adsabs/FOCAL huggingface dataset:
At high redshift, the mapping of the extended Lyα 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 Lyα 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 Lyα haloes around normal star-forming galaxies emitting Lyα emission using narrowband (NB) imaging methods were reported by Møller & Warren (1998) and Fynbo et al. (2001). Later, Hayashino et al. (2004) observed 22 Lyman break galaxies (LBG) and detected extended Lyα emission by stacking the NB images. These authors were followed six years later by Ono et al. (2010) who detected Lyα haloes in their composite NB images of 401 Lyα 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 ≈2000 and ≈4500 LAEs at redshift z ≃ 3 and 2.2 ≤ z ≤ 6.6, respectively. Momose et al. (2014) found typical Lyα halo exponential scale lengths of 5–10 physical kpc. Matsuda et al. (2012) found that Lyα 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 ≈1500 galaxies in two overdense regions at z ≈ 3 and 4. Using stacking methods these authors reported Lyα halo exponential scale lengths of 5–6 physical kpc and found that Lyα halo sizes correlate with the UV continuum and Lyα luminosities, but not with overdensity. Steidel et al. (2011) stacked 92 brighter (RAB ≃ 24.5) and more massive LBGs at z = 2.3−3, finding large Lyα extents of ≈80 physical kpc beyond the mean UV continuum size at a surface brightness level of ~10-19 erg s-1 cm-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: Momose et al. (2014); Momose et al. (2014)
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In equations (2) and (3), $\boldsymbol{f}_{\rm rad}=(f_{{\rm rad},\, r},\, f_{{\rm rad},\theta })$ is the radiation force described as
(6)$$\begin{eqnarray*}
\boldsymbol{f}_{\rm rad}=\frac{\sigma _{\rm e} \boldsymbol{F}_{\rm D}}{c}+\frac{\sigma _{\rm e} \boldsymbol{F}_{\rm line}}{c}M,
\end{eqnarray*}$$where σe 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 200–3200 Å. 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 (200–3200  Å), 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 (200–3200  Å) 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 τe 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 θ-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 φ-directions, we here prepare 12 800 grids whose sizes are determined by Δri/Δri − 1 = 1.0005 in the r-direction. In the φ-direction, we set 1600 uniform grids in the range of 0 ≤ φ  2π. 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–$10^6 \, M_\odot$), 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 ∼1000RS for $M_{\rm BH}=10^8\, M_\odot$ while it reaches as far as 3 × 104RS for $M_{\rm BH}=10^3\, M_\odot$.
Citation: Castor, Abbott & Klein 1975
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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 (





), is assigned to Fe to make FeO until either O or Fe is exhausted. The excess oxygen available to make FeO is obtained using
4





where 





 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, 





, 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 ΔIW expected for the body taken as a whole, including accreted rock and ices.
Citation: Raddi et al. 2015
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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; Öberg 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; Öberg 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; Lefèvre et al. 2014). The growth can also be inferred from the evolution of the silicate-to-K band ratio (τ9,7/AK; e.g. Madden et al. 2022; van Breemen et al. 2011; Chiar et al. 2007).
Citation: Dartois 2005
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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 “critical” 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; Fürst 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: Staubert et al. 2019
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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
(31)$$\begin{eqnarray}
\partial _x f (u) \simeq \frac{F_{i - 1 / 2} - F_{i + 1 / 2}}{\Delta x},
\end{eqnarray}$$where Fi − 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 $F^{\operatorname{LO}}$:
(32)$$\begin{eqnarray}
F_{i + 1 / 2} = F_{i + 1 / 2}^{\operatorname{HO}} - A_{i + 1 / 2} \varphi _{i + 1 / 2} \left(F_{i + 1 / 2}^{\operatorname{HO}} - F_{i + 1 / 2}^{\operatorname{LO}}\right) .
\end{eqnarray}$$The term φi + 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 + 1/2φi + 1/2 = 0 the second-order flux is used, while if Ai + 1/2φi + 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–Friedrichs flux is used for $F^{\operatorname{LO}}$:
(33)$$\begin{eqnarray}
F^{\operatorname{HO}}_{i + 1 / 2} = \frac{f (u_i) + f (u_{i + 1})}{2},
\end{eqnarray}$$(34)$$\begin{eqnarray}
F^{\operatorname{LO}}_{i + 1 / 2} = \frac{1}{2} [f (u_i) + f (u_{i + 1})] - \frac{c_{i + 1 / 2}}{2} [u_{i + 1} - u_i] .
\end{eqnarray}$$The characteristic speed in the Lax–Friedrichs flux ci is taken to be the maximum value of the speed of light between the right and left cells
(35)$$\begin{eqnarray}
c_{i + 1 / 2} = \max _{a \in \lbrace i, i + 1 \rbrace } \left\lbrace \left| \alpha _a \sqrt{\gamma _a^{x x}} \pm \beta _a^x \right| \right\rbrace .
\end{eqnarray}$$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–Friedrichs 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: LeVeque 1992
Citation function: Uses
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’s 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 ∼0.01–0.05 M⊙, and redbacks (RB) with companion star masses ∼0.1–1 M⊙. A large proportion of the spiders, whether BWs or RBs, have been observed to exhibit (quasi-)periodic eclipses of the pulsars’ radio emission (e.g. Fruchter, Stinebring & Taylor 1988; Lyne et al. 1990) that are generally attributed to excess material in the orbits – 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) – 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 – largely as a result of difficult observing requirements and (until recently) a low number of known spider pulsars – 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: Lyne et al. 1990
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VHE gamma-rays from distant blazars suffer serious EBL absorption. For instance, the attenuation factor of the flux at ∼1 TeV emitted at z = 0.6 due to EBL absorption is ∼10−4 (Domínguez 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. Böttcher, Reimer & Marscher 2009; Yan & Zhang 2015), in which VHE emission is attributed to synchrotron emissions of relativistic protons and pair cascades created in proton–photon (pγ) interaction. However, the jet model cannot explain the VHE emission from PKS 1424+240 if its redshift z > 0.7–0.8 (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–Heitler (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: Essey et al. 2010
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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 II region in the proximity to a strong UV field arising from massive star formation. PDRs and H II 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 II regions, and IBLs has improved considerably since the availability of far-infrared spectroscopic data from the Herschel Space Observatory (see Ossenkopf et al. 2013; Köhler 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; Pérez-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 II 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–PDR 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 II] 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 II regions as indicated by weak [C II] emission. However, because [C II] samples both weakly and highly ionized regions, there remains some uncertainty about the relative contributions of the ionized and PDR regions. Furthermore, because [C II] 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 II] by assuming a thermal pressure and its Galactic gradient.
Citation: Joblin et al. 2018
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