Identifier
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82
| Paragraph
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9.23k
| Citation Text
list | Functions Text
list | Functions Label
list | Citation Start End
list | Functions Start End
list |
|---|---|---|---|---|---|---|
2018AandA...615L..16F__Gal_et_al._(2014)_Instance_1
|
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−3, 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 × 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 (≳105 cm−3; e.g., Pontoppidan 2006), causes their abundances to be enhanced (e.g., Aikawa et al. 2005).
|
[
"Le Gal et al. (2014)"
] |
[
"According to the pseudo-time-dependent gas-phase astrochemical model under dense cloud conditions (104 cm−3, 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 × 105 yr, and it occurs after the significant fraction of nitrogen is frozen out."
] |
[
"Compare/Contrast",
"Compare/Contrast"
] |
[
[
634,
654
]
] |
[
[
508,
633
],
[
656,
967
]
] |
2018MNRAS.481..138B__Mullaney_et_al._2013_Instance_1
|
To mitigate this problem, several studies have suggested to use the width of the [O iii]λ5007 Å emission line (hereafter [O iii]) originating in the narrow-line region (NLR) as a surrogate for σ⋆, 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 iii] 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 iii] 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 iii] line (σ[O iii]), 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 ii]λλ6716, 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 iii] line. While all studies confirm the original findings by Nelson & Whittle (1996), i.e. a moderately strong correlation between σ⋆ and σ[O iii] 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).
|
[
"Mullaney et al. 2013"
] |
[
"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 iii] blue wing"
] |
[
"Motivation",
"Motivation"
] |
[
[
969,
989
]
] |
[
[
818,
879
],
[
1009,
1189
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2021AandA...653A..36M__Goulding_&_Alexander_(2009)_Instance_6
|
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 − 36 μm 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 − 35 μm 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 μm. 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 μm for the [NeII]12.8 μm and [NeIII]15.6 μm lines, and the continuum at 25 μm for the [OIV]25.9 μm, [FeII]26 μm, [SIII]33.5 μm, and [SiII]34.8 μm 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 − 36 μm 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 − 205 μm interval were taken from Díaz-Santos et al. (2017). For the GOALS sample, 20 starburst galaxies were taken from Fernández-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’ 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).
|
[
"Goulding & Alexander (2009)"
] |
[
"Both Bernard-Salas et al. (2009) and",
"reported data from the high-resolution Spitzer-IRS spectra."
] |
[
"Similarities",
"Similarities"
] |
[
[
2800,
2827
]
] |
[
[
2763,
2799
],
[
2828,
2887
]
] |
2015ApJ...803...79L__Güsten_et_al._2006_Instance_1
|
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; Güsten et al. 2006) telescope25
25
This publication is based on data acquired with the Atacama Pathfinder Experiment (APEX). APEX is a collaboration between the Max-Planck-Institut für Radioastronomie, the European Southern Observatory, and the Onsala Space Observatory.
and at 2.1 GHz (5″ resolution) with the Australia Telescope Compact Array (ATCA) of a sample of massive SZE-selected galaxy clusters. We call the project “LASCAR,” 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’ 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 ΛCDM cosmology with
,
, and
(Komatsu et al. 2011).
|
[
"Güsten et al. 2006"
] |
[
"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"
] |
[
"Uses",
"Uses",
"Uses"
] |
[
[
734,
752
]
] |
[
[
361,
659
],
[
690,
732
],
[
754,
763
]
] |
2019ApJ...875...90L__Velli_et_al._2015_Instance_2
|
When energy flows from the interior of the Sun outward into the solar atmosphere, why is the Sun’s 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—all 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.
|
[
"Velli et al. 2015"
] |
[
"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"
] |
[
"Motivation"
] |
[
[
1997,
2014
]
] |
[
[
1648,
1918
]
] |
2019MNRAS.488.4638L__Drabek-Maunder_et_al._2016_Instance_1
|
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 – ‘non-outflow contribution’/‘all contributions’. We observe that the outflow has a contribution in the FWHM: about 20 per cent in the local region at the H ii region (non-outflow contribution is about 81 per cent) and about 10 per cent even in the clumps. According to Eturb = (3/16 ln 2)Mcloud × FWHM2, outflow has a contribution in the turbulent energy up to 35 per cent in the local region at the H ii region (1 − 0.812). It has a contribution of at least 15 per cent 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).
|
[
"Drabek-Maunder et al. 2016"
] |
[
"It has a contribution of at least 15 per cent in the clump at early stages of massive star formation, which is lower than that reported in previous studies"
] |
[
"Differences"
] |
[
[
753,
779
]
] |
[
[
579,
734
]
] |
2018AandA...619A.177B__Mayne_(2010)_Instance_1
|
The wealth of precise all-sky data from the Gaia Data Release 2 (DR2) revealed a new feature in the Herzsprung–Russell diagram (HRD), namely a gap in the mid-M dwarf main sequence (Jao et al. 2018). The gap appears at a magnitude MG ∼ 10 and colour GBP – GRP ∼ 2.3−2.5 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 ∼0.31 M⊙−0.34 M⊙ 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.
|
[
"Mayne (2010)"
] |
[
"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."
] |
[
"Background",
"Background"
] |
[
[
615,
627
]
] |
[
[
600,
614
],
[
628,
1086
]
] |
2015AandA...584A..76S__Bernstein_et_al._(1995)_Instance_1
|
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{$\Delta H_{0}^{\circ} = -305$}ΔH0◦=−305 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.
|
[
"Bernstein et al. (1995)"
] |
[
"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"
] |
[
"Compare/Contrast",
"Compare/Contrast"
] |
[
[
1187,
1210
]
] |
[
[
1155,
1186
],
[
1212,
1359
]
] |
2019AandA...629A.134G__Leitherer_et_al._2010_Instance_1
|
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 I-, He I-, and He II-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 M⊙ and 100 M⊙, 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 = 0.014, 0.008, 0.002, and 0.001 together with our models with metallicity of Z = 0.014, 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 ∼5000 Å. 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 ∼5000 Å. 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.
|
[
"Leitherer et al. 2010"
] |
[
"We chose to compare to the",
"spectral models from WM-BASIC for OB-stars"
] |
[
"Uses",
"Uses"
] |
[
[
944,
965
]
] |
[
[
732,
758
],
[
877,
919
]
] |
2019MNRAS.485.3288I__Tsao,_Silberberg_&_Barghouty_1998_Instance_1
|
Coming back to the flux values given in Table 1, it is tempting to explain the deficiency of the measured fluxes at energies ∼68 and ∼78 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 ${}^{44}_{}\mathrm{Ti}$ 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 ${}^{44}_{}\mathrm{Ti}$ or excited ${}^{44}_{}\mathrm{Ca}^{*}$ (Silberberg, Tsao & Barghouty 1998; Tsao, Silberberg & Barghouty 1998), the validity of such an interpretation of the discrepant measurements of ${}^{44}_{}\mathrm{Ti}$ 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 ${}^{44}_{}\mathrm{Ti}$ line flux compared to that at 68 and 78 keV as very problematic.
|
[
"Tsao, Silberberg & Barghouty 1998"
] |
[
"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 ${}^{44}_{}\\mathrm{Ti}$ or excited ${}^{44}_{}\\mathrm{Ca}^{*}$",
"the validity of such an interpretation of the discrepant measurements of ${}^{44}_{}\\mathrm{Ti}$ fluxes at 68 and 78 keV and at 1.157 MeV by different instruments and by the same instrument (SPI) may be doubted."
] |
[
"Compare/Contrast",
"Compare/Contrast"
] |
[
[
850,
883
]
] |
[
[
557,
813
],
[
886,
1098
]
] |
2015AandA...577A..43S__Odstrčil_&_Karlický_(1997)_Instance_1
|
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 Karlický & Jungwirth (1989) and Odstrčil & Karlický (1997). Karlický & 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 (Bárta & Karlický 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. Odstrčil & Karlický (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.
|
[
"Odstrčil & Karlický (1997)"
] |
[
"These triggering processes were numerically studied by Karlický & Jungwirth (1989) and"
] |
[
"Background"
] |
[
[
869,
895
]
] |
[
[
782,
868
]
] |
2019MNRAS.482.3656R__Kashi_&_Soker_2011_Instance_1
|
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-Siguán 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.
|
[
"Kashi & Soker 2011"
] |
[
"Additional evolutionary channels for SNIa include",
"the core-degenerate channel"
] |
[
"Background",
"Background"
] |
[
[
1037,
1055
]
] |
[
[
754,
803
],
[
965,
992
]
] |
2020MNRAS.497.4231R__Brown,_Vasil_&_Zweibel_2012_Instance_1
|
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:
(8)$$\begin{eqnarray*}
\nabla \cdot \overline{\rho } \boldsymbol {v} = 0,
\end{eqnarray*}$$(9)$$\begin{eqnarray*}
\frac{\partial \boldsymbol {v}}{\partial t} = - \nabla \left(\frac{P}{\overline{\rho }}\right) - C \overline{g} \boldsymbol {\hat{r}}
\end{eqnarray*}$$(10)$$\begin{eqnarray*}
\frac{\partial T}{\partial t} = - v_\mathrm{ r} \left(\frac{\partial \overline{T}}{\partial r} - (\gamma - 1) \overline{T} h_\rho \right)
\end{eqnarray*}$$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:
(11)$$\begin{eqnarray*}
0 &=& \frac{\partial ^2 \alpha }{\partial r^2} + \left(\frac{N^2}{\omega ^2} - 1 \right) \frac{m^2}{r^2} \alpha + \left[ -\overline{\rho }^{-1/2}\frac{\partial ^2 \left(\overline{\rho }^{1/2}\right)}{\partial r^2} + \frac{\partial h_{\rho }}{\partial r} \right]\alpha \nonumber \\
&&+\, \frac{1}{4r^2} \alpha,
\end{eqnarray*}$$where $\alpha = v_\mathrm{ r} \overline{\rho }^{1/2} r^{3/2}$. We have used a wave ansatz of the form vr(r, θ, z) ∝ vr(r)eimθe−iωt, where m is the horizontal wavenumber,1 ω is the angular frequency, and θ is the angular coordinate in the equatorial plane. Generally, in the radiation zone, there will be regions where the oscillatory term (OT), (N2/ω2 − 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.
