and doubly deuterated forms of c-C3H2, with isotopic ratios of 5.1 \u00b1 1.5% and 0.97 \u00b1 0.29%, respectively, which are similar to the values found in L1544 (Spezzano et al. 2013) and L1527 (Sakai et al. 2009b; Yoshida et al. 2019). Moreover, c-C3HD and c-C3D2 have been surveyed in a sample of low-mass prestellar and protostellar cores (Chantzos et al. 2018), finding that the corresponding isotopic ratios are relatively uniform, within a factor of a few, and similar to those in L1544 and L483. Thanks to the high sensitivity of our line survey, we detected the different deuterated forms of the two 13C substituted isotopologs of c-C3H2, that is, c-H13CCCD, c-HCC13CD, and c-HC13CCD; the latter only tentatively. This is to our knowledge the first time these species have been detected in space. The deuterium ratios derived for c-H13CCCD and c-HCC13CD are in line with that found for c-C3HD. We also detected the deuterated form of the linear isomer of C3H2, l-C3HD, which was recently observed for the first time toward TMC-1 and L1544 (Spezzano et al. 2016). The isotopic ratio we find in L483, 3.8 \u00b1 1.1%, is similar to the values derived in TMC-1 and L1544. Moreover, it seems that the linear isomer of C3H2 shows very similar levels of deuterium fractionation to the cyclic isomer in both low-mass prestellar and protostellar cores. The two deuterated forms of methyl acetylene, CH2DCCH and CH3CCD, which have previously been observed in TMC-1 with isotopic ratios of a few percent (Gerin et al. 1992a; Markwick et al. 2005), are also detected in L483 with slightly higher deuterium ratios (~6%). The deuterium ratio found for CH2DCCH in L483, 6.5 \u00b1 1.9%, is similar to that derived in L1527 (4.7%; Yoshida et al. 2019).","Citation Text":["Chantzos et al. 2018"],"Functions Text":["Moreover, c-C3HD and c-C3D2 have been surveyed in a sample of low-mass prestellar and protostellar cores","finding that the corresponding isotopic ratios are relatively uniform, within a factor of a few, and similar to those in L1544 and L483."],"Functions Label":["Background","Background"],"Citation Start End":[[612,632]],"Functions Start End":[[506,610],[635,771]]}
{"Identifier":"2021AandA...656A..94G__Gronow_et_al._2021_Instance_2","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"],"Functions Text":["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","for an explanation of the difference)."],"Functions Label":["Differences","Differences"],"Citation Start End":[[1816,1834]],"Functions Start End":[[1509,1815],[1835,1873]]}
{"Identifier":"2021AandA...649A..84H__Geers_et_al._2007_Instance_1","Paragraph":"The objective of this article is to study the spatial distribution and possible changes in the properties of carbon nano-dust in protoplanetary disks (PPDs). Carbon nanodust, detected under more or less organised structures and different ionisation states, constitutes a major component of dust in the interstellar and circumstellar environments. Vibrational emission bands in the near- to mid-IR from nanocarbon dust have been observed towards PPDs around most of the Herbig Ae stars, about half of the Herbig Be stars, and a few T-Tauri stars (e.g. Brooke et al. 1993; Acke & van den Ancker 2004; Acke et al. 2010; Seok & Li 2017). In contrast to large grains, these tiny and numerous carbon grains are well coupled to the gas and do not settle towards disk midplanes. This results in different spatial distributions in which tiny grains are present at the disk surfaces (e.g. Meeus et al. 2001; Habart et al. 2004; Lagage et al. 2006) and in the cavity or gaps from which the pebbles are missing (e.g. Geers et al. 2007; Kraus et al. 2013; Klarmann et al. 2017; Kluska et al. 2018; Maaskant et al. 2013). The very small carbon grains in the irradiated disk layers may have strong consequences (e.g. Gorti & Hollenbach 2008). As in the irradiated regions of the interstellar medium, they are the prevalent contributors to the energetic balance because they are very efficient at absorbing UV photons and heating the gas through the photoelectric effect. The highest fluxes of lines tracing the warm gas (e.g. [OI] 63 and [OI] 145 \u03bcm, H2 0-0 S(1), and high-J CO) are found in PPDs that show a large amount of flaring and high aromatic band strength (e.g. Meeus et al. 2013). Moreover, due to their large effective surface area, they may dominate the catalytic formation of key molecules as H2 and the charge balance. The disk structure may further depend on the level of nanograins that are coupled with the gas. Characterising the size and properties of these tiny grains through the disks, from internal to external regions, is thus of prime importance to understand the structure and evolution of PPDs.","Citation Text":["Geers et al. 2007"],"Functions Text":["This results in different spatial distributions in which tiny grains are present at the disk surfaces","and in the cavity or gaps from which the pebbles are missing (e.g."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1005,1022]],"Functions Start End":[[771,872],[938,1004]]}
{"Identifier":"2021MNRAS.504....1C__Ripepi_et_al._2014_Instance_1","Paragraph":"Our first attempt to derive the $PL_{K_\\mathrm{s}}$ relation has revealed a number of RRLs brighter than the main relation. We investigated their spatial distribution and found that they are mainly located in the central part of the LMC. This is shown in Fig. 5 which presents in the bottom panel the spatial distribution of the whole sample of RRLs considered in this paper (${\\sim}29\\, 000$), in the top panel the distribution of the RRLs which appear to be overluminous in the $PL_{K_\\mathrm{s}}$ relation, and, in the middle panel, the distribution of the RRLs which were actually used to fit our final $PL_{K_\\mathrm{s}}$ relation. Further investigations, performed using the Fourier parameters (\u03d531, \u03d521, R21, and R31) of the light curves available in the OGLE IV catalogue did not show any particular properties of the overluminous RRLs. On the other hand, in the period\u2013amplitude diagram based on the I amplitudes available in the OGLE IV catalogue, the bright RRLs all show small, in some cases near-zero amplitudes, compared with regular RRLs of the same period. The decrease in amplitude at a given period can be owing to these RRLs being blended with non-variable stars. We expect the centroid of a blended source to be determined with poor accuracy (see e.g. Ripepi et al. 2014, 2015). For this reason as a further test we plotted the distribution of distances in arcsec of the VMC sources cross-matched with the OGLE IV RRLs. A clear separation is now seen in the two samples, for 94 per cent of the RRLs lying on the PL\/PW relations the cross-match radius is less than 0.2 arcsec, whereas for 68 per cent of the overluminous RRLs the cross-match radius is larger than 0.2 arcsec. The average accuracy of the VMC astrometry is of the order of 0.080 arcsec both in RA and in Dec. (Cioni et al. 2011). We therefore discarded the RRLs with a cross-match radius larger than 0.2 arcsec. A total of 3252 objects were discarded. This procedure allowed us to significantly reduce the scatter on the PL\/PW relations. We visually inspected the VMC images of some of the discarded RRLs, confirming that they all are clearly blended with stars and\/or background galaxies. Similar investigations were performed in the past and the same effects were noted, e.g. by Ripepi et al. (2015) for Type II Cepheids. The final sample of LMC RRLs after this cleaning procedure contains 25\u2009795 objects. This is the sample that was used as a starting point to investigate the PL and PW relations presented in the following sections. Additional RRLs were later discarded from the final fit based on a 3\u03c3 clipping procedure leading to $PL_{K_\\mathrm{s}}$ relations using a clean sample of ${\\sim}22\\, 000$ RRLs.","Citation Text":["Ripepi et al. 2014"],"Functions Text":["We expect the centroid of a blended source to be determined with poor accuracy (see e.g.","For this reason as a further test we plotted the distribution of distances in arcsec of the VMC sources cross-matched with the OGLE IV RRLs."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1272,1290]],"Functions Start End":[[1183,1271],[1299,1439]]}
{"Identifier":"2021MNRAS.506.1715D__Jones_et_al._2003_Instance_1","Paragraph":"In direct relation to the interplay between the structure of the dark matter distribution and the baryon physics, galaxies are found in a wide range of structural hierarchies, from low-density regions to groups and clusters (see e.g. Tully 1987; Berlind et al. 2006; Yang et al. 2007), and during their lifetime they experience merging events (e.g. Mamon 1988; Tempel et al. 2017). In some cases, the mergers eventually devoid their entire neighbourhood, leaving behind a single elliptical galaxy of group-scale mass, called a fossil group galaxy (Ponman et al. 1994; Jones et al. 2003). Numerical simulations by Barnes (1989) first motivated such a hierarchical merging scenario (see also D\u00edaz-Gim\u00e9nez, Muriel & Mendes de Oliveira 2008). Since then, there have been several supporting reports of X-ray sources identified as fossil groups (Santos, Mendes de Oliveira & Sodr\u00e9 2007; La Barbera et al. 2009). Because most fossil systems found to date lie within z  0.2, fossil galaxy groups most likely are old, undisturbed systems due to the lack of major mergers. While some luminous galaxies experience major merger events in their evolution, fossil group galaxies acquire their mass typically through minor merger events, where the mass ratio stays below 0.3 (DOnghia et al. 2005). Simulations show that a fossil system may assemble half of its mass in dark matter by redshift z > 1, and that the assembled mass at any redshift is generally higher in a fossil than in regular groups (Dariush et al. 2007). Since this merging process is relatively fast compared to the cooling time of the surrounding gas, comparable to one to several Hubble times, fossil groups are usually found embedded in giant, X-ray luminous gas haloes (Mulchaey 2000). If a fossil system has not yet fully merged, it can be identified by another criterion, a gap in brightness of at least 2.5 mag (usually defined in the r band) between the first and fourth brightest galaxies in the group (Dariush et al. 2010; Zarattini et al. 2014) within half its virial radius.","Citation Text":["Jones et al. 2003"],"Functions Text":["In some cases, the mergers eventually devoid their entire neighbourhood, leaving behind a single elliptical galaxy of group-scale mass, called a fossil group galaxy"],"Functions Label":["Background"],"Citation Start End":[[568,585]],"Functions Start End":[[382,546]]}
{"Identifier":"2016AandA...596A.116S__Bochanski_et_al._2010_Instance_1","Paragraph":"At present, however, kinematical studies of the nearby late-type stars have become somewhat foreshadowed with the advent of deep, wide-field SDSS (York et al. 2000), SDSS\/SEGUE (Yanny et al. 2009) and RAVE (Steinmetz et al. 2006) surveys which have made it possible to study K\u2013M dwarfs to much greater distances from the Sun. Using their spectroscopic catalog with radial velocities measured with an external accuracy of 7\u201310 km\u2009s-1 and applying photometric parallax relations, the SDSS teams have traced K\u2013M dwarfs in the distance range up to ~2 kpc, thus providing valuable information on the spatial, velocity, and metallicity distributions with respect to vertical distance from the Galactic plane (Bochanski et al. 2007, 2011; Juri\u0107 et al. 2008; Fuchs et al. 2009; Bond et al. 2010; West et al. 2011; Schlesinger et al. 2012; Zhang et al. 2013), as well as on the luminosity and mass functions of low-mass dwarfs in the Galactic disk (Covey et al. 2008; Bochanski et al. 2010). Relatively local (| z |   500 pc) samples of late-type dwarfs from the RAVE survey, based on radial velocities accurate to ~2 km\u2009s-1 and photometrically determined distances, have been used to deduce the solar space velocity with respect to the Local Standard of Rest (Co\u015fkuno\u01e7lu et al. 2011; Pasetto et al. 2012a; Golubov et al. 2013). The RAVE survey has also resulted in a catalog of ~44\u2009000 candidate active stars (\u017derjal et al. 2013) which makes a major contribution to the data on chromospheric emission of cool dwarfs. In the context of these massive surveys, as well as of present-day models of the Galaxy, the samples of nearby late-type dwarfs with most accurate trigonometric parallax and radial-velocity measurements, as well as other high-quality observational data, still remain important, as they provide a fundamental framework for calibration and tests of relations between low-mass star parameters (such as, e.g., color-luminosity, mass-luminosity, activity-age, chemo-kinematic relations). ","Citation Text":["Bochanski et al. 2010"],"Functions Text":["Using their spectroscopic catalog with radial velocities measured with an external accuracy of 7\u201310 km\u2009s-1 and applying photometric parallax relations, the SDSS teams have traced K\u2013M dwarfs in the distance range up to ~2 kpc, thus providing valuable information on the spatial, velocity, and metallicity distributions with respect to vertical distance from the Galactic plane","as well as on the luminosity and mass functions of low-mass dwarfs in the Galactic disk"],"Functions Label":["Background","Background"],"Citation Start End":[[959,980]],"Functions Start End":[[326,701],[851,938]]}
{"Identifier":"2019ApJ...874..154D__Stanway_et_al._2016_Instance_1","Paragraph":"There are several reasons to suspect that the (galaxy-sourced) ionizing background during reionization may have been somewhat harder than the estimates given here. First, our calculations neglect the filtering effects of optically thick H i in the ISM of the host galaxy, and within the cosmic web. Absorption by this gas would have hardened the spectrum of the ionizing radiation as it escaped the galaxy and traveled through the IGM (e.g., Madau 1995; Faucher-Gigu\u00e8re et al. 2009; Haardt & Madau 2012). These effects were likely strongest during the last stages of reionization, when the radiation typically had to travel large distances to reach the I-fronts. Second, our calculations neglect the effects of binary star systems. Mass transfers and mergers between binary companions can extend the period over which ionizing photons are produced by the stellar population, which would harden the time-integrated spectrum (Eldridge & Stanway 2009; Stanway et al. 2016). Lastly, recent studies have suggested that the IMF in starburst galaxies may be more top-heavy than the IMF assumed here (Baugh et al. 2005; Gunawardhana et al. 2011; Marks et al. 2012; Zhang et al. 2018). Most recently, Schneider et al. (2018) measured a logarithmic slope of \n\n\n\n\n\n in the mass range 15\u2013200 M\u2299, using spectroscopic measurements of the 30 Doradus star-forming region in the Large Magellanic Cloud. (The IMF adopted here has a slope of 2.3 for M > M\u2299, and a cutoff of 120 M\u2299.) Each of the provided effects would work in the direction of making \n\n\n\n\n\n smaller. Based on these considerations, we argue that the lower half of Figure 2, with \n\n\n\n\n\n \u2272 1.5, is likely the most relevant region of parameter space for \n\n\n\n\n\n. In what follows, we adopt \n\n\n\n\n\n = 1.5 as our fiducial value, but we note that \n\n\n\n\n\n is only mildly sensitive to \n\n\n\n\n\n except at the fastest I-front speeds. In the next section we find that \n\n\n\n\n\n = 104 km s\u22121 is close to the upper limit achieved by I-fronts in cosmological simulations, which yields \n\n\n\n\n\n = 26,200 K, assuming \n\n\n\n\n\n = 1.5 (see Figure 2). This result varies by \u0394\n\n\n\n\n\n = \u22124200 (+1600) K if we instead assume \n\n\n\n\n\n = 2.5(0.5).","Citation Text":["Stanway et al. 2016"],"Functions Text":["Second, our calculations neglect the effects of binary star systems. Mass transfers and mergers between binary companions can extend the period over which ionizing photons are produced by the stellar population, which would harden the time-integrated spectrum"],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[949,968]],"Functions Start End":[[663,922]]}
{"Identifier":"2021MNRAS.505..523R__Mo\u015bcibrodzka_et_al._2017_Instance_1","Paragraph":"Recently, the EHT Collaboration completed an analysis of the linear polarization of M87*, providing novel insights into its magnetic field structure in particular (Event Horizon Telescope Collaboration 2021a,b; Goddi et al. 2021). Synchrotron emission, which dominates the millimetre image, is initially generated perpendicular to the local magnetic field, such that its polarization carries with it an imprint of the field geometry (Palumbo, Wong & Prather 2020). Then, as this polarization propagates, it is further modified by Faraday effects, important for depolarizing accretion flows down to observed levels and generating the observed rotation measure (Ballantyne, \u00d6zel & Psaltis 2007; Mo\u015bcibrodzka et al. 2017; Jim\u00e9nez-Rosales & Dexter 2018; Ricarte et al. 2020). Fully polarized radiative transport simulations on EHT scales have been developed in the past few decades, allowing us to link polarized images to the detailed physics of the underlying plasma and the space\u2013time producing them (Bromley, Melia & Liu 2001; Broderick & Loeb 2006; Broderick & McKinney 2010; Porth et al. 2011; Shcherbakov, Penna & McKinney 2012; Dexter 2016; Mo\u015bcibrodzka & Gammie 2018). This enabled (Event Horizon Telescope Collaboration 2021b) to discriminate between two major classes of accretion disc: a \u2018Magnetically Arrested Disk\u2019 (MAD) and \u2018Standard and Normal Evolution\u2019 (SANE). MAD accretion discs have magnetic fields strong enough to affect the disc dynamics and exhibit stronger poloidal (or non-toroidal) magnetic field components (Bisnovatyi-Kogan & Ruzmaikin 1974; Igumenshchev, Narayan & Abramowicz 2003; Narayan, Igumenshchev & Abramowicz 2003; Chael, Narayan & Johnson 2019). Meanwhile, the weaker magnetic fields of a SANE disc are sheared out by the motion of the plasma into a mostly toroidal configuration (Narayan et al. 2012; S\u0105dowski et al. 2013; Ryan et al. 2018). The fractional linear polarization of M87*, the upper limit on its circular polarization, and most importantly the \u2018twisty pattern\u2019 of its spatially resolved linear polarization map favour \u2018MAD\u2019 models of M87* (Event Horizon Telescope Collaboration 2021b).","Citation Text":["Mo\u015bcibrodzka et al. 2017"],"Functions Text":["Then, as this polarization propagates, it is further modified by Faraday effects, important for depolarizing accretion flows down to observed levels and generating the observed rotation measure"],"Functions Label":["Background"],"Citation Start End":[[693,717]],"Functions Start End":[[465,658]]}
{"Identifier":"2019MNRAS.486.4671M__Schwenn_2006_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":["Schwenn 2006"],"Functions Text":["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"],"Functions Label":["Background"],"Citation Start End":[[706,718]],"Functions Start End":[[444,610]]}
{"Identifier":"2021AandA...645A.137A__Gopal-Krishna_et_al._(2011)_Instance_1","Paragraph":"The first INM of PG 1553+113 to our knowledge was reported by Stalin et al. (2005) \u2013 they applied the C-test to the R-band LCs (total duration of about 10 h) and found the source to be variable in one night out of two. Another two nights of INM about 13 h long were presented by Osterman et al. (2006). They found no significant INV, but they did not present statistical tests. To quantify the conclusion of Osterman et al. (2006), we used the C-test in the form C\u2004=\u2004\u03c3\/\u27e8e\u27e9, where \u03c3 is the standard deviation of the source LC and \u27e8e\u27e9 the mean uncertainty of the source photometric measurements. We applied this test to the data listed in Table 4 of Osterman et al. (2006) and found C\u2004=\u20041.52 and C\u2004=\u20041.28 for the respective nights, which means that PG 1553+113 was non-variable. Andruchow et al. (2007, 2011) reported the results from the INM they carried out during four nights in the VR-bands (April 21 to 24, 2007) and five nights in the BR-bands (April 21 to 25, 2009), respectively. Applying the C-test, the authors found no significant variability. Gopal-Krishna et al. (2011) detected INV during three nights out of three using C- and F-tests (a total monitoring duration of 16 h in the R-band). Gaur et al. (2012) found no INV during six nights of monitoring in the BR-bands (C- and F-tests were applied). Gupta et al. (2016) presented the results from the INM during seven nights in the R-band (a total monitoring duration of about 26 h). The authors found the source to be variable in one night, non-variable in another, and probably variable in the remaining nights (F- and \u03c72 tests were applied). We, however, should point out that in the latter nights the variability amplitude seems to be quite close to the magnitude uncertainties. In such cases the usage of the C-test could be more appropriate (see Zibecchi et al. 2017, 2020). Pandey et al. (2019) monitored PG 1553+113 for eight nights in the VR-bands for 2\u20134 h each night. Employing enhanced F- and nested ANOVA tests, the authors found the source to vary on intra-night timescales for three nights. Finally, Pasierb et al. (2020) found no INV in the BVR-bands during a single night of monitoring (duration of about 3.7 h, F-test employed).","Citation Text":["Gopal-Krishna et al. (2011)"],"Functions Text":["detected INV during three nights out of three using C- and F-tests (a total monitoring duration of 16 h in the R-band)."],"Functions Label":["Compare\/Contrast"],"Citation Start End":[[1053,1080]],"Functions Start End":[[1081,1200]]}
{"Identifier":"2016MNRAS.462S.376B__Bowell_et_al._1989_Instance_1","Paragraph":"In order to assess the impact of the brightness enhancement due to the phase function effects in the light scattering of the coma dust, we apply different phase function corrections to the measured coma brightness of 67P for the 10\u2009000 km aperture radius. In general, the phase function reduces the reflected coma magnitude value. For small phase angles \u03b1 (typically up to 30\u00b0), a linear reduction with \u03b1 can be applied with typical parameters \u03b2 for comets between 0.02 and 0.06 (Lamy et al. 2004). Alternatively, the two-parameter phase function of Schleicher, Millis & Birch (1998), based on 1P\/Halley measurements, is available or the phase angle correction by M\u00fcller (1999) that also considers geometric projection effects by a non-spherical coma, symmetric to the radial direction with respect to the Sun. Fornasier et al. (2015) obtained results for the 67P nucleus surface reflectivity applying a HG-type phase function (Bowell et al. 1989) with a best-fitting G of \u22120.13. Fig. 5 shows the phase angle-corrected af\u03c1 values of 67P for the various phase function solutions. In Fig. 5, \u03b2 = 0.04 is used as mean value for the linear phase function that is lower than the \u03b2 value found for the nucleus of 67P (0.059\u20130.076; Tubiana et al. 2011; Lowry et al. 2012). A second HG-type phase function is applied as well with G = 0.15 that is found from many asteroid light curves (Tedesco 1989). Obviously, the af\u03c1 value for zero phase angle depends on the phase darkening and is highest for the Bowell-type phase function with G = \u22120.13 (as obtained for the 67P nucleus by Fornasier et al. 2015). Two phase functions provide a close to linear decrease of af\u03c1 with solar distance over the time interval from the end of 2015 September to the beginning of 2016 April, i.e. the linear phase function with \u03b2 = 0.04 and the Bowell-HG-type phase function with G = 0.15. The phase functions of Schleicher et al. (1998) and of M\u00fcller (1999) can approximate the af\u03c1 data well between the end of September and end of December, but do not describe well the 67P dust activity measured in 2016.","Citation Text":["Bowell et al. 1989"],"Functions Text":["Fornasier et al. (2015) obtained results for the 67P nucleus surface reflectivity applying a HG-type phase function","with a best-fitting G of \u22120.13. Fig. 5 shows the phase angle-corrected af\u03c1 values of 67P for the various phase function solutions."],"Functions Label":["Uses","Uses"],"Citation Start End":[[928,946]],"Functions Start End":[[811,926],[948,1078]]}
{"Identifier":"2015ApJ...805..134M__Yoo_et_al._2014_Instance_1","Paragraph":"The physics of asymmetric inflow reconnection has been investigated in detail for fully ionized plasmas. One of the principal applications of this work has been Earth's dayside magnetopause. Asymmetric inflow reconnection has also been investigated in the context of Earth's magnetotail and elsewhere in the magnetosphere (Oieroset et al. 2004; Muzamil et al. 2014), the solar atmosphere (Murphy et al. 2012; Nakamura et al. 2012; Su & van Ballegooijen 2013; Su et al. 2013), laboratory experiments (Yamada 2007; Murphy & Sovinec 2008; Yoo et al. 2014), and plasma turbulence (Servidio et al. 2009, 2010). Cassak & Shay (2007) performed a scaling analysis for asymmetric inflow reconnection. They found that the outflow is governed by a hybrid upstream Alfv\u00e9n speed that is a function of the magnetic field strength and density in both upstream regions (see also Birn et al. 2010) and that the flow stagnation point in the simulation frame and magnetic field null were not colocated. The structure and dynamics of asymmetric reconnection in fully ionized collisionless and two-fluid plasmas have been studied previously in several works (e.g., Swisdak et al. 2003; Cassak & Shay 2008, 2009; Mozer et al. 2008; Murphy & Sovinec 2008; Pritchett 2008; Mozer & Pritchett 2009; Pritchett & Mozer 2009; Malakit et al. 2010, 2013; Aunai et al. 2013a, 2013b). Simulations of the plasmoid instability during reconnection with asymmetric upstream magnetic fields by Murphy et al. (2013) showed that the resultant magnetic islands developed primarily into the weak-field upstream region. Because the reconnection jets impacted the islands obliquely rather than directly, the islands developed net vorticity. In addition to asymmetric inflow reconnection, several groups have investigated asymmetric outflow reconnection (e.g., Oka et al. 2008; Murphy 2010; Murphy et al. 2010) and reconnection with three-dimensional asymmetry (e.g., Al-Hachami & Pontin 2010; Wyper & Jain 2013).","Citation Text":["Yoo et al. 2014"],"Functions Text":["Asymmetric inflow reconnection has also been investigated in","laboratory experiments"],"Functions Label":["Background","Background"],"Citation Start End":[[536,551]],"Functions Start End":[[191,251],[476,498]]}
