adsabs/FOCAL huggingface dataset:
A black hole mass this large in Leo I is not expected from extrapolation of any of the standard black hole−host galaxy correlations. Of course, these small systems do not necessarily need to follow the trends seen in normal galaxies, but the black hole mass reported here does stand out. Lützgendorf et al. (2015) explore extrapolations of black hole correlations down to globular cluster scales, and using a velocity dispersion of 12 km s−1, Leo I has a black hole mass a factor of 100 more than the extrapolated trends. On the numerical side, van Wassenhove et al. (2010) consider different scenarios for formation of a black hole in Milky Way satellites and place the likelihood of one of them having a black hole around the size found here to be below 1%, but this result also depends on the initial seed mass (see also Bellovary et al. 2021). Runaway mergers of stellar mass black holes are unlikely to produce such a black hole in such a small galaxy, since the required initial mass function to reach the ratios seen in the models might be more top-heavy than what chemical abundances and star formation history studies suggest. An alternative explanation for the abnormally large central black hole may come from the recent study on Leo I’s star formation history from Ruiz-Lara et al. (2020). The authors identify a period of quenching from z = 1–2 followed by reignition until almost the present day, when ram pressure stripping may have shut it down as it fell into the Milky Way. While the authors speculate that this reignition at intermediate redshifts could be due to a past merger with a smaller dwarf, this could also be consistent with gas accretion and potential active galactic nuclei feedback, lending support to the high MBH values presented here. Amaro-Seoane et al. (2014) also suggest that dwarf systems may in fact have significantly larger black holes compared to the host galaxy–black hole relationships. Having a larger sample of black hole limits measured in dwarf galaxies will be important to explore.
Citation: Bellovary et al. 2021
Citation function: Compare/Contrast
As seen in Fig. 6 (right), the large MOS-pn discrepancy in the Ni/Fe abundance ratio prevents us from deriving a precise measurement. This discrepancy is worrying, but can be explained by imperfections in the cross-calibration of the two instruments. Alternatively, and perhaps more likely, the high energy band around the Ni-K transitions is significanly affected by the instrumental background (as the flux of the cluster emission sharply decreases at high energies). This hard particle background (already mentioned in Sect. 3.1) has a different spectral shape in MOS and pn, which might even vary with time, thus between observations. In particular, an instrumental line (Cu Kα) is known to affect pn at a rest-frame energy of ~8 keV (Mernier et al. 2015). Despite our efforts to carefully estimate the background, that line might interfere with the Ni-K line in several observations, making a proper modelling of the Ni-K line impossible, and hence, boosting the Ni absolute abundance in pn. In this context, it can be instructive to compare our Ni/Fe measurements with those of Suzaku, which has a lower relative hard particle background. Sato et al. (2007b) (A 1060) and Tamura et al. (2009) (Perseus) reported ratios of ~1.3 ± 0.4 and ~1.11 ± 0.19, respectively (after rescaling to the proto-solar values). Although these measurements might be also be affected by further uncertainties (e.g. the choice of the spectral modelling, Sect. 4.3), they appear to be consistent with the Ni/Fe average ratio measured with MOS is this work, favouring our above supposition that MOS is more trustworthy than pn for measuring Ni/Fe. However, in order to be conservative, we prefer to retain the pn value as a possible result and, therefore, we keep large systematic uncertainties for Ni/Fe. We finally note that, unsurprisingly, Ni/Fe cannot be constrained in the cool objects (Fig. 6, left) because the gas temperature is too low to excite Ni-K transitions.