|
[
"Brown, Vasil & Zweibel 2012"
] |
[
"We ignore the pressure term in equation (4), which leads to a set of equations that conserve energy"
] |
[
"Uses"
] |
[
[
1049,
1076
]
] |
[
[
948,
1047
]
] |
2019ApJ...886...14F__Urquhart_et_al._2014_Instance_1
|
Figure 6(a) shows the 1.3 mm continuum and C18O(J = 2–1) distribution around the Papillon Nebula. The extended 1.3 mm continuum emission around the Papillon Nebula YSO roughly agrees with the distribution of Hα emission and has no counterparts of molecular gas, indicating that the continuum emission seems to be dominated by free–free 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)–(c). Because we detected extended emission in 1.3 mm and C18O along the north–south 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 ∼1 × 1024 cm−2, and ∼2 × 102 M⊙, respectively, with an assumption of the dust emissivity,
of 1 cm2 g−1 for protostellar envelopes (e.g., Ossenkopf & Henning 1994), a dust-to-gas ratio of 3.0 × 10−3 (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 ∼30 km s−1 in 12CO(J = 2–1) toward MMS-1 and MMS-2. Figures 6(b)–(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 2σ level of the velocity-smoothed spectra with a smoothing kernel of 9 ch (=1.8 km s−1). We extracted the spectra at positions ∼1″ away from the millimeter sources in the east direction as the references without the outflow contaminations and defined the velocities below 2σ 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 ∼0.1 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−1. 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 ∼10−3–10−2 M⊙ km s−1 yr−1, depending on the assumption of the outflow inclination angle (i = 30°–70°). 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 ∼102 M⊙ 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 ≲104 yr. We note that the Papillon Nebula YSO is more evolved than the two sources; however, the age is as young as ∼0.1 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 ∼1–2 pc are formed within the order of ∼0.1 Myr. We discuss the possible processes of high-mass star formation taking place in the N159E-Papillon region in Section 4.1.
|
[
"Urquhart et al. 2014"
] |
[
"The peak column densities and total masses of both sources deduced from the 1.3 mm continuum emission are ∼1 × 1024 cm−2, and ∼2 × 102 M⊙, 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.,"
] |
[
"Uses",
"Uses"
] |
[
[
1325,
1345
]
] |
[
[
839,
1010
],
[
1193,
1324
]
] |
2017ApJ...850...97B__Tamburro_et_al._2009_Instance_2
|
The H i mass fraction of every gas particle in the baryonic runs is calculated based on the particle’s 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
. 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’s 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, ∼10–15 km s−1 (e.g., Stanimirović 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,
, 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 ∼1 kpc resolution in our lowest-mass galaxies up to ∼9 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−1 (two-channel boxcar-smoothed).
|
[
"Tamburro et al. 2009"
] |
[
"Dispersions are thought to be driven at least partially by thermal velocities or supernovae"
] |
[
"Compare/Contrast"
] |
[
[
1100,
1120
]
] |
[
[
1007,
1098
]
] |
2015ApJ...801...88S__Liu_et_al._2008_Instance_1
|
There are several different effects that have been considered to explain the observed offset in the [O iii]λ5007/Hβ vs. [N ii]λ6584/Hα 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]λ5007/Hβ vs. [N ii]λ6584/Hα 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]λ5007/Hβ vs. [N ii]λ6584/Hα diagram toward higher [O iii]/Hβ and [N ii]/Hα 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]/Hβ line ratios in excess of the “maximum starburst” 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 “composite” region of the [O iii]λ5007/Hβ vs. [N ii]λ6584/Hα 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—specifically, the N/O ratio—which can affect where galaxies fall in the [O iii]λ5007/Hβ vs. [N ii]λ6584/Hα BPT diagram. If the relationship between N/O and
evolves out to high redshift, then distant galaxies will shift relative to local ones in the [O iii]λ5007/Hβ vs. [N ii]λ6584/Hα BPT diagram.
|
[
"Liu et al. 2008"
] |
[
"Early work highlighting the issue of the high-redshift offset in the [O iii]λ5007/Hβ vs. [N ii]λ6584/Hα BPT diagram focused on these parameters (e.g.,"
] |
[
"Background"
] |
[
[
631,
646
]
] |
[
[
435,
585
]
] |
2022MNRAS.510.3479B__Perucho_et_al._2010_Instance_1
|
Our simulations, despite probing quite different jet velocities and resolution levels, point to a very unstable nature of the structure resulting from the jet–wind–orbit 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 χj values and longer Porb values than those explored here, we propose that jet–wind–orbit interaction could induce instability growth, which could produce precession–like 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 Φ is initially rather small, the growth of Kelvin–Helmholtz (KH) instabilities is expected (see also, e.g. Perucho 2019, in the context of extragalactic jets). In fact, the jet–wind 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–Taylor and Richtmyer–Meshkov instabilities could also develop at the dynamical jet–wind 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, χj values even smaller than θj (see Bosch-Ramon & Barkov 2016, for a discussion on the χj – θj relation) could also lead to distorted jets on scales ≫a. 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 χj value, if the jet gets disrupted on scales $\lesssim 10\, a$, the strong pressure drop outwards of the embedding medium can make the jet–wind 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 ≫ a in HMMQ, devoted simulations are planned.
|
[
"Perucho et al. 2010"
] |
[
"In fact, the jet–wind interaction alone can already perturb the jet"
] |
[
"Background"
] |
[
[
1208,
1227
]
] |
[
[
1111,
1178
]
] |
2016AandA...589A.132B__Sarangi_&_Cherchneff_2013_Instance_1
|
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).
|
[
"Sarangi & Cherchneff 2013"
] |
[
"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"
] |
[
"Background"
] |
[
[
668,
693
]
] |
[
[
360,
611
]
] |
2018ApJ...866...48U__Plagge_et_al._2013_Instance_1
|
RX J1347.5–1145 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–Zel’dovich 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 ∼6 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., Köhlinger & Schmidt 2014). The total mass of RX J1347.5–1145 within r200 is estimated to be ∼1.5 × 1015 h−1
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−1 Mpc (Lu et al. 2010) for this galaxy cluster.18
18
They adopted the Hubble constant of 70 km s−1 Mpc−1.
|
[
"Plagge et al. 2013"
] |
[
"The presence of the substructure has been confirmed by Chandra and XMM-Newton",
"as well as by more recent SZE measurements"
] |
[
"Background",
"Background"
] |
[
[
915,
933
]
] |
[
[
703,
780
],
[
831,
873
]
] |
2021ApJ...912..106Y__Minchev_et_al._2013_Instance_1
|
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-[α/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 (López-Corredoira & Molgó 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; García de la Cruz et al. 2021), which allows for the low-[α/Fe] stars to exist several kpc above the disk’s midplane. As analyzed by B16, the exponential flaring profiles for the low-[α/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 “churning,” 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-[α/Fe] MAP indicates the radial migration has occurred in the formation of the thick disk as suggested by model simulations (e.g., Schönrich & 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 = 14–15 kpc as stated by Ruphy et al. (1996), which is also discovered by B16.
|
[
"Minchev et al. 2013"
] |
[
"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"
] |
[
"Similarities",
"Similarities"
] |
[
[
789,
808
]
] |
[
[
486,
543
],
[
656,
787
]
] |
2020AandA...640A.121G__Nakajima_1983_Instance_1
|
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 (Bétrémieux & 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 (García Muñoz 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.
|
[
"Nakajima 1983"
] |
[
"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"
] |
[
"Background"
] |
[
[
719,
732
]
] |
[
[
463,
653
]
] |
2018AandA...617A.116L__Hendecourt_&_Allamandola_(1986)_Instance_1
|
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 μm (2833 cm−1) and 9.74 μm (1027 cm−1) bands. They correspond to the CH3 symmetric stretching mode and the CO
stretching mode, respectively. The laboratory data most widely used for comparison to observations are those provided by d’Hendecourt & 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 2.5–20 μm range. The second reference presents a thorough infrared study providing a compendium of band strengths and optical constants in the 2.5–200 μm 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 μm band has been pointed out.
|
[
"d’Hendecourt & Allamandola (1986)"
] |
[
"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 2.5–20 μm range."
] |
[
"Uses",
"Uses"
] |
[
[
640,
673
]
] |
[
[
550,
639
],
[
725,
828
]
] |
2019AandA...622A.106M__Lanz_et_al._(2010)_Instance_1
|
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; González-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; López-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, “multifrequency detection”. 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 μm. As discussed in Hopwood et al. (2010), Lapi et al. (2011), González-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 ∼ 2, that is redshifted from its rest-frame wavelength around 70–100 μm 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 ≳ 4 (Michałowski 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 μm (the so-called “500 μm-risers”), 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 μm 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 μm-riser candidates.
|
[
"Lanz et al. 2010"
] |
[
"However, over the last few years a number of multifrequency compact source detection techniques have been proposed in the literature"
] |
[
"Background"
] |
[
[
1867,
1883
]
] |
[
[
1691,
1823
]
] |
2020ApJ...901...21B__Cane_2000_Instance_1
|
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; Dumbović 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; Dumbović et al. 2018), and the anticoincidence shield of the International Gamma-Ray Astrophysics Laboratory’s 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’s Magnetosphere Imaging Instrument and Low Energy Magnetospheric Measurement System measurements (Roussos et al. 2018) and beyond (Witasse et al. 2017; Winslow et al. 2018).
|
[
"Cane 2000"
] |
[
"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."