{"Identifier":"2021ApJ...910...18C__Rubele_et_al._2018_Instance_1","Paragraph":"The LMC is known to have brought a large population of star clusters as it has been accreted onto the Milky Way (Bica et al. 2008); thus, it is possible that DELVE 2 may share a similar history to the thousands of known star clusters in the Magellanic system. The LMC star cluster formation history is believed to be three-staged, including a period of rapid cluster star formation in the early universe (\u03c4 \u2273 10 Gyr), followed by a long quiescent period between \u223c10 and \u223c2\u20134 Gyr ago and then by a period of rapid star cluster formation extending to the present day, potentially due to the interaction between the LMC and SMC (Harris & Zaritsky 2009; Weisz et al. 2013; Rubele et al. 2018; Ruiz-Lara et al. 2020). One consequence of this period of quiescence in the LMC cluster formation history is the so-called \u201cage gap\u201d in the age distribution of LMC clusters, with a small (N \u2272 20) population of globular clusters with ages comparable to most known Milky Way globular clusters, separated by the gap from a much larger population of less massive young clusters (e.g., Bertelli et al. 1992; Girardi et al. 1995; Olszewski et al. 1996). These two populations of clusters obey an overarching age\u2013metallicity relation, within which the older clusters (\u03c4 > 12 Gyr) are significantly more metal-poor (\u22122.2 \u2272 [Fe\/H] \u2272 \u22121.2) compared to the younger population of clusters ([Fe\/H] \u2273 \u22120.7; Meschin et al. 2014).38\n\n38\nWe note that Gatto et al. (2020) recently discovered 16 cluster candidates believed to be within the LMC cluster age gap (4 Gyr \u2272 \u03c4 \u2272 10 Gyr).\n Therefore, although the photometrically derived metallicity and age for DELVE 2 are limited in accuracy by the small number of red giant branch stars available to precisely constrain these properties through synthetic isochrone fitting, it is clear that DELVE 2 is more consistent with an old, metal-poor stellar population and thus the former class of LMC clusters (provided the system is not a dwarf galaxy, as discussed in the previous subsection).","Citation Text":["Rubele et al. 2018"],"Functions Text":["The LMC star cluster formation history is believed to be three-staged, including a period of rapid cluster star formation in the early universe (\u03c4 \u2273 10 Gyr), followed by a long quiescent period between \u223c10 and \u223c2\u20134 Gyr ago and then by a period of rapid star cluster formation extending to the present day, potentially due to the interaction between the LMC and SMC"],"Functions Label":["Background"],"Citation Start End":[[669,687]],"Functions Start End":[[260,624]]}
{"Identifier":"2017ApJ...834L..21A__Larsen_&_Lane_1994_Instance_1","Paragraph":"The experiment was carried out using the VULCAN laser facility at the Rutherford Appleton Laboratory (Danson et al. 1998), using the setup shown in Figure 1(a). A laser pulse of duration \u223c1 ns and energy \u223c 70 J, was focused onto a 50 \u03bcm thick gold foil, at an incidence angle of \n\n\n\n\n\u223c\n45\n\u00b0\n\n\n to a peak intensity of \u223c1015 W cm\u22122. The interaction of the nanosecond pulse with the gold foil leads to the generation of ablated plasma that is mainly constituted of thermally distributed gold ions and faster lighter ions with an average energy per nucleon of tens of keV (Tan et al. 1984; Gitomer et al. 1986). The lighter ions, such as protons and carbon ions, originate from contaminant layers (water vapor and hydrocarbon) typically present on the surface of the targets (Gitomer et al. 1986). The laser\u2013target interaction was enclosed in a gas cell filled with pure nitrogen at a controlled pressure of \u223c10\u22121 mbar. Hydrodynamic simulations using the code HYADES  (Larsen & Lane 1994) indicate that the gas becomes fully ionized within 100 ps from the start of the laser irradiation by the X-rays emitted from the target (Dean et al. 1971), resulting in a stationary plasma with an electron density and temperature of \u223c3 \u00d7 1016 cm\u22123 and \n\n\n\n\n\n\nT\n\n\ne\n\n\n\u223c\n1\n\n\n  keV, respectively. These values imply an electron Debye length of \n\n\n\n\n\n\n\u03bb\n\n\nD\n\n\n\u223c\n1.4\n\n\u03bc\nm\n\n\n and an ion-acoustic speed of \n\n\n\n\n\n\nC\n\n\ns\n\n\n\u223c\n2.2\n\u00d7\n\n\n10\n\n\n5\n\n\n\nm\n\n\n\ns\n\n\n\u2212\n1\n\n\n\n\n. Moreover, the Coulomb logarithm for electron\u2013electron and ion\u2013ion collisions are of the order of 6 and 11, respectively, indicating a characteristic timescale for collisions of \n\n\n\n\n\n\n\u03c4\n\n\nee\n\n\n\u223c\n36\n\nns\n\n\n and \n\n\n\n\n\n\n\u03c4\n\n\nii\n\n\n\u223c\n600\n\nns\n\n\n, respectively. The PPI technique (Borghesi et al. 2002; Sarri et al. 2010) was employed to investigate the interaction of the ablated plasma with the background plasma. The probe proton beam was generated by focusing a second laser pulse of \u223c1 ps duration and \u223c50 J energy, to an intensity of \n\n\n\n\n\u223c\n\n\n10\n\n\n19\n\n\n\n\n W cm\u22122 onto a thin gold foil (thickness \u223c20 \u03bcm). In the experimental arrangement shown in the Figure 1(a), the distance between proton source and interaction region was \n\n\n\n\nl\n\u2243\n4\n\nmm\n\n\n, and the detector was placed at \n\n\n\n\nL\n\u2243\n38\n\nmm\n\n\n from the interaction region, giving an intrinsic geometrical magnification of \n\n\n\n\nM\n\u2248\n(\nl\n+\nL\n)\n\n\/\n\nl\n\u223c\n10.5\n\n\n (Borghesi et al. 2002). A stack of several layers of dosimetrically calibrated (Kirby et al. 2011) RadioChromic Films (RCF), was used to detect the proton beam. The two lasers were temporally delayed so that a proton with an energy of 13 MeV traverses the interaction region at the time \n\n\n\n\n\n\nt\n\n\n0\n\n\n=\n180\n\u00b1\n20\n\n\n ps after the start of the long-pulse irradiation, where the start of the interaction corresponds to 1\/10 of the peak intensity. The error in defining the beginning of the interaction is mainly due to the systematic error in the synchronization of both laser beams. This does not affect the temporal resolution of the PPI technique, which is on the order of picoseconds (Sarri et al. 2010).","Citation Text":["Larsen & Lane 1994"],"Functions Text":["Hydrodynamic simulations using the code HYADES"],"Functions Label":["Uses"],"Citation Start End":[[965,983]],"Functions Start End":[[916,962]]}
{"Identifier":"2015AandA...580A..71L__Sutton_et_al._(2013)_Instance_1","Paragraph":"The simplest two component model (power law + disk) is a phenomenological model often used to describe the spectra of ULXs as an empirical description of a disk plus corona geometry. In the presence of a cool (kT ~ 0.1\u22120.4 keV) and luminous (L ~ 1039\u22121040 erg\/s) disk, it allows inferring the presence of intermediate-mass black holes (e.g., Makishima et al. 2000). This is not the case for M33 X-8, where the disk component describes the high-energy part of the spectrum well and appears to be hot (kT ~ 1.15 keV), leaving a soft excess that is accounted for by the power law. The overall disk parameters are then inconsistent with a massive black hole, but instead are more typical of an ordinary stellar mass black hole: using the relationship between mass, temperature, and luminosity in a standard disk (see, e.g., Makishima et al. 2000), we derive a mass of ~10 M\u2299 for a nonrotating black hole, consistent with the estimation obtained by data from other satellites (e.g., Foschini et al. 2006; Weng et al. 2009; Isobe et al. 2012). Sutton et al. (2013) developed a classification scheme based on a disk+power law fit, to be applied to ULX spectra, according to which the spectral state of an ULX source can be defined by the disk temperature, the power-law slope, and the ratio between the flux contribution of the two spectral components in the 0.3\u22121 keV band. Our result is consistent with that found by Sutton et al. (2013) using XMM-Newton data, and, according to their classification, it identifies M33 X-8 as a broadened disk source, or in other words, as a source whose spectrum is dominated by emission from a hot disk (see Table 2) and where the additional soft component may be the effect of an unrealistic description of the disk spectrum by the diskbb model. In fact, such hot-disk\/soft power-law spectra are difficult to explain in the context of the analogy of ULXs with GBHs: the thermal state of GBHs is indeed characterized by a hot disk, but the presence of a soft power-law-like component in addition to the disk is unusual, and its physical interpretation is not simple: if this component is due to the presence of a Comptonized corona, we do not expect it to be dominant at energies lower than the temperature of the seed photons that come from the disk. ","Citation Text":["Sutton et al. (2013)"],"Functions Text":["developed a classification scheme based on a disk+power law fit, to be applied to ULX spectra, according to which the spectral state of an ULX source can be defined by the disk temperature, the power-law slope, and the ratio between the flux contribution of the two spectral components in the 0.3\u22121 keV band."],"Functions Label":["Background"],"Citation Start End":[[1038,1058]],"Functions Start End":[[1059,1367]]}
{"Identifier":"2022ApJ...924...97W__Zhang_&_M\u00e9sz\u00e1ros_2001_Instance_1","Paragraph":"Using all GRBs showing X-ray plateau phases, Dainotti et al. (2008) discovered a tight correlation between L0 and tb (the Dainotti relation). Subsequently, the Dainotti relation has been used to measure cosmological parameters (Cardone et al. 2009, 2010; Dainotti et al. 2013; Postnikov et al. 2014). Cardone et al. (2010) relied on this correlation to build a Hubble diagram of 66 GRBs and present a preliminary constraint on cosmological parameters. The errors on the cosmological constraints are large, which may be related to the nature of GRB light curves and in part due to the sample selection. In previous works, GRBs with X-ray plateaus were used to derive the Dainotti relation, and then to constrain cosmological parameters. However, a newborn magnetar can be spun down through a combination of electromagnetic dipole and gravitational-wave quadrupole emission (Shapiro & Teukolsky 1983). The X-ray luminosity of GRBs is given by the energy input from the electromagnetic and gravitational waves into the surrounding medium (Dai & Lu 1998; Zhang & M\u00e9sz\u00e1ros 2001; Metzger et al. 2011). Therefore, in order to standardize GRBs as standard candles through the Dainotti relation, it is recommended that X-ray plateaus caused by the same physical mechanism are used (electromagnetic dipole radiation or gravitational waves). Similar to supernova cosmology, only SNe Ia from accretion channels can be treated as standard candles. Following the Dainotti correlation (Dainotti et al. 2008) between the plateau luminosity and the end time of the plateau in X-ray afterglows, we further confirm this relation out to redshift z = 5.91. In this paper, we perform a first attempt to standardize long GRBs with X-ray plateaus dominated by electromagnetic dipole radiation as reliable standard candles. Using the same method, Hu et al. (2021) considered two other kinds of GRBs to constrain cosmological parameters, i.e., short GRBs with plateau phases dominated by magnetic dipole radiations and long GRBs with gravitational-wave-dominated plateau phases. It is interesting to note that this correlation also holds for the optical sample (Dainotti et al. 2020b).","Citation Text":["Zhang & M\u00e9sz\u00e1ros 2001"],"Functions Text":["The X-ray luminosity of GRBs is given by the energy input from the electromagnetic and gravitational waves into the surrounding medium","Therefore, in order to standardize GRBs as standard candles through the Dainotti relation, it is recommended that X-ray plateaus caused by the same physical mechanism are used (electromagnetic dipole radiation or gravitational waves)."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1051,1072]],"Functions Start End":[[900,1034],[1096,1330]]}
{"Identifier":"2021AandA...653A..85S__However,_Genovali_et_al._(2014)_Instance_2","Paragraph":"Figure 6 shows the orbital eccentricities as a function of [M\/H] for the metal-rich disc sample. The solid lines correspond to the required eccentricity (see Eq. (2)) for different values of ISM radial metallicity gradients: \u22120.10 dex kpc\u22121 (black), \u22120.07 dex kpc\u22121 (our measured gradient for young stars in Table 1; see also Minchev et al. 2018, red), \u22120.04 dex kpc\u22121 (orange), and \u22120.06 dex kpc\u22121 (Cepheids analysis from Genovali et al. 2014, green). For the three first cases, we assumed ISM[M\/H](R\u2299) = 0.0 to estimate Rbirth from the stellar metallicity. However, Genovali et al. (2014) have their own zero point, defined as: [Fe\/H]\u2004=\u2004\u22120.06\u2005*\u2005Rg\u2005+\u20050.57, with a clear shift in the relation compared to the other ones assumed in this work. The impact of the ISM gradient value and the zero-point assumption on the derived Rbirth, and therefore on the required eccentricity to reach the solar vicinity without the need for churning, is clearly observed. As described in Hayden et al. (2020), given the measured [M\/H] and eccentricity, stars lying to the left are able to reach the solar neighbourhood through blurring, while the stars to the right of the line are possible candidates to have migrated through churning. This is the case for most of the SMR stars (70% of the SMR stars lie below the line that corresponds to the Cepheids analysis); they are therefore likely to have been brought to the solar neighbourhood by churning, which is in close agreement with previous studies (e.g., Kordopatis et al. 2015a; Wojno et al. 2016). However, it is worth noting that the observed metallicity distribution function in Fig. 2 peaks around 0.2 dex, which is higher than previous reported solar vicinity MDFs (see e.g., Fuhrmann et al. 2017). A possible ignored bias towards more metal-rich objects in the sample selection could be pulling the percentage of possible migrators to higher values. Among the entire distribution, our churned candidates with [M\/H] > \u2005+\u20050.1 comprise around 17% of the sample. If we constrain the number of migrators to only stars with [M\/H] > \u2005+\u20050.25, the global percentage decreases to 8% of the sample.","Citation Text":["Genovali et al. (2014)"],"Functions Text":["For the three first cases, we assumed ISM[M\/H](R\u2299) = 0.0 to estimate Rbirth from the stellar metallicity. However,","have their own zero point, defined as: [Fe\/H]\u2004=\u2004\u22120.06\u2005*\u2005Rg\u2005+\u20050.57, with a clear shift in the relation compared to the other ones assumed in this work."],"Functions Label":["Differences","Differences"],"Citation Start End":[[568,590]],"Functions Start End":[[453,567],[591,741]]}
{"Identifier":"2015ApJ...799...55G__Klassen_et_al._2000_Instance_2","Paragraph":"While the angular extent of IP shocks can be directly investigated using multi-point in situ measurements, the size of coronal shocks can only be indirectly inferred via remote-sensing observations of the electromagnetic emissions associated with them. According to Nelson & Robinson (1975), the average angle subtended at the solar surface by fundamental metric type II radio emission sources is 43\u00c2\u00b0. Aurass et al. (1994) found particular cases with larger, double type II source structures covering a separation angle beyond 90\u00c2\u00b0. Type II radio sources often show non-radial propagation trajectories (see Mann et al. 2003, and references therein). Wave-like large-scale disturbances propagating over the solar disk in extreme ultraviolet observations (usually referred to as \u00e2\u0080\u009cEIT waves\u00e2\u0080\u009d or \u00e2\u0080\u009cEUV waves\u00e2\u0080\u009d) are in close empirical correlation with type II radio bursts (Klassen et al. 2000). Most EIT waves are accompanied by CMEs, and observations and MHD modeling suggest that they are driven by the lateral expansion of CMEs, while the ultimate nature of the phenomenon remains under debate and could consist of true waves, pseudo waves (e.g., reconnection fronts), or a combination of both (Patsourakos & Vourlidas 2012; Nitta et al. 2013b, and references therein). According to Patsourakos & Vourlidas (2012), EIT waves can reach distances up to 1.3 R (850 Mm) from the source. Single-case studies reported some EIT waves covering a whole solar hemisphere (Klassen et al. 2000; Kienreich et al. 2009; Thompson & Myers 2009). Connections between EIT waves and SEP events have been often suggested (e.g., Bothmer et al. 1997; Krucker et al. 1999), and recently Rouillard et al. (2012) hypothesized that the EIT wave can be used to track the expansion of a coronal shock responsible for particle acceleration. Other authors question the EIT wave acceleration scenario for SEPs, with many EIT waves being observed at well-connected positions having no associated SEP increase (Miteva et al. 2014).","Citation Text":["Klassen et al. 2000"],"Functions Text":["Single-case studies reported some EIT waves covering a whole solar hemisphere"],"Functions Label":["Background"],"Citation Start End":[[1468,1487]],"Functions Start End":[[1389,1466]]}
{"Identifier":"2021AandA...653A..36M__Goulding_&_Alexander_(2009)_Instance_2","Paragraph":"The SFG sample was constructed using the Great Observatories All-Sky LIRG Survey (GOALS sample, Armus et al. 2009), from which we extracted 158 galaxies, with data from Inami et al. (2013), who report the fine-structure lines at high resolution in the 10\u2005\u2212\u200536\u2006\u03bcm interval, and Stierwalt et al. (2014), who include the detections of the H2 molecular lines and the PAH features at low spectral resolution. For those galaxies in the GOALS sample that have a single IRAS counterpart, but more than one source detected in the emission lines, we have added together the line or feature fluxes of all components, to consistently associate the correct line or feature emission to the total IR luminosity computed from the IRAS fluxes. To also cover lower luminosity galaxies, as the GOALS sample only includes luminous IR galaxies (LIRGs) and ultra-luminous IR galaxies (ULIRGs), we included 38 galaxies from Bernard-Salas et al. (2009) and Goulding & Alexander (2009), to reach the total sample of 196 galaxies with IR line fluxes in the 5.5\u2005\u2212\u200535\u2006\u03bcm interval in which an AGN component is not detected. For the Bernard-Salas et al. (2009), Goulding & Alexander (2009), and the GOALS samples, we excluded all the composite starburst-AGN objects identified as those with a detection of [NeV] either at 14.3 or 24.3 \u03bcm. It is worth noting that the original samples from Goulding & Alexander (2009) and Bernard-Salas et al. (2009) have spectra solely covering the central region of the galaxies. To estimate the global SFR, we corrected the published line fluxes of the Spitzer spectra by multiplying them by the ratio of the continuum reported in the IRAS point source catalogue to the continuum measured on the Spitzer spectra extracted from the CASSIS database (Lebouteiller et al. 2015). We assumed here that the line emission scales (at first order) with the IR brightness distribution. In particular, we considered the continuum at 12\u2006\u03bcm for the [NeII]12.8\u2006\u03bcm and [NeIII]15.6\u2006\u03bcm lines, and the continuum at 25 \u03bcm for the [OIV]25.9\u2006\u03bcm, [FeII]26\u2006\u03bcm, [SIII]33.5\u2006\u03bcm, and [SiII]34.8\u2006\u03bcm lines. This correction was not needed for the AGN sample and the GOALS sample because of the greater average redshift of the galaxies in the 12MGS and GOALS samples. In particular, the 12MGS active galaxy sample has a mean redshift of 0.028 (Rush et al. 1993), while the GOALS sample has a mean redshift of 0.026. The galaxies presented by Bernard-Salas et al. (2009) have instead an average redshift of 0.0074, while the sample by Goulding & Alexander (2009) has an average redshift of 0.0044. For the other lines in the 10\u2005\u2212\u200536\u2006\u03bcm interval, Goulding & Alexander (2009) did not report a detection, and we used the data presented in Bernard-Salas et al. (2009) for a total of 15 objects. Both Bernard-Salas et al. (2009) and Goulding & Alexander (2009) reported data from the high-resolution Spitzer-IRS spectra. Data in the 50\u2005\u2212\u2005205\u2006\u03bcm interval were taken from D\u00edaz-Santos et al. (2017). For the GOALS sample, 20 starburst galaxies were taken from Fern\u00e1ndez-Ontiveros et al. (2016), and 23 objects were taken from the ISO-LWS observations of Negishi et al. (2001). As a result, we obtained a total sample of 193 objects. Lastly, the PAH features\u2019 fluxes were measured from the low-resolution Spitzer-IRS spectra by Brandl et al. (2006), including 12 objects from the sample of Bernard-Salas et al. (2009) and 179 objects from Stierwalt et al. (2014).","Citation Text":["Goulding & Alexander (2009)"],"Functions Text":["For the Bernard-Salas et al. (2009),","and the GOALS samples, we excluded all the composite starburst-AGN objects identified as those with a detection of [NeV] either at 14.3 or 24.3 \u03bcm."],"Functions Label":["Uses","Uses"],"Citation Start End":[[1132,1159]],"Functions Start End":[[1095,1131],[1161,1308]]}
{"Identifier":"2018AandA...613A..50C__Reiners_et_al._2013_Instance_1","Paragraph":"As previously mentioned, the small amplitude of the RV modulation in the NIR as measured by GIANO contemporaneously with REM photometry is intriguing. Assuming that the depth of the spectral lines relative to their adjacent continuum is constant and considering a spot with a contrast Cs at latitude \u03d5 with a filling factor fs, we expect an RV modulation approximately of Csfsvsinicos\u03d5 (cf. Saar et al. 1997; Desort et al. 2007) that is \u2248 0.75\u22121.0 km s\u22121 for a spot at \u03d5 = 60\u00b0, which is remarkably higher than what has been observed. A cool spot at a higher latitude would reduce the amplitude of the NIR wide-band flux modulations; a quenching of the convective shifts or the Zeeman effect also do not appear to be viable explanations because they increase the effect of a cool spot on the RV at NIR wavelengths (Reiners et al. 2013). Nevertheless, the variation of the relative depths of the spectral lines in the NIR cannot be neglected and it is the dominant effect in the cool spot responsible for the large photometric modulation in the infrared. The relative line depths are strong functions of the continuum opacity and of the degree of element ionization both remarkably varying in a cool spot area with respect to the unperturbed photosphere. In general, these effects produce a remarkable increase of the relative depths of the spectral lines in the cool spot. This compensates for the decrease of the continuum intensity in the spot, reducing the distortions of the spectral line profiles and yielding an RV variation in the NIR smaller than expected from the wide-band photometric variation in the case of constant relative line depths. In any case, a quantitative analysis is not warranted by our data, since a larger number of observations would be required. This scenario suggests a need to better investigate this kind of target, since they might go through specific activity phases during which the VIS and NIR RV amplitudes are similar, possibly resulting in false positives. Looking at the curve phase shifts might give crucial information in these cases (see e.g. the recent result by Hatzes et al. 2018 for the K-giant \u03b3 Draconis).","Citation Text":["Reiners et al. 2013"],"Functions Text":["A cool spot at a higher latitude would reduce the amplitude of the NIR wide-band flux modulations; a quenching of the convective shifts or the Zeeman effect also do not appear to be viable explanations because they increase the effect of a cool spot on the RV at NIR wavelengths"],"Functions Label":["Uses"],"Citation Start End":[[814,833]],"Functions Start End":[[534,812]]}
{"Identifier":"2021MNRAS.506.5015H__Page_et_al._2004_Instance_1","Paragraph":"Due to their high core temperatures at birth, neutron stars cool pre-dominantly by neutrino emission at ages \u2272 106 yr (Potekhin, Pons & Page 2015). The most rapid changes of surface temperature occur early, first when the temperature of the outer layers achieves equilibrium with the rapidly cooling core at an age of \u2272 100 yr (Lattimer et al. 1994; Gnedin, Yakovlev & Potekhin 2001; see also Nomoto & Tsuruta 1987) and then when the temperature drops below the critical temperature for core neutrons to become superfluid, which activates the efficient neutrino emission process of Cooper pair formation and breaking (Gusakov et al. 2004; Page et al. 2004). The rapid cooling of the CCO in Cassiopeia A (at a ten-year rate of \u22482.2 \u00b1 0.2 or 2.8 \u00b1 0.3 per\u2009cent, depending on whether NH varies between observations) indicates the latter starts to take place at an age of \u223c200 yr (Page et al. 2011; Shternin et al. 2011). Neutron star cooling models predict that by an age of several hundred years, the cooling rate will be 1 per\u2009cent per decade. From the 1\u03c3 temperature uncertainties of our fit results with model parameters linked between observations (see Tables 4 and 7), we estimate upper limits on the ten-year cooling rates of 6 per\u2009cent for XMMU J172054.5\u2212372652 and 17 per\u2009cent for CXOU J160103.1\u2212513353. We measure a possible increase in temperature of \u223c4 \u00b1 2 per\u2009cent (accompanied by a decrease in emission area; see Table 5 and 6) for 1WGA J1713.4\u22123949. We also perform fits of the spectra of XMMU J172054.5\u2212372652 and CXOU J160103.1\u2212513353 which allow the temperature to be different between each observation, and the results are shown in Table 9 and Fig. 8 (analogous results for 1WGA J1713.4\u22123949 are shown in Fig. 3). We point out that we are concerned in Fig. 8 with relative changes in temperature and not in absolute temperature differences between CCOs since absolute temperatures depend on a variety of factors that are intrinsic to each CCO and may be different among CCOs in our sample, e.g. mass (and hence neutrino cooling rate) and radius (and hence gravitational redshift) and envelope composition and thickness. Nevertheless it is noteworthy that the temperatures of all three older CCOs appear to be higher than those of Cassiopeia A. Unlike for the 340 yr old CCO in Cassiopeia A, we do not see that temperatures of the 600 yr old XMMU J172054.5\u2212372652 and 1000 yr old CXOU J160103.1\u2212513353 are changing, at least within measurement uncertainties.","Citation Text":["Page et al. 2004"],"Functions Text":["The most rapid changes of surface temperature occur early,","and then when the temperature drops below the critical temperature for core neutrons to become superfluid, which activates the efficient neutrino emission process of Cooper pair formation and breaking"],"Functions Label":["Background","Background"],"Citation Start End":[[639,655]],"Functions Start End":[[148,206],[416,616]]}
