adsabs/FOCAL huggingface dataset:
In spite of these similarities, the angular profiles of these properties in two-dimensional simulations reveal a totally different accretion pattern. Within the stagnation radius, winds are mainly accreted in the polar region as a result of the low angular momentum there, whereas material with high angular momentum in the mid-plane outflow is subject to outflow, because the angular momentum of the stellar winds at a given radius decreases towards the pole. This ‘funnel’ accretion scenario is also found in other simulations with a spherical-like distribution of rotationally accreting material (Proga & Begelman 2003; Ressler et al. 2018). The density and temperature at the equatorial plane are higher than those in the polar region, and thus the highest mass accretion rate is found at the boundary of the polar and disc regions. Therefore the accretion pattern of the two-dimensional simulations is totally different from the inflow–outflow structure in the one-dimensional calculations. Another difference between our results is that the parameters of our best-fitting model, fq and vw, sn, are generally smaller than those in one-dimensional calculations. The smaller fq can be explained in three ways. First, our BH mass is double that in the one-dimensional case, which results in a deeper gravitational potential that allows for a more massive gas reservoir. Thus fq should be smaller to achieve the observed gas density. Second, our outer boundary is set at 17 arcsec, which is larger than the radius of the one-dimensional outer boundary of 12 arcsec. The larger outer boundary causes the mass injection rate by stellar winds to increase by 50 per cent. Because stellar winds are the only mass source in our models, a smaller fq is required to cancel out this extra mass injection. Third, a smaller vw, sn could also lead to a smaller fq owing to less efficient heating by supernova explosions. The smaller vw, sn is constrained by the temperature at large radii, especially the temperature between 10 and 20 arcsec, which is not fitted in the one-dimensional calculations.
Citation: Proga & Begelman 2003
Citation function: Similarities
The radial distribution of S/H ratios and the corresponding gradient are shown in panel (b) of Fig. 12. We obtain a slope of −0.035 ± 0.006 dex kpc−1 (very similar to the one we obtain with the ICF of ADIS20, as can be seen in Table 8), which is consistent with the one of −0.041 ± 0.014 dex kpc−1 estimated by Rudolph et al. (2006) using FIR lines and also very similar to the slope of our O/H gradient. We report a dispersion around the S/H gradient of 0.10 dex, somewhat larger than the individual observational uncertainties. Recently, Fernández-Martín et al. (2017) reported a slope of −0.108 ± 0.006 dex kpc−1 using optical spectra for H  ii regions located at RG between 5 and 17 kpc. That value of the slope is considerably much steeper than our determination and other previous estimates from the literature (e.g. Shaver et al. 1983; Simpson et al. 1995; Afflerbach, Churchwell & Werner 1997; Rudolph et al. 2006). Esteban & García-Rojas (2018) observed some H ii regions with a very low ionization degree – O2 +/O  0.03 – but measurable [S iii] 6312 or 9069 Å lines. For such nebulae is possible to assume that S/H ≈ S+/H+ + S2 +/H+, since the contribution of S3 +/H+ in those objects is expected to be negligible. The objects with such properties are IC 5146, Sh 2-235, Sh 2-257, and Sh 2-271, whose total abundances of S/H range from 6.70 to 7.01 but, unfortunately, cover a rather narrow range of RG – between 9.3 and 11.7 kpc – and no confident gradient can be estimated with so small baseline. Inspecting Tables 3 and 4 we can see that there is a sizable group of low-ionization degree H ii regions that lack of determination of their S2 +/H+ ratios. This is because the rather faint auroral [S iii] 6312 Å line could not be detected in those objects. We plan to obtain additional optical spectra covering the bright nebular [S iii] 9069, 9532 Å lines of that group of nebulae for trying to increase the number of objects with S/H ratios determined without ICF and estimate a more precise gradient for this element.
Citation: Rudolph et al. 2006
Citation function: Differences
Our simulation results show that the median detectable distances of targeted GW events from BH–NS mergers for a single 2nd generation GW detector and a network of such detectors are ∼300 Mpc and ∼700 Mpc, respectively (see Table 4). For comparison, Figure 12 shows that the detection rate and detectable distance for HLV (O3) are approximately the same as those for the case when only bKAGRA is running. This is basically consistent with the detection rate and the distance distribution of BH–NS merger candidates detected during LVC O3 (e.g., Anand et al. 2020; Antier et al. 2020b; Coughlin et al. 2020b; Gompertz et al. 2020; Kasliwal et al. 2020). In Section 3, we have shown that the kilonova absolute magnitude at 0.5 days after a BH–NS merger is mainly distributed in the range ∼–10 to ∼–15.5. In view of the fact that the limiting magnitude of the follow-up wide-field survey projects is almost ≲21 mag (e.g., Antier et al. 2020b; Gompertz et al. 2020; Coughlin et al. 2020b; Kasliwal et al. 2020; Wyatt et al. 2020), the maximum detectable distance for BH–NS kilonovae would be ≲200 Mpc, which can hardly cover the horizon of GW-triggered