] |
[
"Background",
"Background"
] |
[
[
1219,
1228
]
] |
[
[
990,
1218
],
[
1231,
1288
]
] |
2021ApJ...915...86A__Troja_et_al._2017_Instance_1
|
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 & Mészáros 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.
|
[
"Troja et al. 2017"
] |
[
"Indeed, the scientific debate about the emission profile of the jet and the subsequent gamma-ray production mechanism of GRB 170817A is still ongoing"
] |
[
"Motivation"
] |
[
[
932,
949
]
] |
[
[
714,
863
]
] |
2021ApJ...912..163B__Davis_et_al._2018_Instance_1
|
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 ∼1–1.4; 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 ∼0.8–1.1); 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 ∼0.1–0.7 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 ±5 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.
|
[
"Davis et al. 2018"
] |
[
"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"
] |
[
"Motivation"
] |
[
[
1540,
1557
]
] |
[
[
1137,
1429
]
] |
2022MNRAS.509.3427D__Díaz_et_al._2017_Instance_1
|
Unlike BH–BH 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 & Paczyński 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)] ∼11 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; Díaz 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–infrared (IR) ‘kilonova’ 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 10 000 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 ∼0.3c.
|
[
"Díaz et al. 2017"
] |
[
"GRB 170817A was followed by the discovery of the optical counterpart Swope Supernova Survey 17a [SSS17a (AT 2017gfo)] ∼11 h later in the galaxy NGC 4993",
"and later confirmed by others teams"
] |
[
"Background",
"Background"
] |
[
[
895,
911
]
] |
[
[
622,
774
],
[
797,
832
]
] |
2016AandA...595A..16T__Eggleton_et_al._(2006)_Instance_1
|
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⊙ 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 (μ) inversion in their 1 M⊙ 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 μ-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 μ-inversion is low and negligible compared to a stabilising μ-stratification. Charbonnel & Zahn (2007) computed stellar evolution models including the ideas of Kippenhahn et al. (1980), who extended Ulrich’s 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).
|
[
"Eggleton et al. (2006)"
] |
[
"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"
] |
[
"Background"
] |
[
[
886,
908
]
] |
[
[
731,
885
]
] |
2018AandA...614A..72V__Bocquet_et_al._2015_Instance_1
|
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’eldovich 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.
|
[
"Bocquet et al. 2015"
] |
[
"Cluster number counts from Sunyaev-Z’eldovich 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."
] |
[
"Background"
] |
[
[
844,
863
]
] |
[
[
505,
773
]
] |
2022AandA...661A..71H__Liu_et_al._2020_Instance_1
|
While using the Padé(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 γ is studied. One case is the fitting parameters q0 and j0 in the case of fixed parameters δ, γ, and β. The other is to refit q0 and j0 by changing the setting of the fixed parameter γ within a 1σ error (0.005). The fixed parameters δ, γ, and β are given in terms of the best fitting results from 1598 quasars in a flat ΛCDM model. By substituting Eqs. (25) into (26), and replacing dL with Eq. (6), we obtained the best fits, that is
Ω
m
=
0
.
65
−
0.19
+
0.16
$ \Omega_{m} = 0.65^{+0.16}_{-0.19} $
, δ = 0.23 ± 0.01, γ = 0.62 ± 0.01, and β = 7.60 ± 0.28. This result is consistent with previous research within a 1σ 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 1σ errors of the γ 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 δ = 0.23 ± 0.01, γ = 0.64 ± 0.01,
β
=
7
.
17
−
0.27
+
0.28
$ \beta = 7.17^{+0.28}_{-0.27} $
, q0 = −0.67 ± 0.04, and
j
0
=
2
.
43
−
0.46
+
0.51
$ j_{0} = 2.43^{+0.51}_{-0.46} $
. This result is in line with the previous results obtained by the logarithmic polynomials. The corresponding q0 − j0 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 δ = 0.23, γ = 0.62, and β = 7.60 (best fits from 1598 quasars in a flat ΛCDM model), the results are q0 = −0.82 ± 0.05 and
j
0
=
4
.
99
−
0.63
+
0.69
$ j_{0} = 4.99^{+0.69}_{-0.63} $
, which is represented by the red contours in Fig. 5. There exists more than a 4σ tension with the flat ΛCDM. (2) In only changing γ = 0.625, the new result is q0 = −0.58 ± 0.03 and
j
0
=
1
.
19
−
0.31
+
0.34
$ j_{0} = 1.19^{+0.34}_{-0.31} $
which is described by the blue contours of Fig. 5. We find that the γ value changes from 0.62 to 0.625 (within a 1σ error of 0.01), and the 4σ tension disappears. This result is consistent with the ΛCDM model within a 1σ level. We find that the quasar relation can obviously affect the cosmographic constraint, especially the slope parameter γ.
|
[
"Liu et al. 2020"
] |
[
"This result is consistent with previous research within a 1σ error",
"Here, it is worth noting that the 1σ errors of the γ best fit from previous research are both larger than 0.01."
] |
[
"Similarities",
"Differences"
] |
[
[
1018,
1033
]
] |
[
[
862,
928
],
[
1054,
1165
]
] |
2021MNRAS.507.2766S__Sumiyoshi_et_al._2005_Instance_2
|
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, νe, $\bar{\nu }_\mathrm{ e}$, νμ/τ, and $\bar{\nu }_{\mu /\tau }$, i.e. we implement six species of neutrinos by assuming μ-type and τ-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
(1)$$\begin{eqnarray*}
\mathrm{ d}s^2 = -\mathrm{ e}^{2\Phi (t,m_\mathrm{ b})}\mathrm{ d}t^2 + \mathrm{ e}^{2\Lambda (t,m_\mathrm{ b})}\mathrm{ d}m_\mathrm{ b}^2 + r^2(t,m_\mathrm{ b})(\mathrm{ d}\theta ^2 + \sin ^2\theta \mathrm{ d}\mathrm{ }\phi ^2), \nonumber\\
\end{eqnarray*}$$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, Φ(t, mb) and Λ(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.
|
[
"Sumiyoshi et al. 2005"
] |
[
"For the collision term associated with neutrino emission, absorption, and scattering with leptons, nucleons, and nuclei handles, the basic neutrino reactions are adopted"
] |
[
"Uses"
] |
[
[
1312,
1333
]
] |
[
[
1128,
1297
]
] |
2020AandA...641A.126B__Jones_&_Hardee_1979_Instance_1
|
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 “hot spot” of material, orbiting near the innermost stable circular orbit (ISCO) of Sgr A* in a poloidal magnetic field, emits polarized synchrotron radiation.
|
[
"Jones & Hardee 1979"
] |
[
"In the optically thin limit, the emission is significantly polarized",
"an effect that has been observed in higher-luminosity AGN sources"
] |
[
"Background",
"Similarities"
] |
[
[
728,
747
]
] |
[
[
658,
726
],
[
750,
815
]
] |
2022MNRAS.513.4464T__Hopkins_et_al._2020_Instance_1
|
Galactic winds: Galactic winds driven by CRs have often been simulated in two limits: a diffusion-dominated regime, due possibly to ‘extrinsic confinement’, where CRs are scattered by extrinsic turbulence, and/or due to various wave damping mechanisms (e.g. ion neutral damping) and streaming-dominated ‘self confinement’, where CRs are confined by Alfven waves they produce via the gyroresonant streaming instability. In the diffusive ‘extrinsic confinement’ case, CRs do not heat the gas.19 In the streaming dominated ‘self confinement’ case, CR transport heats gas at a rate vA · ∇Pc. The diffusive case fits γ 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 ∼GeV 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 ∼GeV 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 – in contrast to, for instance, van de Voort et al. (2021) – the plasma β 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.
|
[
"Hopkins et al. 2020"
] |
[
"It is also much better at driving winds, because the CRs do not suffer strong energy losses via Alfven wave heating"
] |
[
"Uses"
] |
[
[
849,
868
]
] |
[
[
711,
826
]
] |
2016ApJ...821..107G__Schwadron_et_al._2011_Instance_1
|
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
j
ENA
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, ΔP × l, in the following way:
3
Δ
P
×
l
=
4
π
3
n
H
m
H
v
j
ENA
(
E
)
σ
(
E
)
Δ
E
c
f
4
c
f
=
(
v
+
u
R
)
2
v
4
(
v
2
+
4
u
R
2
+
2
u
R
v
)
.
In Equation (3), ΔE 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−1 as measured by Voyager 2, whereas uR = 40 km s−1 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−3 is assumed (Schwadron et al. 2011; Gloeckler & Fisk 2015). The charge-exchange cross section between protons and neutral hydrogen decreases from (4 to 2) × 10−15 cm−2 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 × l = 304 pdyn cm−2 au for the downwind hemisphere and 66 pdyn cm−2 au for the Voyager 1 region (1 pdyn cm−2 au = 0.015 N m−1). 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−2 sr−1 s−1 keV−1 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−2 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−2 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.
|
[
"Schwadron et al. (2011)"
] |
[
"We repeated the plasma pressure calculation presented by",
"and Fuselier et al. (2012) for the new ENA energy spectrum."
] |
[
"Uses",
"Uses"
] |
[
[
57,
80
]
] |
[
[
0,
56
],
[
81,
140
]
] |
2018ApJ...859..115W__Luo_et_al._2016_Instance_1
|
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
to
, and with a mean value of
. 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 “strangulation” 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% ∼ 30% satellites reside at the low values of
, 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 ∼5% 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 ∼3% 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.
|
[
"Luo et al. 2016"
] |
[
"Such evolutionary behavior is consistent with “strangulation” wherein cold gas that could be used for star formation is depleted as satellites are accreted (e.g.,"
] |
[
"Similarities"
] |
[
[
685,
700
]
] |
[
[
481,
643
]
] |
2019MNRAS.489.4429R__Lada_et_al._2008_Instance_1
|
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 × 103 M⊙ (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; Román-Zúñiga, Lada & Alves 2009; Román-Zúñiga 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. Román-Zúñiga et al. 2010; Alfaro & Román-Zúñiga 2018, hereafter RAL10 and AR18, respectively). The Pipe was also the first cloud where a ‘core mass function’ (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.
|
[
"Lada et al. 2008"
] |
[
"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."