{"Identifier":"2020AandA...637A..59A__Ziurys_et_al._(2018)_Instance_1","Paragraph":"Several molecules show a large discrepancy between the abundances derived from observations and calculated by chemical equilibrium, although it is not as severe as for the molecules discussed above. We refer to PN in O-rich stars and H2S in C-rich stars, which are indicated by hatched rectangles in Fig. 2. For PN in O-rich AGB atmospheres, the disagreement between the observed abundances, (1\u20132) \u00d7 10\u22128 (Ziurys et al. 2018), and the calculated maximum chemical equilibrium abundance is almost three orders of magnitude. However, uncertainties on the observational and theoretical sides mean that the true level of disagreement is unclear. For example, while Ziurys et al. (2018) derived a PN abundance of 10\u22128 relative to H2 in IK Tau, De Beck et al. (2013) and Velilla Prieto et al. (2017) derived higher abundances, (3\u20137) \u00d7 10\u22127, in this source. When we give preference to these latter abundances, the level of disagreement would be even higher. On the other hand, the formation enthalpy of PN is rather uncertain (see Lodders 1999), which directly translates into the calculated chemical equilibrium abundance. In this study we adopted the thermochemical data for PN from Lodders (1999), who gives preference to a formation enthalpy at 298.15 K of 171.5 kJ mol\u22121, while other compilations such as JANAF use lower values that would result in higher chemical equilibrium abundances for PN. This would reduce the level of disagreement. In the case of H2S in C-rich AGB stars, the calculated maximum chemical equilibrium abundance is 7 \u00d7 10\u221211, while the value derived from observations is ~50 times higher. In this case, the observed abundance is based on the detection of only one line in only one source (see Ag\u00fandez et al. 2012), and thus it has to be viewed with some caution. In summary, the main failures of chemical equilibrium to account for the observed abundances of parent molecules in circumstellar envelopes are NH3, HCN, CS, SO2, and possibly PN in M-type stars, H2O and NH3 in S-type stars, and the hydrides H2O, NH3, SiH4, PH3, and perhaps H2S as well in C-type stars. The large discrepancies between the abundances derived from observations and those calculated with chemical equilibrium necessarily imply that nonequilibrium chemical processes must be at work in AGB atmospheres. Any invoked nonequilibrium scenario must account for all these anomalously overabundant molecules, but must also reproduce the remaining molecular abundances that are reasonably well explained by chemical equilibrium. No scenario currently provides a fully satisfactory agreement with observations, although two mechanisms that can drive the chemical composition out of equilibrium have been proposed.","Citation Text":["Ziurys et al. 2018"],"Functions Text":["For PN in O-rich AGB atmospheres, the disagreement between the observed abundances, (1\u20132) \u00d7 10\u22128","and the calculated maximum chemical equilibrium abundance is almost three orders of magnitude."],"Functions Label":["Differences","Differences"],"Citation Start End":[[406,424]],"Functions Start End":[[308,404],[427,521]]}
{"Identifier":"2018ApJ...864...90M__Shields_1992_Instance_1","Paragraph":"The LLAGN interpretation of LINERs was initially motivated by their significant X-ray emission (Ferland & Netzer 1983; Halpern & Steiner 1983). Although LLAGNs are found via radio and X-ray observations in the majority of LINERs (Dudik et al. 2005, 2009; Nagar et al. 2005; Filho et al. 2006; Flohic et al. 2006; Gonz\u00e1lez-Mart\u00edn et al. 2009), they are not powerful enough to photoionize the gas in their vicinity on \u223c100 pc (i.e., a few arcseconds in nearby galaxies) scales on which the characteristic emission lines are detected (Flohic et al. 2006; Eracleous et al. 2010a, and references therein). Imaging studies of LINERs have found extended line-emitting regions and complex circumnuclear dust morphologies that might obscure and further prevent the LLAGNs from fully ionizing the surrounding gas (Barth et al. 1999; Pogge et al. 2000; Sim\u00f5es Lopes et al. 2007; Gonz\u00e1lez Delgado et al. 2008; Gonz\u00e1lez-Mart\u00edn et al. 2009; Masegosa et al. 2011). Wolf-Rayet stars (i.e.,\u201cwarmers\u201d; Terlevich & Melnick 1985) could successfully mimic the X-ray emission produced by a LLAGN, as well as provide the hard ionizing photons necessary to explain the relative intensities of the observed optical emission lines. If Wolf-Rayet stars were the primary source of ionizing photons for LINER-like emission lines, then most LINERs would be in the immediate post-starburst phase, which is unlikely given their high occurrence rate (Ho et al. 1997b) and stellar populations (Cid Fernandes et al. 2004; Gonz\u00e1lez Delgado et al. 2004). Compact starbursts containing hot O stars offer an alternative explanation for the relative intensities of LINER emission lines (Filippenko & Terlevich 1992; Shields 1992), but have difficulty explaining the broad Balmer emission wings often seen in LINER spectra (Filippenko 1996), since this would require long-lived supernova remnants. Moreover, the census of stellar populations in LINERs by Gonz\u00e1lez Delgado et al. (2004) and Cid Fernandes et al. (2004) does not show a high incidence of compact, nuclear starbursts. pAGB stars and planetary nebulae (Binette et al. 1994; Taniguchi et al. 2000; Yan & Blanton 2012; Belfiore et al. 2016) are more plausible stellar-based models, but are applicable only to a subset of LINERs due to their inability to explain large H\u03b1 equivalent widths (Ho et al. 2003). Shock models can adequately describe the off-nuclear optical and ultraviolet (UV) emission-line spectra of some LINERs, such as M87, obtained with the Hubble Space Telescope (HST; see Dopita et al. 1996, 1997; Sabra et al. 2003). Meanwhile the UV spectra of some LINER nuclei, such as NGC 4579, do not have emission line ratios that are well described by shock excitation, and yet show high ionization lines that would require a hard extreme UV source, such as the continuum from an active galactic nucleus (AGN), or fast shocks (Barth et al. 1996; Dopita et al. 2015). Successful shock models could have a wide range of shock velocities (Filippenko 1996), but the gas must be continuously shocked to maintain LINER-like ratios. The shocks could potentially be driven by radio jets, which are fairly common in LINERs and whose kinetic power is considerably higher than the electromagnetic luminosity of the LLAGNs (Filho et al. 2002; Nagar et al. 2005; Maoz 2007). Alternatively, the shocks could result from supernovae or winds from either young or evolved stars.","Citation Text":["Shields 1992"],"Functions Text":["Compact starbursts containing hot O stars offer an alternative explanation for the relative intensities of LINER emission lines",", but have difficulty explaining the broad Balmer emission wings often seen in LINER spectra",", since this would require long-lived supernova remnants."],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1676,1688]],"Functions Start End":[[1518,1645],[1689,1781],[1799,1856]]}
{"Identifier":"2016ApJ...830...15J__Rice_et_al._2011_Instance_1","Paragraph":"T Tauri stars in general are known to flare (e.g., Gahm 1990; Guenther & Ball 1999). A potentially better analog of the type of variable H\u03b1 emission expected from chromospheric emission and flaring on PTFO 8-8695 is the WTTS V410 Tau, with v sin i = 77.7 km s\u22121 (e.g., Carroll et al. 2012) compared to the measured v sin i = 80.6 \u00b1 8.1 km s\u22121 for PTFO 8-8695 (van Eyken et al. 2012). The H\u03b1 emission equivalent width (3 \u212b with a typical value \u223c1\u20132 \u212b ) on V410 Tau (e.g., Hatzes 1995; Fern\u00e1ndez et al. 2004; Mekkaden et al. 2005) is weaker than seen on PTFO 8-8695, although V410 Tau has an earlier spectral type that raises the continuum level without necessarily affecting the strength of the chromospheric emission. V410 Tau has been observed to flare in a number of studies. Outside of flares, the H\u03b1 line of V410 Tau is fairly symmetric, relatively narrow, and is similar in shape to the chromospheric H\u03b1 profiles for PTFO 8-8695 seen in Figures 3, 5, and 6 (Hatzes 1995; Fern\u00e1ndez et al. 2004; Skelly et al. 2010). The H\u03b1 line of V410 Tau can grow much stronger and broader during a flare, and also show asymmetries; however, the observed asymmetries seen during flares do not show excess emission with apparent peaks shifted out to greater than \u00b1200 km s\u22121 (Hatzes 1995; Rice et al. 2011) as seen here in PTFO 8-8695. The typical pattern in a flare is for the line to very rapidly (timescale of a few minutes) strengthen and broaden with only a slight asymmetry developing. The strength and width of the line then decay exponentially with a timescale of \u223c1 hr for strong flares (e.g., Fern\u00e1ndez et al. 2004). This is not the temporal behavior observed in PTF08-8695. There is at least one additional piece of evidence against the flaring interpretation for the excess H\u03b1 emission seen in PTFO 8-8695. Whenever V410 Tau shows flare emission in H\u03b1, significant He i 5876 \u212b emission also appears. This He i line is covered in the echelle formats of both our McDonald and Kitt Peak data. We have searched both data sets for evidence of this emission, including co-adding the spectra when the H\u03b1 emission appears stationary (UT 9:44 to 11:17 for McDonald; UT 6:45 to 8:25 for Kitt Peak) to increase the signal-to noise. No evidence of He i emission is seen. Lastly, if the observed excess H\u03b1 emission seen in Figures 2 and 7 were the result of a stellar flare, it would be a remarkable coincidence that the flare-induced asymmetry just happened to appear at and move with the same velocity position in the line profile as that expected for the planetary companion. In particular, the motion shown in Figures 2 and 4 where the excess emission first appears strongly on one side of the line profile and then moves to the other side has not to our knowledge been observed in the H\u03b1 emission of flare stars. Flares have been observed on PTFO 8-8695 (van Eyken et al. 2012; Ciardi et al. 2015), and while flares on this star likely will produce changes in the strength and shape of the H\u03b1 emission line, we believe all the points described above argue strongly against a purely stellar origin.","Citation Text":["Rice et al. 2011"],"Functions Text":["The H\u03b1 line of V410 Tau can grow much stronger and broader during a flare, and also show asymmetries; however, the observed asymmetries seen during flares do not show excess emission with apparent peaks shifted out to greater than \u00b1200 km s\u22121","as seen here in PTFO 8-8695."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1277,1293]],"Functions Start End":[[1020,1262],[1295,1323]]}
{"Identifier":"2017AandA...605A...5S__Dokkum_2001_Instance_1","Paragraph":"The H\u03b1 imaging for the galaxies discussed in this paper was generated from two different sources: NGC 3628 was observed on May 8, 1991 using the ESO New Technology Telescope (NTT, red arm of the ESO Multi-Mode Instrument (EMMI)) for 900 s in H\u03b1 and 300 s in R as part of ESO program 047.01-003. The detector was a Ford 20482 single CCD (ESO CCD #24). Both images were taken with subarcsecond seeing. The data and all relevant calibration files taken during the run and a few days before and after the observing run were retrieved from the ESO archive and re-reduced by us using IRAF in the usual manner. L.A. Cosmic (van Dokkum 2001) was used to clear the image of cosmic rays to produce the final H\u03b1 and final R image. The continuum subtracted image was then produced by aligning H\u03b1 and R band, homogenizing the PSF, and subtracting the appropriately scaled R image from the H\u03b1 image (e.g., Skillman et al. 1997). To determine the scaling factor, we measured the apparent fluxes for several stars in the field on both the H\u03b1 and R images. Scaling, subtraction, and correction to reach an optimal correction of the continuum light of the galaxy was performed with our own IRAF scripts. The NGC 4522 H\u03b1 data were taken directly from the work of Koopmann et al. (2001) and the Galaxy On Line Database Milano Network (GOLDMine, Gavazzi et al. 2003). The images were astrometrically calibrated to the Digitized Sky Survey (DSS) system. Flux calibration of the continuum subtracted H\u03b1 images was performed by transferring the total Sloan Digital Sky Survey (SDSS) r\u2032 band flux of each galaxy to the its continuum image. With the derived scaling factor (see above), the flux scale can then be directly transferred to the continuum corrected H\u03b1 image. A conversion from magnitudes to flux was performed using the fact that SDSS is calibrated in AB magnitudes, so that the zero-point flux density of each filter is 3631 Jy (with 1 Jy = 1 Jansky = 10-26 W Hz-1 m-2 = 10-23 erg s-1 Hz-1 cm-2)1. ","Citation Text":["van Dokkum 2001"],"Functions Text":["L.A. Cosmic","was used to clear the image of cosmic rays to produce the final H\u03b1 and final R image."],"Functions Label":["Uses","Uses"],"Citation Start End":[[617,632]],"Functions Start End":[[604,615],[634,719]]}
{"Identifier":"2021MNRAS.503.6155C__Lovisari_et_al._2017_Instance_1","Paragraph":"Galaxy clusters are the traces of the formation of the largest structures in the Universe and so reliable tools to investigate structures formation and evolution. In principle, this is possible only if and when we have full knowledge of the properties of these objects. The total mass (i.e. the total amount of the dark matter (DM), the intracluster medium (ICM), and the stellar components) is an invaluable quantity when exploring the abundances of clusters along the redshift: a standard way to infer cosmological parameters such as the mean matter density \u03a9m and the amplitude of matter perturbations \u03c38(Planck Collaboration XIII 2016). Furthermore, under the assumption of a simple self-similar model (Kaiser 1986; Voit 2005), we could derive the total mass of the clusters from a few observables in optical, X-ray, or millimetre band (Giodini et al. 2013). This approach results in a few scaling relations valuable when we are interested to obtain averaged results based on some statistics. However, it is prone to the assumed simplified approximations: hydrostatic equilibrium and isothermal and spherical distribution for DM and ICM (Bryan & Norman 1998). It is well known that the hydrostatic equilibrium in haloes is not always satisfied, due to non-thermal pressure contributions from internal motions and turbulence (see e.g. Fang, Humphrey & Buote 2009; Lau, Kravtsov & Nagai 2009; Lagan\u00e1, de Souza & Keller 2010; Rasia et al. 2012; Nelson, Lau & Nagai 2014; Yu, Nelson & Nagai 2015; Biffi et al. 2016; Eckert et al. 2019; Angelinelli et al. 2020; Ansarifard et al. 2020; Gianfagna et al. 2020; Green et al. 2020), pointing out the impact that the dynamical state of those large gravitational bounded objects should have. Several attempts have been made to infer clusters dynamical state, using both observational data and simulations, by analysing the images of the emission in optical (see e.g. Ribeiro, Lopes & Rembold 2013; Wen & Han 2013) and in the X-ray band (see e.g. Rasia, Meneghetti & Ettori 2013; Lovisari et al. 2017; Nurgaliev et al. 2017; Bartalucci et al. 2019; Cao, Barnes & Vogelsberger 2020; Yuan & Han 2020) or of the diffusion of the cosmic microwave background (CMB) photons by thermal Sunyaev\u2013Zel\u2019dovich (tSZ) effect in the millimetre band (Cialone et al. 2018; De Luca et al. 2020, hereafter DL20), or a combination of some of them (see e.g. Mann & Ebeling 2012; Molnar, Ueda & Umetsu 2020; Ricci et al. 2020; The CHEX-MATE Collaboration 2020; Zenteno et al. 2020). Among the possibilities, we have to mention the studies of the clusters morphology in X-ray and tSZ maps. Several indicators are commonly used, such as asymmetry parameter (Schade et al. 1995), light concentration (Santos et al. 2008), third-order power ratio (Buote & Tsai 1995; Wei\u00dfmann et al. 2013), centroid shift (Mohr, Fabricant & Geller 1993; O\u2019Hara et al. 2006), strip parameter, Gaussian fit parameter (Cialone et al. 2018), and so on. They exploit the maps with different apertures and efficiencies and are applied individually or combined together, even with different weights (see e.g. B\u00f6hringer et al. 2010; Nurgaliev et al. 2013; Rasia et al. 2013; Wei\u00dfmann et al. 2013; Mantz et al. 2015; Cui et al. 2016; Lovisari et al. 2017; Cialone et al. 2018; Cao et al. 2020; DL20; Yuan & Han 2020). A complementary approach is by applying thresholds on specific thermodynamic variables. Among the others, the central electron gas density and the core entropy are fairly reliable (Hudson et al. 2010). The azimuthal scatter in radial profiles of gas density, temperature, entropy, or surface brightness (Vazza et al. 2011) is also used as a proxy of the ICM inhomogeneities and correlated to the clusters dynamical state (see e.g. Roncarelli et al. 2013; Ansarifard et al. 2020). Alternatively, the projected sky separations between key positions in the images are resulting in reliable estimators of the dynamical state. Interestingly, the offsets between the bright central galaxy (BCG) and the peaks and\/or the centroids of X-ray or tSZ maps are an indication of how much the relaxation condition is satisfied, with different efficiency (see e.g. Jones & Forman 1984; Katayama et al. 2003; Lin & Mohr 2004; Sanderson, Edge & Smith 2009; Mann & Ebeling 2012; Rossetti et al. 2016; Lopes et al. 2018; DL20; Ricci et al. 2020; Zenteno et al. 2020). To be mentioned also other approaches based on wavelets analysis (Pierre & Starck 1998), on the Minkowski functionals (Beisbart, Valdarnini & Buchert 2001), or on machine learning (see e.g. Cohn & Battaglia 2019; Green et al. 2019; Gupta & Reichardt 2020).","Citation Text":["Lovisari et al. 2017"],"Functions Text":["Several attempts have been made to infer clusters dynamical state, using both observational data and simulations, by analysing the images of the emission in optical","and in the X-ray band (see e.g."],"Functions Label":["Background","Background"],"Citation Start End":[[2022,2042]],"Functions Start End":[[1735,1899],[1957,1988]]}
{"Identifier":"2018MNRAS.475.4704R__Smith_et_al._2010_Instance_1","Paragraph":"Understanding the influence of environment is contingent on being able to identify and quantify galaxy environments. Common environmental measures include the projected number density of galaxies out to the Nth nearest neighbour, the halo mass of a host group or cluster, or the projected separation from the centre of a group or cluster. Star formation and morphology of galaxies correlate well with these environment proxies, with galaxies in high densities regions (or alternatively, high halo mass or small group\/cluster-centric radius) being preferentially red, passive, and early type (Dressler 1980; Goto et al. 2003; Poggianti et al. 2008; Kimm et al. 2009; Li, Yee & Ellingson 2009; Wetzel, Tinker & Conroy 2012; Wilman & Erwin 2012; Fasano et al. 2015; Haines et al. 2015). An alternative way to parametrize the environment of a host group or cluster is to classify the degree to which a system is dynamically relaxed. A relaxed, dynamically old group or cluster should be characterized by a central galaxy which is the brightest (most massive) member by a significant margin (e.g. Khosroshahi, Ponman & Jones 2007; Dariush et al. 2010; Smith et al. 2010) and is located near the minimum of the potential well (e.g. George et al. 2012; Zitrin et al. 2012, however also see Skibba et al. 2011), satellite galaxies which are distributed in velocity space according to a Gaussian profile (e.g. Yahil & Vidal 1977; Bird & Beers 1993; Hou et al. 2009; Mart\u00ednez & Zandivarez 2012), and diffuse X-ray emission which is symmetric about the group\/cluster centre (e.g. Rasia, Meneghetti & Ettori 2013; Wei\u00dfmann et al. 2013; Parekh et al. 2015). The dynamical state of clusters is related to the age of the halo and the time since infall for member galaxies, which simulations have shown is an important quantity in determining the degree to which galaxy properties are affected by environment (e.g. Wetzel et al. 2013; Oman & Hudson 2016; Joshi, Wadsley & Parker 2017). Unrelaxed groups and clusters are systems which formed more recently or which have recently experienced a significant merger event, and in either case it would be expected that the time-since-infall on to the current halo for member galaxies will be relatively short. Therefore galaxies in unrelaxed groups may have properties which have been less influenced by environment compared to galaxies in more relaxed systems.","Citation Text":["Smith et al. 2010"],"Functions Text":["A relaxed, dynamically old group or cluster should be characterized by a central galaxy which is the brightest (most massive) member by a significant margin (e.g."],"Functions Label":["Background"],"Citation Start End":[[1147,1164]],"Functions Start End":[[929,1091]]}
{"Identifier":"2021AandA...655A..12T__Tang_et_al._2017b_Instance_2","Paragraph":"Using the RADEX3 non local thermodynamic equilibrium (LTE) modeling program (van der Tak et al. 2007) with collisional rate coefficients from Wiesenfeld & Faure (2013), we modeled the relation between the gas kinetic temperature and the measured average of para-H2CO 0.5 \u00d7 [(322\u2013221 + 321\u2013220)\/303\u2013202] ratios, adopting a 2.73 K background temperature, an average observational linewidth of 4.0 km s\u22121, and column densities N(para-H2CO) = 2.7 \u00d7 1012 and 3.7 \u00d7 1012 cm\u22122 for N113 and N159W, respectively. The results are shown in Fig. 5. The values of the para-H2CO column density were obtained with APEX data (beam size ~30\u2033; Tang et al. 2017b), which cover similar regions. Different column densities of para-H2CO only weakly affect derived kinetic temperatures (see Fig. 3 in Tang et al. 2017b or Fig. 4 in Tang et al. 2018a; this was also shown in Fig. 13 and discussed in Sect. 4.3.1 of Mangum & Wootten 1993) as long as all lines are optically thin. Considering that the relation between the gas temperature and the para-H2CO line ratio may vary at different spatial densities (see Fig. 2 in Tang et al. 2017b), we modeled it at spatial densities 104, 105, and 106 cm\u22123 in Fig. 5. It appears that Tkin at n(H2) = 105 cm\u22123 is consistently lower than values at 104 and 106 cm\u22123 by \u227223% and \u227234%, respectively, for Tkin \u2272 100 K. Local thermodynamic equilibrium (LTE) is a good approximation for the H2CO level populations under optically thin and high-density conditions (Mangum & Wootten 1993; Tang et al. 2017a,b, 2018b). Following the method applied by Tang et al. (2017b) in their Eq. (2), we plot the relation between the LTE kinetic temperature, TLTE, and the para-H2CO (3\u20132) line ratio in Fig. 5. Apparently, TLTE agrees well with Tnon-LTE at volume densities n(H2) ~ 105 cm\u22123 as long as Tkin \u2272 100 K. Previous observations show that para-H2CO (3\u20132) is sensitive to gas temperature at density 105 cm\u22123 (Ginsburg et al. 2016; Immer et al. 2016; Tang et al. 2017b). The spatial density measured with para-H2CO (303\u2013202) and C18O (2\u20131) in N113 and N159W is n(H2) ~ 105 cm\u22123 on a size of ~30\u2033 (Tang et al. 2017b). Therefore, here we adopt 105 cm\u22123 as an averaged spatial gas density in the N113 and N159W regions.","Citation Text":["Tang et al. 2017b"],"Functions Text":["Different column densities of para-H2CO only weakly affect derived kinetic temperatures (see Fig. 3 in","as long as all lines are optically thin."],"Functions Label":["Uses","Uses"],"Citation Start End":[[778,795]],"Functions Start End":[[675,777],[914,954]]}
{"Identifier":"2019AandA...626A.130T__Imanishi_et_al._(2007)_Instance_1","Paragraph":"We detect a deep (\u03c4sil\u2006\u2004\u223c\u2004\u20062.3) absorption feature due to silicate grains at around 10\u2006\u03bcm. We compared the optical depth of the silicate feature in LEDA 1712304 with those in other AGNs in a wide range of IR luminosities (1010\u2006L\u2299\u2004 \u2004LIR\u2004 \u20041013\u2006L\u2299). The spectra of the AGNs to be compared are taken from those of the IR galaxies observed by Spitzer\/IRS (Stierwalt et al. 2013; Imanishi et al. 2007; Imanishi 2009; Roussel et al. 2006) with the threshold that the equivalent width of the PAH 6.2\u2006\u03bcm feature is smaller than 0.27\u2006\u03bcm (Stierwalt et al. 2013). In addition, we also take AGNs with low IR luminosities (LIR\u2004 \u20041011\u2006L\u2299) from Wu et al. (2010). Figure 3 shows the relation between the IR luminosity and the optical depth of the silicate feature for LEDA 1712304 and the AGNs. We estimated the optical depths of the AGNs of Stierwalt et al. (2013), Wu et al. (2010), and Roussel et al. (2006) by ourselves in the same manner as was performed for LEDA 1712304 with the method defined by Imanishi et al. (2007), the spectra of which were obtained from the NASA\/IPAC IR Science Archive (IRSA). On the other hand, we adopted the values given in each reference paper for the other AGNs. The blue dotted lines in Fig. 3 show the differences between our estimates and those by Stierwalt et al. (2013), which confirms that the differences between the different methods are not large enough to change a global relation. Figure 3 shows that galaxies with low IR luminosities (LIR\u2004 \u20041011\u2006L\u2299) show significantly shallow silicate absorption features, as already pointed out by Stierwalt et al. (2013). Imanishi (2009) suggested that the number of heavily obscured AGNs that have deep silicate absorption features (\u03c4sil\u2004> \u20042) increases with the IR luminosities of the host galaxies. Therefore, LEDA 1712304 may be a rare galaxy from the aspect of having both deep absorption feature \u03c4sil\u2004\u223c\u20042.3 and low IR luminosity (5\u2005\u00b1\u20051)\u00d71010\u2006L\u2299. Such galaxies have hardly been observed; an exception is NGC 1377 (Roussel et al. 2006), as can be seen in Fig. 3. NGC 1377 is a lenticular galaxy (de Vaucouleurs et al. 1991), the stellar mass of which is 109.3\u2005\u00b1\u20050.1\u2006M\u2299 (Skibba et al. 2011). The IR spectrum of NGC 1377 shows a featureless continuum except the silicate feature due to circumnuclear dust (Imanishi 2006; Roussel et al. 2006).","Citation Text":["Imanishi et al. 2007"],"Functions Text":["The spectra of the AGNs to be compared are taken from those of the IR galaxies observed by Spitzer\/IRS"],"Functions Label":["Uses"],"Citation Start End":[[375,395]],"Functions Start End":[[248,350]]}