Citation: Mernier et al. 2015
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In recent years, type Ia supernovae (SNe Ia; Amanullah et al. 2010; Suzuki et al. 2012; Betoule et al. 2014; Scolnic et al. 2018) have been widely employed to test cosmic isotropy. Antoniou & Perivolaropoulos (2010) searched for the preferred direction of anisotropy for the Union2 sample by adopting the hemisphere comparison (HC) method (Schwarz & Weinhorst 2007). They found a maximum accelerating expansion rate, which corresponds to a preferred direction of anisotropy. After that, Mariano and Perivolaropoulos (Mariano & Perivolaropoulos 2012) found a possible preferred anisotropic direction at the 2σ level using the Union2 sample, but by employing the dipole fitting (DF) method. Since then, these two methods have been widely used to explore the cosmic anisotropy (Cai & Tuo 2012; Cai et al. 2013; Zhao et al. 2013; Li et al. 2013; Heneka et al. 2014; Bengaly et al. 2015; Andrade et al. 2018; Sun & Wang 2019) by investigating observational data of, for instance, the Union2.1 sample (Yang et al. 2014; Javanmardi et al. 2015; Lin et al. 2016a), the Joint Light-Curve Analysis (JLA) sample (Lin et al. 2016b; Chang et al. 2018a; Wang & Wang 2018), the Pantheon sample (Sun & Wang 2018), gamma-ray bursts (GRBs; Wang & Wang 2014), galaxies (Zhou et al. 2017), as well as gravitational wave and fast radio bursts (Qiang et al. 2019; Cai et al. 2019). Using HC and DF methods, Zhao et al. (2019) studied the cosmic anisotropy via the Pantheon sample. They found that the SDSS sample plays a decisive role in the Pantheon sample. It may imply that the inhomogeneous distribution has a significant effect on the cosmic anisotropy (Chang et al. 2018b). This opinion was also presented by Sun & Wang (2019). Their conclusions show that the effect of redshift on the result is weak and there is a negligible anisotropy when making a redshift tomography. Deng & Wei (2018a) tested the cosmic anisotropy with the Pantheon sample, but by using the following three methods: the HC method, the DF method, and Healpix1 (Górski et al. 2005). They also performed a cross check. There are two preferred directions from the HC method. In adopting the DF method and Healpix, they found no noticeable anisotropy. They also compared the HC method with the DF method by using the JLA sample (Deng & Wei 2018b) and found that the results of these two methods have not always been approximately coincident with each other. In order to better test the cosmic isotropy, the best way would be to add new samples with a relatively homogeneous distribution.
Citation: Lin et al. 2016a
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The broad blue WR bump feature around 4686 Å was searched in all the star-forming regions in the galaxies. A clear detection of the broad blue WR bump was made only in two star-forming regions hosted by SBS 1222+614 as shown in Fig. 8. This detection is in good agreement with Shirazi & Brinchmann (2012), who have also reported the WR features in SBS 1222+614. The broad blue bump consists of a blend of C iii/C iv λ4650, 4658, N iii λ4634, 4640, [Ar iv] λ4711, 4740, and He ii λ4686 emission lines. This detection generally indicates a good number (102–105) of young WR stars in the galaxy (e.g. Kunth & Sargent 1981; Kunth & Schild 1986). The blue bump appears mainly due to the presence of late-type WN (WNL) and early-type WC (WCE) stars (Schaerer & Vacca 1998). The red WR bump around 5808 Å is also expected in the WR galaxies. We also possibly identified the red bump feature in the SBS 1222+614 (# a + b) region as shown in Fig. 9. The red WR bump appears mainly due to the presence of emission lines [N ii] λ5795 and He i λ5875 from the WCE stars. The red bump is rarely detected in WR galaxies as it contains very weak emission line and is expected in high-mettalicity region (Guseva, Izotov & Thuan 2000). Although we could not detect the WR features in the star-forming regions of other galaxies in our sample, the optical SDSS spectrum for the brightest regions of knot #b and #a in IC 3521 and CGCG 0410-023, respectively, shows a detection of WR features (Brinchmann et al. 2008). The SDSS spectrum also shows the detection of the broad WR blue bump having a relatively lower strength in the knot #b of CGCG 038-051. These detections are missed out in this study most likely due to the low SNR in the blue part of the spectra. The starburst ages estimated to be very young (≤6 Myr; see Table 3) for these star-forming regions including those present in SBS 1222+614 are consistent with the detection of the WR features, which appears only during the very early periods of star formation. Although a few other star-forming regions such as the knot #a in IC 3521 and CGCG 038-051 and knot #c in CGCG 041-023 showing a very young starburst of ≤6 Myr are expected to have the WR features, it is not clear from the present data if these regions also host WR stars or not. Some star-forming regions such as knot #c and b in CGCG 038-051 and CGCG 041-023, respectively, show the starburst of ages of ∼10 Myr, indicating that they have probably completed their WR phases. Overall, the detections of the WR features from our own observations and those from the SDSS data from at least one star-forming region in each galaxy in the sample suggest that the sample galaxies are undergoing young massive star formation phase having a significant population of WR stars.