] |
[
"Background",
"Background",
"Background"
] |
[
[
914,
930
]
] |
[
[
574,
590
],
[
770,
835
],
[
1079,
1248
]
] |
2021ApJ...916...70Z__Millon_et_al._2020_Instance_1
|
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 ∼0.8% uncertainty on DΔt (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 ∼2% 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–lens 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.
|
[
"Millon et al. 2020"
] |
[
"For instance, the TDCOSMO collaboration",
"has achieved ∼2% 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."
] |
[
"Background",
"Background",
"Background"
] |
[
[
1041,
1059
]
] |
[
[
1000,
1039
],
[
1061,
1277
],
[
1300,
1318
]
] |
2017AandA...602A..26L__Robrade_&_Schmitt_2005_Instance_1
|
Along with the flare related changes, we also examine possible orbital variations by dividing the overall quiescent spectra in periods dubbed as “Quiet 1” and “Quiet 2” (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 (Güdel 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 0.2–10 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.
|
[
"Robrade & Schmitt 2005"
] |
[
"As is often observed",
", we require multi-temperature components to achieve an adequate description of the observed coronal spectra."
] |
[
"Similarities",
"Similarities"
] |
[
[
607,
629
]
] |
[
[
566,
586
],
[
651,
760
]
] |
2021MNRAS.504.3316B__than_2000_Instance_5
|
WASP-43b is the most heavily scrutinized phase curve, with four analyses of this data set already published (Stevenson et al. 2017; Mendonça et al. 2018; Morello et al. 2019; May & Stevenson 2020). Our phase curve semi-amplitude, eclipse depth, and radius are consistent with all of these works. The more contentious issue is that of the phase curve’s phase offset and nightside temperature. Stevenson et al. (2017) initially reported only a 2σ upper limit on the nightside temperature of 650 K, while all subsequent reanalyses (including ours) favour a significantly detectable nightside temperature of ∼800 K. As for the planet’s phase offset, Stevenson et al. (2017) and May & Stevenson (2020) favour a larger phase offset (21 ± 2 °E) than Mendonça et al. (2018) and Morello et al. (2019) (12 ± 3 °E and 11 ± 2 °E). May & Stevenson (2020) claimed that the differences between the retrieved phase offsets is the result of temporal binning which was not used by Stevenson et al. (2017) and May & Stevenson (2020) but was used by Mendonça et al. (2018), Morello et al. (2019), and this work. Fitting the temporally binned photometry for all 17 phase curves with each of our detector models already required more than 2000 CPU hours, and expanding this to unbinned photometry for all phase curve fits would require more than 125 000 CPU hours (or 434 d using our 12× multithreading computer) optimistically assuming all of detector models scaled linearly with the number of input data. However, we did try fitting just the WASP-43b unbinned phase curve with our preferred detector model (BLISS) and found that our phase offset and nightside temperature was unchanged. Including a linear slope in time also did not affect our phase offset or nightside temperature. Instead, we find that the phase offset inferred by our models depends on the choice of phase curve model, as our 4-parameter (v2) phase curve models are consistent with those of Stevenson et al. (2017) and May & Stevenson (2020), while our 2-parameter phase curve models (v1) are consistent with Mendonça et al. (2018) and Morello et al. (2019). Ultimately, we cannot decide between these two discrepant offsets as the ΔBIC between the two phase curve models for our preferred BLISS detector model is only 3.7 (insignificantly favouring the 20.4 ± 3.6 offset from the v2 model). For reference, Stevenson et al. (2014b) found phase offsets ranging from roughly −6 to 17 deg east in the Hubble/WFC3 bandpass.
|
[
"Stevenson et al. (2017)"
] |
[
"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 Mendonça et al. (2018), Morello et al. (2019), and this work."
] |
[
"Compare/Contrast",
"Compare/Contrast"
] |
[
[
963,
986
]
] |
[
[
819,
962
],
[
987,
1091
]
] |
2018MNRAS.478.4986K__Greig_&_Mesinger_2017_Instance_1
|
Various observational techniques, such as measuring the Gunn–Peterson optical depth from QSO spectra or the prevalence of Ly α 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∼ 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’Odorico 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).
|
[
"Greig & Mesinger 2017"
] |
[
"Various observational techniques, such as measuring the Gunn–Peterson optical depth from QSO spectra or the prevalence of Ly α 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∼ 6"
] |
[
"Background"
] |
[
[
829,
850
]
] |
[
[
0,
322
]
] |
2021AandA...649A.126T__Luck_(2018b)_Instance_1
|
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 ≤4 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−1. 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−1). 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 α-process elements by Li et al. (2018), even though the production sites of α-elements and s-processes dominated elements are quite different.
|
[
"Luck (2018b)"
] |
[
"More recently,",
"also investigated the gradients of n-capture element abundance-to-iron ratios with respect to Rgc for a sample of 435 Cepheids."
] |
[
"Background",
"Background"
] |
[
[
1020,
1032
]
] |
[
[
1005,
1019
],
[
1033,
1160
]
] |
2015MNRAS.451.2544P__Leitet_et_al._2011_Instance_1
|
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 per cent. At z ∼ 1, no objects with leaking ionizing photons have been detected (Bridge et al. 2010; Siana et al. 2010), but at z ∼ 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 ∼10 per cent of the observed objects (Nestor et al. 2013). Attempts to constrain the escape fraction with numerical simulations find ranges between fesc 10 per cent (Razoumov & Sommer-Larsen 2006, 2007; Gnedin, Kravtsov & Chen 2008; Paardekooper et al. 2011; Kim et al. 2013) and fesc > 80 per cent (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.
|
[
"Leitet et al. 2011"
] |
[
"Except for two objects",
"in the local Universe no ionizing radiation has been detected directly",
"although some objects show indirect evidence of photon leakage"
] |
[
"Background",
"Background",
"Background"
] |
[
[
1124,
1142
]
] |
[
[
1100,
1122
],
[
1151,
1221
],
[
1270,
1332
]
] |
2019MNRAS.482.5651M__Schweizer_&_Middleditch_1980_Instance_1
|
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 arcmin 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 ϖ = 0.0736 ± 0.1244, and then σϖ/ϖ = 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 ± 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 ∼6.67 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.
|
[
"Schweizer & Middleditch 1980"
] |
[
"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"
] |
[
"Similarities"
] |
[
[
868,
896
]
] |
[
[
467,
866
]
] |
2022AandA...666A.134S__Rodríguez-Almeida_et_al._2021b_Instance_1
|
A more puzzling task was the resolution of the fine and hyperfine structure. First, the A–E split structure attributed to the methyl internal rotation motion can be generally used as a molecular “fingerprint” 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+, Rodríguez-Almeida et al. 2021b) which has been detected toward G+0693-0.027 along with ethyl isocyante (C2H5NCO, Kolesniková 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 (León 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.
|
[
"Rodríguez-Almeida et al. 2021b"
] |
[
"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."
] |
[
"Background",
"Background",
"Background"
] |
[
[
776,
806
]
] |
[
[
646,
775
],
[
806,
851
],
[
915,
975
]
] |
2021ApJ...912...92F__Feng_et_al._2021_Instance_1
|
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)
2
where 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⊙, 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 ∼ 4–5 is obtained for most reverberation-mapped AGNs with the Hβ 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 Hα and Hβ, 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 “disappearance” of BELs. A long-term monitoring with a good sampling, e.g., twice a year, may give a reliable duration of “disappearance” of BELs in the CL processes, especially, the continuous and complete process of changing of Seyfert 1 → Seyfert 2 → Seyfert 1. This will be important to test the obscuration model of CL-AGNs.
|
[
"Feng et al. 2021"
] |
[
"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"
] |
[
"Uses"
] |
[
[
939,
955
]
] |
[
[
777,
937
]
] |
2021MNRAS.506.3330W__Gronow_et_al._2020_Instance_1
|
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 ‘normal’ SNe Ia as well as ‘peculiar’ 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 ‘normal’ SNe Ia; detailed observations are thus required in order to place constraints on the relative fractions of each possible progenitor channel.
|
[
"Gronow et al. 2020"
] |
[
"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."
] |
[
"Background"
] |
[
[
966,
984
]
] |
[
[
725,
898
]
] |
2021AandA...653A.129C__Coutens_et_al._(2016)_Instance_1
|
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 (López-Sepulcre et al. 2015, 2019; Allen et al. 2020). In the top-left panel of Fig. 15 the best power-law fits derived by López-Sepulcre et al. (2015), X[HC(O)NH2] = 0.04 × X[HNCO]0.93, and by Quénard et al. (2018), X[HC(O)NH2] = 32.14 × X[HNCO]1.29, are compared with the one derived here, X[HC(O)NH2] = 0.006 × 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 Martín-Doménech 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 Quénard 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 → 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.
|
[
"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."