{"Identifier":"2020AandA...641A..29G__Kotera_et_al._(2015)_Instance_1","Paragraph":"Our case studies focus on several regions of the parameter space of explosive transients that may be related to specific source categories. From Sect. 3, different types of transient emissions from highly magnetized pulsars (also magnetars) can be affected by secondary acceleration. As mentioned previously, magnetars have been identified in many studies as promising candidates for the acceleration of cosmic rays and the production of secondary high-energy neutrinos, for instance Blasi et al. (2000), Fang et al. (2012, 2013), Lemoine et al. (2015), Kotera et al. (2015). The case of newborn magnetars with millisecond periods illustrates a non-relativistic source class, whereas the case of magnetar giant flares involves relativistic outflows. Moreover, tidal disruption events, low-luminosity gamma-ray bursts and blazar flares are examples of relativistic outflows, whose properties partially overlap in the parameter space of explosive transients. In these overlapping regions, they can be similarly affected by secondary acceleration. Therefore, we chose to describe the case of jetted tidal disruptions, while keeping in mind that this case study can be used as a benchmark example for low-luminosity gamma-ray bursts and blazar flares. We note that beyond standard scenarios involving gamma-ray bursts (e.g., Waxman & Bahcall 1997; Murase & Nagataki 2006; Murase et al. 2008; M\u00e9sz\u00e1ros 2015) and active galactic nuclei (e.g., Bednarek & Protheroe 1999; Atoyan & Dermer 2001; Halzen & Hooper 2005; Dermer et al. 2014; Petropoulou et al. 2016; Murase et al. 2018; Gao et al. 2019), jetted tidal disruptions have also been proposed as candidate sources for the production of high-energy cosmic rays and neutrinos (Wang et al. 2011; Senno et al. 2016; Dai & Fang 2017; Lunardini & Winter 2017; Wang & Liu 2016; Zhang et al. 2017; Biehl et al. 2018; Gu\u00e9pin et al. 2018). These case studies are associated with different types of photon fields that we simply model by hard or soft broken power laws and we can thus assess their impact on the high-energy neutrino spectrum. From Sect. 3, we can see that all these source categories are affected by strong secondary synchrotron losses and should be affected differently by secondary acceleration.","Citation Text":["Kotera et al. (2015)"],"Functions Text":["As mentioned previously, magnetars have been identified in many studies as promising candidates for the acceleration of cosmic rays and the production of secondary high-energy neutrinos, for instance Blasi et al. (2000), Fang et al. (2012, 2013), Lemoine et al. (2015),"],"Functions Label":["Motivation"],"Citation Start End":[[554,574]],"Functions Start End":[[284,553]]}
{"Identifier":"2019AandA...626A.130T__Imanishi_et_al._(2007)_Instance_2","Paragraph":"We detect a deep (\u03c4sil\u2006\u2004\u223c\u2004\u20062.3) absorption feature due to silicate grains at around 10\u2006\u03bcm. We compared the optical depth of the silicate feature in LEDA 1712304 with those in other AGNs in a wide range of IR luminosities (1010\u2006L\u2299\u2004 \u2004LIR\u2004 \u20041013\u2006L\u2299). The spectra of the AGNs to be compared are taken from those of the IR galaxies observed by Spitzer\/IRS (Stierwalt et al. 2013; Imanishi et al. 2007; Imanishi 2009; Roussel et al. 2006) with the threshold that the equivalent width of the PAH 6.2\u2006\u03bcm feature is smaller than 0.27\u2006\u03bcm (Stierwalt et al. 2013). In addition, we also take AGNs with low IR luminosities (LIR\u2004 \u20041011\u2006L\u2299) from Wu et al. (2010). Figure 3 shows the relation between the IR luminosity and the optical depth of the silicate feature for LEDA 1712304 and the AGNs. We estimated the optical depths of the AGNs of Stierwalt et al. (2013), Wu et al. (2010), and Roussel et al. (2006) by ourselves in the same manner as was performed for LEDA 1712304 with the method defined by Imanishi et al. (2007), the spectra of which were obtained from the NASA\/IPAC IR Science Archive (IRSA). On the other hand, we adopted the values given in each reference paper for the other AGNs. The blue dotted lines in Fig. 3 show the differences between our estimates and those by Stierwalt et al. (2013), which confirms that the differences between the different methods are not large enough to change a global relation. Figure 3 shows that galaxies with low IR luminosities (LIR\u2004 \u20041011\u2006L\u2299) show significantly shallow silicate absorption features, as already pointed out by Stierwalt et al. (2013). Imanishi (2009) suggested that the number of heavily obscured AGNs that have deep silicate absorption features (\u03c4sil\u2004> \u20042) increases with the IR luminosities of the host galaxies. Therefore, LEDA 1712304 may be a rare galaxy from the aspect of having both deep absorption feature \u03c4sil\u2004\u223c\u20042.3 and low IR luminosity (5\u2005\u00b1\u20051)\u00d71010\u2006L\u2299. Such galaxies have hardly been observed; an exception is NGC 1377 (Roussel et al. 2006), as can be seen in Fig. 3. NGC 1377 is a lenticular galaxy (de Vaucouleurs et al. 1991), the stellar mass of which is 109.3\u2005\u00b1\u20050.1\u2006M\u2299 (Skibba et al. 2011). The IR spectrum of NGC 1377 shows a featureless continuum except the silicate feature due to circumnuclear dust (Imanishi 2006; Roussel et al. 2006).","Citation Text":["Imanishi et al. (2007)"],"Functions Text":["We estimated the optical depths of the AGNs of Stierwalt et al. (2013), Wu et al. (2010), and Roussel et al. (2006) by ourselves in the same manner as was performed for LEDA 1712304 with the method defined by"],"Functions Label":["Uses"],"Citation Start End":[[988,1010]],"Functions Start End":[[779,987]]}
{"Identifier":"2021MNRAS.507.6012Z__Kendrick,_Hazra_&_Balakrishnan_2015_Instance_1","Paragraph":"The interaction of H2 and HD with atomic hydrogen is among the most widely investigated and important processes in elementary chemical reactions. The H + H2, H + D2, and H + HD reactions serve as benchmarks for experimental and theoretical investigations of bimolecular processes (Marinero, Rettner & Zare 1984; Zhang & Miller 1989; D\u2019Mello, Manolopoulos & Wyatt 1991; Fern\u00e1ndez-Alonso & Zare 2002; Harich et al. 2002; Aoiz, Ba\u00f1ares & Herrero 2005; Yang 2007; Gao et al. 2015; Karandashev et al. 2017; Yuan et al. 2018a,b, 2020; Goswami et al. 2020) and continue to attract much attention in the quest for unraveling quantum effects such as the geometric phase (GP) in chemical reactions (Kendrick, Hazra & Balakrishnan 2015; Hazra, Kendrick & Balakrishnan 2016; Croft et al. 2017; Kendrick 2018; Yuan et al. 2018a,b; Kendrick 2019). These elementary reactions are also of considerable interest in early universe chemistry models in determining H2 and HD column densities and the relative abundance of H\/D in the interstellar medium (Flower 1999, 2000; Flower & Roueff 1999; Neufeld et al. 2006; Wrathmall, Gusdorf & Flower 2007; Gay et al. 2011; Nolte et al. 2011; Desrousseaux et al. 2018; Walker, Porter & Stancil 2018; Neufeld et al. 2019; Zhou et al. 2020). The H + HD \u2194 D + H2 chemical reaction is especially important in this context as it cycles the heavier isotope between molecular form (HD) and purely atomic form. The HD molecule by virtue of its small dipole moment also serves as a tracer of H2 through its j = 1 \u2192 0 rotational transition at 112 \u03bcm (Neufeld et al. 2006; Desrousseaux et al. 2018). Observations (Howat et al. 2002; Yuan et al. 2012) using the Infrared Spectrograph on the Spitzer Space Telescope have reported HD emissions from excited rovibrational level v = 1,  j = 5, as well as the pure rotational R(3) and R(4) lines, but relevant reaction rate coefficients are still limited in terms of initial HD rovibrational levels and gas temperature.","Citation Text":["Kendrick, Hazra & Balakrishnan 2015"],"Functions Text":["and continue to attract much attention in the quest for unraveling quantum effects such as the geometric phase (GP) in chemical reactions"],"Functions Label":["Background"],"Citation Start End":[[689,724]],"Functions Start End":[[550,687]]}
{"Identifier":"2022MNRAS.516..731B__Wang,_Hammer_&_Yang_2022_Instance_1","Paragraph":"Now with the recent availability of high quality full 6D phase-space information for large numbers of sources, much effort has been made to decrease the uncertainties in the Milky Way mass estimate. Recent works using a tracer mass estimator with 6D phase-space information include Sohn et al. (2018, globular clusters), Watkins et al. (2019, globular clusters), and Fritz et al. (2020, satellites). The most recent work using the spherical Jeans equation by Zhai et al. (2018) is very similar to our current investigation in method and data (LAMOST K giants) although only line-of-sight velocities were included, whereas we additionally make use of proper motions from Gaia to obtain the stellar tangential velocities. Using Bayesian analysis to fit a distribution function to full 6D phase-space data (globular clusters, satellites, and halo stars) has been a recent popular choice among many works (Eadie & Juri\u0107 2019; Posti & Helmi 2019; Vasiliev 2019; Deason et al. 2021; Correa Magnus & Vasiliev 2022; Shen et al. 2022; Slizewski et al. 2022; Wang, Hammer & Yang 2022) and a similar distribution function analysis using 5D phase-space data from Gaia (Hattori, Valluri & Vasiliev 2021). In addition to fitting the observational data with a distribution function, several works have incorporated into the fitting a comparison of the observed data with Milky Way-type galaxies from cosmological simulations (Callingham et al. 2019; Li et al. 2020). Newly discovered high velocity stars with full 6D phase-space information have been used to estimate the mass of the Milky Way (Hattori et al. 2018; Monari et al. 2018; Deason et al. 2019; Grand et al. 2019; Koppelman & Helmi 2021; Necib & Lin 2022). Vasiliev, Belokurov & Erkal (2021) and Craig et al. (2021) have estimated the Milky Way mass by fitting models for the Sagittarius and Magellanic Steams, respectively. Several recent studies have estimated the Milky Way mass using measurements of the rotation curve (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). Other works have used 6D satellite phenomenology, characterizing simulated Milky Way-type satellite populations and comparing to the observations of satellites in the Milky Way, to estimate the mass of the Milky Way (Patel et al. 2018; Villanueva-Domingo et al. 2021; Rodriguez Wimberly et al. 2022). Zaritsky et al. (2020) apply the timing argument to distant Milky Way halo stars to derive a lower limit to the Milky Way mass.","Citation Text":["Wang, Hammer & Yang 2022"],"Functions Text":["Using Bayesian analysis to fit a distribution function to full 6D phase-space data (globular clusters, satellites, and halo stars) has been a recent popular choice among many works"],"Functions Label":["Background"],"Citation Start End":[[1049,1073]],"Functions Start End":[[720,900]]}
{"Identifier":"2018MNRAS.475.1646F__Ji_et_al._2003_Instance_1","Paragraph":"Solar filaments, or prominences as they are called when observed above the solar limb, can be observed in a stable state for many days or weeks. Sometimes, they suddenly start to ascend as a whole (full eruptions) (Joshi & Srivastava 2011; Holman & Foord 2015) or within limited sections of their length (partial eruptions) (Gibson & Fan 2006; Kliem et al. 2014). The ascending of a filament can go on high into the corona (successful eruptions) and gives rise to a coronal mass ejection (CME) or can stop at some greater height in the corona (confined or failed eruptions) (Ji et al. 2003; T\u00f6r\u00f6k & Kliem 2005; Alexander, Liu & Gilbert 2006; Kuridze et al. 2013; Kushwaha et al. 2015). Occasionally two-step eruptions are observed. A filament after the first jump decelerates and stops at a greater height as in failed eruptions, but after a rather short period of time it starts to rise again and develops into the successful eruption with a CME formation. Byrne et al. (2014) observed on 2011 March 8 at the solar limb the erupting loop system that stayed in a metastable intermediate position for an hour and then proceeded and formed the core of a CME. Gosain et al. (2016) analysed observations of the eruption of a long quiescent filament on 2011 October 22 observed from three viewpoints by space observatories. A two-ribbon flare and the onset of a CME appeared 15\u2009h after the filament disappearance on the disc. The filament was not observed at the high metastable position but some coronal structures that can be attributed to a corresponding flux rope were recognized. A clear example of the two-step filament eruption on 2015 March 14\u201315 was reported by Wang et al. (2016) and Chandra et al. (2017). In this event, a part of a large filament separated from the main body of the filament at the height of \u224830 Mm and rose upwards to the height of \u224880 Mm, where it stayed for 12\u2009h clearly visible in chromospheric and coronal spectral lines. Finally, it erupted and produced a halo CME.","Citation Text":["Ji et al. 2003"],"Functions Text":["The ascending of a filament can go on high into the corona (successful eruptions) and gives rise to a coronal mass ejection (CME) or can stop at some greater height in the corona (confined or failed eruptions)"],"Functions Label":["Background"],"Citation Start End":[[575,589]],"Functions Start End":[[364,573]]}
{"Identifier":"2018ApJ...864..165W__Golub_et_al._1974_Instance_1","Paragraph":"In configuration 1S, the field close to the separatrix was sheared, producing steady interchange reconnection modulated by quasi-periodic reconnection bursts. We can roughly estimate the free energy release rate of the steady component from the energy injected before the onset of reconnection in the first 200 time units of the simulation. By t = 200, around 2 units of free magnetic energy are injected into the closed field; see Figure 16(a). Accounting for the ramp up of the driver and scaling the values, this corresponds to an energy injection rate of \u22485.6 \u00d7 1023 erg s\u22121 at the maximum driving speed. During the quasi-steady phase, this injection is balanced by losses to numerical diffusion and equates to roughly the free energy available for heating the plasma. Even after accounting for the unrealistically fast driving speed (see below), this energy release rate compares well with the observed values of 1023\u20131024 erg s\u22121 for bright points (Golub et al. 1974; Priest et al. 1994). The energy released by the bursts was a small fraction of the stored free magnetic energy\u2014\u22480.5 \u00d7 Es = 4.9 \u00d7 1025 erg occurring with a period of \u2248240 \u00d7 ts = 8 minutes\u2014while the outflow speeds reached typical values of \u22480.05 \u00d7 Vs \u2248 60 km s\u22121 along the outer spine. The energy released in each burst corresponds to \u224818% of the energy released over the same period by the steady component. Many bright points exhibit quasi-periodic intensity increases, with periods ranging from a few minutes to a couple of hours (Kariyappa & Varghese 2008; Tian et al. 2008; Zhang et al. 2012). Our results demonstrate that some of this periodicity can be explained by the natural modulation of the interchange reconnection that occurs as minority-polarity elements are moved by surface motions. The predicted outflow speeds, and certainly the periods of the reconnection cycles, are likely too fast because the driving speed (12.5 km s\u22121) employed in our simulations is too high. However, configuration 1F demonstrated that the cycle period is mainly set by the displacement of the minority polarity. We speculate that, at more typical photospheric speeds (\u22481.5 km s\u22121; e.g., Brandt et al. 1988), the reconnection cycle period would increase by a factor of 12.5\/1.5 \u00d7 8 minutes \u224867 minutes, corresponding to the longest observed oscillations in brightness. Without a full treatment of the thermodynamics, however, it is not clear whether the repetitive, low-intensity reconnection jets in this case would be observable.","Citation Text":["Golub et al. 1974"],"Functions Text":["Accounting for the ramp up of the driver and scaling the values, this corresponds to an energy injection rate of \u22485.6 \u00d7 1023 erg s\u22121 at the maximum driving speed. During the quasi-steady phase, this injection is balanced by losses to numerical diffusion and equates to roughly the free energy available for heating the plasma. Even after accounting for the unrealistically fast driving speed (see below), this energy release rate compares well with the observed values of 1023\u20131024 erg s\u22121 for bright points"],"Functions Label":["Similarities"],"Citation Start End":[[955,972]],"Functions Start End":[[446,953]]}
{"Identifier":"2020ApJ...899L...6L__Margalit_et_al._2019_Instance_1","Paragraph":"The leading FRB source model invokes magnetars as the power source to produce repeating bursts. There are two versions of this model. One version invokes rapidly spinning young magnetars that are produced in extreme stellar transients such as GRBs and SLSNe. The main motivation is that the host galaxy of FRB 121102 resembles those of LGRBs and SLSNe (Metzger et al. 2017; Nicholl et al. 2017; Wadiasingh & Timokhin 2019). The fact that the hosts of all other FRBs do not resemble that of FRB 121102 disfavors the simplest version of this proposal. A possible fix of this proposal is to introduce rapidly spinning magnetars born from binary neutron star (BNS) mergers (Margalit et al. 2019; Wang et al. 2020). In order to make this scenario work, one needs to require that rapidly spinning magnetars made from BNS mergers should be much more abundant than those made from LGRBs and SLSNe. Comparing the event rate densities of BNS mergers, LGRBs, and SLSNe (e.g., Sun et al. 2015; Abbott et al. 2017; Nicholl et al. 2017), this may be possible if a significant fraction of BNS mergers leave behind stable neutron stars (e.g., Gao et al. 2016). However, if this fraction is very low, as required if GW170817 leaves behind a black hole (Margalit et al. 2019), the fast magnetar model may fail to explain the small fraction of LGRB\/SLSN-like hosts in FRB samples. The second version of the magnetar model invokes emission (e.g., giant flares) from slowly rotating magnetars like the ones observed in the Galaxy (e.g., Popov & Postnov 2010; Katz 2014; Kulkarni et al. 2014). The births of these magnetars do not require extreme explosions such as GRBs and SLSNe (e.g., Beniamini et al. 2019). If this is the case, the host galaxy distribution may be more analogous to that of SNe II. All FRBs but FRB 121102 are consistent with this scenario (Figure 4). In order to interpret FRB 121102, the more extreme channel of forming rapid magnetars is still needed. So we conclude that the magnetar model would work only if both fast magnetars produced in extreme explosions and slow magnetars produced in regular channels (Beniamini et al. 2019) can produce FRBs. In any case, since the birth rate of these magnetars is very high (Beniamini et al. 2019), an additional factor is needed to select a small fraction of magnetars to produce FRBs (e.g., Ioka & Zhang 2020).","Citation Text":["Margalit et al. 2019"],"Functions Text":["A possible fix of this proposal is to introduce rapidly spinning magnetars born from binary neutron star (BNS) mergers"],"Functions Label":["Background"],"Citation Start End":[[670,690]],"Functions Start End":[[550,668]]}
{"Identifier":"2021MNRAS.503.3629D__Kelley,_Blecha_&_Hernquist_2017a_Instance_1","Paragraph":"Having found evidence that galaxies show morphological signatures of a recent galaxy merger over time-scales on the order of 300\u2013500\u2009Myr, we now consider the typical BH merger time-scales and the impact that may have on our results. Within the Illustris simulation (and indeed many similar simulations), a pair of BHs merge as soon as their separation is less than the particle\u2019s smoothing length, rather than incorporating a coalescence time for the binary. Recently, several works have attempted to estimate the expected coalescence time for binary BHs by post-processing cosmological simulations, finding time-scales on the order of 100s of Myr to Gyr for the binary coalescence time (e.g. Blecha et al. 2016; Rantala et al. 2017; Kelley, Blecha & Hernquist 2017a; Mannerkoski et al. 2019; Sayeb et al. 2021). Additionally, the time for a satellite BH to reach the centre of a galaxy and form a binary with the central BH can also be on the order of 100s of Myr to Gyr, based on the dynamical friction time-scale for infall to the galactic centre (e.g. Volonteri et al. 2020). The galaxy structure, e.g. the existence or lack of a dense stellar core, can affect the time satellite BHs spend at large radii (Tremmel et al. 2018a,b; Barausse et al. 2020), and directly incorporating dynamical friction into cosmological volumes suggests that the orbital decay time-scale may be both substantial and redshift dependent (Bartlett et al. 2021). The Illustris simulation uses a re-centring scheme whereby BHs are re-positioned towards the local potential minimum; this prevents numerical wandering of BHs, but also means the the full infall time to reach the galaxy centre may be notably underestimated. Furthermore, both simulations (e.g. Bellovary et al. 2019) and observations (e.g. Reines et al. 2020) suggest that BHs in dwarf galaxies may frequently be located offset from the galaxy centre, which could further delay any mergers involving low-mass BHs seeded into Illustris (which are initially placed at the galaxy centre). Overall, this suggests that the Illustris simulation likely overestimates the speed with which BHs merge following the merging of their host galaxies, and properly accounting for this has the potential to impact the expected GW detection rate, shift the peak detection time to lower redshift, and prevent GW hosts from being visibly disturbed. Although a complete investigation into accurately estimating the time delays remains beyond the scope of this paper, we address this, by imposing a delay between when the BH particles merge in the simulation (which is closer to when the BH binary may form) and when the final coalescence and GW emission occur.","Citation Text":["Kelley, Blecha & Hernquist 2017a"],"Functions Text":["Recently, several works have attempted to estimate the expected coalescence time for binary BHs by post-processing cosmological simulations, finding time-scales on the order of 100s of Myr to Gyr for the binary coalescence time (e.g."],"Functions Label":["Background"],"Citation Start End":[[734,766]],"Functions Start End":[[459,692]]}
{"Identifier":"2016ApJ...821..107G__Gloeckler_&_Fisk_2015_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":["Gloeckler & Fisk 2015"],"Functions Text":["For the density of neutral hydrogen in the inner heliosheath a constant nH = 0.1 cm\u22123 is assumed"],"Functions Label":["Uses"],"Citation Start End":[[1174,1195]],"Functions Start End":[[1053,1149]]}
{"Identifier":"2018ApJ...853...34Z__Giebels_et_al._2007_Instance_1","Paragraph":"Several well-studied TeV blazars show rich spectral behavior in X-rays, which may represent the general behavior of the synchrotron peak of all AGN jets. The X-ray spectra are usually curved (Massaro et al. 2004) and can only locally be fitted by a power law. The spectral variation with flux can be complex (Zhang et al. 2002; Cui 2004). Generally, the spectrum hardens when the flux increases (e.g., Gliozzi et al. 2006; Xue et al. 2006; Tramacere et al. 2009), but photon indices can saturate at higher fluxes (Xue & Cui 2005; Giebels et al. 2007). The synchrotron peak usually moves to higher frequencies with increasing flux during outbursts (e.g., Pian et al. 1998), but no correlation between the break energy and the flux exists when a broken power law is adopted to fit the X-ray spectra (Xue & Cui 2005; Giebels et al. 2007; Garson et al. 2010). A cooling break in the spectrum of emitting particles cannot explain these features (Wierzcholska & Wagner 2016), and some special particle acceleration processes may be involved (Madejski & Sikora 2016). There are also energy-dependent lags between the variations of different energy bands. In some flares, soft bands lag behind hard bands (e.g., Zhang et al. 2002), while lags in the opposite direction can also happen (e.g., Ravasio et al. 2004; Sato et al. 2008). Hysteresis in the HR (hardness ratio)\u2013flux diagram is often used as a diagnostic of lags. Clockwise loops (e.g., Acciari et al. 2009; Kapanadze et al. 2016) in the HR\u2013flux plane are a sign of soft lags while counterclockwise loops (e.g., Tramacere et al. 2009) are a sign of hard lags. The same source can exhibit both clockwise and counterclockwise loops; the observed patterns are further complicated by the superposition of flares at different timescales (Cui 2004). The above knowledge of TeV blazars in the X-ray regime comes from studies focusing on timescales of hours to weeks. We will extend this kind of analysis to much smaller timescales in this paper.","Citation Text":["Giebels et al. 2007"],"Functions Text":["Generally, the spectrum hardens when the flux increases","but photon indices can saturate at higher fluxes"],"Functions Label":["Background","Background"],"Citation Start End":[[530,549]],"Functions Start End":[[339,394],[464,512]]}