Citation: Schaerer & Vacca 1998
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The full sample for which the circumstellar CO line emission will be modeled consists of the ∼180 C-, M-, and S-type AGB stars analyzed in Schöier & Olofsson (2001), González Delgado et al. (2003), and Ramstedt et al. (2006) together with additional sources presented in Danilovich et al. (2015). In this initial paper, the new interferometric data for the southern M- and C-type stars are presented. Some of the available sample statistics for the full ∼180 star DEATHSTAR sample are shown in Fig. 1. The distance distribution (Fig. 1, middle) is compared with the estimated distribution in the solar neighborhood (Jura 1990; Jura & Kleinmann 1992; Jura et al. 1993). The estimated distribution is derived from 2MASS and ground-based observations (Jura & Kleinmann 1990) and assumes a smooth distribution of 40 C-type stars kpc−2, a scale height of 200 pc, and that there are a third as many S-type as C-type stars. For the full ∼180 star DEATHSTAR sample, the C-type stars from Schöier & Olofsson (2001) are all brighter than K = 2 mag. The M-type sample consists of the non-Mira stars from the General Catalog of Variable Stars (GCVS; Samus’ et al. 2017) with quality flag 3 in the IRAS 12, 25, and 60 μm bands and 60 μm flux ≳3 Jy with the addition of the Mira stars in González Delgado et al. (2003). The S-type sample also consists of stars that have good quality flux measurements in the IRAS 12, 25, and 60 μm bands, that are found in the General Catalog of Galactic S stars, and that are detected in Tc and are hence intrinsic. The completeness of the S-type sample is discussed in Ramstedt et al. (2009) and it is thought to be complete out to 600 pc. Furthermore, stars of all three spectral types are only included in the sample if they are detected in CO radio line emission, which could be reproduced under the assumption of spherical symmetry. Sources that show strongly asymmetric line profiles when observed with single-dish telescopes, or with known CO-detached shells, are hence not included (e.g., R Scl, U Ant, EP Aqr, and π1 Gru). In this paper, we also exclude stars that have previously been observed with ALMA; however, they will be included in the future analysis. The lower panel of Fig. 1 shows that the stellar and wind parameters of the full ∼180 star sample cover the ranges expected for AGB stars. As expected, fewer stars are found at the high end. The sample is biased to mass-losing stars since only stars that are previously detected in CO radio line emission are included. It is also likely that the full range of AGB masses is not covered simply because higher mass stars are rare. It is our assessment that the full ∼180 star DEATHSTAR sample is representative of Galactic mass-losing AGB stars and covers the relevant ranges of wind and stellar parameters to provide the necessary constraints for theoretical models.
Citation: Samus’ et al. 2017
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Finally, we attempt to characterize the reflection component using the full 2–35 keV spectra with a sophisticated model [M4: ${\tt{\rm constant}}$*${\tt{\rm TBabs}}$*(${\tt{\rm simplr}}$*${\tt {\rm kerrbb2}}$+${\tt{\rm kerrconv}}$*(${\tt{\rm ireflect}}$*${\tt{\rm simplc}}$)), to evaluate the impact on the spin measurement and understand the changes of the accretion flow and the interaction between the disc and the corona. This model features a self-consistent treatment of the thermal, Compton scattering and the reflection component: ${\tt {\rm kerrbb2}}$ describes the thermal component and supplies the seed photons for ${\tt{\rm simplr}}$ (a modified version of ${\tt{\rm simpl}}$, Steiner et al. 2011) to generate the Compton component; while a portion of the Compton component will escape to reach an observer, the remains (refer as ${\tt{\rm simplc}}$, Steiner et al. 2011) will strike back to the disc to generate the reflected component. The reflection fraction Rref in ${\tt{\rm ireflect}}$ (Magdziarz & Zdziarski 1995), defined as the ratio of the Compton photons striking back to the disc to that escaping to infinity, is restricted to negative value thereby only the reflected component is returned by ${\tt{\rm ireflect}}$. It is linked to the reflection constant parameter x in ${\tt{\rm simplr}}$ via the relation x = 1 + |Rref| (Gou et al. 2011). We set the elemental abundance to unity and the iron abundance AFe to five times the solar abundance (Bharali et al. 2019; Buisson et al. 2019; Xu et al. 2020). The disc temperature Tin is fixed at the value returned by ${\tt{\rm diskbb}}$ (M1, refer to Gou et al. 2011). The ionization parameter ξ is fixed at 1000 (i.e. log(ξ) = 3, Xu et al. 2020; Buisson et al. 2019), as it is difficult to be constrained. Finally we use ${\tt{\rm kerrconv}}$ (Brenneman & Reynolds 2006) to apply relativistic effects assuming an unbroken emissivity profile with index q = 3. The key parameters in ${\tt {\rm kerrbb2}}$ and ${\tt{\rm kerrconv}}$ are linked together.