] |
[
"Background"
] |
[
[
1816,
1837
]
] |
[
[
1838,
2035
]
] |
2019ApJ...875...68A__Clemens_&_Alexander_2002_Instance_1
|
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 ∼4.7 M yr−1 (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 0.5–8 keV band) fall into the range (2.6–4.9) × 1039 erg s−1. 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 (0.6–3) × 1039 erg s−1. 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—this source was ignored because of its low luminosity of LX ∼ 6 × 1038 erg s−1, 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 (0.9–4) × 1039 erg s−1 within the 0.5−8 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).
|
[
"Clemens & Alexander 2002"
] |
[
"Also, NGC 4490 has a giant H i envelope that probably originated from the star formation"
] |
[
"Background"
] |
[
[
1199,
1223
]
] |
[
[
1109,
1197
]
] |
2019AandA...628A.110M__Kryukova_et_al._(2012)_Instance_1
|
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 1–3 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⊙ by Whitney et al. (2004), it can be seen that the completeness limits at 3.6 and 4.6μm 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 22μm. Furthermore, the completeness limit at 22 μm 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 μm, after conservatively dereddening it by AV = 20, we assumed a spectral index γ = −2 (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 ~ 1.7–2.8 L⊙, depending on whether the NIR flux is neglected (which may be the case) or extrapolated from γ = −2. A comparison with the birthline of Palla & Stahler (1993) indicates a mass of ~ 0.4–0.5 M⊙ 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 μm 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 μm translates into a bolometric luminosity of the central (proto)star Lbol ~ 0.28 L⊙ (we note that Dunham et al. 2008 indicate this luminosity as Lint). We highlight the fact that the 70 μm 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.
|
[
"Kryukova et al. (2012)"
] |
[
"Alternatively, one can compute the bolometric luminosity following"
] |
[
"Uses"
] |
[
[
1171,
1193
]
] |
[
[
1104,
1170
]
] |
2021AandA...647A.186S__Shebalin_et_al._1983_Instance_1
|
Chandrasekhar & Fermi (1953) assumed that “turbulent motions are isotropic”, and they adopted
$f\,{=}\,1/\sqrt{3}$f = 1/3
. If the field strength is weak, turbulent motions will drag the field lines in random directions and turbulence will be isotropic (super-Alfvénic turbulence). However, there is overwhelming observational evidence that magnetic fields in the ISM have well-defined directions, indicating that turbulence is sub, trans-Alfvénic 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-Alfvénic 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-Alfvénic 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-Alfvénic as well. All these lines of evidence indicate that ISM turbulence is sub, trans-Alfvénic, and hence anisotropic.
|
[
"Shebalin et al. 1983"
] |
[
"However, there is overwhelming observational evidence that magnetic fields in the ISM have well-defined directions, indicating that turbulence is sub, trans-Alfvénic and hence turbulent properties are highly anisotropic"
] |
[
"Background"
] |
[
[
549,
569
]
] |
[
[
284,
503
]
] |
2021AandA...650A.155Z__Oh_et_al._2012_Instance_3
|
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.
|
[
"Oh et al. 2012"
] |
[
"However, the question of whether galactic bars can significantly affect AGN activity is still under debate"
] |
[
"Background"
] |
[
[
903,
917
]
] |
[
[
756,
862
]
] |
2015ApJ...811..129B__Pasachoff_et_al._2009_Instance_1
|
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 “classical” 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 (∼5 minutes and 25 km s−1). (N. E. Raouafi et al. 2015, in preparation) suggest that the term “classical spicules” 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 (30–110 km s−1) and shorter lifetimes (50–150 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 Å 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 Å 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).
|
[
"Pasachoff et al. 2009"
] |
[
"N. E. Raouafi et al. (2015, in preparation) have argued that type II spicules have much more in common with so-called “classical” 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."
] |
[
"Similarities",
"Similarities"
] |
[
[
772,
793
]
] |
[
[
612,
750
],
[
795,
1008
]
] |
2019MNRAS.488.1728N__Kaiser_et_al._1995_Instance_1
|
The mass estimate for each cluster was obtained by fitting a reduced tangential shear profile predicted by a projected Navarro–Frenk–White (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
(4)
\begin{eqnarray*}
\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}}}.
\end{eqnarray*}
Here gi(θi, Σcrit, 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 ∼5 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 0.2–4.2 Mpc for the fitting procedure. We minimized the χ2 on a grid of R200 and c200. Finally, we used the mass–concentration 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 β 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 β distribution that are averaged over when calculating βi from a reference catalogue (equation 3), we find a biased point estimate for βi. 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 β 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 θ from the cluster centre, and βk. We then introduce $\beta _i^{\prime }$, which satisfies the equation:
(5)
\begin{eqnarray*}
g(\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})}.
\end{eqnarray*}
|
[
"Kaiser et al. 1995"
] |
[
"The factor 1.08 is the multiplicative shear calibration bias of the used KSB+ pipeline",
"to convert from measured to true ellipticity."
] |
[
"Uses",
"Uses"
] |
[
[
886,
904
]
] |
[
[
798,
884
],
[
925,
970
]
] |
2018AandA...618L...3M__Greaves_&_Rice_(2010)_Instance_1
|
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 ∼1–3 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 ∼350 M⊕ of pebbles needed to grow the core of Jupiter from half a lunar mass to 20 M⊕ (e.g., Morbidelli et al. 2016). Thus, such an inefficiency would imply that disks were initially ∼10–100 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 ∼10–100 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).
|
[
"Greaves & Rice (2010)"
] |
[
"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."
] |
[
"Compare/Contrast",
"Compare/Contrast"
] |
[
[
525,
546
]
] |
[
[
395,
524
],
[
548,
888
]
] |
2021MNRAS.507.5882S__Mackereth_et_al._2018_Instance_1
|
Cosmological hydro dynamical N-body simulations offer another possibility to investigate the origin of the bimodality in the ([Fe/H], [α/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], [α/Fe]) plane, with the thick disc being old metal poor and rich in [α/Fe] and the thin disc being young, metal-rich and poor in [α/Fe]. They also showed that migration was important to generate the two discs. However, a clear bimodality in the ([Fe/H], [α/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], [α/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-[α/Fe] sequence a gas-rich merger dilutes the metallicity of the ISM leading to the formation of the low-[α/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 Schönrich & Binney (2009a) paradigm.
|
[
"Mackereth et al. 2018"
] |
[
"The bimodality has now been observed in some simulations"
] |
[
"Similarities"
] |
[
[
983,
1004
]
] |
[
[
906,
962
]
] |
2020MNRAS.498.5116I__Andrews_et_al._2012_Instance_1
|
It is now largely understood that the dust-to-gas mass ratio ϵ 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 ϵ values to be much higher (∼0.2; see Ansdell et al. 2016), and simulations have demonstrated that discs can be formed with significant dust enrichment (Lebreuilly, Commerçon & 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; Pérez 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
(1)$$\begin{eqnarray*}
v_{\mathrm{r}}=v_{\mathrm{drag}}+v_{\mathrm{drift}}=\frac{v_{\mathrm{g}, \mathrm{r}}}{1+{St}^{2}}+\frac{1}{{St}+{St}^{-1}} \frac{1}{\rho _{\mathrm{g}} \Omega } \frac{\partial P}{\partial R},
\end{eqnarray*}$$where vg, r is the radial velocity of the gas, ρg is the gas mass surface density, Ω 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
(2)$$\begin{eqnarray*}
St = \frac{\pi a\rho _\mathrm{s}}{2\Sigma _\mathrm{g}},
\end{eqnarray*}$$where a is grain size, ρs is intrinsic bulk grain density, and Σg 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.
|
[
"Andrews et al. 2012"
] |
[
"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"
] |
[
"Background"
] |
[
[
664,
683
]
] |
[
[
493,
662
]
] |
2018ApJ...858...13H___2009b_Instance_1
|
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 (
) white dwarf (WD) (Hoyle & Fowler 1960) near the Chandrasekhar mass (
). 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; Lorén-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-
WD with accretion from a helium (
) 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
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
. (Hoyle & Fowler 1960; Sugimoto & Nomoto 1980; Nomoto 1982b; Hoeflich & Stein 2002; Piersanti et al. 2003). It is likely that the dynamical merger,
explosion, and double-detonation channels all contribute to the SN Ia population because of the “stellar amnesia” 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).
|
[
"Wang et al.",
"2009b"
] |
[
"This process can end up as a SN Ia, a highly magnetized WD (MWD), or an accretion-induced collapse"
] |
[
"Background"
] |
[
[
992,
1003
],
[
1011,
1016
]
] |
[
[
744,
842
]
] |
2018ApJ...853...34Z__Giebels_et_al._2007_Instance_2
|
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)–flux diagram is often used as a diagnostic of lags. Clockwise loops (e.g., Acciari et al. 2009; Kapanadze et al. 2016) in the HR–flux 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.
|
[
"Giebels et al. 2007"
] |
[
"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"
] |
[
"Background",
"Background"
] |
[
[
814,
833
]
] |
[
[
552,
646
],
[
673,
796
]
] |
2022AandARv..30....6M__Magliocchetti_et_al._2018b_Instance_2
|
AGN While—starting 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)—virtually 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 20–30% of them inhabit rich (group- and cluster-like) structures (e.g., Best et al. 2007; Magliocchetti and Brüggen 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—under 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)—to 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 ∼60\documentclass[12pt]{minimal}
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\begin{document}$$\sim 60$$\end{document} Myr up to a few Gyr (e.g., Lin and Mohr 2007; Smolčić et al. 2011; Hatch et al. 2014; Magliocchetti et al. 2017). We also note that recent studies based on direct LOFAR observations which—we remind—sample 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 (τ>200\documentclass[12pt]{minimal}
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\begin{document}$$\tau > 200$$\end{document} Myr – Heesen et al. 2018).
|
[
"Magliocchetti et al. 2018b"
] |
[
"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—under 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)"
] |
[
"Compare/Contrast",
"Compare/Contrast"
] |
[
[
1773,
1799
]
] |
[
[
1451,
1750
],
[
1838,
1899
]
] |
2016ApJ...817...12P__Kleeorin_et_al._2000_Instance_1
|
Large-scale magnetic fields with strength of the order of 1–10 μG 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 α-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 (α-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 α-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.
|
[
"Kleeorin et al. 2000",
"Kleeorin et al. 2000"
] |
[
"To avoid catastrophic suppression of the dynamo action (α-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"
] |
[
"Background",
"Background",
"Background"
] |
[
[
946,
966
],
[
1432,
1452
]
] |
[
[
752,
915
],
[
969,
1048
],
[
1404,
1430
]
] |
2020ApJ...896L..21R___2015_Instance_1
|
Top left: the revised, CO(J = 1 → 0)-based Mgas from VLASPECS confirm that z = 2–3 galaxies detected in the ASPECS survey (green circles; tentative detections are marked with a plus sign) closely follow the “star formation law” (i.e., Mgas–SFR relation) at high redshift. CO-detected main-sequence galaxies at similar redshifts from the PHIBBS1/2 surveys (typically based on CO J = 3 → 2, but using a metallicity-dependent conversion factor; Tacconi et al. 2018) and local galaxies from the xCOLD GASS CO(J = 1 → 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 ∼ 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; Gómez-Guijarro et al. 2019), but ∼2 times higher on average than BzK-selected main-sequence galaxies at z ∼ 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–dotted 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.