{"Identifier":"2022AandA...665A..46M__C\u00f4t\u00e9_et_al._2019_Instance_1","Paragraph":"Several nucleosynthesis processes have been proposed as production sites of these light neutron-capture elements, including r-process in neutron star mergers (Wanajo et al. 2014; Watson et al. 2019), in magneto-rotational supernovae (Winteler et al. 2012), or in collapsars (Siegel et al. 2019), s-process in low- to intermediate-mass stars (Karakas & Lattanzio 2014), weak s-process in rapidly rotating massive stars (Frischknecht et al. 2012; Choplin et al. 2018), and weak r-process in electron-capture supernovae (Wanajo et al. 2011). It is not yet clear which process is the dominant source of the elements in the early Universe (see discussions by, e.g., C\u00f4t\u00e9 et al. 2019; Prantzos et al. 2018; Kobayashi et al. 2020). Since we here discuss the low-metallicity end of the sample, the production of neutron-capture elements would not be dominated by low- to intermediate-mass stars (e.g., de los Reyes et al. 2022). One possible explanation for the low light neutron-capture element abundances of the Helmi streams is that, as a result of the low stellar mass of the galaxy, the progenitor did not experience rare r-process nucleosynthesis events, such as neutron star mergers, electron capture supernovae, and magneto-rotational supernovae. In this case, a small amount of light neutron-capture elements could be produced by rapidly rotating massive stars (Hirai et al. 2019; Tarumi et al. 2021). Another explanation is that the progenitor dwarf galaxy had a small number of rotating massive stars. The small number of rotating massive stars might be a result of the top-light initial mass function in dwarf galaxies (Weidner & Kroupa 2005), or different distribution of initial rotation velocity of stars. The observational indication by Gull et al. (2021) that metal-poor stars of the Helmi streams show r-process abundance pattern in neutron-capture elements heavier than Ba might favor the second possibility. However, it is necessary to investigate the abundance pattern of light neutron-capture elements in order to understand the cause of the low light neutron-capture element abundance of the Helmi streams. A larger sample of low-metallicity Helmi stream stars with neutron-capture element abundances would also be welcomed. They would enable us to constrain the property of the nucleosynthesis processes, such as their event rates, by studying how neutron-capture elements were enriched as a function of metallicity (e.g., Tsujimoto et al. 2017).","Citation Text":["C\u00f4t\u00e9 et al. 2019"],"Functions Text":["It is not yet clear which process is the dominant source of the elements in the early Universe (see discussions by, e.g.,"],"Functions Label":["Motivation"],"Citation Start End":[[661,677]],"Functions Start End":[[539,660]]}
{"Identifier":"2022MNRAS.513.1459M__Conselice,_Yang_&_Bluck_2009_Instance_1","Paragraph":"Hierarchical structure formation scenarios (e.g. Fall & Efstathiou 1980; van den Bosch et al. 2002; Agertz, Teyssier & Moore 2011) predict that massive galaxies acquire much of their stellar mass through a combination of continuous cold gas accretion and mergers with smaller objects (e.g. Press & Schechter 1974; Moster, Naab & White 2013; Kaviraj et al. 2015; Rodriguez-Gomez et al. 2016; Martin et al. 2018b; Davison et al. 2020; Martin et al. 2021). As a consequence, mergers are also expected to play a significant role in driving the evolution of galaxy properties, for example, by triggering (Schweizer 1982; Mihos & Hernquist 1996; Duc et al. 1997; Elbaz & Cesarsky 2003; Kaviraj et al. 2011; Lofthouse et al. 2017; Martin et al. 2017) or quenching (Schawinski et al. 2014; Barro et al. 2017; Kawinwanichakij et al. 2017; Pontzen et al. 2017) star formation in the host galaxy or by driving its morphological evolution (e.g. Toomre 1977; Conselice, Yang & Bluck 2009; Dekel, Sari & Ceverino 2009; Taranu, Dubinski & Yee 2013; Naab et al. 2014; Fiacconi, Feldmann & Mayer 2015; Graham, Dullo & Savorgnan 2015; Deeley et al. 2017; G\u00f3mez et al. 2017; Welker et al. 2017; Martin et al. 2018a; Jackson et al. 2019). Signatures of past mergers take the form of faint extended tidal features such as tails (e.g. Pfleiderer 1963; Toomre & Toomre 1972; Peirani et al. 2010; Kaviraj 2014; Kaviraj, Martin & Silk 2019), or plumes (e.g. Lauer 1988) \u2013 which are typically produced by major mergers \u2013 and streams (e.g. Johnston, Sigurdsson & Hernquist 1999; Shipp et al. 2018; Martinez-Delgado et al. 2021) or shells (e.g. Malin & Carter 1983; Quinn 1984) \u2013 which mainly arise from minor interactions \u2013 as well as in the structure of the surrounding diffuse light (e.g. Choi, Guhathakurta & Johnston 2002; Graham 2002; Johnston, Choi & Guhathakurta 2002; Seigar, Graham & Jerjen 2007; Kaviraj et al. 2012; Monachesi et al. 2016, 2019; Iodice et al. 2019; Montes 2019). These features, which arise from many different types of encounter, hold a fossil record of the host galaxy\u2019s past interactions and mergers which can be used to reconstruct its assembly history and dynamical history (Johnston et al. 2008; Mart\u00ednez-Delgado et al. 2009; Belokurov et al. 2017; Montes et al. 2020; Ren et al. 2020; Spavone et al. 2020; Vera-Casanova et al. 2021). However, the majority of tidal features are expected to have surface brightnesses fainter than 30 mag arcsec\u22122 in the r-band (Johnston et al. 2008). Although pushing towards these kinds of limiting surface brightnesses remains extremely challenging, it is nevertheless desirable to do so, being necessary to uncover a more detailed history of local Universe. This is not only vital for our understanding of hierarchical galaxy assembly (e.g. Johnston, Sackett & Bullock 2001; Wang et al. 2012), but also serves as a novel galactic scale probe of more fundamental physics such as theories of gravity (e.g. Gentile et al. 2007; Renaud, Famaey & Kroupa 2016) and dark matter (Dubinski, Mihos & Hernquist 1996; Kesden & Kamionkowski 2006; Dumas et al. 2015; van Dokkum et al. 2018; Montes et al. 2020). In particular, tidal structure is a powerful tracer of the underlying galactic halo potential (e.g. Dubinski, Mihos & Hernquist 1999; Varghese, Ibata & Lewis 2011; Bovy et al. 2016; Ibata et al. 2020; Malhan, Valluri & Freese 2021).","Citation Text":["Conselice, Yang & Bluck 2009"],"Functions Text":["As a consequence, mergers are also expected to play a significant role in driving the evolution of galaxy properties, for example,","star formation in the host galaxy or by driving its morphological evolution (e.g."],"Functions Label":["Background","Background"],"Citation Start End":[[946,974]],"Functions Start End":[[454,584],[851,932]]}
{"Identifier":"2019MNRAS.486....2M__Yamamoto_et_al._2014_Instance_1","Paragraph":"The discovery of ultraluminous pulsars (ULPs or PULXs; Bachetti et al. 2014; F\u00fcrst et al. 2016; Israel et al. 2017a,b; Carpano et al. 2018) has revolutionised the field of ultraluminous X-ray sources (ULXs; see the review of Kaaret, Feng & Roberts 2017). Whilst ULXs were long considered to be possible candidates for hosting intermediate-mass black holes (IMBHs), it was immediately apparent that the explanation for the extreme observed luminosities (>1039 erg s\u22121) in at least some ULXs was accretion in excess of the classical Eddington limit on to common-place primary objects \u2013 in this case neutron stars. However, whilst the mass regime of the compact object in ULXs has been at least partly resolved (we note that candidate IMBHs still remain, e.g. Farrell et al. 2009), the relative number of black hole to neutron star primaries in ULXs (see King, Lasota & Klu\u017aniak 2017; Middleton & King 2017) and the nature of the accretion flow in ULPs remain outstanding puzzles. At the centre of the debate is the strength of the surface dipole field and any multipolar component. Should the dipole field strength be similar to that of Galactic HMXBs (1012 G \u2013 e.g. Bellm et al. 2014; F\u00fcrst et al. 2014; Tendulkar et al. 2014; Yamamoto et al. 2014) then it is quite plausible that the flow will be supercritical ($\\dot{m}\/\\dot{m}_{\\rm Edd} \\gt $ 1 where $\\dot{m}_{\\rm Edd}$ is the Eddington accretion rate) at radii greater than the magnetospheric truncation radius (rM). In this case, the supercritical portion of the disc will have a large \u2013 close to unity \u2013 vertical scale height and winds will be launched from the surface (see Shakura & Sunyaev 1973; Poutanen et al. 2007 and the simulations of Ohsuga et al. 2009; Jiang, Stone & Davis 2014; Sa\u0327dowski et al. 2014). Within rM, the flow will take the form of an accretion curtain (Mushtukov et al. 2017, 2019) and shock-heated column as material falls on to the magnetic poles. Due to collimation by the disc and outflows beyond rM, it is expected that the intrinsic luminosity is then partially geometrically beamed (see King 2009). Conversely, should the dipole field strength be very high (typically > 1013 G) then it is quite probable that the disc will truncate before becoming locally supercritical. The geometry in this case is then expected to take the form of a geometrically thin disc down to rM, an accretion curtain and shock-heated column. Rather than a supercritical disc and geometrical beaming, super-Eddington luminosities can then be explained by a magnetic pressure supported accretion column (Basko & Sunyaev 1976) and high field strength, the latter allowing for a substantially increased luminosity from a reduction in the electron scattering cross-section (e.g. Herold 1979; Paczynski 1992; Thompson & Duncan 1995; Mushtukov et al. 2015).","Citation Text":["Yamamoto et al. 2014"],"Functions Text":["At the centre of the debate is the strength of the surface dipole field and any multipolar component. Should the dipole field strength be similar to that of Galactic HMXBs (1012 G \u2013 e.g.","then it is quite plausible that the flow will be supercritical ($\\dot{m}\/\\dot{m}_{\\rm Edd} \\gt $ 1 where $\\dot{m}_{\\rm Edd}$ is the Eddington accretion rate) at radii greater than the magnetospheric truncation radius (rM)."],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[1226,1246]],"Functions Start End":[[978,1164],[1248,1470]]}
{"Identifier":"2016AandA...593A..22R__Shibuya_et_al._2015_Instance_2","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"],"Functions Text":["We have now access to the size evolution up to z ~ 10 from the deepest HST imaging data"],"Functions Label":["Background"],"Citation Start End":[[1614,1633]],"Functions Start End":[[1420,1507]]}
{"Identifier":"2021ApJ...922..131T__Kumari_et_al._2019_Instance_1","Paragraph":"DIG contributes with a fraction of 20% to 90% of the total H\u03b1 flux in galaxy disks, with a mean fraction around 50%\u201360% (Hoopes & Walterbos 2003; Oey et al. 2007; Sanders et al. 2017; Tomi\u010di\u0107 et al. 2017; Poetrodjojo et al. 2019; Della Bruna et al. 2020; Tomi\u010di\u0107 et al. 2021). This large contribution may cause star formation rates (SFRs) to be overestimated as H\u03b1 flux from DIG may be wrongly associated with star formation. There is a debate about the extent to which DIG affects measurements of gas-phase metallicity and its radial slope (Searle 1971; Vila-Costas & Edmunds 1992; S\u00e1nchez et al. 2014; Belfiore et al. 2017b; Sanders et al. 2017; S\u00e1nchez-Menguiano et al. 2018; Zhang et al. 2017; Vale Asari et al.2019; Kumari et al. 2019; Poetrodjojo et al. 2019) as some observations indicate lower metallicity (up to 1 dex) in DIG compared to nearby H ii regions. DIG may also exhibit different values of line ratios and ionizing parameter log(q), further affecting observations and analysis of ISM characteristics as well as adding scatter in the distribution of galaxy properties measured from unresolved observations (Martin 1997; Flores-Fajardo et al. 2011; Dopita et al. 2014; Zhang et al. 2017; Poetrodjojo et al. 2018; Mingozzi et al. 2020). Furthermore, the detection of gas that shows different line ratios and ionization parameters located at large distances from H ii regions\u2014larger than the thickness of a typical galactic disk ( \u2248 1 kpc)\u2014would indicate that sources other than star-forming (SF) regions are ionizing such gas (for example HOLMES, shocks, or mixing of different gas layers; Flores-Fajardo et al. 2011; Zhang et al. 2017; Poetrodjojo et al. 2018; Poggianti et al. 2019a). Different galactic characteristics (like mass, SFR, age, etc.) and external physical processes, such as galaxy interactions and gas stripping caused by ram pressure (Gunn et al. 1972; Toomre & Toomre 1972), may affect the ionization parameter and various line ratios (Maier et al. 2006; Nagao et al. 2006; Flores-Fajardo et al. 2011; Zhang et al. 2017; S\u00e1nchez 2020).","Citation Text":["Kumari et al. 2019"],"Functions Text":["There is a debate about the extent to which DIG affects measurements of gas-phase metallicity and its radial slope","as some observations indicate lower metallicity (up to 1 dex) in DIG compared to nearby H ii regions."],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[721,739]],"Functions Start End":[[426,540],[766,867]]}
{"Identifier":"2022AandA...667A..90M__Priestley_&_Whitworth_2022_Instance_1","Paragraph":"The line mass (mass per unit length) of G1.75-0.08 is Mline = 1011 \u00b1 146 M\u2299 pc\u22121. The dynamical state of the filament can be addressed by comparing its line mass with the virial or critical line mass (e.g. Fiege & Pudritz 2000, Eq. (12) therein). To calculate the latter quantity, we used the total (thermal+non-thermal) velocity dispersion, where the observed spectral line width (FWHM) was taken to be the FWHM of the HCN(1 \u2013 0) line detected towards clump B in G1.75-0.08 from Miettinen (2014), that is, 13.50 \u00b1 0.38 km s\u22121. We note that this is a very broad line width (and it is even broader (by a factor of 1.73) in clump A; Sect. 3.5), which could be attributed to multiple factors. From an observational point of view, the hyperfine structure of HCN was not resolved in the MALT90 spectra, which can lead to an overestimation of the line width although it was derived through fitting the hyperfine structure of the transition (Miettinen 2014). Moreover, the angular resolution of the Mopra telescope observations employed by Miettinen (2014) is 38\u2033 (HPBW), which corresponds to 1.5 pc at the cloud distance, and hence the beam might have captured emission from the turbulent outer parts of the cloud. For example, if G1.75-0.08 follows the line width\u2013size relation in the central molecular zone, or CMZ, which for HCN is found to be \u03c3 \u221d R0.62 (Shetty et al. 2012; Table 2 therein), the aforementioned HCN(1 \u2013 0) line width would be expected to be only ~2.5 km s\u22121 on the ~0.1 pc scale, which is a typical inner width (FWHM) of filamentary molecular clouds (e.g. Arzoumanian et al. 2019; Priestley & Whitworth 2022, and references therein). We note that Henshaw et al. (2016a) derived a velocity dispersion of 11 km s\u22121 (~26 km s\u22121 FWHM for a Gaussian profile) for another Galactic centre region IRDC, namely G0.253+0.016 or the Brick from spectral line observations with Mopra, which is comparable to the observational results for G1.75-0.08 with the same telescope (Miettinen 2014). However, based on much higher angular resolution (1\u2033.7) observations with the Atacama Large Millimetre\/submillimetre Array (ALMA), Henshaw et al. (2019) derived an average velocity dispersion of 4.4 km s\u22121 in the Brick, which demonstrates that higher resolution spectral line observations towards G1.75-0.08 are also required. The effect of the angular resolution of the observations on the derived spectral line widths was also demonstrated by Hacar et al. (2018) in the case of the Orion integral filament (see e.g. Fig. 6 therein). From a physical point of view, there are several effects that can lead to line broadening. First, the HCN lines detected towards the G1.75-0.08 clumps were not optically thin, and hence the optical thickness effects might contribute to the broadening of the lines (e.g. Hacar et al. 2016). On the other hand, even the optically thin transitions detected towards the clumps had line widths that are comparable to those of HCN(1 \u2013 0) (Miettinen 2014, Table 3 therein). Second, G1.75-0.08 might be associated with high-velocity gradients along the filament (e.g. Federrath et al. 2016; Gong et al. 2018) that would be blended in the MALT90 spectra and hence lead to overestimated line widths. Third, star formation driven outflows and shocks could lead to broad line profiles. Fourth, G1.75-0.08 is located close to the Galactic centre (RGC \u2243 270 pc), where the interstellar medium is highly turbulent (e.g. Salas et al. 2021 and references therein). At least part of this turbulent gas could contribute to our single-dish-observed spectral line widths. Obviously, further spectral line observations are needed to study the gas kinematics of G1.75-0.08 in more detail.","Citation Text":["Priestley & Whitworth 2022"],"Functions Text":["For example, if G1.75-0.08 follows the line width\u2013size relation in the central molecular zone, or CMZ, which for HCN is found to be \u03c3 \u221d R0.62",", the aforementioned HCN(1 \u2013 0) line width would be expected to be only ~2.5 km s\u22121 on the ~0.1 pc scale, which is a typical inner width (FWHM) of filamentary molecular clouds"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[1595,1621]],"Functions Start End":[[1209,1350],[1388,1563]]}
{"Identifier":"2022ApJ...937...97C__Bottrell_et_al._2019_Instance_1","Paragraph":"Recently, machine learning (ML) has been applied to derive various physical parameters of galaxies (e.g., Masters et al. 2015; Krakowski et al. 2016; D\u2019Isanto & Polsterer 2018; Bonjean et al. 2019; Davidzon et al. 2019; Hemmati et al. 2019; Chang et al. 2021) and improves on linear combinations through nonlinear activations (e.g., Ackermann et al. 2018; Walmsley et al. 2019; Ferreira et al. 2020; Bickley et al. 2021; Bickley et al. 2022; Ferreira et al. 2022). In particular, classification by ML (e.g., Banerji et al. 2010; Huertas-Company et al. 2015; Dom\u00ednguez S\u00e1nchez et al. 2018; Bottrell et al. 2019; Pearson et al. 2019; Barchi et al. 2020; Chang et al. 2021) can avoid time-consuming visual inspections and will be helpful for the visual classification of galaxy\u2212galaxy interactions from the forthcoming large surveys. For instance, Pearson et al. (2019) developed a convolutional neural network (CNN) architecture with observational SDSS and simulated EAGLE images to identify galaxy mergers. They showed that the networks achieve better performance in observational data than in simulations. Ferreira et al. (2020) achieve 0.90 accuracy to classify major mergers and measure galaxy mergers in all five CANDELS fields using CNN trained with simulated galaxies from the IllustrisTNG simulation and separate star-forming galaxies from post-mergers in a following work (Ferreira et al. 2022). Bickley et al. (2022) deployed a CNN and evaluated mock observations of simulated galaxies from the IllustrisTNG simulations to identify post-mergers. Bottrell et al. (2022) examine both the morphological and kinematic features of merger remnants from the TNG100 and show that the stellar kinematic data have few contributions. Moreover, it has been discussed whether ancillary information such as kinematics and spectroscopic information in addition to the images may provide an additional basis for classification (e.g., Nevin et al. 2019; Pan et al. 2019; Bottrell et al. 2022; McElroy et al. 2022). Therefore, it is important to identify specific features for classification with ML from both photometric and spectroscopic observables.","Citation Text":["Bottrell et al. 2019"],"Functions Text":["In particular, classification by ML (e.g.,","can avoid time-consuming visual inspections and will be helpful for the visual classification of galaxy\u2212galaxy interactions from the forthcoming large surveys."],"Functions Label":["Future Work","Future Work"],"Citation Start End":[[589,609]],"Functions Start End":[[465,507],[671,830]]}
{"Identifier":"2018ApJ...864...31A__Smith_2007_Instance_1","Paragraph":"The FORCAST images of HD 168625 are shown in Figure 4. The nebula is clearly resolved, with a partially complete ring structure that has two peaks almost symmetric around the star. We concur with Meixner et al. (1999) and O\u2019Hara et al. (2003) in their interpretation of these two peaks as limb-brightened peaks of a torus of dust with a radius of \u223c10\u2033. The appearance is consistent with previously obtained images at 8.8, 12.5, and 20.6 \u03bcm from Meixner et al. (1999) and PACS 70 \u03bcm images (Groenewegen et al. 2011). We stress that our SOFIA\/FORCAST images do not detect the outer polar rings seen in Spitzer IRAC images (Smith 2007), suggesting that the rings must be cold and below the sensitivity limits of SOFIA\/FORCAST. The emission detected with the ring morphology in the IRAC band 4 image was probably polycyclic aromatic hydrocarbon (PAH) emission or atomic line emission, not thermal emission from warm dust. Figure 5 shows the temperature map that was derived from stacking the \u03bbF\u03bb SOFIA\/FORCAST 7.7\u201337.1 \u03bcm images and performing a least-squares fit of the dust temperature, Td, using the best-fit modified blackbody of the SED (i.e., B\u03bb(Td) \u00b7 \u03bb0.33) at each pixel location. The images were centered relative to one another by comparing the locations of the limb brightness peaks, and the 7.7\u201333.6 \u03bcm images were convolved with a 2D Gaussian kernel with an FWHM of 35 to match the resolution of the 34.8 and 37.1 \u03bcm images. Our temperature map shows a large gradient in dust temperatures with inner torus temperatures of \u223c180 K and outer temperatures of \u223c80 K, which is in agreement with the estimate of 170 \u00b1 40 K obtained from our least-squares fit to the IR excess but is slightly higher than the equilibrium temperature estimates made by Pasquali et al. (2002; 113 K), Robberto & Herbst (1998; 135 K), and O\u2019Hara et al. (2003; 130 K). The inaccuracies of this temperature map are due in large part to the fact that the method used to create it assumes that emission is purely thermal and that the dust shell is in thermal equilibrium. As noted previously, much of the 8.8\u201312.5 \u03bcm flux arises from transient, nonequilibrium emission from PAH grains. Therefore, using images at these wavelengths to derive quantitative conclusions from the temperature map has some limitations; however, we can interpret the maps qualitatively as discussed in Section 4.","Citation Text":["Smith 2007"],"Functions Text":["We stress that our SOFIA\/FORCAST images do not detect the outer polar rings seen in Spitzer IRAC images"],"Functions Label":["Differences"],"Citation Start End":[[621,631]],"Functions Start End":[[516,619]]}
{"Identifier":"2018AandA...620L...8H__Gundlach_&_Blum_2013_Instance_1","Paragraph":"We adopted the model developed by \u010eapek et al. (2005), in which 1D thermal conduction below each of the surface facets is solved numerically with the nonlinear Robin boundary condition at the surface, and the assumption of an isothermal core at a sufficient depth is made. A temperature-dependence of the thermal conductivity following Eq. (2) was used. For the sake of simplicity, the specific heat capacity c was assumed constant, c\u2004=\u2004560 J kg\u22121 K\u22121, and the regolith grain density obtained for C-type meteorites was used, \u03c1\u2004=\u20043.11 g cm\u22123 (both from Gundlach & Blum 2013). We ran solutions for four values of the packing factor \u03d5 in the range between 0.3 and 0.6. Each time, the parameters of the thermal conductivity were adjusted to satisfy the constraints from thermal observations described in Sect. 3.3 (see Table 2). The time domain of one revolution about the Sun was divided into steps of 60 s, short enough when compared to the \u22433.6 h rotation period, and the space grid describing the depth below each of the surface increased exponentially, as described in \u010eapek et al. (2005). We ensured that at each depth, the von Neumann stability condition was satisfied. Typically, ten iterative steps of the algorithm provide the temperature with an accuracy of one degree or better in the whole space and time domain of the solution. The shape and spin state of Phaethon was taken from the modeling in Sect. 3.1. Similarly, the volume-equivalent size of 5.1 km from Sect. 3.2 was used as an implicit value. The last parameter required to compute the thermal recoil acceleration (the Yarkovsky effect) is the bulk density of Phaethon. Our nominal models use 1 g cm\u22123 for the clarity, but we treated this value as a free parameter in the orbit determination process (similarly to what was done for asteroid Bennu in Chesley et al. 2014). Scaling to different densities is easily implemented by using the inverse-proportional dependence of the thermal acceleration on the bulk density. In our analysis we neglected the enhancement of the Yarkovsky effect that is due to surface roughness (Rozitis & Green 2012). This effect could cause an increase in our bulk density estimate of less than 10%, which is well within the formal uncertainty.","Citation Text":["Gundlach & Blum 2013"],"Functions Text":["For the sake of simplicity, the specific heat capacity c was assumed constant, c\u2004=\u2004560 J kg\u22121 K\u22121, and the regolith grain density obtained for C-type meteorites was used, \u03c1\u2004=\u20043.11 g cm\u22123 (both from"],"Functions Label":["Uses"],"Citation Start End":[[552,572]],"Functions Start End":[[354,551]]}