BH–NS merger events that O3 found (as shown in Figure 10). However, although BH–NS merger kilonovae can hardly be detected for the present search depths, Figures 9 and 11 reveal that there are great opportunities to discover on-axis afterglows associated with sGRBs or orphan afterglows if the BH components have a high-spin distribution. In order to cover the distance range for searching for BH–NS kilonovae for the network of 2nd generation GW detectors as completely as possible, a search limiting magnitude mlimit ≳ 23–24 is required as shown in Figure 10. Present survey projects could reach this search limiting magnitude by increasing exposure times and the number of simultaneous exposures. However, the GW candidates during O3 had very large localization areas with an average of thousands of square degrees (Antier et al. 2020b). Increasing exposure times makes it hard for the present survey projects to cover such large localization areas. Therefore, during the HLVK era, we recommend that survey projects may search for jet afterglows after GW triggers with a relatively shallow search limiting magnitude. If BH–NS mergers have a high location precision, a limiting magnitude of mlimit ≳ 23–24 can be reached, which gives a higher probability of discovering associated kilonovae.
Citation: Coughlin et al. 2020b
Citation function: Similarities
Citation: Coughlin et al. 2020b
Citation function: Uses; Uses
To fully describe a population of streams, we need a realistic model of the Galaxy. The latest generation of hydrodynamical simulations have produced models that well reproduce a multitude of features observed in galaxies, both as individual objects (e.g., Wetzel et al. 2016) and as a population (e.g., Nelson et al. 2018; Pillepich et al. 2018). Part of their success is in the achieved high resolution, so at the present a Milky-Way-like galaxy is modeled by up to 140 million particles (Wetzel et al. 2016). Ultimately, we would like to have a description of a galaxy that is representative of these models, but at a fraction of the numerical cost. In the classic study, Hernquist & Ostriker (1992) developed a set of basis functions for density and gravitational potential that can reproduce complex morphology of galaxies with a small number of terms (e.g., Lowing et al. 2011; Lilley et al. 2018a, 2018b). These expansions reproduce the force field of an N-body simulation with a precision of a few percent, so representing the gravitational potential in our model with an expansion of basis functions is tempting. However, a truly realistic solution needs to accurately capture not only the current structure of the Galaxy but also its evolution in time. Even though the Milky Way has had a relatively quiet recent merger history, it is currently undergoing a major merger with the Large Magellanic Cloud (e.g., Besla et al. 2007; Peñarrubia et al. 2016), which has been a source of gravitational perturbation for at least a billion years—sufficiently long to affect stellar streams. Specifically, Law & Majewski (2010) showed that the only static, ellipsoidal halo that reproduces both the positions and radial velocities along the Sagittarius stream is triaxial. This model correctly predicted proper motions along the stream (Sohn et al. 2015), so it appears to be describing the effective potential well, even though it is cosmologically improbable (Debattista et al. 2013). On the other hand, modeling Sagittarius in a combined system of the Milky Way and the LMC relaxes the requirement for the dark matter halo to be triaxial (Vera-Ciro & Helmi 2013), signaling that having a model of the potential that is correct on average is no guarantee of recovering the true halo shape. To ensure that the complexities added to the model are realistic, the basis function expansion should thus be time dependent, simultaneously describing the interaction between the Milky Way and the LMC, while maintaining a thin and old disk. We relegate the development of such a model and its implementation for mapping the dark matter in the Galaxy to a future study.
Citation: Lilley et al.; 2018b
Citation function: Background; Motivation
It is still debated if and how all the stellar debris efficiently circularizes via the stream–stream collision. Some hydrodynamical simulations show that the TDE disk retains a significantly elliptical shape because the orbital energy is not dissipated efficiently enough to reduce the eccentricity of the entire disk to zero in a reasonable time (Guillochon et al. 2014; Shiokawa et al. 2015; Sądowski et al. 2016). Lu & Bonnerot (2020) show that a significant fraction of the debris can become unbound causing an outflow from the self-interaction region. Nevertheless, the debris that remains bound eventually contributes to the accretion flow around the SMBH. This part of the debris stream will finally be circularized by energy dissipation, leading to the formation of a small, initially ring-like, accretion disk around the black hole (Hayasaki et al. 2013; Bonnerot et al. 2016; Hayasaki et al. 2016). Note that, in an inefficient debris circularization case, the subsequent fallback material interacts with the outer elliptical debris so that their combined effect on the subsequent evolution of the initial ring is negligible. Angular momentum conservation allows us to estimate the circularization radius of the stellar debris, which is given by
2