|
[
"Sharon et al.",
"2015"
] |
[
"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"
] |
[
"Similarities",
"Similarities"
] |
[
[
1291,
1304
],
[
1311,
1315
]
] |
[
[
780,
886
],
[
1058,
1174
]
] |
2021MNRAS.502.5038N__Gosset_et_al._1994_Instance_1
|
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 40. However, at the corresponding frequency of 138 d−1, 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 d−1 signal in WR 66 (Antokhin et al. 1995), but with strong daily aliasing. It was subsequently confirmed in an independent, less aliased dataset (largest peak at 5.815 d−1: 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 78 but considered as a transient event, as it was observed only during one night. Because WR 78 was not observed with a 2-min cadence, we cannot check for the presence of such a frequency in the TESS data.
|
[
"Gosset et al. 1994"
] |
[
"However, at the corresponding frequency of 138 d−1, the TESS high-cadence data (sector 10) show no signs of such a signal, confirming previous negative reports"
] |
[
"Similarities"
] |
[
[
603,
621
]
] |
[
[
442,
601
]
] |
2018MNRAS.475.1160H__Werk_et_al._2013_Instance_1
|
Galaxies are surrounded by vast gaseous haloes which extend well beyond the hosts’ 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 & Boissé 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 ∼ 104 − 5 K (clumpy in nature) embedded within a hot diffuse T ∼ 106 K 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 i, Si ii, Si iii, C ii, C iv), the hot medium is home for the most highly ionized species (e.g. O vi, O vii). 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).
|
[
"Werk et al. 2013"
] |
[
"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."
] |
[
"Background"
] |
[
[
1392,
1408
]
] |
[
[
1155,
1391
]
] |
2015MNRAS.454..193B__Langhoff_et_al._1998_Instance_1
|
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 6–311G** 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. 6–311G**, 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–H out-of-plane (oop) mode, 0.972 for the C–H in-plane and C-C stretching modes and 0.965 for the C–H stretching mode. Gaussian line shapes of 30 cm− 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).
|
[
"Langhoff et al. 1998"
] |
[
"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"
] |
[
"Background"
] |
[
[
610,
630
]
] |
[
[
335,
537
]
] |
2015MNRAS.450.2404G__Fornasa_&_Sánchez-Conde_2015_Instance_1
|
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 γ-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 & Sánchez-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).
|
[
"Fornasa & Sánchez-Conde 2015"
] |
[
"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."
] |
[
"Background"
] |
[
[
1096,
1124
]
] |
[
[
877,
1063
]
] |
2017MNRAS.468.2590S__Fassnacht_et_al._2002_Instance_2
|
We initiated the H0LiCOW (H0 Lenses in COSMOGRAIL's Wellspring) program with the aim of measuring the Hubble constant with better than 3.5 per cent 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 ∼3 per cent accuracy (Fassnacht et al. 2002; Courbin et al. 2011; Tewes et al. 2013b). In addition, we obtain HST imaging that reveal the ‘Einstein ring’ 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.
|
[
"Fassnacht et al. 2002"
] |
[
"The unprecedented quality of the light curves combined with new curve-shifting algorithms",
"lead to time delays with typically ∼3 per cent accuracy"
] |
[
"Uses",
"Uses"
] |
[
[
996,
1017
]
] |
[
[
817,
906
],
[
939,
994
]
] |
2017AandA...599A..13Y__Wimmer-Schweingruber_et_al._1997_Instance_1
|
Figure 2 shows solar wind plasma and magnetic field measurements for a CIR that occurred between July 26 and 27, 2003 (days of year 207–208). 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′), the compressed fast wind region (F′), 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′-F′) 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′) and trailing edge (F′-F) of the CIR were determined by the total pressure (Jian et al. 2006). Bučík 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 20−30 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+B2/2μ0, 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′ from F). We clearly see that the suprathermal He++ intensity peaks inside the decelerated and compressed fast-wind region (F′), close to the reverse shock. In contrast, suprathermal particles are very rare in the S and S′ 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′ regions is higher than 100, confirming that our observations are due to real He++ particles.
|
[
"Wimmer-Schweingruber et al. 1997"
] |
[
"The stream interface (S′-F′) 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"
] |
[
"Uses"
] |
[
[
712,
744
]
] |
[
[
533,
710
]
] |
2019AandA...631A..88Y__Bohren_&_Huffman_(1998)_Instance_3
|
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 Köhler 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 = ⟨cosθ⟩. 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 μm grain and to the mid-infrared(MIR) for a 10 μm 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).
|
[
"Bohren & Huffman (1998)"
] |
[
"For grains surrounded by an ice mantle, the optical properties are derived with the core-mantle Mie theory using the BHCOAT routine given in"
] |
[
"Uses"
] |
[
[
3177,
3200
]
] |
[
[
3036,
3176
]
] |
2021MNRAS.504.3316B__than_2000_Instance_2
|
WASP-43b is the most heavily scrutinized phase curve, with four analyses of this data set already published (Stevenson et al. 2017; Mendonça et al. 2018; Morello et al. 2019; May & Stevenson 2020). Our phase curve semi-amplitude, eclipse depth, and radius are consistent with all of these works. The more contentious issue is that of the phase curve’s phase offset and nightside temperature. Stevenson et al. (2017) initially reported only a 2σ upper limit on the nightside temperature of 650 K, while all subsequent reanalyses (including ours) favour a significantly detectable nightside temperature of ∼800 K. As for the planet’s phase offset, Stevenson et al. (2017) and May & Stevenson (2020) favour a larger phase offset (21 ± 2 °E) than Mendonça et al. (2018) and Morello et al. (2019) (12 ± 3 °E and 11 ± 2 °E). May & Stevenson (2020) claimed that the differences between the retrieved phase offsets is the result of temporal binning which was not used by Stevenson et al. (2017) and May & Stevenson (2020) but was used by Mendonça et al. (2018), Morello et al. (2019), and this work. Fitting the temporally binned photometry for all 17 phase curves with each of our detector models already required more than 2000 CPU hours, and expanding this to unbinned photometry for all phase curve fits would require more than 125 000 CPU hours (or 434 d using our 12× multithreading computer) optimistically assuming all of detector models scaled linearly with the number of input data. However, we did try fitting just the WASP-43b unbinned phase curve with our preferred detector model (BLISS) and found that our phase offset and nightside temperature was unchanged. Including a linear slope in time also did not affect our phase offset or nightside temperature. Instead, we find that the phase offset inferred by our models depends on the choice of phase curve model, as our 4-parameter (v2) phase curve models are consistent with those of Stevenson et al. (2017) and May & Stevenson (2020), while our 2-parameter phase curve models (v1) are consistent with Mendonça et al. (2018) and Morello et al. (2019). Ultimately, we cannot decide between these two discrepant offsets as the ΔBIC between the two phase curve models for our preferred BLISS detector model is only 3.7 (insignificantly favouring the 20.4 ± 3.6 offset from the v2 model). For reference, Stevenson et al. (2014b) found phase offsets ranging from roughly −6 to 17 deg east in the Hubble/WFC3 bandpass.
|
[
"Stevenson et al. (2017)"
] |
[
"The more contentious issue is that of the phase curve’s phase offset and nightside temperature.",
"initially reported only a 2σ upper limit on the nightside temperature of 650 K, while all subsequent reanalyses (including ours) favour a significantly detectable nightside temperature of ∼800 K."
] |
[
"Differences",
"Differences"
] |
[
[
392,
415
]
] |
[
[
296,
391
],
[
416,
611
]
] |
2022MNRAS.515.3299G__Tagawa_et_al._2020a_Instance_1
|
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-Escudé & Gould 2000; Freitag, Amaro-Seoane & Kalogera 2006; Hopman & Alexander 2006; O’Leary 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 ∼20 to ∼300 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).
|
[
"Tagawa et al. 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",
"then get captured in the disc by hydrodynamic drag as they cross the disc (e.g."
] |
[
"Background",
"Background"
] |
[
[
759,
778
]
] |
[
[
281,
462
],
[
607,
686
]
] |
2019ApJ...873...89M__Bourrier_et_al._2016_Instance_1
|
Ultraviolet observations of the hot Jupiter HD 209458 b have found up to a 15% occultation in the wings of hydrogen Lyα, with effective Doppler shifts of up to ±150 km s−1, 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
)5
5
Along with a nondetection for a super Earth, 55 Cnc e, placing an upper limit on its mass loss.
, 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
6
Redshifted occultation of (0.7 ± 3.6)% pretransit and (8.0 ± 3.1)% posttransit. Blueshifted occultation of (17.6 ± 5.2)% pretransit and (47.2 ± 4.1)% posttransit (Ehrenreich et al. 2015).
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 Hα absorption has been seen in HD 189733 b’s 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 Oklopčić & Hirata (2018), provides an opportunity for ground-based observations.
|
[
"Bourrier et al. 2016"
] |
[
"Similar outflows have been reported for",
"and for one hot Neptunian planet, GJ 435 b"
] |
[
"Similarities",
"Similarities"
] |
[
[
835,
855
]
] |
[
[
463,
502
],
[
748,
790
]
] |
2016MNRAS.463.3637R__Mirabel,_Dottori_&_Lutz_1992_Instance_1
|
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.
|
[
"Mirabel, Dottori & Lutz 1992"
] |
[
"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"
] |
[
"Motivation",
"Motivation"
] |
[
[
758,
786
]
] |
[
[
656,
756
],
[
789,
847
]
] |
2015ApJ...805..115V__Robitaille_et_al._2006_Instance_1
|
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., André et al. 1993; Chen et al. 1995). Classifications relying upon physical properties are usually referred in the literature as “stages”, whereas those using observational signatures are called “classes”. For simplicity here we use the term “class” 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
g cm−2 and the velocity field in the infalling envelope. Varying the value of
by a factor of 5 results in changes of the estimated onset time of different phases by only a few per cent.
|
[
"Robitaille et al. 2006"
] |
[
"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"
] |
[
"Uses",
"Differences"
] |
[
[
592,
614
]
] |
[
[
459,
590
],
[
643,
744
]
] |
2017MNRAS.464..183N__Iannuzzi_&_Dolag_2012_Instance_1
|
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 ∼ 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 ∼ 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 ∼3 Gyr, part of these infalling galaxies may reach the R 2R200 region, at z ∼ 0.4, mixing with virialized and backsplash objects, and thus presenting a less pronounced radial segregation between red and blue galaxies.
|
[
"Iannuzzi & Dolag 2012"
] |
[
"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."