{"Identifier":"2022MNRAS.517.4202X__Lister_et_al._2018_Instance_1","Paragraph":"Cavagnolo et al. (2010) suggested that the jet kinetic power is able to inflate the X-ray cavities or bubbles in different systems, including giant elliptical galaxies and cD galaxies (Type cD galaxy, a subtype of type-D giant elliptical galaxy), and proposed to evaluate the kinetic power Pkin = Pcav. However, this method is only limited to a small number of sources at present. It is known that the luminosity of extended region of radio jet, which is believed to be less Doppler-boosted, is related to jet kinetic power (Rawlings & Saunders 1991; Willott et al. 1999; Cavagnolo et al. 2010; Meyer et al. 2011), $P_{\\rm rad} = \\eta \\, L_{\\rm 5GHz}^{\\rm ext}{}^{\\kappa }$. Though the factor \u03ba and \u03b7 are given in discrepancy in literature due to the different sizes and source types of sample (Cavagnolo et al. 2010; Meyer et al. 2011), the ${\\rm log}\\, L^{\\rm ext}_{\\rm 5GHz}$ scales with the ${\\rm log}\\, P_{\\rm rad}$ in the logarithmic space. We collect the total radio flux density from literature (Taylor et al. 1996 at 5 GHz; Piner & Edwards 2014 at 8.4 GHz; and Lister et al. 2018 at 15 GHz), and we convert the data at other frequencies to 5 GHz by assuming that\n(14)$$\\begin{eqnarray*}\r\nS_{\\rm 5GHz}^{\\rm core} = S_{\\rm \\nu }^{\\rm core} \\, \\, {\\rm and} \\, \\, S_{\\rm 5GHz}^{\\rm ext} = S_{\\rm \\nu }^{\\rm ext} \\left(\\frac{\\nu }{\\rm 5 \\, GHz} \\right)^{\\alpha _{\\rm ext}},\r\n\\end{eqnarray*}$$where the total radio flux is the sum of the flux of core and the flux of the extended region, Stot = Score + Sext, the \u03b1ext = 0.75 and \u03b1core = 0 (Fan et al. 2011; Pei et al. 2016, 2019, 2020). Together with the radio-core dominance parameter at 5 GHz\n(15)$$\\begin{eqnarray*}\r\nR = \\left(\\frac{S^{\\rm core}}{S^{\\rm ext}} \\right) (1+z)^{\\alpha _{\\rm core}-\\alpha _{\\rm ext}},\r\n\\end{eqnarray*}$$that we collect from Pei et al. (2020), we obtain $S_{\\rm 5GHz}^{\\rm ext}$ and calculate ${\\rm log}\\, L^{\\rm ext}_{\\rm 5GHz}$. The correlations of ${\\rm log}\\, P_{\\rm rad}$ and ${\\rm log}\\, L_{\\rm 5GHz}^{\\rm ext}$ against ${\\rm log}\\, \\beta _{\\rm app}^{\\rm max}$ are illustrated in Fig. 6 and linear regression results are listed in Table 2. Positive correlations of ${\\rm log}\\, P_{\\rm rad}\\, versus \\, {\\rm log}\\, \\beta _{\\rm app}^{\\rm max}$ and ${\\rm log}\\, L_{\\rm 5GHz}^{\\rm ext}\\, versus \\, {\\rm log}\\, \\beta _{\\rm app}^{\\rm max}$ are found for blazars. The positive correlation of ${\\rm log}\\, P_{\\rm rad}\\, versus \\, {\\rm log}\\, \\beta _{\\rm app}^{\\rm max}$ holds for both FSRQs and BL Lacs when we consider them independently, while the positive correlation of ${\\rm log}\\, L_{\\rm 5GHz}^{\\rm ext}\\, versus \\, {\\rm log}\\, \\beta _{\\rm app}^{\\rm max}$ only hold for BL Lacs. It is found that the motion of jet knots is significantly correlated with jet radiation power for both FSRQs and BL Lacs, however, the motion of jet knots is correlated with the kinetic power only for BL Lacs.","Citation Text":["Lister et al. 2018"],"Functions Text":["We collect the total radio flux density from literature"],"Functions Label":["Uses"],"Citation Start End":[[1070,1088]],"Functions Start End":[[947,1002]]}
{"Identifier":"2015MNRAS.453.3414A__the_1999_Instance_2","Paragraph":"Filippenko & Chornock (2001) first presented the dynamical estimate of mass of the source to be around 7.4 \u00b1 1.1\u2009M\u2299. Recently, Radhika & Nandi (2014) claimed that the mass of XTE J1859+226 is perhaps in between 6.58 and 8.84\u2009M\u2299 which is similar to the prediction of Shaposhnikov & Titarchuk (2009), although the lower mass limit is estimated as 5.4\u2009M\u2299 by Corral-Santana et al. (2011). However, we consider the typical mass of the source as 7\u2009M\u2299. The distance of this source is around d \u223c 11\u2009kpc (Filippenko & Chornock 2001). Steiner et al. (2013) measured the spin as ak \u223c 0.4; however, Motta et al. (2014b) recently reported that the spin of the source is ak \u223c 0.34. Since the spin predictions are quite close, we use ak \u223c 0.4 for this analysis. We estimate the fluxes Fx (see Table 1) of LHS and HIMS of the 1999 outburst of the source (Radhika & Nandi 2014). The corresponding disc luminosities are calculated as $L_{\\rm disc}^{{\\rm LHS}}=8.26 \\times 10^{37}\\ {\\rm erg\\ s^{-1}}$ and $L_{\\rm disc}^{{\\rm HIMS}}=1.85 \\times 10^{38}\\ {\\rm erg\\ s^{-1}}$, respectively. Now, it is reasonable to assume the accretion efficiency for rotating BH as \u03b7 = 0.3 which corresponds to the accretion rate of the inflowing matter as ${\\dot{M}}_{{\\rm acc}}^{{\\rm LHS}} = 0.304 {\\dot{M}}_{{\\rm Edd}}$ in LHS and ${\\dot{M}}_{{\\rm acc}}^{{\\rm HIMS}} = 0.680 {\\dot{M}}_{{\\rm Edd}}$ in HIMS. For LHS, we use $R_{\\dot{m}}=9.83$\u2009per\u2009cent following our theoretical estimate where xs = 64.6rg for ak = 0.4, ${\\mathcal {E}}=0.001\\,98$ and \u03bb = 3.18. Incorporating these inputs in equation (15), we obtain the jet kinetic power as $L^{{\\rm LHS}}_{{\\rm jet}} = 2.52\\times 10^{37}\\ {\\rm erg\\ s^{-1}}$. The maximum mass outflow rate for HIMS corresponding to ak = 0.4 is obtained from Fig. 9 as $R^{\\rm max}_{\\dot{m}}=17.5$\u2009per\u2009cent for ${\\mathcal {E}}=0.005\\,47$ and \u03bb = 3.1, where the shock transition occurs at 21.9rg. Using these values in equation (15), we obtain the maximum jet kinetic power as $L^{{\\rm HIMS}}_{{\\rm jet}} = 1.08\\times 10^{38}\\ {\\rm erg\\ s^{-1}}$ which we regard to be associated with the HIMS of this source.","Citation Text":["Radhika & Nandi 2014"],"Functions Text":["We estimate the fluxes Fx (see Table 1) of LHS and HIMS of the 1999 outburst of the source"],"Functions Label":["Uses"],"Citation Start End":[[839,859]],"Functions Start End":[[747,837]]}
{"Identifier":"2017MNRAS.472..940K__Malu_et_al._2010_Instance_1","Paragraph":"Cluster mergers can stir the intra-cluster medium (ICM) and lead to complex distributions of density and temperatures. The X-ray surface brightness traces regions of high electron densities ($\\propto n_e^{2}$) and the SZ is sensitive to the pressure (\u221dneT) along the line of sight. An offset in the peaks of these signals can be used as an indicator of the density and temperature distribution in the disturbed ICM. There are examples of merging clusters that show presence of X-ray\u2013SZ offsets, such as Abell 2146 (AMI Consortium et al. 2011) and Bullet cluster (Malu et al. 2010). Simulations have shown that the offset is sensitive to initial relative velocity of the merging clusters and the mass ratio (e.g. Molnar, Hearn & Stadel 2012; Zhang, Yu & Lu 2014). The cluster PLCK\u2009G200.9\u221228.2\u2009\u2009discussed in this work has its X-ray peak and the Planck detection peak offset by 3.4\u2009arcmin, which is the extreme in the Planck\u2009\u2009sample (PC12). The Planck\u2009\u2009SZ positions have mean and median errors of 1.5 and 1.3\u2009arcmin, respectively (Planck Collaboration IV 2013). The X-ray peak is located at the northern sub-cluster [Fig. 1(a) and Fig. 6(a)] and the Planck\u2009\u2009SZ position is separated from it by 3.4\u2009arcmin (700\u2009kpc) in the direction of the radio relic. Due to the presence of a shock at the relic, the region is expected to be overpressured and thus may result in shifting the peak of the SZ-signal. The offset in the direction of the relic indicates possible physical origin for the offset in addition to the position reconstruction uncertainty of Planck. Based on the results of simulations, the offset can be explained as a result of two comparable mass sub-clusters with mass ratio between 1 and 3 (Zhang et al. 2014). Deep optical observations tracing the galaxy distribution in this cluster will be useful to measure the mass ratios of the sub-cluster masses. Due to the large uncertainty in the Planck\u2009\u2009position we cannot analyse the offsets for a statistical sample of merging clusters. However, the offset in the X-ray and SZ positions opens an additional probe for understanding the properties of merging galaxy clusters.","Citation Text":["Malu et al. 2010"],"Functions Text":["There are examples of merging clusters that show presence of X-ray\u2013SZ offsets, such as","and Bullet cluster"],"Functions Label":["Background","Background"],"Citation Start End":[[563,579]],"Functions Start End":[[416,502],[543,561]]}
{"Identifier":"2016ApJ...821..107G__Schwadron_et_al._2011_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":["Schwadron et al. 2011"],"Functions Text":["For the density of neutral hydrogen in the inner heliosheath a constant nH = 0.1 cm\u22123 is assumed"],"Functions Label":["Uses"],"Citation Start End":[[1151,1172]],"Functions Start End":[[1053,1149]]}
{"Identifier":"2021MNRAS.500.5614A__Adri\u00e1n-Mart\u00ednez_et_al._2013_Instance_1","Paragraph":"Using ANTARES data from the end of 2007\u20132017, a search for upward-going muon neutrinos and antineutrinos in spatial and temporal coincidence with 784 GRBs has been performed. The numerical model NeuCosmA was used to estimate the expected neutrino flux from each burst individually, in the context of one-zone internal shock model. A novel aspect of the search here presented is the inclusion in the data analysis chain of the uncertainty that possible unknown parameters, related to the characteristic activity of the central engine, can introduce in the neutrino flux evaluation. This is crucial in order to correctly interpret the validity of model-dependent results, in terms of upper limits set by non-detections of neutrinos in coincidence with GRBs (Adri\u00e1n-Mart\u00ednez et al. 2013; Aartsen et al. 2017). These parameters have been identified in the bulk Lorentz factor, variability time-scale, and source redshift, all of which are affecting the so-called dissipation radius, where shell collisions are realized. Among these parameters, the former was shown to impact the most GRB-neutrino flux predictions. At the same time, it is also possible to marginalize the uncertainty related to it by assuming a correlation with the source isotropic gamma-ray luminosity (which is in turn a physical observable). This was realized by relying upon the observational correlation found by L\u00fc et al. (2012). As a result of such procedure, the minimum variability time-scale was found to contribute more than redshift to the uncertainty on the neutrino flux predictions from GRBs. Indeed, when letting tv free to vary, the estimated uncertainty on the neutrino flux expected from the model is observed to span up to several orders of magnitude. As a consequence, the expected \u03bd-fluxes are provided with an uncertainty band of \u00b12\u03c3. Analogously to previous ANTARES searches (Adri\u00e1n-Mart\u00ednez et al. 2013, 2017b; Celli et al. 2017), MC simulations of the signal predicted by NeuCosmA were performed, while the respective background was estimated directly from off-source data collected by ANTARES. Only track-like events reconstructed within 10\u00b0 in radius from the expected GRB position were selected and in temporal correlation with the prompt gamma-ray emission.","Citation Text":["Adri\u00e1n-Mart\u00ednez et al. 2013","Adri\u00e1n-Mart\u00ednez et al. 2013"],"Functions Text":["A novel aspect of the search here presented is the inclusion in the data analysis chain of the uncertainty that possible unknown parameters, related to the characteristic activity of the central engine, can introduce in the neutrino flux evaluation. This is crucial in order to correctly interpret the validity of model-dependent results, in terms of upper limits set by non-detections of neutrinos in coincidence with GRBs","Analogously to previous ANTARES searches","MC simulations of the signal predicted by NeuCosmA were performed, while the respective background was estimated directly from off-source data collected by ANTARES."],"Functions Label":["Compare\/Contrast","Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[756,783],[1864,1891]],"Functions Start End":[[331,754],[1822,1862],[1920,2084]]}
{"Identifier":"2019MNRAS.487.5902K__Behroozi,_Wechsler_&_Conroy_2013_Instance_1","Paragraph":"In the 50 Myr after a star particle is formed, it can undergo supernova (SN) explosions. These are randomly drawn from a delay-time distribution (Kimm et al. 2015). We employ the mechanical feedback model described in Kimm et al. (2015, 2017) and Rosdahl et al. (2018) and the amount of momentum injected into the gas depends on the phase of the SN that is resolved by the simulation as to capture the final momentum of the snowplow phase. The equivalent of 1051\u2009ergs is injected into the gas for each SN. The maximum momentum that we inject is boosted according to Geen et al. (2015) to account for unresolved H\u2009ii regions (Kimm et al. 2017). For each massive star that explodes, 20 per\u2009cent of the mass is recycled back into the gas. This gas is metal enriched assuming a metallicity of 0.075. Following Rosdahl et al. (2018), we have calibrated the SN feedback in order to reproduce the high-redshift stellar mass\u2013halo mass relation from abundance matching (Behroozi, Wechsler & Conroy 2013). For our simulation, this requires assuming that the mean SN progenitor mass is 5\u2009M\u2299 which leads to 4\u00d7 more SN on average compared to a standard Kroupa IMF (Kroupa 2001). While not ideal, this results in galaxies that fall nicely on the stellar mass\u2013halo mass relation (see fig. 1 of Katz et al. 2018) and produces a UV luminosity function consistent with observations for a similar set of simulations at approximately the same resolution (Rosdahl et al. 2018). While we cannot be certain that even our calibrated simulations have the correct stellar mass\u2013halo mass relation as we are using a high-redshift extrapolation and there is intrinsically a lot of scatter, we have made an effort to calibrate on one of the best available predictions as any significant offset from the stellar mass\u2013halo mass relation may lead to large systematic offsets in the SFR-line luminosity relations. In summary, our star formation and stellar feedback models are based on the work of Rosdahl et al. (2018) which were chosen to reproduce both a reasonable reionization history and UV luminosity function.","Citation Text":["Behroozi, Wechsler & Conroy 2013"],"Functions Text":["Following Rosdahl et al. (2018), we have calibrated the SN feedback in order to reproduce the high-redshift stellar mass\u2013halo mass relation from abundance matching"],"Functions Label":["Uses"],"Citation Start End":[[961,993]],"Functions Start End":[[796,959]]}
{"Identifier":"2017ApJ...850L..40A__Yang_et_al._2017_Instance_2","Paragraph":"Aided by the tight localization constraints of the three-detector network and the proximity of the GW source, multiple independent surveys across the EM spectrum were launched in search of a counterpart beyond the sGRB (Abbott et al. 2017c). Such a counterpart, SSS17a (later IAU-designated AT 2017gfo), was first discovered in the optical less than 11 hours after merger, associated with the galaxy NGC 4993 (Coulter et al. 2017a, 2017b), a nearby early-type E\/S0 galaxy (Lauberts 1982). Five other teams made independent detections of the same optical transient and host galaxy all within about one hour and reported their results within about five hours of one another (Allam et al. 2017; Arcavi et al. 2017a, 2017b; Lipunov 2017b; Tanvir & Levan 2017; Yang et al. 2017; Soares-Santos et al. 2017; Lipunov et al. 2017a). The same source was followed up and consistently localized at other wavelengths (e.g., Corsi et al. 2017; Deller et al. 2017a, 2017b, 2017c; Goldstein et al. 2017; Haggard et al. 2017a, 2017b; Mooley et al. 2017; Savchenko et al. 2017; Alexander et al. 2017; Haggard et al. 2017c; Goldstein et al. 2017; Savchenko et al. 2017). The source was reported to be offset from the center of the galaxy by a projected distance of about 10\u2033 (e.g., Coulter et al. 2017a, 2017b; Haggard et al. 2017a, 2017b; Kasliwal et al. 2017; Yang et al. 2017; Yu et al. 2017). NGC 4993 has a Tully\u2013Fisher distance of \u223c40 Mpc (Freedman et al. 2001; NASA\/IPAC Extragalactic Database164\n\n164\nThe NASA\/IPAC Extragalactic Database (NED) is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration.\n), which is consistent with the luminosity distance measurement from gravitational waves (\n\n\n\n\n\n Mpc). Using the Tully\u2013Fisher distance, the \u223c10\u2033 offset corresponds to a physical offset of \u22432.0 kpc. This value is consistent with offset measurements of sGRBs in other galaxies, though below the median value of \u223c3\u20134 kpc (Fong et al. 2010; Fong & Berger 2013; Berger 2014).","Citation Text":["Yang et al. 2017"],"Functions Text":["The same source was followed up and consistently localized at other wavelengths (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[1343,1359]],"Functions Start End":[[824,910]]}
{"Identifier":"2017ApJ...850...20G__Lonardoni_et_al._2015_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":["Lonardoni et al. 2015"],"Functions Text":["On the one hand, it is more energetically favorable for the system to populate new degrees of freedom, such as hyperons","in order to lower its Fermi energy (starting at about two times the saturation density)."],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[814,835]],"Functions Start End":[[542,661],[1345,1433]]}
{"Identifier":"2019MNRAS.484.4083H__Kappenman_2006_Instance_1","Paragraph":"The primary impact of the 1859 storm was on telegraphy (e.g. Boteler 2006). Today the principal \u2018space weather\u2019 threat is to the electric power grid (Baker et al. 2008; Hapgood, 2011, 2012; Oughton et al. 2016; Dyer et al. 2018). Because of this threat, several studies have been carried out to estimate how frequently such extreme space weather events may occur (e.g. Tsubouchi & Omura 2007; Riley 2012; Curto, Castell & Del Moral 2016; Riley & Love 2017). Such studies are dependent on the observations of a handful of storms that have approached or rivalled the Carrington geomagnetic storm in observed\/inferred intensity and\/or auroral extent. These include great storms in February 1872 (Chapman 1957a,b; Silverman 2008; Hayakawa et al. 2018c); September 1909 (Silverman 1995; Willis, Stephenson & Fang 2007; Love et al. 2019); May 1921 (Silverman & Cliver 2001; Kappenman 2006; Cliver & Dietrich 2013). The only such storm observed during the space era occurred in March 1989 (Allen et al. 1989; Silverman 2006a; Pulkkinen et al. 2012; Cid et al. 2014). More recently, in July 2012, a major backside eruption on the Sun was observed both remotely and in situ by the STEREO spacecraft (Kaiser et al. 2008; Russell et al. 2013; Liu et al. 2014; Riley et al. 2016). The interplanetary coronal mass ejection (CME) propagated to 1\u2009AU in \u223c20\u2009h. Baker et al. (2013) calculated that had the eruption occurred on the frontside of the Sun \u2013 with optimal seasonal and local timing to maximize solar wind \u2013 magnetosphere coupling (worst case scenario (see e.g. Temerin & Li 2002) \u2500 it might have produced a storm greater than that inferred for the Carrington event. In addition to these storms, auroral evidence has recently been provided for two pre-1859 storms that may have ranked with the Carrington event: February 1730 (Hayakawa et al. 2018a), and, particularly, September 1770 (Willis et al. 1996; Nakazawa, Okada & Shiokawa 2004; Ebihara et al. 2017; Hayakawa et al. 2017e).","Citation Text":["Kappenman 2006"],"Functions Text":["Such studies are dependent on the observations of a handful of storms that have approached or rivalled the Carrington geomagnetic storm in observed\/inferred intensity and\/or auroral extent. These include great storms in","May 1921"],"Functions Label":["Background","Background"],"Citation Start End":[[868,882]],"Functions Start End":[[458,677],[833,841]]}
{"Identifier":"2020MNRAS.499.3792B__Pimbblet_2011_Instance_1","Paragraph":"For consistency, we adopt the translation of this to the absolute velocities of cluster galaxies normalized by their respective galaxy cluster velocity dispersions into the range $0.3 \\lt |\\Delta \\mathrm{V}|\/\\sigma _{r_{200}} \\lt 0.5$ as deduced by Pimbblet (2011). Thus, if the mode of the standardised velocities for a sub-population has its foci at around $0.3 \\lt |\\Delta \\mathrm{V}|\/\\sigma _{r_{200}} \\lt 0.5$ for values around the virial radius, which we assume to be Rvirial \u223c r200, said sub-population would be classified as infalling. In contrast, a sub-population of backsplash cluster galaxies would be expected to peak significantly at $|\\Delta \\mathrm{V}|\/\\sigma _{r_{200}}\\sim 0$ for values at or beyond our definition of the virial radius, with their fraction reaching zero at some upper limit (e.g. Mamon et al. 2004; Pimbblet 2011; Bah\u00e9 et al. 2013; Haggar et al. 2020). Therefore, with respect to Fig. 5, we see that the column of our non-merging sub-populations across both bins of radius do not show any significant difference in the distributions of velocities with the exception of those that lie \u2264r200, which show the non-AGN sub-population to occupy a mode within the range that nominally represents infallers, most likely for cluster galaxies 0.5 \u2264 r200  1.0 (Gill et al. 2005). Additionally, the AGN sub-population slightly deviates from the non-AGN velocity distribution with a mode centred at $|\\Delta \\mathrm{\\it V}|\/\\sigma _{r_{200}}\\sim 0.8$, which could indicate stronger infalling. In contrast, the column of our merging AGN sub-populations shows the strongest deviations from the distribution of non-AGN, especially with the >r200 bin showing a significant centrally dominated AGN sub-population, where such a central dominance in relative velocity corresponds to a sub-population that were predominantly backsplash cluster galaxies. However, the dependence of this being the true nature of the sub-population relies upon more precise definitions of the radii since there is a natural upper limit a bound cluster galaxy can extend outward to with respect to its galaxy cluster\u2019s potential, known as the splashback radius (More et al. 2015, 2016). In addition, Haggar et al. (2020) show that the fraction of backsplash galaxies diminishes by 2r200 and 2.5r200 for massive (\u223c\u00d71015M\u2299) merging and non-merging cluster systems, respectively, thus demonstrating that merging cluster environments experience a greater decrease in the fraction of harbouring backsplash galaxies as one continues to extend beyond r200. Indeed, the sub-populations of the merging cluster galaxies present in the \u2264r200 bin show more variations in their general distributions with the modes of both the AGN and non-AGN sub-populations lying around $0.3 \\lt |\\Delta \\mathrm{V}|\/\\sigma _{r_{200}} \\lt 0.5$, which eludes to mostly infalling sub-populations rather than those associated with backsplash. Finally, if one considers the equivalent peak of the AGN density histogram at $\\Delta \\mathrm{V}|\/\\sigma _{r_{200}}\\sim 1.7$ it could be possible there is a mix of recently accreted cluster galaxies and those that are relaxing on to a common potential. Although it should be noted that not much information can be confidently derived from the AGN sub-populations within the bins that possess small samples size (N \u2272 100), especially with the merging AGN-hosting cluster galaxies at \u2264r200 that only has N = 15.","Citation Text":["Pimbblet (2011)"],"Functions Text":["For consistency, we adopt the translation of this to the absolute velocities of cluster galaxies normalized by their respective galaxy cluster velocity dispersions into the range $0.3 \\lt |\\Delta \\mathrm{V}|\/\\sigma _{r_{200}} \\lt 0.5$ as deduced by"],"Functions Label":["Uses"],"Citation Start End":[[249,264]],"Functions Start End":[[0,248]]}
{"Identifier":"2021AandA...651A.111P__Herrera-Camus_et_al._2018_Instance_1","Paragraph":"Irrespective of its origin, the [C\u202fII] emission is linked to the presence of stellar far-ultraviolet (FUV) photons (E 13.6 eV). As FUV photons are tied to the presence of massive O and B stars that have short lifetimes, the [C\u202fII] 158 \u03bcm line is also astar formation rate (SFR) indicator. Indeed, ISO, Herschel and SOFIA observations have demonstrated the good correlation between the [C\u202fII] luminosity and the SFR in the Milky Way and in regions of massive star formation in other galaxies (e.g., Kramer et al. 2013, 2020; Pineda et al. 2014, 2018; Herrera-Camus et al. 2015, 2018; De Looze et al. 2011). With ALMA and NOEMA, ground-based observations of the [C\u202fII] 158 \u03bcm line in high redshift galaxies have come into reach and such data are routinely used to infer SFRs (e.g., Walter et al. 2012; Venemans et al. 2012; Knudsen et al. 2016; Bischetti et al. 2018; Khusanova et al. 2021) based upon validations of this relationship in the nearby Universe (Herrera-Camus et al. 2018; De Looze et al. 2011). However, it is well-understood that the intensity of the [C\u202fII] line depends on the local physical conditions (Hollenbach & Tielens 1999). Observationally, the presence of the so-called [C\u202fII]-deficit \u2013 a decreased ratio of [C\u202fII] 158 \u03bcm luminosity to FIR dust continuum with an increasing dust color temperature and also with FIR luminosity \u2013 is well established (Malhotra et al. 2001; D\u00edaz-Santos et al. 2013; Magdis et al. 2014; Smith et al. 2017). This deficit is particularly pronounced in (local) ultraluminous infrared galaxies (ULIRGs), very dusty galaxies characterized by vigorous embedded star formation (e.g., Luhman et al. 2003; Abel et al. 2009; Graci\u00e1-Carpio et al. 2011). This deficit, however, does not necessarily hold in the early Universe at high redshift (e.g., Stacey et al. 2010; Brisbin et al. 2015; Capak et al. 2015). Some studies have indicated that not only [C\u202fII] emission is deficient in some sources, but other FIR cooling lines ([O\u202fI], [O\u202fIII], [N\u202fII]), as well (e.g., Graci\u00e1-Carpio et al. 2011; Herrera-Camus et al. 2018). These deficits must be linked to the global ISM properties and star-formation characteristics in these galaxies.","Citation Text":["Herrera-Camus et al. 2018"],"Functions Text":["With ALMA and NOEMA, ground-based observations of the [C\u202fII] 158 \u03bcm line in high redshift galaxies have come into reach and such data are routinely used to infer SFRs","based upon validations of this relationship in the nearby Universe"],"Functions Label":["Background","Background"],"Citation Start End":[[957,982]],"Functions Start End":[[606,772],[889,955]]}
{"Identifier":"2018ApJ...853..131L__Renzini_2009_Instance_1","Paragraph":"Classifying galaxies into different star formation regimes at high redshift is facilitated by the fact that star formation rate (SFR) and stellar mass (M*) of star-forming galaxies (SFGs) are strongly correlated out to at least \n\n\n\n\n\n (Daddi et al. 2007; Noeske et al. 2007; Pannella et al. 2009, 2015; Elbaz et al. 2011; Whitaker et al. 2012, 2014; Lee et al. 2015; Salmon et al. 2015; Schreiber et al. 2015; Tomczak et al. 2016). This correlation is commonly called the \u201cmain sequence of star formation\u201d (MS). A common interpretation of the MS is that the location of galaxies relative to the MS follows a different time evolution of SFR (Renzini 2009; Daddi et al. 2010; Rodighiero et al. 2011; Sargent et al. 2012; Renzini & Peng 2015). The tight MS with near unity slope reflects that the majority of SFGs follow a steadily increasing star formation history governed by a set of gradual physical processes like gas exhaustion (Noeske et al. 2007). A small fraction of galaxies exhibit quasi-exponential mass and SFR growth, either through major mergers or through strong bursts of star formation in the densest regions (Elbaz et al. 2011; Sargent et al. 2012). While typical galaxies therefore spend most of their time on the MS prior to additional quenching processes, these starburst galaxies are located above the MS and play a relatively minor role in the star formation history of the universe (Rodighiero et al. 2011). Galaxies located below the MS include quiescent galaxies (QGs), with spheroidal-like structures and little star formation activity, as well as fading SFGs with diminishing star formation activity. The transient galaxies, such as those in the green valley, can dominate the lower region of the MS. At \n\n\n\n\n\n, green valley galaxies are known to be off the MS (Schawinski et al. 2014), and they have intermediate morphologies combining disk-dominated and bulge-dominated systems (Salim et al. 2009; Mendez et al. 2011; Pandya et al. 2017).","Citation Text":["Renzini 2009"],"Functions Text":["A common interpretation of the MS is that the location of galaxies relative to the MS follows a different time evolution of SFR"],"Functions Label":["Background"],"Citation Start End":[[641,653]],"Functions Start End":[[512,639]]}