Citation: Steiner et al. 2011
Citation function: Uses; Uses
Citation: Steiner et al. 2011
Citation function: Uses; Uses
The stream velocity is proportional to the halo virial velocity, which can be related to the sound speed of gas at the virial temperature, yielding Mb ∼ 1. The density contrast is obtained by assuming pressure equilibrium between hot gas at Tb ∼ Tv, the virial temperature of an NFW halo, and cold gas at Ts ∼ 104K, set by the steep drop in the cooling rate below that temperature (Sutherland & Dopita 1993). If both the halo and the stream are roughly isothermal, then this ratio is constant throughout the halo. In practice, the stream temperature can be as high as ∼3 × 104 K due to photoheating from the UV background (e.g. Goerdt et al. 2010), while the post-shock temperature in the hot halo near Rv may only be ∼0.5Tv (Dekel & Birnboim 2006, P18). Finally, the stream radius is related to its velocity and density through the mass accretion rate along the stream, ${\dot{M}}_{\rm s}\simeq \pi R_{\rm s}^2 \rho _{\rm s}V_{\rm s}$. Cosmological simulations suggest that this is typically a fixed fraction of the total accretion rate on to the halo virial radius (Danovich et al. 2012), which is constrained by cosmology (e.g. Dekel et al. 2013). The final expressions, including uncertainties in model parameters, are
(42)
\begin{eqnarray*}
M_{\rm b}\simeq 0.75-2.25,
\end{eqnarray*}
(43)
\begin{eqnarray*}
\delta \simeq (10-100)\times M_{12}^{2/3} (1+z)_3^{-1},
\end{eqnarray*}
(44)
\begin{eqnarray*}
\frac{R_{\rm s}}{R_{\rm v}} \simeq (0.01-0.1)\times \left(\delta _{75} M_{\rm b,1.5}\right)^{-1/2},
\end{eqnarray*}
where $M_{12}=M_{\rm v}/10^{12}\, \mathrm{M}_\odot$, (1 + $z$)3 = (1 + $z$)/3, δ75 = δ/75, and Mb,1.5 = Mb/1.5. We note that δ ∼ 10 implies that the background gas has a temperature of ${\sim } 3\times 10^5\, {\rm K}$, which would make it thermally unstable unless it is illuminated by a very hard ionizing background (e.g. Efstathiou 1992, though see also Binney, Nipoti & Fraternali 2009, who argue that buoyancy effects may stabilize the halo against thermal instability even without a photoionizing source). This implies that δ ≳ 30 is likely the most physically plausible regime.