] |
[
"Uses"
] |
[
[
854,
875
]
] |
[
[
609,
829
]
] |
2015ApJ...807...91C__Moffat_1969_Instance_1
|
The standard photometric analysis (see Dalessandro et al. 2008a, 2008b) has been performed on the “flc” images, which are corrected for flat field, bias, dark counts, and charge transfer efficiency. These images have been further corrected for “Pixel-Area-Map”10
10
For more details see the ACS Data Handbook.
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 ∼200 bright but not saturated stars. The model has been chosen on the basis of a χ2 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 3σ detection limit, where σ 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
11
We used CataXcorr, a code that is specifically developed to perform accurate astrometric solutions. It has been developed by P. Montegriffo at INAF–Osservatorio 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.
our catalog with that by Anderson et al. (2008), using the ∼7600 stars in common.
|
[
"Moffat 1969"
] |
[
"The model has been chosen on the basis of a χ2 test and, in every image, the best fit is provided by a Moffat function"
] |
[
"Uses"
] |
[
[
709,
720
]
] |
[
[
589,
707
]
] |
2018ApJ...864...51W__Lee_et_al._2014_Instance_1
|
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; Kereš 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.
|
[
"Lee et al. 2014"
] |
[
"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."
] |
[
"Background",
"Background"
] |
[
[
521,
536
]
] |
[
[
142,
314
],
[
807,
1177
]
] |
2020MNRAS.493.2373D__Haghi_&_Amiri_2016_Instance_1
|
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–Zeldovich 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 ∼100 more diffuse than ultra diffuse galaxies. It has a very wide core ∼3 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.
|
[
"Haghi & Amiri 2016"
] |
[
"Nevertheless, a recent study of the dynamics of dwarf spheroidals orbiting around the Milky Way has shown some inconsistencies with previous results"
] |
[
"Differences"
] |
[
[
909,
927
]
] |
[
[
759,
907
]
] |
2015MNRAS.451.2544P__Mesinger,_Furlanetto_&_Cen_2011_Instance_1
|
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 per cent. At z ∼ 1, no objects with leaking ionizing photons have been detected (Bridge et al. 2010; Siana et al. 2010), but at z ∼ 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 ∼10 per cent of the observed objects (Nestor et al. 2013). Attempts to constrain the escape fraction with numerical simulations find ranges between fesc 10 per cent (Razoumov & Sommer-Larsen 2006, 2007; Gnedin, Kravtsov & Chen 2008; Paardekooper et al. 2011; Kim et al. 2013) and fesc > 80 per cent (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.
|
[
"Mesinger, Furlanetto & Cen 2011"
] |
[
"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."
] |
[
"Background",
"Background"
] |
[
[
810,
841
]
] |
[
[
508,
627
],
[
679,
724
]
] |
2021MNRAS.500.1817L__Abbott_et_al._2020b_Instance_1
|
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–time 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’s 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–time volumes $VT=1.2/\bar{\mathcal {R}}$ for each of these three events (‘1.2’ 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
(1)$$\begin{eqnarray*}
{\mathrm{d}P\over \mathrm{d}\mathcal {R}_{\rm bns}} = \int _0^{\mathcal {R}_{\rm bns}} \mathrm{d}\mathcal {R}_1 {\mathrm{d}P\over \mathrm{d}\mathcal {R}_1} \left.{\mathrm{d}P\over \mathrm{d}\mathcal {R}_2}\right|_{\mathcal {R}_{\rm bns}-\mathcal {R}_1},
\end{eqnarray*}$$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
(2)$$\begin{eqnarray*}
{\mathrm{d}P\over \mathrm{d}\beta } = \int _0^\infty {\mathrm{d}\mathcal {R}_{3}\over \mathcal {R}_3} {\mathrm{d}P\over \mathrm{d}\mathcal {R}_{3}} \left.{\mathrm{d}P\over \mathrm{d}\mathcal {R}_{\rm bns}^{-1}}\right|_{\beta /\mathcal {R}_3},
\end{eqnarray*}$$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 $0.064\, \rm {per\, cent}\lt \beta \lt 2.8\, \rm {per\, cent}$.
|
[
"Abbott et al. 2020",
"b"
] |
[
"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–time 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’s prior"
] |
[
"Uses",
"Uses"
] |
[
[
82,
100
],
[
102,
103
]
] |
[
[
0,
80
],
[
106,
523
]
] |
2016AandA...591A..91P__Halpern_et_al._2014_Instance_2
|
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 ≳25 M⊙; Filipović et al. 2005; Reynoso et al. 2006; García et al. 2012; Kamitsukasa et al. 2015; Auchettl et al. 2015a). Following these authors, the distance to the SNR is in the range 6−11 kpc; the most recently determined values converge towards 7 ± 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 Ṗ = (8.56 ± 0.51) × 10-15 s s-1. The estimated spin-down energy is Ė = 1.36 × 1036 erg s-1 and the surface dipolar magnetic field is 7.4 × 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′ 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″ 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 3.7σ 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).
|
[
"Halpern et al. 2014"
] |
[
"The estimated spin-down energy is Ė = 1.36 × 1036 erg s-1 and the surface dipolar magnetic field is 7.4 × 1011 G"
] |
[
"Background"
] |
[
[
991,
1010
]
] |
[
[
877,
989
]
] |
2015ApJ...804..130C__Bertschinger_1985_Instance_3
|
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’s gravity. The embedded lens theory is based on the Swiss cheese cosmologies (Einstein & Straus 1945; Schücking 1954; Kantowski 1969). The idea of embedding (or Swiss cheese) is to remove a co-moving sphere of homogeneous dust from the background Friedmann–Lemaître–Robertson–Walker (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 “compensated” by an overdense bounding ridge, i.e., a compensated void (Sato & Maeda 1983; Bertschinger 1985; Thompson & Vishniac 1987; Martínez-González 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 “uncompensated” or “undercompensated” void (Fillmore & Goldreich 1984; Bertschinger 1985; Sheth & van de Weygaert 2004; Das & Spergel 2009).3
3
This 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 “largely” 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.
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’s equations. The critical difference between an embedded lens and a traditional lens lies in the fact that embedding effectively reduces the gravitational potential’s 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.
|
[
"Bertschinger 1985"
] |
[
"The initial perturbation can be compensated or uncompensated, which leads to different void growth scenarios"
] |
[
"Background"
] |
[
[
2188,
2205
]
] |
[
[
2078,
2186
]
] |
2015ApJ...807..101K__Sui_et_al._2005_Instance_1
|
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
2
where ELC is the low-energy cutoff, Fe is the total number of electrons per second above ELC in units of 1035 electrons s−1, and δ 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 20–32 keV with relatively higher values (∼32 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.
|
[
"Sui et al. 2005"
] |
[
"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"
] |
[
"Motivation"
] |
[
[
848,
863
]
] |
[
[
666,
846
]
] |
2015AandA...583A..77L__Busso_et_al._2001_Instance_1
|
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 “pocket”); (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 −1.15 to −1.34), 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 1–1.5 M⊙. 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] ≃ 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.
|
[
"Busso et al. 2001"
] |
[
"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."
] |
[
"Background",
"Background"
] |
[
[
598,
615
]
] |
[
[
446,
597
],
[
618,
777
]
] |
2020AandA...643A.148G__Martin_et_al._1998_Instance_1
|
Figure 11 shows that the distribution of the spectral index α (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⊙ is of the order of 1 Myr which could explain the early disapperance of circumstellar discs for some stars in our sample.
|
[
"Martin et al. 1998"
] |
[
"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."
] |
[
"Background",
"Differences"
] |
[
[
827,
845
]
] |
[
[
561,
806
],
[
848,
1020
]
] |
2016MNRAS.457.2569M__Binney_&_Piffl_2015_Instance_1
|
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’ 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–Planck 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.
|
[
"Binney & Piffl 2015"
] |
[
"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."
] |
[
"Uses"
] |
[
[
875,
894
]
] |
[
[
588,
874
]
] |
2019AandA...631A..35B__Bridges_et_al._(1996)_Instance_4
|
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., Ramírez et al. 1999; Zhang & Vu-Quoc 2002). However, the experiments by Heißelmann 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 (≤ 1 cm s−1). 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 Heißelmann 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).
|
[
"Bridges et al. (1996)"
] |
[
"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)."
] |
[
"Compare/Contrast",
"Compare/Contrast"
] |
[
[
1639,
1660
]
] |
[
[
1573,
1638
],
[
1661,
2065
]
] |
2022AandA...667A.131B__Izumi_et_al._(2016)_Instance_1
|
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; García-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 “submillimeter-HCN diagram”, 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(4−3)/HCO+(4−3) and HCN(4−3)/CS(7−6), 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 μm infrared photons due to the presence of hot dust around AGN. These photons vibrationally excite HCN to the ν2 = 1 state. Upon de-exciting from this state back to the vibrational ground state, ν = 0, 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 μm infrared photons can similarly vibrationally excite HCO+, thus nullifying the extent of this effect (Imanishi et al. 2016).
|
[
"Izumi et al. (2016)"
] |
[
"In particular, the “submillimeter-HCN diagram”, 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(4−3)/HCO+(4−3) and HCN(4−3)/CS(7−6), where all of the molecules involved are considered tracers of dense gas."