{"Identifier":"2017MNRAS.464.2120S__Diego_et_al._2015_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":["Diego et al. 2015"],"Functions Text":["Several studies have already measured the distance ratios (e.g.","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."],"Functions Label":["Background","Differences"],"Citation Start End":[[931,948]],"Functions Start End":[[847,910],[1018,1333]]}
{"Identifier":"2021AandA...645A..96P__Marcantonio_et_al._(2018)_Instance_1","Paragraph":"The ESPRESSO DFS concept (Di Marcantonio et al. 2018) was conceived during its preliminary design phases with the goal of maximizing operational efficiency, flexibility, and scientific output while complying with the standard Paranal Observatory operational scheme. The main challenge derives from the requirement to operate ESPRESSO in a seamless way with any of the UT\u2019s or with all four UT\u2019s simultaneously. This must be possible not only with a predetermined schedule, but also \u201con the fly\u201d. The flexibility in ESPRESSO\u2019s operations has been tackled by adopting a new DFS deployment plan described in Di Marcantonio et al. (2018) that is exceptional under various aspects because it has to cope with various telescope and instrument configurations while remaining operationally simple. Figure 8 shows the main ESPRESSO DFS elements and their final deployment. Besides the software packages already described, part of the software for the control of the CT devices has been incorporated into the VLT telescope CS to allow CT operations even when ESPRESSO is offline (thus avoiding conflicts, e.g., with instruments of the VLT Interferometer operations). In addition to the standard DFS software packages, ESPRESSO is the first instrument to also provide a data analysis package that is able to extract relevant astronomical observables from the reduced data. The following DFS subsystems are specific to ESPRESSO: (1) the ETC hosted on the ESO web page9; (2) the CS with the full suite of acquisition, observation, and calibration templates that are able to control all vital parts of the instrument and the CT (Calderone et al. 2018); (3) the data reduction software (DRS) package (or \u201cpipeline\u201d) capable of providing \u201cscience-ready\u201d reduced data only minutes after the end of the individual observation; (4) the data analysis software (DAS) package that produces higher-level astronomical observables with no or limited supervision; (5) the DRS and DAS are distributed to the community10.","Citation Text":["Di Marcantonio et al. 2018"],"Functions Text":["The ESPRESSO DFS concept","was conceived during its preliminary design phases with the goal of maximizing operational efficiency, flexibility, and scientific output while complying with the standard Paranal Observatory operational scheme."],"Functions Label":["Background","Background"],"Citation Start End":[[26,52]],"Functions Start End":[[0,24],[54,265]]}
{"Identifier":"2019MNRAS.490.3588M__Burgio_et_al._2018_Instance_1","Paragraph":"In the era of gravitational wave and multimessenger astronomy of binary neutron stars accurate numerical modelling of neutron-star mergers and their remnants on long time-scales $\\simeq 1\\, \\rm s$ has never been more important. The coincident detection of a gravitational-wave signal from a neutron-star merger (The LIGO Scientific Collaboration & The Virgo Collaboration 2017) and an accompanying electromagnetic counterpart in form of a short gamma-ray burst (LIGO Scientific Collaboration et al. 2017) and kilonova afterglow (Kasen et al. 2017; Drout et al. 2017) has established a firm connection with electromagnetic counterparts and highlights the need for multiphysics modelling of neutron-star mergers. A neutron-star merger consists of several stages starting from the late inspiral, where accurate numerical waveforms are needed to calibrate analytical models for waveforms, through merger (Kawaguchi et al. 2018; Nagar et al. 2018; Dietrich et al. 2019), which requires sophisticated microphysics in terms of finite-temperature equations of state (EOS) satisfying recent observational constraints (Annala et al. 2018; Most et al. 2018; Tews, Margueron & Reddy 2018; Burgio et al. 2018; Raithel, \u00d6zel & Psaltis 2018) until the post-merger phase where neutrino and magnetic viscosity can drive large amounts of mass ejection (Just et al. 2015; Siegel & Metzger 2017, 2018; Fern\u00e1ndez et al. 2019; Fujibayashi et al. 2018), that are needed to make a connection to the kilonova afterglow produced by the decay of heavy elements in the matter outflow. The modelling of this complicated multiphysics system requires both the use of numerical relativity (Baiotti & Rezzolla 2017; Duez & Zlochower 2019) and of an accurate modelling of the fluid, the electromagnetic fields as well as the microphysics. Considerable effort has been placed on improved and highly accurate methods to solve Einstein field equations numerically (Baumgarte & Shapiro 2010; Shibata 2016) and to couple them with high-order methods for relativistic hydrodynamics (Radice, Rezzolla & Galeazzi 2014a; Bernuzzi & Dietrich 2016). At the same time, mainly driven by an effort to model core-collapse supernovae, very sophisticated numerical schemes for neutrino transport have been developed (Ruffert, Janka & Schaefer 1996; Buras et al. 2006; Shibata et al. 2011; Sumiyoshi & Yamada 2012; Foucart et al. 2015; Just et al. 2015). When considering the late stages of the evolution of the system not only is it important to account for the various relevant physics contributions, such as neutrino interactions, but it is also crucial to understand how numerical errors at finite-numerical resolution accumulate over time. This is even highlighted by the fact that current simulations of neutron-star mergers sometimes show non-convergent behaviour in the magnetic field evolution (Endrizzi et al. 2016; Ciolfi et al. 2017) even when small resolution changes are used. Notwithstanding, that with current computational efficiencies and available resources not even all relevant physical scales involving magnetic turbulence can be resolved (Kiuchi et al. 2015b, 2018), studying the late-time evolution of the remnants accretion disc is not only feasible but has been the subject of recent investigations (Siegel & Metzger 2017, 2018; Fern\u00e1ndez et al. 2019). While all such simulations so far have used traditional second-order accurate finite-volume schemes to model the evolution of the general-relativistic magnetohydrodynamics system (GRMHD), earlier works have already indicated the benefit of using more accurate high-order methods in this context (Del Zanna et al. 2007; Tchekhovskoy, McKinney & Narayan 2007; Radice et al. 2014a), while even more recent studies have already started to consider advanced finite-element approaches (Kidder et al. 2017; Fambri et al. 2018). Taking an intermediate approach similar to Del Zanna et al. (2007), McCorquodale & Colella (2011), Chen, T\u00f3th & Gombosi (2016), and Felker & Stone (2018), we will consider the impact of using a fourth-order accurate numerical scheme to model the merger of magnetized binary neutron stars and show the advantages gained when additionally finite-temperature effects and neutrino cooling are included.","Citation Text":["Burgio et al. 2018"],"Functions Text":["A neutron-star merger consists of several stages starting from the late inspiral, where accurate numerical waveforms are needed to calibrate analytical models for waveforms, through merger","which requires sophisticated microphysics in terms of finite-temperature equations of state (EOS) satisfying recent observational constraints"],"Functions Label":["Motivation","Motivation"],"Citation Start End":[[1177,1195]],"Functions Start End":[[711,899],[966,1107]]}
{"Identifier":"2022AandA...659A..21P__Khadka_et_al._2021_Instance_1","Paragraph":"Using Eq. (16), we derive the probability density in the parameter space (\u03c3DRW, RBLR) for a given set of parameters \u27e8M\u27e9, fBLR, ve, and R0. By marginalising on ve and R0 we obtain the measurement of RBLR shown in Fig. 10. One could argue that we obtain a bi-modal distribution in the posterior probability for RBLR. This observation can be explained by the fact that, as shown in Sect. 4, the R-band encapsulates the Mg\u202fII and Fe\u202fII emission lines which can arise from two distinct regions of the BLR. Indeed, the H\u03b2 (used in Mosquera & Kochanek 2011) and Mg\u202fII lines seem to arise from the same part of the BLR in various quasars (e.g., Karouzos et al. 2015; Khadka et al. 2021) and should both yield similar sizes; whereas the Fe\u202fII line is thought to arise from a larger part of the BLR (e.g., Sluse et al. 2007; Hu et al. 2015; Zhang et al. 2019; Li et al. 2021). Therefore, the combination of the two signals modelled as a single BLR emission could broaden our measurement and induce its slight bi-modality. Still, the core of the probability lies in the [0.1\u20131.5] RBLRMK11 range and the second mode observed for higher values of RBLR rises only for the highest values of \u03c3DRW. The marginalisation of this posterior over \u03c3DRW yields a probability distribution for RBLR and by taking its 16th, 50th, and 84th percentiles we measure \n\n\n\n\nR\nBLR\n\n=\n1\n.\n\n6\n\n\u2212\n0.8\n\n\n+\n1.5\n\n\n\u00d7\n\n10\n17\n\n\n\n$ R_{\\mathrm{BLR}} = 1.6^{+1.5}_{-0.8}\\times 10^{17} $\n\n\n cm. With a relative precision of \u224880% our method is less precise than recent spectroscopical reverberation mapping measurements (e.g., Grier et al. 2019; Penton et al. 2022 have around 30% relative precision for quasars with z\u2004>\u20041.3) but is more precise than photometric reverberation mapping (e.g., Kaspi et al. 2021 have above 100% relative precision when using a cross-correlation function with R and B filter light curves). The value of RBLRMK11 predicted by the luminosity\u2013size relation is in agreement with our measurement at the 1\u2212\u03c3 level.","Citation Text":["Khadka et al. 2021"],"Functions Text":["Indeed, the H\u03b2","and Mg\u202fII lines seem to arise from the same part of the BLR in various quasars (e.g.,","and should both yield similar sizes; whereas the Fe\u202fII line is thought to arise from a larger part of the BLR","Therefore, the combination of the two signals modelled as a single BLR emission could broaden our measurement and induce its slight bi-modality."],"Functions Label":["Uses","Uses","Uses","Uses"],"Citation Start End":[[659,677]],"Functions Start End":[[501,515],[551,636],[679,788],[867,1011]]}
{"Identifier":"2018ApJ...854..137S__Zank_&_Matthaeus_1993_Instance_1","Paragraph":"Modulation in steady-state has been well studied by previous works (Potgieter 2013; Potgieter et al. 2014; Zhao et al. 2014). Potgieter et al. (2014) studied the modulation of proton spectra with the PAMELA data from 2006 to 2009 July, and they concluded that the recent solar minimum was \u201cdiffusion dominated.\u201d In the work of Potgieter et al. (2014), parameters in diffusion coefficients and drift coefficients were adapted to observations (see also Potgieter et al. 2015; Raath et al. 2016). Zhao et al. (2014) studied the modulation of the GCR energy spectra during the past three solar minima using an empirical diffusion coefficient model according to Zhang (1999). They found that decreased perpendicular diffusion in polar direction, which is in contrast to the assumption of enhanced diffusion in polar regions that was used to explain the observed Ulysses CR gradients (see, e.g., Potgieter 2000), and increased parallel diffusion might be the reason for the record high-level of GCR intensity measured at Earth. Since the diffusion coefficients describe the scattering of GCRs by random fluctuations in the IMF, turbulence quantities are needed in diffusion theory. In the solar wind, the evaluation of turbulence is well described by magnetohydrodynamic (MHD) theory (Marsch & Tu 1989; Zhou & Matthaeus 1990), and the turbulence transport throughout the heliosphere has been studied over the years (e.g., Zank & Matthaeus 1993; Zank et al. 1996, 2012, 2017; Matthaeus et al. 1999; Smith et al. 2001; Breech et al. 2008; Hunana & Zank 2010; Oughton et al. 2011; Wiengarten et al. 2016). With the turbulence transport models (TTMs), the diffusion tensor can be calculated (e.g., Zank et al. 1998; Pei et al. 2010a; Engelbrecht & Burger 2013; Zhao et al. 2017). Theoretical and numerical works have shown that drift coefficients can be reduced in the presence of turbulence (e.g., Jokipii 1993; Fisk & Schwadron 1995; Giacalone & Jokipii 1999; Candia & Roulet 2004; Stawicki 2005; Minnie et al. 2007; Tautz & Shalchi 2012); see also the first-order approach in Engelbrecht et al. (2017). The reduced drift coefficients, which are obtained from fitting the simulation results (Burger & Visser 2010; Tautz & Shalchi 2012), are also used together with the turbulence transport theory to study the modulation of GCRs (e.g., Engelbrecht & Burger 2013). Using the nearly incompressible (NI) MHD TTM developed by Zank et al. (2017), Zhao et al. (2017) also showed the effect of both weak and moderately strong turbulence on drift coefficients. Since there exists close coupling between turbulence, solar wind, and energetic particles, some work combined large-scale solar wind flow with small-scale fluctuations in a self-consistent way (see, e.g., Usmanov et al. 2011, 2014, 2016; Wiengarten et al. 2015; Shiota et al. 2017) to study the spatial variations of the diffusion coefficients (see, e.g., Chhiber et al. 2017). Furthermore, using a diffusion coefficient model according to Giacalone & Jokipii (1999), Guo & Florinski (2016) studied the modulation of GCRs by CIRs at 1 au. They combined the small-scale turbulence transport with the MHD background for the simulation of cosmic-ray transport to show short-term modulation effects.","Citation Text":["Zank & Matthaeus 1993"],"Functions Text":["In the solar wind, the evaluation of turbulence is well described by magnetohydrodynamic (MHD) theory","and the turbulence transport throughout the heliosphere has been studied over the years (e.g.,"],"Functions Label":["Background","Background"],"Citation Start End":[[1416,1437]],"Functions Start End":[[1176,1277],[1321,1415]]}
{"Identifier":"2018AandA...614A...9J__in_2017_Instance_1","Paragraph":"During our monitoring period V2492 Cyg remained always undetected at our sensitivity, therefore the light curve depicted in Fig. 1 displays only upper limit values at different levels, as explained in Sect. 3. As a consequence, no fading or outbursting event can be detected, nevertheless some useful information can be derived. As mentioned above, after its discovery in 2010, V2492 underwent a long-lasting period of strong activity with intermittent burst and fading events (see Hillenbrand et al. 2013 and AAVSO2 data) and reached its maximum recorded brightness in 2017 Giannini et al. (2018). During most of this period, the source was sampled with an almost daily cadence and, for long time intervals, remained brighter than the following values: B  18, V  16, R  15, and I  14 mag. In comparison, our plate measurements are largely undersampled, presenting long periods (up to a decade) without any data. In any case (not considering the I band, which is practically uncovered, and the B band, which presents no significant upper limits), our V and R-band upper limitstentatively suggest that duringa period of about 30 yr from 1958 to 1987 an activity similar (both in duration and in brightness) to that more recently (2010\u20132017) monitored, did not occur. Indeed, we note that for a significant amount of time the source is brighter than the level indicated by our upper limits, thus suggesting that, in the past, the activity of V2492 Cyg was not as strong as it is now. The above scenario, if correctly described, means that an enhanced brightness variability could be an infrequent feature of V2492 Cyg. Such circumstances favour an accretion- more than an extinction-driven origin for the bursts. Indeed, the former is expected to occur with a long and irregular cadence related to the viscous motion of the matter toward the inner edge of the disk, while the latter should occur more frequently and regularly, according to the orbital motion of the obscuring matter along the line of sight.","Citation Text":["Giannini et al. (2018)"],"Functions Text":["As mentioned above, after its discovery in 2010, V2492 underwent a long-lasting period of strong activity with intermittent burst and fading events","and reached its maximum recorded brightness in 2017"],"Functions Label":["Background","Background"],"Citation Start End":[[575,597]],"Functions Start End":[[329,476],[523,574]]}
{"Identifier":"2018ApJ...861....2T__Poisson_1999_Instance_1","Paragraph":"However, one can reduce the order of the LD equation using the method proposed in Landau & Lifshitz (1975), i.e., by rewriting the self-force in terms of the external force and the four-velocity of a particle. Substituting the higher-order terms in Equation (3) with the derivatives of the Lorentz force, we get the equation in the following form:\n5\n\n\n\n\n\nThis equation, usually referred to as the Landau\u2013Lifshitz (LL) equation, has important consequences: it is of the second order, does not violate the principle of inertia, and the self-force vanishes in the absence of the external (Lorentz) force (Rohrlich 2001; Poisson 1999). The self-contained derivation of Equation (5) in terms of retarded potentials is given in Poisson et al. (2011). Equation (5) can be applied to cases with any external forces acting on a charged particle instead of the Lorentz force. In the case where \n\n\n\n\n\n, the radiation-reaction force can be rewritten in the form\n6\n\n\n\n\n\nwhere \n\n\n\n\n\n is the specific charge of the particle, \n\n\n\n\n\n, and the comma in the first term denotes the partial derivative with respect to the coordinate x\u03b1. Spohn (2000) concluded that using the LL equation is identical to imposing Dirac\u2019s asymptotic condition \n\n\n\n\n\n on the LD equation. It was later confirmed by Rohrlich (2001) that the reduced form of the equation of motion is exact, rather than approximative, though the LL equation was proposed in Landau & Lifshitz (1975) as an approximative solution to the third-order LD equation. More details on the treatment of the radiation reaction of charged particles in flat spacetime can be found in the book by Spohn (2004). In our numerical study, we found that the LL approximation is perfectly applicable if the Schott term is small with respect to the radiation recoil term, which is the case we consider here. Below we show a representative example of charged particle motion in an external uniform magnetic field, integrating both LD and LL equations. The results of the numerical studies of LD and LL equations for the motion of a charged particle in a uniform magnetic field in flat spacetime are in accord with the analytical treatment of the radiation-reaction force performed in Spohn (2000).","Citation Text":["Poisson 1999"],"Functions Text":["This equation, usually referred to as the Landau\u2013Lifshitz (LL) equation, has important consequences: it is of the second order, does not violate the principle of inertia, and the self-force vanishes in the absence of the external (Lorentz) force"],"Functions Label":["Background"],"Citation Start End":[[617,629]],"Functions Start End":[[355,600]]}
{"Identifier":"2020ApJ...904...11F__Zamirri_et_al._2019a_Instance_1","Paragraph":"BE values can also be obtained by means of computational approaches that, in some situations, can overcome the experimental limitations. Many computational works have so far focused on a few important astrochemical species like H, H2, N, O, CO, and CO2, in which BEs are calculated on periodic\/cluster models of crystalline\/amorphous structural states using different computational techniques (e.g., Al-Halabi & Van Dishoeck 2007; Karssemeijer et al. 2014; Karssemeijer & Cuppen 2014; \u00c1sgeirsson et al. 2017; Senevirathne et al. 2017; Shimonishi et al. 2018; Zamirri et al. 2019a). In addition, other works have computed BEs in a larger number of species but with a very approximate model of the substrate. For example, in a recent work by Wakelam et al. (2017) BE values of more than 100 species are calculated by approximating the ASW surface with a single water molecule. The authors then fitted the most reliable BE measurements (16 cases) against the corresponding computed ones, obtaining a good correlation between the two data sets. In this way, all the errors in the computational methods and limitations due to the adoption of a single water molecule are compensated by the fitting with the experimental values, in the view of the authors. The resulting parameters are then used to scale all the remaining computed BEs to improve their accuracy. This clever procedure does, however, consider the proposed scaling universal, leaving aside the complexity of the real ice surface and the specific features of the various adsorbates. In a similar work, Das et al. (2018) have calculated the BEs of 100 species by increasing the size of a water cluster from one to six molecules, noticing that the calculated BE approaches the experimental value when the cluster size is increased. As we will show in the present work, these approaches, relying on an arbitrary and very limited number of water molecules, cannot, however, mimic a surface of icy grain. Furthermore, the strength of interaction between icy water molecules, as well as with respect to the adsorbates, depends on the hydrogen bond cooperativity, which is underestimated in small water clusters.","Citation Text":["Zamirri et al. 2019a"],"Functions Text":["Many computational works have so far focused on a few important astrochemical species like H, H2, N, O, CO, and CO2, in which BEs are calculated on periodic\/cluster models of crystalline\/amorphous structural states using different computational techniques (e.g.,"],"Functions Label":["Background"],"Citation Start End":[[559,579]],"Functions Start End":[[137,399]]}
{"Identifier":"2019ApJ...876...85R__Shappee_et_al._2014_Instance_1","Paragraph":"Next, we applied a correction for the expected difference between the magnitude of each Cepheid at the observed phase and the magnitude at the epoch of mean intensity of its light curve. These phase corrections are derived from ground-based light curves of each Cepheid in filters with wavelengths best corresponding to the WFC3 filters. Because the phase corrections are relative quantities, they do not change the zero-point of the light curves, which remain on the HST WFC3 natural system.6\n\n6\nIn practice, the ground-based light curves are transformed to the HST system using color terms. While an uncertainty in color term could produce systematic errors, these are negligible. We determined empirically that a 10% error in the color terms changes the mean phase correction by \u22640.1 mmag.\n We derived and applied these phase corrections following the same methodology described in R18b. The periods and phases for F555W and F814W were determined using the V- and I-band light curves from OGLE surveys (Szymanski 2005; Udalski et al. 2008, 2015). For some Cepheids (OGL0434, OGL0501, OGL0510, OGL0512, OGL0528, OGL0545, OGL0590, OGL0712, OGL0757, OGL0966, and OGL0992), we also included V-band light curves from the ASAS survey (Pojmanski 1997) and\/or ASAS-SN survey (Shappee et al. 2014; Kochanek et al. 2017) to increase the baseline coverage. We made use of the J- and H-band light curves from M15 and Persson et al. (2004) to correct the F160W random phased measurements to mean intensity. The standard deviations of these corrections are 0.29, 0.17, and 0.11 mag in F555W, F814W, and F160W, respectively, decreasing with the smaller light-curve amplitudes at redder wavelengths. Phase corrections also account for the difference between the Cepheid light-curve magnitude mean (the average of many measured magnitudes) and the magnitude at the epoch of mean intensity (the standard convention for distance measurements). This expected difference is consistent with our sample average correction of \u22120.048, \u22120.013, and \u22120.001 mag, in F555W, F814W, and F160W, respectively. The uncertainties in these phase corrections depend on the quality of the ground-based light curves; the average uncertainty is 0.013, 0.008, and 0.029 mag per epoch in F555W, F814W, and F160W, respectively, which dominates over the statistical photometry errors (i.e., photon statistics) in a single epoch. The differences between repeat measurements for the same target, available for a subset of 19 epochs and filters, is consistent with these uncertainties. The final mean individual uncertainty for these 70 Cepheids is 0.016, 0.012, and 0.029 mag in F555W, F814W, and F160W, respectively. The final mean photometry for each Cepheid in three colors is given in Table 2.","Citation Text":["Shappee et al. 2014"],"Functions Text":["For some Cepheids (OGL0434, OGL0501, OGL0510, OGL0512, OGL0528, OGL0545, OGL0590, OGL0712, OGL0757, OGL0966, and OGL0992), we also included V-band light curves from the","and\/or ASAS-SN survey","to increase the baseline coverage."],"Functions Label":["Uses","Uses","Uses"],"Citation Start End":[[1271,1290]],"Functions Start End":[[1050,1218],[1248,1269],[1314,1348]]}