r


c


=
(
1
+


e


*


)


r


p


,
=




1
+


e


*




β





r


t


,


where e* is the orbital eccentricity of the stellar orbit, rp = rt/β is the pericenter distance radius. If debris circularization takes place only through dissipation at the self-interaction shock, the circularization timescale for the non-magnetized, most tightly bound debris can be estimated based on the ballistic approximation (Bonnerot et al. 2017) as
3









t


circ




≈


8.3



η


−
1





β


−
3




M


bh
,
6


−
5

/

3





t


mtb







∼


0.93







η


1.0







−
1




β


−
3




M


bh
,
6


−
7

/

6





m


*
,
1


−
1





r


*
,
1


3

/

2



yr
,





where the orbital period of the stellar debris on the most tightly bound orbit:
4









t


mtb




=





π




2








1




Ω


*














M


bh






m


*









1

/

2







≈


0.11



M


bh
,
6


1

/

2





m


*
,
1


−
1





r


*
,
1


3

/

2



yr
,













Ω


*


=




Gm


*



/



r


*


3






 is the dynamical angular frequency of the star, and we introduce η(≤1) as the circularization efficiency which represents how efficiently the kinetic energy at the stream–stream collision is dissipated and the most efficient (η = 1) case corresponds to that of Bonnerot et al. (2017). Note that tcirc is not the circularization timescale of all the stellar debris. Our interest here is in the circularization timescale and radius of the most tightly bound debris because the accretion of this debris contributes most to the delayed X-ray peak luminosity in terms of the emitted energy.
Citation: Hayasaki et al. 2013
Citation function: Background
Some previous observations have shown that surface motions (e.g., shearing motions and converging motions) acting on preexisting coronal fields to form filaments and filament channels always act over a short period of a few days (Gaizauskas et al. 1997; Schmieder et al. 2004; Wang & Muglach 2007; Yan et al. 2015; Yang et al. 2015) to a period of months (Gaizauskas et al. 2001). However, in this event, the filament is formed within about 20 minutes. On this timescale, only the reformation of the filament in the same filament channel after the partial filament eruption has been observed (Tripathi et al. 2009a; Joshi et al. 2014). The rapid formation of the filament is a rare observation that has not been previously reported. Even though our event follows the “head-to-tail” scenario, there is a little difference. It is worth noting that the “head-to-tail” scenario seems to support the idea that the filament is supported by sheared arcades (Litvinenko & Wheatland 2005; Welsch et al. 2005). However, our event may present a picture in which the filament is supported by a twisted flux rope. Therefore, we could not fully exclude any other explanations, although we are inclined toward the “head-to-tail” scenario of Martens & Zwaan. Twisted structures in filaments have been previously observed during their eruptions (Wang et al. 1996; Bi et al. 2012, 2015; Li & Zhang 2013; Yan et al. 2014; Yang et al. 2014; Filippov et al. 2015; Raouafi 2009). In particular, Bi et al. (2012, 2015) found that the erupted filaments appear to be composed of two intertwined twisted flux ropes. However, in the present case, it is found that the flux rope appeared in the course of the formation of the filament, and as reported by Li & Zhang (2013; Yang et al. 2014), it can be observed in all seven AIA EUV lines formed from 0.05 MK to 11 MK. The AIA 304 Å observations (see panel (e) in Figure 2) also indicate that the filament may be composed of two sets of intertwined dark threadlike structures. These observations may be direct evidence that the magnetic structure of the formed filament is a flux rope. As mentioned before, the sigmoids and the elongated EUV channel-like structures are evidences for the existence