] |
[
"Background",
"Background"
] |
[
[
883,
902
]
] |
[
[
770,
882
],
[
904,
1150
]
] |
2018MNRAS.475.3419A__Davis_et_al._1999_Instance_1
|
If we consider for the bulk density the value 4500 kg m−3, 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 km (Shepard et al. 2017), the Psyche parent body (PPB) was supposed to be ∼500 km in diameter and have suffered severe ‘hit-and-run’ impact events capable of removing all crust and mantle, exposing the core (Elkins-Tanton et al. 2016). In addition, Psyche should have ∼40 per cent macroporosity, if we assume that it is made of blocks of iron/nickel with a density around 7500 kg m−3. 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 ≤50 per 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 Myr 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 km and above the detection limit. There is a lack of primordial asteroid families in the Main Belt (Brož et al. 2013; Spoto, Milani & Knežević 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’ 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).
|
[
"Davis et al. 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 Myr 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 km and above the detection limit."
] |
[
"Compare/Contrast",
"Compare/Contrast"
] |
[
[
1531,
1548
]
] |
[
[
1365,
1529
],
[
1551,
1815
]
] |
2019ApJ...886...34F__Bonal_et_al._2010_Instance_1
|
Here, we report Li–Be–B and Al–Mg 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–CB) chondrites and their components are characterized by enrichments in 15N (up to 1,100‰; Murty et al. 2007; Ivanova et al. 2008; Briani et al. 2009; Bonal et al. 2010), which is similar to the characteristics of comets (Füri & Marty 2015). In addition, Van Kooten et al. (2016) and Olsen et al. (2016) proposed that Mg and Cr isotopic compositions of bulk CH–CB chondrites require significant amounts (20%–50%) of primordial molecular cloud matter in their precursor material. Hence, CH–CB chondrites may have accreted a significant amount of outer solar system materials. Most CAIs in CH–CB 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–CB 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–CB 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.
|
[
"Bonal et al. 2010"
] |
[
"The CH and CH/CB (hereafter: CH–CB) chondrites and their components are characterized by enrichments in 15N (up to 1,100‰;"
] |
[
"Background"
] |
[
[
488,
505
]
] |
[
[
305,
427
]
] |
2021AandA...655A..99D__Carigi_et_al._2005_Instance_1
|
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 I 6158 Å 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 I ] 6300 Å 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 Å thin-disk stars with [Fe/H] ≲ –0.2 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.
|
[
"Carigi et al. 2005"
] |
[
"Once metallicity starts to increase, low- and intermediate-mass stars release carbon and massive stars start to eject more carbon than oxygen"
] |
[
"Background"
] |
[
[
1301,
1319
]
] |
[
[
1158,
1299
]
] |
2019ApJ...875L..31H__Leary_et_al._2006_Instance_2
|
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’Leary 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’Leary et al. 2009; Cholis et al. 2016; Gondán 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’Leary et al. 2006; Breivik et al. 2016; Nishizawa et al. 2016; Chen & Amaro-Seoane 2017; Nishizawa et al. 2017; D’Orazio & Samsing 2018; Kremer et al. 2019; Samsing & D’Orazio 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.
|
[
"O’Leary et al. 2006"
] |
[
"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.,"
] |
[
"Compare/Contrast"
] |
[
[
1971,
1990
]
] |
[
[
1750,
1970
]
] |
2022AandA...659A...5Y__Mills_et_al._2018_Instance_2
|
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; Galván-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.
|
[
"Mills et al. 2018"
] |
[
"These have led to the discovery of metastable maser lines in 22 different regions"
] |
[
"Background"
] |
[
[
1445,
1462
]
] |
[
[
933,
1014
]
] |
2016ApJ...832...52F__Harris_et_al._2012_Instance_1
|
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 μm (
and S500 S350; i.e., “350 μm peakers”), (2) S500 > 20 mJy, and (3) >3σ 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 μm if redshifted to z ∼ 2.5. This is confirmed by the blind carbon monoxide (CO J = 1–0) survey of a subsample of the brightest 350 μm peakers (S350 ≥ 115 mJy), which has shown a strikingly similar redshift distribution as 850 μm selected SMGs (zCO = 2.5 ± 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–Jeans 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 (
). 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 σ500 = 15 mJy beam−1 for the large HeLMS and HerS fields (confusion noise included; Oliver et al. 2012; Viero et al. 2014) to confusion limited with σ500 = 6.8 mJy beam−1 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 μm flux density ensure that ∼95% of our sample would be classified as SMGs if they were observed at 870 μm. 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 μm 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−2 is five times lower than the observed 870 μm source count above S870 ≳ 3 mJy (∼500 deg−2; Coppin et al. 2006; Weiß et al. 2009). This is not surprising given that almost half of the total Herschel area is only 10%–20% complete at S500 = 20 mJy. Note that this incompleteness in the Herschel catalogs is not a concern for compiling a sample of SMG−QSO pairs.
|
[
"Harris et al. 2012"
] |
[
"This is confirmed by the blind carbon monoxide (CO J = 1–0) survey of a subsample of the brightest 350 μm peakers (S350 ≥ 115 mJy), which has shown a strikingly similar redshift distribution as 850 μm selected SMGs (zCO = 2.5 ± 0.8;"
] |
[
"Similarities"
] |
[
[
808,
826
]
] |
[
[
575,
807
]
] |
2019ApJ...885...81S__Grogin_et_al._2011_Instance_1
|
We measured the apparent axial ratio on the IF814W-band data, which correspond to the rest-frame V band for galaxies at z ∼ 0.4 and the rest-frame B band for those at z ∼ 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 − (b/a)F814W are −0.006 and −0.007 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.
|
[
"Grogin et al. 2011"
] |
[
"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"
] |
[
"Uses"
] |
[
[
727,
745
]
] |
[
[
541,
725
]
] |
2021ApJ...912..163B__Kleine_et_al._2020_Instance_1
|
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’s lifetime. One key observation supporting this dichotomy is the ε94Mo and ε95Mo values (where ε 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 ε48Ca, ε50Ti, ε54Cr, ε62Ni, ε100Ru, and △17O values, where △ 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).
|
[
"Kleine et al. 2020"
] |
[
"For instance, a number of measurements have revealed that the stable isotopic compositions of all meteorites fall into two distinct families"
] |
[
"Background"
] |
[
[
1354,
1372
]
] |
[
[
1159,
1299
]
] |
2016MNRAS.457.2433P__Nolan_et_al._2012_Instance_2
|
From the result of the χ2 minimization, we found that the minimized χ2 values agree with the expected values, i.e. the computed χ2 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 χ2 ≈ 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 γ-ray source with a photon index of ≈1.77 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 χ2 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 χ2 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 ≈1.5.
|
[
"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"
] |
[
"Compare/Contrast"
] |
[
[
987,
1004
]
] |
[
[
830,
985
]
] |
2016MNRAS.463.2716M__Cho_&_Lazarian_2007_Instance_1
|
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 λ (e.g. Cho & Lazarian 2007; for λ = 1.25 mm, amax = 0.2 mm), 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 → 0 as the grain axis ratio → 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 mm and 10 per cent at 10 mm, with a peak fractional polarization of ∼20 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 ∼10 lower value of P measured in this source at 1.25 mm8 in terms of a dilution of the polarized emission from a ≲ 0.1 mm 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 ≳ 1 mm.9 Although a fraction of these grains may have sizes in excess of 1 mm and would therefore emit less efficiently at that wavelength than a ≲ 1 mm grains (e.g. Miyake & Nakagawa 1993), the mid-plane region should still dominate the total mm-wavelength flux if most of the a ≳ 1 mm grains are concentrated there. Grains of size a ≲ 0.1 mm 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 ( ≳ 1 cm) 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 Λ0 ≪ 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.
|
[
"Cho & Lazarian 2007"
] |
[
"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 λ (e.g.",
"; for λ = 1.25 mm, amax = 0.2 mm)"
] |
[
"Uses",
"Uses"
] |
[
[
1082,
1101
]
] |
[
[
786,
1081
],
[
1101,
1134
]
] |
2015MNRAS.454.2003P__Paardekooper_&_Papaloizou_2009_Instance_1
|
In both of these approaches, disc–planet 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).
|
[
"Paardekooper & Papaloizou 2009"
] |
[
"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"
] |
[
"Background"
] |
[
[
1604,
1634
]
] |
[
[
1321,
1572
]
] |
2021AandA...651A..87O__Brunthaler_et_al._2021_Instance_2
|
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″ and 1.5″ 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 1σ values in the range from 0.056 to 0.43 mJy beam−1 for B-configuration images. For the D configuration, the 1σ rms noise ranges from 0.10 to 2.6 mJy beam−1. 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.
|
[
"Brunthaler et al. 2021"
] |
[
"Continuum and methanol line maps from Effelsberg observations have also been obtained as part of the GLOSTAR survey"
] |
[
"Uses"
] |
[
[
1004,
1026
]
] |
[
[
887,
1002
]
] |
2021AandA...652A..98G__Buta_et_al._(2009)_Instance_1
|
On the CND scales displayed in Fig. 12, the CO(3–2) emission in this highly inclined (i = 59°; Appendix D) barred Seyfert 2 galaxy is detected at every single position inside r 200 pc. However, a sizeable fraction of the molecular gas appears concentrated in a nuclear ring of ∼200 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 ∼140″ (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 50 pc), molecular gas probed by CO is concentrated in an asymmetric ringed disk of ∼30 − 40 pc radius located around the AGN and oriented along ∼160°. The molecular disk shows a similar orientation to the extended component of the 351 GHz continuum emission detected by ALMA (PAGauss ∼ 160° −162°; 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 − H 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 ≤ 7 − 9 pc] ∼ (1.5 − 2.9)×105 M⊙ and N(H2)∼(3.5 − 3.6)×1022 cm−2, respectively. The galaxy shows a bright ionization cone southwest of the nucleus oriented along PAout ∼ 235° −245° 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 − 90 pc-size disk detected by ALMA in continuum and CO is therefore equatorial (torus-like).
|
[
"Buta et al. (2009)"
] |
[
"The bar, detected in the NIR by Quillen et al. (1997) and",
"shows a prominent boxy-shape morphology."
] |
[
"Background",
"Background"
] |
[
[
630,
648
]
] |
[
[
572,
629
],
[
650,
690
]
] |
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