{"Identifier":"2021AandA...654A.141L__Li_2019_Instance_1","Paragraph":"The situation seems to be more complicated in luminous AGNs, which can be divided into radio-quiet (RQ) and radio-loud (RL) AGNs according to their radio loudness R (R = f5\u2004GHz\/f4400\u2004\u00c5, where f5\u2006GHz and f4400\u2004\u00c5 are the radio flux at 5 GHz and optical flux at 4400 \u00c5, respectively, Kellermann et al. 1989). The accretion process can be well described by a disk-corona model in RQAGNs (Jin et al. 2012; Lusso & Risaliti 2017; Qiao & Liu 2018), where the optical-UV flux is emitted from an optically thick, geometrically thin accretion disk and the hard X-ray flux comes from the inverse Compton scattering of optical-UV photons by the hot corona. This model can naturally fit the observational positive relationships between \u03b1OX and \u03bbO as well as \u0393 and \u03bbO in RQAGNs (Shemmer et al. 2006; Risaliti et al. 2009; Lusso et al. 2010; Brandt & Alexander 2015; Li 2019), where the optical to X-ray spectral index \u03b1OX is defined as \u03b1OX = 0.384 log[L\u03bd(2500\u2004\u00c5)\/L\u03bd(2\u2004keV) (e.g., Lusso et al. 2010), and \u03bbO = 8.1\u03bdL\u03bd(5100\u2004\u00c5)\/LEdd is the Eddington ratio based on the optical luminosity at 5100 \u00c5 (Runnoe et al. 2012a), with LEdd being the Eddington luminosity. In contrast, the origin of X-ray emission is still unclear in RLAGNs. In observations, RLAGNs appear to have different properties than RQAGNs, for instance the positive relationships listed above in RQAGNs have not been found in RLAGNs (Li 2019; Zhou & Gu 2020). Furthermore, the average X-ray luminosity in RLAGNs is found to be 2\u20133 times higher than that in RQAGNs (Zamorani et al. 1981; Wilkes & Elvis 1987; Wu et al. 2013; Gupta et al. 2018). This seems to indicate that the X-ray emission from jets is important in RLAGNs. However, Gupta et al. (2018) recently compiled an excellent sample to investigate the differences of X-ray properties between luminous radio galaxies and their radio-quiet counterparts, where the black hole mass and Eddington ratio were well selected and the bolometric luminosity were calculated from mid-infrared emission observed by the Wide-field Infrared Survey Explorer (WISE) mission. They argue that the X-ray emission in radio-loud radio galaxies should also come from a disk-corona system because their distribution of the X-ray slope is very similar with those of radio-quiet counterparts (see Gupta et al. 2018 for details, see also Gupta et al. 2020).","Citation Text":["Li 2019","Li 2019"],"Functions Text":["This model can naturally fit the observational positive relationships between \u03b1OX and \u03bbO as well as \u0393 and \u03bbO in RQAGNs","In observations, RLAGNs appear to have different properties than RQAGNs, for instance the positive relationships listed above in RQAGNs have not been found in RLAGNs"],"Functions Label":["Compare\/Contrast","Compare\/Contrast"],"Citation Start End":[[852,859],[1382,1389]],"Functions Start End":[[645,763],[1215,1380]]}
{"Identifier":"2018ApJ...857...79J__Fadda_et_al._1996_Instance_1","Paragraph":"Finally, in order to measure the ICL fraction once we have a background-free ICL map, we created an image of the cluster removing the CHEF models of the foreground and background galaxies. As described in Jim\u00e9nez-Teja & Ben\u00edtez (2012), the cluster membership is determined in a two-step process, the PEAK+GAP algorithm (Owers et al. 2011), using the spectroscopic data available for each system. This composite method first identifies the peak of the cluster in the redshift space and selects a redshift window wide enough to contain the whole distribution of velocities assigned to that peak. Implicitly, the size of this window is proportional to the velocity dispersion of the clusters: merging clusters, with a more scattered velocity distribution, will need a wider window, compared to relaxed systems. This crude selection of cluster member candidates is obviously prone to contamination by interlopers. So, we further refine it using the shifting gapper method (Fadda et al. 1996; Girardi et al. 1996; Boschin et al. 2006; Owers et al. 2011), which uses velocity and spatial information on the candidates simultaneously. The shifting gapper method spatially distributes the candidates according to their clustercentric distance in radial bins. The mean velocity of the candidates within each bin is calculated, and those candidates with velocities that are too far from the others are rejected. As unrelaxed clusters are more likely to have a broader spatial distribution, this procedure naturally allows candidates at larger distances to be identified as cluster members for these systems. These two steps are, thus, essential to guarantee that our cluster membership algorithm implicitly takes into account the dynamical state of the systems and does not bias the measurement of their total luminosity, while minimizing contamination by interlopers at the same time. We refer the reader to Jim\u00e9nez-Teja & Ben\u00edtez (2012) for further information on the cluster membership selection algorithm.","Citation Text":["Fadda et al. 1996"],"Functions Text":["This crude selection of cluster member candidates is obviously prone to contamination by interlopers. So, we further refine it using the shifting gapper method","which uses velocity and spatial information on the candidates simultaneously. The shifting gapper method spatially distributes the candidates according to their clustercentric distance in radial bins. The mean velocity of the candidates within each bin is calculated, and those candidates with velocities that are too far from the others are rejected. As unrelaxed clusters are more likely to have a broader spatial distribution, this procedure naturally allows candidates at larger distances to be identified as cluster members for these systems. These two steps are, thus, essential to guarantee that our cluster membership algorithm implicitly takes into account the dynamical state of the systems and does not bias the measurement of their total luminosity, while minimizing contamination by interlopers at the same time."],"Functions Label":["Uses","Background"],"Citation Start End":[[969,986]],"Functions Start End":[[808,967],[1050,1875]]}
{"Identifier":"2017ApJ...850..195E__Lacy_et_al._1982_Instance_1","Paragraph":"RCW 57A (also known as NGC 3576, G291.27\u20130.70, or IRAS 11097\u20136102) is a H ii region associated with a filament and bipolar bubble, and is located at a distance of 2.4\u20132.8 kpc (Persi et al. 1994; de Pree et al. 1999). We adopt 2.4 kpc, which is within uncertainties of both the kinematic and spectroscopic determinations (see Persi et al. 1994). Figure 1 depicts the overall morphology of RCW 57A. It contains optically bright nebulosity with several dark globules and luminous arcs (Persi et al. 1994). It is one of the massive star-forming regions in the southern sky, hosting a H ii region (cyan contour) embedded in a filament (white contour) from which a widely extended bipolar bubble (yellow contours) is emerging. A deeply embedded near-IR cluster, consisting of more than 130 young stellar objects (YSOs), is associated with this region (Persi et al. 1994). The observed ratios of the infrared fine-structure ionic lines (Ne ii, Ar iii, and S iv; Lacy et al. 1982) indicate that at least eight O7.5V stars are necessary to account for the ionization of the region. However, even these stars may not be sufficient to account for the Ly\u03b1 ionizing photons inferred from radio data (Figuer\u00eado et al. 2002; Barbosa et al. 2003; Townsley 2009). Based on the newly discovered cluster of stars, using X-ray data, Townsley et al. (2014) suggested that an additional cluster of OB stars that were not known before might be deeply embedded. This cluster is located slightly southwest of the center of the near-IR cluster. The 10 \u03bcm map (cf., Frogel & Persson 1974) reveals the presence of five infrared sources (IRS; black squares in Figure 1) near the center of the H ii region. These, together with water and methanol maser sources (green crosses in Figure 1) distributed along the filament, are indicative of active ongoing star formation in RCW 57A (Frogel & Persson 1974; Caswell 2004; Purcell et al. 2009). Therefore, RCW 57A is an ideal target to investigate the morphological links among filaments, bipolar bubbles, and B-fields so as to understand the star formation history.","Citation Text":["Lacy et al. 1982"],"Functions Text":["The observed ratios of the infrared fine-structure ionic lines (Ne ii, Ar iii, and S iv;","indicate that at least eight O7.5V stars are necessary to account for the ionization of the region."],"Functions Label":["Uses","Uses"],"Citation Start End":[[955,971]],"Functions Start End":[[866,954],[973,1072]]}
{"Identifier":"2020ApJ...889...42B__Ricker_et_al._2014_Instance_1","Paragraph":"The characterization of the interior of observed exoplanets is one of the main goals in current exoplanetary science. With the large number of newly discovered exoplanets expected in the next 10 years by ground-based surveys such as the Wide Angle Search for Planets (WASP; Pollacco et al. 2006), the Next-Generation Transit Survey (NGTS; Wheatley et al. 2017), and the Hungarian Automated Telescope Network\/Hungarian Automated Telescope Network-South (HATNet\/HATSouth; Hartman et al. 2004; Bakos et al. 2013), as well as the ongoing Transiting Exoplanet Survey Satellite (TESS) space survey (Ricker et al. 2014) and the upcoming PLAnetary Transits and Oscillations (PLATO) mission (Rauer et al. 2014), a rapid characterization scheme of the interior structure of these planets will become increasingly necessary to further our understanding of planetary populations. The vast majority of the confirmed exoplanets has been identified either through transits or radial velocities surveys. Planets identified with both techniques are characterized by their mass and radius, which, combined, provide a first indication of the bulk composition through comparison with theoretical mass\u2013radius curves (e.g., Valencia et al. 2006; Sotin et al. 2007; Zeng & Sasselov 2013). A common approach to the interior characterization of exoplanets is the use of numerical models to compute interior structures that comply with the measured mass and radius of the planet (e.g., Fortney et al. 2007; Sotin et al. 2007; Valencia et al. 2007; Wagner et al. 2011; Zeng & Sasselov 2013; Unterborn & Panero 2019). As this is an inverse problem, it requires the calculation of a large number of interior models to obtain an overview over possible interior structures (Rogers & Seager 2010a; Brugger et al. 2017; Dorn et al. 2017). If other observables are used in addition to mass and radius, the number of samples needed for an accurate inference of possible interior structures increases drastically, due to the increase in dimensionality (e.g., James et al. 2013). Thus, the inference can quickly become computationally expensive. Moreover, with only mass and radius, possible solutions tend to be highly degenerate, with multiple, qualitatively different interior compositions that can match the observations equally well (e.g., Rogers & Seager 2010a, 2010b).","Citation Text":["Ricker et al. 2014"],"Functions Text":["With the large number of newly discovered exoplanets expected in the next 10 years by ground-based surveys such as","as well as the ongoing Transiting Exoplanet Survey Satellite (TESS) space survey","a rapid characterization scheme of the interior structure of these planets will become increasingly necessary to further our understanding of planetary populations."],"Functions Label":["Background","Background","Motivation"],"Citation Start End":[[593,611]],"Functions Start End":[[118,232],[511,591],[703,867]]}
{"Identifier":"2020MNRAS.497..687S__Dere_et_al._1997_Instance_1","Paragraph":"The object J004415.00 required our special consideration. We notice large variations of spectra at the hydrogen series limits. We designate the limits of the Balmer, Paschen, Brackett and Pfund series in Fig. 4 for J004415.00 with vertical arrows. A jump at the Brackett series limit in the IR spectrum can be clearly seen. The Balmer and Paschen jumps can also be noticed by photometric brightness changes. These inverse jumps at the hydrogen series limits attest to the noticeable contribution of free\u2013free (f\u2013f) and free\u2013bound (f\u2013b) emissions to the spectrum. This confirms the presence of an ionized circumstellar envelope where the f\u2013f and f\u2013b emissions originate, which is typical for B[e]SGs (Zickgraf et al. 1985, 1986, 1989). To take the contribution of f\u2013f and f\u2013b radiations into account and estimate the corresponding spectra, we used the\u2009chianti package (Dere et al. 1997; Landi et al. 2013). We consider the case of isothermal pure hydrogen plasma at a temperature of Te = 10\u2009000\u2009K (Lamers et al. 1998). From the spectrum of the star we assume that its effective temperature is Tstar =15\u2009000\u2009K and match the model spectrum to the observations with the following parameters: Te = 10\u2009000\u2009K, EM = 1.37 \u00d7 1039\u2009cm\u22125 (emission measure), Tstar = 15\u2009000\u2009K, Rstar = 43R\u2299, AV = 1.1. In order to describe the IR excess of the SED we also add hot dust emission to the model spectrum by assuming it to be effective blackbody radiation with a temperature of about 1000\u2009K (the best fit obtained with Tdust = 1050\u2009K), which is typical for B[e]SG hot circumstellar dust (Lamers et al. 1998). The result of the SED fitting is shown in Fig. 4: the dashed line shows the blackbody radiation, the dotted line designates the contribution of free\u2013free and free\u2013bound emissions, the dash-dotted line demonstrates the dust emission contribution, the solid line indicates the total model spectrum. All spectra shown include the interstellar extinction (Fitzpatrick 1999, assuming RV = 3.1) with AV = 1.1 obtained as the best-fitting parameter. Since the object J004415.00 does not show considerable photometric variability, we incorporate some more published photometric data points for fitting (see Table 3) in addition to the BTA data. For better illustration of the matching between the model and observed spectra, the synthetic photometry of the total model continuum multiplied by the normalized spectrum is shown with filled circles. Note that since the parameters are obtained in the isothermal plasma approximation with fixed temperatures of the star and circumstellar plasma, we report the other parameters derived from them as estimates and do not assess their uncertainties. Nevertheless, our estimates of the model parameters are good enough and allow us to constrain the dust extinction and luminosity of the object.","Citation Text":["Dere et al. 1997"],"Functions Text":["To take the contribution of f\u2013f and f\u2013b radiations into account and estimate the corresponding spectra, we used the\u2009chianti package"],"Functions Label":["Uses"],"Citation Start End":[[868,884]],"Functions Start End":[[735,866]]}
{"Identifier":"2022AandA...663A.105P__Weeren_et_al._2009b_Instance_1","Paragraph":"Cluster radio relics are usually found in the outskirts of merging galaxy clusters. They exhibit elongated morphologies and high degrees of polarisation above 1 GHz (up to 70%, Ensslin et al. 1998; Bonafede et al. 2014; Loi et al. 2019; de Gasperin et al. 2022). The resolved spectral index in radio relics shows a gradient: it steepens towards the cluster centre and flattens towards the outskirts. Their size can reach up to \u223c2 Mpc, and high-resolution observations have revealed filamentary structures within relics themselves (Di Gennaro et al. 2018; Rajpurohit et al. 2020, 2022a,b; de Gasperin et al. 2022). The Largest Linear Sizes (LLS) and radio powers of relics are correlated, as well as the integrated spectral index and the radio power (van Weeren et al. 2009b; Bonafede et al. 2012; de Gasperin et al. 2014). Relics trace ICM shock waves with relatively low (M\u2004 \u20043) Mach numbers (Finoguenov et al. 2010; Akamatsu et al. 2013; Shimwell et al. 2015; Botteon et al. 2016). The acceleration of electrons is believed to proceed via diffusive shock acceleration (DSA) in the ICM (Ensslin et al. 1998; Roettiger et al. 1999), in which particles scatter back and forth across the shock front gaining energy at every crossing. Nevertheless, this mechanism has been shown to be rather inefficient in accelerating electrons from the thermal pool (Vazza & Br\u00fcggen 2014; Vazza et al. 2016; Botteon et al. 2020a; Br\u00fcggen & Vazza 2020; see Brunetti & Jones 2014 for a review). Recently, it has been suggested that seed electrons could originate from the tails and lobes (driven by AGN outflows) of cluster radio galaxies (Bonafede et al. 2014; van Weeren et al. 2017; Stuardi et al. 2019), which alleviates the requirements of high acceleration efficiencies at cluster shocks (e.g., Markevitch et al. 2005; Kang et al. 2012, 2017; Botteon et al. 2016; Eckert et al. 2016). In some cases, double relics have been detected on opposite sides of the cluster centre (e.g., Rottgering et al. 1997; van Weeren et al. 2010, 2012b; Bonafede et al. 2012; de Gasperin et al. 2015a). In these clusters it is possible to constrain the merger history, providing important information about the formation processes of relics.","Citation Text":["van Weeren et al. 2009b"],"Functions Text":["The Largest Linear Sizes (LLS) and radio powers of relics are correlated, as well as the integrated spectral index and the radio power"],"Functions Label":["Background"],"Citation Start End":[[750,773]],"Functions Start End":[[614,748]]}
{"Identifier":"2019AandA...627A.172R__Rozitis_&_Green_(2013)_Instance_2","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"],"Functions Text":["Since the light curve inversion produced convex shape models only, then shadowing and self-heating effects inside global-scale concavities (see",") were not possible to model."],"Functions Label":["Differences","Differences"],"Citation Start End":[[872,892]],"Functions Start End":[[728,871],[892,921]]}
{"Identifier":"2016ApJ...832..128C___1983a_Instance_1","Paragraph":"The understanding of proton acceleration at the Sun in large solar energetic particle (SEP) events has oscillated between flare and shock pictures (Cliver 2009b; Reames 2015). The earliest picture following the discovery of ground-level events (GLEs; major SEP events requiring >500 MeV protons) by Forbush (1946) was that protons were accelerated in flares, the clear choice in the absence of other observations. Subsequently, Wild et al. (1963) conjectured, mainly on the basis of radio observations, that large SEP events required coronal shock waves as manifested by type II solar radio bursts. Smaller electron-dominated SEP events were linked to metric type III bursts. Early observational support for this view on the SEP side was provided by Lin (1970). Through the work of \u0160vestka & Fritzov\u00e1-\u0160vestkov\u00e1 (1974), Kahler et al. (1978, 1984), Cliver et al. (1982, 1983a, 1983b), Cane & Stone (1984), Klecker et al. (1984), Mason et al. (1984, 1986), Meyer (1985), Reames et al. (1985, 1994, 1996), Cane et al. (1986, 1988), Luhn et al. (1987), Reames (1990, 1999), Kahler (1992, 1994), Gosling (1993), and others, involving various comparisons of SEP events with flare electromagnetic emissions and coronal mass ejections (CMEs), as well as considerations of SEP composition, charge states, and the longitude distribution of SEP-associated flares, the two-class picture of SEP acceleration presciently proposed by Wild et al. (1963) became established. The new consensus view was almost immediately challenged by observations of the first large (\u201cgradual\u201d; Reames 1993) proton events observed by the Advanced Composition Explorer (ACE). Mazur et al. (1999), Cohen et al. (1999), Mason et al. (1999a), and Mason et al. (1999b) reported that large SEP events, including GLEs, recorded by ACE and SAMPEX in 1997 and 1998 had elemental composition and charge states at >10 MeV\/nuc that were similar to those found in small (\u201cimpulsive\u201d) SEP events (e.g., Mason et al. 1986; Luhn et al. 1987) at lower energies. Subsequently, Cane et al. (2002, 2003, 2006) presented evidence based on low-frequency radio observations, SEP composition data, and flare location to argue for the presence of a flare-accelerated high-energy (>25 MeV) proton component in large SEP events to augment that produced by coronal\/interplanetary shock waves driven by CMEs. The relative importance of flare and shock components was left as an open question.","Citation Text":["Cliver et al.","1983a"],"Functions Text":["Through the work of \u0160vestka & Fritzov\u00e1-\u0160vestkov\u00e1 (1974), Kahler et al. (1978, 1984)","(1982","1983b), Cane & Stone (1984), Klecker et al. (1984), Mason et al. (1984, 1986), Meyer (1985), Reames et al. (1985, 1994, 1996), Cane et al. (1986, 1988), Luhn et al. (1987), Reames (1990, 1999), Kahler (1992, 1994), Gosling (1993), and others, involving various comparisons of SEP events with flare electromagnetic emissions and coronal mass ejections (CMEs), as well as considerations of SEP composition, charge states, and the longitude distribution of SEP-associated flares, the two-class picture of SEP acceleration presciently proposed by Wild et al. (1963) became established."],"Functions Label":["Background","Background","Background"],"Citation Start End":[[847,860],[868,873]],"Functions Start End":[[762,845],[861,866],[875,1456]]}
{"Identifier":"2018ApJ...865...60V__Baym_et_al._1969_Instance_1","Paragraph":"The core of a neutron star is composed primarily of neutrons, with \u223c5%\u201310% of the mass in protons; for the electrically neutral medium, the number density of electrons is equal to that of the protons. At the supranuclear densities of the outer core, the Fermi energy for protons and neutrons is well above the typical temperature of a mature neutron star, and both the neutrons and protons are expected to condense into Bardeen\u2013Cooper\u2013Schrieffer superfluids, with 3PF2 and 1S0 Cooper pairing, respectively (Migdal 1959; Baym et al. 1969). To support rotation, the neutron superfluid forms an array of quantized vortices, filaments of microscopic cross section, each carrying one quantum of circulation. The superconductivity of the protons is predicted to be type II, and the magnetic field is supported by an array of quantized flux tubes, each carrying one quantum of magnetic flux. Fermi-liquid interactions between the two condensates results in a nondissipative coupling between the mass currents of the two species (Andreev & Bashkin 1975; Chamel & Haensel 2006), so the neutron vortices are magnetized by entrained proton currents (Alpar et al. 1984a). Electron scattering from magnetized vortices and flux tubes produces dissipative and nondissipative forces on the vortices and flux tubes. The magnetic interaction at junctions between magnetized neutron vortices and flux tubes is energetic enough to produce pinning, wherein the neutron vortices pin to the dense array of flux tubes in the outer core (Srinivasan et al. 1990; Jones 1991; Chau et al. 1992; Ruderman et al. 1998; Link 2012b), similar to the predicted pinning of the vortices to the nuclear lattice of the crust (Anderson & Itoh 1975; Alpar 1977; Epstein & Baym 1988; Donati & Pizzochero 2006; Avogadro et al. 2007; Link 2009). Thermal fluctuations stochastically excite vortex motion, causing the neutron vortices to slip with respect to the flux tubes (Ding et al. 1993; Sidery & Alpar 2009; Link 2014).","Citation Text":["Baym et al. 1969"],"Functions Text":["At the supranuclear densities of the outer core, the Fermi energy for protons and neutrons is well above the typical temperature of a mature neutron star, and both the neutrons and protons are expected to condense into Bardeen\u2013Cooper\u2013Schrieffer superfluids, with 3PF2 and 1S0 Cooper pairing, respectively"],"Functions Label":["Uses"],"Citation Start End":[[520,536]],"Functions Start End":[[201,505]]}
{"Identifier":"2022ApJ...937L..34K__Narayan_et_al._2012_Instance_1","Paragraph":"We consider a radio galaxy of SMBH mass M = 109\nM\n9\nM\n\u2299 with a mass accretion rate of \n\n\n\nM\u0307=m\u0307LEdd\/c2\u22431.4\u00d71022M9m\u0307\u22124gcm\u22122\n\n, where c is the speed of light and L\nEdd is the Eddington luminosity. The gravitational radius of the BH is r\n\ng\n = GM\/c\n2 \u2243 1.5 \u00d7 1014\nM\n9 cm. We consider that the accretion flow is in the MAD state, and then the magnetic field strength around the SMBH is estimated to be \n\n\n\nBmad=M\u0307c\u03a6mad2\/(4\u03c02rg2)\u22431.1\u00d7103M9\u22121\/2m\u0307\u221241\/2\u03a6mad,1.7G\n\n (e.g., Yuan & Narayan 2014), where \u03a6mad \u2248 50\u03a6mad,1.7 is the saturated magnetic flux (Tchekhovskoy et al. 2011; Narayan et al. 2012; McKinney et al. 2012; White et al. 2019). The high-resolution GRMHD simulation with a BH spin parameter a = 0.9375 suggests that magnetic reconnection occurs at a distance of r\nrec \u223c 2r\n\ng\n (Ripperda et al. 2022). The value of r\nrec could depend on a or other parameters, but we fix r\nrec = 2r\n\ng\n throughout this paper for simplicity. We estimate the reconnecting magnetic field strength to be (see Appendix A)\n1\n\n\n\nBrec\u22482Bmadrrecrg\u22122\u22433.9\u00d7102M9\u22121\/2m\u0307\u221241\/2\u03a6rec,1.2G,\n\nwhere \n\n\n\n\u03a6rec=2\u03a6mad(rrec\/rg)\u22122\n\n is the effective magnetic flux at the reconnection region. The magnetosphere will be formed around the SMBH. The minimum number density of the magnetosphere that can maintain the electric current for the BZ process is (Goldreich & Julian 1969; Levinson & Cerutti 2018)\n2\n\n\n\nnGJ=Brec\u03a9F2\u03c0ec\u2248Brec8\u03c0erg\u22432.2\u00d710\u22124M9\u22123\/2m\u0307\u221241\/2\u03a6rec,1.2cm\u22123,\n\nwhere \u03a9\nF\n \u2248 ac\/(4r\n\ng\n) is the field line angular velocity (Tchekhovskoy et al. 2010; Nathanail & Contopoulos 2014; Ogihara et al. 2021; Camilloni et al. 2022) and we assume the BH spin parameter as a \u223c 1. For the magnetosphere, which consists of e\n+\ne\n\u2212 pair plasma with the density n\nGJ, the magnetization parameter is\n3\n\n\n\n\u03c3B,GJ=Brec24\u03c0nGJmec2\u22486.8\u00d71013M91\/2m\u0307\u221241\/2\u03a6rec,1.2.\n\nThis value should be regarded as an upper limit, because the number density of the magnetosphere can be higher than n\nGJ. Various mechanisms of particle injection into the BH magnetosphere have been proposed (see Appendix B), which can lead to multiplicity of \u03ba\n\u00b1 \u2261 n\/n\nGJ \u223c 1 \u2212 103. This results in the magnetization parameter of \u03c3\n\nB\n \u2273 1010.","Citation Text":["Narayan et al. 2012"],"Functions Text":["where \u03a6mad \u2248 50\u03a6mad,1.7 is the saturated magnetic flux"],"Functions Label":["Uses"],"Citation Start End":[[568,587]],"Functions Start End":[[486,540]]}