of flux ropes in the higher solar atmosphere. More recently, observational evidence of the detailed structure and evolution of flux ropes in the lower solar atmosphere has been presented by Wang et al. (2015) and Kumar et al. (2015). The flux rope in our case can be observed in all of  the AIA EUV passbands, but it is still in the lower solar atmosphere. Our observation is quite similar to the formation of a chromospheric flux rope as a result of magnetic reconnection between two sheared cool Hα loops (Wang et al. 2015) in the lower atmosphere, but sheds more light on the formation of the filament. This event may also be a distinct example revealing the evolution of the flux rope in the lower solar atmosphere.
Citation: Litvinenko & Wheatland 2005
Citation function: Compare/Contrast
Solar flares are powerful eruption events on the Sun associated with a rapid and violent release of magnetic free energy through a reconnection process. A typical flare can radiate at almost all wavelengths constituting the solar spectrum, ranging from radio through optical and ultraviolet (UV) to soft/hard X-ray (SXR/HXR) and even γ-rays (e.g., Benz 2017; Tan et al. 2020). Only a small part of the flare radiation is emitted at the shortest wavelengths in the X-ray and extreme-UV (EUV) ranges (Emslie et al. 2012). The quantitative estimation of the radiated flare energy partition suggested about 70% in white light (WL) for solar flares (e.g., Kretzschmar 2011) and 55%–80% in WL for stellar flares (e.g., Kuznetsov & Kolotkov 2021). In other words, most of the flare energy is radiated in the longer wavelengths (Kleint et al. 2016). Between those extremes, the solar UV spectrum from 1000 to 3000 Å, which can be further split into the far-ultraviolet (FUV), the mid-ultraviolet (MUV), and the near-ultraviolet (NUV), is thought to provide an important contribution to the flare radiation (Woods et al. 2006; Milligan et al. 2014; Dominique et al. 2018). For instance, the Lyα spectral line produced by the chromospheric neutral hydrogen, which is centered at 1216 Å (in the FUV spectrum), is among the spectral lines in which flares radiate the most (Allred et al. 2005; Curdt et al. 2001; Lu et al. 2021a). The hydrogen Balmer continuum emitted during flares, which is thought to be generated during the recombination of flare-produced free electrons in the chromosphere, is often detected in the MUV and NUV ranges, as well as close to the Balmer recombination edge at 3646 Å (Heinzel & Kleint 2014; Kotrč et al. 2016; Dominique et al. 2018). Both the Lyα and hydrogen Balmer continuum emissions during solar flares are expected to be nonthermal profiles, i.e., similar to the HXR radiation that is produced by the beam of electrons that are accelerated by the magnetic reconnection during the solar flare (e.g., Avrett et al. 1986; Rubio da Costa et al. 2009; Heinzel & Kleint 2014).
Citation: Kuznetsov & Kolotkov 2021
Citation function: Background; Background
As described in Sect. 1, the galaxy protocluster associated with 7C 1756+6520 is characterized by a high fraction of AGN protocluster members: seven AGN, including the central radio galaxy, have been spectroscopically confirmed in close proximity both spatially and in redshift space of the protocluster. This high AGN fraction detected so far, ~23%, makes the overdensity around 7C 1756+6520 similar to the interesting and well-studied cluster around the radio galaxy PKS 1138−262 from this point of view as well (Pentericci et al. 2002; Croft et al. 2005), in addition to the extension (see Sect. 5.2). The source PKS 1138−262 is a massive forming radio galaxy at z ~ 2.16 (Pentericci et al. 1998) that is surrounded by overdensities of Lyα emitters (Pentericci et al. 2000), extremely red objects (Kurk et al. 2004a; Koyama et al. 2013), Hα emitters (Kurk et al. 2004b), X-ray emitters (Pentericci et al. 