Citation: Dekel & Birnboim 2006
Citation function: Uses; Uses
where a is the area of the i, j pixel (constant for a regular grid), μ is the anglebetween the surface normal and the line of sight to the observer, ϕ is the polar angle of a system with the origin at the disk center, F is the total stellar flux, and q and u are normalized Stokes parameters. We note that the relative flux is normalized to the total flux of an unspotted photosphere in our analysis. The center-to-limb variations ofintensity I(μij) and polarization P(μij) were found through trilinear interpolation using the look-up tables from Kostogryz & Berdyugina (2015) and Kostogryz et al. (2016), according to the selected wavelength, surface gravity, and temperature of a star. Kostogryz & Berdyugina (2015) calculated the tables for continuum spectra of FGK stars (Teff = 4500 K–6900 K, log g = 2.0–5.0, and wavelength range 4000–7000 Å) for the Phoenix grid of plane-parallel models, and in Kostogryz et al. (2016) similar calculations were made for a wider range of models (Teff = 4000 K–7000 K and log g = 1.0–5.5) assuming a spherical stellar atmosphere. In addition, we used unpublished data on intensity and polarization variations for both plane-parallel and spherical cooler models with temperatures down to 3000 K, obtained with the same code by Kostogryz (2016, priv. comm.). It should be noted that these data for cooler stars do not include atomic and molecular absorption lines, which can lead to overestimated polarization values, especially for blue wavelengths and for dwarfs with higher surface gravities (depending on the selected wavelength). Evidently, a proper spectrum synthesis code is needed to calculate the intrinsic polarization for these cases. However, the depolarizing effect from absorption lines may not be very significant because when the absorption line suppresses the polarization, the intensity is also reduced. In turn, this will affect the normalized polarization parameters that we calculate less strongly. Additionally, coherent scattering processes, especially in molecular bands, as seen on the Sun (e.g., Berdyugina et al. 2002) probably increase polarization further, which counteracts the effect of the absorption lines.
Citation: Kostogryz et al. (2016)
Citation function: Uses; Uses
There are many methods, beyond those already employed here, that can be used to identify, and probe the masses of, black holes. This includes reverberation mappings of AGNs (e.g. Bahcall, Kozlovsky & Salpeter 1972; Blandford & McKee 1982; Netzer & Peterson 1997), the ‘fundamental plane of black hole activity’ (Merloni et al. 2003; Falcke, Körding & Markoff 2004), spectral modelling of the high-energy X-ray photon coming from the hot accretion discs around IMBHs (Pringle & Rees 1972; Narayan & Yi 1995), high-ionization optical emission lines (Baldwin et al. 1981; Kewley et al. 2001), and high-spatial-resolution observations of maser emission using radio and millimetre/submillimetre interferometry (e.g. Miyoshi et al. 1995; Greenhill et al. 2003; Humphreys et al. 2016; Asada et al. 2017). In addition, the merging of black holes is now quite famously known to produce gravitational radiation during their orbital decay (Abbott et al. 2016). The merging of galaxies containing their own central IMBH is similarly expected to result in the eventual merging of these black holes. The Kamioka Gravitational Wave Detector (KAGRA; Aso et al. 2013) will be a 3-km-long underground interferometer in Japan capable of detecting the gravitational radiation emanating from collisions involving black holes with masses up to 200 M⊙ (Tápai, Zoltán & László 2015). The planned Deci-Hertz Interferometer Gravitational wave Observatory (DECIGO; Kawamura et al. 2011) and the European, Laser Interferometer Space Antenna (LISA) Pathfinder mission15 (Anza et al. 2005; McNamara 2013), with their greater separation of mirrors, will be able to detect longer wavelength gravitational waves and thus better reach into the domain of intermediate-mass and supermassive black hole mergers, the latter of which are currently being searched for via ‘pulsar timing arrays’ (PTAs) (e.g. Hobbs et al. 2010; Kramer & Champion 2013; Shannon et al. 2015). A key constraint to the expected detection threshold of such signals from PTAs – in particular the background of cosmic ripples from the merger of massive black holes (themselves arising from the merger of galaxies) – is the (black hole)-to-(host galaxy/bulge) mass ratio (see equation 4 for spiral galaxies). An additional source of long-wavelength gravitational radiation will arise from the inspiral of compact stellar-mass objects, such as neutron stars and black holes, around these IMBHs (Mapelli et al. 2012). It is reasonable to expect that the densely packed nuclear star clusters, which coexist with low-mass SMBHs (e.g. González Delgado et al. 2008; Seth et al. 2008; Graham & Spitler 2009), will similarly surround many IMBHs. Gravitational radiation and the gravitational tidal disruption of ill-fated stars that venture too close to these black holes (Komossa et al. 2009, Komossa 2013, and references therein; Zhong, Berczik & Spurzem 2015; Stone & Metzger 2016; Lin et al. 2018) are therefore expected from these astrophysical entities. There is, therefore, an array of future observations that could yield further confidence and insight into the realm of IMBHs.
Citation: Tápai, Zoltán & László 2015
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