2002), and an overdensity of dusty starbursts (Dannerbauer et al. 2014; Rigby et al. 2014), several of which are spectroscopically confirmed to be close to the radio galaxy redshift. Five of the 18 X-ray sources (~28%) detected by Pentericci et al. (2002) are AGN, including the central radio galaxy. From the soft X-ray luminosity function of AGN, Pentericci et al. (2002) estimated how many sources are expected in a given region of the cluster PKS 1138−262, finding that it contains about twice the number of expected AGN. More recently, Pentericci et al. (2013) also found high AGN fractions by studying eight galaxy groups from z ~ 0.5 to z ~ 1.1. They found that the fraction of AGN with Log LH > 42 erg s−1 in galaxies with MR  −20 varies from less than 5% to 22%, with an average value of 6.3%, which is more than double the fraction for massive cluster at similar high redshifts (e.g., Overzier et al. 2005). Martini et al. (2013) estimated that the cluster AGN fraction in a sample of 13 clusters of galaxies at 1  z  1.5 is ~3% for AGN with rest-frame, hard X-ray luminosity greater than LX, H ≥ 1044 erg s−1. Based on these findings, the galaxy protocluster around 7C 1756+6520 seems to be a particularly interesting object.
Citation: Croft et al. 2005
Citation function: Motivation
We assumed a screen geometry and used PYNEB3 (Luridiana et al. 2015) to correct line fluxes for dust extinction via the Hα/Hβ ratio, adopting the O’Donnell (1994) reddening law with RV = 3.1 and a theoretical Hα/Hβ = 2.86. The extinction-corrected emission line luminosities of the H II regions were then computed using the distances reported in Table 1. For every H II region in our catalog, we also estimated the gas-phase metallicity and the gas ionization parameter by using extinction-corrected emission line fluxes. The gas-phase metallicity O/H was calculated using the Pilyugin & Grebel (2016) S calibration (Scal hereafter). This calibration relies on three diagnostic line ratios (i.e., [N II]/Hβ, [S II]/Hβ, and [O III]/Hβ) and provides an empirical calibration against H II regions that have direct constraints on their nebular temperatures and hence their abundances. This calibration is relatively insensitive to changes in gas pressure or ionization parameter, and we adopted it as fiducial approach in this paper (see Kreckel et al. 2019, for a discussion). In addition, for each galaxy we fit the radial metallicity gradient by using an unweighted least-square linear fitting of the trend between 12 + log(O/H) and the deprojected galactocentric radius (see Fig. A.19). The gas ionization parameter represents the ratio between the ionizing photon flux density and the gas hydrogen density. In this paper, we express the ionization parameter as q = U × c = Q(H0)/4πR2n, where c is the speed of light, U is the dimensionless ionization parameter, Q(H0) is the number of hydrogen ionizing photons (E >  13.6 eV) emitted per second, R is the empty (wind-blown) radius of the H II region, and n its hydrogen density. The ionization parameter is ultimately defined by the structure of an H II region (e.g., size, gas density, filling factor) and the properties of its ionizing source. Photoionization models show that it can be extracted via different diagnostic lines (Kewley & Dopita 2002; Dors et al. 2011). In this paper we use the calibration proposed by Diaz et al. (1991) based on the [S III](9069+9532)/[S II](6717+6713) line ratio. As the [S III]λ9532 Å line falls outside the wavelength range covered by MUSE, we assumed that [S III]λ9532 Å = 2.47[S III]λ9069 Å according to default atomic data in PYNEB (Luridiana et al. 2015). It should be noted that for about 3000 H II regions we are not able to estimate the ionization parameter due to lack of detection of the [S III]λ9532 Å emission line.
Citation: Kreckel et al. 2019
Citation function: Background; Background